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1607

1607.01049_arXiv.txt

The annihilation of dark matter particles is expected to yield a broad radiation spectrum via the production of Standard Model particles in astrophysical environments. In particular, electrons and positrons from dark matter annihilation produce synchrotron radiation in the presence of magnetic fields. Galaxy clusters are the most massive collapsed structures in the universe, and are known to host $\sim\mu$G-scale magnetic fields. They are therefore ideal targets to search for, or to constrain the synchrotron signal from dark matter annihilation. In this work we use the expected sensitivities of several planned surveys from the next generation of radio telescopes to predict the constraints on dark matter annihilation models which will be achieved in the case of non-detections of diffuse radio emission from galaxy clusters. Specifically, we consider the Tier 1 survey planned for the Low Frequency Array (LOFAR) at 120 MHz, the EMU survey planned for the Australian Square Kilometre Array Pathfinder (ASKAP) at 1.4 GHz, and planned surveys for APERTIF at 1.4 GHz. We find that, for massive clusters and dark matter masses $\lesssim 100$ GeV, the predicted limits on the annihilation cross section would rule out vanilla thermal relic models for even the shallow LOFAR Tier 1, ASKAP, and APERTIF surveys.

While astronomical observations have revealed much about the amount and macroscopic properties of dark matter in our universe, we still do not know what particle or particles constitute the dark matter, nor what beyond the Standard Model particle physics it points to. A leading class of dark matter particle candidates are weakly interacting massive particles (WIMPs) \citep{1996PhR...267..195J, 2000RPPh...63..793B, 2005PhR...405..279B}. WIMPs may be thermally produced in the early universe with the right relic density to explain the present dark matter density if they have an annihilation cross section of $\langle \sigma v \rangle \sim 2.2 \times 10^{-26}$ \citep[e.g.][]{2012PhRvD..86b3506S}. The pair annihilation of WIMPs to Standard Model particles in dark matter dense structures would then give a broad range of potentially observable signatures; the expected products of WIMP annihilation include gamma-ray photons, high-energy electrons and positrons, and neutrinos. Many recent indirect dark matter searches have focused on potential gamma-ray emission from dark matter annihilation (see, e.g., \citealt{Charles2016,Conrad2015,Porter2011,Feng2010} for reviews), and in particular, some of the strongest limits on dark matter annihilation for supersymmetric dark matter models come from the non-detection of local dwarf spheroidal galaxies with the Fermi Large Area Telescope (Fermi-LAT) \citep{Ackermann2015}. However, dark matter annihilation generically produces emission across the entire electromagnetic spectrum \citep[e.g.,][]{Colafrancesco2006, Profumo:2010ya, Profumo:2013yn}. For example, in the presence of magnetic fields, the high-energy electrons/positrons produced will radiate via synchrotron in radio. These same high-energy particles inverse-Compton (IC) up-scatter background radiation, such as the cosmic microwave background (CMB) or starlight, to X-ray or soft gamma-ray frequencies. Clusters of galaxies represent particularly promising targets for searches for secondary IC and synchrotron emission resulting from dark matter annihilation or decay \citep[e.g.,][]{Colafrancesco2006, 2012MNRAS.421.1215J, Storm2013}. In addition to being dark matter dominated, clusters contain large-scale magnetic fields, and the sheer size of clusters enables them to confine $e^\pm$ produced by dark matter long enough for them to radiate. In \cite{Storm2013}, we demonstrated the potential of radio observations of clusters for dark matter studies: the non-detection of radio emission from nearby galaxy clusters strongly constrains the dark matter annihilation cross-section, and in many cases radio observations place better limits than current gamma-ray cluster data. Some clusters are observed to host Mpc scale diffuse radio synchrotron emission in the form of radio halos or radio relics, but most do not \citep[e.g.][]{2012A&ARv..20...54F, 2014IJMPD..2330007B}. This radio emission is thought to stem from high-energy electrons accelerated in cluster mergers. For example, radio halos are primarily observed in disturbed clusters and may result from the turbulent acceleration of mildly relativistic electrons in cluster mergers \citep[e.g.][]{2010ApJ...721L..82C, 2014IJMPD..2330007B, 2016MNRAS.458.2584B}. A second hypothesis is that radio halos are hadronically generated in the collisions of cosmic ray protons with the intracluster medium. Such possibility is however disfavored by the non-detection of gamma-ray emission from clusters and by the spatial distribution of the radio emission \citep[e.g.][]{2011ApJ...728...53J,2012MNRAS.426..956B, Zandanel2013, 2015MNRAS.448.2495S}. In setting limits on dark matter models, one is free to choose target clusters which lack contaminating radio signals from either central AGN or radio halo/relic emission. A radio signal from dark matter, on the other hand, would appear as diffuse radio emission with a consistent spectrum across targets, and with a brightness scaling with cluster mass rather than cluster dynamical state or other factors. Radio astronomy is entering a new era with several new observatories coming online recently or in the near future, and offering dramatically increased sensitivity, new low frequency capabilities, and improved spatial resolution. The spectrum expected for dark matter annihilation/decay rises sharply at low frequencies, particularly for low dark matter particle masses, making low-frequency radio observations ideal for dark matter searches. At the same time, the Square Kilometer Array pathfinders will provide novel sensitivity at GHz frequencies. In this paper, we investigate the sensitivity of current and near-term radio surveys to dark matter annihilation showing that these have the potential to surpass all current indirect detection limits for a large range of particle masses and annihilation final states. In particular, we focus on planned surveys with the Low Frequency Array (LOFAR) \citep{vanHaarlem2013}, the Australian Square Kilometer Array Pathfinder (ASKAP) \citep{Johnston2008}, and a new Phased Array Feed system, APERTIF, which will be installed on the Westerbork Synthesis Radio Telescope (WSRT) \citep{Verheijen2008}. This paper is organized as follows. In Section~\ref{sec:synch}, we derive the signal from synchrotron emission due to dark matter annihilation and describe the models for the dark matter spatial profile and magnetic fields in clusters we employ here. In Section~\ref{sec:rad}, we briefly summarize the objectives and relevant details of the upcoming surveys for ASKAP, APERTIF, and LOFAR. In Section~\ref{sec:pred}, we dicuss how we estimate upper limits for these surveys. We also present our predictions for constraining dark matter annihilation by non-detections of clusters with these upcoming surveys. Finally, we summarize our findings in Section~\ref{sec:end}. Throughout the paper, we assume a $\Lambda$CDM cosmology with H$_0=70~$km~s$^{-1}$~Mpc$^{-1}$, $\Omega_m=0.27$, and $\Omega_{\Lambda}=0.73$.

\label{sec:end} In this paper, we estimate the synchrotron signal from dark matter annihilation in a galaxy cluster, and make predictions for constraints on dark matter annihilation from the non-detection of radio emission from clusters with upcoming radio surveys. We explore which factors are most important in determining a sample for future indirect detection studies using the non-detection of radio emission from galaxy clusters, including the optimal redshift range, region of interest, frequency of observation, and magnetic field profile. In summary, we find that the synchrotron signal as a function of radius starts to fall off in the $100-300$~kpc range for all redshift ranges and different magnetic field configurations; this is approximately the size of the core radius of the cluster (and the characteristic radius of the dark matter halo). If the upper limit on the minimum detectable integrated flux for a given survey increases linearly with area, as we have assumed in this paper, then this region of interest of $\sim300$~kpc used to determine upper limits will maximize the signal-to-noise. Over this region of interest, the differences between magnetic field profiles are also minimal. We find that increasing the redshift does not have a strong effect on the predicted limits for the annihilation cross section. This is not necessarily surprising, as we use a brightness threshold to derive upper limits on the predicted radio emission. Therefore, there is a larger pool of clusters from which to select a sample for future indirect detection studies in the radio band, than in, e.g., the gamma-ray band, where the best constraints result from samples of local objects. The best constraints on dark matter annihilation from indirect detection to-date are those from the non-detection of gamma-rays from dwarf galaxies with \textit{Fermi} \citep{Ackermann2015}. However, future radio surveys have the potential to surpass those constraints, by an order of magnitude or more for dark matter masses $\lesssim 100$~GeV and for annihilation to muons. For annihilation to other channels, the limits from dwarfs remain competitive with those predicted for the shallow, wide-area ASKAP, APERTIF, and the LOFAR Tier 1 surveys. Deeper planned surveys can potentially yield even more stringent upper limits on radio emission from clusters, and thus tighter constraints on dark matter models. For example, the LOFAR Tier 2 survey is expected to be approximately 5 times deeper than the Tier 1 survey \citep{vanHaarlem2013}, so the upper limits on dark matter annihilation could be, correspondingly, up to 5 times deeper. Further into the future, anticipated surveys from SKA will be deeper than AKSAP/APERTIF and LOFAR by a factor of 3 or more \citep[][e.g.,]{Prandoni2014}. Similarly, pointed observations of well-selected target clusters with the instruments studied here would yield novel indirect constraints on dark matter models. Future radio observations might then reveal the first signal from dark matter annihilation. The difficulty in this case would be to disentangle a dark matter signal from astrophysical emission like radio halos. The fact that the dark matter signal is expected to be similar for clusters over a large redshift range and is not highly magnetic-field dependent at some frequencies gives a method of testing the origin of a potential dark matter signature. \begin{figure} \centering \includegraphics[scale=0.5]{f4.eps} \caption{Comparison between predicted limits with radio surveys and best-fit models to the Galactic Center GeV excess for annihilation to taus. Closed contours are $95\%$ confidence limits from the following studies: magenta: \citet{Abazajian2016}, red: \citet{Daylan2016}, cyan: \citet{Calore2015a}. The green lines represent limits for the LOFAR Tier 1 survey, at $120$~MHz, sensitivity of $70~\mu$Jy per beam with a $25$~arcsec beam. The solid green line corresponds to $z=0.023$ and a ``cool-core'' magnetic field profile for Tier 1. Note that a non-cool-core magnetic field model gives almost identical constraints as shown in Figures~\ref{fig:mxsvpanel}(b,c). The black line represents the constraints from the non-detection of gamma rays in dwarf galaxies, from Figure 8 in \citet{Ackermann2015}. The gray horizontal line indicates the benchmark ``thermal relic'' cross-section \citep{2012PhRvD..86b3506S}.} \label{fig:mxsvGCE} \end{figure}

1607.01049

1607

1607.06175_arXiv.txt

Spatial distributions of the dominant oscillation frequency obtained for four sunspots show a feature shared by all the analysed levels of the solar atmosphere in these sunspots. This feature located in the inner penumbrae indicates that this region has favourable conditions for 2.5--4\,mHz oscillation propagation. This agrees with the fact that the spectral composition of the oscillations at three atmospheric heights (Fe\,\textsc{i}\,6173\,\AA, 1700\,\AA, and He\,\textsc{ii}\,304\,\AA) in this region are similar. There have been previous evidence of particular similarities along height of photospheric magnetic field strength, line-of-sight velocity, and temperature profile in the inner penumbra, where the internal boundary of the Evershed flow is located. The finding of the same dominant oscillation frequency at a range of altitudes from the chromosphere up to the transition region extends the height range, suggesting similarities in physical conditions.

\label{S-Introduction} \par Sunspots have long been an important object for the study of oscillations and waves in the solar atmosphere. Sunspots provide a diverse range of interactions between the solar magnetic field and matter. The umbra, where the vertical magnetic field prevails, shows signatures of downward motion and weak five-minute oscillations of the whole umbra at the photospheric level \citep{Lites, Kobanov90}. In the chromosphere, strong three-minute oscillations dominate in the umbra; these oscillations at first were considered as standing acoustic waves \citep{Lites, Georgakilas}. Later, these waves were shown to be moving upward \citep{Rouppe}; also they were shown not to be a simple continuation of the running penumbral waves (RPW) in the horizontal direction \citep{Kobanov06}. New observations with ever-increasing spatial and temporal resolution capabilities reveal many new facts about the sunspot's fine structure \citep{Jess15, Khomenko, Yuan, Sych14, Sych15}. \par Penumbra is a complex part of a sunspot, whose understanding and modelling are largely complicated by horizontal---or rather a mixture of differently inclined---flows, steep temperature gradient, and rapid change of the magnetic field strength towards the outer boundary. Further complexity arises due to the inhomogeneity in the azimuthal direction: the bright and dark filaments are associated with different physical parameters. The magnetic field inclination has been shown to be larger in the dark filaments forming the so-called uncombed penumbra \citep{Title, Solanki, BellotRubio03}. Recently, with the help of high-resolution instruments, two components of the magnetic field were observed in a sunspot penumbra. Both components have an inclination close to 40\degree\ in the inner penumbra, and in the outer penumbra the dark-filament field grow horizontal, while the bright-filament field only reaches 60\degree\ inclination \citep{Langhans}. \par It is necessary to analyse the oscillations in the penumbra to form a comprehensive picture of the waves propagating in sunspots. A wide range of frequencies is observed in sunspot penumbrae: \citet{Lites, Brisken, Zirin} registered oscillations in intensity and Doppler velocity signals. Different frequencies tend to appear at different regions of the penumbra: typically, the longer the period, the farther it is observed from the sunspot centre. \citet{Sigwarth} noted that this change in the frequency is more pronounced in the chromosphere than in the photosphere. Such a distribution is explained by the increase in the magnetic field inclination closer to the boundaries of a sunspot, and thus the decrease in the cut-off frequency \citep{Reznikova, Reznikovaetal, Kobanov13}. \par In a sunspot there is a phenomenon called running penumbral waves (RPW) --- observed increase in brightness travelling outwards the penumbra \citep{Beckers, Giovanelli, Zirin}. These waves span a range of frequencies from 1 to 4~mHz \citep{Lites}. \citet{Lites, Brisken, Jess} showed that the periods of the waves increase closer to the boundaries of a sunspot. RPWs were found at the photospheric heights as well \citep{Musman,Lohner}. Naturally, the question was raised on the origin of RPWs: two concepts were proposed. The first one is that RPWs are real waves propagating horizontally across the penumbra \citep{Alissandrakis, Tsiropoula, Tziotziou04, Tziotziou06}; and the second concept implies them to be an apparent effect---a result of waves rising to the surface along inclined magnetic tubes, thus appearing first at the inner penumbra and then farther from the sunspot centre \citep{Rouppe, Bogdan, Kobanov06, Bloomfield, Kobanov08,cho2015apj}. The same effect is also deemed to be responsible for umbral flashes; and indeed, recently, researchers tend to consider umbral flashes and RPWs as manifestations of one phenomenon \citep{Madsen}. This second explanation raises a consequent question: what is the origin of the waves responsible for the observed effect? A number of authors come to the conclusion that these waves are slow magnetoacoustic modes resulting from photospheric p-mode oscillations \citep{Bloomfield, Madsen} or a broadband energy deposition process, \textit{e.g.} granulation motions \citep{bot2011apj}. \citet{Jess} studied the influence of the magnetic field inclination on the RPW. They concluded that the increase in the inclination leads to the increase in the dominant periodicity due to the dropping of the cut-off frequency. However, due to the complicated dynamical properties, reliable estimations of the RPW parameters are difficult to carry out. \par The important feature of the oscillations discussed above is their relation to the Evershed flow. \citet{Kobanov04a} found three ranges of oscillations that most likely connect the direct Evershed flow and the inverse---the so-called St.~John's---flow. The 20--35-minute oscillations have the most consistent phase difference between the photospheric and chromospheric heights. \par In our previous works \citep{Kobanov15,Kolobov}, we revealed that the dominant frequency spatial distribution shows a peculiarity in sunspots' penumbra. We constructed plots showing the dominant frequencies averaged in the azimuthal direction as a function of the distance to sunspot's barycentre, and these plots converged in the inner penumbra. This feature corresponds to the 3--4 mHz frequency range, which indicates that five-minute oscillations dominate above the inner penumbra at all the heights from the photosphere to the transition region. One can assume that in this ring-shaped region above the inner penumbra, favourable physical conditions for five-minute wave propagation exist.

% \par Oscillations of different frequencies are observed in sunspots, and within a sunspot they are distributed non-uniformly. First, different frequency bands occupy different regions of a sunspot. High-frequency oscillations tend to be located within the umbra boundaries, while lower frequencies form circles, whose radii increase with the decreasing frequency. Such distributions are believed to be related to the magnetic field configuration, namely, magnetic field inclination angle: the location of high frequencies coincides with that of the vertical magnetic field, while low frequencies are located at the high-inclination outer penumbra \citep{Reznikova, Kobanov13}. Second, the circles of oscillation distributions grow with the height. Such a pattern, again, is deemed to be related to the magnetic field configuration: waves propagating along magnetic field lines inclined from the spot centre approach the spot's outer boundaries at each consecutive height level \citep{Kobanov13}. \begin{figure} \centerline{ \includegraphics[width=10cm]{1.pdf} } \caption{Dominant frequency distributions in sunspots at different height levels, from the photosphere to the corona. The black lines show the inner and outer penumbra boundaries as seen in the 1700\,\AA\ band AIA images} \label{fig:dominant} \end{figure} \par Figure\,\ref{fig:dominant} shows spatial distributions of the dominant frequencies based on FFT power spectra. The distributions are reconstructed from the SDO data. They follow the two aforementioned features of the oscillation behaviour in sunspots. Based on these distributions we plotted profiles of dominant frequencies as functions of distance from a sunspot centre \textit{r} presented in Figure\,\ref{fig:radial}. These profiles were plotted for three height levels---from the photosphere to the transition region---in the four sunspots. \par These plots behave as we expected---high-frequency range in the centre and gradual decrease towards the sunspot boundaries. The interesting feature in all the studied sunspots that caught our attention is the convergence of all the profiles in the inner penumbrae (see Figure\,\ref{fig:radial}). In case of NOAA\,11711 the profiles intersect in this region. Figure\,\ref{fig:circularspec} shows intensity oscillation power spectra azimuthally averaged over the region marked in Figure\,\ref{fig:radial}. The panels for NOAA\,12149 show that in the transition region (He\,\textsc{ii}\,304\,\AA) the highest peaks are shifted to the higher frequencies. This can be explained by the inhomogeneity of the penumbra or by the fact that sunspot has not purely circular shape. There are five-minute oscillations in the narrow region of the penumbra that dominate at all the heights (Figure\,\ref{fig:circularspec}). We consider this region to be a channel transporting five-minute oscillations from the photosphere up through the chromosphere. \begin{figure} \centerline{ \includegraphics[width=12.5cm]{2.pdf} } \caption{Radial distributions of the dominant frequencies at three heights in sunspots as functions of the distance from the barycentre. The vertical dashed lines mark the penumbra boundaries. The region of interest is marked with grey area, which corresponds to the averaging over the circular-shaped domain in the penumbra} \label{fig:radial} \end{figure} \begin{figure} \centerline{ \includegraphics[width=13cm]{3.pdf} } \caption{Spectra azimuthally averaged over the region within the penumbrae marked grey in Figure\,\ref{fig:radial}} \label{fig:circularspec} \end{figure} \par The interesting behaviour in this penumbra region motivated a more detailed study of the distributions of several parameters there, including those found in earlier works by other researchers. \par Figure\,\ref{fig:304mf} shows the azimuthally averaged magnetic field inclination at the He\,\textsc{ii}\,304\,\AA\ line formation level estimated from the dominant frequency distributions. The details of the estimation procedure can be found in \citet{Kobanov15}. The inner penumbra in these distributions is characterized by the steepest inclination angle gradient; the inclination angle there being 60--65\degree. \begin{figure}[t] \centerline{\includegraphics[width=13cm]{4.pdf} } \caption{Estimation of the magnetic field inclination at the He\,\textsc{ii}\,304\,\AA\ line height based on the dominant frequency distributions. The region of interest is marked with grey area} \label{fig:304mf} \end{figure} \par Various signatures indicating peculiar properties of the inner penumbrae of sunspots have been found previously by other authors. Identical field strength was found at three photospheric heights (deep photosphere, log $\tau_{500} = 0$; mid photosphere, log $\tau_{500} = -1.5$; and top of the photosphere, log $\tau_{500} = -3$) by \citet{Borrero}. At about the same distance from the sunspot centre, the magnetic field profiles of the three photospheric levels show the same value, and in the outer penumbra the order of the profiles is reversed (Figure\,\ref{fig:BorreroMF}). \cite{BellotRubio06}, see Figure~3 therein, noted a hump in the inner penumbra in the azimuthally averaged temperatures at all the photospheric heights ($-3 \leq \tau_{500} \leq 0$) The amplitude of the hump was found to decrease with height. The authors suggested that these temperature enhancements could be due to hot penumbral tubes, by which plasma emerges from the sub-photospheric layers. Based on these two works, one can conclude that the unique properties of the inner penumbrae seem to originate in deeper levels than that we study here. \begin{figure} \centerline{\mbox{\hskip1cm} \includegraphics[width=8cm]{5.pdf} } \caption{Magnetic field vertical component averaged over azimuth in sunspot NOAA 10923 in the deep photosphere at the continuum level, in the mid-photosphere, and in the upper photosphere. The vertical lines mark the penumbra boundaries. The region of interest is marked with grey area (courtesy of Juan M. Borrero, see \citep{Borrero})} \label{fig:BorreroMF} \end{figure} \par These tubes are probably related to the photospheric and chromospheric Evershed flow, which peaks at the outer penumbra boundary and sharply terminates at the same region of the inner penumbra % (see Figure~8 in \cite{BellotRubio06}, and Figure~4 in \cite{Georgakilas03}). Also, \citet{BellotRubio06} showed that the microturbulence velocity rapidly drops to zero in the inner penumbra. \par Today, we lack a comprehensive explanation for this phenomenon. Probably, the key for understanding such a behaviour of the distributions is a model that describes penumbra magnetic field as a series of two types of interlocking-comb filaments \citep{Weiss}. The first one having a large inclination angle are located at low heights and dive beneath the photosphere at the outer penumbra boundary. These filaments are associated with the Evershed flow. The second filament type is closer to vertical in orientation. Their magnetic field lines rise high in the atmosphere and either return to the surface far from the spot or form an open field line. \par As follows from the aforesaid, interesting peculiarities are observed in the behaviour of a number of physical parameters in the sunspots' inner penumbra. In this paper, we hope to draw attention to this problem, solving which requires widening the height range of data analysis and modelling. { \footnotesize \textbf{Acknowledgements}. The study was performed with partial support of the Project No.\,16.3.2 of ISTP SB RAS, by the Russian Foundation for Basic Research under grants No. 15-32-20504 mol\_a\_ved and 16-32-00268 mol\_a. We acknowledge the NASA/SDO science team for providing the data. We are grateful to an anonymous referee for the helpful remarks and suggestions. }

1607.06175

1607

1607.06496_arXiv.txt

One of the key science goals for extremely large telescopes (ELTs) is the detailed characterization of already known directly imaged exoplanets. The typical adaptive optics (AO) Nyquist control region for ELTs is $\sim$0.4 arcseconds, placing many already known directly imaged planets outside the DM control region and not allowing any standard wavefront control scheme to remove speckles that would allow higher SNR images/spectra to be acquired. This can be fixed with super-Nyquist wavefront control (SNWFC), using a sine wave phase plate to allow for wavefront control outside the central DM Nyquist region. We demonstrate that SNWFC is feasible through a simple, deterministic, non-coronagraphic, super-Nyquist speckle nulling technique in the adaptive optics laboratory at the National Research Council of Canada. We also present results in simulation of how SNWFC using the self coherent camera (SCC) can be used for high contrast imaging. This technique could be implemented on future high contrast imaging instruments to improve contrast outside the standard central dark hole for higher SNR characterization of exoplanets.

\label{sec:intro} The direct imaging of exoplanets is more sensitive to planets beyond $\sim$5-10 AU. Although direct imaging has seen less planet detections than radial velocity or transit techniques, the past eight years have revealed a handful of directly imaged planets, including multiple planets around HR 8799 \cite{8799_1, 8799_2}, HD 95086 b\cite{95086}, Beta Pic b\cite{beta_pic}, and most recently 51 Eri b\cite{51eri}. Detailed characterization of these existing planets with extremely large telescopes (ELTs) may be a difficult task, since these systems may lie either at the edge of or outside of the typical $\sim$0.4 arcsecond ELT adaptive optics (AO) Nyquist control region when observing in the near infrared\cite{nfiraos}. This region is set by the deformable mirror (DM) actuator pitch projected onto the telescope pupil, and for a square grid DM is a $(N_\text{act})(\lambda/D)\times(N_\text{act})(\lambda/D)$ region around the on-axis PSF, where $N_\text{act}$ is the number of actuators in width across the telescope pupil, $\lambda$ is the wavelength of light, and $D$ is the telescope diameter\cite{snwfc}. Thus, with classical, single conjugate AO (SCAO), uncorrected atmospheric turbulence and and quasi-static speckles will lower the planet signal to noise ratio (SNR). A recent technique has recently been proposed to allow wavefront control outside the Nyquist control region, called super-Nyquist wavefront control (SNWFC)\cite{snwfc}. The main hardware component in this technique requires the use of a super-Nyquist element in an AO system, such as a mild pupil plane diffraction grating with a spacing between lines that is smaller than the DM actuator pitch relative to the pupil size. The pupil plane imprint creates a PSF copy in the focal plane that is outside the DM Nyquist region, allowing wavefront control to work in a similar $(N_\text{act})(\lambda/D)\times(N_\text{act})(\lambda/D)$ control region around this super-Nyquist PSF copy. In this paper, we present the results of a laboratory experiment and simulations for a future experiment to show that it is possible to use SNWFC on an ELT AO system system to directly image already known and new exoplanets. In \S\ref{lab} we describe the laboratory experiment design (\S\ref{lab_design}), simulations of expected lab performance (\S\ref{lab_sims}), and results in the lab (\S\ref{lab_results}). In \S\ref{scc} we describe the setup and results of our simulation using the self coherent camera (SCC)\cite{scc1,scc2}. In \S\ref{conclusion}, we summarize our results and discuss future work.

SNWFC allows already known and new exoplanets to be directly imaged at a higher SNR than is otherwise possible with the current AO design for ELTs. Our main conclusions in testing this technique are as follows: \begin{itemize} \item We have demonstrated in the lab that a deterministic SNWFC speckle nulling scheme increases performance, consistent with simulations. \item We demonstrated in simulation that SNWFC performance improvement is possible using the SCC after only one iteration, which is of particular interest to wavefront control on shorter timescales for ground-based telescopes. \end{itemize} There is still a lot of unexplored parameter space to get a better understanding of how well SCC-based SNWFC would perform on a telescope. Some additional factors in our simulations that are beyond the scope of this initial exploratory paper but worth further research are performance in polychromatic light, using a segmented pupil, using different coronagraphs and/or apodization, dependence on higher super-Nyquist phase plate frequencies, dependence on using an inner and outer scale in chosen wavefront error power law, dependence on a higher zero pad sampling for the off-axis SCC Lyot hole, and additional improvement based on diffraction suppression techniques\cite{efc_l}. Ultimately, the next step before this technique can be tested on-sky is to demonstrate it in the lab. \label{conclusion}

1607.06496

1607

1607.02971.txt

\vspace*{-0.2cm} We compare common star-formation rate (SFR) indicators in the local Universe in the GAMA equatorial fields ($\sim160$ deg$^2$), using ultraviolet (UV) photometry from GALEX, far-infrared (FIR) and sub-millimetre (sub-mm) photometry from H-ATLAS, and H$\alpha$ spectroscopy from the GAMA survey. With a high-quality sample of 745 galaxies (median redshift $\left<z\right>=0.08$), we consider three SFR tracers: UV luminosity corrected for dust attenuation using the UV spectral slope $\beta$ (SFR$_{\rm UV, corr}$), H$\alpha$ line luminosity corrected for dust using the Balmer decrement (BD) (SFR$_{\rm H\alpha, corr}$), and the combination of UV and IR emission (SFR$_{\rm UV + IR}$). We demonstrate that SFR$_{\rm UV, corr}$ can be reconciled with the other two tracers after applying attenuation corrections by calibrating IRX (i.e. the IR to UV luminosity ratio) and attenuation in the H$\alpha$ (derived from BD) against $\beta$. However, $\beta$ on its own is very unlikely to be a reliable attenuation indicator. We find that attenuation correction factors depend on parameters such as stellar mass ($M_*$), $z$ and dust temperature ($T_{\rm dust}$), but not on H$\alpha$ equivalent width (EW) or Sersic index. Due to the large scatter in the IRX vs $\beta$ correlation, when compared to SFR$_{\rm UV + IR}$, the $\beta$-corrected SFR$_{\rm UV, corr}$ exhibits systematic deviations as a function of IRX, BD and $T_{\rm dust}$.

The distribution functions of star-formation rates (SFR) at different cosmic epochs (or more commonly its integrated form, the evolution of the cosmic star-formation rate density) provide fundamental observational tests for theoretical models of galaxy formation and evolution (e.g., Hopkins \& Beacom 2006; Madau \& Dickinson 2014). Observationally, one can use a variety of indicators at different wavelengths to measure the level of star-formation activity in galaxies. The rest-frame ultraviolet (UV) non-ionising stellar continuum luminosity, where newly formed massive stars emit the bulk of their energy, is often used as a direct SFR indicator, especially at high redshift. H$\alpha$ nebular recombination emission line luminosity, which probes the hydrogen-ionising photons produced by the most massive and short-lived stars, is another commonly used SFR indicator when spectroscopy is available. One of the most significant challenges when using UV or H$\alpha$ line emission as a direct star-formation tracer is the effect of dust attenuation as the process of star formation takes place in dense, cold and often dusty molecular gas clouds. Indeed, some of the most intensely star-forming galaxies are extremely UV-faint, e.g., those selected in the sub-millimetre (sub-mm). To overcome this problem, various empirical or semi-empirical correction methods have been developed to determine the amount of dust attenuation in a galaxy. For example, the power-law spectral slope of the rest-frame UV continuum $\beta$ ($f^{\lambda}\propto \lambda^{\beta}$), or its proxy the FUV - NUV colour, has been widely used as a practical method for estimating the global attenuation corrections (e.g., Meurer et al. 1995, 1997, 1999; Burgarella et al. 2005; Laird et al. 2005; Reddy et al. 2006; Salim et al. 2007; Treyer et al. 2007; Wijesinghe et al. 2011). However, the spectral slope $\beta$ shows a wide dispersion with varying dust properties, dust/star geometry and redshift (e.g., Witt \& Gordon 2000; Granato et al. 2000; Oteo et al. 2014). In addition, $\beta$ is sensitive to the intrinsic UV spectral slope (determined by properties such as age of the stellar populations, star-formation history, metallicity, etc.) and as such is dependent on a number of parameters that are not solely related to dust attenuation (e.g., Kong et al. 2004; Buat et al. 2005). Another independent method to correct for dust is to use Balmer decrement ratio measurement (i.e. the observed flux ratio of the H$\alpha$ and H$\beta$ nebular emission lines) to estimate the amount of dust attenuation at H$\alpha$ (e.g., Kennicutt 1992; Brinchmann et al. 2004; Moustakas et al. 2006; Garn \& Best 2010). However, Balmer decrement measurements are generally only available for bright H II regions within the galaxies, and so can be problematic when applying to the whole galaxy. Also, H$\beta$ is considerably weaker than H$\alpha$. The Balmer decrement method is also found to be a poor estimator for dust attenuation in dusty starbursts (e.g., Moustakas, Kennicutt \& Tremonti 2006). From an energy conservation point of view, one can also derive SFR from dust emission as dust absorbs the UV and optical light from newly formed stars and re-emit predominantly in the far-infrared (FIR) and sub-mm. The main advantage of inferring SFR from the IR emission is that it is not affected by dust attenuation. However, some of the IR emission could be caused by heating from the old/evolved stellar populations or AGN and thus is not related to recent star formation (e.g., Helou 1986; Popescu et al. 2000; Bell et al. 2003; Natale et al. 2015). {The aim of this paper is to take advantage of the Galaxy and Mass Assembly (GAMA)\footnote{\url{http://www.gama-survey.org}} survey (Driver et al. 2009, 2011; Liske et al. 2015) and associated multi-wavelength surveys to carefully examine some of the most commonly used SFR indicators. GAMA provides a large sample of galaxies in the local Universe where photometric information in the UV and IR as well as measurements of key spectral lines such as H$\alpha$ and H$\beta$ are available. The paper is organised as follows. In Section 2, we describe the various surveys and derived data products used in our analysis. In Section 3, we give a brief overview of the three SFR indicators investigated in this paper and our selection criteria in the UV bands, optical emission lines, and IR and sub-mm bands. In Section 4, we study in detail the properties of our galaxy samples selected at different wavelengths and construct a joint UV-H$\alpha$-IR sample. Then focusing on the joint sample, we examine the dust attenuations derived using different methods, and the correlations between various SFR indicators as a function of galaxy physical parameters such as stellar mass, redshift, UV continuum slope, Balmer decrement, IRX (i.e. the total IR to UV luminosity ratio), S\'ersic index, H$\alpha$ emission line equivalent width, and dust temperature. Finally, we give conclusions and discussions in Section 5. In an upcoming GAMA paper (Luke et al., in prep.), 12 SFR metrics are examined and calibrated to a mean relation. In this paper, we focus on just three commonly used SFR indicators and inter-compare them using a very high quality galaxy sample. With random statistical errors minimised, we investigate the influence of different systematic errors on these SFR indicators. Throughout the paper, we assume a flat $\Lambda$CDM cosmological model with $\Omega_M=0.3$, $\Omega_{\Lambda}=0.7$, $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Flux densities are corrected for Galactic extinction using the $E(B-V)$ values provided by Schlegel, Finkbeiner \& Davis (1998). We use the AB magnitude scale and the Kroupa (Kroupa \& Weidner 2003) initial mass function (IMF) unless otherwise stated.

In this paper, we compare multi-wavelength star-formation rate (SFR) indicators in the local Universe in the three GAMA equatorial fields. Our analysis uses ultraviolet (UV) photometry from GALEX, far-infrared (FIR) and sub-millimetre (sub-mm) photometry from {\it Herschel} H-ATLAS, and H$\alpha$ spectroscopy from the GAMA redshift survey. To minimise random statistical errors, we construct a very high quality sample of 745 objects (median redshift $\left<z\right>=0.08$). We consider three commonly used SFR indicators: UV continuum luminosity corrected for dust attenuation using the UV spectral slope (SFR$_{\rm UV, corr}$), H$\alpha$ emission line luminosity corrected for dust attenuation using the Balmer decrement (SFR$_{\rm H\alpha, corr}$), and the combination of UV and infrared dust emission (SFR$_{\rm UV + IR}$). We find a good linear correlation between SFR$_{\rm UV, corr}$ and SFR$_{\rm UV + IR}$ but with a $\sim0.3$ dex offset when using the UV spectral slope $\beta$ and the Hao et al. (2011) $A_{\rm FUV}$-$\beta$ relation for deriving the dust attenuation correction. This offset is removed when we replace the Hao et al. relation with our new $A_{\rm FUV}$-$\beta$ relation based on calibrating IRX and the attenuation in H$\alpha$ against $\beta$. The $A_{\rm FUV}$-$\beta$ relation is slightly different depending on whether $\beta_{\rm fit}$ or $\beta_{\rm colour}$ is used and the choice of IR SED library. The difference between the Hao et al. (2011) $A_{\rm FUV}$-$\beta$ relation based on a nearby star-forming sample and the new relation derived in this paper is due to the difference in the galaxy samples. In addition to being at higher redshifts, our galaxy sample corresponds to much lower survey flux limits in the IR and UV and therefore contains many more quiescent star-forming galaxies with redder UV spectral slopes and lower IRX values. We also find a good linear correlation between SFR$_{\rm H\alpha, corr}$ and SFR$_{\rm UV + IR}$. There is a small median offset of around 0.1 dex. But we demonstrate that this offset can be entirely explained by systematic effects in deriving the infrared luminosity $L_{\rm IR}$ and/or other systematic errors in the H$\alpha$-based SFR tracer. Moreover, the correlation between SFR$_{\rm H\alpha, corr}$ and SFR$_{\rm UV + IR}$ has a similar scatter (0.2 dex) as the correlation between SFR$_{\rm UV, corr}$ and SFR$_{\rm UV + IR}$. The ratios between different SFR indicators and the dust attenuation correction factors applied in the UV (using $\beta$) and H$\alpha$ (using the Balmer decrement) are examined as a function of various galaxy physical parameters. The attenuation factor applied in SFR$_{\rm H\alpha, corr}$ which is uniquely determined by Balmer decrement increases with increasing values of IRX and $\beta$. Similarly, the attenuation factor applied in SFR$_{\rm UV, \beta}$ which is uniquely determined by $\beta$ increases with increasing values of Balmer decrement and IRX. These trends are consistent with the broad correlations between Balmer decrement, $\beta$, and IRX seen in Fig. 7. We also find that attenuation correction factors depends on stellar mass, redshift and dust temperature, but not on the H$\alpha$ equivalent width or Sersic index in the SDSS optical bands. After applying corrections for dust attenuation, we find that the SFR$_{\rm UV, corr}$/SFR$_{\rm UV + IR}$ ratio does not depend significantly on stellar mass, redshift, UV spectral slope $\beta$, H$\alpha$ equivalent width, or structural parameters such as S\'ersic index. However, the SFR$_{\rm UV, corr}$/SFR$_{\rm UV + IR}$ ratio does systematically decrease with increasing values of IRX, Balmer decrement, and dust temperature. The dependence on IRX is caused by the large scatter in the IRX vs $\beta$ relation. For objects with high IRX values, the dust attenuation correction factor $A_{\rm FUV}$ based on the IRX - $\beta$ correlation derived for the whole sample will underestimate the true level of attenuation. Also, there is a positive correlation between IRX and Balmer decrement and between IRX and dust temperature which explains the systematic trend in the SFR$_{\rm UV, corr}$/SFR$_{\rm UV + IR}$ ratio as a function of Balmer decrement and dust temperature. In contrast, the SFR$_{\rm H\alpha, corr}$/SFR$_{\rm UV + IR}$ ratio does not show any systematic trend as a function of various physical parameters except Balmer decrement and H$\alpha$ equivalent width, which is most likely caused by the fact that both Balmer decrement and H$\alpha$ equivalent width directly determine the dust-corrected H$\alpha$ line luminosity.

1607.02971

1607

1607.08133_arXiv.txt

{ In this work, we investigated the electroweak vacuum instability during or after inflation. In the inflationary Universe, i.e., de Sitter space, the vacuum field fluctuations $\left< {\delta \phi }^{ 2 } \right>$ enlarge in proportion to the Hubble scale $H^{2}$. Therefore, the large inflationary vacuum fluctuations of the Higgs field $\left< {\delta \phi }^{ 2 } \right>$ are potentially catastrophic to trigger the vacuum transition to the negative-energy Planck-scale vacuum state and cause an immediate collapse of the Universe. However, the vacuum field fluctuations $\left< {\delta \phi }^{ 2 } \right>$, i.e., the vacuum expectation values have an ultraviolet divergence, and therefore a renormalization is necessary to estimate the physical effects of the vacuum transition. Thus, in this paper, we revisit the electroweak vacuum instability from the perspective of quantum field theory (QFT) in curved space-time, and discuss the dynamical behavior of the homogeneous Higgs field $\phi$ determined by the effective potential ${ V }_{\rm eff}\left( \phi \right)$ in curved space-time and the renormalized vacuum fluctuations $\left< {\delta \phi }^{ 2 } \right>_{\rm ren}$ via adiabatic regularization and point-splitting regularization. We simply suppose that the Higgs field only couples the gravity via the non-minimal Higgs-gravity coupling $\xi(\mu)$. In this scenario, the electroweak vacuum stability is inevitably threatened by the dynamical behavior of the homogeneous Higgs field $\phi$, or the formations of AdS domains or bubbles unless the Hubble scale is small enough $H< \Lambda_{I} $. } \begin{document}

\label{sec:intro} The recent measurements of the Higgs boson mass $m_{h}=125.09\ \pm \ 0.21\ ({\rm stat})\ \pm \ 0.11\ ({\rm syst})\ {\rm GeV}$~\cite{Aad:2015zhl,Aad:2013wqa,Chatrchyan:2013mxa,Giardino:2013bma} and the top quark mass $m_{t}=173.34\pm 0.27\ ({\rm stat})\ {\rm GeV}$~\cite{2014arXiv1403.4427A} suggest that the current electroweak vacuum state of the Universe is not stable, and finally cause a catastrophic vacuum decay through quantum tunneling~\cite{Kobzarev:1974cp,Coleman:1977py,Callan:1977pt} although the cosmological timescale for the quantum tunneling decay is longer than the age of the Universe ~\cite{Degrassi:2012ry,Isidori:2001bm,Ellis:2009tp,EliasMiro:2011aa}. In de Sitter space, especially the inflationary Universe, however, the curved background enlarges the vacuum field fluctuations $\left<{ \delta \phi }^{ 2 } \right>$ in proportion to the Hubble scale $H^{2}$. Therefore, if the large inflationary vacuum fluctuations $\left<{ \delta \phi }^{ 2 } \right>$ of the Higgs field overcomes the barrier of the standard model Higgs effective potential $V_{\rm eff}\left( \phi \right)$, it triggers off a catastrophic vacuum transition to the negative Planck-energy true vacuum and cause an immediate collapse of the Universe ~\cite{Espinosa:2007qp,Fairbairn:2014zia,Lebedev:2012sy,Kobakhidze:2013tn,Enqvist:2013kaa,Herranen:2014cua,Kobakhidze:2014xda,Kamada:2014ufa,Enqvist:2014bua,Hook:2014uia,Kearney:2015vba,Espinosa:2015qea,Kohri:2016wof,East:2016anr}. The vacuum field fluctuations $\left<{ \delta \phi }^{ 2 } \right>$, i.e., the vacuum expectation values are formally given by \begin{eqnarray} \left< { \delta \phi^{2} } \right>=\int { { d }^{ 3 }k{ \left| \delta { \phi }_{ k }\left( \eta ,x \right) \right| }^{ 2 } } =\int _{ 0 }^{ \infty }{ \frac { dk }{ k } } { \Delta }_{ \delta \phi }^{2}\left(\eta ,k \right)\label{eq:hfhfhedg} . \end{eqnarray} where $ { \Delta }_{ \delta \phi }^{2}\left(\eta ,k \right)$ is defined as the power spectrum of quantum vacuum fluctuations. As well-known facts in quantum field theory (QFT), the vacuum expectation values $\left<{ \delta \phi }^{ 2 } \right>$ have an ultraviolet divergence (quadratic or logarithmic) and therefore a regularization is necessary. The quadratic divergence corresponds to the normal contribution from the fluctuations of the vacuum in Minkowski space, and it can be eliminated by standard renormalization in flat spacetime. The logarithmic divergence, however, appears as a consequence of the expansion of the Universe, and has the physical contributions to the origin of the primordial perturbations or the backreaction of the inflaton field. We usually eliminate this logarithmic ultraviolet divergence by simply neglecting the modes with $k > aH$. That corresponds to the stochastic Fokker-Planck (FP) equation, which treats the inflationary field fluctuations generally originating from long wave modes, i.e., the IR parts~\cite{Linde:1993xx}. Recent works~\cite{Hook:2014uia,Kearney:2015vba,Espinosa:2015qea,Kohri:2016wof,East:2016anr} for the electroweak vacuum stability during inflation based on the stochastic Fokker-Planck (FP) equation. However, from the viewpoint of QFT, we must treat carefully short wave modes as well as long wave modes~\cite{Parker:2007ni,Agullo:2011qg} and it is necessary to renormalize the vacuum expectation values $\left<{ \delta \phi }^{ 2 } \right>$ in the curved space-time in order to obtain exact physical contributions~\cite{Vilenkin:1982wt,Paz:1988mt}. Thus, in this paper, we revisit the electroweak vacuum instability from the legitimate perspective of QFT in curved space-time. In the first part of this paper, we derive the one-loop Higgs effective potential in curved space-time via the adiabatic expansion method. In the second part, we discuss the renormalized field vacuum fluctuations $\left<{ \delta \phi }^{ 2 } \right>$ in de Sitter space by using adiabatic regularization and point-splitting regularization. In the third part, we investigate the electroweak vacuum instability during or after inflation from the global and homogeneous Higgs field $\phi$, and the renormalized vacuum fluctuations $\left<{ \delta \phi }^{ 2 } \right>_{\rm ren}$ corresponding to the local and inhomogeneous Higgs field fluctuations. The behavior of the homogeneous Higgs field $\phi$ is determined by the effective potential $V_{\rm eff}\left( \phi \right)$ in curved space-time, and then, the excursion of the homogeneous Higgs field $\phi$ to the negative Planck-energy vacuum state can terminate inflation and triggers off a catastrophic collapse of the Universe. The local and inhomogeneous Higgs fluctuations described by $\left<{ \delta \phi }^{ 2 } \right>_{\rm ren}$ generate catastrophic Anti-de Sitter (AdS) domains or bubbles and finally cause a vacuum transition. In this work, we improve our previous work~\cite{Kohri:2016wof}, provide a comprehensive study of the phenomenon and reach new conclusions. In addition, we persist in the zero-temperature field theory, leaving the generalization to the finite-temperature case and discussion of the thermal History of the metastable Universe for a forthcoming work. This paper is organized as follows. In Section~\ref{sec:effective} we derive the Higgs effective potential in curved space-time by using the adiabatic expansion method. In Section~\ref{sec:adiabatic} we discuss the problem of renormalization to the vacuum fluctuations in de Sitter space by using adiabatic regularization. In Section~\ref{sec:point} we consider the renormalized vacuum fluctuations via point-splitting regularization and show that the renormalized expectation values via point-splitting regularization is consistent with the previous results via adiabatic regularization. In Section~\ref{sec:electroweak} we discuss the behavior of the global Higgs field $\phi$ and the vacuum transitions via the renormalized Higgs field vacuum fluctuations $\left<{ \delta \phi }^{ 2 } \right>_{\rm ren}$, and investigate the electroweak vacuum instability during inflation or after inflation. Finally, in Section~\ref{sec:Conclusion} we draw the conclusion of our work.

\label{sec:Conclusion} In this work, we have investigated the electroweak vacuum instability during or after inflation. In the inflationary Universe, i.e., de Sitter space, the vacuum field fluctuations $\left< {\delta \phi }^{ 2 } \right>$ enlarge in proportion to the Hubble scale $H^{2}$. Therefore, the large inflationary vacuum fluctuations of the Higgs field $\left< {\delta \phi }^{ 2 } \right>$ is potentially catastrophic to trigger the vacuum decay to a negative-energy Planck-scale vacuum state and cause an immediate collapse of the Universe. However, the vacuum field fluctuations $\left< {\delta \phi }^{ 2 } \right>$, i.e., the vacuum expectation values have a ultraviolet divergence, which is a well-known fact in quantum field theory, and therefore a renormalization is necessary to estimate the physical effects of the vacuum transition. Thus, we have revisited the electroweak vacuum instability during or after inflation from the legitimate perspective of QFT in curved space-time. We have discussed dynamics of homogeneous Higgs field $\phi$ determined by the effective potential ${ V }_{\rm eff}\left( \phi \right)$ in curved space-time and the renormalized vacuum fluctuations $\left< {\delta \phi }^{ 2 } \right>_{\rm ren}$ by using adiabatic regularization and point-splitting regularization, where we assumed the simple scenario that the Higgs field only couples the gravity via the non-minimal Higgs-gravity coupling $\xi(\mu)$. In this scenario, we conclude that the Hubble scale must be smaller than $H<\Lambda_{I} $, or the Higgs effective potential in curved space-time is stabilized below the Planck scale by a new physics beyond the standard model. Otherwise, our electroweak vacuum is inevitably threatened by the catastrophic behavior of the homogeneous Higgs field $\phi$ or the formations of AdS domains or bubbles during or after inflation.

1607.08133

1607

1607.04895_arXiv.txt

We present a method to identify distant solar system objects in long-term wide-field asteroid survey data, and conduct a search for them in the Pan-STARRS1 (PS1) image data acquired from 2010 to mid-2015. We demonstrate that our method is able to find multi-opposition orbital links, and we present the resulting orbital distributions which consist of $154$ Centaurs, $255$ classical Trans-Neptunian Objects (TNOs), $121$ resonant TNOs, $89$ Scattered Disc Objects (SDOs) and $10$ comets. Our results show more than half of these are new discoveries, including a newly discovered 19th magnitude TNO. Our identified objects do not show clustering in their argument of perihelia, which if present, might support the existence of a large unknown planetary-sized object in the outer solar system.

\subsection{Background and Importance} The minor planets orbiting beyond Neptune provide valuable insight on our solar system's formation and evolution, but they have only been studied since 1992 when the first Trans-Neptunian Object (TNO) after Pluto was discovered \citep{Luu1993}. Almost a quarter century later, $\sim2000$ TNOs and Centaurs are known (see Figure \ref{real}) and they are revealing their properties slowly because of the difficulties involved with detecting the faint, slow-moving members of this distant population. \begin{figure}[p] \centering \includegraphics[width=0.9\textwidth]{real.pdf} \caption{(A) Eccentricity, (B) inclination, and (C) absolute magnitude, versus perihelion distance of known TNOs (including SDOs; orange, magenta, and blue marks), Centaurs (brown), and Inner Oort Cloud objects (IOCs; red dots). Objects previously discovered by PS1 through MOPS and IASC (as mentioned in the text) are depicted as yellow triangles. The limiting absolute magnitude is shown as a dashed line in panel (C) for $V=22.5$.} \label{real} \end{figure} Multiple dedicated TNO surveys have been conducted over the years \citep[e.g.,][]{Larsen2001,Gladman2001, Bernstein2004,Elliot2005,Petit2006,Petit2008,Sheppard2011,Gladman2012,Alexandersen2014,Brown2015} including stellar occultation surveys \citep[e.g.,][]{Schlichting2012} which focussed on the discovery of sub-km objects below the sensitivity limit of optical telescopes. Thanks to these studies, the large $100-1000$~km objects have a well characterized size-frequency distribution \citep[SFD;][] {Petit2008,Fuentes2008} while TNOs smaller than $100$~km have only more recently been studied \citep{Fraser2009,Sheppard2010b,Gladman2012}. However, there is a need for more observational data to confirm the apparent transition from a steep to shallow SFD slope among the Neptune Trojans \citep{Sheppard2010b} and SDOs \citep{Shankman2013} around $D\sim100$~km (corresponding to absolute magnitude $H\sim8.5$) \citep{Alexandersen2014}. If the transition is present within all TNO sub-populations it would suggest the formation scenario in which ``asteroids were born big'' \citep{Morbidelli2009} and imply that objects smaller than $100$~km are dominantly the result of collisional evolution. Most of the known TNOs were discovered in `deep and narrow' observing campaigns using large telescopes with small fields of view. Current Near Earth Object (NEO) surveys \citep{Larson1998, Kaiser2004} have the advantage of continuously monitoring large portions of the sky over several years, but are disadvantaged because they use smaller telescopes with cadences designed to identify NEOs that move more than $10\times$ faster than TNOs. \citet{Brown2015} searched archival data from the Catalina Sky Survey \citep{Larson2003} and Siding Spring Survey \citep{Larson2003} and independently identified the eight brightest known TNOs. Even though they did not discover any new objects they predicted a 32\% chance that an object having magnitude $V<19.1$ remains undiscovered in the unsurveyed region of the sky. Evidence has been mounting in the past few years that there is a large planetary-sized distant object in our solar system whose gravitational perturbations influence the orbits of Scattered-Disc Objects (SDOs), particularly those on orbits similar to the dwarf planet (90377) Sedna \citep[e.g.,][]{Trujillo2014,DeFuMarcos2015,Batygin2016}. These works suggest that all currently known extreme TNOs with semi-major axis greater than $150$~AU (including the only other Sedna-like object: 2012~VP$_{113}$) show a pronounced clustering in their arguments of perihelia ($\omega$) not present in the closer TNO population. \citet{Trujillo2014} suggest this clustering is centred at $\omega\sim0^{\circ}$ and that it is due to the Lidov---Kozai effect, a three-body interaction capable of constraining $\omega$ \citep{Kozai1962}. They propose that a super-Earth mass body located at $\sim250$~AU would be capable of restricting $\omega$ for these objects and be stable for billions of years. However, \citet{Batygin2016} made a similar calculation, but excluded orbits which do not demonstrate long term stability because of Neptune, and found that the distant TNOs cluster around $\omega\sim318^{\circ} \pm 8^{\circ}$ which is inconsistent with the Kozai mechanism. They suggest instead that the clustering can be maintained by a distant ten Earth-mass planet on an eccentric orbit with semi-major axis $700$~AU, nearly co-planar with the distant TNOs, but with $\omega$ shifted by $180^{\circ}$. In addition, such a planet might explain the presence of highly inclined TNOs whose existence has not yet been explained \citep{Gladman2009b}. \citet{Trujillo2014} also state that another plausible explanation for such a peculiar asymmetric $\omega$ configuration would be a strong stellar encounter with the Oort cloud in the past. Increasing the number of known retrograde TNOs and Sedna-like SDOs is needed to further test these hypotheses and to constrain the orbital elements and mass of any potentially undiscovered planet. Towards that end, in this work we report on the discovery and detection of the largest number of TNOs by a single asteroid survey, which due to its long-duration and wide-field coverage, provides an excellent complement to targeted deep-and-narrow surveys, resulting in a relatively unbiased TNO sample. \subsection{Pan-STARRS} The prototype telescope for the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS1, hereafter referred to as PS1) located in the United States on Haleakala, Maui, Hawaii, has been surveying the sky since 2010. Many of the observations by PS1 were taken as a sequence of four exposures, each separated by a Transient Time Interval (TTI) of $\sim20$ minutes. This cadence was selected to optimise detection of Near Earth Objects (NEOs) --- objects which have perihelia $q<1.3$~AU. Observations from each night are rapidly processed by the Image Processing Pipeline \citep{Magnier2006}, and all detected moving objects identified by the Moving Object Processing System \citep[MOPS;][]{Denneau2013} are reported to the Minor Planet Center. PS1 has become the leading discovery telescope for NEOs, discovering almost half of the new Near Earth Asteroids in 2015, and discovering more than half of the new comets in 2015 \citep{Wainscoat2015}. The detection of NEOs is done using subtraction of image pairs which have a TTI of $\sim20$ minutes and are well matched in image quality and telescope pointing. This TTI spacing produces a lower limit on the rate of motion for detection of moving objects, below which moving objects are self-subtracted in their image pairs. The lower limit is typically $\sim0.04^{\circ}$ per day ($=2\arcsec$ in $20$ ~minutes), and is seeing dependent. A substantial number of Centaurs (which we define as having perihelia between Jupiter and Neptune) have been discovered from the pair-subtracted images, but only a few more-distant objects have been reported from PS1 (see Figure \ref{real}), some of which were discovered via the International Astronomical Search Collaboration (IASC\footnote{http://iasc.hsutx.edu/}), an educational outreach program in which images were blinked manually. Now that PS1 has thoroughly surveyed the sky north of $-30^{\circ}$ declination, other methods become viable for object detection that are potentially more sensitive to both fainter and slower moving objects. One method uses subtraction of a high-quality static sky image, derived from the cumulative survey data. The other method uses the historical survey to establish a catalogue of stationary objects, and compares catalogues of new detections in new images to the static sky, to reveal moving objects. Over the course of the PS1 survey, image quality has improved, but the grid structure in the PS1 CCDs requires many dithered images to produce a clean static sky image. And although good images in the \textit{gri} passbands are now available for much of the sky north of $-30^{\circ}$ declination, the coverage in the more sensitive \textit{w} passband is more sparse, because surveying in that band has been more focused on the ecliptic for the purpose of NEO discovery. The PS1 survey has also only recently been extended south to $-49^{\circ}$ declination. For these reasons, we have focussed our initial exploration of methods to extract fainter moving objects on the catalogue based approach.

A search for distant objects was made using the archival PS1 data, with $1420$ objects identified, consisting of $255$ classical TNOs, $121$ resonant TNOs, $89$ SDOs, $154$ Centaurs, and $789$ Jupiter Trojans. Excluding the trojans, $371$ of these are new objects which we could not link to known objects. While our identified objects do not show a clustering in their arguments of perihelia, increasing the number of known retrograde TNOs and Sedna-like SDOs is important in better understanding the distant population in our solar system, especially to constrain the orbital elements and mass of any potentially undiscovered large planetary-sized objects \citep{Trujillo2014,Batygin2016}. Future work will focus on validating the detection efficiency, as well as optimising our detection parameters to work well beyond the classical TNO regime.

1607.04895

1607

1607.08844_arXiv.txt

Typical galaxies emit about one third of their energy in the infrared. The origin of this emission reprocessed starlight absorbed by interstellar dust grains and reradiated as thermal emission in the infrared. In particularly dusty galaxies, such as starburst galaxies, the fraction of energy emitted in the infrared can be as high as 90\%. Dust emission is found to be an excellent tracer of the beginning and end stages of a star's life, where dust is being produced by post-main-sequence stars, subsequently added to the interstellar dust reservoir, and eventually being consumed by star and planet formation. This work reviews the current understanding of the size and properties of this interstellar dust reservoir, by using the Large Magellanic Cloud as an example, and what can be learned about the dust properties and star formation in galaxies from this dust reservoir, using SPICA, building on previous work performed with the \emph{Herschel} and \emph{Spitzer} Space Telescopes, as well as the Infrared Space Observatory.

\subsection{Sources and sinks in the Large Magellanic Cloud} The Large Magellanic Cloud (LMC) was imaged using \emph{Spitzer} \citep[SAGE;][]{Meixner_06_Spitzer} and \emph{Herschel} \citep[HERITAGE;][]{Meixner_13_Herschel} using all non-overlapping photometric bands available on these telescopes. The combined SAGE-HERITAGE survey, often referred to as Mega-SAGE, covers the wavelength range from 3.6 to 500 $\mu$m. The four \emph{Spitzer}-IRAC bands (3.6--8.0 $\mu$m) mainly show circumstellar dust emission from stellar point sources, and extended emission due to polycyclic aromatic hydrocarbons (PAHs), thus forming a good tracer of the production and consumption of dust in post- and pre-main-sequence stars, while the \emph{Spitzer}-MIPS and \emph{Herschel} images show the dust mass in the interstellar medium (ISM). The MIPS 24 $\mu$m band represents a transition, while dominated by extended emission due to interstellar dust, extremely dusty point sources are clearly detectable in this band too. It is these point sources with 24-$\mu$m detections that dominated the dust production and consumption budget. Building on a series of previous estimates and methods to determine dust production rates by evolved stars, \citet{Riebel_12_Mass} estimate a total dust production rate of $\sim 2.1 \times 10^{-5}$ \mdot from the \emph{Spitzer}-IRAC and MIPS data. This rate should be contrasted against the dust mass in the interstellar medium of the LMC, $\sim 10^{6}$ \mdot, derived from the \emph{Herschel} observations \citet{Skibba_12_Spatial}, and the dust consumption in star formation \citep[$\sim 1.9 \times 10^{-3}$;][]{Skibba_12_Spatial}, assuming a dust-to-gas ratio of 1/200. Taking these numbers at face value, it becomes clear that evolved stars do not produce sufficient dust to replenish the interstellar medium, on the residence time scale of dust in the interstellar medium \citep[$\sim$ 2.5 Gyr;][]{Tielens_90_Towards}. Moreover, the much higher dust consumption rate in comparison to the dust production rate even suggests that the dust reservoir in the interstellar medium of the LMC is currently being depleted. This would imply that until fairly recently, the dust production rate was much higher than at the present day, or, alternatively that we are missing a significant source of dust production in this equation. Three alternatives suggest themselves: {\bf i)} The dust production in supernovae has not been considered in the estimates of dust production by evolved stars in the LMC \citep[e.g.][]{Riebel_12_Mass}, due to low number statistics, but estimates suggest that it could be similar to the production by Asymptotic Giant Branch stars \citep[e.g.~][]{Whittet_03_dust}. Indeed, \citet{Dunne_03_Type} argued that supernova remnant Cas A has produced 2-4 \msun\, of dust, although it has since been shown that the bulk of the dust mass detected is part of a foreground cloud \citep{Krause_04_No}. From observations of contemporary supernovae in external galaxies, \citet{Sugerman_06_Massive} showed that a more realistic number of dust production may lie around $10^{-3}-10^{-2}$ \msun, although \emph{Herschel} observations of SN 1987A show that it has already produced 0.5 \msun in the first 25 years \citep{Matsuura_11_Herschel}. {\bf ii)} Dust production in the interstellar medium is efficient, and the source of interstellar dust is not stellar. This has been explored by \citet{Zhukovska_08_Evolution} and by \citet{Jones_05_Dust}, and seems to be a viable option. \citet{Zhukovska_08_Evolution} estimate that the dust production in the interstellar medium in the Solar Neighborhood exceeds the dust production by evolved stars after as little as 1 Gyr. {\bf iii)} Dust production by \emph{extreme} Asymptotic Giant Branch (AGB) stars may be overlooked. Searches for dusty evolved stars rely on classifying point infrared (IRAC, 2MASS) point sources \citep[e.g.][]{Boyer_11_Surveying}, which are subsequently being fitted against a dust shell model grid to determine the integrated dust production rate \cite[e.g.][]{Sargent_11_Mass,Srinivasan_11_mass}. From this procedure it becomes clear that the dust production is dominated by the reddest objects, the so-called extreme AGB stars. Indeed, \citet{Riebel_12_Mass} found that this small number (4\%) of extremely red sources account for 75\% of the total dust production. This immediately raises the question how many even dustier and redder sources are present, and how they can be detected. With the peak of their spectral energy distributions (SEDs) shifting to even longer wavelengths they can no longer be picked up in the near- or mid-infrared, and far-infrared detections are required. At these wavelengths, the point spread function (PSF) becomes very large, and the point sources need to be bright to stand out over the extended emission due to interstellar dust within the beam. At the distance of the LMC, this is generally not the case, and \citet{Boyer_10_Cold} demonstrated that only a small number of known evolved stars in the LMC have counterparts in the \emph{Herschel} data. \subsection{The composition of stellar dust} Including spectroscopy in the analysis of evolved stars enables us to dissect the dust production not only by type of AGB star (carbon-rich or oxygen-rich), but to narrow down the mineralogical components in the freshly produced dust. Although a huge amount of variation exists, and still needs to be further explored, studies like those done by \citet{Sargent_10_Mass} and \citet{Srinivasan_10_mass} modeled representative oxygen-rich and carbon-rich AGB star spectra for the LMC, and determined typical compositions. Combining this with the dust production rate derived from model grid fitting as described above, allows for a determination of the total dust production split out by mineral. It is found that in the LMC, of the dust produced by AGB stars and Red Supergiants, 77 wt.\% is in the form of amorphous carbon, 11\% in the form of silicon carbide (SiC), 12\% in amorphous silicates, and a small fraction of $<1$\% of the freshly produced dust may be in the form of crystalline silicates \citep{Kemper_13_Stellar}. Although most extreme AGB stars turn out to be carbon-rich, the exact composition of dust produced by these stars is hard to determine due to the high optical depth in the circumstellar dust shells. The composition of only the optically thin outer layer can be optimally probed (Speck et al.~\emph{in prep.}), so the SiC/amorphous carbon ratio in the vast majority of the dust mass produced by extreme AGB stars remains unknown, although one may assume that it remains constant throughout the mass losing phase of the star. As stated before, supernovae may be important dust sources too, but so far little is known about the composition of the dust in these environments. One of the few studies that have looked into this is done by \citet{Rho_08_Freshly}, who derived a multi-component composition dominated by amorphous silicates for Cas A in an attempt to fit the main spectral feature at 21 $\mu$m. The interstellar dust of our own Milky Way shows a different composition from what is produced by evolved stars; the majority of the dust is in the form of amorphous silicates, and a smaller fraction in amorphous carbon \citep{Tielens_05_Origin}. Only trace amounts of SiC \citep[$\sim 2.6 - 4.2$\%;][]{Min_07_shape} and crystalline silicates \citep[$<2$\%][]{Kemper_04_Absence} can be present. The LMC ISM is similar in composition to the Galactic ISM, because the ultraviolet extinction curve between the Milky Way and the LMC is virtually identical, including the relative strength of the 2175 \AA\, bump \citep{Fitzpatrick_86_average}, meaning similar ratios between carbon-rich and oxygen-rich dust in both galaxies. The discrepancy between ISM dust and stellar dust composition is consistent with the idea that in both galaxies dust formation in denser parts of the ISM itself may dominate the overall dust production, however, there is a contribution from evolved stars that needs to be considered.

1607.08844

1607

1607.03151_arXiv.txt

We analyse the broad-range shape of the monopole and quadrupole correlation functions of the BOSS Data Release 12 (DR12) CMASS and LOWZ galaxy sample to obtain constraints on the Hubble expansion rate $H(z)$, the angular-diameter distance $D_A(z)$, the normalised growth rate $f(z)\sigma_8(z)$, and the physical matter density $\Omega_mh^2$. We adopt wide and flat priors on all model parameters in order to ensure the results are those of a `single-probe' galaxy clustering analysis. We also marginalise over three nuisance terms that account for potential observational systematics affecting the measured monopole. The Monte Carlo Markov Chain analysis with such wide priors and additional polynomial functions is computationally expensive for advanced theoretical models. We develop a new methodology to speed up by scanning the parameter space using a fast model first and then applying importance sampling using a slower but more accurate model. Our measurements for DR12 galaxy sample, using the range $40h^{-1}$Mpc $<s<180h^{-1}$Mpc, are $\{D_A(z)r_{s,fid}/r_s$, $H(z)r_s/r_{s,fid}$, $f(z)\sigma_8(z)$, $\Omega_m h^2\}$ = $\{956\pm28$ $\rm Mpc$, $75.0\pm4.0$ $\Hunit$, $0.397 \pm 0.073$, $0.143\pm0.017\}$ at $z=0.32$ and $\{1421\pm23$ $\rm Mpc$, $96.7\pm2.7$ $\Hunit$, $0.497 \pm 0.058$, $0.137\pm0.015\}$ at $z=0.59$ where $r_s$ is the comoving sound horizon at the drag epoch and $r_{s,fid}=147.66$ Mpc is the sound scale of the fiducial cosmology used in this study. In addition, we divide each sample (CMASS and LOWZ) into two redshift bins (four bins in total) to increase the sensitivity of redshift evolution. However, we do not find improvements in terms of constraining dark energy model parameters. Combining our measurements with Planck data, we obtain $\Omega_m=0.306\pm0.009$, $H_0=67.9\pm0.7\Hunit$, and $\sigma_8=0.815\pm0.009$ assuming $\Lambda$CDM; $\Omega_k=0.000\pm0.003$ assuming oCDM; $w=-1.01\pm0.06$ assuming $w$CDM; and $w_0=-0.95\pm0.22$ and $w_a=-0.22\pm0.63$ assuming $w_0w_a$CDM. Our results show no tension with the flat $\Lambda$CDM cosmological paradigm. This paper is part of a set that analyses the final galaxy clustering dataset from BOSS.

\label{sec:intro} The cosmic large-scale structure from galaxy redshift surveys provides a powerful probe of the properties of dark energy and the time dependence of any cosmological model in a manner that is highly complementary to measurements of the cosmic microwave background (CMB) (e.g., \citealt{Bennett:2012zja,Ade:2013sjv}), supernovae (SNe) \citep{Riess:1998cb,Perlmutter:1998np}, and weak lensing (see e.g. \citealt{VanWaerbeke:2003uq} for a review). The amount of galaxy redshift surveys has dramatically increased in the last decades. The 2dF Galaxy Redshift Survey (2dFGRS) obtained 221,414 galaxy redshifts at $z<0.3$ \citep{Colless:2001gk,Colless:2003wz}, and the Sloan Digital Sky Survey (SDSS, \citealt{York:2000gk}) collected 930,000 galaxy spectra in the Seventh Data Release (DR7) at $z<0.5$ \citep{Abazajian:2008wr}. WiggleZ collected spectra of 240,000 emission-line galaxies at $0.5<z<1$ over 1000 square degrees \citep{Drinkwater:2009sd, Parkinson:2012vd}, and the Baryon Oscillation Spectroscopic Survey (BOSS, \citealt{Dawson:2012va}) of the SDSS-III \citep{Eisenstein:2011sa} has observed 1.5 million luminous red galaxies (LRGs) at $0.1<z<0.7$ over 10,000 square degrees. The newest BOSS data set has been made publicly available in SDSS Data Release 12 (DR12, \citealt{Alam:2015mbd}). The planned space mission Euclid\footnote{http://sci.esa.int/euclid} will survey over 30 million emission-line galaxies at $0.7<z<2$ over 15,000 deg$^2$ (e.g. \citealt{Laureijs:2011gra}), and the upcoming ground-based experiment DESI\footnote{http://desi.lbl.gov/} (Dark Energy Spectroscopic Instrument) will survey 20 million galaxy redshifts up to $z=1.7$ and 600,000 quasars ($2.2 < z < 3.5$) over 14,000 deg$^2$ \citep{Schlegel:2011zz}. The proposed WFIRST\footnote{http://wfirst.gsfc.nasa.gov/} satellite would map 17 million galaxies in the redshift range $1.3 < z < 2.7$ over 3400 deg$^2$, with a larger area possible with an extended mission \citep{Green:2012mj}. The methodologies of the data analyses of galaxy clustering have also developed along with the growing survey volumes. The observed galaxy data have been analysed, and the cosmological results delivered, using both the power spectrum (see, e.g., \citealt{Tegmark:2003uf,Hutsi:2005qv,Padmanabhan:2006cia,Blake:2006kv,Percival:2007yw,Percival:2009xn,Reid:2009xm,Montesano:2011bp}), and the correlation function (see, e.g., \citealt{Eisenstein:2005su,Okumura:2007br,Cabre:2008sz,Martinez:2008iu,Sanchez:2009jq,Kazin:2009cj,Chuang:2010dv,Samushia:2011cs,Padmanabhan:2012hf,Xu:2012fw,Oka:2013cba,Hemantha:2013sea}). Similar analyses have been also applied to the SDSS-III BOSS \citep{Ahn:2012fh} galaxy sample \citep{Anderson:2012sa,Manera:2012sc,Nuza:2012mw,Reid:2012sw,Samushia:2012iq,Tojeiro:2012rp, Anderson:2013oza, Chuang:2013hya, Sanchez:2013uxa, Kazin:2013rxa,Wang:2014qoa,Anderson:2013zyy,Beutler:2013yhm,Samushia:2013yga,Chuang:2013wga,Sanchez:2013tga,Ross:2013vla,Tojeiro:2014eea,Reid:2014iaa,Alam:2015qta,Gil-Marin:2015nqa,Gil-Marin:2015sqa,Cuesta:2015mqa}. In principle, the Hubble expansion rate $H(z)$, the angular-diameter distance $D_A(z)$, the normalized growth rate $f(z)\sigma_8(z)$, and the physical matter density $\Omega_mh^2$ can be well constrained by analysing the galaxy clustering data alone. \cite{Eisenstein:2005su} demonstrated the feasibility of measuring $\Omega_mh^2$ and an effective distance, $D_V(z)$, from the SDSS DR3 \citep{Abazajian:2004it} LRGs, where $D_V(z)$ corresponds to a combination of $H(z)$ and $D_A(z)$. \cite{Chuang:2011fy} measured $H(z)$ and $D_A(z)$ simultaneously using the galaxy clustering data from the two dimensional two-point correlation function of SDSS DR7 \citep{Abazajian:2008wr} LRGs. The methodology has been commonly known as the application of Alcock-Paczynski effect \citep{Alcock:1979mp} on large-scale structure. The methodology has been improved and also applied to different galaxy samples, e.g., see \cite{Chuang:2012qt,Chuang:2012ad,Reid:2012sw,Blake:2012pj,Xu:2012fw}. Galaxy clustering allows us to differentiate between smooth dark energy and modified gravity as the cause for cosmic acceleration through the simultaneous measurements of the cosmic expansion history $H(z)$ and the growth rate of cosmic large scale structure, $f(z)$ \citep{Guzzo:2008ac,Wang:2007ht,Blake:2012pj}. However, to measure $f(z)$, one must determine the galaxy bias $b$, which requires measuring higher-order statistics of the galaxy clustering (see \citealt{Verde:2001sf}). \cite{Song:2008qt} proposed using the normalized growth rate, $f(z)\sigma_8(z)$, which would avoid the uncertainties from the galaxy bias. \cite{Percival:2008sh} developed a method to measure $f(z)\sigma_8(z)$ and applied it on simulations. \cite{Wang:2012rn} estimated expected statistical constraints on dark energy and modified gravity, including redshift-space distortions and other constraints from galaxy clustering, using a Fisher matrix formalism. $f(z)\sigma_8(z)$ has been measured from observed data in addition to $H(z)$ and $D_A(z)$ (e.g., see \citealt{Samushia:2011cs,Blake:2012pj,Reid:2012sw,Chuang:2013hya,Wang:2014qoa,Anderson:2013zyy,Beutler:2013yhm,Chuang:2013wga,Samushia:2013yga}) determined $f(z)\sigma_8(z)$ from the SDSS DR7 LRGs. \cite{Blake:2012pj} measured $H(z)$, $D_A(z)$, and $f(z)\sigma_8(z)$ from the WiggleZ Dark Energy Survey galaxy sample. Analyses have been performed to measure $H(z)$, $D_A(z)$, and $f(z)\sigma_8(z)$ from the SDSS BOSS galaxy sample \citep{Reid:2012sw,Chuang:2013hya,Wang:2014qoa,Anderson:2013zyy,Beutler:2013yhm,Chuang:2013wga,Samushia:2013yga,Alam:2015qta,Gil-Marin:2015sqa}. In \cite{Chuang:2013wga}, we minimize the potential bias introduced by priors and observational systematics, so that one can safely combine our single-probe measurements with other data sets (i.e. CMB, SNe, etc.) to constrain the cosmological parameters of a given dark energy model. However, due to the large parameter space, the Monte Carlo Markov Chain analysis becomes expensive and makes it difficult to use more advanced/slow models. In this study, we develop a new methodology to speed up the analysis with two steps: 1) generate Marcos chains with a fast model (less accurate); 2) replace/calibrate the likelihoods with a accurate model (slower). For convenience, we use the "Gaussian streaming model" described in \cite{Reid:2011ar}, while there have been more developments, e.g. \cite{Carlson:2012bu,Wang:2013hwa,Taruya:2013my,Vlah:2013lia,White:2014gfa,Taruya:2014faa,Bianchi:2014kba,Vlah:2015sea,Okumura:2015fga}. Although the model we use might not be the most accurate model to date, it is good enough for our purpose and the scale ranges used in this study as we will demonstrate. This paper is organized as follows. In Section \ref{sec:data}, we introduce the SDSS-III/BOSS DR12 galaxy sample and mock catalogues used in our study. In Section \ref{sec:method}, we describe the details of the methodology that constrains cosmological parameters from our galaxy clustering analysis. In Section \ref{sec:results}, we present our single-probe cosmological measurements. In Section \ref{sec:application}, given some simple dark energy models, we present the cosmological constraints from our measurements and the combination with other data sets. We compare our results with other studies in \ref{sec:compare}. We summarize and conclude in Section \ref{sec:conclusion}.

\label{sec:conclusion} We present measurements of the anisotropic galaxy clustering from the CMASS and LOWZ samples of the final date release (DR12) of the SDSS-III Baryon Oscillation Spectroscopic Survey (BOSS) and obtain constraints on the Hubble expansion rate $H(z)$, the angular-diameter distance $D_A(z)$, the normalized growth rate $f(z)\sigma_8(z)$, and the physical matter density $\Omega_mh^2$. We analyse the broad-range shape of quasi-linear scales of the monopole and quadrupole correlation functions to obtain cosmological constraints at different redshift bins. In addition to the two redshift bins, i.e. LOWZ ($z_\textrm{LOWZ}=0.32$) and CMASS ($z_\textrm{CMASS}=0.59$), we split each galaxy sample into 2 bins (for a total of 4 redshift bins) and obtain the measurements at $z=\{0.24$, $0.37$, $0.49$, $0.64\}$ to increase the sensitivity of redshift evolution. However, we do not find improvement in terms of constraining different dark energy model parameters. It might indicate that the dark energy component is stable in the redshift range considered. We adopt wide and flat priors on all model parameters in order to ensure the results are those of a `single-probe' galaxy clustering analysis. We also marginalize over three nuisance terms that account for potential observational systematics affecting the measured monopole. The Monte Carlo Markov Chain analysis with such wide priors and additional polynomial functions is computationally expensive for advanced theoretical models. We have developed and validated a new methodology to speed this up by scanning the parameter space using a fast model first and then applying importance sampling using a slower but more accurate model. Our measurements for DR12 galaxy sample, using the range $40h^{-1}$Mpc $<s<180h^{-1}$Mpc, are $\{D_A(z)r_{s,fid}/r_s$, $H(z)r_s/r_{s,fid}$, $f(z)\sigma_8(z)$, $\Omega_m h^2\}$ = $\{956\pm28$ $\rm Mpc$, $75.0\pm4.0$ $\Hunit$, $0.397 \pm 0.073$, $0.143\pm0.017\}$ at $z=0.32$ and $\{1421\pm23$ $\rm Mpc$, $96.7\pm2.7$ $\Hunit$, $0.497 \pm 0.058$, $0.137\pm0.015\}$ at $z=0.59$ where $r_s$ is the comoving sound horizon at the drag epoch and $r_{s,fid}=147.66$ Mpc is the sound scale of the fiducial cosmology used in this study. Combining our measurements with Planck data, we obtain $\Omega_m=0.306\pm0.009$, $H_0=67.9\pm0.7\Hunit$, and $\sigma_8=0.815\pm0.009$ assuming $\Lambda$CDM; $\Omega_k=0.000\pm0.003$ assuming oCDM; $w=-1.01\pm0.06$ assuming $w$CDM; and $w_0=-0.95\pm0.22$ and $w_a=-0.22\pm0.63$ assuming $w_0w_a$CDM. The results show no tension with the flat $\Lambda$CDM cosmological paradigm.

1607.03151

1607

1607.03684_arXiv.txt

We have detected excess absorption in the emission cores of \cahk during transits of HD 189733b for the first time. Using observations of three transits we investigate the origin of the absorption, which is also seen in \ha and the \ionjb{Na}{1} D lines. Applying differential spectrophotometry methods to the \ionjb{Ca}{2} H and \ionjb{Ca}{2} K lines combined, using respective passband widths of \hbox{$\Delta\lambda =$} \hbox{$0.4\ \&\ 0.6$ \AA}\ yields excess absorption of ${ t_d =}\ 0.0074$\ $\pm$\ $0.0044$ ($1.7\sigma$; Transit 1) and \hbox{$0.0214$\ $\pm$\ $0.0022$} ($9.8\sigma$; Transit 2). Similarly, we detect excess \ha absorption in a passband of width $\Delta\lambda = 0.7$ \AA, with \hbox{${ t_d =\,0.0084\,\pm\,0.0016}$} ($ 5.2\sigma$) and \hbox{${ 0.0121\,\pm\,0.0012}$} ($ 9.9\sigma$). For both lines, Transit 2 is thus significantly deeper. Combining all three transits for the \ionjb{Na}{1} D lines yields excess absorption of \hbox{${ t_d =\,0.0041\,\pm\,0.0006}$} ($ 6.5\sigma$). By considering the time series observations of each line, we find that the excess apparent absorption is best recovered in the stellar reference frame. These findings lead us to postulate that the main contribution to the excess transit absorption in the differential light curves arises because the normalising continuum bands form in the photosphere, whereas the line cores contain a chromospheric component. We can not rule out that part of the excess absorption signature arises from the planetary atmosphere, but we present evidence which casts doubt on recent claims to have detected wind motions in the planet's atmosphere in these data.

\label{section:intro} Appropriately, the first known transiting exoplanet, \hbox{HD 209458}b \citep{charbonneau00transit,henry00hd209458}, also yielded the first detection of an atmosphere outside the solar system. The excess absorption of $(2.32 \pm 0.57) \times 10^{-4}$ in the 5893 \AA\ Na resonance doublet compared with neighbouring continuum passbands suggested that Na was present at a lower level than predicted \citep{seager00transmission}. Clouds, hazes and photo-ionisation of Na were investigated by \cite{seager03hd209458}, \cite{fortney03hd209458} and \cite{barman07} as possible causes for the low observed Na absorption. A complete optical transmission spectrum of \hdtwo\ spanning 3000\,-\,7000 \AA\ was observed for the first time with the Hubble Space Telescope (HST), additionally revealing strong Rayleigh scattering \citep{etangs08rayleigh} at wavelengths of 3000\,-\,5000 \AA. With a later spectral type of K1V\,-\,K2V, \hd causes lower incident radiation of the orbiting hot Jupiter despite its smaller orbital radius. HST transmission spectroscopy of \hdone\ have revealed that Na is observed as a weak feature \citep{pont08hd189733,pont13hd189733} while strong Rayleigh scattering dominates at shorter optical and UV wavelengths \citep{pont08hd189733,sing11hd189733,pont13hd189733}. \begin{table*} \begin{center} \caption{HARPS observations during transit of \hbox{HD 189733b}. Columns 5-9 list S/N ratios for the \cahk, the mean of the nearby Mount Wilson S index V \& R continuum passbands, \ha and the mean of the \ha A \& B continuum passbands. \label{table1}} \begin{tabular}{cccccccccc} \hline & & & & & & & Passband S/N ratios & & \\ \cline{5-9} \vspace{-2mm} \\ & Date & Seeing (\arcsec) & Exptime (s) & Airmass & \ionjb{Ca}{2} K & \ionjb{Ca}{2} H & $\overline{V + R}$ & \ha & $\overline{A + B}$ \\ \vspace{-2mm} \\ \transiti & 2006/09/07 & $0.85 \pm 0.22$ & 600 - 900 & 2.1 - 1.6 &21.4 & 28.0 & 28.3 & 91.0 & 162.4 \\ \transitii & 2007/07/19 & $0.68 \pm 0.11$ & 300 & 2.4 - 1.6 &14.6 & 18.6 & 19.8 & 57.8 & 106.1 \\ \transitiii\ & 2007/08/28 & $1.17 \pm 0.51$ & 300 & 2.3 - 1.6 &12.5 & 16.5 & 16.9 & 52.8 & 94.7 \\ \hline \end{tabular} \end{center} \end{table*} Ground-based transmission spectroscopy of transiting exoplanets is notoriously difficult owing to the small signal and variable nature of the Earth's atmosphere. When spectroscopy is employed, variable seeing changes the slit and hence grating illumination, leading to systematic changes in the recorded spectrum. These are generally second order effects, but are crucial when searching for precision radial velocity signatures or subtle changes in the spectrum at levels of order a few per cent or less. Spectrophotometry has proven to be one successful method enabling multi-wavelength transit radius measurements in the atmospheres of close orbiting planets. For example, \cite{sing12xo2b} used a wide 5\arcsec\ long slit to carry out differential spectrophotometry of the planet hosting star, XO-2A, and its planet, XO-2b, yielding a 5.2-$\sigma$ detection of Na. Extending ground-based characterisation to higher resolution is clearly desirable if we wish to detect individual transitions and measure abundances in detail. \'{E}chelle spectroscopy offers the wavelength range needed, but with many optical components, including a grating and a cross-disperser, this presents greater challenges, especially when narrow slit widths of typically $\lesssim 1$\arcsec\ are commonly employed. Fibre fed instruments offer the best internal stability since the light is largely scrambled by the fibre and hence the illumination of the slit on exit from the fibre is less sensitive to seeing changes and guiding errors. By normalising each observation to nearby continuum regions, excess absorption in individual lines can be probed. This was first applied specifically to \hdone\ by \cite{redfield08} to detect excess Na absorption during transit and more recently using data from the High Accuracy Radial Velocity Planet Searcher (\harps) by \cite{wyttenbach15}. \cite{louden15hd189733} have also applied a model that accounts for the differences in the Doppler shift of absorption from opposite sides of the planet. The planetary absorption line profiles were thus modelled as a function of time during the transit to infer an equatorial eastward jet in the planet's atmosphere. \cite{fossati13wasp12} hypothesised that extrinsic absorbing gas local to the WASP-12 system and arising from WASP-12b is responsible for the anomalously low flux in the \cahk cores of WASP-12. If gas is uniformly distributed around the star, we would expect the absorption to be constant with time. \cite{fossati10wasp12metals} and \cite{haswell12} showed that there is more near ultraviolet absorption around the phases of transit, so the absorbing gas in the WASP-12 system is densest there. Because WASP-12 is faint, and there is very little flux in the line cores \citep{fossati13wasp12}, we decided to study \hdone to search for time variable \cahk absorption. \hd is a K1.5V star with a relatively small radius of \hbox{$R_* =$} \hbox{$0.756 \pm 0.018$ \rsun} \citep{torres08}. Hence a large transit signature for a given planet radius is expected compared with a star such as WASP-12b with an estimated radius of \hbox{$R_* =$} \hbox{1.6 \rsun}. Conversely, the pressure scale height of $H_{eq} = 201$ km is not particularly high compared with \hdtwo, with $H_{eq} = 553$ km \citep{sing16nature}. In this paper, we examine \cahk and \ha absorption during the transit of \hdone. We also re-examine \nad absorption following a recent analysis by \cite{wyttenbach15} and also use the \ionjb{Ca}{1} 6122 \AA\ line as a control. The archival \harps\ data exhibit clear evidence of excess absorption during transit, thanks largely to the fact that the star is active and hence exhibits \cahk cores with sufficient flux to reliably detect the transit. In \S \ref{section:analysis_hd189733}, we discuss the activity and chromospheric variability of \hbox{HD 189733}. The observations and the methods used to derive the transit light curves is presented \S \ref{section:observation}. We present the transit light curves in \S \ref{section:results_lightcurves} and the residual line profiles during transit in \S \ref{section:results_transmission}. We discuss the possible origin of the excess absorption during transit in light of our finding that the absorption is both variable and more sharply defined in the stellar reference frame than the planet reference frame \S \ref{section:discussion}. A summary and concluding remarks are found in \S \ref{section:conclusion}.

\protect\label{section:conclusion} We have detected excess absorption in the \cahk lines of \hd during transit of its close orbiting giant exoplanet. We also detected excess absorption in \ha, and re-examined previous results for \nad. The transit depths are significantly variable { in \cahk and \ha, with \transitii exhibiting the strongest line absorption. We} studied the time series spectra of each line during the three transits and find evidence that the excess absorption is located in the stellar reference frame. We postulate that the chromosphere of \hd is the dominant contributor to the excess transit absorption and responsible for its variable nature { since the flux measured in the cores of the lines we have studied contains a significant chromospheric contribution. A systematic affect is introduced because for any given transit, the chromospheric emission will be slightly different, and localised to specific regions on the stellar surface, whereas the continuum emission used for normalisation originates in the stellar photosphere. Temporal variability (from one transit to the next) in chromospheric emission will lead to changes in the depth of the apparent absorption, while localised systematic effects, where the planet transits active regions, will also result in variability. The \nad lines core are less sensitive to chromospheric variability than \cahk and \ha and may contain systematics from both chromospheric features such as plage and filaments, and photospheric variability from starspots.} Observations aimed at detecting { specific lines in} the atmospheres of transiting exoplanets with ground-based observations thus likely yield systematically biased transit signatures { if the lines being probed are sensitive to stellar chromospheric emission}. Space-based observations can, in principle, measure the flux time series for all observed wavelengths so that (singly) differential transmission spectroscopy reveals the difference between IT and OOT spectra. In practice however, often the OOT { spectra} are used to remove systematic effects, or detrending is necessary \citep{pont07hd189733} so even space-based transmission spectroscopy may be affected by this systematic bias when observing chromospherically sensitive regions \citep{pont13hd189733}. It is worth noting also that { since active regions vary in time, the issue is likely to be} more important for planet hosting stars with higher activity levels, such as \hd. Our analysis shows that multiple transits and careful scrutiny of time series spectra are desirable to assess the nature of excess absorption during exoplanetary transits. It would be interesting to obtain further transit observations of both \hd and \hbox{HD 209458} in combination with long term activity monitoring. Study of \cahk is both difficult and risky because it requires that the star be active, with sufficient line reversal flux to obtain reliable transit light curves. \hd is the most favourable target for study of \cahk transit effects; even \hbox{HD 209458}, being much less active, does not possess significant \cahk line core flux. The \nad and \ha lines occur in regions with better continuum flux and thus offer the potential to investigate transit and stellar activity effects on a wider sample of systems.

1607.03684

1607

1607.03367_arXiv.txt

We report on a dark matter search for a Weakly Interacting Massive Particle (WIMP) in the mass range $m_{\chi} \in [4, 30]\,\gev$ with the EDELWEISS-III experiment. A 2D profile likelihood analysis is performed on data from eight selected detectors with the lowest energy thresholds leading to a combined fiducial exposure of 496\,kg-days. External backgrounds from $\gamma$- and $\beta$-radiation, recoils from $^{206}\mathrm{Pb}$ and neutrons as well as detector intrinsic backgrounds were modelled from data outside the region of interest and constrained in the analysis. The basic data selection and most of the background models are the same as those used in a previously published analysis based on Boosted Decision Trees (BDT)~\cite{Armengaud:2016cvl}. For the likelihood approach applied in the analysis presented here, a larger signal efficiency and a subtraction of the expected background lead to a higher sensitivity, especially for the lowest WIMP masses probed. No statistically significant signal was found and upper limits on the spin-independent WIMP-nucleon scattering cross section can be set with a hypothesis test based on the profile likelihood test statistics. The 90\% C.L. exclusion limit set for WIMPs with $m_\chi = 4\,\gev$ is $1.6 \times 10^{-39}\,\mathrm{cm^2}$, which is an improvement of a factor of seven with respect to the BDT-based analysis. For WIMP masses above $15\,\gev$ the exclusion limits found with both analyses are in good agreement.

\label{sec:introduction} Through different astrophysical observations on a wide range of cosmological scales, it is well established that $\sim 27\%$ of the energy density in the Universe is made up of an unknown dark matter~\cite{Adam:2015rua}. A well-motivated class of particles proposed to solve the dark matter problem are \textit{Weakly Interacting Massive Particles} (WIMPs) with masses of the order of $\mathrm{GeV}/c^2$ to $\mathrm{TeV}/c^2$ and an extremely low scattering cross section with ordinary matter. Direct detection experiments search for the elastic scattering of a WIMP from the galactic dark matter halo in detectors on Earth-based experiments. The nuclear recoils from such interactions would have an exponentially falling energy spectrum up to a few keV, depending on the mass $m_\chi$ of WIMPs. In addition, the expected rate is smaller than one interaction per kg of target material per year. To minimize the background for this rare event search, the EDELWEISS experiment is located in the Modane Underground Laboratory (LSM) in the French-Italian Alps, where a rock overburden of $4800\,\mathrm{m\ w.e.}$ reduces the cosmic muon flux down to $5\,\rm{muons/m^2/day}$. Remaining muons are tagged with an active muon veto system surrounding the experiment~\cite{Schmidt:2013gdc}, followed by $50\,\mathrm{cm}$ of polyethylene and $20\,\mathrm{cm}$ of lead to suppress neutrons and gammas. Inside these layers of shielding a cryostat made of ultra-pure copper houses germanium monocrystals which are cooled down to $18\,\mathrm{mK}$. A simultaneous measurement of the heat and ionization energies produced in a recoil allows to discriminate between the dominant \textit{electron recoils} (ER) from radioactivity and \textit{nuclear recoils} (NR), which at low energies are only caused by neutrons and the expected WIMP signal.\newline Other direct detection dark matter experiments use similar approaches based on the same principle to discriminate between different backgrounds and a possible signal from WIMPs. Exclusion limits on the WIMP-nucleon spin-independent scattering cross section from LUX~\cite{Akerib:2015rjg} and SuperCDMS~\cite{Agnese:2014aze} are in strong tension with favoured parameter regions based on observations by DAMA/LIBRA~\cite{Bernabei:2013xsa}, CoGeNT~\cite{Aalseth:2014eft} and CDMSII-Si~\cite{Agnese:2013rvf}.\newline Almost all existing signal claims for low-mass WIMPs can be excluded at 90\% C.L. with the improved limits that were recently published by the EDELWEISS-III collaboration~\cite{Armengaud:2016cvl} considering standard assumptions about the WIMP-nucleus interaction and the galactic halo model. Data from a 10-month WIMP search run were analysed in terms of low-mass WIMPs with masses $m_\chi \in [4, 30] \, \gev$ using a method based on Boosted Decision Trees (BDT). No statistically significant excess of events was observed for eight selected detectors, resulting in exclusion limits up to a factor 40 stronger at $m_\chi = 7 \, \gev$, compared to results from previous EDELWEISS-II~\cite{Armengaud:2012pfa} low-energy data. Such a cut-based analysis performs well when the separation of signal and background is sufficient, as is the case for higher WIMP masses. However, at low energy, the finite resolutions of the detectors cause the electron and nuclear recoils to have overlapping populations in the distributions of the variables that serve as discriminator. A separaration thus requires a cut at lower energy, resulting in a severely reduced efficiency. To overcome this problem, the analysis presented here uses an alternative approach which is based on the maximum likelihood, similar to e.g.~\cite{Akerib:2015rjg, Aprile:2011hx}. It is an unblind analysis performed on a similar data sample that was recorded with the same detectors as in~\cite{Armengaud:2016cvl}. With its completely different analysis approach it improves the sensitivity for low-mass WIMPs and allows to cross-check the results of the BDT-based analysis. Instead of extracting limits without background subtraction from a smaller signal region with optimized signal-to-background ratio, the maximum likelihood method is used to model and fit the data in the entire region of interest (RoI). Thus, the remaining WIMP signal after detector efficiency corrections is not further reduced, while expected backgrounds are fitted and can be subtracted. The systematic uncertainties of the background predictions are taken into account by constraints in the likelihood fit and the calculation of exclusion limits.\newline The operating principle of the EDELWEISS-III detectors and the selection criteria for the analysed data are detailed in Sec.~\ref{sec:detectors}, while a description of the different background components is presented in Sec.~\ref{sec:models}. The formalism of the likelihood model for the analysis is explained in Sec.~\ref{sec:likelihood}, both for fitting the data to individual detectors, as well as for a combined fit of a common signal to all detectors. We also detail how the exclusion limit is set using a hypothesis test based on the profile likelihood test statistics. A discussion of the fit results and a comparison with the result achieved with the BDT method follows in Sec.~\ref{sec:results}. \begin{figure} \includegraphics[width=1\linewidth]{figures/figure1.png} \caption{WIMP search data in the RoI accumulated in eight selected detectors with a fiducial exposure of $496\,\textrm{kg-days}$ in ionization vs. heat energy (black markers). Events before the fiducial cut and in the extended energy range are shown as gray points. Coloured lines indicate the detector-averaged positions that are expected for different background components depending on their ionization yields and collection voltage biases (see text). From top to bottom: electron recoils from tritium decay as well as Compton and cosmogenic gammas in the fiducial volume (blue), surface gammas (dashed blue), nuclear recoils from neutron scattering (magenta), surface betas (dashed green) and $^{206}\text{Pb}$-recoils (dashed brown). Heat-only events have only noise on the ionization channels and no ionization signal on average (red). The coloured contour indicates an $m_\chi = 10\,\gev$ WIMP signal.} \label{fig:alldata} \end{figure}

\label{sec:conclusion} We have presented a search for low-mass WIMPs with the EDELWEISS-III experiment, using eight selected detectors and data taken with a total fiducial exposure of 496\,kg-days after all cuts. A data-driven approach was used to model relevant backgrounds from sideband data. For each detector a likelihood function describing the data in heat and fiducial ionization energies was constructed, with constraint terms for each of the nuisance parameters taking into account systematic uncertainties. No statistically significant signal was found, neither for the fit of data from single detectors, nor for a combined fit over all detectors with a common signal cross section. Exclusion limits were set with a hypothesis test using a profile likelihood based test statistics, including corrections for under-fluctuations of the background. At 90\%\,C.L. limit we exclude spin-independent WIMP-nucleon scattering cross sections of $\sigma = 1.6 \times 10^{-39}\,\textrm{cm}^2$ ($6.9 \times 10^{-44}\,\textrm{cm}^2$) for a WIMP mass of $m_\chi = 4\,\textrm{GeV}$ ($m_\chi = 30\,\textrm{GeV}$). Thanks to the higher signal efficiency and a subtraction of the expected backgrounds, the likelihood analysis shows an improvement of a factor of $\sim 7$ for $4\,\gev$ WIMPs compared to a BDT based analysis while reproducing the limit at $15\,\gev$ and above. The results and achieved sensitivity underline the power of a maximum likelihood analysis based on detailed background models.

1607.03367

1607

1607.04292_arXiv.txt

We present a new set of optical polarization plane rotations in blazars, observed during the third year of operation of {\em RoboPol}. The entire set of rotation events discovered during three years of observations is analysed with the aim of determining whether these events are inherent in all blazars. It is found that the frequency of the polarization plane rotations varies widely among blazars. This variation cannot be explained either by a difference in the relativistic boosting or by selection effects caused by a difference in the average fractional polarization. We conclude that the rotations are characteristic of a subset of blazars and that they occur as a consequence of their intrinsic properties.

\label{sec:introduction} Blazars are active galactic nuclei with relativistic jets oriented toward the observer. Relativistic boosting causes synchrotron radiation from the jet to dominate the blazar spectra at low frequencies \citep{Blandford1979}. Consequently, the optical emission of blazars often has high and variable polarization. Commonly, the polarization fraction and the electric vector position angle (EVPA) in the optical band show irregular variations \citep[e.g.][]{Brindle1985}. However, a number of events have been detected in which the EVPA traces continuous, smooth rotations that in some cases occur contemporaneously with flares in the total broadband emission \citep{Marscher2008}. It has been suggested that at least some large amplitude EVPA swings can be physically associated with gamma-ray flares \citep[e.g.][]{Larionov2013,Blinov2015,Zhang2016}. The {\em RoboPol} programme\footnote{\url{http://robopol.org}} has been designed for efficient detection of EVPA rotations in statistically rigorously defined samples of gamma-ray--loud and gamma-ray--quiet blazars and to investigate possible correlations between their gamma-ray activity and optical EVPA variability \citep{Pavlidou2014}. {\em RoboPol} started observations at Skinakas observatory, Greece, in May 2013. The EVPA rotations detected during its first two years of operation were presented in \citet[hereafter Papers~I and II]{Blinov2015,Blinov2016}. In Paper~I we presented evidence that at least some EVPA rotations must be physically connected to the gamma-ray flaring activity. We also found that the most prominent gamma-ray flares occur simultaneously with EVPA rotations, while fainter flares may be non-contemporaneous with the rotations. This was taken as evidence for the co-existence of two separate mechanisms producing the EVPA rotations. In Paper~II we showed that the polarization degree decreases during the EVPA rotation events. The magnitude of this decrease is related to the rotation rate in the jet reference frame. Moreover we presented indications that the EVPA rotations cannot be of arbitrary duration and amplitude. In this paper, we present a new set of EVPA rotations that were detected during the third {\em RoboPol} observing season in 2015. Then, using data from all three seasons we study the occurrence frequency of the EVPA rotations in blazars. We aim to determine whether EVPA rotations occur in all blazars with the same frequency and to investigate whether the rotation events are related to the activity of the sources in the gamma-ray band.

\label{sec:conclusion} We have presented a set of EVPA rotations detected by {\em RoboPol} during the 2015 observing season. After three years of operation we have detected 40 EVPA rotations, and thereby more than tripled the list of known events of this type. Our monitoring sample was constructed on the basis of statistically robust and bias-free criteria. It included both gamma-ray--loud and gamma-ray--quiet blazars that were monitored with equal cadence. This allowed us to perform statistical studies of the frequency of EVPA rotations in blazars for the first time. We have shown that the frequency of rotations varies significantly among blazars. None of the control sample blazars displayed a rotation during the monitoring period. Moreover, the EVPA rotations occur with significantly different frequency in different blazars in the main sample. There is a subset of blazars that show the events much more frequently than others. This result is consistent with our analysis in Paper~I, where we showed that rotators have higher EVPA variability than non-rotators even outside the rotating periods. This is a major result of the {\em RoboPol} project: only a fraction of blazars ($\sim$ 28\% of sources in both samples) exhibit EVPA rotations with rates $\le 20$ deg d$^{-1}$ in the optical band, with an average frequency of 1/232\,d$^{-1}$ (in the observer frame). The remaining $\sim$ 72\% of sources did not show any rotations. If they do exhibit rotations, this should happen with a frequency less than $\sim$ 1/3230\,d$^{-1}$. The analysis of Sect.~\ref{sec:reasons} shows that the difference in the frequencies of EVPA rotations cannot be explained either by the difference in the EVPA measurement uncertainties or by differences in redshifts and/or Doppler factors among the blazars. This result should be confirmed using a larger number of objects with known $\delta$. Only a small fraction of blazars in our monitoring sample have Doppler factor estimates available. The ongoing analysis of variability in the radio band will allow us to increase the sample of blazars with known Doppler factors and allow to verify our results with better statistics. The tendency for EVPA rotations to occur in LSP blazars found in Sect.~\ref{sec:rot_classes} can be explained in the same way as higher variability of LSP sources in the total optical flux. It has been shown by \cite{Hovatta2014} that LSP blazars are more variable than HSP in the optical band. This was attributed to the fact that, in the optical band, LSP sources are observed near their electron energy peak, which causes stronger variations of the emission compared to HSP sources, where the lower energy electrons cool down slowly and produce mild variability. For the same reason, the polarized flux density as well as the EVPA must be more variable in LSP sources compared to HSP when observed in the optical band. If this interpretation is correct, then HSP blazars must exhibit EVPA rotations more frequently at higher frequencies (UV and X-ray bands). The dependence of the optical EVPA behaviour on the synchrotron peak position is also reported in two other papers based on RoboPol data. \cite{Angelakis2016} have shown that the EVPA in HSP sources centers around a preferred direction, while in LSP blazars it follows a more uniform distributions. \cite{Hovatta2016} have shown that the scatter in the $Q$-$U$ plane is smaller for HSP blazars than for ISP. This also indicates that the polarization plane direction is more stable in HSP sources. We also found that the rotators seem to be more luminous and more variable in the gamma-ray band than non-rotators. This difference can also be explained by the tendency of the EVPA rotations to occur in LSP sources. These sources have higher luminosities on average than ISP and HSP in the 3FGL because of an instrumental selection effect. The same reason can also explain the increase of their variability indices \citep{Ackermann2015}. For this reason, the optical polarimetry monitoring programmes that select their observing samples on the basis of high variability in the gamma-ray band will observe EVPA rotations more frequently than among blazars on average. The 180$^{\circ}$ EVPA ambiguity sets a fundamental limitation on the rate of EVPA rotations that can be detected under a given cadence of observations. So far we have been able to study rotations with rates $\le 20$\,deg\,d$^{-1}$. There was only one rotation with a rate $\sim$ 50\,deg\,d$^{-1}$ detected by {\em RoboPol}. However, there is an indication in the {\em RoboPol} data as well as in the literature that fast EVPA rotations with rates 60 -- 130\,deg\,d$^{-1}$ do occur in blazars \citep[e.g.][]{Larionov2013}. We plan to extend our studies to higher rotation rates by increasing our cadence for future monitoring.

1607.04292

1607

1607.07258_arXiv.txt

Observations of massive outflows with detectable central AGN typically find them within radii $\lesssim 10$\,kpc. We show that this apparent size restriction is a natural result of AGN driving if this process injects total energy only of order the gas binding energy to the outflow, and the AGN varies over time (`flickers') as suggested in recent work. After the end of all AGN activity the outflow continues to expand to larger radii, powered by the thermal expansion of the remnant shocked AGN wind. We suggest that on average, outflows should be detected further from the nucleus in more massive galaxies. In massive gas--rich galaxies these could be several tens of kpc in radius. We also consider the effect that pressure of such outflows has on a galaxy disc. In moderately gas--rich discs, with gas-to-baryon fraction $< 0.2$, the outflow may induce star formation significant enough to be distinguished from quiescent by an apparently different normalisation of the Kennicutt-Schmidt law. The star formation enhancement is probably stronger in the outskirts of galaxy discs, so coasting outflows might be detected by their effects upon the disc even after the driving AGN has shut off. We compare our results to the recent inference of inside--out quenching of star formation in galaxy discs.

Modern galaxy evolution models typically include feedback from active galactic nuclei (AGN) in order to explain the drop-off in the galaxy mass function compared with the expected halo mass function above $M_* \simeq 10^{11} \msun$, prevent the cooling catastrophe in galaxy clusters and produce the scaling relations between galaxies and their central supermassive black holes (SMBH). The existence of a feedback link has been all but confirmed observationally, with pc-scale relativistic winds \citep[e.g.,][]{Tombesi2010A&A,Tombesi2010ApJ} and massive outflows on scales from sub-kpc \citep{Alatalo2011ApJ} to several kpc \citep{Feruglio2010A&A, Sturm2011ApJ, Rupke2011ApJ, Cicone2014A&A} detected in a number of galaxies, sometimes both types seen in the same object \citep{Tombesi2015Natur}. Despite this robust general picture, a number of questions remain regarding the effects of AGN feedback upon the host galaxies. One question is the range of spatial scales over which outflows are found. Simple models \citep[e.g.,][]{Zubovas2012ApJ} predict outflows propagating with roughly constant velocity out to very large radii. However, AGN outflows are only detected within the central $\sim 10$~kpc of the nucleus \citep[e.g.,][]{Spence2016MNRAS}. The fact that AGN outflows should expand for a long time after the driving AGN switches off \citep[e.g.,][]{King2011MNRAS} might offer an explanation to this problem. Another question is the possibility of the outflow triggering star formation in the galaxy disc. Such a process has been proposed and investigated before \citep{Silk2005MNRAS, Silk2009ApJ, Gaibler2012MNRAS, Zubovas2013MNRASb, Silk2013ApJ, Bieri2016MNRAS}, but there is still no consensus. Observations do not provide a unique answer either, with AGN activity found associated with both elevated \citep[e.g.,][]{Santini2012A&A, Bernhard2016MNRAS} and suppressed \citep[e.g.,][]{Page2012Natur, Carniani2016arXiv} star formation in the host galaxy, while sometimes no connection is evident at all \citep[e.g.,][]{Bongiorno2012MNRAS, Mullaney2012MNRAS}. One interesting piece of evidence is recent detection of star formation being quenched from the inside out in galaxy discs \citep{Tacchella2015Sci}. We suggest that similar behaviour may be expected from disc galaxies affected by AGN outflows. In this paper, we investigate both the propagation of AGN-driven outflows and their effect on galaxy discs by means of a semi-analytical model. We track the expansion of an outflow driven by a flickering AGN by numerically integrating the equation of motion and find that the outflow is more likely to be detected close to the centre of the galaxy than far away. We assume that the growth of the central black hole and the accompanying AGN activity switches off once the total energy injected into the outflow is of order the binding energy of the gas, because this starves the central black hole. With this assumption we show that outflows are only detectable within $\sim 10$~kpc of the centre of the host galaxy while the galaxy nucleus is active. Next we use a simple prescription based on the KS law to calculate the expected star formation rate and dynamical pressure in the galactic disc and compare this with the outflow pressure. We find that AGN outflows can produce a significant enhancement of star formation, especially in the outskirts of galaxy discs. This process helps eventually quench star formation from inside out, similar to the suggestion in recent work \citep{Tacchella2015Sci}, by exhausting the available gas supply. It offers a diagnostic tool for detecting remnant outflows which might be unobservable directly. We structure this paper as follows. In Section \ref{sec:windmodel}, we briefly describe the basics of the AGN wind outflow model. In Section \ref{sec:expansion}, we describe the numerical integrator used to follow the evolution of the outflow driven by a flickering AGN and present the results of outflow size. In Section \ref{sec:burst}, we consider the effect of such an outflow upon the galactic disc and estimate the strength of the starburst. We discuss in Section \ref{sec:discuss} and conclude in Section \ref{sec:concl}.

\label{sec:concl} We have used a largely analytic model to follow the evolution of a spherically--symmetric, energy-driven AGN outflow expanding in a realistic galaxy environment with an NFW halo and a Hernquist bulge. We showed that if that the nucleus remains active only long enough to inject enough energy into the gas to unbind it, then the outflow can only be detected simultaneously with the AGN when its radius is $\lesssim 10$~kpc, consistent with observations. Outflows detected at greater distances in inactive galaxies may point to recent episodes of nuclear activity which have already ended. We also predict that outflows in more massive active galaxies will typically be detected at larger distances from the nucleus than those in lower--mass galaxies. We have shown that the outflow pressure can enhance star formation in a galaxy disc, producing observable differences in the normalisation of the KS law for the host, at least if the disc is not extremely gas--rich. The SFR enhancement is strongest in the centre and the outskirts of the disc ($<1$~kpc and $>10$~kpc from the centre, respectively), so these regions are the best places to look for the effects of recent AGN activity. In particular, enhanced star formation efficiency in the outskirts of a galaxy's disc can be used as another diagnostic of an AGN phase within the last few times $10^7$~yr.

1607.07258

1607

1607.00152_arXiv.txt

{B[e] supergiants are evolved massive stars, enshrouded in a dense wind and surrounded by a molecular and dusty disk. The mechanisms that drive phases of enhanced mass loss and mass ejections, responsible for the shaping of the circumstellar material of these objects, are still unclear.} {We aim to improve our knowledge on the structure and dynamics of the circumstellar disk of the Large Magellanic Cloud B[e] supergiant LHA\,120-S\,73.} {High-resolution optical and near-infrared spectroscopic data were obtained over a period of 16 and 7 years, respectively. The spectra cover the diagnostic emission lines from [\ion{Ca}{ii}] and [\ion{O}{i}], as well as the CO bands. These features trace the disk at different distances from the star. We analyzed the kinematics of the individual emission regions by modeling their emission profiles. A low-resolution mid-infrared spectrum was obtained as well, which provides information on the composition of the dusty disk.} {All diagnostic emission features display double-peaked line profiles, which we interpret as due to Keplerian rotation. We find that the profile of each forbidden line contains contributions from two spatially clearly distinct rings. In total, we find that LHA\,120-S\,73 is surrounded by at least four individual rings of material with alternating densities (or by a disk with strongly non-monotonic radial density distribution). Moreover, we find that the molecular ring must have gaps or at least strong density inhomogeneities, or in other words, a clumpy structure. The optical spectra additionally display a broad emission feature at 6160--6180\,\AA, which we interpret as molecular emission from TiO. The mid-infrared spectrum displays features of oxygen- and carbon-rich grain species, which indicates a long-lived, stable dusty disk. We cannot confirm the previously reported high value for the stellar rotation velocity. \ion{He}{i} $\lambda$ 5876 is the only clearly detectable pure atmospheric absorption line in our data. Its line profile is strongly variable in both width and shape and resembles of those seen in non-radially pulsating stars. A proper determination of the real underlying stellar rotation velocity is hence not possible.} {The existence of multiple stable and clumpy rings of alternating density recalls ring structures around planets. Although there is currently insufficient observational evidence, it is tempting to propose a scenario with one (or more) minor bodies or planets revolving around LHA\,120-S\,73 and stabilizing the ring system, in analogy to the shepherd moons in planetary systems.}

\label{sec:intro} B[e] supergiants (B[e]SGs) are evolved massive stars \citepads{1986A&A...163..119Z}. They form a subgroup of objects displaying the B[e] phenomenon \citepads{Lamers1998}. Their optical spectra display a hybrid character, meaning that they present broad and intense Balmer emission lines and simultaneously narrow emission lines of low-excitation permitted and forbidden transitions of low-ionized and neutral elements (i.e., \ion{Fe}{ii}, [\ion{Fe}{ii}] and [\ion{O}{i}]). In addition, the stars exhibit a very strong near- or mid-infrared excess. \begin{table*} \caption{Observation summary.} \label{tab:obs} \centering \begin{tabular}{lccccccr} \hline \hline Date & $t_{\rm exp}$ & Instrument/ & Resolution & Wavelength \\ (UT) & (s) & Observatory & & Range \\ \hline 1999-12-17 & 4500 & FEROS/ESO-La Silla & 55 000 & $3600-9200\,\AA$ \\ 2005-12-12 & 1500 & FEROS/ESO-La Silla & 55 000 & $3600-9200\,\AA$ \\ 2014-11-28 & 1600 & FEROS/ESO-La Silla & 55 000 & $3600-9200\,\AA$ \\ 2015-10-12 & 1400 & FEROS/ESO-La Silla & 55 000 & $3600-9200\,\AA$ \\ 2008-11-13 & 1800 & Echelle-du Pont/LCO & 45 000 & $3480-10150\,\AA$ \\ 2006-01-16 & 2700 & REOSC/CASLEO & 13 000 & $4560-7100\,\AA$ \\ 2012-11-26 & 2700 & REOSC/CASLEO & 13 000 & $5780-9100\,\AA$ \\ 2012-11-27 & 1800 & REOSC/CASLEO & 13 000 & $5780-9100\,\AA$ \\ 2004-12-17 & 1800 & Phoenix/GEMINI-South & 50 000 & $2.290-2.299\,\mu$m \\ 2010-12-24 & 5920 & Phoenix/GEMINI-South & 50 000 & $2.290-2.299\,\mu$m \\ 2011-01-04 & 5920 & Phoenix/GEMINI-South & 50 000 & $2.319-2.330\,\mu$m \\ 2012-10-03 & 400 & T-ReCS/GEMINI-South & 100 & $8.0-13.0\,\mu$m \\ \hline \end{tabular} \end{table*} The B[e] phenomenon is typically related to the physical conditions of the gaseous and dusty shells, rings, or disks surrounding a hot star and not to the properties of the star itself. The origin of circumstellar envelopes of B[e]SGs is attributed to the mass lost from the star either through dense, aspherical stellar winds or through sudden mass ejections expected to occur during short-lived phases in the post-main-sequence evolution of the stars. The possible structure of the B[e]SG circumstellar envelopes was described by \citetads{1985A&A...143..421Z}, who proposed an empirical model, consisting of a hot and fast line-driven wind in the polar regions, and a slow, much cooler and denser (by a factor of $10^2-10^3$) wind in the equatorial region. In this cool and dense equatorial region, the material condenses into molecules and dust. Molecular emission from carbon monoxide (CO), tracing the inner rim of the molecular region, was found in many B[e]SGs \citepads{1988ApJ...324.1071M, 1988ApJ...334..639M, 1989A&A...223..237M, 1996ApJ...470..597M, 2010MNRAS.408L...6L, 2012A&A...548A..72C, 2012MNRAS.426L..56O, 2013A&A...558A..17O, 2012A&A...543A..77W, 2013A&A...549A..28K}. Emission from less prominent molecules such as titanium oxide \citepads[TiO,][]{1989A&A...220..206Z, 2012MNRAS.427L..80T} and silicon oxide \citepads[SiO,][]{2015ApJ...800L..20K} was also reported for a few objects. In the outer parts of these equatorial regions, where the temperature drops below the dust condensation value, dust forms. The large amount of dust connected with B[e]SGs is visible as strong near- and mid-infrared excess emission \citepads{1986A&A...163..119Z, 2009AJ....138.1003B, 2010AJ....140..416B} and resolved spectral dust features \citepads{2006ApJ...638L..29K, Kastner2010}. Interferometric observations of some Galactic B[e]SGs revealed that the dust is concentrated in a ring \citepads{2007A&A...464...81D, 2011A&A...525A..22D, 2012A&A...548A..72C}. In a few cases, the central object turned out to be a close binary, and the dusty rings are circumbinary \citepads{2011A&A...526A.107M, 2012A&A...538A...6W, 2012A&A...545L..10W}. Investigations of the kinematics within the gaseous (atomic and molecular) disk regions often revealed that it is consistent with Keplerian rotation \citepads{2011A&A...526A.107M, 2012MNRAS.423..284A, 2012A&A...548A..72C, 2012A&A...543A..77W, 2010A&A...517A..30K, 2013A&A...549A..28K, 2014ApJ...780L..10K, 2015ApJ...800L..20K, 2015AJ....149...13M}. In some cases, observations support evidence of disk variability as seen in LHA\,115-S\,18 \citepads{2012MNRAS.427L..80T} and HD\,327083 \citepads{2013msao.confE.160K}, sudden disk formation as in LHA\,115-S\,65 \citepads{2012MNRAS.426L..56O}, and disk dissipation as in CI\,Cam \citepads{2014MNRAS.443..947L}. To improve our knowledge on the disk formation process, the physical properties of the disks, and the timescales of possible variability and its origin, extensive observational monitoring of individual objects is indispensable. In the present work, we focus on the Large Magellanic Cloud (LMC) B[e]SG star LHA\,120-S\,73. We investigate its gaseous (atomic and molecular) environment based on multi-epoch high-resolution optical and near-infrared spectroscopic observations and its dusty disk based on low-resolution mid-infrared data.

\label{sec:concl} We investigated the structure and kinematics of the circumstellar environment of the B[e]SG star LHA\,120-S\,73 in the LMC based on combined sets of high-resolution optical and near-infrared spectroscopic data collected over a time interval of 16 and 7 years, respectively. The near-infrared spectra cover the region of the first and second CO band heads. The high spectral resolution revealed rotational broadening of the individual rotation-vibrational CO lines with a deprojected rotation velocity of 34\,km\,s$^{-1}$. In addition, we found that the CO band head emission displays intensity variations, which we interpreted as density inhomogeneities within the molecular ring. In the optical spectra we discovered an emission feature that we identified as molecular emission from TiO. LHA\,120-S\,73 is hence the fourth B[e]SG with confirmed TiO band emission. In general, the optical emission features show no or only little indication for variability, suggesting that both the wind and the circumstellar material are in quasi-stationary state. The spectra display emission of the forbidden lines of [\ion{Ca}{ii}] and [\ion{O}{i}], which are diagnostic tracers for circumstellar rings or disks. The lines have double-peaked profiles, which we interpreted, in analogy to previous studies, as due to Keplerian rotation. Modeling of the profile shapes revealed that each line profile consists of emission from two individual, spatially clearly distinct rings. The optical spectra also reveal one clear absorption line, identified as \ion{He}{i} $\lambda$ 5876. This line was formerly used to determine $\varv\sin i$ of the star and to demonstrate that the star is rapidly rotating. However, from our data we could not confirm the previously reported high value of $\varv\sin i$. Instead, our spectra show that this line is clearly variable and highly asymmetric, reminding of line profiles of non-radially pulsating stars. Although the $S/N$ of our data is too poor to favor a pulsational origin of the line asymmetries over line pollution due to emission, the presence of at least two periodic variabilities determined from the photometric light curve of LHA\,120-S\,73 provides evidence that pulsations might indeed play a non-negligible role in the dynamics of the stellar atmosphere. In summary, we found that the global structure of the circumstellar material, traced by the forbidden emission lines and the CO molecular bands, consists of at least four distinct rings of alternating densities. Although we have currently no observational proof, we might speculate whether such a ripple structure of the circumstellar material might be caused by the presence of minor bodies (e.g. in the form of planets) that evacuated the space in between the rings, caused density inhomogeneities, and stabilize the currently observable ring structures, in analogy to the shepherd moons within the ring systems of planets or the rings observed around pre-main sequence objects. Follow-up observations with reasonable time resolution are clearly needed to investigate in more detail the variability, in particular of the molecular emission, to obtain better constraints on the density structure of the environment of LHA\,120-S\,73.

1607.00152

1607

1607.07919_arXiv.txt

The central star of the planetary nebula NGC 1514 is among the visually brightest central stars in the sky (V=9.5). It has long been known to show a composite spectrum, consisting of an A-type star and a much hotter star responsible for the ionization of the surrounding nebula. These two stars have always been assumed to form a binary system. High-resolution spectrograms obtained with Espadons at the CFHT on Mauna Kea have allowed to measure good radial velocities for both stars. They differ by 13 $\pm$ 2 km s$^{-1}$. The stellar velocities have not changed after 500 days. We have also estimated the metallicity of the cooler star. Combining these data with other information available in the literature, \edit1{we conclude that, unless all the published nebular radial velocities are systematically wrong,} the cooler star is just a chance alignment, and the two stars are not orbiting each other. The cooler star cannot have played any role in the formation of NGC 1514.

\label{sec:intro} NGC 1514 belongs to a small group of planetary nebulae with A-type central stars. Obviously in each of these cases a hotter star needs to be present, to explain the ionized state of the surrounding nebula. Thus the central star of NGC 1514 has always been assumed to be a binary system. We will not give an extended historical introduction on NGC 1514; the reader is directed to a recent paper by Aller et al. (2015). That the central star spectrum is composite was shown by Kohoutek (1967), and confirmed spectroscopically by Greenstein (1972). Greenstein's spectra did not show any convincing velocity variation, but he did note that He {\sc ii} $\lambda$4686 from the hot star showed a mean velocity difference of $-7 \pm 5$ km s$^{-1}$ relative to the H and metal lines of the A-type star. Other observers have given conflicting reports on radial velocity variations. According to Seaton (1980), this was not satisfactorily resolved. Since then, it appears that nobody ever tried again. We have obtained new high-resolution spectrograms of this central star, and report the resulting radial velocities in what follows. Section 2 describes the observations. Section 3 explains how the velocities were measured, and presents the results. \edit1{In Section 4 we give a preliminary estimate of the metallicity of the cooler star. In Section 5 we discuss all the information we have collected, and Section 6 recapitulates the conclusions}.

The average velocity of the cooler star in Table 3 is 44 $\pm$ 2 km s{$^{-1}$}. The average velocity of the hotter star in Table 4 is 57 $\pm$ 1 km s{$^{-1}$}. The small differences between the average velocities from the different lines are probably due to small errors in the laboratory wavelengths. The velocities of the two stars differ significantly, by 13 $\pm$ 2 km s{$^{-1}$}. None of the two stars have shown substantial velocity variations. Assume the two stars are orbiting each other. The different velocities indicate that we cannot be observing the orbits pole-on. Then the lack of variations of the order of 10 km s{$^{-1}$} in almost two years would require a long orbital period. The results in Tables 3 and 4 can be used to reject all previous claims or suggestions of high-amplitude, short-period radial velocity variations (e.g. as discussed by Muthu \& Anandarao 2003 in their section 5). \edit1{This implies that the cooler star cannot have played any direct role in the formation of NGC 1514, e.g. through a common envelope episode.} There is some additional information to be extracted from the existing literature. \subsection{Radial velocity of the cooler star} The only previous study at high spectral resolution is by Greenstein (1972). There are coud\'e spectrograms from 1949 to 1971. The average radial velocity (47.6 $\pm$ 1.6 km s{$^{-1}$}) looks quite in agreement with the values reported in Table 3. The Greenstein velocities are not directly comparable to those in Table 3, because he included several lines we have rejected, like the Balmer lines; did not give the velocities for each spectral feature separately; nor did he report what laboratory wavelengths he used. But given the dispersion of values listed in his Table 1, and his conclusion (no sign of velocity variation, in his own words), it is reasonable to conclude that there is no reliable evidence of variations in the velocity of the cooler star in six different epochs from 1949 to 2016. \subsection{Radial velocity of the hotter star} In Greenstein's paper the velocity of He {\sc ii} 4686 is mentioned to be more negative than that of the Balmer lines and the metals by 7 $\pm$ 5 km s{$^{-1}$}. He did not elaborate on this, so probably he did not consider the difference to be significant. The velocity we measured on the co-added Espadons spectrum for He {\sc ii} 4686 is 52 km s{$^{-1}$}. We did not include velocities from this line in Table 4 because the profile is affected by what appears to be an unidentified emission line at 4688 \AA \ (see Figure 1). Another reason to avoid this line as a radial velocity indicator is that it is among the first, in the visible spectral range, to suffer the effects of a stellar wind. Even a very incipient P Cygni-like profile would induce a shortward wavelength displacement in the observed absorption (Kudritzki et al.\ 2006). \edit1{Continuous monitoring for at least a decade would be required to verify if Greenstein's data were implying long-period variability of the hotter star's velocity.} \subsection{Radial velocity of the nebula} According to Schneider et al. (1983), the heliocentric radial velocity of NGC 1514 is 60 $\pm$ 4 km s{$^{-1}$} (an average of six different determinations, each with a reported uncertainty of between 3 and 9 km s{$^{-1}$}). This is perfectly compatible with the velocity of the hotter star in Table 4. Two velocities in Schneider et al. are discrepant: one substantially above the average, and one below it. The one above can be ignored, because of its lower accuracy. The one below is more interesting: it is 41 $\pm$ 5 km s{$^{-1}$, from Greenstein (1972). There is a way of understanding the discrepancy. Greenstein reports that the spectrograms he used to measure the nebular velocity were taken with the slit located one arc minute south of the central star. There is a spatiokinematic study of NGC 1514 by Muthu \& Anandarao (2003), made with a Fabry-Perot spectrometer. Unfortunately they did not attempt to measure the absolute radial velocity; but they do show velocity channel maps (their Figure 4). In this figure we find that slightly less than one arc minute south of the central star there is a concentration of gas with a velocity of -13 km s{$^{-1}$. There is no corresponding gas at the opposite velocity (+13). That is conceivably the reason why Greenstein's nebular velocity was more negative. Anyway, it is fair to say that systematic errors in the published nebular velocities cannot be completely discarded. Integral field spectroscopy of NGC 1514 at high spectral resolution would permit to accurately map the whole velocity field, and produce a more reliable number for the systemic radial velocity of this planetary nebula. \subsection{Binary system or chance alignment?} Since the radial velocities of the two central stars differ, \edit1{and the nebular velocity agrees with that of the hotter star,} there is a distinct possibility of chance alignment \edit1{with the cooler star}. This would be extremely surprising, because Ciardullo et al. (1999) were unable to resolve the pair on Hubble Space Telescope (HST) images. The probability of such a perfect chance alignment is extremely low, of the order of $10^{-8}$ (using an area of 0.01 square arc second, and an estimated surface density of 10 stars of 10th visual magnitude per square degree at the galactic latitude of NGC 1514, which is 15 degrees). See e.g. Figure 4 of Bahcall \& Soneira (1980). \edit1{ In dealing with such a low probability, every possible alternative, however unlikely, must be considered. Our anonymous referee has kindly provided a few.} \edit1{ Let us first consider if the velocity of the hotter star could be higher because of gravitational redshift. This would require its log $g$ to be around 7.5. But this replaces one problem with another: since the hotter star is rather bright, its distance from us would have to be much less than 50 pc, ruling out any association with the cooler star.} \edit1{ Next, imagine that the velocity of the cooler star is different because of a recoil effect produced in a binary system when NGC 1514 was formed. The problem here is that, in order to produce a substantial recoil effect, the orbital period would have to be longer than the time scale for nebular ejection. But for such long orbital periods, the orbital velocities would not be enough to explain the observed velocity difference between the two stars (a few examples are given in Table 5, to be explained below).} \edit1{ Finally, under what conditions can the two central stars of NGC 1514 be members of a binary system? In Section 4 we concluded that the cooler star must be rather massive, around 3 M$\odot$. Assume a typical central star mass of 0.6 M$\odot$ for the hotter star. If these two stars are orbiting each other, the nebular velocity must be close to the velocity of the cooler star. In other words, all the nebular measurements discussed in the previous subsection would have to be systematically wrong. } Consider the value of the sum of the radial velocity semiamplitudes as a function of the orbital period. Table 5 shows the values of the sum of the orbital semimajor axes (a$_1$ + a$_2$) and the sum of the radial velocity semiamplitudes (K$_1$ + K$_2$) as functions of the orbital period, using Newton's version of Kepler's 3rd law, with a total mass of \edit1{3.6} M$\odot$, and using the relation between semiamplitude, orbital period, semimajor axis and eccentricity. Taking an orbital inclination of 53 degrees and an eccentricity e=0.6, the factors sin $i$ and $(1-e^2)^{0.5}$ cancel each other. \edit1{Unfavorable (smaller) values of inclination} would produce smaller numbers in the third column of Table 5. \edit1{An orbital eccentricity $e$=0.7 would increase those third-column numbers by a factor 1.4. But on the other hand, most of the time, the stars would be far from periastron, and the velocity difference between the two stars would be actually smaller.} \floattable \begin{deluxetable}{rrr} \tablecaption{Radial velocity amplitudes for different periods \label{tab:table5}} \tablehead{ \colhead{Period} & \colhead{a$_1$ + a$_2$} & \colhead{K$_1$ + K$_2$} \\ \colhead{(years)} & \colhead{(R$\odot$)} & \colhead{(km s{$^{-1}$)}} } \startdata 1 & 329 & 46 \\ 10 & 1529 & 21 \\ 30 & 3180 & 15 \\ 100 & 7096 & 10 \\ \enddata \end{deluxetable} Table 5 \edit1{suggests an upper limit of about 50} years on the possible orbital period. Any longer period would require a value of K$_1$ + K$_2$ smaller than the observed velocity difference of 13 km s{$^{-1}$. \edit1{It will not take more than a few decades to decide conclusively if these two stars can actually be related.} \subsection{Radial velocities: summary and consequences} The Espadons spectrograms show that the radial velocities of the two central stars of NGC 1514 are clearly different. The constancy of these high-quality velocities over a time of two years permits to reject all previous claims or suggestions of substantial short-term changes in the stellar velocity. \edit1{Even if they orbit each other, the implied separation between the two stars precludes any past evolutionary interaction like a common-envelope episode.} For all practical purposes, these two stars are \edit1{evolutionarily} unrelated. \edit1{On the other hand, if they orbit each other, arguing from Kepler's 3rd law, given the observed velocity difference, we do not expect an orbital period longer than about 50 years.} \edit1{If new measurements of the nebular radial velocity confirm the values reported in the literature, which agree with the radial velocity of the hotter star,} we will have to conclude that, unlikely as it may seem, the cooler star is a chance alignment. There is the problem that Ciardullo et al. (1999) could not resolve the two stars on their HST images. If there is a significant difference in the proper motions, then we would expect that sooner or later the two stars will become resolvable. In 1791 William Herschel reported that the star was perfectly in the center of NGC 1514 (see e.g.\ Greenstein 1972), and it still is today. \edit1{The proper motion of the cooler star, measured by the Hipparcos satellite, is 8.4 milliarcsec/yr. This amounts to about 1.7 arcsec since the time of Herschel. On the other hand, the galactic longitude of NGC 1514 is 165 degrees, i.e., almost in the anticenter direction; so it would not be terribly surprising to find the two stars, even if at somewhat different distances from us, having similar proper motions. What is required now is a new observation with the highest possible angular resolution, which we expect to make in the near future.} The results reported here do not preclude the existence of low-amplitude velocity variations; both stars could have planetary systems or low-mass companions. Only future measurements of comparable or higher quality will tell.

1607.07919

1607

1607.00014_arXiv.txt

We investigate the star formation rate (SFR) dependence on the stellar mass and gas-phase metallicity relation at $z=2$ with MOSFIRE/Keck as part of the ZFIRE survey. We have identified 117 galaxies (1.98$\leq z\leq$2.56), with $8.9\leq$log(M/M$_{\odot}$)$\leq11.0$, for which we can measure gas-phase metallicities. For the first time, we show discernible difference between the mass--metallicity relation, using individual galaxies, when deviding the sample by low ($<10$~M$_{\odot}$yr$^{-1}$) and high ($>10$~M$_{\odot}$yr$^{-1}$) SFRs. At fixed mass, low star-forming galaxies tend to have higher metallicity than high star-forming galaxies. Using a few basic assumptions, we further show that the gas masses and metallicities required to produce the fundamental mass--metallicity relation, and its intrinsic scatter, are consistent with cold-mode accretion predictions obtained from the OWLS hydrodynamical simulations. Our results from both simulations and observations are suggestive that cold-mode accretion is responsible for the fundamental mass--metallicity relation at $z=2$ and demonstrates the direct relationship between cosmological accretion and the fundamental properties of galaxies.

\label{sec:intro} Mass and metallicity are arguably the most fundamental properties of galaxies since they reflect stellar build-up/evolution and the cycling of baryons through outflows and accretion. A clear consequence of these processes is the galaxy stellar mass--metallicity relation, which has been observed up to $z\sim4$ \citep[e.g.,][]{tremonti04,erb06,steidel14,troncoso14,zahid14}. Cosmological simulations continue to show that cosmic accretion provides significant fuel for star formation resulting in galaxy mass growth and chemical enrichment via outflows \citep{dekel09,freeke12}. They also show that most of this gas accretes in the `cold mode', i.e., without being heated by a virial shock \citep[e.g.,][]{keres05,fg11,freeke11}. Metal-poor cold-gas accretion has been observed on the periphery of galaxies, using absorption lines detected in background quasar spectra \citep[e.g.,][]{kacprzak12b,crighton13,bouche13}, which are also dynamically consistent with large-scale cosmic web accretion \citep[e.g.,][]{steidel02,kacprzak10}. These inward flows are likely why, at a given stellar mass, metallicity is dependent on the star formation rate (SFR) \citep{mannucci10,bothwell13,forbes14,maier15}. It has been found that at a fixed mass, galaxies with high SFRs have lower gas-phase metallicities than galaxies with low SFRs. This fundamental mass--metallicity relation \citep[FMR: see][]{mannucci10} exhibits little scatter at $z=0$, which implies an equilibrium between inflowing and outflowing gas and star formation. At higher redshifts, it is still questionable if a fundamental mass--metallicity relation exists since at $z\sim2$, the mass--metallicity relation is highly scattered for stellar masses less than 10$^{10}$M$_{\odot}$ \citep{sanders14,kacprzak15,tran15}. This observed scatter may not be surprising given that if a significant metal-poor accretion event occurs, which is especially common at high redshifts, it is expected that the galaxy's integrated metallicity decreases and that it experiences a boost in star formation. Thus, we should expect to see a fundamental mass--metallicity relation at higher redshifts as well. The time-scale between accretion and induced star formation is still uncertain. However, direct observations of metallicity and SFR profiles of local metal-poor galaxies have shown that highly star-forming regions within the galaxy have roughly a factor of 10 lower metallicity than the total galaxy metallicity; indicating that accreted gas is converted into stars within $\sim100$~Myr \citep{sanchez15}. Thus, this direct relation between accretion, SFR and metallicity is critical for understanding galaxy formation and evolution. We further explore the relationship between SFR and the mass--metallicity relation at $z\sim2$, and examine the hypothesis that accretion could cause the scatter in the mass--metallicity relation. Here, we present Keck/MOSFIRE observations from the ZFIRE Survey (Nanayakkara et al.\ submitted) of 117 galaxies for which we have gas-phase metallicities, stellar masses, SFRs and gas masses described in \S~\ref{sec:data}. In \S~\ref{sec:results}, we show that there is a fundamental mass--metallicity relation driven by SFR. We further use our data, with some underlying basic assumptions, to compute the gas accretion masses and metallicities, which are shown to be in agreement with cosmological simulations. We adopt a $h=0.70$, $\Omega_{\rm M}=0.3$, $\Omega_{\Lambda}=0.7$ cosmology for our ZFIRE Survey.

\label{sec:conclusion} We have identified a fundamental mass--metallicity relation at $z\sim2$ ($4.4\sigma$), whereby galaxies with low star formation rates ($<10$~M$_{\odot}$yr$^{-1}$) exhibit higher metallicities than high star-forming ($<10$~M$_{\odot}$yr$^{-1}$). We have examined whether the activity of cold accretion could drive this scatter in the fundamental mass--metallicity relation. Given the simplicity of our assumptions, it is interesting to see that the metallicities and accreted masses required to reduce the scatter in the mass--metallicity relation at $z\sim2$ are consistent with and show similar trends as in our simulations. Therefore, it is tempting to suggest that gas accretion is solely responsible for the existence of the fundamental mass--metallicity relation. Our conclusions are consistent with previous works at low redshift that attribute the observed scatter in the mass--metallicity relation seen at very low stellar masses to accretion and/or mergers \citep{bothwell13,forbes14}. These results demonstrate the direct relationship between cosmological accretion and the fundamental properties of galaxies. We note that at a fixed metallicity, a higher ionization parameter produces lower {\NII}/{\Ha} ratios and thus lower N2-metallicities \citep[e.g.,][]{kewley02}. The existence of shock will elevate {\NII}/{\Ha} ratios \citep[e.g.,][]{rich11,yuan12}. However, we do not think ionization parameters/shocks are the main driver for the scatter of the MZ relation because it is difficult to explain the bifurcation in SFR since ionization parameter is usually positively correlated with specific SFR. We will study the effect of ionization parameter and/or shocks in future work. Our interpretation that accretion drives the fundamental mass--metallicity relation may not be unexpected given that $z\sim2$ is near the peak epoch of star formation, where outflows are ubiquitous \citep[e.g.,][]{steidel10} and cold accretion likely occurs \citep[e.g.,][]{bouche13} delivering a significant amount of metal-poor gas to galaxies. These objects are ideal targets for ALMA to determine whether high star-forming galaxies indeed have significantly higher gas fractions and vice versa.

1607.00014

1607

1607.08239_arXiv.txt

{Radial velocity and transit methods are effective for the study of short orbital period exoplanets but they hardly probe objects at large separations for which direct imaging can be used.}{We carried out the international deep planet survey of 292 young nearby stars to search for giant exoplanets and determine their frequency.}{We developed a pipeline for a uniform processing of all the data that we have recorded with NIRC2/Keck II, NIRI/Gemini North, NICI/Gemini South, and NACO/VLT for 14 years. The pipeline first applies cosmetic corrections and then reduces the speckle intensity to enhance the contrast in the images.}{The main result of the international deep planet survey is the discovery of the HR\,8799 exoplanets. We also detected 59 visual multiple systems including 16 new binary stars and 2 new triple stellar systems, as well as $2,279$ point-like sources. We used Monte Carlo simulations and the Bayesian theorem to determine that $1.05\rlap{\textsuperscript{\tiny+2.80}}\textsubscript{\tiny-0.70}\ \%$ of stars harbor at least one giant planet between $0.5$ and $14$\,M$_\mathrm{J}$ and between $20$ and $300$\,AU. This result is obtained assuming uniform distributions of planet masses and semi-major axes. If we consider power law distributions as measured for close-in planets instead, the derived frequency is $2.30\rlap{\textsuperscript{\tiny+5.95}}\textsubscript{\tiny-1.55}\ \%$, recalling the strong impact of assumptions on Monte Carlo output distributions. We also find no evidence that the derived frequency depends on the mass of the hosting star, whereas it does for close-in planets.}{The international deep planet survey provides a database of confirmed background sources that may be useful for other exoplanet direct imaging surveys. It also puts new constraints on the number of stars with at least one giant planet reducing by a factor of two the frequencies derived by almost all previous works.}

After the first discoveries of exoplanets by indirect detections in the late 80s and 90s, several teams performed surveys to obtain direct images of substellar objects in the optical and near-infrared. As the instruments were not optimized for high contrast imaging, the first surveys probed brown dwarfs that are brighter than planets and thus, easier to detect. For example, \citet{becklin88}, \citet{schroeder00}, \citet{gizis01}, \citet{oppenheimer01}, and \citet{carson09} observed nearby stars, while \citet{neuhauser04} and \citet{lowrance05} surveyed young systems, and \citet{chauvin06} targeted stars harboring planets detected by stellar radial velocity. Taking advantage of new observing modes, such as spectral differential imaging \citep[SDI;][]{marois00} and angular differential imaging \citep[ADI;][]{marois06} and using more powerful adaptive optics systems, later surveys of youthful stars probed fainter and fainter substellar companions down to the young gas-giant extrasolar planet regime \citep{biller07,kasper07,lafreniere07a,nielsen08,metchev09,chauvin10,heinze10,rameau13a,yamamoto13,biller13,nielsen13,wahhaj13,brandt14}. For 14 years, we had been using Keck II, Gemini North and South, and VLT to run two near-infrared imaging surveys: a Keck adaptive optics search for young exoplanets \citep{kaisler03} and the international deep planet survey \citep[IDPS;][]{marois10c}. In this paper, we merge the two and call IDPS the resulting survey. The targets were 292 young and nearby stars of A to M spectral types with a majority of massive stars. The main objectives of the IDPS were the detection and spectral characterization of new exoplanets, and the determination of the frequency of stars harboring giant planets with long orbital periods. A first paper \citep{vigan12} presented a fraction of the A stars of the IDPS. In the current paper, we present the complete survey. The target sample is described in \S\,\ref{sec:target}. The observations and instruments that we used are presented in \S\,\ref{sec:obs}. Then, \S\,\ref{sec:pipeline} details the pipeline that we developed for the uniform data processing of the $\sim$30,000 frames. The characteristics of all the detected sources (multistellar systems, exoplanets, and field stars) are listed in \S\,\ref{sec:results}. Finally, we run a Monte Carlo analysis in \S\,\ref{sec: plafreq} to constrain the frequency of stars with giant planets.

We have completed the IDPS for giant planets around $292$ young nearby stars of all spectral types with a majority of massive stars. We developed a pipeline for a uniform processing of the data that have been recorded for $14$ years using several instruments: NIRC2/Keck II, NIRI/Gemini north, NICI/Gemini South, and NACO/VLT. We achieved contrasts of $\sim12.5\pm2.5$ magnitude at $1''$ at H, CH4, K and Lp bands. We detected a total of $2,279$ point-like sources. Most of these sources were confirmed to be background objects. Four were confirmed to be exoplanets. They are the now well-characterized HR\,8799 planets \citep{marois08,marois10}. We also discovered $16$ stellar binary systems and $2$ triple stars. We used the Bayesian formalism developed in \citet{lafreniere07a} to derive the frequency of stars with giant planets from the detection limits of the survey as well as the confirmed exoplanets. To complete the IDPS sample, we combined our results with the \citet{lafreniere07a} and \citet{chauvin10} surveys that mainly observed low-mass stars (G, K and M stars). The complete sample includes $356$ stars of all spectral types with a median age of $100$\,Myr and a median distance of $37$\,pc. We determined that $1.05\rlap{\textsuperscript{\tiny+2.80}}\textsubscript{\tiny-0.70}\ \%$ of stars harbor at least one giant planet of $0.5$-$14$\,M$_\mathrm{J}$ between $20$ and $300$\,AU. We also found no evidence that giant planets at large separations are more likely formed around BAF stars than around GKM stars. We confirm previous results reducing the measured frequencies of stars with at least one giant planet by a factor of two or more in almost all cases. The fact that the results of the different studies are consistent is encouraging, but we should keep in mind that each team uses different assumptions (planet atmosphere models, orbital parameter distributions, star aging, etc.) and one might wonder on which assumption(s) the derived frequencies mainly depend. For example, we showed in this paper that the assumption on the planet mass and semimajor axis distributions can change the conclusions and it will be essential to address this question to prepare the interpretation of the Gemini Planet Imager \citep{macintosh08} and the Spectro-Polarimetric High-contrast Exoplanet REsearch \citep{beuzit08} surveys. Most of their targets are part of the IDPS sample but the extreme adaptive optics systems will probe lighter and closer-to-their-star exoplanets, which will complete our knowledge of giant planets at large separations.

1607.08239

1607

1607.06154_arXiv.txt

We use the halo occupation distribution (HOD) framework to characterise the predictions from two independent galaxy formation models for the galactic content of dark matter haloes and its evolution with redshift. Our galaxy samples correspond to a range of fixed number densities defined by stellar mass and span $0 \le z \le 3$. We find remarkable similarities between the model predictions. Differences arise at low galaxy number densities which are sensitive to the treatment of heating of the hot halo by active galactic nuclei. The evolution of the form of the HOD can be described in a relatively simple way, and we model each HOD parameter using its value at $z=0$ and an additional evolutionary parameter. In particular, we find that the ratio between the characteristic halo masses for hosting central and satellite galaxies can serve as a sensitive diagnostic for galaxy evolution models. Our results can be used to test and develop empirical studies of galaxy evolution and can facilitate the construction of mock galaxy catalogues for future surveys.

\label{Sec:Intro} In the standard cosmological framework, galaxies form, evolve and reside in dark matter haloes. A key requirement of this framework is to understand how galaxies populate dark matter haloes. What determines how many galaxies are hosted by a dark matter halo? How do the properties of galaxies depend on the mass of the halo? These questions lie at the core of galaxy formation theory. The answers are also crucial if we are to take full advantage of the next generation of galaxy surveys, which aim to make pristine clustering measurements to pin down the nature of dark energy. The extraction of cosmological inferences from these data will no longer be dominated by statistical errors but instead will be limited by the accuracy of our theoretical models. Understanding how galaxies relate to the underlying dark matter is thus essential for optimally utilizing the large-scale distribution of galaxies as a cosmological probe. The clustering of dark matter is dominated by gravity and can be computed reliably with cosmological N-body simulations. However, the detailed physics of galaxy formation -- gas cooling, star formation, and feedback effects -- is only partially understood, so that the relation between galaxies and the underlying dark matter cannot be predicted robustly from first principles. A useful approach to study this is semi-analytic modeling (SAM) of galaxy formation (for reviews see, e.g., \citealt{Cole:2000,Baugh:2006,Benson:2010,Somerville:2015}). In such models, haloes identified in N-body simulations are ``populated'' with galaxies using analytical prescriptions for the baryonic processes. Following the dark matter merger trees, galaxies merge and evolve as new stars form and previous generations of stars change. Different feedback or heating mechanisms, such as those caused by star formation, active galactic nuclei, or the photo-ionizing ultra-violet background, are also incorporated. SAMs have been successful in reproducing a range of observed properties including stellar mass functions and galaxy luminosity functions (see, e.g., \citealt{Bower:2006,Croton:2006,Fontanot:2009,Guo:2011,Guo:2013a,Gonzalez:2014,Padilla:2014,Henriques:2015,Lacey:2016,Croton:2016}). The connection between the mass of a dark matter halo and the galaxies which populate it is often expressed through the halo occupation distribution (HOD) framework (e.g., \citealt{Jing:1998a,Benson:2000,Peacock:2000,Seljak:2000,Scoccimarro:2001,Berlind:2002,Berlind:2003,Cooray:2002,Yang:2003,Kravtsov:2004,Zheng:2005,Conroy:2006}). The HOD formalism describes the ``bias'' relation between galaxies and mass at the level of individual dark matter haloes, in terms of the probability distribution that a halo of virial mass $M_{\rm h}$ contains $N$ galaxies which satisfy a particular selection criterion. It transforms measures of galaxy clustering into a physical relation between galaxies and dark matter haloes, setting the stage for detailed tests of galaxy formation models. The HOD approach has proven to be a very powerful theoretical tool to constrain the galaxy-halo connection and has been applied to interpret clustering data from numerous surveys at low and high redshifts (e.g., \citealt{Jing:1998b,Jing:2002,Bullock:2002,Moustakas:2002,vandenBosch:2003,Magliocchetti:2003,Yan:2003,Zheng:2004,Yang:2005,Zehavi:2005,Zehavi:2011,Cooray:2006,Hamana:2006,Lee:2006,Lee:2009,Phleps:2006,White:2007,White:2011,Zheng:2007,Zheng:2009,Blake:2008,Brown:2008,Quadri:2008,Wake:2008,Wake:2011,Kim:2009,Simon:2009,Ross:2010,Coupon:2012,Coupon:2015,delaTorre:2013,Krause:2013,Parejko:2013,Guo:2014b,Durkalec:2015,Kim:2015,McCracken:2015,Skibba:2015}). HOD models have mostly been used to constrain the relation between galaxies and haloes at a fixed epoch, as the HOD approach by itself does not offer any guidance as to how to treat the evolution of the galaxy population over cosmic time. Attempts to study galaxy evolution using this framework have for the most part explored ``snapshots'' of clustering at different epochs to empirically constrain the evolution (e.g., \citealt{Zheng:2007,White:2007,Wake:2008,Wake:2011,Abbas:2010,Coupon:2012,delaTorre:2013,Guo:2014b,Manera:2015,Skibba:2015}). but a complete model for the evolution of the HOD is still missing. Our goal is to remedy this situation and develop a theoretical model for this evolution by studying how the HOD changes with time. A simplified approach in this vein was taken by \citet{Seo:2008} who studied the predictions for passive evolution of the HOD by populating simulations with galaxies according to a range of assumed HOD and tracking their evolution with time. That work is of limited use due to the unphysical assumption of passive evolution. The form of the HOD at different redshifts has also been studied in the context of abundance matching modeling \cite{Kravtsov:2004,Conroy:2006}. Here we will perform a comprehensive study of the evolution of the HOD using the predictions of semi-analytic modeling which captures the important galaxy formation physics. The present paper builds upon our work exploring the predictions of SAM of galaxy formation, focusing on the connection between galaxies and their host dark matter haloes. \cite{C13} demonstrated that SAMs from different groups give consistent predictions for the galaxy correlation function on large scales, for samples constructed to have the same abundance of galaxies, and that the differences on small scales (the so-called one-halo term) can be readily understood in terms of the choices made about the placement of galaxies within dark matter haloes. In a second paper, we examined the connection between different galaxy properties and the mass of the dark matter halo hosting the galaxy \citep{C15}. We found that some properties, such as stellar mass, depend on subhalo mass in a monotonic fashion (albeit with a scatter), whereas others, such as the cold gas mass, have a more complex dependence on halo mass. Here we use the HOD formalism to compare how different models populate haloes with galaxies over cosmic time. We study the output of two independently developed SAMs, originally from the Durham and Munich groups, at different number densities. This allows us to assess which features of the predicted HODs are generic and which are sensitive to the details of the modelling of the various physical processes, and how best to describe the evolution of the HOD at a given number density. We also consider some simplified empirical models for the evolution of the galaxy distribution and show how these differ from the predictions of the SAMs. This study will enable the incorporation of evolution into the halo models, an aspect that is absent from the standard implementation. The applications of this are two fold. First, from the observational side, it will allow for a consistent combined analysis of clustering measures over a range of epochs, in order to constrain galaxy formation and evolution. A second application of the HOD is to quantify evolution in the galaxy population, which will facilitate the creation of realistic mock galaxy catalogues for surveys which span a large range of lookback times. Accurate estimates exist for the form of the HOD using measurements of the galaxy clustering in, for example, local surveys (e.g. \citealt{Zehavi:2011}). These can be used, in conjunction with a sample of dark matter haloes extracted from an N-body simulation to build a mock catalogue with the same clustering properties and abundance of galaxies. The problem then is how to extend this approach to build a catalogue that expands beyond the redshift interval covered by the original survey, for use with upcoming surveys. Our evolution study presented here is an essential input for such efforts. The outline of this paper is as follows: in Section~\ref{Sec:GC} we introduce the SAMs used, along with the N-body simulation the models are grafted onto, and we describe the HOD characterisation of the galaxy population. In Section~\ref{Sec:HOD_Ev} we show the evolution of the HODs for a wide range of number densities and redshifts, we fit the HODs predicted by the SAMs and we show the evolution of the best-fitting parameters. In Section~\ref{Sec:PC2} we compare our results with simple models for the evolution of galaxy clustering. Finally in Section~\ref{Sec:Concl} we present our conclusions.

\label{Sec:Concl} The halo occupation distribution (HOD) framework has proven to be a useful theoretical tool to interpret galaxy clustering measurements and describe the relation between galaxies and dark matter haloes. Here we set to study how the halo occupation models evolve with time, an aspect that is missing from standard applications, using the outputs of semi-analytic models (SAMs) that capture the galaxy formation physics. It is important to recall that the SAMs predict the galaxy content of dark matter haloes along with the properties of these galaxies. The halo occupation functions are used here as a useful approach to characterise how the haloes are populated by galaxies in the SAMs. The halo occupation function has the attraction that it can be readily written in terms of the contribution from the main (e.g. most luminous or the galaxy from the most massive progenitor halo) or central galaxy, and satellite galaxies, which were once central galaxies in their host haloes but have subsequently merged with more massive dark matter haloes. Furthermore, the halo occupation function in itself is not dependent on the radial distribution of galaxies within haloes (though an assumption about this is required to predict the correlation function from the HOD). This is appealing for our purposes, as different SAMs handle the placement of galaxies within haloes in different ways (see the discussion in \citealt{C13} and \citealt{Campbell:2015}). The SAMs we consider use different implementations of the physical processes involved in galaxy formation and set the values of the model parameters in different ways, putting emphasis on different observables \citep{Henriques:2015, Lacey:2016}. We compare the model output at a series of number densities for galaxies ranked by their stellar mass. (Note that we also show the stellar mass functions so the reader can see how closely these agree with one another.) The HODs look remarkably similar until the samples characterised by the lowest number densities. In this case, the details of the suppression of gas cooling by heating by accretion onto active galactic nuclei become important and introduce differences in the HOD of central galaxies. The main aim for this study is to characterise the evolution of the HOD at a fixed number density, and we explore the evolution of the HOD best-fitting parameters over the redshift range $0 \le z \le 3$. As always, it is important to first assess which features of the SAMs are robust to the details of the implementation of the physics and the setting of the model parameters. Four out of the five parameters in the HOD parametrization that we used displayed remarkably similar behaviour. As before this similarity was strained when comparing the lowest density samples or the parameter which describes the transition from zero to one galaxy for the central HOD. Three of the HOD parameters are masses (see Fig.~1 for an illustration of how the parameters control the shape of the HOD). The evolution of the best-fitting values of these masses for samples of fixed number density is much weaker than the evolution in the characteristic halo mass (roughly speaking the mass at which there is a break from a power-law in the halo mass function). We found that the evolution is well described by a single parameter describing a power law in redshift and the $z=0$ value of the parameter. We also compared the evolution predicted in the SAMs to simplified evolution models that have been used to model galaxy clustering and evolution. These models make different assumptions about the fate of ``galaxies'' identified at some redshift. None of these models behave in the same way as the output of the SAMs, giving very different predictions for the evolution of the HOD parameters and the fraction of satellite galaxies in the sample. We find, in particular, that the ratio between the characteristic halo mass for hosting a satellite galaxy to that of hosting a central galaxy and its change with redshift can serve as a sensitive diagnostic for different galaxy formation and evolution scenarios. In so far as the models describe the clustering of stellar mass selected samples and its evolution, our results can be used to build mock catalogues for surveys from $z=0$ to $z=3$. Typically, an observational determination of the HOD may exist for one redshift, eg the low redshift results for $r$-band selected galaxies from \cite{Zehavi:2011}. The problem becomes how to extend these best-fitting parameters to other redshifts where there may not be an equivalent determination of the HOD parameters. For example, one might want to build a mock catalogue for the Euclid redshift survey, which will recover emission line galaxies over the redshift range $z \approx 0.5 - 2$ from a measurement of the clustering of H-alpha emitters at a different redshift (e.g. \citealt{Geach:2012}). We plan to pursue such efforts in future work.

1607.06154

1607

1607.03771_arXiv.txt

Evidence for the presence of quasi-periodic oscillations (QPOs) in the optical polarization of the blazar \pks, during a period of enhanced gamma-ray brightness, is presented. The periodogram of the polarized flux revealed the existence of a prominent peak at $T\sim 13$ min, detected at $>99.7$\% significance, and $T\sim 30$ min, which was nominally significant at $>99$\%. This is the first evidence of QPOs in the polarization of an active galactic nucleus, potentially opening up a new avenue of studying this phenomenon.

\pks{} is a blazar, a subclass of radio-loud active galactic nucleus (AGN) for which the relativistic jet is oriented close to the line-of-sight of the observer \citep{Urry2009}. Active galactic nuclei are powered by accretion of matter onto a $10^6-10^9$\msun{} supermassive black hole (SMBH). The observed emission is dominated by the jet and is characterized by intense and rapid brightness fluctuations across the electromagnetic spectrum, and the presence of polarization at optical to radio wavelengths. \pks{} exhibits significant variability on timescales ranging from minutes \citep{HESS2007} to years \citep{Fan2000, Kastendieck2011}. Short timescale variations, from minutes to hours, probe emission regions with sizes comparable to the gravitational radius of the central SMBH, while variations on timescales of months to years probe the jet structure. The timing signatures indicate that the emission is governed by correlated noise processes \citep{Vaughan2003, Edelson1999, Vaughan2005, Chatterjee2008, Chatterjee2012, Kastendieck2011}, possibly originating from instabilities in the accretion rate and jet. Quasi-periodic variations have been reported for a small number of AGN. The clearest detection is a $\sim 1$ hour QPO in the \xray{} light curve of the radio-quiet Seyfert galaxy RE J1034+396 \citep{GMWD2008}. \pks{} is one of only a few blazars for which convincing evidence of quasi-periodic brightness variations have been documented. Long time-scale periodicities were discovered in the optical light curve of the source. Amplitude modulations with periods of 4 and 7 years \citep{Fan2000}, and 315 days \citep{Zhang2014, Sandrinelli2014} were identified, while a short time-scale oscillation of \til 4.6 hours was observed at \xray{}energies \citep{Lachowicz2009}. Evidence of quasi-periodic oscillations in the optical polarization of \pks{} is presented here. The polarization of the source was monitored from 25 -- 27 July, 2009 with the HIgh Speed Photo-POlarimeter (HIPPO) of the Southern African Astronomical Society (SAAO). Gamma-ray observations were obtained independently with the High Energy Stereoscopic System (\hess) over the same period, with indications of two \gray{} flares which bracket the occurrence of the optical polarization fluctuations. The acquisition of the optical polarization measurements and a description of the polarization and \gray{} observations is given in section~\ref{sect:obs}. Section~\ref{sect:analysis} presents an analysis of the observations. Discussion of the results follows in section~\ref{sect:discussion}, while the conclusions are presented in section~\ref{sect:conclusions}.

\label{sect:conclusions} Monitoring of the intra-day optical polarization of \pks{} revealed the first evidence for quasi-periodic oscillations in the polarized flux of an AGN. The periodogram showed the existence of a periodic component at $T\sim 13$ min, detected at $>99.7$\% significance, and $T\sim 30$ min, which was found to be nominally significant at $>99$\%. The period of these oscillations are similar to the $T\approx 15$ min QPO seen in the optical light curve of the blazar S5 0716+714 \citep{Rani2010}, which provides further support for the existence of such short timescale periodicities in these sources. Gamma-ray observations of the source over the same period showed that it was in a high state of activity, experiencing two \gray{} flares at very high energies (VHE, photons with energies exceeding a few GeV). The first simultaneous optical polarization measurements of the source during a high-state was recorded for the latter flare. Although the physical cause of this QPO is unclear, comparison with the VHE light curve showed that \pks{} experienced a \gray{} flare two days before the appearance of the QPOs and another peak two days later. The oscillations could therefore be related either to the late phase of post flare activity of the first flare, which is supported by the decreasing trend observed over the night (as opposed to increasing trend seen on the following nights, see Fig.~\ref{fig:polvar}) or due to transition between two different gamma-ray flares. Since blazar emission is dominated by synchrotron emission at low energies, which also leads to polarization, and inverse Compton emission at high energies, which both arise from the relativistic jet, this could suggest that the observed QPOs are part of a longer-lived phenomenon within the jet of \pks.

1607.03771

1607

1607.00772_arXiv.txt

{Properties of magnetic field in the internetwork regions are still fairly unknown due to rather weak spectropolarimetric signals.} {We address the matter by using the 2D inversion code that is able to retrieve the information on smallest spatial scales, up to the diffraction limit, while being less susceptible to noise than most of the previous methods used.} {Performance of the code and the impact of the various effects on the retrieved field distribution is tested first on the realistic MHD simulations. The best inversion scenario is then applied to the real Hinode/SP data.} {Tests on simulations show: (1) the best choice of node position ensures a decent retrieval of all parameters, (2) code performs well for different configurations of magnetic field, (3) slightly different noise level or slightly different defocus included in the spatial PSF produces no significant effect on the results and (4) temporal integration shifts the field distribution to the stronger, more horizontally inclined field.} {Although the contribution of the weak field is slightly overestimated due to noise, the 2D inversions are able to recover well the overall distribution of the magnetic field strength. Application of the 2D inversion code on the Hinode/SP internetwork observations reveals a monotonic field strength distribution. The mean field strength at optical depth unity is $\sim 130$~G. At the higher layers, field strength drops as the field becomes more horizontal. Regarding the distribution of the field inclination, tests show that we cannot directly retrieve it with the observations/tools at hand, however the obtained distributions are consistent with those expected from simulations with a quasi-isotropic field inclination after accounting for observational effects.}

Determining the magnetic properties of the internetwork has always been an important task \citep{Almeida:Marian:2011}, because it carries a substantial fraction of the solar magnetic flux and has a large impact on the energy budget in the solar atmosphere. Since decades, it has been clear that the internetwork is filled with magnetic field elements of opposite polarities organized on small-scales \citep{Livingston:Harvey:1975,Livi:etal:1985,Martin:1988}. Yet, whether their field strength is predominantly hG or kG, i.e. how much flux they carry, has been the subject of debate \citep{Keller:etal:1994,Lin:1995,Almeida:Lites:2000,Socas:Lites:2004,Cerdena:2006,Marian:2006}. As the polarimetric sensitivity and spatial resolution of the observations increased, it became possible to characterize the full magnetic field vector. The debate then further expanded on the inclination of the magnetic field. Currently, three hypothesis remain: that internetwork magnetic field is predominantly horizontal \citep{david07a,david07b,lites08}, predominantly vertical \citep{reza2009,stenflo10,ryuko11} or quasi-isotropic \citep{asensio2009,asensio2014}. The basic problem lies in retrieving information on magnetic field based on the rather weak internetwork spectropolarimetric signals. The influence of noise on the result is overwhelming and leads to a systematic overestimation of the inclination of the magnetic field vector \citep{Borrero:Kobel:2011}. One way of getting round that problem is to limit the analysis only to pixels where the signal to noise ratio is high enough. However, these selection criteria tend to exclude a significant portion of the internetwork surface and to bias the retrieved distributions of magnetic field strength and inclination in different ways \citep{Borrero:Kobel:2012}. In this paper, we address the issue by using the 2D inversion technique \citep{Michiel2012} that accounts for spatial coupling between the neighbouring pixels and simultaneously and self-consistently fits the observed spectra, which makes it less susceptible to noise than most of the previous methods used. The code is first tested on 'synthesized observations' produced from realistic magnetohydrodynamic simulations and then the best inversion strategy is applied to real observations. We limit our study to the general properties of the distributions of the magnetic field strength and inclination.

An inversion code that self-consistently accounts for the effects that the instrumental PSF has on data, was applied, for the first time, to quiet Sun observations. Extensive testing on MHD simulations confirmed that the inverted results make sense and also provided an inversion strategy which we then applied to the solar observations. In studies like the one presented here, typically only pixels are selected that contain significant signals in Q,U and/or V, to ensure that the result returned by the inversion code used to recover the magnetic field strength is reliable. However, this selection disregards signals that in individual pixels are below the noise level, but averaged over a large number of pixels are statistically significant. This not only results in loss of signal, it may also introduce a bias in the results towards the properties that are particular of strong magnetic fields only. The spatially coupled inversion method used here is able to constrain a result using such signals, and can therefore make use of {\em all} pixels in the FOV. The inversions return a distribution with mainly weak field, without any secondary peak at kG field strengths, as in \cite{stenflo10} and \cite{lites2011}. It shows no peak either at hG values as detected by \cite{david07a}, but monotonously increases towards the smallest field strengths. This however does not exclude the possibility that real distribution of the field strength does not have a peak at $5-10$~G as local dynamo simulations show. The tests on the simulations demonstrate that the code tends to set the field strength to zero when signals are too weak. The code also tends to retrieve mostly horizontal field in the regions which harbour very weak field. This will then produce large differences between the original and retrieved distribution of the field inclinations in the case of local dynamo simulations which show salt and pepper pattern even in these regions. Due to the noise, the code tends to overestimate the hG field which results in a slight overestimation of the mean field strength. Nevertheless the retrieved value comes close to the original. The mean field strength $>100$~G at optical depth unity retrieved from the observations, puts ours results closer to the results based on the Hanle effect \citep{trujillo04}, although we cannot confirm that the field strength in the upper photosphere is also over 100~G, as found by \cite{Shchukina2011}. A mean field strength at the solar surface of the same magnitude was also retrieved by \cite{david2012a} and \cite{Luis2012}. Their results, however, show much more horizontal field. Our tests show that this can be explained as an artefact, produced by prolonged temporal averaging, since increasing the integration time beyond the evolution time scale of the solar scene (2-3~min.), artificially increases the apparent contribution of the horizontal fields significantly. In the case that the distribution of magnetic field is isotropic, this can be easily understood, since in that case integration beyond the evolution timescale will superpose a statistically independent realization of the magnetic field distribution, which will decrease the measured polarimetric signal at the same rate as the photon noise, so that no nett improvement of the S/N ratio can be obtained by continued integration. The scaling properties of the noise equivalent horizontal and magnetic fields, however, continue to favour a more and more horizontally inclined field configuration. At the same time, the granular motions produce wider line profiles which results in larger retrieved field strengths. The distribution of the field inclination has a maximum at $90^{o}$, which confirms the results given by \cite{david07b} and \cite{lites08}. However, we cannot interpret this as evidence for a predominantly horizontal field, nor can we state that this is in agreement with the results obtained by \cite{asensio2014}, who claim that the distribution is quasi-isotropic. As shown in Fig.~\ref{hist_diffsim}, similar results can be obtained from both predominantly horizontal and a quasi-isotropic distribution of the magnetic field. Re-examining the lower left panel of Fig.~\ref{hist_diffsim}, one can conclude that even if the photospheric magnetic field is isotropic, it might not be possible to recover it as such. We are currently limited, not only by our inversion tools, but also by the resolution limit of our instruments. Tests of the 2D inversion technique on the simulated Hinode/SP data at the disc center suggest that the information for discerning between the two distributions is just not there.

1607.00772

1607

1607.07687_arXiv.txt

We show that the widely used density dependent magnetic field prescriptions, necessary to account for the variation of the field intensity from the crust to the core of neutron stars violate one of the Maxwell equations. We estimate how strong the violation is when different equations of state are used and check for which cases the pathological problem can be cured.

\paragraph*{} The physics underlying the quantum chromodynamics (QCD) phase diagram has still not been probed at all temperatures and densities. While some aspects can be confirmed either by lattice QCD or experimentally in heavy ion collisions for instance, other aspects depend on extraterrestrial information. One of them is the possible constitution and nature of neutron stars (NS), which are compact objects related to the low temperature and very high density portion of the QCD phase diagram. Astronomers and astrophysicists can provide a handful of information obtained from observations and infer some macroscopic properties, namely NS masses, radii, rotation period and external magnetic fields, which have been guiding the theory involving microscopic equations of state (EOS) aimed to describe this specific region of the QCD phase diagram. \paragraph*{} There are different classes of NS and three of them have shown to be compatible with highly magnetised compact objects, known as magnetars \cite{Duncan, Usov}, namely, the soft gamma-ray repeaters, the anomalous X-ray pulsars and more recently, the repeating fast radio burst \cite{frb}. The quest towards explaining these NS with strong surface magnetic fields, has led to a prescription that certainly violates Maxwell equations \cite{Chakrabarty}. The aim of this letter is to show how strong this violation is and check whether the density dependent magnetic field prescriptions can or cannot be justified. Magnetars are likely to bear magnetic fields of the order of $10^{15}$ G on their surfaces, which are three orders of magnitude larger than magnetic fields in standard NS. In the last years, many papers dedicated to the study of these objects have shown that the EOS are only sensitive to magnetic fields as large as $10^{18}$ G or stronger \cite{originalB,originalNS}. The Virial theorem and the fact that some NS can be quark (also known as strange) stars allow these objects to support central magnetic fields as high as $3 \times 10^{18}$ G if they contain hadronic constituents and up to $10^{19}$ G if they are quark stars. To take into account these possibly varying magnetic field strength that increases towards the centre of the stars, the following proposition was made \cite{Chakrabarty}: \begin{equation} B_z(n) = B^{surf} + B_0\bigg [ 1 - \exp \bigg \{ - \beta \bigg ( \frac{n}{n_0} \bigg )^{\gamma} \bigg \} \bigg ], \label{brho} \end{equation} where $B^{surf}$ is the magnetic field on the surface of the neutron stars taken as $10^{14}G$ in the original paper, $n$ is the total number density, and $n_0$ is the nuclear saturation density In subsequent papers \cite{Mao,Rabhi,Menezes1,Ryu,Rabhi2,Mallick,Lopes1,Dex,Benito1,Benito2,Mallick2,Ro,Dex2}, the above prescription was extensively used, with many variations in the values of $B^{surf}$, generally taken as $10^{15}$G on the surface, $\beta$ and $\gamma$, arbitrary parameters that cannot be tested by astronomical observations. The high degree of arbitrariness was checked already in \cite{Chakrabarty} and later in \cite{Lopes2015} and more than 50\% variation in the maximum stellar mass and 25\% variation in the corresponding radius was found. \paragraph*{} With the purpose of reducing the number of free parameters from two to one and consequently the arbitrariness in the results, another prescription was then proposed in \cite{Lopes2015}: \begin{equation} B_z(\epsilon) = B^{surf} + B_0 \bigg ({\frac{\epsilon}{\epsilon_c}} \bigg )^{\alpha}, \label{bepsilon} \end{equation} where $\epsilon_c$ is the energy density at the centrer of the maximum mass neutron star with zero magnetic field, $\alpha$ is any positive number and $B_0$ is the fixed value of the magnetic field. With this recipe, all magnetic fields converge to a certain value at some large energy density, despite the $\alpha$ value used. \paragraph*{} One of the Maxwell equations tells us that $\nabla\cdot\vec B=0$. In most neutron star calculations, the magnetic field is chosen as static and constant in the $z$ direction, as proposed in the first application to quark matter \cite{originalB}. In this case, the energy associated with the circular motion in the $x-y$ plane is quantised (in units of $2qB$, $q$ being the electric charge) and the energy along $z$ is continuous. The desired EOS is then obtained and the energy density, pressure and number density depend on the filling of the Landau levels. Apart from the original works \cite{originalB, originalNS}, detailed calculations for different models can also be found, for instance, in \cite{Menezes1,Benito1} and we will not enter into details here. However, it is important to stress that the magnetic field is taken as constant in the $z$ direction, what results in different contributions in the $r$ and $\theta$ directions when one calculates the magnetic field in spherical coordinates. \paragraph*{} According to \cite{Goldreich,Danielle}, and assuming a perfectly conducting neutron star ($B_r(R)=0$) that bears a magnetic dipole moment aligned with the rotation axis such that $\mu=B_p R^3/2$, where $R$ is the radius of the star and $B_p$ the magnetic field intensity at the pole, the components of the magnetic field in spherical coordinates are given by: \begin{equation} B_r=B_P \cos \theta \left(\frac{R}{r}\right)^3, \quad B_\theta=\frac{B_P}{2} \sin \theta \left(\frac{R}{r}\right)^3, \label{daniele} \end{equation} and in this case, it is straightforward to show that $\nabla \cdot \vec B=0$. \paragraph*{} If one cast eqs. (\ref{brho}) and (\ref{bepsilon}) in spherical coordinates, from now on called respectively original and LL's prescriptions, they acquire the form $B_r=\cos \theta B_z$, $B_\theta=-\sin \theta B_z$ and $B_\phi=0$ and the resulting divergent reads \begin{equation} \nabla \cdot \vec B= \cos \theta \frac{\partial B}{\partial r}, \label{divB} \end{equation} where the magnetic field in the radial direction can be obtained from the solution of the TOV equations \cite{tov}, where $r$ runs from the centre to the radius of the star. As a simple conclusion, $\nabla\cdot\vec B$ is generally not zero, except for some specific values of the parameters that we discuss in the next Section.

\paragraph*{} In what follows we analyse how much $\nabla\cdot\vec B$ deviates from zero when $B$ is allowed to vary either with the original prescription as in eq. (\ref{brho}) or with LL's proposal, as in eq. (\ref{bepsilon}) with two different models, the NJL and the MIT bag model. These two models have been extensively used to describe stellar matter in the interior of quark stars. It is important to point out that the same test could be performed with hadronic models used to account for magnetised NS with hadronic constituents, as in \cite{ Mao,Rabhi,Ryu,Rabhi2,Mallick,Lopes1,Dex,Benito1,Mallick2,Ro}, for instance. \paragraph*{} Before we analyse the behaviour of $\nabla\cdot\vec B /B$ that depends on $\frac{\partial B}{\partial r} = \frac{\partial B}{\partial n} \frac{\partial n}{\partial r}$, we show in Fig. \ref{fig0} how the magnetic field varies with the star radius for one specific case, i.e., the MIT bag model with $B=10^{18}$ G and the prescription given in eq.(\ref{bepsilon}). All other cases studied next present a very similar behaviour. It is interesting to notice that the curve changes concavity around half the stellar radius. \begin{figure}[ht] \includegraphics[scale=0.40]{b-r.eps} \caption{ Magnetic field versus r for the MIT model, bag constant 148 MeV$^{1/4}$, $M_{max}=1.4~M_\odot$ and $R=10.04$ km.} \label{fig0} \end{figure} \paragraph*{} In Fig. \ref{fig1} we plot $\nabla\cdot\vec B /B$ as a function of the star radius for different latitude ($\theta$) angles and a magnetic field equal to $B_0=10^{18} G$. In both cases the equations of state were obtained with the Nambu-Jona-Lasinio model \cite{Menezes1, Luiz2016}. The violation is quite strong and $\nabla\cdot\vec B /B$ can reach 70\% for small angles. \begin{figure}[ht] \includegraphics[scale=0.44]{NJL-c-LL-1e18.eps} \caption{Equations of state obtained with the NJL model and $B_0=10^{18}$ G. Chakrabarty's prescription was calculated with $M_{max}=1.44~M_\odot$ and $R=8.88$ km, $\beta=5 \times 10^{-4}$ and $\gamma=3$. LL's prescription was calculated with $M_{max}=1.46~M_\odot$, $R=8.83$ km and $\zeta=3$.} \label{fig1} \end{figure} \paragraph*{} In Figure \ref{fig2} we plot the same quantity as in Figure \ref{fig1} for the LL's prescription and the much simpler and also more used MIT bag model for different latitude angles. Again the violation amounts to the same values as the ones obtained within the NJL model. Finally, in Figure \ref{fig3}, we show how large the deviation can be for different values of the magnetic field intensity and a fixed angle $\theta=45$ degrees and the original prescription. We see that the deviation reaches approximately the same percentages, independently of the field intensity. \begin{figure}[ht] \includegraphics[scale=0.39]{mit1e18a.eps} \caption{Equations of state obtained with the MIT model, bag constant 148 MeV$^{1/4}$, $M_{max}=1.4~M_\odot$ and $R=10.04$ km for different latitude angles. } \label{fig2} \end{figure} \begin{figure}[ht] \includegraphics[scale=0.5]{NJL-c-45.eps} \caption{ Quark stars described by NJL model with different values of $B_0$ and $\theta=45$ degrees. For $B_0=10^{18}$ G, $M_{max}=1.44_\odot$, $R=8.88$ km and central energy density $\epsilon_c=7.67~fm^{-4}$. For $B_0=3\times10^{18}$ G: $M_{max}=1.45_\odot$, $R=8.88$ km and central energy density $\epsilon_c=7.65~fm^{-4}$. For $B_0=10^{19} G$: $M_{max}=1.50_\odot$, $R=8.80$ km and central energy density $\epsilon_c=8.11~fm^{-4}$. } \label{fig3} \end{figure} \paragraph*{} Now, we turn our attention to a possible generalisation of eqs. (\ref{brho}) and (\ref{bepsilon}) in spherical coordinates in order to verify for which conditions the divergent becomes zero and check whether the situation can be circumvented. The generalized magnetic field components is given by \begin{equation} B_r= B_0 \cos \theta \left( f(r)\right)^\eta, \quad B_\theta=-\frac{B_0}{\zeta} \sin \theta \left( f(r)\right)^\eta, \end{equation} and the corresponding divergent reads: \begin{equation} \nabla \cdot \vec B= B_0 \cos \theta \left[ \frac{2 f(r)^{\eta}}{r} + \eta f(r)^{\eta-1} \frac{df}{dr} - \frac{2 f(r)^{\eta}}{r \zeta} \right], \label{divBLL} \end{equation} which is zero either for the trivial solution $\cos \theta=0$ or if $$\frac{2}{r} + \frac{\eta}{f(r)} \frac{df}{dr} - \frac{2}{r \zeta}=0, $$ for which a general solution has the form $f(r)=A r^{\frac{2-2 \zeta}{\eta \zeta}}$ . When $\zeta=-2$ and $\eta=3$, $A=R$ and $B_0=B_p$, eqs.(\ref{daniele}) are recovered. When $\zeta=1$, the numerator of the exponent becomes zero and $f(r)$ is simply a constant ($A$), does not depending on $r$. Another possibility is the assumption that $f(r)$ is a function of the density, for instance, as $f(n(r))$ or of the energy density as $f(\varepsilon(r))$. In these cases, $$ \frac{df}{d n} \frac{d n}{dr} = \frac{2~f}{r~\eta} \left( \frac{1-\zeta}{\zeta}\right), \quad \frac{df}{d \varepsilon} \frac{d \varepsilon}{dr} = \frac{2~f}{r~\eta} \left( \frac{1-\zeta}{\zeta}\right).$$ If we take $\zeta=1$, the result resembles the original (LL's) prescription, $ \frac{df}{d n} =0$ ($\frac{df}{d \varepsilon} =0$) because $\frac{d n}{dr}$ ($\frac{d \varepsilon}{dr}$) is obtained from the TOV equations and is never zero. For $\zeta \ne 1$, solutions can be obtained from numerical integration. \\ \\

1607.07687

1607

1607.02541_arXiv.txt

{} {Galaxy mass models based on simple and analytical functions for the density and potential pairs have been widely proposed in the literature. Disk models constrained by kinematic data alone give information on the global disk structure only very near the Galactic plane. We attempt to circumvent this issue by constructing disk mass models whose three-dimensional structures are constrained by a recent Galactic star counts model in the near-infrared and also by observations of the hydrogen distribution in the disk. Our main aim is to provide models for the gravitational potential of the Galaxy that are fully analytical but also with a more realistic description of the density distribution in the disk component.} {From the disk model directly based on the observations (here divided into the thin and thick stellar disks and the H\,{\scriptsize I} and H$_2$ disks subcomponents), we produce fitted mass models by combining three Miyamoto-Nagai disk profiles of any ``model order'' (1, 2, or 3) for each disk subcomponent. The Miyamoto-Nagai disks are combined with models for the bulge and ``dark halo'' components and the total set of parameters is adjusted by observational kinematic constraints. A model which includes a ring density structure in the disk, beyond the solar Galactic radius, is also investigated.} {The Galactic mass models return very good matches to the imposed observational constraints. In particular, the model with the ring density structure provides a greater contribution of the disk to the rotational support inside the solar circle. The gravitational potential models and their associated force-fields are described in analytically closed forms.} {The simple and analytical models for the mass distribution in the Milky Way and their associated three-dimensional gravitational potential are able to reproduce the observed kinematic constraints, and in addition, they are also compatible with our best knowledge of the stellar and gas distributions in the disk component. The gravitational potential models are suited for investigations of orbits in the Galactic disk.}

\label{intro} Reliable models for the gravitational potential of the Galaxy are mandatory when studies of the structure and evolution of the Galactic mass components rely upon the characteristics of the orbits of their stellar content. In this sense, Galaxy mass models are regarded as the simplest way of assessing and understanding the global structure of the main Galactic components, providing a great insight into their mass distribution once a good agreement between the model predictions and the observations is obtained. A pioneer Galactic mass model was that of \citet{Schmidt1956}, contemporary of the early years of the development of radio astronomy and the first studies of the large-scale structure of the Milky Way. With the subsequent improvement of the observational data, updated mass models have been undertaken by several authors (e.g., \citealt[among others]{Bahcall_Soneira1980,Caldwell_Ostriker1981,Rohlfs_Kreitschmann1988}), and with the advent of the Hipparcos mission and large-scale surveys in the optical and near-infrared, new observational constraints have been adopted in the more recent Galaxy mass models (e.g., \citealt{Dehnen_Binney1998,Lepine_Leroy2000,Robin2003,Polido2013}). In order to evaluate the capability of a given mass model in reproducing some observables, the force-field associated with the resulting gravitational potential has to be compared with available dynamical constraints such as the radial force in the plane given by the rotation curve, as well as the force perpendicular to the plane of the disk along a given range of Galactic radii. Regarding the latter one, the associated mass-surface density, to our better knowledge, is the one integrated up to the height of 1.1 kpc of the Galactic mid-plane (\citealt{Kuijken_Gilmore1991,Bovy_Rix2013}), and as pointed out by \citet{Binney_Merrifield1998}, this constraint is not able to provide much information about the mass distribution some kiloparsecs above the plane. Due to these shortcomings, a degeneracy in the set of best models is observed, which means that different mass models are able to reproduce the kinematic information of the observed data equally well. As stated by \citet{McMillan2011}, one possible way of circumventing such obstacles is by combining the kinematic data with star counts to improve the Galactic potential models and its force-field above the plane. Regarding the use of a Galactic potential model for the purpose of orbit calculations, one which has been widely adopted is that of \citet{Allen_Santillan1991}. Such model has the attractive characteristics of being mathematically simple and completely analytical, with closed forms for the potential and density, assuring both fast and accurate orbit calculations. \citet{Irrgang2013} have recalibrated the \citet{Allen_Santillan1991} model parameters using new and improved observational constraints. The main goal of the present work is to provide a fully-analytical, three-dimensional description of the gravitational potential of the Galaxy, but with the novelty of expending considerable efforts in a detailed modelling of the disk component. The basic new aspects of the present Galactic mass model, all of which related to the disk modelling process, can be summarized in the following way: \begin{itemize} \item the structural parameters of the disk - scale-length, scale-height, radial scale of the central `hole' - are based on the Galactic star counts model in the near-infrared developed by \citet*[hereafter PJL]{Polido2013} for the case of the stellar disk component; for the gaseous disk counterpart, we adopt recent values returned by surveys of the distribution of hydrogen, atomic H\,{\scriptsize I} and molecular H$_2$, in the disk; \item the density and potential of the disk components are modelled by the commonly used Miyamoto-Nagai disk profiles (equations 4 and 5 of \citealt{Miyamoto_Nagai1975}), but here we also attempt to make use of the higher ``model orders'' 2 and 3 of Miyamoto-Nagai disks (equations 6, 7, 8 and 9 of the above-referred paper) in order to better fit some of the disk subcomponents. The approach followed for the construction of the Miyamoto-Nagai disks is based on the one presented by \citet{Smith2015}; \item a model with a ring density structure added to the disk density profile is studied, with the ring feature placed somewhat beyond the solar Galactic radius. The inclusion of such ring structure is motivated by the attempt of modelling the local dip in the observed Galactic rotation curve also placed a little beyond the solar orbit radius. An explanation for the existence of such ring density structure is given by \citet*[hereafter BLJ]{Barros2013}. \end{itemize} The organization of this paper is as follows: in Sect.~\ref{MW_disk_mass_models}, we present the details of the mass models of the Galactic disk and the steps through the construction of Miyamoto-Nagai disks versions of the `observed' ones. In Sect.~\ref{MW_models}, we give the expressions for the bulge and dark halo components, as well as the functional form for the gravitational potential associated with the ring density structure. The group of observational constraints adopted for the fitting of the models are presented in Sect.~\ref{obs_constr}, while the fitting scheme and the estimation of uncertainties are presented in Sect.~\ref{fit_proc}. In Sect.~\ref{result_discuss}, we analyse the results of each mass model by a direct comparison with other models in the literature. Concluding remarks are drawn in the closing Sect.~\ref{conclusions}.

\label{conclusions} We have developed models for the axisymmetric mass distribution of the Galaxy with the aim to derive fully-analytical descriptions of its associated three-dimensional gravitational potential. We have followed an approach which intentionally expends more efforts in a detailed modelling of the disk component. Based on photometric constraints for the stellar distribution in the disk given by the Galactic infrared star counts model of PJL, as well as on the observed distribution of atomic and molecular hydrogen gas, we derived an ``empirical basis" for the structural parameters (scale-length, scale-height, and radial scale of the disk central hole) of the thin and thick stellar disks and the H\,{\scriptsize I} and H$_2$ disks subcomponents. With {\it a priori} values for the masses of each disk, based on the most recent determinations of the local stellar and gaseous disk surface densities, we have constructed versions of Miyamoto-Nagai disks for each disk subcomponent mass model. The method follows the approach developed by \citet{Smith2015}, but here with the allowance of using the Miyamoto-Nagai disk models of higher orders 2 and 3 beside the commonly used one of order 1 (\citealt{Miyamoto_Nagai1975}). Along with parametric models for the bulge and an extra unknown spherical mass component to which for ``convenience'' we refer by the often-used term {\it dark halo}, we searched for the dynamical mass of each Galactic component by fitting the models to the kinematic constraints given by the observed rotation curve and some local Galactic measured properties. We have shown that a disk model which includes a ring density pattern beyond but very close to the solar orbit radius is able to better reproduce an observed local dip in the Galactic rotation curve centered at $R\sim 9.0$ kpc; such dip is naturally explained by a ring density structure composed by a minimum followed by a maximum density of similar amplitudes. Furthermore, the model with the ring structure allows a more massive disk to increasingly contribute to the rotational support of the Galaxy inside the solar circle, helping the Milky Way to satisfy the condition to be considered a ``maximal disk'' galaxy. The model with the ring structure also relies on a numerical study of the stellar disk where a ring density structure develops as a consequence of interactions between the stars and the galactic spiral arms in resonance at the co-rotation radius (BLJ). Since we still have no information about the three-dimensional structure of this ring density feature, we had to make crude approximations for the vertical profile of its associated gravitational potential. We believe that, with the forthcoming data of the GAIA mission (\citealt{Perryman2001}), the ring structure in the disk may become an object of examination, and in the case of confirmation of its existence, constraints on its three-dimensional distribution in the global disk structure will help us to create more realistic disk models. The method we have applied for the construction of the disk mass model and its associated gravitational potential is quite flexible in the sense that for any set of structural parameters, disks of different masses can be generated. We emphasize at this point that the models of Galactic gravitational potential presented in this work are aimed for being used in studies of orbits in the disk that do not extend too far away in Galactic radii as well as do not reach great heights above the disk mid-plane. These limitations are imposed by the spatial coverage of the observational data used to constrain the models, in the sense that for radii greater than $\sim 2R_{0}$ (the maximum radius of the data used for the rotation curve) and heights $|z|\gtrsim 3$ kpc, we do not guarantee that our models return confident representations of the Galactic potential, but maybe reasonable ones at least. Also in this respect, no constraint on the mass at large radii is adopted here, as has been made by some studies that derive properties of the dark halo by requiring distant halo stars to be bound to the Milky Way potential. We have included a density component associated with a spherical logarithmic potential just to explain the observed rotation curve data at radii $R_{0}\lesssim R \lesssim 2R_{0}$. However, as pointed out by \citet{Dehnen_Binney1998}, instead of a distinct physical component, it is possible that we are measuring the dynamical effects of a disk and/or a bulge in which the mass-to-light ratio of their content strongly increases from the center to the Galactic outskirts. The models for the three-dimensional gravitational potential of the Galaxy presented in this work, being fully analytical and easy to obtain the associated gravitational force-field at any point, are suited for fast and accurate calculations of orbits of stellar-like objects belonging to the main populations of the Galactic disk. In a forthcoming study, we aim to present an application of these new potential models to a description of the distributions of orbital parameters for samples of Galactic open clusters and some expected correlations with their chemical abundance patterns.

1607.02541

1607

1607.01074.txt

The anisotropic galaxy clustering of large scale structure observed by the Baryon Oscillation Spectroscopic Survey Data Release 11 is analyzed to probe the sum of neutrino mass in the small $m_\nu\la 1\eV$ limit in which the early broadband shape determined before the last scattering surface is immune from the variation of $m_\nu$. The signature of $m_\nu$ is imprinted on the altered shape of the power spectrum at later epoch, which provides an opportunity to access the non--trivial $m_\nu$ through the measured anisotropic correlation function in redshift space (hereafter RSD instead of Redshift Space Distortion). The non--linear RSD corrections with massive neutrinos in the quasi linear regime are approximately estimated using one-loop order terms computed by tomographic linear solutions. We suggest a new approach to probe $m_\nu$ simultaneously with all other distance measures and coherent growth functions, exploiting this deformation of the early broadband shape of the spectrum at later epoch. If the origin of cosmic acceleration is unknown, $m_\nu$ is poorly determined after marginalising over all other observables. However, we find that the measured distances and coherent growth functions are minimally affected by the presence of mild neutrino mass. Although the standard model of cosmic acceleration is assumed to be the cosmological constant, the constraint on $m_\nu$ is little improved. Interestingly, the measured CMB distance to the last scattering surface sharply slices the degeneracy between the matter content and $m_\nu$, and the hidden $m_\nu$ is excavated to be $m_\nu=0.19^{+0.28}_{-0.17} \eV$ which is different from massless neutrino more than 68\% confidence.

A neutrino is an elementary particle in the Standard Model (hereafter SM) of particle physics with three flavours whose existence was suggested as a ``desperate remedy" to explain the violation of energy conservation in nuclear beta decay~\cite{2007ConPh..48..195K}. The theoretically predicted particle was discovered~\cite{1956Sci...124..103C}, and became a fundamental building block of SM of particle physics. The detection of neutrino flavour oscillations also provides the indirect evidence of a non--trivial neutrino mass through the solar neutrino experiments~\cite{2004PhRvL..92r1301A, 2003PhRvL..90b1802E, 2005ConPh..46....1M}. The detection of possible neutrino mass constrains the lower bound on the sum of neutrino mass as $\sum m_{\nu}\ga 0.06\eV$ and $\sum m_{\nu}\ga 0.1\eV$ while assuming a normal mass hierarchy and an inverted hierarchy respectively~\cite{2016arXiv160503159C}. The tightest neutrino mass upper bound from laboratory has been reported to $\la 2\eV$~\cite{2002NuPhS.110..395B}. While this constraint bounded by particle physics experiments is expected to be improved by a few orders of magnitude by seeking for neutrinoless double beta decay events in the future~\cite{2002RvMP...74..663Z}, the more stringent upper neutrino mass bound is imposed through cosmological observations \cite{2006JCAP...10..014S, 2009ApJ...692.1060V, 2009A&A...500..657T, 2010JCAP...01..003R, 2010PhRvL.105c1301T, 2010MNRAS.406.1805M, 2010MNRAS.409.1100S, 2011ApJS..192...18K, 2014MNRAS.444.3501B, 2016PDU....13...77C}. Alternatively, the neutrino mass can be probed by cosmological observations through the distinct clustering caused by the neutrino damping effect. The effect of massive neutrinos is imprinted on the recombination history through the alternated expansion history, which influences the shape of spectra determined at the last scattering surface. However, if $\sum m_{\nu}\la 1\eV$, the cosmic neutrinos become non--relativistic after the last scattering surface, and the transfer function with all massless neutrinos remains unchanged~\cite{2014PhRvD..89j3541S}. Light massive neutrinos are detectable by the Cosmic Microwave Background (hereafter CMB) experiment through the early Integrated Sachs--Wolfe (ISW) effect due to their being less relativistic around the last scattering surface and through the gravitational lensing effect developed at later epoch~\cite{2003PhRvL..91x1301K}. The Planck experiment constrains the neutrino mass upper bound as $\sum m_{\nu}\la 0.68\eV$ at 95\% confidence level~\cite{2015arXiv150201589P}. The signature of neutrino mass is also imprinted on the large scale structure of the universe which is traced by galaxy distribution~\cite{1998PhRvL..80.5255H}. The observed galaxy clustering is plagued by uncertainties due to the non--linear mapping from the real to redshift spaces~\cite{Kaiser:1900zz, 1998ASSL..231..185H, 2012ApJ...748...78K, 2013PhRvD..88j3510Z}. This mapping is intrinsically non--Gaussian which leads to the infinite tower of cross--correlation pairs between density and velocity fields, and the non--perturbative suppression due to the randomness of infalling velocity is not theoretical predictable~\cite{2010PhRvD..82f3522T, 2016arXiv160300101Z}. Those all non--trivial corrections are not easily formulated with the cosmological model in which the neutrino becomes non--relativistic. But, if the targeted range of scale remains in the quasi linear regime, the known linear solutions can be fed tomographically at each epoch where the large scale structure evolves non--linearly. This approximation is proved to be applicable in our interesting scales~\cite{2016PhRvD..93f3515U}. Constraint on the neutrino mass is studied using the Baryon Oscillation Spectroscopic Survey Data Release 11 (hereafter BOSS DR 11) catalogue in this manuscript. Unlike the more conventional methodology to probe $m_\nu$ using galaxy clustering data, we suggest a new approach. The early broadband shape of the power spectrum determined at last scattering surface is altered when the neutrinos become non--relativistic~\cite{2011ARNPS..61...69W}. This shape departure is a unique signature of the non--trivial $m_\nu$, and it needs to be disclosed simultaneously with all other distance measures and coherent growth functions. When we apply the minimal theoretical prior for cosmic acceleration physics in which the structure evolves coherently after the last scattering surface, the $m_\nu$ constraint becomes weak and no confirming upper bound is discovered at small $m_\nu\la 1\eV$ limit. Although we impose the $\Lambda$CDM prior, the $m_\nu$ constraint little improves. With the $\Lambda$CDM prior, the CMB distance measure can be combined. This combination breaks the degeneracy between the matter content $\Omega_m$ and the neutrino mass $m_\nu$ significantly, and the massless limit of $m_\nu$ becomes distinguishable. The $m_\nu$ is measured to be $m_\nu=0.19^{+0.28}_{-0.17} \eV$ with the assumption of Gaussian probability distribution. We present the details in the following sections.

%We have introduced a massive neutrino of mass ranging between 0.0 eV and 1.0 eV into our RSD analysis to study its effect on the large-scale structure in the Universe. %First, we constrained distances, growth functions, velocity dispersion, and the neutrino mass in a model-independent way. %Compared with the previous study without massive neutrino, the constraint on $(D_A, H^{-1}, G_\Theta)$ is indistinguishable, although the measurements of the astrophysical parameters $\sigma_z$ and $G_b$ slightly changed due to the massive neutrino. %This result indicates that our study with massive neutrino is consistent with the previous one, which implies that some proper bound on total neutrino mass is not provided in the model-independent case. The early broadband shape of the power spectrum is altered by the damping effect due to the massive neutrino with $m_\nu\la 1\eV$ at later epoch. The observed galaxy clustering in redshift space can be exploited to observe this signature imprinted by non--trivial neutrino mass. We provide the methodology to detect $m_\nu$ with probing this altered broadband shape of power spectrum. This damping effect can be disclosed simultaneously with distance measures and coherent growth factors which are dependent on the unknown dark energy models. When the origins of cosmic acceleration is completely unknown, i.e. neither distance measures nor coherent growth functions are known, the neutrino mass is not measured in precision using BOSS DR11 catalogues. But we find that both distance measures and growth functions are minimally affected by $m_\nu$ marginalisation. It justifies our previous results without including $m_\nu$ damping effect. Although there is no confirming evidence of cosmological constant, the most reliable and simplest theoretical model to explain the cosmic acceleration is $\Lambda$CDM model. We continue our test with a theoretical prior that the cosmic expansion is accelerated by the existence of cosmological constant. Then all distance measures and growth functions vary coherently with one single cosmological paramter $\Omega_m$. Despite this reduction of parameter space, the neutrino mass is not clearly distinguishable from the massless limit, and the upper bound is poorly set by $m_\nu\la 0.8\eV$. However, the hidden $m_\nu$ can be revealed by ``smoking gun" from the combination of CMB distance measure at the last scattering surface. The angular extension of sound horizon constrain the $\Omega_m$ and $m_\nu$ correlation pattern. The likelihood plane of ($\Omega_m,m_\nu$) probed by BOSS DR11 is sharply sliced by the determined trajectory of $\Omega_m$ and $m_\nu$ correlation by CMB distance measure, presented in Fig.~\ref{fig:result2}. This combination excavates the hidden neutrino mass as $m_{\nu} = 0.19^{ +0.28}_{ -0.17}\eV$. The statistical error is reported with assuming Gaussianity of probability distribution. Although this assumption is not made, both likelihood function and $\Delta \chi^2$ results in Fig.~\ref{fig:result2} confirms the fact that $m_\nu$ is probed to be different from massless limit, which was also reported in \cite{2014MNRAS.444.3501B} using different approach. While we focus on probing neutrino mass through galaxy clustering observed by BOSS DR11 with the assistance of CMB distance measures, CMB also contains an alternative clustering information imprinted by lensing effect. The signature of neutrino mass is also preserved in the distorted CMB anisotropy maps at smaller scales. Thus, instead of taking only CMB distance measure constraint, the full combination between BOSS DR11 and CMB measurements is an interesting topic to be investigated. Considering the targeted volume of BOSS DR11, the observed clustering by BOSS DR11 affects the measured CMB lensing as well about 20\%. We would like to address the full covariance approach in the following work, we plan to probe $m_\nu$ using not only galaxy clustering but also CMB lensing effect. %\minji{As -$>$ Regarding} our study's limit and future direction at the same time, we would like to list up two points. %In our research, we combine only peak location information in CMB data with the galaxy clustering data. %In principle, however, information from all scales in CMB should be utilized, especially small scales where the damping effect from the massive neutrino is more dominant, to be analyzed with BOSS data together. %To make this full combination methodology available, we need to understand how galaxy distribution and free-streaming cosmic photon from the last scattering surface are cross-correlated to treat covariance matrix in a proper way. %In addition, the lensed information in CMB could contribute to our neutrino mass study as an independent data between the last scattering surface and the galaxies, which would be also under our consideration for the future work. The measured constraints on $m_\nu$ presented in this manuscript is trustable with low resolution catalogues provided by BOSS DR11. But we do not think that the same theoretical templates are applicable for the future precision experiments. The same methodology can be used with more rigorously developed RSD theoretical templates. We work towards to improve those using the detailed perturbation theory and the believable neutrino simulations in near future. %From the aspect of precision both in data and the theoretical template, we have another challenging point. %Since the galaxy data is from the low resolution experiment, our theoretical template based on the 1-loop perturbation calculation was available. To prepare upcoming next generation survey with large volume and deep redshift, however, full perturbation approach should be made as a preparation. As an intermediate goal, testing convergence to check if up to which order of the perturbation theory can provide stable and reasonable approximation would be enough.

1607.01074

1607

1607.07825_arXiv.txt

In cosmology, phenomenologically motivated expressions for running vacuum are commonly parametrized as linear functions $\Lambda(H^2)$ or $\Lambda(R)$. Such kind of models assume an equation of state for vacuum given by $\,\overline P_\Lambda=-\,\overline\rho_\Lambda$, relating their background pressure $\,\overline P_\Lambda$ and mean energy density $\,\overline\rho_\Lambda\equiv\Lambda/8\pi G$. This equation of state requires that the dynamic for vacuum is due to the energy exchange with the material species. Most of the approaches to background level consider only the energy exchange between vacuum and the transient dominant material component of the universe. We extend such models assuming the running vacuum as the sum of independent contributions $\,\overline\rho_{\Lambda} =\sum_i\,\overline\rho_{\Lambda i}$, associated with (and interacting with) each of the $i$ material species. We derive the linear scalar perturbations for two running scenarios, modeling its cosmic evolution and identifying their different imprints on the cosmic microwave background anisotropies and the matter power spectrum. In the $\Lambda(H^2)$ scenario the running vacuum are coupled with all the material species in the universe, whereas the $\Lambda(R)$ description only leads to coupling between vacuum and the non-relativistic matter components; which produces different imprints of the two models on the matter power spectrum. A comparison with the Planck 2015 data was made in order to constrain the free parameters of the models. In the case of the $\Lambda(H^2)$ model, it was found that $\Omega_\Lambda=0.705\pm0.027$ and $H_0=69.6\pm2.9\, km\, Mpc^{-1}\, s^{-1}$, which diminish the tension with the low redshift expectations.

One of the most important discoveries of the 20th century is that the universe is expanding \cite{Hubble}, and more surprising is that it is accelerating \cite{RiessETAJ98,PerlmuterETNature98}. At large scales gravity is the dominant force and, because its attractive nature in the context of General Relativity, the only way to have an accelerating universe is to assume a new cosmic component with the special feature to be gravitationally repulsive. This component is known as dark energy (DE), and it must have some exotic characteristics like negative pressure, and permeate every part of the universe to have a global repulsive effect. The cosmological constant (CC) of the Einstein's field equations can account for such acceleration, and jointly with another unknown component referred as cold dark matter (CDM), turn out to be a remarkably good and widely accepted cosmological model known as $\Lambda$CDM \cite{BAO,WMAP,PlanckCosmology}. However, this framework has some important theoretical problems. In the general relativity context, a bare CC needs a fine-tuning of about 100 orders of magnitude, so that combined with the expected value of the vacuum energy density in quantum field theory, reproduce the effective dark energy density estimated from astronomical observations. This theoretical conundrum is known as the CC problem \cite{WeinbergRMP89, Sahni:1999gb, Peebles:2002gy, Padmanabhan:2002ji, Sola:2013gha}. Another hassle is that in spite of behaving quite differently with respect to the cosmic expansion the CDM and DE are found to contribute to the energy content of the universe today with amounts of the same order, this riddle is known as the cosmic coincidence problem. Astronomical observations of different types support the existence of DE \cite{RiessETAJ98,PerlmuterETNature98, BAO, DWeinberg2013, WMAP, PlanckCosmology}, but do not provide a single clue about its origin from fundamental physics. This allows the proposal for a wide range of types of DE candidates besides the pure positive CC like: quintessence, K-essence, chameleon field, $f(R)$ gravity and others; for a review see \cite{Yoo:2012ug,Joyce2015} and references therein. Another dynamical DE formulation consider it as a decaying entity, which can be modeled as an effective interaction with the material components. Since the lack of information on the nature of the DE it is difficult to describe these interactions from first principles; therefore, the interactions are often described phenomenologically. The most worked approaches of interacting DE are the interacting dark sector models (DE and CDM), see \cite{Chimento2009, Koshelev2011, Abdalla} among many others. In this vein, the model proposed in the present work extends this idea, taking into account the DE interact with all the other components, and assuming that DE can be separated into a sum of different contributions, each one associated only with one material component (say photons, baryons, CDM, etc.). A recent motivation for the DE dynamics has emerged in the context of the renormalization group approach, where the simultaneous running of the CC and the gravitational coupling constant, due to quantum effects, has been considered \cite{BauerCQG05, ShapiroETJCAP05, GrandeETPLB07, GrandeETCQG10, Grande:2011xf}. In those and in the present work, the \textit{running CC} is identified as the renormalized vacuum energy density. Those renormalization group studies have shown that the corrections of the gravitational constant vary logarithmically with the scale of energy, therefore very slowly; while it is expected that the corrections to the CC are described by a power series. That has led to increased interest in the study of CC running, leaving constant G \cite{Shapiro:2009dh, Sola:2013gha, SolaAIP14}, and this is the class of models in which we are interested. As done by \cite{SolaAIP14} in the cosmological context, we identify the typical scale of energy of the process as the Hubble parameter, or the scalar curvature. It is worth noting that before the renormalization group formulation of the CC running, decaying DE models were studied by several authors from the phenomenological point of view \cite{Lima1992, Overduin:1998zv}. Most of the work done on this subject focuses mainly on the study of evolution, cosmological consequences and observational constraints of the running vacuum at the background level. Effects of these models at perturbative level have been studied shallowly, but recently have begun to receive more attention \cite{Fabris07,Borges2008,Borges2008b,Borges2011,Basilakos2012,Toribio2012,borges2015,Gomez-Valent2015,Gomez-Valent2015b,GengLee2016,GengLee2016b}. In such perturbative studies, the running vacuum is often modeled as decaying into the dominant material component of each cosmic era, i. e. without considering the contributions of the other components. Thanks to the quantity and quality of current observational data, such an approximation may not be appropriate when modeling the evolution of linear perturbations. However, expressions suggested by \cite{SolaAIP14} for the mean vacuum energy density of the form $\Lambda(H^2)$ and $\Lambda(R)$ lack a Lagrangian origin, the existence or the explicit form of the interaction to first order in perturbations is not given by itself. The aim of the present work is to find a consistent formulation for the running vacuum perturbations and its material sources, as well to identify their observational imprints on the cosmic background radiation (CMB) and matter power spectrum. The organization of the paper is as follows. In section \ref{sec_renormalization_group}, we review two types of running vacuum models, one with $\Lambda(H^2)$ and another with $\Lambda(R)$. In section \ref{sec_scalar_perturbations} we show the fluid conservation equations for coupled species, where the coupling terms still remain unspecified. In sections \ref{sec_L(H2)_pert_equations} and \ref{sec_L(R)_pert_equations}, we apply the linear scalar perturbation theory to the running vacuum models described in section \ref{sec_renormalization_group}. % Using the Boltzmann equation, we find the coupling terms of the running vacuum with the material components for each model. The behavior of the vacuum perturbations for sub-horizon modes is described, and the super-horizon initial conditions are founded. In section \ref{sec_numerical_results} we show and discuss the result of integrating numerically the complete set of cosmological equations, for which was made use the free code CLASS \cite{class}. Moreover, we use Planck 2015 data set \cite{PlanckData} and the statistical analysis package MontePython \cite{MontePython} to derive observational constraints. Finally, in section \ref{sec_conclusions} we present our conclusions and some important remarks.

This paper seeks to advance the linear perturbations study of phenomenological models for a running cosmological constant built from the main ideas of the renormalization group. We generalized two running vacuum modeling, considering the energy exchange between the vacuum and each material species, not only with the dominant component as is often done in interacting dark sector models. This improvement in such approach is necessary in order to analyze the effects of the model on perturbations. We considered two representative classes models, $\Lambda(H^2)$ and $\Lambda(R)$, derived and solved numerically the perturbation equations. In addition, we plotted two important cosmological observables: CMB temperature and matter power spectrum; as well as the evolution of the perturbations. The hypothesis of vacuum perturbations is viable if we consider a dynamical vacuum modeled as a perfect fluid. This could possibly indicate which observables are the most appropriate to track with the purpose to give more conclusive observational supports to the models. Although the both $\Lambda(H^2)$ and $\Lambda(R)$ models studied in this paper lacking of a Lagrangian formulation, the use of the first order collision terms inherited from the Boltzmann ones to zero order allows a self-consistent treatment for perturbations. As a result, the matter perturbations are only affected by its vacuum counterparts through the metric perturbations, leaving unmodified the usual transformation between synchronous and Newtonian gauges. Unlike the material components, the equations for vacuum perturbations have a non-zero collision term. Besides being able to consistently formulate a set of equations for the evolution of vacuum perturbations, in the context of the model presented here, we have put them in the middle of the cosmic inventory and estimate their impact on material components. In principle, the contribution of vacuum perturbations to the gravitational potential, throughout the cosmic history, allow that the small scales coming on the horizon are the most affected by the interaction. The numerical solutions show that this effect results in a decrease/increase of matter power spectrum in small scales, also as a leftward/rightward shift on the peak positions in the $TT$ spectrum of the CMB depending on the sign of the running parameters $\alpha$ and $\beta$. In addition, the vacuum coupling also modifies the slope of the matter spectrum on large scales; which becomes a potential source of observational distinction between the two models. The study developed in this work, complements similar ones concerning running vacuum, vacuum scalar perturbations and dark sector interaction; establishing a detailed derivation of the linear scalar perturbation in two well-motivated classes of running vacuum ($\Lambda(H^2)$ and $\Lambda(R)$), which interacts with all the material species. The vacuum perturbations were tracked in detail. The Planck constraints obtained in this work for the $\Lambda(H^2)$ model gives $10^4\alpha=-4.7\pm 6.5$, which is consistent with the $\Lambda$CDM model in $\sim0.72\sigma$, while $10^4\beta=-1.4\pm5.6$ is almost indistinguishable of the the $\Lambda$CDM model. Remarkable, it has that $\Lambda(H^2)$ calculated parameters constrained by the CMB power spectrum, $\Omega_\Lambda$ and $H_0$, seems to alleviate the tension with low-redshift observations: they present a positive sifts with respect to the standard case. Even so, the sifts in this parameters --due to the vacuum running-- are compatible with the $\Lambda$CDM case, because the constrains becomes weaker due to the degeneracies between the running parameters and the standard cosmological parameters. The enhancement of the statistical analysis with the use of additional observations such as BAO, SNe and $H(z)$ measurements could both improve the constraints in $\Omega_\Lambda$ and $H_0$, and confirm the reduction of the low and hight redshift stress present in the $\Lambda$CDM studies. \subsubsection*{Acknowledgements} ELDP. is supported by FAPESP under grants 2015/01721-9. DAT was partially supported by CAPES. This work has made use of the computing facilities of the Laboratory of Astroinformatics (IAG/USP, NAT/Unicsul), whose purchase was made possible by the Brazilian agency FAPESP (grant 2009/54006-4) and the INCT-A.

1607.07825

1607

1607.00416_arXiv.txt

The merger of a neutron star binary may result in the formation of a rapidly-spinning magnetar. The magnetar can potentially survive for seconds or longer as a supramassive neutron star before collapsing to a black hole if, indeed, it collapses at all. During this process, a fraction of the magnetar's rotational energy of $\sim10^{53}$ erg is transferred via magnetic spin-down to the surrounding ejecta. The resulting interaction between the ejecta and the surrounding circumburst medium powers a $\gtrsim$ year-long synchrotron radio transient. We present a search for radio emission with the Very Large Array following nine short-duration gamma-ray bursts (GRBs) at rest-frame times of $\approx1.3-7.6$~years after the bursts, focusing on those events which exhibit early-time excess X-ray emission that may signify the presence of magnetars. We place upper limits of $\lesssim 18-32\,\mu$Jy on the 6.0 GHz radio emission, corresponding to spectral luminosities of $\lesssim (0.05-8.3) \times 10^{39}$~erg~s$^{-1}$. Comparing these limits to the predicted radio emission from a long-lived remnant and incorporating measurements of the circumburst densities from broad-band modeling of short GRB afterglows, we rule out a stable magnetar with an energy of $10^{53}$~erg for half of the events in our sample. A supramassive remnant that injects a lower rotational energy of $10^{52}$~erg is ruled out for a single event, GRB\,050724A. This study represents the deepest and most extensive search for long-term radio emission following short GRBs to date, and thus the most stringent limits placed on the physical properties of magnetars associated with short GRBs from radio observations.

The merger of two neutron stars (NSs) in a compact binary can result in the formation of a massive NS remnant, which is generally assumed to collapse subsequently to a black hole. Accretion onto the black hole then powers a relativistic transient, a short-duration gamma-ray burst (GRB; \citealt{npp92,rj99a,ajm+05,rgb+11,bfc13,tlf+13,ber14,rlp+16}), with a prompt gamma-ray emission duration of $\lesssim 2$~sec. One of the biggest uncertainties in this canonical picture is how long the NS remnant survives prior to collapse. This depends on the mass of the final remnant and the highly uncertain Equation of State (EoS) of dense nuclear matter \citep{obg10,lhr+14,fbr+15,ltb+15,opg+16}. A massive NS remnant, which is supported against gravity exclusively by its differential rotation, is known as a {\it hypermassive} NS. Somewhat less massive NSs, which can be supported even by their solid body rotation, are known as {\it supramassive}. A hypermassive NS can survive for at most a few hundred milliseconds after the merger, before collapsing due to the loss of its differential rotation by internal electromagnetic torques and gravitational wave radiation (e.g.,~\citealt{st06}). In contrast, supramassive remnants spin-down to the point of collapse through less efficient processes, such as magnetic dipole radiation, and hence can remain stable for $\gtrsim$seconds to minutes. The discovery of NSs with masses $\approx 2~M_{\odot}$ \citep{dpr+10,afw+13} places a lower limit on the maximum NS mass, making it likely that the remnants produced in at least some NS mergers are supramassive (e.g.,~\citealt{opr+10}). The mergers of particularly low mass binaries may even produce {\it indefinitely stable} remnants, from which a black hole never forms (e.g., ~\citealt{gp13}). The high angular momentum of a merging binary guarantees that the NS remnant will be born rotating rapidly, with a spin period close to the break-up value of $\sim\!1$~ms. The remnant may also acquire a strong magnetic field, $\gtrsim\!10^{14}-10^{15}$~G, as a result of shear-induced instabilities and dynamo activity \citep{dt92,uso92,pr06,zm13}. Such a supramassive ``magnetar'' remnant possesses a reservoir of rotational energy up to $\approx 10^{53}$~erg \citep{mb14,mmk+15}, which is not available in cases where the NS promptly collapses to a black hole. Through its magnetic dipole spin-down, a magnetar remnant serves as a continuous power source, with the exact evolution of its spin-down luminosity dependent on the birth period and dipole magnetic field strength of the remnant \citep{zm01,mqt08,bmt+12,scr14,sc16}. \tabletypesize{\small} \begin{deluxetable*}{lcccccc} \tablecolumns{7} \tablewidth{0pc} \tablecaption{Log of VLA 6.0 GHz Observations \label{tab:obs}} \tablehead { \colhead {GRB} & \colhead {$z$} & \colhead {UT Date} & \colhead {$\delta t_{\rm rest}$} & \colhead {$F_{\nu}$} & \colhead {$\nu L_{\nu}$} & \colhead {X-ray behavior} \\ \colhead {} & \colhead {} & \colhead {} & \colhead {(yr)} & \colhead {($\mu$Jy)} & \colhead {(erg s$^{-1}$)} & \colhead {} } \startdata GRB\,050724A & 0.257 & 2015 Feb 22.472 & 7.629 & $<22.1$ & $2.7 \times 10^{38}$ & Extended emission \\ GRB\,051221A & 0.546 & 2015 Feb 22.718 & 5.936 & $<19.5$ & $1.4 \times 10^{39}$ & Plateau \\ GRB\,070724A & 0.457 & 2015 Feb 21.058 & 5.206 & $<19.1$ & $9.1 \times 10^{38}$ & Plateau \\ GRB\,080905A & 0.122 & 2015 Feb 23.723 & 5.769 & $<22.2$ & $5.2 \times 10^{37}$ & Plateau \\ GRB\,090510 & 0.903 & 2015 Mar 6.625 & 3.062 & $<26.5$ & $6.6 \times 10^{39}$ & Extended emission \\ GRB\,090515 & 0.403 & 2015 Mar 2.427 & 4.135 & $<22.7$ & $8.0 \times 10^{38}$ & Plateau \\ GRB\,100117A & 0.915 & 2015 Feb 27.674 & 2.671 & $<32.0$ & $8.3 \times 10^{39}$ & Plateau$^{a}$ \\ GRB\,101219A & 0.718 & 2015 Feb 24.011 & 2.437 & $<17.5$ & $2.5 \times 10^{39}$ & Plateau \\ GRB\,130603B & 0.356 & 2015 Mar 5.451 & 1.292 & $<20.6$ & $5.4 \times 10^{38}$ & Late-time excess \enddata \tablecomments{Upper limits correspond to $3\sigma$ confidence. \\ $^{a}$ The X-ray afterglow of GRB\,100117A also exhibited flaring activity \citep{mcg+11}. } \end{deluxetable*} \citet{mqt08} showed that the ongoing energy input from a long-lived magnetar is well-matched to a puzzling feature which distinguishes a subset of short GRBs: $\approx\!$ 1/4$-$1/2 of short bursts discovered with the \swift\ satellite \citep{ggg+04} have excess emission in their light curves when compared to the standard synchrotron model for afterglows. Indeed, $\approx\!15$-$20\%$ of \swift\ short GRBs have prolonged X-ray activity for tens to hundreds of seconds following the bursts themselves (``extended emission''; \citealt{nb06,pmg+09}). Other events have a temporary flattening or ``plateau'' in the flux decline rate of their X-ray afterglows for $\approx 10^2-10^3$ seconds after the burst \citep{nkg+06}. Still others have late-time excess X-ray emission on timescales of $\sim$few days \citep{pmg+09,fbm+14}. Other than the anomalous X-ray behavior, there are no obvious differences between these bursts and the normal population of short GRBs in any of their host galaxy properties \citep{fbc+13}, suggesting an origin intrinsic to the burst central engine. Long-lived magnetar remnants have been commonly invoked to explain the excess emission observed following short GRBs. Several studies have fit magnetar models to short GRBs with extended emission \citep{gow+13}, X-ray and optical plateaus \citep{rot+10,rom+13,gvo+15}, and late-time excess emission \citep{fyx+13,fbm+14}, resulting in inferred spin periods of $\approx 1-10$~ms and large magnetic fields of $\approx (2-40) \times 10^{15}$~G. As an alternative to the magnetar model, other energy sources remain energetically viable: most notably, late-time ``fall-back'' accretion onto the remnant black hole \citep{ros07,pnj08,ctg11}. In order to substantiate the magnetar scenario for this subset of short GRBs, and also provide crucial insight on the NS EoS, it is necessary to test additional predictions of the magnetar model. Synchrotron radio emission is expected from the interaction of the ejecta in a NS merger and the surrounding circumburst medium \citep{np11}, similar to a young supernova remnant. \citet{mb14} pointed out that the radio brightness of this interaction would be significantly enhanced in the case of a supramassive or stable magnetar remnant, due to the additional energy imparted to the ejecta by the injected rotational energy, which can exceed that of the dynamical ejecta by three or four orders of magnitude. Since there is substantial observational evidence linking short GRBs to NS mergers \citep{fb13,bfc13,tlf+13,ber14}, radio observations following short GRBs offer an independent way to explore the existence of long-lived (supramassive or stable) magnetar remnants. Using radio observations of seven short GRBs on $\sim 1$-$3$~yr timescales, \citet{mb14} placed constraints on the circumburst density of $\lesssim 0.1-1$~cm$^{-3}$ assuming an energy reservoir of $3 \times 10^{52}$~erg. Similarly, \citet{hhp+16} analyzed two bursts and placed limits of $\lesssim 0.001-5$~cm$^{-3}$ depending on the value of the assumed ejecta mass, and assumed the same energy of $3 \times 10^{52}$~erg. Here, we present radio observations of nine short GRBs on rest-frame timescales of $\sim2-8$ years after the bursts, focusing on those events which exhibit excess X-ray emission at early times that may signify the presence of magnetars. This sample represents the largest and deepest survey for long-timescale radio emission of short GRBs to date, and provides a unique test of the magnetar model. We utilize this data set to constrain the presence of magnetars formed as a result of the mergers. In Section~\ref{sec:obs}, we outline the sample and radio observations. In Section~\ref{sec:model}, we describe the magnetar model and in Section~\ref{sec:ar}, we present the analysis and results, including the constraints on the magnetar rotational energies and environment circumburst densities. In Section~\ref{sec:disc}, we compare this work to previous studies of emission from magnetars in short GRBs, and we conclude in Section~\ref{sec:conc}.

\label{sec:conc} We study the long-term radio behavior of nine short GRBs with early-time excess emission in the X-ray band that may signify the presence of magnetars. Through our VLA observations on rest-frame timescales of $\approx 2-8$~yr after the bursts, we find no radio emission to luminosity limits of $\lesssim (0.05-8 )\times 10^{39}$~erg~s$^{-1}$ at $6.0$~GHz. Our study demonstrates that a significant fraction of short GRBs with anomalous X-ray behavior do not have the associated radio emission predicted from long-lived magnetars with energy reservoirs of $10^{53}$~erg. We also rule out a stable magnetar with an energy reservoir of $10^{52}$~erg in a single case, GRB\,050724A. These radio observations, together with the known X-ray behavior, imply that supramassive magnetars which inject $\lesssim 10^{53}$~erg of energy are common relative to stable magnetars. Our study shows that a stiff NS EoS, corresponding to a maximum stable (non-rotating) NS mass of $M_{\rm ns} \gtrsim 2.3-2.4~M_{\odot}$, is disfavored, unless the NS mergers which give rise to short GRBs are particularly massive. However, population synthesis models suggest that such massive binaries only comprise a small fraction of all NS mergers \citep{bok+08}. A comparison of the observed rate of short GRBs to constraints on the NS merger rate from Advanced LIGO/Virgo will soon provide insight on the fraction of NS mergers which give rise to short GRBs, and thus additional insight on the NS EoS (e.g., \citealt{fbr+15}). Upcoming wide-field radio surveys will also constrain the population of long-lived magnetars \citep{mwb15}, independent of an association with short GRBs. We cannot rule out that a long-lived magnetar is responsible for the extended X-ray activity after some short GRBs. However, we can conclude that most such remnants should be supramassive and hence should collapse to black holes on timescales which are comparable to or shorter than their magnetic dipole spin-down timescales. If future radio observations can uniformly constrain the total available energy from a magnetar to $\lesssim 10^{52}$~erg, we should expect more abrupt collapse signatures in the X-ray light curves of short GRB afterglows. The lack of evidence for stable, long-lived magnetars may impact observational signatures from NS mergers at other wavelengths. For example, neutron-rich outflows from the NS merger form heavy elements via the $r$-process and undergo radioactive decay, resulting in a kilonova transient \citep{lp98}. In the absence of a long-lived magnetar, the signal is expected to peak in the near-IR band on $\sim$week timescales due to the large opacities of the heavy elements produced \citep{bk13,kbb13,th13,gkr+14,ffh+15}. In contrast, the large neutrino luminosity from a long-lived magnetar may inhibit the formation of very heavy elements, resulting in a bluer kilonova which peaks at optical wavelengths on $\sim$day timescales \citep{mf14,kfm15}. If the sample in this paper is representative of all NS mergers, this supports the idea that kilonovae associated with NS mergers peak in the redder bands. Since NS mergers are expected to be strong sources of gravitational waves, similar searches for long-term radio emission following NS mergers detected within the Advanced LIGO/Virgo horizon distance of $200$~Mpc will be able to place limits of $\lesssim 6 \times 10^{36}$~erg~s$^{-1}$. Thus, such searches will be crucial in constraining the fraction of mergers that lead to magnetars with delayed or no collapse to a black hole, to significantly higher confidence than is possible with the cosmological sample.

1607.00416

1607

1607.02625_arXiv.txt

{Theoretical modelling of time-lags between variations in the Fe K$\alpha$ emission and the X-ray continuum might shed light on the physics and geometry of the X-ray emitting region in active galaxies (AGN) and X-ray binaries. We here present the results from a systematic analysis of time-lags between variations in two energy bands ($5-7$ vs $2-4\,\mathrm{keV}$) for seven X-ray bright and variable AGN.} {We estimate time-lags as accurately as possible and fit them with theoretical models in the context of the lamp-post geometry. We also constrain the geometry of the X-ray emitting region in AGN.} {We used all available archival \textit{XMM-Newton} data for the sources in our sample and extracted light curves in the $5-7$ and $2-4\,\mathrm{keV}$ energy bands. We used these light curves and applied a thoroughly tested (through extensive numerical simulations) recipe to estimate time-lags that have minimal bias, approximately follow a Gaussian distribution, and have known errors. Using traditional $\chi^2$ minimisation techniques, we then fitted the observed time-lags with two different models: a phenomenological model where the time-lags have a power-law dependence on frequency, and a physical model, using the reverberation time-lags expected in the lamp-post geometry. The latter were computed assuming a point-like primary X-ray source above a black hole surrounded by a neutral and prograde accretion disc with solar iron abundance. We took all relativistic effects into account for various X-ray source heights, inclination angles, and black hole spin values.} {Given the available data, time-lags between the two energy bands can only be reliably measured at frequencies between $\sim5\times10^{-5}\,\mathrm{Hz}$ and $\sim10^{-3}\,\mathrm{Hz}$. The power-law and reverberation time-lag models can both fit the data well in terms of formal statistical characteristics. When fitting the observed time-lags to the lamp-post reverberation scenario, we can only constrain the height of the X-ray source. The data require, or are consistent with, a small ($\lesssim10$ gravitational radii) X-ray source height.} {In principle, the $5-7\,\mathrm{keV}$ band, which contains most of the Fe K$\alpha$ line emission, could be an ideal band for studying reverberation effects, as it is expected to be dominated by the X-ray reflection component. We here carried out the best possible analysis with \textit{XMM-Newton} data. Time-lags can be reliably estimated over a relatively narrow frequency range, and their errors are rather large. Nevertheless, our results are consistent with the hypothesis of X-ray reflection from the inner accretion disc.}

\label{sec1} According to the currently accepted paradigm, active galactic nuclei (AGN) contain a central, super-massive ($M_{\mathrm{BH}}\sim10^{6-9}\,\mathrm{M}_{\odot}$) black hole (BH), onto which matter accretes in a disc-like configuration. In the standard $\alpha$-disc model \citep{1973A&A....24..337S}, this accretion disc is optically thick and releases part of its gravitational energy in the form of black-body radiation, which peaks at optical to ultra-violet wavelengths. A fraction of these low-energy thermal photons is assumed to be Compton up-scattered by a population of high-energy ($\sim100\,\mathrm{keV}$) electrons, which is often referred to as the corona. The Compton up-scattered disc photons form a power-law spectrum that is observed in the X-ray spectra of AGN at energies $\sim2-10\,\mathrm{keV}$ \citep[e.g.][]{1991ApJ...380L..51H}. We here refer to this source as the X-ray source and to its spectrum as continuum emission. Depending on the X-ray source and disc geometry, a significant amount of continuum emission may illuminate disc and be reflected towards a distant observer. The strongest observable features of such a reflection spectrum from neutral material are the fluorescent Fe K$\alpha$ emission line at $\sim6.4\,\mathrm{keV}$ and the so-called Compton hump, which is an excess of emission at energies $\sim10-30\,\mathrm{keV}$ \citep[e.g.][]{1991MNRAS.249..352G}. Additionally, if the disc is mildly ionised, an excess of emission at energies $\sim0.3-1\,\mathrm{keV}$ can be observed \citep[e.g.][]{2005MNRAS.358..211R}. In addition to these spectral features, the X-ray reflection scenario also predicts unique timing signatures. For example, \citet{2016A&A...588A..13P} showed that X-ray reflection should leave its imprint in the X-ray power spectra. Owing to X-ray illumination, the observed power spectra should show a prominent dip at high frequencies, and an oscillatory behaviour, with decreasing amplitude, at higher frequencies. These reverberation echo features should be more prominent in energy bands where the reflection component is more pronounced. Furthermore, as a result of the different light travel paths between photons arriving directly at a distant observer and those reflected off the surface of the disc, variations in the reprocessed disc emission are expected to be delayed with respect to continuum variations. The magnitude of these delays will depend on the size and location (with respect to the disc) of the X-ray source, the viewing angle, the mass, and spin of the BH. Hints for such reverberation delays were first reported by \citet{2007MNRAS.382..985M} in Ark 564. The first statistically robust detection was later reported by \citet{2009Natur.459..540F} in 1H 0707--495, where variations in the $0.3-1\,\mathrm{keV}$ band (henceforth, the soft band) were found to lag behind variations in the $1-4\,\mathrm{keV}$ band by $\sim30\,\mathrm{sec}$ on timescales shorter than $\sim30\,\mathrm{min}$. The discovery of these time-lags, commonly referred to as soft lags in the literature, has triggered a significant amount of research over the past few years. Soft lags have been discovered in $\sim20$ AGN so far \citep[see e.g.][for a review]{2014A&ARv..22...72U}. A growing number of AGN show evidence of reverberation time-lags between the Fe K$\alpha$ emission line and the continuum \citep[e.g.][]{2012MNRAS.422..129Z,2013MNRAS.428.2795K,2013MNRAS.430.1408K,2013MNRAS.434.1129K,2013ApJ...767..121Z,2014MNRAS.440.2347M}, and between the Compton hump and the continuum \citep[e.g.][]{2014ApJ...789...56Z,2015MNRAS.446..737K}. Detecting them is a particularly difficult task because of the low sensitivity of most current detectors and the intrinsically low brightness of AGN at Fe K$\alpha$ line and Compton hump energies. Theoretical modelling of X-ray time-lags can elucidate the physical and geometrical nature of the X-ray emitting region in AGN. This requires knowledge of how the disc responds to the continuum emission, and the construction of theoretical time-lag spectra, which can then be fitted to the observed ones. Initial modelling attempts were based on the assumption that this response is a simple top-hat function \citep[e.g.][]{2011MNRAS.412...59Z,2011MNRAS.416L..94E}. \citet{2012MNRAS.420.1145C} were the first to consider a more realistic scenario, in which relativistic effects and a moving X-ray source were considered to quantify the response of the disc. They deduced that, for 1H 0707--495, a more complex physical model is required to explain both the source geometry and intrinsic variability. More recently, \citet{2013MNRAS.430..247W} considered a variety of different geometries for the primary X-ray source and deduced that, in 1H 0707--495, it has a radial extent of $\sim35r_{\mathrm{g}}$ (where $r_{\mathrm{g}}\equiv GM_{\mathrm{BH}}/c^2$ is the gravitational radius) and is located at a height of $\sim2r_{\mathrm{g}}$ above the disc plane. \citet[][; E14 hereafter]{2014MNRAS.439.3931E} were the first to perform systematic model fitting of the time-lags between the $0.3-1$ and $1.5-4\,\mathrm{keV}$ bands (henceforth, the soft excess vs continuum time-lags) for 12 AGN. They assumed the X-ray source to be point-like and located above the BH \citep[the so-called lamp-post geometry; e.g.][]{1991A&A...247...25M}, and calculated the response of the disc taking all relativistic effects into account. They deduced that the average X-ray source height is $\sim4r_{\mathrm{g}}$ with little scatter. \citet{2014MNRAS.438.2980C} were the first to model the time-lags between the $5-6\,\mathrm{keV}$ (which contains most of the photons from the red wing of a relativistically broadened Fe K$\alpha$ line) and $2-3\,\mathrm{keV}$ bands in the AGN NGC 4151. They used a similar procedure to E14, and deduced that the X-ray source height is $\sim7r_{\mathrm{g}}$, while the viewing angle of the system is $<30^{\circ}$. More recently, \citet[][; CY15, hereafter]{2015MNRAS.452..333C} simultaneously fitted, for the first time, the $4-6.5$ vs $2.5-4\,\mathrm{keV}$ time-lags and the $2-10\,\mathrm{keV}$ spectrum of Mrk 335. They found that the X-ray source is located very close to the central BH, at a height of $\sim2r_{\mathrm{g}}$. Our main aim is to study the iron line vs continuum time-lag spectra (hereafter, the iron line vs continuum time-lags), within the context of the lamp-post geometry, similarly to E14, C14, and CY15. To this end, we chose the $5-7\,\rm{keV}$ band as representative of the energy band where most of the iron line photons will be (henceforth, the iron line band), and the $2-4\,\rm{keV}$ band as the energy band where the primary X-ray continuum dominates (henceforth, the continuum band). In our case, the exact choice of these two energy bands is relatively unimportant since, contrary to previous works (with the exception of CY15), we take into account the full disc reflection spectrum in both the iron line and continuum bands when constructing the theoretical lamp-post time-lag models, which we subsequently fitted to the observed iron line vs continuum time-lag spectra. Our sample consists of seven AGN. We chose these objects because they are X-ray bright and have been observed many times by \textit{XMM-Newton}. We used all the existing \textit{XMM-Newton} archival data for these objects to estimate their iron line vs continuum time-lags. Our work improves significantly on the estimation of time-lags. We relied on \citet[][; EP16, hereafter]{2016A&A...591A.113E} to calculate time-lag estimates that are minimally biased, have known errors, and are approximately distributed as Gaussian variables. These properties render them appropriate for model fitting using traditional $\chi^2$ minimisation techniques. Our results indicate that the data are consistent with a reverberation scenario, although the quality of the data is not high enough to estimate the various model parameters with high accuracy, except for the X-ray source height.

\label{sec7} We performed a systematic analysis of the iron line vs continuum ($5-7$ vs $2-4\,\mathrm{keV}$) time-lags in seven AGN. The AGN we studied are X-ray bright and highly variable. The BH mass estimates for these sources are $\lesssim5\times10^6\,\mathrm{M}_{\odot}$, except for Mrk 335, which has a corresponding estimate of $\sim3\times10^7\,\mathrm{M}_{\odot}$ (note that these mass estimates are determined from optical techniques like reverberation mapping, and are not derived here). Our measurements are among the best that can currently be achieved and are able to be obtained for many years to come (with current X-ray satellites). Our choice of focusing on the iron line band was motivated by the simple fact that its existence indicates the presence of an X-ray reflection component (either from the disc or from distant material) in this band. It is thus is a clean band, ideal for investigating whether X-ray reflection operates in the inner part of the putative accretion disc. However, the low number of photons in this band undermines this advantage. Nevertheless, we found that the iron line vs continuum time-lags are consistent with the simplest X-ray reflection scenario. They also imply X-ray source heights that are close to those derived using data from lower energy bands. This result supports the hypothesis that the X-ray soft excess in these sources is a reflection component (see the relevant discussion in Sect.\,\ref{subsec73}). \subsection{Estimation of time-lag spectra} \label{subsec71} We used all the available archival \textit{XMM-Newton} data for seven X-ray bright and highly variable Seyfert galaxies and employed standard Fourier techniques to estimate the iron line vs continuum time-lag spectrum of each source. These sources have a large ($\gtrsim0.3\,\mathrm{Ms}$) amount of archival \textit{XMM-Newton} data. We also took the results obtained from extensive numerical simulations performed by EP16 into account, who studied the effects of the light curve characteristics (duration, time bin size, and Poisson noise) on the statistical properties of the traditional time-lag estimators assuming various intrinsic time-lag spectra commonly observed between X-ray light curves of accreting systems. EP16 found the following: \noindent \begin{itemize} \item[a)] Time-lag estimates should be computed at frequencies lower than half the Nyquist frequency. This minimises the effects of light-curve binning on their mean values. \item[b)] The cross-periodogram should not be binned over neighbouring frequencies, as this may introduce significant bias that can only be taken into account when a model CS (and not just a model time-lag spectrum) is assumed. \item[c)] Time-lags should be estimated from a cross-periodogram that is averaged over pairs of continuous light-curve segments with the same duration. \item[d)] If the number of segments, $m$, is $\gtrsim10$, the time-lag estimates will have known errors and approximately follow a Gaussian distribution, provided they are estimated at frequencies at which the sample coherence is $\gtrsim1.2/(1+0.2m)$. This minimises the effects of Poisson noise on their mean values. \end{itemize} \noindent Following these results, we chose the segment duration to be $\sim20\,\mathrm{ks}$. This limits the minimum frequency that can be reliably probed to be $\sim5\times10^{-5}\,\mathrm{Hz}$. A longer segment duration would allow us to probe even lower frequencies, but at the same time it would decrease the number of the available segments, and, consequently, increase the error of the resulting time-lag estimates. According to EP16, if the segment duration is $\gtrsim20\,\mathrm{ks}$, then the time-lag bias should be $\lesssim15\%$ compared to their intrinsic values for the model CS they considered. In Appendix \ref{appe}, we demonstrate that we do not expect the time-lag bias to be a problem in our study. The maximum frequency that can be reliably probed by the current data is set by point (d) above. The frequency at which the coherence becomes lower than the critical value of $\sim1.2/(1+0.2m)$ depends on the number of segments and is mainly determined by the energy band with the lowest average count rate. This is the iron line band in all cases; the mean count rate of all light curves in our sample is $0.38\pm0.27\,\mathrm{cts/sec}$ and $1.49\pm0.98\,\mathrm{cts/sec}$ for the iron line and continuum band, respectively. We found that the maximum frequency is $\lesssim10^{-3}\,\mathrm{Hz}$ for all sources. Given that the sources in our sample are X-ray bright and have a large amount of archival data, the available \textit{XMM-Newton} data allow for the reliable estimation of iron line vs continuum time-lags at frequencies between $\sim5\times10^{-5}\,\mathrm{Hz}$ and $\sim10^{-3}\,\mathrm{Hz}$. A direct comparison with published iron line vs continuum time-lags for the sources in our sample is complicated by three factors: the choice of energy bands, the \textit{XMM-Newton} observations used to estimate them, and the cross-periodogram smoothing and/or averaging scheme employed to estimate the time-lags. Similar energy bands to ours have been used for Mrk 335, NGC 7314, NGC 4151, and MCG--5-23-16. For Mrk 335, the time-lag magnitudes and errors we find are consisted with those reported by CY15, although they only used data from a single \textit{XMM-Newton} observation, which corresponds to $\sim40\%$ of the data we used. The iron line vs continuum time-lags reported by \citet{2013ApJ...767..121Z} for NGC 7314 are also roughly consistent in magnitude with our findings. They used data from only two \textit{XMM-Newton} observations, which corresponds to $\sim30\%$ of the data we used. Their time-lags are larger (in magnitude) than ours at low frequencies. They provide time-lag estimates at frequencies lower than ours. Owing to the limited length of the data sets they used, their low-frequency estimates must have been obtained from averaging a small number of cross-spectral estimates at neighbouring frequencies. As a result, according to the EP16, these estimates should be far from being Gaussian-distributed, and the frequently used time-lag error prescription of \citet{1999ApJ...510..874N} should severely underestimate the true scatter of these estimates around their mean. We did not estimate time-lags for NGC 4151 and MCG--5-23-16, since the available \textit{XMM-Newton} archival light curves at the time we were analysing the data were not long enough to obtain reliable (in the sense explained in Sect.\,\ref{sec3}) time-lag estimates. \subsection{Modelling the observed time-lag spectra} \label{subsec72} We considered two different model time-lag spectra: (a) a power-law time-lag spectrum that describes delays between X-ray continuum variations in different energy bands (model A), and (b) a reverberation time-lag spectrum that describes delays between the X-ray continuum and reprocessed disc emission in a lamp-post geometry (model B). The first is a phenomenological model, while the second is a physical model that depends on the central source geometry. We calculated the model B time-lag spectra by determining accurate disc response functions in the iron line and continuum bands. We fixed the photon index of the X-ray source at a value of 2 and assumed a neutral, prograde disc with an iron abundance equal to the solar value, around a spinning BH. The inner disc radius was set to the location of the ISCO, and the outer radius was fixed at $10^3r_{\mathrm{g}}$. We took all relativistic effects into account and considered the total reprocessed disc emission (and not just the photons initially emitted by the disc at $6.4\,\mathrm{keV}$) in both the iron line and continuum bands. In this respect, our modelling is more accurate than previous attempts (e.g. E14 and C14). We found that the model B time-lag spectra have a weak dependence on the BH spin and viewing angle. On the other hand, they depend strongly on the BH mass and X-ray source height. These parameters affect the model B time-lag spectra in a similar way. As the height increases, the model B time-lag spectra flatten at lower frequencies, and to a lower level; the same effect can also be produced by a higher BH mass for the same height (in units of $r_{\mathrm{g}}$). In addition, the characteristic flattening of the reverberation time-lag spectra to a constant value at low frequencies also depends on the outer disc radius. Therefore, the magnitude of this constant level cannot be used in a straight-forward way to determine either the X-ray source height or the outer disc radius, even when the BH mass is known. Our modelling can be improved in many ways. For example, we could let the slope of the X-ray continuum spectrum, as well as the iron abundance, be free parameters. These parameters mainly influence the amplitude of the disc response function (as they affect the reflection fraction in each energy band). In this case, these parameters should affect the response functions similarly to the BH spin (at small heights). Consequently, we do not expect the difference in the resulting model time-lag spectra to be significant (see the bottom left panel in Fig.\,\ref{figb1}). As shown by CY15, for instance, disc ionisation also affects the model time-lag spectra and should be included in the determination of the response functions. More importantly, however, the main limitation of our modelling is the adopted geometry. The lamp-post geometry is a simplification of the AGN X-ray emitting region. A different geometry can significantly affect the shape and amplitude of the disc response function, and as a result, it can significantly alter the resulting model time-lag spectrum (see the discussion in Appendices \ref{appb} and \ref{appd}). We adopted it (as has been done by many authors in the past) because the estimation of the disc response is relatively straightforward in this case. Furthermore, our intention was to investigate whether the observed iron line vs continuum time-lag spectra are consistent with the simplest theoretical reverberation model, and to see which constraints they can impose on the X-ray source and disc geometry. In retrospect, given the results of our study (see the discussion below), the current data sets fail to distinguish between the predictions of the lamp-post model and those from a more detailed approach. \subsection{Model-fit results} \label{subsec73} We fitted models A and B separately to the observed time-lag spectra because given their quality (limited frequency range and large errors), we would not have been able to constrain the lamp-post parameters by fitting a combined model A+B to the data. Both models provide statistically acceptable fits. We therefore cannot prefer one model based on the quality of the model fits. However, our best-fit results do provide useful hints. For example, the best-fit model A power-law index values for 1H 0707--495, MCG--6-30-15, and Mrk 766 are consistent with zero. The observed time-lags in these sources are flat, and the best-fit model A reduces to just a constant. This result (i.e. that the best-fit power-law model to the data is a horizontal line) leads us to believe that the case for X-ray reverberation time-lags is strong, at least in these three sources. If the observed time-lags were indeed representative of continuum time-lags, we would expect a non-zero best-fit slope. As we showed in Sect.\,\ref{subsec43}, the observed time-lags should have both a continuum and a reverberation component. The lack of a significant detection of the expected continuum component for these three sources (at least) is not surprising and can be explained physically. The continuum time-lags depend on the energy separation between the chosen energy bands, which is small in our case. Our best-fit model A amplitude values are systematically lower than the respective best-fit values found by E14. This is what we should expect for continuum time-lags, as the energy separation between the iron line and continuum bands is smaller than the separation between the $1.5-4$ and $0.3-1\,\mathrm{keV}$ bands used by E14. When fitting the observed time-lags to the model B time-lag spectrum, we found that the BH spin and inclination cannot be constrained. This is due to the large errors of the time-lag estimates and the weak dependence of the model B time-lag spectra on these parameters. Furthermore, there is a degeneracy between the X-ray source height and the BH mass that is due to the similar dependence of the model B time-lag spectrum on these parameters. We thus froze the BH mass value for each source to the most accurate and reliable values we could find in the literature and managed to constrain the X-ray source height. The observed iron line vs continuum time-lag spectra either require, or are consistent with, small X-ray source heights. For example, the best-fit height estimates are $\lesssim10r_g$ in three sources. The best-fit height for NGC 4051 is also consistent with such a low value. Even for Ark 564 and NGC 7314, the data are consistent with an X-ray source height as small as $\sim4r_{\mathrm{g}}$ when we considered a combined model A+B. Figure\,\ref{fig4} shows the our best-fit $h$ values versus the E14 best-fit results. The red dashed line indicates the one-to-one relation. Although most of the points are located above this line, given the large uncertainties, the plot suggests a broad agreement with the results of E14. The direct comparison is complicated because we considered more data sets than E14 for some sources. \citet{2013MNRAS.435.1511A} showed that the soft lags of NGC 4051 vary significantly and systematically with source flux. In our case, we cannot fit model B to time-lag spectra estimated from low- and high-flux segments, as the uncertainty on the model parameters will be significantly larger than what we obtain when we fit the overall time-lags. Nevertheless, if this trend is present in all AGN and in the iron line vs continuum time-lags as well, then when we average over data with a wide flux range, segments with the highest flux may dominate the cross-periodogram, as they may be associated with higher amplitude variations \citep[due to the rms-flux relation;][]{2001MNRAS.323L..26U}. If the data sets we considered exhibit a wider flux variability range than the one in the E14 data sets, differences in the best-fit results may be easier to explain. In conclusion, the soft excess vs continuum time-lags are consistent with the iron line vs continuum time-lags we presented here, in that they both support the hypothesis of disc reflection from an X-ray source that is located very close to the disc and the central BH. \begin{figure}[h!] \centering \includegraphics[width=\hsize]{figures/fig4.pdf} \caption{Comparison between the best-fit X-ray source height obtained by fitting the iron line vs continuum time-lags (vertical axis; this work) with those obtained by fitting the soft excess vs continuum time-lags (horizontal axis; E14).} \label{fig4} \end{figure} \subsection{Implications for the X-ray reflection scenario} \label{subsec74} Except for the source height, we are unable to constrain additional reverberation model parameters such as the BH mass and spin, viewing angle, and the outer disc radius. Accurate determination of these parameters would require a significant reduction in the errors of the time-lag estimates and/or an increase in the frequency range that can be reliable probed. However, this requires a substantial increase in the number of X-ray observations of AGN. For example, to probe frequencies lower by a factor of $\sim5$ (i.e. to reach a low limit of $\sim10^{-5}\,\mathrm{Hz}$), segments with a duration of $\sim100\,\mathrm{ks}$ are required. Assuming the number of segments used for the time-lag estimation remains the same as in the present work, this would require the net \textit{XMM-Newton} exposure times to increase by a factor of $\sim5$ for each source (on average). This will, however, neither decrease the error of the time-lag estimates nor allow allow us to probe higher frequencies, since both require an increase in the number of segments. Extending the high-frequency limit requires an increase of $\nu_{\mathrm{max}}$, which can only be achieved by increasing the number of segments. For example, to probe frequencies $\sim2\times10^3\,\mathrm{Hz}$ for MCG--6-30-15, the critical coherence value has to decrease from its present value of $\sim0.18$ to $\sim0.05$ (see Fig.\,\ref{fig1}). This requires the number of segments to increase from 28 to 115, which corresponds to an increase in the net \textit{XMM-Newton} exposure times by a factor of $\sim4$. This would, in turn, reduce the errors of the time-lag estimates by a factor of $\sim2$. In this case, however, we would be unable to probe lower frequencies, since this requires segments of longer duration. One possibility to extend the frequency range of the observed time-lag spectra would be to use the large volume of available archival data from past and current low-Earth orbit satellites (e.g. \textit{ASCA}, \textit{Chandra}, and \textit{Suzaku}). The idea would be to bin the respective light curves at one orbital period ($\sim96\,\mathrm{min}$) to probe low frequencies, although this requires a large number of long observations. For instance, estimating time-lags at frequencies lower than $\sim10^{-5}\,\mathrm{Hz}$ requires an ensemble of at least ten observations, which will be longer than at least a few days. We are currently investigating this possibility to estimate time-lag spectra over a wider frequency range. Given the quality of the present data sets in the iron line band and the resulting iron line vs continuum time-lag spectra, the need for constructing more sophisticated theoretical disc response functions is questionable. It seems that the best way to test the X-ray reverberation scenario and significantly constrain the model parameters is to focus on the soft excess vs continuum time-lag modelling, where the S/N of the existing light curves in the soft band is much higher than those in the iron line band. This would require considering the ionisation structure of the disc in the construction of appropriate disc response functions. Modelling the energy dependence of the time-lag spectra is another possibility. However, we note that the errors of the resulting time-lag estimates are dictated by the energy band with the lower average count rate. As such, the use of light curves over a broad energy band as a reference should not significantly lower the errors of the resulting time-lag estimates, even at the lowest possible frequencies. We plan to model the energy dependence of the observed time-lag spectra in a future work, where we will also consider \textit{NuSTAR} data to study time-lags between the Compton hump and the X-ray continuum.

1607.02625

1607

1607.03553_arXiv.txt

We present the first study of GALEX far ultra-violet (FUV) luminosity functions of individual star-forming regions within a sample of 258 nearby galaxies spanning a large range in total stellar mass and star formation properties. We identify $\sim$65,000 star-forming regions (i.e., FUV sources), measure each galaxy's luminosity function, and characterize the relationships between the luminosity function slope ($\alpha$) and several global galaxy properties. A final sample of \ntot galaxies with reliable luminosity functions are used to define these relationships and represent the largest sample of galaxies with the largest range of galaxy properties used to study the connection between luminosity function properties and galaxy environment. We find that $\alpha$ correlates with global star formation properties, where galaxies with higher star formation rates and star formation rate densities ($\Sigma_{\rm{SFR}}$) tend to have flatter luminosity function slopes. In addition, we find that neither stochastic sampling of the luminosity function in galaxies with low-number statistics nor the effects of blending due to distance can fully account for these trends. We hypothesize that the flatter slopes in high $\Sigma_{\rm{SFR}}$ galaxies is due to higher gas densities and higher star formation efficiencies which result in proportionally greater numbers of bright star-forming regions. Finally, we create a composite luminosity function composed of star-forming regions from many galaxies and find a break in the luminosity function at brighter luminosities. However, we find that this break is an artifact of varying detection limits for galaxies at different distances.

A galaxy's ultra-violet (UV) flux traces young, massive stars and thus is a good indicator of recent star formation \citep[i.e., far ultra-violet (FUV) traces $t < 100$~Myr;][]{kennicutt98,murphy11}. Furthermore, clustering studies of young stars and star clusters have shown that star formation occurs in a clustered environment \citep[$t < 100$~Myr;][]{lada03,gouliermis10,gouliermis15}, thus FUV sources with relatively low angular resolution (a few arcseconds) will trace groups of young, massive stars that formed at similar times and physical locations (i.e., star-forming regions). The star formation process shows evidence of a fractal, or scale-free, structure where the distributions of stars, star clusters, and stellar complexes show similar patterns on different physical scales \citep[see review and references within][]{elmegreen10}. As a consequence of this fractal picture, the mass function (and similarly the luminosity function; LF) of these distributions can be represented as a power-law characterized by a slope of -2 \citep{elmegreen06b}. Observationally, the mass and luminosity functions of star-forming regions (i.e., star clusters and \hii regions) have shown to be adequately approximated as a power-law with a slope of -2 $\pm$ 0.2 \citep[][Adamo et al. 2016; submitted]{zhang99,larsen02,hunter03,bik03,degrijs03b,mccrady07,cook12,whitmore14,chandar15}. However, it is not clear if galaxy environment has a significant effect on $\alpha$, nor is the scatter well understood. Many star cluster and \hii region studies that measure MF/LFs do so with small galaxy samples (usually only a few), and the methods used to identify regions and to generate MF/LFs can vary from study to study. The inhomogeneity of data and methods can add scatter to any relationship of MF/LF slope with galaxy properties and possibly mask any correlations. However, mass and luminosity functions derived for multiple galaxies with uniform data and methods have found hints of systematic trends between $\alpha$ and global galaxy properties \citep[e.g., SFR, M$_B$, galaxy type, etc. ][]{kennicutt89,elmegreen1999,youngblood99,vanzee00,thilker02,whitmore14}. These trends suggest that environment may play a role in the formation of stars and the quantification of such trends would provide clues into the star formation process. In addition to environmental affects on the star formation process, mass and luminosity functions have shown evidence for a break in their fitted power-laws at higher masses \citep[e.g., M$_{star} \sim 10^5 - 10^6$ M$_{\odot}$][Adamo et al. 2016; submitted]{gieles06a,bastian12a} and luminosities \citep[e.g., L$_{\rm{H}\alpha}\sim$38.6 erg/s][see also Whitmore et al. 2014, Adamo et al. 2016; submitted]{kennicutt89,pleuss00,bradley06}, where the slope is steeper at higher masses/luminosities. This break has been interpreted as a truncation in the mass function due to density bounded high luminosity regions suggesting a possible upper limit on the mass of star-forming regions \citep{beckman00}. Since the measurements of the masses and luminosities of bright star-forming regions required to accurately characterize this break are relatively rare even in normal star-forming galaxies \citep{larsen02,gieles06a,bastian08}, studies have constructed composite LFs that include the regions of many galaxies to increase the number statistics of bright, star-forming regions \citep{bradley06,whitmore14}. However, the galaxy samples used span a range of distance ($\sim$10s of Mpc) which may contribute to an artificial break due to different luminosity detection limits of galaxies at different distances. This paper is the first of two papers which aim to test the universal nature of the MF/LF of star-forming regions in a sample of 258 nearby galaxies whose properties span a wide range in SFR, metallicity, luminosity, and Hubble type. The current paper will look at the FUV LFs of star-forming regions while the second paper will look at the MFs. Specifically, this paper measures the LFs, quantifies and trends between LF slope and global galaxy properties, and investigates a break in the composite LF containing star-forming regions from many galaxies.

1607.03553

1607

1607.01884_arXiv.txt

{ We explore the distance-redshift relation using a cosmographic methodology, and show how the cosmographic parameters can be used to determine the redshift of transition from deceleration to acceleration. Such a transition at a low redshift occupies only a small region of the available parameter space, and the prior assumption of an early period of deceleration can significantly change the posterior constraints. We use available type Ia Supernovae (SN-Ia) and Baryon Acoustic Oscillation (BAO) data sets to determine the cosmographic deceleration $q_0$, jerk $j_0$, snap $s_0$ and lerk $l_0$ parameters. The parameters are consistent with the $\Lambda$CDM model for a flat universe within 2-sigma. We derive constraints on the redshift of transition from deceleration to acceleration for the different expansions, and find $z_{\rm acc} > 0.14$ at 95\% confidence in the most conservative case.}

While the notion of an expanding universe has been well known for over eighty years, the 1998 discovery \cite{1998ApJ...507...46S,1998AJ....116.1009R,1999ApJ...517..565P} that this expansion rate was accelerating was a major challenge to our understanding of the composition of the universe. For this reason, it is imperative that this discovery is confirmed directly with a range of models and datasets. Confirmation of the acceleration through a model-independent analysis of the data was suggested early on, through the statefinder approach \cite{StatefinderA,StatefinderB}. The cosmographic approach, where the connection between the dynamical history of the Universe and its material components (the Einstein equations) is abandoned, was first suggested in \cite{2004CQGra..21.2603V}, with subsequent work in \cite{2005GReGr..37.1541V,2008PhRvD..78f3501C,2011PhRvD..84l4061C,2015mgm..conf.1574V,Lazkoz:2013by,2008MNRAS.390..210C,2011APh....35...17S,Sollerman09}. In this approach, the expansion history $a(t)$ is reconstructed kinematically in terms of the cosmographic parameters, assuming a Taylor-series expansion of the expansion rate (Hubble parameter today) and its time derivatives. Such a model-independent analysis is more rigorous in testing the statistical significance of the accelerated expansion, and recovering details of the expansion history. One advantage of a purely kinematic analysis is that the derived quantities will also be similarly model-independent. A good example is the redshift at which the transition from deceleration to acceleration ($z_{\rm acc}$) occurs. There has already been some interest in a fully model-independent determination of this quantity, either through some limited cosmographic analysis \cite{2008MNRAS.390..210C,2011APh....35...17S} or a different presentation of the data through a remapping of the redshift $z$ \cite{2015MNRAS.446.3863S}. Other papers have addressed constraints on the transition redshift when considering models to $\Lambda$CDM \cite{Farooq2013,Farooq2013b,TransitionfR,TransitionfT}. In this paper we use both SN1a and BAO datasets to map the distance-redshift relation expansions in terms of a variable $\zeta$, first introduced in \cite{Cattoen2007}. We then use the statistical constraints on the cosmographic parameters to reconstruct the acceleration history of the Universe as a function of redshift $q(z)$, and determine limits on the redshift of acceleration $z_{acc}$. While this paper was in preparation, another paper was released claiming strong limits on the redshift of transition \cite{Moresco2016}. However, these constraints exist solely in the context of a $\Lambda$CDM analysis, assuming the Einstein equations. We would argue that this prior (that the $\Lambda$CDM model is the correct model) provides a high degree of constraining power by itself, beyond that of the data. In contrast, a model independent analysis would provide rigorous constraints on the transition redshift.

We have undertaken a cosmographic analysis of standard candle and standard ruler data in order to determine a model-independent estimate of the epoch of the onset of acceleration, as given by the redshift $z_{acc}$. We demonstrate that a large volume of the available parameter space is inconsistent with a decelerating universe at high-redshift. Even the modest prior of deceleration by $z=1$ severely limits the available parameter space. In our likelihood analysis, we have used type-Ia supernovae data from the Union 2.1 data compilation \cite{2012ApJ...746...85S} as our standard candle data set, and Baryon Acoustic Oscillation data from the 6dFGS, \cite{beutler:2011}, WiggleZ Dark Energy Survey \cite{blake:2011dm} and SDSS \cite{Percival:2010} surveys as our standard ruler data set. For the BAO dataset we standardize our rulers relative to the lowest redshift measurement (6dfGS), rather than the sound horizon at high redshift (as measured by the CMB), so as to remove any assumptions regarding the details of the expansion history at high-redshift. We perform a Markov Chain Monte Carlo analysis, marginalising over the value of $H_0$ when using the SN-Ia data. We performed three different analyses, increasing the number of terms in the expansion of the scale factor $a(t)$ in each case. In the first case with only two cosmographic parameters, we find that $q_0$ and $j_0$ are well constrained by a combination of the BAO and SN-Ia data. As we increase to three parameters $q_0$, $j_0$ and $s_0$, we find that the constraints on $q_0$ remain roughly the same, while the constraints on $j_0$ widen substantially. This is is also the case when we increase to four parameters $q_0$, $j_0$, $s_0$ and $l_0$. All our constraints are consistent with the values of the cosmographic parameters as predicted by the $\Lambda$CDM cosmological model at the 95\% confidence level. By sampling from our MCMC chains, we show the that acceleration is well constrained at very low redshift ($z<0.1$), independent to the order of the expansion. However, for $z>0.1$, the statistical limits on the deceleration $q(z)$ are very model dependent, and the data allows for a very abrupt transition to acceleration at low-redshift for the higher order models. This rapid transition models are driven by very large values of the snap $s_0$ and lerk $l_0$ parameters, that may be ruled out by future data.

1607.01884

1607

1607.03415_arXiv.txt

{The first \fermi-LAT High-energy source catalog (1FHL), containing $\gamma$-ray sources detected above 10\,GeV, is an ideal sample to characterize the physical properties of the most extreme $\gamma$-ray sources. % } {We investigate the pc scale properties of a sub-sample of radio faint 1FHL sources with the aim to confirm the proposed blazar associations, by revealing a compact high brightness temperature radio core, and we propose new low-frequency counterparts for the unassociated $\gamma$-ray sources (UGS). Moreover, we increase the number of 1FHL sources with high resolution observations to explore the possible connection between radio and $\gamma$ rays at E >10\,GeV. } {We observed 84 1FHL sources, mostly blazars of High Synchrotron Peaked (HSP) type, in the northern sky with the Very Long Baseline Array (VLBA) at 5\,GHz. These sources lack high resolution radio observations and have at least one NRAO VLA sky survey counterpart within the 95\% confidence radius. For those sources without a well identified radio counterpart we exploit the VLBA multiple phase-center correlation capability to discern among the possible low-frequency candidates. } {For $\sim$93\% of the sources of our sample we reveal a compact high brightness temperature radio core, supporting their proposed blazar association. The vast majority of the detected sources are radio weak, with a median VLBI flux density value of 16.3 mJy. For the detected sources we obtain an average brightness temperature of the order of 2 $\times 10^{10}$ K.% We find a compact component for 16 UGS, for which we propose a new low-frequency association. } { We find brightness temperature values which do not require high Doppler factors, and are in agreement with the expected values for the equipartition of energy between particles and magnetic field. We find strong indications about the blazar nature of all of the detected UGS, for which we propose new low-frequency associations. The characterization of the physical properties of this emerging population is relevant in view of the construction of the new generation Cherenkov Telescope Array. }

The vast majority of high energy (HE, 100\,MeV $<E<$ 100\,GeV) and very high energy (VHE, $E>0.1$\,TeV) $\gamma$-ray sources are associated with radio loud objects, typically blazars, i.e.\ flat spectrum radio quasars (FSRQs) or BL Lac type objects (BL Lacs) \citep[e.g. ][]{Acero2015, Ackermann2015}. Depending on the position of the peak of the synchrotron component of their Spectral Energy Distribution (SED), blazars are further classified as Low-, Intermediate-, or High-Synchrotron Peaked (LSP, ISP, HSP, respectively) sources, characterized by a synchrotron peak frequency $\nu_\mathrm{peak}$ (Hz) such that $\log \nu_\mathrm{peak}<14$, $14<\log \nu_\mathrm{peak}<15$, and $\log \nu_\mathrm{peak}>15$, respectively \citep{Abdo2010a}. In the MeV/GeV domain, the Large Area Telescope (LAT) on board the \textit{Fermi Gamma-ray telescope} (\fermi) has provided a deep, uniform sky survey detecting as many as 3,033 sources in the \fermi\ third source catalog \citep[3FGL, ][]{Acero2015}. The results from \fermi\ show that both types of blazars are common HE emitters, with FSRQs having softer spectra and being generally more luminous than BL Lacs. Moreover, \fermi\ clearly revealed the existence of a highly significant correlation between the radio flux density and the $\gamma$-ray energy flux \citep{Ackermann2011}. However, observations from Imaging Atmospheric Cherenkov Telescopes (IACTs) detect preferentially BL Lac sources, in particular of the HSP type, and no evidence of a correlation between radio and VHE emission has been reported so far. The different demographics of the VHE population is explained by two main factors: FSRQs have softer spectra and are more distant, which decreases their VHE emission due to extragalactic background light (EBL) absorption. The lack of a VHE-radio correlation may be due to the IACTs operational mode. IACTs operate in pointing mode with a limited sky coverage, and they preferentially observe sources in a peculiar state. All of these limitations introduce a strong bias in VHE catalogs, and it is difficult to assess any possible radio-VHE correlation. Nonetheless, it is likely that some physical effect may also be at work. The First \fermi-LAT Catalog of Sources above 10\,GeV \citep[1FHL,][]{Ackermann2013} is an ideal resource for addressing the connection between HE and VHE emission. The 1FHL is based on LAT data accumulated during the first 3 years of the \fermi\ mission, providing a deep and uniform all-sky survey. It contains 514 sources, % of which $\sim$76\% are statistically associated with Active Galactic Nuclei (AGN), $\sim$11\% are sources of Galactic nature (pulsars, supernova remnants, and pulsar wind nebulae), while $\sim$13\% remain unassociated. Various observational campaigns were dedicated to search for the low-frequency counterpart of the unassociated $\gamma$-ray sources (UGS) detected by \fermi\ \citep{Nori2014, Massaro2013b, Giroletti2016}. Recently, \citet{Schinzel2015} found 76 new low-frequency associations of $\gamma$-ray sources between 5 and 9\,GHz, by detecting parsec-scale emission through Very Long Baseline Interferometric (VLBI) observations. We are now investigating 1FHL sources at $\delta>0^\circ$ with the use of high angular resolution VLBI radio observations. The goals of this project are multi-fold: support the proposed blazar associations by revealing a compact high brightness temperature radio core; search for new counterparts for UGS; increase the size of the population of $E>10$\,GeV sources with high angular resolution observations; explore the existence of a correlation between VLBI and $E>10$\,GeV emission on a sample as large and unbiased as possible. In this paper we present the first results of this project. From the whole 1FHL catalog we extracted a sample of 269 sources in the northern sky with a radio counterpart in the NRAO VLA Sky Survey \citep[NVSS, ][]{Condon1998} within the 95\% confidence radius (r95), by using TOPCAT software \citep{Taylor2005}. We call this the 1FHL-n sample. For 185 out of these 269 sources Very Long Baseline Array (VLBA) archival observations are already available. Eighty-four sources ($\sim$31\%) have never been observed with the VLBA, and 21 of them are UGS. We focus our attention on this these 84 1FHL sources (see Table~\ref{tab_sources_log}). For the sources with an association we aim to confirm their low-frequency association and to study their parsec scale properties. For the UGS we want to investigate their nature, possibly detecting a compact radio source associated with them. The paper is organized as follows: in Sect.~\ref{sec.observations}, we present the observations and data reduction procedures; we show the results in Sect.~\ref{sec.results}; we discuss and summarize the results in Sect.~\ref{sec.discussion} and in Sect.~\ref{sec.conclusions}, respectively. % In a following paper, we will present a detailed statistical analysis of the entire 1FHL sample, based on new and archival data. Throughout the paper, we use a $\Lambda$CDM cosmology with $h = 0.71$, $\Omega_m = 0.27$, and $\Omega_\Lambda=0.73$ \citep{Komatsu2011}. The radio spectral index is defined such that $S(\nu) \propto \nu^{-\alpha}$ and the $\gamma$-ray photon index $\Gamma$ such that $dN_{\rm photon}/dE \propto E^{-\Gamma}$. \addtocounter{table}{1} \begin{table} \centering \caption{Observations details.} \label{obs_log} \begin{tabular}{lll} \hline \hline Observing date & Experiment Code & Stations log \\ \hline 2013 Dec 9 & S6340A & No FD \\ 2013 Oct 3 & S6340B & No LA \\ 2013 Oct 4 & S6340C & No LA \\ 2013 Sep 30 & S6340D & No LA \\ 2013 Nov 22 & S6340E & No FD \\ 2013 Dec 7 & S6340F & No FD \\ 2015 Jul 10 & S6340G & No PT \\ 2015 Dec 8 & S6340H & - \\ \hline \end{tabular} \tablefoot{Station codes: FD -- Fort Davis, LA -- Los Alamos, PT -- Pie Town.} \end{table} \addtocounter{table}{1} \begin{figure*} \includegraphics[bb=70 154 470 641, width=0.85\columnwidth, angle=-90, clip]{J0043.ps} \includegraphics[bb=70 154 470 641, width=0.85\columnwidth, angle=-90 ,clip]{J0648.ps} \\ \caption{5\,GHz VLBA image of the radio counterparts of the $\gamma$-ray sources 1FHL J0043.7+3425 (left panel), classified as a FSRQ, and 1FHL J0648.9+1516 (right panel), classified as a BL Lac. The beam size is 1.6 mas $\times$ 4.1 mas and 1.3 mas $\times$ 3.2 mas, respectively. Levels are drawn at $(-1, 1, 2, 4...) \times$ the lowest contour (that is, at 1.2 mJy/beam for the source 1FHL J0043.7+3425 and at 1.4 mJy/beam for the source 1FHL J0648.9+1516) in steps of 2. The noise level is 0.31 mJy and 0.37 mJy, respectively.} \label{maps} \end{figure*}

\label{sec.conclusions} In the present work we targeted and characterized the 1FHL sources in the northern sky with the faintest radio flux densities, which represent an important extension to an unexplored region of the parameter space, and are essential to gather a sample as large and unbiased as possible to explore the possible correlation between radio VLBI and E >10\,GeV emission. \begin{itemize} \item % For all of the detected sources we reveal a compact high brightness temperature VLBI radio core and we confirm their blazar nature. The vast majority of them are radio weak sources, with a VLBI flux density median value of 16.3 mJy, consistently with the predictions of the blazar sequence \citep[e.g.][]{Fossati1998, Ghisellini1998}. \item % Thanks to the new VLBI observations we propose new low-frequency counterparts for 16 1FHL UGS. We note that for 12 out of the 16 detected UGS our proposed low-frequency counterparts are in agreement with those proposed by using indirect and complementary association methods based on the analysis of their MWL properties, and with those included in the 2FHL and 3FGL catalogs. For the remaining four 1FHL UGS we propose for the first time a low-frequency counterpart. \item % For the sources whose radio core is resolved we obtain brightness temperature values of the order of $2 \times 10^{10}$ K, which are close to the value expected for the equipartition of energy between particles and magnetic field. Therefore, there is no evidence of a strong beaming and no high Doppler factor values are required. This result increases the statistical significance of the similar brightness temperature values obtained for other HSP sources \citep[e.g. ][]{Giroletti2004, Piner2014}. \end{itemize} While for blazars belonging to the first catalog of AGNs detected by \fermi\ \citep[1LAC, ][]{Abdo2010b}, \citet{Ackermann2011} found a strong and significant correlation between radio and $\gamma$ rays in the 0.1-100\,GeV energy range, no systematic investigation for the existence of a possible correlation between radio and VHE $\gamma$ rays has been made. The main reason is that at VHE there is no homogeneous, deep and large survey which allows us to perform a similar correlation analysis. This is because the IACTs have small field of view and they mainly observe in pointing mode. All of these limitations will be overcome in the next years thanks to the new generation Cherenkov Telescope Array \citep[CTA, ][]{Dubus2013}. The investigation in the VHE domain is important because it provides us with details about the blazar sequence (i.e. the anti-correlation between the synchrotron luminosity and the SED peak frequency \citep{Fossati1998}), and the interaction of VHE photons with the EBL \citep[e.g. ][]{Ackermann2012}. In a following paper, by using the VLBI flux densities presented in this work complemented by VLBI archival observations, we will explore and quantify the possible radio-VHE correlation for the 1FHL AGN sample. \begin{acknowledgement} \begin{small} We thank Leonid Petrov for his support and the fruitful discussions. We thank the anonymous referee for the helpful comments and discussions which improved the manuscript. This work is based on observations obtained through the S6340 VLBA project, in the framework of the \fermi -NRAO cooperative agreement. The VLBA makes use of the Swinburne University of Technology software correlator, developed as part of the Australian Major National Research Facilities Programme and operated under licence \citep{Deller2011}. We acknowledge financial contribution from grant PRIN-INAF-2011 and PRIN-INAF-2014. This research has made use of NASA's Astrophysics Data System, of the VizieR catalog access tool, CDS, Strasbourg, France and the TOPCAT software \citep{Taylor2005}. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. For this paper we made use of the NASA/IPAC Extragalactic Database NED which is operated by the JPL, Californian Institute of Technology, under contract with the National Aeronautics and Space Administration. The \textit{Fermi}-LAT Collaboration acknowledges support for LAT development, operation and data analysis from NASA and DOE (United States), CEA/Irfu and IN2P3/CNRS (France), ASI and INFN (Italy), MEXT, KEK, and JAXA (Japan), and the K.A.~Wallenberg Foundation, the Swedish Research Council and the National Space Board (Sweden). Science analysis support in the operations phase from INAF (Italy) and CNES (France) is also gratefully acknowledged. \end{small} \end{acknowledgement}

1607.03415

1607

1607.02469_arXiv.txt

A new family of inflationary models is introduced and analysed. The behaviour of the parameters characterising the models suggest preferred values, which generate the most interesting testable predictions. Results are further improved if late reheating and/or a subsequent period of thermal inflation is taken into account. Specific model realisations consider a sub-Planckian inflaton variation or a potential without fine-tuning of mass scales, based on the Planck and grand unified theory scales. A toy model realisation in the context of global and local supersymmetry is examined and results fitting the Planck observations are determined.

\leavevmode\\ Different classes of inflationary models produce wildly varying predictions for observables with some models being % ruled out by recent advances in the precision of cosmic microwave background observations. The Planck 2015 results~\cite{Ade:2015lrj} found the following value for the spectral index of the curvature perturbation, $n_s$, and the following upper bound on the ratio of the power spectra of the tensor to scalar perturbations, $r$: $n_{s} = 0.968 \pm 0.006$ at 1-$\sigma$ with negligible tensors and $r<0.11$. These (and the fact that non-Gaussian and isocurvature perturbations have not been observed) strongly suggest that primordial inflation is single-field with a concave scalar potential featuring an inflationary plateau. Indeed, while the minimal versions of chaotic and hybrid inflation are all but ruled out (but see Refs.~\cite{dihybrid,apostolos}), models with such an inflationary plateau (e.g. Starobinsky $R^2$ Inflation \cite{R2}, or Higgs Inflation \cite{Higgs}) have received enormous attention. In these models (and similar other such models, e.g. $\alpha$-attractors \cite{alpha}), the approach to the inflationary plateau is exponential. In this case, however, distinguishing between models is difficult \cite{allthesame}. In contrast, a power-law inflationary plateau was considered in Ref.~\cite{shaft}, where shaft inflation was introduced, based in global supersymmetry with a deceptively simple but highly non-perturbative superpotential \mbox{$W\propto(\Phi^n+m^n)^{1/n}$}, where $m$ is a mass scale. In this paper, we consider a new family of single-field inflationary models, which feature a power-law approach to the inflationary plateau but are not based on an exotic superpotential like in Shaft Inflation. We call this family of models power-law plateau inflation.% \footnote{This should not be confused with plateau inflation in Ref.\cite{plateau}.} Power-law plateau inflation is characterised by a simple two-mass-scale potential, which, for large values of the inflaton field, features the inflationary plateau but for small values, after the end of inflation, the potential is approximately monomial. The family is parameterised by two real parameters, whose optimal values are determined by contrasting the model with observations. We find that the predicted spectral index of the curvature perturbation is near the sweet spot of the Planck observations. For a sub-Planckian inflaton, the Lyth bound does not allow for a large value of the tensor to scalar ratio $r$. However, for mildly super-Planckian inflaton values we obtain sizeable $r$, which is easily testable in the near future. Even though our treatment is phenomenological, and the form of the potential of power-law plateau inflation is data driven, we develop a simple toy model in global and local supersymmetry to demonstrate how power-law plateau inflation may well be realised in the context of fundamental theory. We use natural units, where $c=\hbar=1$ and Newton's gravitational constant is $8\pi G=m_{P}^{-2}$, with $m_{P}=2.43\times 10^{18}\mathrm{GeV}$ being the reduced Planck mass.

We have studied in detail a new family of inflationary models called power-law plateau inflation. The models feature an inflationary plateau, which is approached in a power-law manner, in contrast to the popular Starobinsky/Higgs inflation models (and their variants) but similarly to Shaft Inflation. We have shown that power-law plateau inflation is in excellent agreement with Planck observations. To avoid supergravity corrections we mostly considered a sub-Planckian excursion for the inflaton in field space. As expected, this resulted in very small values for the ratio of the spectra of tensor to scalar curvature perturbation $r$. In an attempt to improve our results and produce observable $r$ we have considered minimising the remaining number of e-folds of primordial inflation when the cosmological scales exit the horizon. To this end, we assumed late reheating as well as a subsequent period of thermal inflation, driven by a suitable flaton field. We have managed to achieve \mbox{$r\simeq 3\times 10^{-4}$} which might be observable in the future (see Table~\ref{tab:running2}). For economy we have also investigated the possibility that our model is characterised by a single mass scale. We have found that the spectral index of the scalar curvature perturbation $n_s$ satisfies well the Planck observations but the model produces unobservable $r$. Abandoning sub-Planckian requirements allows the model to achieve much larger values of $r$. Indeed, for natural values of the mass scales (Planck and GUT scale), i.e. without fine-tuning, we easily obtain $r$ as large as a few percent (up to 9\%, see Table~\ref{tab:Maximising r}), which is testable in the near future. Our predicted values for $r$ and $n_s$ fall comfortably within the 1-$\sigma$ bounds of the Planck observations, while different models of the power-law plateau inflation family are clearly distinguishable by future observations (see Figs.~\ref{fig:OurResultsonPlanck} and \ref{fig:OurResultsonPlanck2}). From our analysis, we have found that the best choice of model in the power-law plateau inflation family has the scalar potential \mbox{$V=V_0\varphi^2/(m^2+\varphi^2)$}, which is also a member of the shaft inflation family of models \cite{shaft}\footnote{% However, in general, shaft inflation and power-law plateau inflation are different.}. Such a potential was originally introduced by S-dual inflation in Ref.~\cite{Sdual}, where, however, the inflaton was non-canonically normalised so the predicted value for $n_s$ was too large and incompatible with the Planck data. Following Ref.~\cite{Sdual} but crucially considering canonically normalised fields (i.e. minimal K\"{a}hler potential) we have constructed a toy-model realisation in global and local supersymmetry for our preferred power-law plateau inflation model. All in all, the level of success of power-law plateau inflation, and the fact that it offers distinct and testable predictions make this a worthy candidate for primordial inflation, which may well be accommodated in a suitable theoretical framework, as our toy models suggest. \paragraph{Acknowledgements}\leavevmode\\ CO is supported by the FST of Lancaster University. KD is supported (in part) by the Lancaster-Manchester-Sheffield Consortium for Fundamental Physics under STFC grant: ST/L000520/1.

1607.02469

1607

1607.07882_arXiv.txt

We run simulations to determine the expected distribution of orbital elements of nearly isotropic comets (NICs) in the outer solar system, assuming that these comets originate in the Oort Cloud at thousands of AU and are perturbed into the planetary region by the Galactic tide. We show that the Large Synoptic Survey Telescope (LSST) should detect and characterize the orbits of hundreds to thousands of NICs with perihelion distance outside 5 AU. Observing NICs in the outer solar system is our only way of directly detecting comets from the inner Oort Cloud, as these comets are dynamically excluded from the inner solar system by the giant planets. Thus the distribution of orbital elements constrains the spatial distribution of comets in the Oort cloud and the environment in which the solar system formed. Additionally, comet orbits can be characterized more precisely when they are seen far from the Sun as they have not been affected by non-gravitational forces.

The distribution of cometary orbital elements in the Oort cloud depends on the dynamical evolution of the solar system's planetesimal disk and the environment in which the solar system formed. Unfortunately, the vast majority of Oort cloud comets are unobservable, usually being seen only when they are perturbed onto orbits with perihelion $\lesssim 5$ AU. Furthermore, the orbits of visible comets have generally been modified by poorly understood non-gravitational forces \citep{Yeomans04}. For both reasons it is difficult to infer the properties of the Oort cloud from the statistical distribution of comet orbits. \par This paper describes the expected distribution of orbital elements of nearly isotropic comets (NICs). We define these to be comets that have been perturbed into the planetary region from the Oort cloud. Theoretical models of the distribution of NICs have been constructed by others e.g., \citet{Wiegert99}, \citet{Levison01} and \citet{Fouchard1, Fouchard2, Fouchard3}; the novel feature of the present study is that we focus on much larger heliocentric distances (up to 45 AU) in anticipation of deep wide-angle sky surveys that are currently under development. We use a simple physical model that assumes a static spherical Oort cloud with comets uniformly distributed on a surface of constant energy in phase space. Models of the formation of the Oort cloud \citep[e.g.,][]{Dones04} suggest that this approximation is reasonable except perhaps at the smallest semi-major axis we examine, $a=5,\!000$ AU, where the cloud is somewhat flattened. Furthermore, we assume that these comets evolve solely under the influence of the Galactic tide and perturbations from the giant planets. We do not consider the stochastic effects of passing stars, as they have effects qualitatively similar to the effects of the Galactic tide \citep{Heisler87, Collins10}; see \citet{Fouchard1, Fouchard2, Fouchard3} for a detailed comparison of the influence of these two agents on the Oort Cloud. We follow the evolution of these comets through N-body simulations for up to 4.5 Gyr and use the results of these orbit integrations to construct a simulated comet catalog. \par This topic is of special interest now as the Large Synoptic Survey Telescope (LSST) will likely see many of these NICs in the outer solar system. LSST will survey 20,000 square degrees of sky (48\% of the sphere) about 2,000 times over 10 years, down to an $r$-band magnitude of 24.5 \citep{LSST09}. LSST has a flux limit 3.2 magnitudes fainter than, and more than three times the area of, the current leader in finding faint distant solar system objects --- the Palomar Distant Solar System Survey \citep{Schwamb10}. It is expected to find tens of thousands of trans-Neptunian objects \citep{LSST09}; however we are not aware of predictions made specifically for objects originating in the Oort cloud. \par \citet{Francis05} studied the long-period ($P > 200$ years) comet population using the Lincoln Near-Earth Asteroid Research (LINEAR) survey \citep{Stokes00}. Most observed long-period comets likely originated in the Oort cloud. He found a sample of 51 long-period comets which were either discovered by LINEAR or would have been, had they not previously been discovered by another group. Fifteen of these had perihelion distances beyond 5 AU, but none beyond 10 AU. He used this sample to estimate properties of the Oort cloud. He found a``suggestive" discrepancy between the distribution of cometary perihelion distances in the observed sample and in theoretical models \citep{Tsujii92, Wiegert99}, but cautioned that the difference could be the result of a poor understanding of the rate at which comets brighten as they approach the Sun due to cometary activity. LSST will address this question by observing many comets at large heliocentric radii where they are inactive (see Section \ref{sect:disrupt}). \par \citet{Hills81} proposed that the apparent inner edge of the Oort cloud at around 10,000 AU is not due to a lack of comets at smaller semi-major axes, but rather because the perihelion distances of those comets evolve slowly, so they are ejected or evolve to even smaller semi-major axes due to perturbations from the outer planets before they become visible from Earth. In contrast, comets with semi-major axis $a \gtrsim$ 10,000 AU have their perihelion distance changed by more than the radius of Saturn's orbit in one orbital period, so they are able to jump across the dynamical barrier of the outer planets, and be seen in the inner solar system \citep{Hills81}. This barrier is not 100\% leak-proof, but as is demonstrated later in the paper, one expects the number density of comets with initial semi-major axes of 10,000 AU to decline by over two orders of magnitude interior to 10 AU. LSST should detect NICs at distances $>10\,$-15 AU and so will enable us to estimate the population of this inner Oort cloud directly, because we will be able to see NICs outside the region of phase space from which they are excluded by the giant planets. The properties of this cloud may contain information about the density and mass distribution in the Sun's birth cluster \citep{Brasser12}. \par Observing NICs far from the Sun also probes in unique ways the parts of the Oort cloud that do send comets near Earth. Non-gravitational forces due to outgassing when the comet comes near the Sun are the primary source of error in determining the original orbits of these comets \citep{Yeomans04}. It is somewhat uncertain at what radius outgassing begins, but a reasonable estimate would be around 5 AU (see discussion in Section \ref{sect:disrupt}). Therefore, astrometric observations of comets beyond $\sim \!10$ AU should allow much more precise determination of their original orbits (see discussion at the end of Section \ref{sect:orbel}).

We simulated the evolution of NICs originating in the Oort cloud as they interact with the giant planets. We used these simulations to create a catalog of simulated comet positions and velocities. We observe different distributions of orbital elements including perihelion distance, semi-major axis and inclination depending on the semi-major axes at which the NICs originate. Observations by LSST will therefore let us determine the absolute numbers of comets in the Oort cloud as a function of semi-major axis, and test Oort's standard model for the origin of comets. The distribution of NICs outside the orbits of Jupiter and Saturn will provide direct evidence for the presence or absence of the hypothesized ``inner Oort cloud" corresponding to semi-major axes between $5,\!000$ and $20,\!000$ AU. \par One surprising result is that we expect at least tens of percent of the comets observed by LSST in the outer solar system to have been interacting with the giant planets for more than 1 Gyr. This result makes the interpretation of the comet population detected by LSST more difficult and interesting, since the population and spatial distribution of comets in the Oort cloud almost certainly evolves on timescales of a few Gyr. \par We will also get a better measurement of the Oort spike --- the excess of comets in nearly parabolic orbits --- as we will be able to measure high-precision orbits in a regime where comets are likely unaffected by non-gravitational forces. This will put constraints on models of comet fading, as well as the original semi-major axes of comets. \par We will furthermore be able to constrain the size distribution of Oort cloud comets out conservatively to several tens of kilometers, and perhaps even to larger bodies, depending on the size distribution and number density of comets. Our results are based on a relatively steep slope for the size distribution of comets, $dN\propto r^{-3.76}$ (Equation \ref{dNdr}) and may substantially underestimate the total number of comets that will be discovered at large distances by LSST. \par We thank the referee, Ramon Brasser, for helpful and constructive comments and advice.

1607.07882

1607

1607.02143_arXiv.txt

We present coarse but robust star-formation histories (SFHs) derived from spectrophotometric data of the Carnegie-Spitzer-IMACS Survey, for 22,494 galaxies at $0.3<z<0.9$ with stellar masses of 10$^9$\,\Msun\ to 10$^{12}$\,\Msun. Our study moves beyond ``average" SFHs and distribution functions of specific star-formation rates (sSFRs) to individually measured SFHs for tens of thousands of galaxies. By comparing star-formation rates (SFRs) with timescales of $10^{10}, 10^9$, and $10^8$ years, we find a wide \emph{diversity} of SFHs: ``old galaxies" that formed most or all of their stars early; galaxies that formed stars with declining or constant SFRs over a Hubble time, and genuinely ``young galaxies" that formed most of their stars since $z=1$. This sequence is one of decreasing stellar mass, but, remarkably, each type is found over a mass range of a factor of 10. Conversely, galaxies at any given mass follow a wide range of SFHs, leading us to conclude that (1) halo mass does not uniquely determine SFHs, (2) there is no ``typical'' evolutionary track, and (3) ``abundance matching" has limitations as a tool for inferring physics. Our observations imply that SFHs are set at an early epoch, and that---for most galaxies---the decline and cessation of star formation occurs over a Hubble time, without distinct ``quenching" events. SFH diversity is inconsistent with models where galaxy mass, at any given epoch, grows simply along relations between SFR and stellar mass, but is consistent with a two-parameter lognormal form, lending credence to this model from a new and independent perspective.

Large data samples of galaxy photometry are now available from the present epoch back to $\sim$1 Gyr after the Big Bang. Hundreds of studies have described and analyzed these data in terms of luminosity, mass, and structural evolution, relying on trends between such quantities that assume a considerable uniformity of the growth of stellar populations and, by implication, of dark-matter halos. Several considerations, including \emph{N}-body simulations, the well-populated trend of cosmic SFR density as a function of redshift (Lanzetta, Wolfe, \& Turnshek 1995, Lilly \etal\ 1996, Pei \& Fall 1995, Madau \& Dickenson 2014), and the characterization of the controversially-named ``star-formation main sequence" (SFMS: Noeske \etal\ 2007; Whitaker \etal\ 2012)---showing a correlation between stellar mass and SFR at every epoch---have guided many studies in the crafting of mean evolutionary tracks, whose nature might lend insight into the phenomena driving the evolution of individual galaxies (e.g., Whitaker \etal\ 2014). However, it has been difficult to use such data to go beyond average properties and average evolution, to measuring the star formation histories (SFHs) of individual or classes of galaxies. Such data would inform to what degree galaxies follow similar growth histories, offset or scaled in cosmic time, or whether there is a genuine diversity in SFHs that is non-conformal. That is, they could demonstrate whether \emph{measured} SFHs fail to conform to scenarios wherein the evolution of scaling relations {\it controls} (as opposed to \emph{reflects}) galaxy growth (Peng \etal\ 2010; Leitner \etal\ 2012; Behroozi \etal\ 2013), by crossing or not appearing as offset/scaled versions of each other in mass or time. Such data could also determine if the significant scatter in the SFMS represents fundamental, long-term diversity in SFHs---generating an illusion of uniform growth patterns---or only a distracting perturbation to physically informative ``average tracks.'' The fundamental problem is that the available data, including integrated mass functions over most of cosmic time, are unable to uniquely ``connect the dots" between one epoch and another: the galaxies at later epochs are not necessarily the decendants of earlier galaxies observed to follow a similar trend. Abramson \etal\ (2016) have in particular emphasized the ambiguity of the presently available diagnostics by showing that models in which galaxy growth is conformal over mass and those that show great diversity of SFHs are both able to pass the observational tests that the present data provide. The promise of ``average SFHs" to elucidate important physical processes in galaxy evolution has arguably blinded us to the possibility that more physics will be learned from their diversity than from their sameness. This ambiguity can be broken by measuring the SFHs of individual galaxies, but this has proven very difficult to do, particularly because---from our vantage point in the local universe---stellar populations more than 2 Gyr old are essentially indistinguishable from each other. For this reason, ``population-synthesis models'' have been ineffective in describing the build-up of stellar mass even in relatively nearby, well-observed galaxies outside the Local Group, although considerable progress is coming from \emph{Hubble Space Telescope}-WFC3 observations, for example, the study of resolved stellar populations in M31 (Dalcanton \etal\ 2012). Applying the population-synthesis technique at higher redshift would allow a better resolution of the ages of older populations, but spectral observations with sufficient resolution and depth are costly. One might argue that we already know that SFHs are diverse. The iconic elliptical galaxy NGC3379 in the Leo Group likely formed most of its stars before $z=2$, while common galaxies of the same mass, for example, the Milky Way and Andromeda, have been puttering along for the full age of the universe. However, the proponents of the ``conformal" model suggest that both NGC3379 and these starforming spirals are following the same growth history, but that star formation was \emph{quenched} in the former by some mechanism that rendered inaccessible the considerable amount of gas still in its vicitiny. Some method of forcing long cooling times by heating the gas is the likely process. Suggested mechanisms include a transition from cold accretion to a hot halo driven by dark-matter halo growth (Kere{\v s} \etal\ 2005, Dekel and Birnboim 2006), heating from an active galactic nucleus (AGN) or supernovae feedback (Voit \etal\ 2015b), or the transition from low-entropy to high-entropy gas that is unable to cool onto the host galaxy (Voit \etal\ 2015a; Voit \etal\ 2015c). For galaxies like NGC3379 ``quenching'' could have been rapid---likely a gigayear or less. But what, then, of galaxies like the Milky Way and Andromeda? Both could have been forming stars consistently at a few \MsunYr\ for most of a Hubble time, or they could have had SFRs of $\sim$10 \MsunYr\ in the first few billion years of their history, falling steadily since $z\sim1$ to reach their present rates of $\sim$1\,\MsunYr.\footnote{Basic observational data for our Galaxy cannot distinguish between the two, if one is willing to accept that the present trickle of star formation is a temporary lull.} It is clear from recent studies measuring ultraviolet flux and Balmer absorption lines (Schawinski \etal\ 2014; Dressler \& Abramson 2015, p140; Vulcani \etal\ 2015a) that galaxies traversing the color space between the ``blue cloud'' and the ``red sequence'' are mostly doing so over billions of years, not in a $\ls$1\,Gyr timescale following an abrupt termination of star formation. Is such a slow decline a ``quenching?'' Has it been triggered by an event, such as an AGN or intense feedback from star formation, or is it instead nothing more than a slow exhaustion of gas suitable for star formation? Rather than ``slow quenching," as it is now being called (e.g., Barro \etal\ 2016), perhaps this is just the normal course of galaxy evolution. Both pictures have been valid, given the tools we have used to describe galaxy evolution. Gladders \etal\ (2013, hereinafter G13) developed the fast-clock/slow-clock model by assuming that galaxy SFHs are lognormal in cosmic time. This work was motivated by two observations: \begin{itemize} \item{Oemler \etal\ (2013b, hereinafter O13) studied distribution functions of specific star formation rates (sSFRs) for a sample of galaxies with redshifts $0.0 < z < 0.8$ and found that an increasing fraction had rising SFRs earlier in this epoch. The presence of such galaxies, whose abundance increasse steadily from essentially zero today to $\sim$20\% at $z\sim1$, obviates the ``$\tau$-model," in which the most aggressive SFH is constant in time (Tinsley 1972)} \item{G13, noted that the evolution of the cosmic star formation rate density (SFRD, the `Madau-Lilly' Diagram) from the present day back to $z\sim6$ is very well fit by a lognormal in time with two parameters, \T0---the half-mass time in the production of stellar mass, and $\tau$---the width of the lognormal. \begin{equation} \color{black} {\rm SFR} \propto \frac{1}{\sqrt{2\pi\tau^2}} \frac{\exp\left[-\frac{(\ln t - t_0)^2}{2\tau^2}\right]}{t} \color{black} \end{equation} G13 adopted this as a parameterization of SFHs of \emph{individual} galaxies, and showed an existence proof that the distribution of sSFRs for 2094 present-epoch galaxies, and the SFRD could be simulataneously fit by the sum of 2094 lognormal SFHs. Moreover, this model described the sSFR distributions at $z=0.2, 0.4$ from the \emph{ICBS} survey (Oemler \etal\ 2013a) and sSFR distributions for samples at $z=0.6, 0.8$ from data of the \emph{AEGIS} survey (Noeske \etal\ 2007). This is the model that Abramson (2015) and Abramson \etal\ (2015, 2016) have found to be successful in fitting a variety of other data, including mass functions and the SFMS (zero-point, slope, and scatter) back to $z\sim2$, and the zero-point and slope back to $z\sim8$.} \end{itemize} Encouraging as these and other results may be, these tests are incapable of distinguishing between conformity and diversity in SFHs.\footnote{Another non-conformal approach, the stochastic SFH model explored by Kelson (2014b), is also able to pass a wide range of observational tests.} Specifically, the work on lognormal SFHs based on the model in G13 demonstrates only that a large set of lognormal SFHs can be constructed that reproduces the available data well. But, the individual SFHs in this model cannot be tagged to specific galaxies: it is only the distribution function of the \T0\ and $\tau$ parameters, not the assignment to individual galaxies, that is robust in this approach. \begin{figure*}[t] \centerline{ \includegraphics[width=7.0in, angle=0]{f1.pdf} } \caption{Star-formation histories from the CSI Survey, plotted as log SFR in \MsunYr\ against cosmic time \emph{T}, the age of the universe. The four panels are for different ranges in total stellar mass: (upper left) log M$_*$ = 11.0--12.1; (upper right) log M$_*$ = 10.5--11.0; (lower left) log M$_*$ = 10.0--10.5; (lower right) log M$_*$ = 9.0--10.0. Each panel shows SFHs characterized by log SFRs at three epochs, for 60 randomly selected galaxies (in steps of 0.01 in redshift), out of the thousands that span the sample's redshift range, $0.3<z<0.9$. Each SFH has three measurements of SFR: (SFR1, red dots) $z=5$ (T=1.2 Gyr) to 1 Gyr before \Tobs\ (star formation not normally filling the full time); (SFR2, green dots) 1 Gyr to 200 Myr before \Tobs; and (SFR3, blue dots) 200 Myr to \Tobs. The dots are placed at the appropriate age of the universe for the mean of each star-formation epoch; for example, the green dot of SFR2 is placed at T = \Tobs - 600 Myr Typical 1$\sigma$ errors in log SFR are a 0.17, 0.27, and 0.22 for SFR1, SFR2, SFR3, respectively. As described in the text, the progression from the highest-mass (upper left) to the lowest-mass (lower right) galaxies shows a clear trend of SFHs that start early, with high SFRs that soon decline, toward an increasing fraction of galaxies with little or no early star formation that we refer to (following O13) as ``young galaxies" (not just the frosting, but the cake!). The latter dominate the SFHs in the lower right panel of log M = 9.0--10.0\,\Msun: these are observed to have rising SFRs since $z=1$, during which time most of their stellar mass was produced. There is considerable diversity in each mass range; in fact, examples of falling, constant, and rising SFHs can be found in every panel.} \end{figure*} In this paper we take a next step, presenting what we believe is compelling evidence for SFH diversity, using well-measured though coarse SFRs and mass build-up for over 22,494 galaxies over the epoch $0.30<z<0.90$ from the Carnegie-Spitzer-IMACS (CSI) Survey. Following the recent custom in colloquia of presenting conclusions at the start, we move immediately to a graphic presentation of our principal result on the diversity of SFHs. The customary discussion of the data and the analysis that compels this result \emph{follow}, and this is used to develop a different, more quantitative description of the data based on the mass of ``old'' and ``young'' stellar populations.

Our principle conclusions are as follows: \begin{enumerate} \item{Coarse SFHs from the CSI Survey exhibit a diversity of SFHs that are nonconformal. That is, they are neither replicas nor appear organized by the behavior of macroscopic/population-level scaling laws. Rather, a two-parameter description, such as an SFR that is lognormal in time---with its double timescales of \T0 and $\tau$ (Figures 9 and 10)---seems to be the minimum required to reproduce the diversity of SFHs found in our study of the CSI Survey data for galaxies of stellar mass log\,M/\Msun \,= $9-12$ and $0.3 < z < 0.9$.} \item{Our analysis of the CSI Survey data demonstrates that broad-wavelength coverage and accurate spectrophotometry can constrain the populations of stars formed over ``natural'' timescales of stellar evolution, of $10^{10}, 10^9$, and $10^8$ years. SFRs calculated from a sophisticated fitting of spectral energy distributions, based on well-understood signatures of stellar populations, are able to distinguish galaxies that formed all their stars early---in a few billion years---from those that formed stars continuously, and from those that formed most of their stars late---after $z=1$.} \item{We confirm predictions by O13 and modeling by G13 by identifying a substantial population of genuinely young galaxies, those that formed most of their stellar mass after $z=1$ and within 1 Gyr of the epoch of observation.} \item{We demonstrate through duplicate measurements of galaxy stellar mass that a parameter \emph{z5fract}---the fraction of total stellar mass that was formed in the epoch starting at $z=5$ and lasting $\sim$5 Gyr---is a simple but robust indicator of SFHs. We quantify the relative proportions of diverse SFHs through \emph{z5fract} and show that, at $z\sim0.6$ and log M$_* \sim 10$, genuinely young galaxies---those with greater stellar mass made in the final gigayear than in the time prior---are $\sim$50\% of the population, and that---though they are rare---there are young galaxies with masses of the Milky Way, $4 \times 10^{10}$\,\Msun, and higher.} \item{The different forms of SFHs are functions of total mass, but there is a large diversity at any given mass and, conversely, a wide range of mass ($\sim$1.0 dex) over which a particular SFH can be found. Assuming that stellar mass is strongly coupled to dark-halo mass, this means that galaxies with the same halo mass at any given epoch can host very different SFHs. This suggests another parameter, in addition to halo \emph{mass}, that controls when star formation begins and how it progresses. This could be a property of the halo, for example, its density or turnaround time, or of baryonic physics, such as star-formation or black hole growth feedback. These factors could drive the ``efficiency" of turning baryonic mass into stars that is offered as an explanation of why the halo mass function and the stellar mass function of galaxies do not ``track.''} \item{Our sample is dominated by ``field" galaxies, which means isolated galaxies or those in moderate-mass groups. The most massive galaxies, and those galaxies most exposed to `environment effects,' are not represented in the present study because the CSI Survey includes no rich clusters. That being said, the SFHs we present account for $\sim$95\% of all galaxies with masses larger than 10$^9$\,\Msun. The conformal approach to SFHs explicitly includes ``quenching mechanisms" to alter the course of star formation, specifically, to stop it, but only a small fraction of our sample are good candidates for ``mass quenching" or ``environmental quenching,'' two popular generic concepts. Nevertheless, a very large fraction of our 22,494 galaxy sample have evolved from SFRs of tens of \MsunYr\ early in their lifetime to $\ls$1 \MsunYr\ by the present epoch. Our data therefore suggest that natural processes working on the timescale of a Hubble time are able to regulate rising and then falling star formation without discrete, short-timescale quenching events. The picture that emerges is one of predestined paths for galaxies, perhaps modulated by mergers in their early history, with a ``clock'' and as-yet unidentified intrinsic processes that propel a galaxy---\emph{sooner or later}---to a state where halo gas is insufficient or in a unusuable state for continued star formation.} \item{These and other data support a picture where SFHs are generally locked in at an early epoch, determined by specific properties of the environment or the galaxy itself. Such built-in trajectories can provide an alternative explanation for supposedly environment-related correlations, like galaxy morphology with local density or the satellite/central paradigm, as expressions of \emph{initial conditions}. Dark-matter halos must play a role in this process, but a simple correlation of SFH and halo mass is clearly incompatible with the diverse histories we present here. We look to theoretical work to explore halo or baryonic properties that are responsible for this manifold variety of SFHs.} \end{enumerate}

1607.02143

1607

1607.06760_arXiv.txt

We have studied how the energy spectrum of muons at production affects some of the most common measurements related to muons in extensive air shower studies, namely, the number of muons at the ground, the slope of the lateral distribution of muons, the apparent muon production depth, and the arrival time delay of muons at ground. We found that by changing the energy spectrum by an amount consistent with the difference between current models (namely \epos and \qgs), the muon surface density at ground increases $5\%$ at $20^\circ$ zenith angle and $17\%$ at $60^\circ$ zenith angle. This effect introduces a zenith angle dependence on the reconstructed number of muons which might be experimentally observed. The maximum of the muon production depth distribution at $40^\circ$ increases $\sim10\text{ g/cm}^2$ and $\sim0\text{ g/cm}^2$ at $60^\circ$, which, from pure geometrical considerations, increases the arrival time delay of muons. There is an extra contribution to the delay due to the subluminal velocities of muons of the order of $\sim3$ ns at all zenith angles. Finally, changes introduced in the logarithmic slope of the lateral density function are less than 2\%.

In recent years, a new generation of experiments has measured with unprecedented precision the energy spectrum of \gls{UHECR}, demonstrating the existence of a cutoff at the GZK energies \cite{AugerSpectrum,Hires}. However, it is not yet possible to determine the mass of those particles with enough precision to constrain the astrophysical scenarios for their origin, and thus distinguish the origin of such a cutoff, which can be due to the exhaustion of the energy at the sources or to the interaction of the \gls{UHECR} with the \gls{CMB}. \gls{UHECR} interact with atmospheric atomic nuclei producing a cascading process of particle reactions commonly designated by Extensive Air Shower (EAS). EAS can be measured by detecting the radiation emitted by their passage through the atmosphere or by sampling the secondary particles at ground. The reconstruction of the primary \gls{UHECR} properties relies on our understanding of EAS physics and, in particular, of hadronic particle interactions. Our knowledge of hadronic interactions is limited, which is translated into phenomenological models including parameters supported by experiments only up to LHC energies. Moreover, the available accelerator data do not cover the full kinematic region of interest, namely the forward region, nor the full possible interaction systems such as pion-Nitrogen interactions, which are of the utmost importance for the description of the shower development. In summary, the different compositions of the UHECR, uncertainties in the hadronic interactions extrapolations, and the possibility of new physics scenarios at the highest energies often share the same phase-space of EAS observables, making it difficult to disentangle hadronic physics effects from UHECR mass determination. % The Pierre Auger Collaboration\cite{PAO_0} interprets the evolution of the electromagnetic shower maximum with increasing energy as an indication towards a heavier primary composition \cite{Xmax2014,XmaxInt2014}. On the other hand, hadronic interaction models fail to accurately predict the number of muons that reach the ground. At $10^{19}$ eV the mean muon number in simulations would have to be increased by $30\%$ (\epos) to $80\%$ (\qgs) \cite{MuonN2014} to match the number of muons deduced from the depth of the electromagnetic shower maximum. Moreover, measurements of the depth at which the muon production rate reaches its maximum at energies above $10^{19.2}$ \cite{MPD2014} are strongly inconsistent with \epos predictions. In each hadronic interaction, the produced $\pi^0$ carry around $25\%$ of the energy. They decay almost instantly into photons feeding the electromagnetic cascade. The rest of the energy is carried mainly by charged pions and also other mesons and barions, which continue the hadronic cascade. They eventually decay into muons whenever their energy drops below their critical energy, $ \mathcal{O}$(100 GeV). In this way, the muons act as tracers of the hadronic development: their study gives access to the details of the hadronic cascade. Assuming that the composition can be properly understood, a discrepancy between observed and predicted muon numbers, might indicate a change in the fraction of energy carried by the neutral pions, hadronic cross sections, or hadronic multi-particle production. This paper explores a complementary scenario: a change in the energy distribution of the muons at production (leaving the total number of produced muons unaltered) might produce changes in the muon measurements performed at ground. In particular, it could mimic (or contribute to changes in) the number of muons arriving at ground. Note that the energy distribution of the muons is directly linked to the angular distribution of muon trajectories, the probability of decay, the interactions with the geomagnetic field and Coulomb scattering. In summary, we study the sensitivity of muon observables at ground to changes in the energy of muons at production, keeping their overall number constant. The paper is organised as follows. In section \ref{section:Procedure}, we describe the muon distributions at production. In section \ref{subsection: Muons Energy}, we present the details regarding the modification of the muon energy spectrum and its propagation to the ground. The results obtained for the muon observables at ground are shown and discussed in sections \ref{section: Number of muons} and \ref{section: Signal of muons}. Finally, we summarise all results in section \ref{section: Sensitivity}.

\label{section:Conclusions} We have studied the impact of changing the muon energy spectrum at production on EAS observables related to the muon shower component. The muon content at ground increases between 5\% at $20^\circ$ zenith angle to 17\% at $60^\circ$ zenith angle when the muon spectrum is changed by $E^{\delta}$, with $\delta=+0.1$. This introduces an additional dependence on the observed number of muons at ground with respect to the zenith angle, that might be measured by experiments. The slope of the average LDF $\beta_{\mu}$, measured through a modified NKG function for $r \in [500;2000]$ m, is relatively insensitive to variations in the energy spectrum, changing by around 2\% for $\delta=+0.1$. The maximum of the apparent MPD $X^\mu_{max}$ changes by $10\text{ g/cm}^{2}$, and the kinematic delay changes by 3 ns independently of the distance to the core. In the MPD reconstruction algorithms, the kinematic delay is usually parametrized from models and subtracted from the total arrival time delay in order to access the pure geometric transformation from arrival times into depth. This could introduce a bias on the reconstructed apparent $X^\mu_{max}$ of the order of $\sim 5\text{ g/cm}^{2}$ at 60 degrees and 1700 m from the core, or $10\text{ g/cm}^{2}$ at 1000 m. The parameters studied in this work are summarised in table \ref{tab: Results}, for proton \epos simulated at $\theta=40^\circ$. Since the effect of the modified spectrum is different with zenith angle, we define the variable $\Delta_{20^\circ}^{60^\circ}(\delta)=\frac{x_{\delta}}{x_{\delta=0}}(60^\circ) - \frac{x_{\delta}}{x_{\delta=0}}(20^\circ)$, where $x_{\delta=0}$ is the parameter without modification, and it is also given in table\ref{tab: Results}.\\ It should be kept in mind, that the large number of muons on data might be caused simply by an increase in the overall number of produced muons. In that case, the shape of the muon content should be just a normalization factor at some energy for all angles. Nonetheless, if some different effect plays a role (and also due to the propagation effects in the atmosphere) the final muon normalization factor might depend on the shower zenith angle. According to \cite{GlennysICRC13}, the excess on the total signal seen in data of the Auger Observatory changes with zenith angle. This effect contains also the information about different muonic/electromagnetic ratios per angle, but it might suggest a difference in the muon normalization. In the near future, with the upgrade of the Auger Observatory(AugerPrime)\cite{AugerPrime}, it might be possible to clarify this feature and understand where the differences in the muon component of air showers come from. \begin{table}[h] \centering{} \caption[]{Summary of the variations of the parameters for proton \epos$\left(\frac{x_{\delta}}{x_{\delta=0}}\right)$, at $\theta=40^\circ$ under $\delta=\pm0.1$ modification of the spectrum, and the angular dependence of such variations as $\Delta_{20^\circ}^{60^\circ}( \delta=\pm0.1)=\frac{x_{\delta}}{x_{\delta=0}}(60^\circ)-\frac{x_{\delta}}{x_{\delta=0}}(20^\circ)$ for all variables except MPD and time, which is calculated between 40$^\circ$ and 60$^\circ$. } \begin{tabular}{ll|l|ll} \toprule \toprule & Values & Variation ($\delta=\pm 0.1$) & Angular dependence ($\delta=\pm0.1$) & Range\\ \tabularnewline \bottomrule $N_\mu$ & $1.57\times10^{7}$ & $\pm$8\%& $\pm$11.0\% & 20-60$^\circ$\\ $S_{1000}[\text{a.u.}]$ & 1.44 & $\pm8\%$ & $\pm$10\% & 20-60$^\circ$\\ $\beta_\mu$ & 1.99 & $\pm$1.6\% & $\mp$0.7\% & 20-60$^\circ$\\ $\beta_S$ & 2.06 & $\pm$1.5\% & $\pm$ 0.6\% & 20-60$^\circ$\\ $X_{max}^{\mu}[\text{g/cm}^{2}]$& 611 & $\mp$10 & $\pm$10 & 40-60$^\circ$\\ $t_\epsilon$ [ns] & 140 & $\mp$3 & $\pm$0 & 40-60$^\circ$\\ \tabularnewline \bottomrule \bottomrule \end{tabular} \label{tab: Results} \end{table} \normalsize

1607.06760

1607

1607.08257_arXiv.txt

A fading radio source, coincident in time and position with the fast radio burst FRB\,150418, has been associated with the galaxy WISE\,J071634.59$-$190039.2. Subsequent observations of this galaxy have revealed that it contains a persistent, but variable, radio source. We present e-MERLIN, VLBA, and ATCA radio observations and Subaru optical observations of WISE\,J071634.59$-$190039.2 and find that the persistent radio source is unresolved and must be compact ($<0.01$\,kpc), and that its location is consistent with the optical centre of the galaxy. We conclude that it is likely that WISE\,J071634.59$-$190039.2 contains a weak radio AGN.

Fast radio bursts (FRBs, see e.g.~\citealt{pbj+16} and references therein) are millisecond-duration bursts of radio emission that have been observed at the Parkes, Arecibo, and Green Bank radio telescopes~(\citealt{lbm+07,sch+14,mls+15}). FRBs have dispersion measures (DMs), a measure of the electron column density, that range from $1.4$ to $33$ times the maximum Galactic contribution~\citep{cl02}, thought to be attributable to free electrons in the intergalactic medium. With this interpretation the distances to FRBs are cosmological~\citep{lbm+07,tsb+13}, and the corresponding luminosities of the FRB signals are thus many orders of magnitude higher than typical pulsar luminosities. Non-cosmological explanations have been put forward (e.g.\,\citealt{bbe+11,lsm14,kon+14}) but a cosmological interpretation remains favored, based on current observational evidence. While the extragalactic interpretation of FRBs currently prevails, their progenitor(s) are as yet unknown. In an effort to determine the nature of FRBs, the SUrvey for Pulsars and Extragalactic Transients (SUPERB) performs real time FRB searches at the Parkes telescope, and employs an array of multi-wavelength telescopes to follow up FRB discoveries. Multi-wavelength follow-up of FRB\,150418 led, for the first time, to the detection of a fading radio source that was associated with a galaxy at $z=0.49$~\citep{kjb+16}. This galaxy is also detected in the mid-infrared by WISE \citep{wem+10} and cataloged as WISE\,J071634.59$-$190039.2. Radio imaging observations of WISE\,J071634.59$-$190039.2 with the Australia Telescope Compact Array (ATCA) showed a source declining by a factor $\sim3$ in brightness at 5.5\,GHz during the first 6\,d after the FRB. This source subsequently settled at a persistent brightness of approximately 100\,$\upmu$Jy\,beam$^{-1}$. Comparing this behaviour to the results of transient surveys \citep{bhh+15,mhb+16} led \citet{kjb+16} to argue in favour of the association between FRB\,150418 and the fading radio source, and hence with the galaxy. The association of FRB\,150418 with the fading radio source and hence with WISE\,J071634.59$-$190039.2 met with criticism. JVLA observations of the persistent radio source in WISE\,J071634.59$-$190039.2 showed rapid variability, which led \citet{wb16} to argue that the variability of the fading radio source is consistent with the intrinsic or scintillating behaviour of a compact, weak active galactic nucleus (AGN). Similarly, \citet{aj16} show that the variability of the persistent source may be extrinsic and attributable to refractive scintillation in the Milky Way, assuming a compact radio source is present in the galaxy. \citet{vrm+16} find a flat radio spectrum for the persistent source and suggest it is consistent with the properties of an AGN. Since FRB\,150418 is the first FRB for which a radio counterpart and host galaxy have been suggested, the host galaxy warrants closer study. In this Letter we report on an astrometric radio and optical analysis that establishes that WISE\,J071634.59$-$190039.2 currently contains a single weak, compact radio source consistent with an AGN located at the centre of the galaxy.

\label{sec:discussion} A single unresolved radio source is detected in the ATCA, e-MERLIN and VLBA observations of WISE\,J071634.59$-$190039.2. The source has a brightness of 130 to 151\,$\upmu$Jy\,beam$^{-1}$ at frequencies between approximately 4.8 and 5.3\,GHz, consistent within the uncertainties. The coordinates derived from the ATCA, e-MERLIN and VLBA observations are in agreement and are consistent with previously reported ATCA \citep{kjb+16} coordinates. These detections are consistent with the low significance detection at the same position with the EVN, reported by \citet{mgg+16a,mgg+16b}, after our high resolution astrometry was first reported as an ATel \citep{bbt+16a}. We note that the JVLA position by \citet{vrm+16} is inconsistent in right ascension with our e-MERLIN and VLBA positions. The coordinates of the compact radio source are plotted in Figure\,\ref{fig:chart} and are overlaid on the Subaru Suprime-Cam $i^\prime$-band image of WISE\,J071634.59$-$190039.2 that was obtained on 2015, April 19, as well as the VLBA radio image. We conservatively plot 95\% confidence uncertainty regions of the radio source positions. The optical centre of light of WISE\,J071634.59$-$190039.2 is offset from the C-band VLBA and e-MERLIN positions by $\Delta \alpha=0\farcs01(10)$ and $\Delta \delta=-0\farcs05(8)$. On this basis, our astrometry shows that, within the uncertainties quoted above, the location of the compact radio source is consistent with the optical centre of light of WISE\,J071634.59$-$190039.2 in the Subaru $i^\prime$-band image. The compact radio source detected recovers a very high percentage of the persistent flux density reported by \citet{kjb+16} and \citet{vrm+16}, indicating that little, if any, extended radio emission exists. The upper limit to the size of the radio source of $<1.5$\,mas (implying a brightness temperature in excess of $5\times10^6$\,K), corresponds to a physical size of less than 0.01\,kpc at a redshift of $z=0.49$. An interpretation of the radio emission as due to a star formation region therefore appears highly unlikely as circumnuclear star formation regions (e.g. NGC\,253; \citealt{lt06}), those associated with merger activity \citep{edg+11}, or jet induced star formation regions \citep{ssc14}, are generally two orders of magnitude larger than our upper limit. Furthermore, the observed brightness temperature exceeds that expected from thermal radio emission processes associated with star formation regions. The location of the compact radio source is consistent with our best estimate of the centre of its host galaxy. Thus, our data are consistent with the existence of a weak radio AGN within the galaxy (see \citealt{gbp+13} for a discussion of the existence of AGN in `radio quiet' galaxies). While the emission of the persistent source can be shown to be compact on VLBI angular scales (milliarcseconds), interstellar scintillation requires structure which is compact on far smaller scales (microarcseconds; \citealt{mgbh13}). Our ATCA, e-MERLIN and VLBA observations cannot directly probe the required angular scales. Hence, the presence of a compact persistent radio source such as an AGN in WISE\,J071634.59$-$190039.2 allows the scenario proposed by \citet{wb16} and \citet{aj16}; that the fading radio source coincident in position with WISE\,J071634.59$-$190039.2 and coincident in time with FRB\,150418 could be due to interstellar scintillation. The best probe of the scintillation interpretation will come from extensive radio photometry measurements, as the signature of scintillation is well known and can be tested against the data \citep{aj16}. A careful analysis of all available flux density measurements of WISE\,J071634.59$-$190039.2 from the ATCA, JVLA, e-MERLIN, VLBA, EVN, and other facilities should reveal whether the variability properties of the compact radio source are consistent with intrinsic AGN variability or scintillation.

1607.08257

1607

1607.04079_arXiv.txt

Prestellar cores, by definition, are gravitationally bound but starless pockets of dense gas. Physical conditions that could render a core starless(in the local Universe) is the subject of investigation in this work. To this end we studied the evolution of four starless cores, {\small B68, L694-2, L1517B, L1689}, and {\small L1521F}, a {\small VeLLO}. The density profile of a typical core extracted from an earlier simulation developed to study core-formation in a molecular cloud was used for the purpose. We demonstrate - (\textbf{i}) cores contracted in quasistatic manner over a timescale on the order of $\sim 10^{5}$ years. Those that remained starless did briefly acquire a centrally concentrated density configuration that mimicked the density profile of a unstable Bonnor Ebert sphere before rebounding, (\textbf{ii}) three of our test cores viz. L694-2, L1689-SMM16 and L1521F remained starless despite becoming thermally super-critical. On the contrary B68 and L1517B remained sub-critical; L1521F collapsed to become a VeLLO only when gas-cooling was enhanced by increasing the size of dust-grains. This result is robust, for other cores viz. {\small B68, L694-2, L1517B} and \small{L1689} that previously remained starless could also be similarly induced to collapse. Our principle conclusions are : (\textbf{a}) acquiring the thermally super-critical state does not ensure that a core will necessarily become protostellar, (\textbf{b}) potentially star-forming cores (the {\small VeLLO} {\small L1521F} here), could be experiencing coagulation of dust-grains that must enhance the gas-dust coupling and in turn lower the gas temperature, thereby assisting collapse. This hypothesis appears to have some observational support, and (\textbf{c}) depending on its dynamic state at any given epoch, a core could appear to be pressure-confined, gravitationally/virially bound, suggesting that gravitational/virial boundedness of a core is insufficient to ensure it will form stars, though it is crucial for gas in a contracting core to cool efficiently so it can collapse further to become protostellar. Gas temperature in these purely hydrodynamic calculations was calculated by explicitly solving the respective equations of thermal-balance for gas and dust; the attenuation factor for the interstellar radiation-field, $\chi$, which in literature has been well constrained to $\sim 10^{-4}$ for putative star-forming clumps and cores was adopted in these calculations.

\label{sec:intro} It is well-known that a number of cores do in fact remain starless. Key physical characteristics about these cores to have emerged from various detailed observational surveys are - \textbf{(a)} their longevity or in other words, the fact that some cores remain starless for a period significantly longer than others do (e.g. Ward-Thompson \emph{et al.} 2007), and \textbf{(b)} that cores typically are not pockets of isothermal molecular gas, but are colder close to their respective centres. Also, gas within cores is usually sub-sonic or transonic at best (e.g. Evans \emph{et al.} 2001, Crapsi \emph{et al.} 2004; 2007, di Francesco \emph{et al.} 2007, Schnee \emph{et al.} 2013). Observational estimates of the typical age of cores derived from detailed studies of their chemical composition and the strength of the magnetic field threading them are on the order of $\sim 10^{6}$ years, which is on the order of a free-fall time for a typical prestellar core (e.g. Beichman \emph{et al.}1986, Kirk \emph{et al.} 2005, Ward-Thompson \emph{et al.}2007; see e.g. Maret \emph{et al.} 2013 for an estimate of the chemical age). Broadly speaking, models attempting to explain the evolution of prestellar cores may be classified into two paradigms : \textbf{(1)} those that propose quasistatic evolution of magnetised cores, and \textbf{(2)} those in which cores form and evolve on a relatively short time-scale in turbulent clouds, the so-called paradigm of rapid star-formation. \\ \\ Depending on the strength of magnetic field threading a core, it may be further classified as magnetically sub-critical(mass to flux ratio less than unity), or super-critical(mass to flux ratio greater than unity). Sub-critical cores collapse on the ambipolar diffusion timescale (e.g. Shu \emph{et al.} 1987, Mouschovias 1991), while those that are super-critical, in the absence of turbulence, collapse essentially on the free-fall timescale (e.g. Nakano 1998). This model, however, encounters a number of difficulties : \textbf{(i)} the sub-critical nature of cores is inconsistent with the findings of a detailed study of a handful of isolated cores that show these cores are threaded by only a relatively weak magnetic field (e.g. Kirk \emph{et al.} 2006), \textbf{(ii)} the ambipolar diffusion timescale, typically on the order of $\sim 10^{7}$ yrs - the lifetime of a molecular cloud (e.g. Nakano 1998, Ciolek \& Basu 2000), is much higher than the estimates of typical core lifetime, and \textbf{(iii)} it is also inconsistent with the paradigm of rapid star-formation, reinforced by the observations of several nearby star-forming regions where stellar populations are found to have ages significantly smaller than the crossing-time for their respective host cloud (e.g. Hartmann \emph{et al.} 2001, Elmegreen 2000). \\ \\ In the dynamic picture of star-formation, on the other hand, supersonic turbulence, perhaps moderated by a weak magnetic field, largely regulates formation of prestellar cores as they condense out via amplification of local perturbations in the density field of molecular clouds. Also, since turbulence decays rapidly, over less than a free-fall time, cores supported by turbulence could possibly explain various features of starless cores (e.g. Mac Low \emph{et al.}1998, Ballesteros-Paredes \emph{et al.} 2003; 2006). Cores, according to this model, are in approximate (thermal) pressure-equilibrium with the ambient medium and have a density distribution that reflects the profile of a pressure-confined {\small BES}. One of the difficulties with this model of core-formation is the transient nature of resulting cores on account of being merely in hydrostatic equilibrium (e.g. Mac Low \& Klessen 2004), which is inconsistent with typical properties of starless cores. Furthermore, simulations such as those by Walch \emph{et al.}(2010), for instance, demonstrate that turbulence can only delay collapse in gravitationally super-critical cores, but cannot arrest it. Evidently a solution to the problem of starless nature of cores must be sought beyond these tested ideas such that their physical properties noted above can be reconciled. \\ \\ Evolution of cores has been studied analytically by a number of authors over the last four decades. Some of the earliest models suggested by e.g., Bodenheimer \& Sweigert (1968), Larson (1969) and Penston(1969) were criticised later by Shu (1977) on grounds of idealised(artificial) initial conditions. The Shu model that envisaged a test core modelled as a unstable Bonnor-Ebert sphere({\small BES}) which would then collapse in the \emph{inside-out} manner was itself later deemed physically unrealisable, for it is difficult to reconcile the assembly of a core that is critically stable at best, and unstable at worst (Whitworth \& Summers 1985; hereafter WS). WS demonstrated analytically that a pressure-confined {\small BES} could indeed be driven to collapse in \emph{outside-in} fashion by a compressional wave triggered by raising sufficiently the magnitude of the externally confining pressure, $P_{ext}$. This was later demonstrated numerically by Hennebelle \emph{et al.}(2003). Physically, this would be the case in regions of active star-formation where energy and momentum is dumped into the ambient medium by ongoing episodes of star-formation. Similar arguments were presented by G{\' o}mez \emph{et al.}(2007) who demonstrated that a sufficiently strong inwardly propagating compressional wave could induce a {\small BES} to collapse. More recently, Anathpindika \& Di Francesco (2013) showed that a core modelled as a non-isothermal {\small BES} and with sufficiently cold interiors can also collapse in this manner to become protostellar. Other cores in their work where gas close to the core-centre was not sufficiently cold remained starless and in fact, oscillated radially. Similarly, Kaminski \emph{et al.} (2014) demonstrated that {\small BES}s could be similarly induced to collapse by increasing the magnitude of $P_{ext}$, effected by raising the density of the confining medium. They also showed that {\small BES}s enveloped by a hot tenuous medium would likely evolve in quasistatic manner before eventually collapsing. \\ \\ Models invoking detailed radiative transfer modelling of starless cores include those suggested by Keto \& Caselli (2008, 2010). These authors classified starless cores into two classes viz., thermally sub-critical(central density, $n_{c}\lesssim 10^{5}$ cm$^{-3}$) and super-critical($n_{c}\gtrsim 10^{5}$ cm$^{-3}$) cores. They argue, gas temperature in centrally dense cores is lowered more efficiently via collisional coupling between gas and dust which renders such cores thermally super-critical and therefore, more susceptible to collapse which is likely to proceed in the outside-in manner (e.g. Williams \emph{et al.} 1999, Caselli \emph{et al.} 2002, Schnee \emph{et al.} 2007). On the other hand, thermally sub-critical cores are likely to be relatively stable against self-gravity and may exhibit radial oscillations (e.g. Lada \emph{et al.} 2003, Aguti \emph{et al.}2007). These models usually assume a density distribution to model the temperature profile of a typical starless core. Other relatively simple analytic models such as the Plummer-type suggested by Whitworth \& Ward-Thompson (2001) produces rapid density enhancement on a timescale much shorter than the observational estimates of contraction timescale. \\ \\ In the present work we therefore propose to numerically investigate the evolution of a set of cores whose physical characteristics are well-known. We are particularly interested in studying the conditions under which a core is likely to be rendered unstable against self-gravity. To this end we will self-consistently deduce the temperature profile for a core that remains starless and one that collapses. Also, we will study the temporal excursion of the virial components of our test cores to examine if gravitational and/or virial boundedness is sufficient to ensure that a prestellar core becomes protostellar. This article is organised as follows : in \S 2 we will discuss the choice of our initial conditions and the numerical scheme employed to investigate the problem. We will present our results and discuss them in respectively \S 3, and \S 4 before concluding in \S 5.

In light of the results obtained from numerical simulations discussed in this work we argue, the fate of a prestellar core bears largely upon its temperature profile. Our realisation of the core {\small B68} with different choices of the initial distribution of gas shows that the likely fate of a test core bears more strongly on the temperature profile that it develops over the course of its evolution. For even the test core {\small B68} modelled as a unstable {\small BES}, like other test cores that remained starless in this numerical exercise, rebounded soon after acquiring a centrally peaked density profile. Furthermore, other cores that remained starless were seen to mimic the profile of a unstable {\small BES} at the epoch when they were centrally most concentrated and yet, remained starless. We note, some of our test cores viz., {\small L694-2, L1689-SMM2} and {\small L1521F}(with standard dust grain-size) did in fact become thermally super-critical before rebounding, while the remaining two viz., {\small B68} and {\small L1517B} remained sub-critical with a peak density less than 10$^{5}$ cm$^{-3}$. It must be noted here that the core {\small L1521F} became super-critical and eventually a {\small VeLLO} only with a larger dust grain-size. In fact, with a larger grain-size all other starless cores also developed approximately uniform cold interiors before eventually collapsing. These observations from the present work lead to the conclusion that thermal super-criticality is probably only a necessary condition, though not sufficient, to ensure that a core becomes protostellar. \\ \\ We suggest, the strength of gas-dust coupling in a core likely holds the key to the eventual fate that befalls the respective core. In this case, we demonstrate, gas temperature within a core is lowered more uniformly. With the choice of grain-size for dust typically found in the {\small ISM} we observed that the core {\small L1521F} developed a temperature profile similar to that observed in a typical starless core such as the {\small B68}. In this latter case gas-temperature in the core was approximately uniform, on the order of 12 K-14 K and showed a gradual reduction to $\lesssim$10 K only very close to its centre (see plots in Fig. 10). The conclusion about fluffy dust grains is consistent with the earlier suggestion by Keto \& Caselli (2008) in this regard. This idea is presently supported by at least tentative evidence from a few detailed studies of dust properties towards some prestellar cores and clumps. Our result here therefore calls for more studies about the nature of dust grains in prestellar cores and in cores that remain starless. \\ \\ Furthermore, realisations of cores that remained starless in this work also demonstrate that these cores can indeed survive over at least a few free-fall times (see Table 2). In fact realisations discussed here reproduce the typical core contraction time, the in-fall velocities as well as the rate of accretion for a typical {\small VeLLO}. These simulations also show that the virial state of a core changes over the course of its evolution. We tracked this change on the so-called virial chart that marks the temporal variation of virial components of a core. Cores were seen to evolve from a state that was pressure-confined to one that was gravitationally bound. Similarly, we have also shown that cores make a transition from a state that is less virially bound to one that is more strongly virially bound. Listed in Table 3 are the magnitudes of virial coefficients $Q$, and $X$, for respective cores modelled in this work. The magnitudes of these coefficients for cores that remained starless in this work call for attention; not only are these cores virially bound ($X >$ 1), but also relatively strongly bound ($Q >$1), except {\small L1517B}. This suggests, gravitational and/or virial boundedness of a core is not a sufficient condition for it to collapse, though it is probably necessary. A prospective protostellar core must continue to cool efficiently as it continues to contract in quasistatic manner. In the foreseeable future we propose to expand the remit of this work by coupling the hydrodynamic calculations discussed here by including the effects of depletion of various molecular species as a core continues to acquire higher densities during its collapse. Also, we should like to examine the impact of ambient environment on the evolution of cores and furthermore, include the magnetic field in these calculations to study how it affects the dynamic evolution of a core.

1607.04079

1607

1607.03145_arXiv.txt

We present baryon acoustic oscillation (BAO) scale measurements determined from the clustering of 1.2 million massive galaxies with redshifts $0.2 < z < 0.75$ distributed over 9300 square degrees, as quantified by their redshift-space correlation function. In order to facilitate these measurements, we define, describe, and motivate the selection function for galaxies in the final data release (DR12) of the SDSS III Baryon Oscillation Spectroscopic Survey (BOSS). This includes the observational footprint, masks for image quality and Galactic extinction, and weights to account for density relationships intrinsic to the imaging and spectroscopic portions of the survey. We simulate the observed systematic trends in mock galaxy samples and demonstrate that they impart no bias on baryon acoustic oscillation (BAO) scale measurements and have a minor impact on the recovered statistical uncertainty. We measure transverse and radial BAO distance measurements in $0.2 < z < 0.5$, $0.5 < z < 0.75$, and (overlapping) $0.4 < z < 0.6$ redshift bins. In each redshift bin, we obtain a precision that is 2.7 per cent or better on the radial distance and 1.6 per cent or better on the transverse distance. The combination of the redshift bins represents 1.8 per cent precision on the radial distance and 1.1 per cent precision on the transverse distance. This paper is part of a set that analyses the final galaxy clustering dataset from BOSS. The measurements and likelihoods presented here are combined with others in \cite{Acacia} to produce the final cosmological constraints from BOSS.

The Baryon Oscillation Spectroscopic Survey (BOSS) has built on the legacy of previous wide-field surveys such as Two Degree Field Galaxy Redshift Survey (2dFGRS; \citealt{2df}) and the Sloan Digital Sky Survey I-II (SDSS; \citealt{York00}) to amass a sample (\citealt{DR12,Reid15}) of more than 1 million spectroscopic redshifts of the galaxies with the greatest stellar mass to $z < 0.75$. This final BOSS data set represents the premier large-scale structure catalog for use in measuring cosmologic distances based on the baryon acoustic oscillation (BAO) feature and the rate of structure growth via the signature of redshift-space distortions (RSD). Previous results have demonstrated that the current and previous BOSS data sets produce precise and robust BAO and RSD measurements (c.f., \citealt{Reid12,alphDR9,Chuang13,Kazin13,Sanchez13,Anderson14DR9,alph,Sanchez14,Samushia14,CuestaDR12,Gil15BAO,Gil15RSD}). The results of \cite{Ross12,Ross14,Alam15,Osumi15} have demonstrated that the BOSS results are robust to observational systematic concerns and details of sample selection related to galaxy evolution. This paper represents a final, detailed, investigation of observational systematic concerns in the BOSS sample. We detail how the angular selection functions of the BOSS galaxy samples are defined and test for any systematic uncertainty that is imparted into BAO measurements based on this process. The work we present details how BOSS galaxy data can be combined into one BOSS galaxy catalog, and that robust BAO distance and RSD growth measurements can be obtained from the data set. This work uses the `combined' BOSS galaxy catalog to determine BAO scale distance measurements, making use of density field `reconstruction' (c.f., \citealt{Pad12}). Following \cite{Xu13,Anderson14DR9,alph,Ross152D,CuestaDR12}, we use the monopole and quadrupole of the correlation function to measure the expansion rate, $H(z)$, and the angular diameter distance, $D_A(z)$, at the redshift of BOSS galaxies. BAO measurements obtained using the monopole and quadrupole of the power spectrum are presented in \cite{BeutlerDR12BAO}, while \cite{VargasDR12BAO} diagnoses the level of theoretical systematic uncertainty in the BOSS BAO measurements. Measurements of the rate of structure growth from the RSD signal are presented in \cite{BeutlerDR12RSD,GriebDR12RSD,SanchezDR12RSD,SatpathyDR12RSD}. \cite{Acacia} combines the results of these seven (including this work) results together into a single likelihood that can be used to test cosmological models. The paper is outlined as follows: In Section \ref{sec:analysis} we describe how clustering measurements and their covariance are determined, and how these measurements are used to determine the distance to BOSS galaxies using the BAO feature; in Section \ref{sec:data}, we describe how BOSS galaxies are selected, masked, and simulated. In section \ref{sec:weights}, we describe how weights that correct for observational systematic relationships with galaxy density are determined and applied to clustering measurements. In Section \ref{sec:clus}, we present the configuration-space clustering of BOSS galaxies, demonstrating the effect of systematic weights, comparing the clustering of different BOSS selections and showing that the clustering in the independent NGC and SGC hemispheres is consistent and that the separate BOSS selections can be combined into one BOSS sample to be used for clustering measurements. In Section \ref{sec:BAOrob}, we show that the BOSS BAO measurements are robust to observational systematics (both for data and mock samples). In Section \ref{sec:BAOres}, we present the BAO measurements of the BOSS combined sample; these measurements are used in \cite{Acacia}, combined with the BAO distance measurements and RSD growth measurements of \cite{BeutlerDR12BAO,BeutlerDR12RSD,GriebDR12RSD,SanchezDR12RSD,SatpathyDR12RSD,VargasDR12BAO} and using the methods described in \cite{SanchezDR12comb} to constrain cosmological models. In Section \ref{sec:disc}, we compare our BAO results with those obtained from other BOSS studies and make general recommendations for how to consider any residual observation systematic uncertainty when using BOSS clustering results. Unless otherwise noted, we use a flat $\Lambda$CDM cosmology given by $\Omega_m = 0.31$, $\Omega_bh^2 = 0.0220$, $h=0.676$. This is consistent with \cite{Planck2015} and is the same as used in the companion papers studying the BOSS combined sample.

\label{sec:disc} \begin{figure} \includegraphics[width=84mm]{DA-H-BAODR12_zbin1.pdf}\vspace*{-0.3em} \includegraphics[width=84mm]{DA-H-BAODR12_zbin3.pdf} \caption{The allowed 1 and 2$\sigma$ regions (black ellipses) in the Hubble parameter, $H$, and the angular diameter distance, $D_A$, determined from our post-reconstruction anisotropic BAO scale measurements using BOSS galaxies with $0.2 < z < 0.5$ (top panel) and with $0.5 < z < 0.75$ (bottom panel). The colored points represent the 2$\sigma$ allowed region when assuming a flat $\Lambda$CDM cosmology and the the Planck 2015 results, with different colors representing the value of $H$ at $z=0$ (as indicated by the color bar on the right). } \label{fig:DAH} \end{figure} \subsection{Comparison to other DR12 BAO measurements} The final output of this work is the BAO measurements using the post-reconstruction, anisotropic correlation function measurements of the BOSS DR12 galaxy sample in redshift bins $0.2 < z < 0.5$, $0.4 < z < 0.6$, and $0.5 < z < 0.75$. Other studies have made similar measurements using DR12 data. \cite{CuestaDR12} obtained BAO measurements using the post-reconstruction anisotropic correlation function of the DR12 CMASS and LOWZ samples. In our robustness checks, we made the same measurements for the CMASS sample. Accounting for the difference in the fiducial cosmologies assumed by each analysis, the differences between \cite{CuestaDR12} and ours are 0.018 for $\alpha_{||}$ and -0.011 for $\alpha_{\perp}$. However, once we adjust to use the same bin size (8$h^{-1}$Mpc) as \cite{CuestaDR12}, the differences reduce to 0.011 for $\alpha_{||}$ and -0.004 for $\alpha_{\perp}$. Each of these represent a difference of less than 0.5$\sigma$ and are likely due to small methodological differences in the BAO fitting. We find smaller uncertainties on $\alpha_{||}$ (for both the data and the mocks) due to these differences. Both \cite{BeutlerDR12BAO} and \cite{VargasDR12BAO} obtain BAO measurements for the same post-reconstruction data set and redshift bins as we use. \cite{BeutlerDR12BAO} is a Fourier space analysis. Analyzing the same set of mocks, we find our results are correlated with a factor 0.9 and that the differences we obtain on the BOSS data are consistent with this high level of correlation. Both recover nearly identical uncertainties on the anisotropic BAO parameters, for both the data and the mock samples. \cite{VargasDR12BAO} uses the same configuration space data as presented in this study, but apply slightly different methodology to obtain their BAO measurements; they recover results that are consistent with ours. A more detailed comparison of these results is presented in \cite{Acacia}, where consensus sets of BOSS DR12 BAO and BOSS DR12 BAO + RSD measurements, combined as described in \cite{SanchezDR12comb}, are presented. \subsection{Comparison with $\Lambda$CDM} Our measurements of $\alpha_{||}$ and $\alpha_{\perp}$ can be translated into constraints on $D_A(z)(r_{\rm d}^{\rm fid}/r_{\rm d})$ and $H(z)(r_{\rm d}/r_{\rm d}^{\rm fid})$ and thereby test cosmological models. Here, we simply compare our measurements with the allowed parameter space in $\Lambda$CDM as determined by \cite{Planck2015}\footnote{Specifically, the results from the 'base\_plikHM\_TT\_lowTEB\_lensing' chains.}. This is show in Fig. \ref{fig:DAH} for the $0.2 < z < 0.5$ and $0.5 < z < 0.75$ redshift bins. Our low redshift result is fully consistent with the Planck $\Lambda$CDM prediction. Our high redshift result is in slight tension, as the 1$\sigma$ contours just barely overlap; this is mostly driven by the $H(z)$ measurement. This is similar to what was found in \cite{alph} for the DR11 CMASS data; the agreement is slightly better in \cite{BeutlerDR12BAO} and significantly better (to the level there is no tension) when these two post reconstruction results are optimally combined with pre-reconstruction full-shape results in \cite{Acacia}. Our results for the $0.4 < z < 0.6$ redshift slice (not plotted) are consistent with the Planck $\Lambda$CDM prediction, as one would predict based on the mean of the $0.2 < z < 0.5$ and $0.5 < z < 0.75$ results. The full cosmological context of our measurements, when combined with other BOSS DR12 results, is explored in detail in \cite{Acacia}.

1607.03145

1607

1607.08061_arXiv.txt

{The origin of the high-energy flux of neutrinos detected by IceCube is still unknown. Recent works report the evidence for a possible positional correlation between the reconstructed neutrino arrival directions and the positions in the sky of low power, high-energy emitting BL Lac objects (HBL).} {Assuming that $\gamma$-ray emitting HBL form the bulk of the sources of high-energy neutrinos above 100 TeV, we intend to calculate the number of events expected to be detected for each source by IceCube and KM3NeT.} {Based on a simple theoretically-motivated framework inspired by the structured jet scenario for these sources, we postulate a direct proportionality between high-energy $\gamma$-ray and neutrino fluxes. We calculate the expected neutrino event rate for the HBL sources of the Second Fermi-LAT Catalog of High-Energy Sources (2FHL) for IceCube and the presently under construction KM3NeT using declination-dependent and exposure-weighted effective areas.} {We provide a list of 2FHL HBL with the calculated number of events. For IceCube, the derived count rate for several sources is relatively high, of the order of $\lesssim$1 yr$^{-1}$, consistently with the recent findings of a possible positional correlation. For KM3NeT the calculated rates are higher, with several sources with expected rate exceeding 1 yr$^{-1}$. This, coupled with the improved angular resolution, implies that the HBL origin can be effectively tested with few years of observation of KM3NeT (and IceCube Gen2, for which similar performances are foreseen) through the direct association of neutrinos and single HBL.} {Our results show that if -- as hinted by recent works -- HBL represent a possible population of high-energy neutrino emitters, several single sources should be identified in few years of exposure of KM3NeT, highlighting the importance of the improved angular resolution anticipated for KM3NeT and IceCube Gen2.}

The cosmic sources responsible for the extraterrestrial neutrino flux detected by IceCube at PeV energies (Aartsen et al. 2013, 2014, 2015a) are still unknown. The substantial isotropy of the flux (with only a non significant small excess in the direction of the galactic center) is consistent with an extragalactic origin, although a slight north-south intensity and hardness asymmetry could hint to a possible contribution from a soft galactic component, superseded by a harder extragalactic component emission above $\approx 100$ TeV (e.g. Ahlers \& Murase 2014, Neronov \& Semikoz 2015, Aartsen et al. 2015, Palladino \& Vissani 2016). Among the possible extragalactic astrophysical sources there are propagating comic rays (e.g., Essey et al. 2010, Kalashev et al. 2013), star-forming and starburst galaxies (e.g., Loeb \& Waxman 2006, Wang, Zhao, \& Li 2014, Tamborra et al. 2014), galaxy clusters (e.g., Murase \& Beacom 2013, Zandanel et al. 2015), $\gamma$--ray burst (e.g., Waxman \& Bahcall 1997, Petropoulou et al. 2014) and active galactic nuclei (AGN, e.g., Mannheim 1995, Atoyan \& Dermer 2003, Kimura et al. 2014, Kalashev et al. 2014, Petropoulou et al. 2015, 2016). Among AGN, blazars (e.g. Urry \& Padovani 1995) are often considered the most probable candidates. Blazars (further divided in flat spectrum radio quasars, FSRQ, and BL Lac objects) are AGN presenting two jets ejected at relativistic speeds in opposite directions, one of which is well aligned with the line of sight to the Earth. In this geometry, relativistic effects greatly enhance the observed intensity (relativistic beaming), making these sources among the brightest extragalactic sources. Because of the beaming, the emission observed from blazars (extending over the entire electromagnetic spectrum with a characteristic double-humped shape when plotted in the $\nu F_{\nu }$ representation - the so-called spectral energy distribution, SED) is dominated by the non-thermal continuum produced in the jet. Leptonic models attribute the entire emission to relativistic electrons/pairs radiating through synchrotron and inverse Compton mechanisms, responsible for the low and the high energy SED peaks, respectively (e.g. Ghisellini et al. 1998). In the hadronic scenario, instead, the high-energy peak is linked to high-energy hadrons co-accelerated with electrons, cooling through the synchrotron or the photo-meson channel (e.g. B{\"o}ttcher et al. 2013). Blazars jets appear ideal sites to accelerate hadrons (protons, for simplicity) to the energy $E_{\rm p}\approx 10^{17}$ eV required to produce PeV neutrinos, most likely via the photomeson reaction ($p+\gamma \to X +\pi$), followed by the prompt decay of the charged pions ($\pi ^{\pm}\to \mu^{\pm}+\nu_{\mu}\to e^{\pm} + 2\nu_{\mu} +\nu_{\rm e}$; hereafter we do not distinguish among $\nu$ and $\bar{\nu}$). In fact the possible role of blazars has been recently highlighted by the results of Kadler et al. (2016) and Padovani et al. (2016). Kadler et al. (2016) report a tempting correlation between the arrival time of one of the neutrino with the highest reconstructed energy ($\sim 2$ PeV) and an exceptional outburst phase of the FSRQ PKS B1414-418, which lies in the (large, radius $\sim 16^{\circ}$) IceCube uncertainty region for this event. Padovani et al. (2016) on the other hand, improving a previous work by Padovani \& Resconi (2014), have presented the evidence for a significant (random expectation level of $\approx 0.4\%$) spatial correlation between the reconstructed arrival direction of neutrinos (including both hemispheres) and BL Lac objects emitting very high-energy $\gamma$ rays ($>50$ GeV). No correlation is instead found with other classes of blazars, such as FSRQ or BL Lacs with larger luminosity. Taken together, these two results are quite intriguing and puzzling, since powerful FSRQ and high-energy emitting BL Lacs (hereafter HBL, standing for highly peaked BL Lac objects), are objects characterized by rather different physical properties that lie at the opposite sides of the so-called blazar sequence (Fossati et al. 1998), relating the spectral properties of the emission of blazars with their luminosity. At a first sight, the FSRQ environment seems to offer the best conditions (high jet power, dense target radiation fields) to account for neutrino production through the photomeson channel ($pp$ reactions are unlikely in the low-energy jet environment), while BL Lac seem disfavored, mainly because their low luminosity hints to inefficient photomeson production (e.g. Murase et al. 2014). However, in a previous paper (Tavecchio, Ghisellini \& Guetta 2014, hereafter Paper I), we showed that, if the jet is characterized by a velocity structure, i.e. the flow is composed by a fast spine surrounded by a slower sheath (or layer), the neutrino output from HBL can be highly boosted with respect to the one-zone models and the cumulative emission of the HBL population could match the observed intensity with an acceptable value of the cosmic ray power for the jet. The existence of a velocity structure of the jet has been previously considered as a possible solution for several issues related to TeV emitting BL Lacs and to unify the BL Lacs and radiogalaxy populations (e.g. Chiaberge et al. 2000, Meyer et al. 2011, Sbarrato et al. 2014). Direct radio VLBI imaging of jets both in low-power radiogalaxies (e.g. Nagai et al. 2014, M{\"u}ller et al. 2014) and BL Lac (e.g., Giroletti et al. 2004, Piner \& Edwards 2014), often showing a ``limb brightening" transverse structure, provides a convincing observational support to this idea, also corroborated by numerical simulations (e.g. McKinney 2006, Rossi et al. 2008). The increased neutrino (and inverse Compton $\gamma$-ray) production efficiency in the spine--layer structure is based on the fact that for particles flowing in the faster region the radiation field produced in the layer is amplified by the relative motion between the two structures (e.g. Ghisellini et al. 2005, Tavecchio \& Ghisellini 2008). In this condition, the density of the soft photons in the spine rest frame -- determining the proton energy loss rate and hence the neutrino luminosity -- can easily exceed that of the radiation produced locally, the only radiative component considered in the one--zone modeling of BL Lacs (for FSRQ, instead, the photon field is thought to be dominated by the radiation coming from the external environment). In Tavecchio \& Ghisellini (2015) we relaxed the condition that only HBL jets are able develop an important layer, assuming that all BL Lacs jets are characterized by a structure region and that the layer radiative luminosity and the cosmic ray power are both proportional to the jet power. In searching for a {\it direct} association between neutrinos and possible sources one can exploit the temporal coincidence between the neutrino detection and high-state/flares of a source (e.g., Kadler et al. 2016, Halzen \& Kheirandish 2016) and/or the coincidence between the reconstructed arrival direction of neutrinos and the position of a putative source in the sky (e.g. Padovani et al. 2016). The latter works best when applied to the events detected through up-going muons, which provide the best angular resolution. The practical application of this methods is generally based on the use of a pre-selected list of possible neutrino source candidates (e.g. Adrian-Martinez et al. 2016a). Along these lines, in this paper, motivated by the recent results by Padovani et al. (2016), we aim at reconsidering the possible production of neutrinos by HBL, focusing in particular on the fluxes expected for present (IceCube) and future (KM3NeT) neutrinos observatories. Assuming a simple phenomenological framework inspired by the spine-layer scenario and supported by the Padovani et al. findings (\S2), we connect the putative neutrinos fluxes to the observed high-energy gamma-ray fluxes (\S3) and then we provide the expected neutrinos counts (\S4). Throughout the paper, the following cosmological parameters are assumed: $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\rm M}=0.3$, $\Omega_{\Lambda}=0.7$. We use the notation $Q=Q_X \, 10^X $ in cgs units.

In this work we have presented a heuristic framework to connect the $\gamma$-ray flux produced through the inverse Compton inside a structured jet of an HBL and the (hypotetical) neutrino flux. The scheme is motivated by the findings of Padovani et al. (2016) and is somewhat different from that discussed in Tavecchio et al. (2015). The latter was based on the use of the low-energy gamma-ray band as a proxy for both the cosmic ray power and the density of the photon targets, resulting in a quadratic dependence of the neutrino luminosity on the gamma-ray luminosity. We remark that an important feature of the present scheme, i.e. the dependence on the gamma-ray {\it flux} makes it possible to derive neutrino fluxes also for BL Lacs without a secure redshift measurements. This is quite important, since about 50\% of the HBL of the 2FHL have uncertain $z$. We also note that, although based on a specific model -- assuming a structured BL Lac jet --, the linear correlation found between $\gamma$-ray and neutrino fluxes have been already suggested in the past for blazars (e.g. Halzen \& Hooper 2005, Neronov \& Ribordy 2009). We have derived the expected number of muon neutrino for the HBL of the 2FHL catalogue for both IceCube and KM3NeT. We have provided a list of sources and expected numbers. Our study is focused on the through-going $\nu_\mu$ because of the angular resolution is well defined in the detectors. Our analysis takes into account the structural differences between the detectors. We have used the effective area at different declinations for IceCube and the effective area for muon neutrino for all declinations for KM3NeT. A great difference between the two detector is their latitude; IceCube is located at the South Pole, therefore the sources always have the same visibility throughout the year. Different is the case of KM3NeT, that have a range of declinations for which the sources are only partially visible during the year. For this reason we have considered for our calculation the visibility as a function of source declination for the muon-track analysis for tracks below the horizon and up to $10^\circ$ above the horizon, given by KM3NeT collaboration. We have calculated the expected number of neutrino from HBL both for tracks below the horizon and for tracks up to $10^\circ$ above the horizon; the difference between the two values is roughly a factor 1.2. A more detailed study will be done when the effective area to the various declinations for KM3NeT will be available. With our calculation we derive for IceCube fluxes consistent with the observations, predicting that only for few, $\gamma$-ray bright, BL Lacs we expect a handful of neutrinos detectable in few years of operation. The majority of the sources, instead, have fluxes implying rates of the order of $\lesssim$0.1 events yr$^{-1}$, for which a clear association is thus problematic. For KM3NeT, on the other hand, we foresee an appreciable neutrino flux for several sources. We report 20 BL Lacs for which the expected rate is $>0.3$ events yr$^{-1}$. For the brightest sources (Mkn 421, PKS 2155--304, Mkn 501), the event rate would likely be high enough to allow a firm identification. By construction, our method provides {\it average} fluxes. However, considering a typical flare of HBL, with a factor $\alpha \gtrsim$10 increase of the gamma (and thus neutrino) flux, lasting for $T_{\rm flare}\sim $1-2 weeks, the neutrino expected to be detected during the flare will be $ N_\nu = \alpha \bigl ( \frac{T_{\rm flare}}{1 \, {\rm yr}} \bigr )\tilde N_{\nu}$ where $\tilde N_{\nu}$ is the annual neutrino counts. This implies that for the handful of sources with an annual count $\tilde N_{\nu} \sim 1$ it would thus be possible to obtain one or more neutrinos concomitantly with the $\gamma$-ray flare. This detection would provide a clear signature that HBL can produce neutrinos. We would like to remark that, as far as the identification of the sources is concerned, KM3NeT and the proposed upgraded IceCube Gen2 (Aartsen et al. 2014b) are expected to play a quite valuable role. In particular, both are expected to have an improved (sub-degree) angular resolution for through-going muon neutrinos\footnote{The preliminary estimate for Km3NeT is ($<0.2^\circ$) (Adri{\'a}n-Mart{\'{\i}}nez et al.2016b).}, which will greatly help studies of correlation between the direction of the neutrino revealed and an extragalactic (or galactic) source. Moreover, having two instruments covering both hemispheres it will be possible to investigate better possible south-north anisotropies and spectral differences. The structured jet model that we adopt is based on the assumption that the emission we observe from HBL is (almost) totally produced by leptons through synchrotron and IC mechanisms (although it is applicable to all cases in which one predicts a linear relation between neutrino and $\gamma$-ray fluxes). Protons (or hadrons) are only responsible for the observed neutrino flux. The accompanying UHE $\gamma$-ray photons (from $\pi^0$ decay and emitted by the $e^{\pm}$ pairs from the charged pions decay) are readily reprocessed through electromagnetic cascades, leaving the sources as a low-level MeV-GeV component. This is different from what is instead envisaged in lepto-hadronic models (e.g. Petropoulou et al. 2015, 2016), predicting a luminous, hard MeV-GeV emission. Indeed, observations in the hard-X-ray band by the {\it NuSTAR}) satellite (sensitive up to 80 keV), revealing a steep continuum up to the highest energies, seem to leave small room for this bright hard-X/soft gamma component (e.g. Balokovi{\'c} et al. 2016 for Mkn 421), expected to be have a luminosity not much below that of the observed high-energy peak. Future instruments sensitive in the MeV band will play a key role in clarifying this issue. In particular, the proposed {\it e-ASTROGAM} mission\footnote{\tt http://astrogam.iaps.inaf.it} is foreseen to provide a sensitivity of the order of $\lesssim10^{-12}$ erg cm$^{-2}$ s$^{-1}$ in the band 0.3 MeV--3 GeV, where the bulk of the reprocessed emission is predicted. With such a sensitivity {\it e-ASTROGAM} would be able to detect the reprocessed emission even in the case of moderately bright HBLs.

1607.08061

1607

1607.03235_arXiv.txt

In this paper we present a simple discussion on the properties of compact stars using an EoS obtained in effective field theory anchored on scale and hidden-local symmetric Lagrangian endowed with topology change and a unequivocal prediction on the deformation of the compact star, that could be measured in gravitational waves. The objective is not to offer a superior or improved EoS for compact stars but to confront with a forthcoming astrophysical observable the given model formulated in what is considered to be consistent with the premise of QCD. The model so obtained is found to satisfactorily describe the observation of a 2-solar mass neutron star \cite{demorest,antoniadis} with a minimum number of parameters. Specifically the observable we are considering in this paper is the tidal deformability parameter $\lambda$ (equivalently the Love number $k_2$), which affects gravitational wave forms at the late period of inspiral stage. The forthcoming aLIGO and aVirgo observations of gravitational waves from binary neutron star system will provide a valuable guidance for arriving at a better understanding of highly compressed baryonic matter.

The observation of 2-solar mass neutron stars\cite{demorest,antoniadis} seems to indicate that the equation of state(EoS) for compact stars needs to be sufficiently stiffer to accommodate the mass larger than 1.5-solar mass. Moreover it requires the detailed information on the hadronic matter at higher density than the normal nuclear density, $n_0$, which seems to be however much higher beyond the reach of presently planned terrestrial laboratories. After the recent detections of gravitational waves from binary black holes\cite{PRL}, the expectation of detecting gravitational waves from a binary neutron stars and/or a black hole-neutron star binary becomes very optimistic than ever. The gravitational waves emitted during the binary inspiral phase up to merging can provide us the information on the dense hadronic matter of higher density at the core of compact stars. The nuclei with large atomic numbers are already obvious examples of highly dense matters composed of nucleons, $n \sim n_0 = 0.16/\textrm{fm}^{-3}$. The effective theories of nucleons have been developed and constrained by experimental data available up to and slightly above the normal nuclear density, $n_0$. Hence they are fairly well controlled theoretically and experimentally. but the high density regime much above $n_0$ is more or less uncharted both experimentally and theoretically. Compact objects such as neutron stars are supposed to have higher core density than the normal nuclear density, $n_0$. Roughly for neutron star with mass $\sim 1.5 M_{\odot}$ the relevant density at the center is likely around $2 n_0-3 n_0$ and for the mass $\sim 2 M_{\odot}$ the density is supposed to be larger than $\sim 5n_0$. On top of the possibilities of getting high density nuclear matter at terrestrial laboratories, for example FRIB(USA), FAIR(Germany), J-Parc(Japan) and RAON(Korea) in near future, the possible detections of gravitational waves from binary neutron stars(or binary neutron star black hole) are believed to be promising probes of the high density interior of neutron stars. Theoretically the success of low density effective theories univocally up to the normal nuclear density seems not to guarantee the similar success at higher density hadronic matter since the predictions of mass, radius, symmetry energy and deformability to name a few diverse from each other drastically beyond the normal nuclear density. In this sense we are now at the very exciting period of foreseeing the opportunity of constraining theories at higher density by experiments and observations. For the highly dense hadronic matter, recently we proposed a new approach\cite{Dong,LPR,Paeng} in which we try to formulate a field theory framework wherein both low and high density regimes are treated on the same footing. To cover both regimes in a consistent way in a unified field theoretic approach seems like a tall order. In this work, we discuss the physical properties of stellar matter using the newly proposed scheme of a new scaling law(BLPR) in medium\cite{byp,lr} all the way from normal nuclear density up to the higher density at the core of a neutron star. With a confrontation with the observed massive stars\cite{Dong}, we discuss the physical properties of stellar matter with the EoS obtained therein. The relevant quantities are mass, radius and deformability parameters, which could be constrained by the gravitational wave forms emitted during the binary inspiraling phase. The aim here is then to confirm or falsify the strategies taken and assumptions made in Dong {\it et al.}\cite{Dong} and Lee {\it et al.}\cite{LPR} and also to find the directions to be taken in constructing the correct effective theory at higher density. In section II, the basic concept of unified approach in this work is discussed. Using the minimal effective lagrangian in the frame work of relativistic mean field, the equation of state of compact star with neutrons, protons, electrons and muon in weak equilibrium and charge neutrality condition is discussed in section III and the mass and radius are estimated. Tidal deformation with new stiffer EoS is discussed with the observational possibility in gravitational wave detections at aLIGO\cite{aligo} and aVirgo\cite{avirgo} in section IV. The results are summarized in section V. We use units in which $c=G=1$ and the notation in which Minkowski metric $\eta_{\mu\nu} = \mathrm{diag}[-1,\ 1,\ 1,\ 1]$.

We discussed the physical properties of stellar matter with a new stiffer EoS, which has been proposed recently using a new scaling law (new-BR/ BLPR) in medium caused by topology change at high density~\cite{Dong}, by extending Dong {\it et al.}'s work for pure neutron matter to a realistic nuclear matter of $n$, $p$, $e$ and $\mu$. The mass- radius and the tidal deformability were calculated. The calculated maximum mass of compact star is found to be about $2 M_\odot$ with its radius about 11km. The radius for the mass range of $1M_{\odot}$--$ 2 M_{\odot}$ is found to be $11.2$--$12.2 $km. The calculated deformability parameter for the stiffer EoS employed in this work is in the range $ 4.0$--$ 0.68$. What characterizes the approach presented in this work is the stiffening of the EoS due to topology change predicted in the description of baryonic matter with skyrmions put on crystal background to access high density. The change is implemented in the properties of the parameters of the effective Lagrangian anchored on chiral symmetry and manifests in nuclear EFT formulated in terms of RG-implemented $V_{lowk}$. Given that the approach describes fairly well the baryonic matter up to normal nuclear density, it is the changeover of skyrmions to half-skyrmions at a density $\sim$ (2--3)$n_0$ that is distinctive of the model used. This topology change involves no change of symmetries -- and hence no order parameters, therefore it does not belong to the conventional paradigm of phase transitions. But it impacts importantly on physical properties as described in various places in a way that is not present in standard nuclear physics approaches available in the literature. It is interesting to note that, as has been discussed recently~\cite{hatsuda,baym}, there is another way to produce the stiffening in EoS to access the massive compact stars. It is to implement a smooth changeover from hadronic matter -- more or less well-described -- to strongly correlated quark matter, typically described in NJL model. By tuning the parameters of the quark model so as to produce a changeover at a density $\gsim 2 n_0$, it has been possible to reproduce the features compatible with the properties of observed massive stars. After the detections of gravitational waves binary black holes, the possible detection of gravitational wave signals from coalescing binary neutron stars are well expected during the next run of aLIGO and aVirgo and the detection will inform us of the detailed effect of the tidal deformation\cite{PRD/77/021502,PRD/81/123016,read,HKSS}. Recently the tidally modified waveforms have been developed up to the high frequency of merger\cite{read,bernucci}, such that the deformability parameter $\lambda$ , a function of the neutron-star EOS and mass, is measurable within the frequency range of the projected design sensitivity of aLIGO and aVirgo. It has been also demonstrated in Bayesian analysis that the tidal deformability can be measured to better than $\pm 1 \times 10^{36}$ g cm$^2$ s$^2$ when multiple inspiral events from three detectors of aLIGO-aVirgo network~\cite{aLIGO,aVirgo} are analyzed\cite{Lackey}. They also show that the neutron star radius can be measured to better than $\pm 1$ km. Thus the simultaneous measurement of mass, radius and deformability using gravitational wave detectors could present an exciting possibility to eventually pin down the highly uncertain EoS for the nuclear matter in the mass range of $1M_{\odot}$--$2 M_{\odot}$. This would provide a probe for the state of baryonic matter at the high density that is theoretically the most uncertain. And the hope is whether one can confirm or falsify the strategies taken and assumptions made in \cite{Dong,LPR} and whether the result would then help point the directions to be taken in the efforts described in \cite{Dong,Paeng}. {\it Hyun Kyu Lee}: ``When Gerry invited me to Stony Brook in 1998 for my sabbatical year, he put me in a house just next to his. One late afternoon he came to our place with a big smile and a basket of potatoes he just dug out in his yard. He had keen interest in hearing the news of detecting gravitational waves, which was one of his favorite laboratories up in the sky. I am now missing his big smile and a basket of comments on recent detections of gravitational waves, GW150914 and GW151226."

1607.03235

1607

1607.03003_arXiv.txt

Dark energy models and modified gravity theories have been actively studied and the behaviors in the solar system have been also carefully investigated in a part of the models. However, the isotropic solutions of the field equations in the simple models of dark energy, e.g. quintessence model without matter coupling, have not been well investigated. One of the reason would be the nonlinearity of the field equations. In this paper, a method to evaluate the solution of the field equations is constructed, and it is shown that there is a model that can easily pass the solar system tests, whereas, there is also a model that is constrained from the solar system tests.

} The accelerated expansion of the Universe was discovered in the late 1990s from the observations of type Ia supernovae \cite{Riess:1998cb,Perlmutter:1998np}. Many hypotheses about the cause of the accelerated expansion have been proposed, however, it is still unclear what is the real cause of it. The representative way to explain the accelerated expansion of the Universe are the following; introducing the cosmological constant $\Lambda$, introducing a dynamical scalar field instead of $\Lambda$, modifying the geometry part of the Einstein equations. They are called dark energy models or modified gravity theories. We will treat $k$-essence model \cite{ArmendarizPicon:1999rj,Garriga:1999vw,ArmendarizPicon:2000ah}, which is a scalar field model of dark energy and contains quintessence model \cite{Peebles:1987ek,Ratra:1987rm,Chiba:1997ej,Zlatev:1998tr} as a special case, in this paper. $k$-essence model has only minimal coupling with gravity, the potential term of the field, and the kinetic terms of the field. Scalar field models of dark energy are defined not by the field equations but by their action. Therefore, not only homogeneous, isotropic, and expanding solutions but also static spherical solutions are affected by the scalar field. The influences for the static spherical solution from dark energy have been considered in the models that have a nonminimal coupling with the usual matter \cite{Khoury:2003rn} or have a nonminimal coupling with gravity \cite{Will:2014kxa}, however, they have not been enough investigated in $k$-essence model without matter coupling. One of the reason may come from an assumption that there are no serious influences to gravity from a minimal couping. In fact, the gravitational effects from the cosmological constant in the solar system certainly exist, however, they are too small to be observed \cite{Kagramanova:2006ax}. While, the dynamical scalar field under the Friedmann-Lemeitre-Robertson-Walker (FLRW) space-time varies depending on the radius $r$ in a spherically symmetric background. Therefore, we should carefully take into account the influences from the scalar field because they can be large enough to be observed in the solar system even if they are same order as those from the cosmological constant at the horizon scale of the Universe. The other reason would be caused from the nonlinearity of the field equation. The analytic solutions for the nonlinear differential equation do not generally exist. The reason why the investigations in more complex theories are possible is that many assumptions are applied. It is important to inspect whether or not the assumptions are valid in a simpler model. In this paper, we will consider the behavior of the scalar field in a static and spherical space-time in $k$-essence model without matter coupling. The contents of the paper are as follows. We deduce the general equations in a static and spherical background, and construct a general formalism to evaluate the solution in Sec.~\ref{sec2}. The solutions of the equations in quintessence model are investigated in Sec.~\ref{sec3}, and those in $k$-essence model are considered in Sec.~\ref{sec4}. Conclusions are in Sec.~\ref{sec5}. The units of $k_\mathrm{B} = c = \hbar = 1$ are used and gravitational constant $8 \pi G$ is denoted by ${\kappa}^2 \equiv 8\pi/{M_{\mathrm{Pl}}}^2$ with the Planck mass of $M_{\mathrm{Pl}} = G^{-1/2} = 1.2 \times 10^{19}$GeV in this paper.

} We have considered the behavior of scalar field dark energy with a minimal coupling in a static isotropic background. In Sec.~\ref{sec2}, we have derived the general equations in $k$-essence model, and have constructed a general formalism to investigate the behaviors of the scalar field and the metric functions in the case $\rho _\mathrm{m}=p_\mathrm{m}=0 $. In Sec.~\ref{sec3}, we have considered quintessence model $K(\phi, X)=-X +V(\phi)$ and demonstrated the procedure shown in Sec.~\ref{sec2} in the case $w=const.$ and in the case of negative power law potential. The case $w=const.$ is approximately same as the cosmological constant deep inside the horizon scale of the Universe. Therefore, it has been shown that this model can easily pass the solar system tests. In the case of negative power law potential $V(\phi) \propto \phi ^{-n}$ $(n>0)$, we have obtained the particular solutions for the field equation in the region $r \ll 1/H_0$. The existence condition for the particular solutions and the constraint from the solar system test have shown that the power of $\phi$ should be $-1$, $-3$, or $-5$. While, the constraint for $n$ from Hubble parameter mesurement was investigated by O. Farooq \textit{et al.} \cite{Farooq:2012ev}. They showed that $n \lesssim 1$ at $3 \sigma$ level. Combining our result with their result give a tight constraint on $n$: $n=1$. If $n \neq 1$, the model is observationally rejected or does not have a spherically symmetric solution. In Sec.~\ref{sec4}, we have shortly considered $k$-essence model that only consists of kinetic terms of the scalar field. The field equation can be strictly solved in the limit that the higher derivative terms are dominant or the lower derivative terms are dominant. However, the expansion parameter $s_2$ and $s_3$ are severely constrained by the conditions $\vert \delta \lambda \vert , \vert \delta \Phi \vert \ll 1$ and $\vert s_2 X \vert , \vert s_3 X/s_2 \vert \gg 1$. If we only use either $\vert \delta \lambda \vert , \vert \delta \Phi \vert \ll 1$ or $\vert s_2 X \vert , \vert s_3 X/s_2 \vert \gg 1$, there is no constraint on $s_2$ and $s_3$. Therefore, evaluating the conditions $\vert \delta \lambda \vert , \vert \delta \Phi \vert \ll 1$ is imperative.

1607.03003

1607

1607.04415_arXiv.txt

This is the second paper in a series on a new luminous $z \sim 5$ quasar survey using optical and near-infrared colors. Here we present a new determination of the bright end of the quasar luminosity function (QLF) at $z \sim$ 5. Combined our 45 new quasars with previously known quasars that satisfy our selections, we construct the largest uniform luminous $z \sim$ 5 quasar sample to date, with 99 quasars in the range 4.7 $\le z <$ 5.4 and $-29 < M_{1450} \le -26.8$, within the Sloan Digital Sky Survey (SDSS) footprint. We use a modified 1/V$_{\rm a}$ method including flux limit correction to derive a binned QLF, and we model the parametric QLF using maximum likelihood estimation. With the faint-end slope of the QLF fixed as $\alpha = -2.03$ from previous deeper samples, the best fit of our QLF gives a flatter bright end slope $\beta = -3.58\pm0.24$ and a fainter break magnitude $M^{*}_{1450} = -26.98\pm0.23$ than previous studies at similar redshift. Combined with previous work at lower and higher redshifts, our result is consistent with a luminosity evolution and density evolution (LEDE) model. Using the best fit QLF, the contribution of quasars to the ionizing background at $z \sim$ 5 is found to be 18\% $-$ 45\% with a clumping factor $C$ of 2 $-$ 5. Our sample suggests an evolution of radio loud fraction with optical luminosity but no obvious evolution with redshift.

Quasars comprise the most luminous class of non-transient objects in the universe. Characterizing their population and evolution is the critical tool to constrain directly the formation and evolution of supermassive black holes (SMBHs) across cosmic time. The fundamental way to characterize these objects is through the evolution of their number density with luminosity and redshift, namely the quasar luminosity function (QLF). The QLF and its cosmological evolution have been a key focus of quasar studies for half a century. \cite{schmidt68} first determined the evolution of the quasar population and found the first evidence for a significant increase of the quasar number density with redshift in both radio and optical bands. More recently, based on measurements of the QLF from several successful surveys, such as the 2dF Quasar Redshift Survey \citep{boyle00, croom04}, COMBO-17 \citep{wolf03}, the 2dF-SDSS LRG and QSO survey \cite[2SLAQ;][]{richards05}, the SDSS Faint Quasar Survey \citep{jiang06}, the VIMOS-VLT Deep Survey \cite[VVDS;][]{bongiorno07}, SDSS and 2SLAQ \citep{croom09} and BOSS DR9 \citep{ross13}, the QLF, especially in optical bands, has been well characterized at low to intermediate redshifts. The QLF can be parameterized with a double power law shape and pure luminosity evolution for quasars at redshifts up to $z$ = 2 \citep{boyle00, croom04}. The bright end slope at low redshift, the effect of $''$cosmic downsizing$''$ and the density peak of quasars at 2 $< z <$ 3 \citep{richards06, brown06, jiang06, croom09} have been confirmed by many subsequent investigations. The measurements based on large samples from BOSS yield a QLF evolution best fit by a luminosity evolution and density evolution (LEDE) model at 2 $< z <$ 3.5 \citep{ross13}. In their work, the bright end slope does not evolve with redshift and is different from the result of \cite{richards06}, which suggested a flatter bright end slope at high redshift than that at low redshift. To better determine the evolution of QLF parameters, a wider redshift range is needed. Towards higher redshift, quasars are important tracers of the structure and evolution of the early Universe, the evolution of the intergalactic medium (IGM), the growth of SMBHs and co-evolution of SMBHs and host galaxies at early epochs. Observations of the Gunn-Peterson effect using absorption spectra of quasars at $z \gtrsim$ 5.7 have established $z \sim$ 6 as the end of cosmic reionization, when the IGM is rapidly transforming from largely neutral to completely ionized \citep{fan06}. \cite{becker15} find evidence for UV background fluctuations at $z\sim$ 5.7 in excess of predictions from a single mean-free-path model, which indicates that reionization is not fully complete at that redshift. \cite{mcgreer15} suggest that reionization is just completing at $z\sim$ 6, possibly with a tail to $z\sim$ 5.5. Therefore, in the post-reionization epoch, the QLF at $z \gtrsim 5$ is needed to estimate the contribution of quasars to the ionizing background during and after the reionization epoch. Although quasars are not likely to be the dominant source of ionizing photons \citep{fan01a, willott10b, mcgreer13}, their exact contribution is still highly uncertain. In addition, $z\sim 5$ quasar absorption spectra can be used to constrain the physical conditions of the IGM in this key redshift range, and provide the basic boundary conditions for models of reionization, such as the evolution of IGM temperature, photon mean free path, metallicity and the impact of helium reionization \citep{bolton12}. However, high redshift quasars are very rare, especially at $z >$ 5. Although more than 300,000 quasars are known, only $\sim$200 of them are at $z >$ 5. Therefore, QLF measurements at high redshift still have large uncertainties. From the combination of SDSS DR7 quasars and the Stripe 82 (S82) faint quasar sample, \cite{mcgreer13} provided the most complete measurement of the $z \sim$5 QLF so far, especially at the faint end. A factor of 2 greater decrease in the number density of luminous quasars from $z$ = 5 to 6 than that from $z$ = 4 to 5 was claimed \citep[][hereafter M13]{mcgreer13}. However, their work focused on the faint end; there are only 8 quasars with $M_{1450} < -27.3$ in the sample. A survey described in this series of papers aims at finding more luminous quasars at 4.7 $< z <$ 5.5, which allows a better determination of the bright end QLF and a better constraint on the quasar evolution model at high redshift. \citet[][hereafter Paper I]{wang16} presented a new selection using SDSS and the Wide-field Infrared Survey Explorer (WISE) optical/NIR colors. In this followup paper, we report our measurement of the bright end $z \sim$ 5 QLF using the quasar sample selected by the method presented in Paper I. The outline of our paper is as follows. In Section 2, we briefly review the quasar candidate selection and the spectroscopic observations of these candidates. The survey completeness will be presented in Section 3. In this section, we use a quasar color model (M13) to quantify our selection completeness and to correct the incompleteness due to the ALLWISE detection flux limit and spectral coverage. We then calculate the binned luminosity function and fit our data using a maximum likelihood estimator in Section 4. We also study the evolution of the QLF and compare our results with previous work in this section. In Section 5, we discuss the contribution of $z \sim$ 5 quasars to the ionizing background and the radio loud fraction of our quasar sample. We summarize our main results in Section 6. In this paper, we adopt a $\Lambda$CDM cosmology with parameters $\Omega_{\Lambda}$ = 0.728, $\Omega_{m}$ = 0.272, $\Omega_{b}$ = 0.0456, and H$_{0}$ = 70 $km s^{-1} Mpc^{-1}$ \citep{komatsu09} for direct comparison with the result in M13. Photometric data from the SDSS are in the SDSS photometric system \citep{Lupton99}, which is almost identical to the AB system at bright magnitudes; photometric data from ALLWISE are in the Vega system. All SDSS data shown in this paper are corrected for Galactic extinction.

We establish a highly effective $z \sim$ 5 quasar selection method based on SDSS and ALLWISE optical/near-infrared colors. We relax the traditional $r-i/i-z$ color limit by including color cuts in the ALLWISE W1 and W2 data. We selected 110 quasar candidates that satisfied our selection criteria with good optical image quality and obtained spectroscopic observations for 99 candidates. 64 new quasars have been discovered in the redshift range $4.4 < z < 5.5$ and magnitude range $-29 < M_{1450} < $-$26.4$. We restrict our luminous quasar sample to $4.7 \le z < 5.4$ and $ M_{1450} \le $-$26.8$ for the QLF calculation. Combining all previously known quasars in this range, we construct the largest luminous quasar sample at $z \sim$ 5 and determine the QLF, covering a sky area of 14555 deg$^2$. Here we list our main conclusions. \begin{itemize} \item Within the redshift range $4.7 \le z < 5.4$ and magnitude range $ M_{1450} \le $-$26.8$, there are 45 newly identified quasars and 54 known quasars. Our new discovery successfully extends the population of luminous quasars at $z \sim$ 5, especially at $ M_{1450} \le $-$27.3$, where we discovered 27 new quasars and increased the number of known quasars by a factor of 1.5 in this luminosity range. Our final sample including 99 quasars is the largest sample of luminous $z \sim$ 5 quasars (Fig. 1). \item We derive the selection function of our color-color selection by using 311,000 simulated quasars in the redshift range $z$ = 4 to 6 and luminosity range $-$29.5 $< M_{1450} < $-25.5. The selection function shows that by relaxing the traditional $r-i$/$i-z$ color cut and adding the W1-W2 color, our color selection criteria extend the selection fuction to a higher redshift z$\sim$ 5.4 than previous work (Fig. 5). \item Using this sample, we calculate the binned QLF and fit the parametric QLF by using maximum likelihood fitting at $z = 5.05$ (Fig. 8 \& 9). For the parametric QLF, we fix the faint end slope $\alpha$ =-2.03, which is measured by using the S82 and DR7 quasar samples (M13), and find the best fit result of the bright end slope $\beta$ = $-$3.58$\pm$0.25 and break magnitude of $M_{1450}^{*}$ = $-$26.99 $\pm$0.23. \item We compare parameters of our best fit QLF with previous work at different redshifts and use all points to fit an LEDE model. Our result is consistent with the previous LEDE model but prefers a steeper slope of log$\Phi^{*}(z)$ evolution and a flatter slope of break magnitude evolution. The comparison for $\beta$ shows no clear evolution with redshift (Fig. 11). \item We calculate the contribution of quasars to the ionizing background at $z \sim$ 5 based on our QLF. Integrating our best fit QLF, we find that quasars are able to provide $\sim$ 18\% $-$ 45\% of the required photons based on a clumping factor $C \sim$ 2 $-$ 5. \item We use FIRST and NVSS data to calculate the radio loud fraction of our sample and give a lower limit for the RLF of $\sim$ 7.1\% which agrees with the result at $z \sim$ 6 of \cite{banados15}. In comparison with the predicted evolution function of the RLF with $M_{2500}$ and $z$ proposed by \cite{jiang07}, our result shows evolution with optical luminosity but no obvious evolution with redshift (Fig. 12). \end{itemize}

1607.04415

1607

1607.01837_arXiv.txt

As an alternative explanation for the cosmic acceleration, $f(R)$ theories of gravity can predict an almost identical expansion history to standard $\Lambda$CDM, yet make very different predictions for the growth of cosmological structures. Measurements of the cosmic bulk flow provides a method for determining the strength of gravity over the history of structure formation. We use the modified gravity N-body code ECOSMOG to simulate dark matter particles and make predictions for the bulk flow magnitude in both $\Lambda$CDM and $f(R)$ gravity. With the peculiar velocities output by ECOSMOG we determine the bulk flow at depths ranging from $20h^{-1}$Mpc to $50h^{-1}$Mpc, following the redshift and sky distribution of the 2MASS Tully-Fisher survey (2MTF). At each depth, we find that the $\Lambda$CDM and $f_{R0} = 10^{-5}$ simulations produce bulk flow measurements that are consistent with $\Lambda$CDM predictions and the 2MTF survey at a $1\sigma$ level. We also find that adopting an $f(R)$ strength of $f_{R0} = 10^{-3}$ predict a much larger value for the bulk flow, which disagree with $\Lambda$CDM predictions at all depths considered. We conclude that $f_{R0}$ must be constrained to a level no greater than $10^{-4}$ to agree with bulk flow measurements.

One of the most inescapable facts in recent cosmology is that the Universe is undergoing a period of accelerated expansion. The effect of this acceleration was observed through measurements of supernovae \citep{b1,b2}, confirming previous indications from large-scale structure and galaxy surveys \citep{Efstathiou,OstrikerSteinhardt,KraussTurner,YoshiiPetersen}. The source of this late-time acceleration has been named `Dark Energy' which exerts a negative pressure to combat the attractive force of gravity. Currently the simplest candidate for dark energy is the cosmological constant $\Lambda$. However theoretical calculations yield a value of $\Lambda$ at least 120 orders of magnitude larger than observations \citep{b4,b5}. As a result, cosmological models that do not include an explicit cosmological constant form an appealing alternative. These alternatives are usually categorized depending upon which side of the Einstein equations they alter. The first category adds to or alters the energy-momentum tensor $T_{\mu\nu}$ to yield a negative pressure (dark fluid models), while the second category alters the Einstein tensor $G_{\mu\nu}$ to generate the acceleration (modified gravity models). Throughout this paper we focus on a specific modified theory, $f\left(R\right)$ gravity. $f\left(R\right)$ gravity changes the gravitational theory by modifying the action, from the standard Einstein-Hilbert action, to be some new function of the Ricci scalar $R$ \citep{Nojiri,Carroll}. Given the freedom to choose the function $f\left(R\right)$, the expansion histories of both $\Lambda$CDM and $f\left(R\right)$ models can be very similar, or even identical \citep{Song}. Therefore we must consider alternate methods to observationally differentiate between the models. One such approach is to study the peculiar velocity of galaxies which result from the gravitational interaction between a galaxy and the surrounding matter, causing the galaxy redshift to deviate from Hubble's Law. In essence, the peculiar velocity of a galaxy is an integrated history of its gravitational interactions, and thus provides a tool to differentiate between $\Lambda$CDM and $f\left(R\right)$ models. Measuring peculiar velocities offers an observational difficulty as such measurements must be performed using redshift independent distance indicators such as type Ia Supernovae \citep{b12}, the Tully-Fisher relation \citep{b13} and the Fundamental Plane relation \citep{b14}. A common parameter that many peculiar velocity surveys quote is the net dipole, or the `bulk flow', of the peculiar velocity field. There has been much debate over whether the measured bulk flows are consistent with the $\Lambda$CDM model. \cite{b15} analyzed 2,018 galaxies from the 2MASS Tully-Fisher survey (2MTF) utilizing both $\chi^2$ and minimum variance methods, finding a bulk flow that is consistent with the $\Lambda$CDM model to a $1\sigma$ level. Conversely, \cite{b16} utilized a catalogue of 4,481 peculiar velocity measurements with a characteristic depth of $33h^{-1}$Mpc and claim that the resulting bulk flow is inconsistent with the $\Lambda$CDM model at a $>$98\% confidence level. \newline \indent A possible solution to these anomolaous bulk flow measurements is to adopt a modified theory of gravity. To this end, N-body simulations can be employed to evolve particles under both $\Lambda$CDM and modified gravity models. The results of these simulations can then be compared to surveys such as 2MTF. An added benefit of utilizing N-body simulations to measure bulk flow is the lack of underlying systematic biases that most surveys are subject to. This is especially important as \cite{b17} has shown that unaccounted systematic uncertainty could explain the discrepancies between surveys agreeing/disagreeing with the $\Lambda$CDM model. \newline \indent In this paper we utilize N-body simulations to measure bulk flow in both $\Lambda$CDM and $f\left(R\right)$ regimes. In Section 2 we outline $f\left(R\right)$ gravity and show how we quantify the deviation from the $\Lambda$CDM model. In Section 3 we give an overview of the simulations we use, how the output is utilized to calculate bulk flow and a brief outline of the 2MTF survey. In Section 4 we present the results of the simulations and compare them to the 2MTF survey. We conclude in Section 5. \newline \indent Throughout the paper we adopt a standard cosmology of $\Omega_m = 0.30$, $\Omega_\Lambda = 0.70$ and $H_0 = 100h$ km s$^{-1}$Mpc$^{-1}$. Whilst our results are $h$ independent, we use a value of $h = 0.70$ in our simulations.

} Modified gravity theories provide an appealing alternative to the $\Lambda$CDM model by providing a model that does not include an explicit cosmological constant. One such theory, $f\left(R\right)$ gravity, involves changing the Einstein-Hilbert action by altering the functional dependence upon the Ricci scalar. We consider one particular $f(R)$ model, the Hu \& Sawicki model \citep{b20}, and parameterise the degree of deviation from the predictions of standard $\Lambda$CDM through the gradient of the function today $f_{R0}=\frac{\partial f(R)}{\partial R}|_{t=t_0}$. The bulk flow is the net dipole moment of the cosmological peculiar velocity field, which is the result of gravitational influence on the motions of particles on large scales. As the bulk flow is sensitive to the gravitational theory considered, this results in a measurable difference in the predicted value of the bulk flow between $\Lambda$CDM and $f\left(R\right)$ gravity. We utilized N-body simulations to create a set of mock surveys under both $\Lambda$CDM and $f\left(R\right)$ gravity. These mocks were analyzed assuming the redshift and sky distribution of 2MTF, a previous survey that studied the bulk flow in the local Universe \citep{b15}. We found that the simulations under $\Lambda$CDM and $f_{R0} = 10^{-5}$ gravity produced bulk flows that were consistent with both 2MTF and $\Lambda$CDM predictions. Choosing $f_{R0} = 10^{-4}$ gave bulk flows that agreed with the $\Lambda$CDM predictions on small scales with weak agreement on large scales. Finally, setting $f_{R0} = 10^{-3}$ resulted in bulk flows that did not agree with $\Lambda$CDM predictions to a $1\sigma$ level at all scales considered. From these results we conclude that in order to obtain bulk flow measurements that match previous work \citep{b30, b28,b29,b27} in addition to theoretical predictions, the upper limit on $f_{R0}$ lies somewhere in the range $10^{-4}$ and $10^{-5}$. We finally note that, given the agreement between the $\Lambda$CDM and $f_{R0} = 10^{-5}$ predictions, it seems unlikely that bulk flow measurements can be of any further use in constraining the parameters in an $f(R)$ theory. We have already reached the theoretical limit in which the bulk flow amplitude will provide useful information. Instead to make further progress in this area, it is more advantageous to use the full velocity power spectrum \citep[e.g.][]{Johnson}.

1607.01837

1607

1607.08483_arXiv.txt

Valuable theoretical predictions of nuclear dipole excitations in the whole chart are of great interest for different nuclear applications, including in particular nuclear astrophysics. Here we present large-scale calculations of the $E1$ $\gamma$-ray strength function obtained in the framework of the axially-symmetric deformed quasiparticle random phase approximation based on the finite-range Gogny force. This approach is applied to even-even nuclei, the strength function for odd nuclei being derived by interpolation. The convergence with respect to the adopted number of harmonic oscillator shells and the cut-off energy introduced in the 2-quasiparticle (2-$qp$) excitation space is analyzed. The calculations performed with two different Gogny interactions, namely D1S and D1M, are compared. A systematic energy shift of the $E1$ strength is found for D1M relative to D1S, leading to a lower energy centroid and a smaller energy-weighted sum rule for D1M. When comparing with experimental photoabsorption data, the Gogny-QRPA predictions are found to overestimate the giant dipole energy by typically $\sim$2 MeV. Despite the microscopic nature of our self-consistent Hartree-Fock-Bogoliubov plus QRPA calculation, some phenomenological corrections need to be included to take into account the effects beyond the standard 2-$qp$ QRPA excitations and the coupling between the single-particle and low-lying collective phonon degrees of freedom. For this purpose, three prescriptions of folding procedure are considered and adjusted to reproduce experimental photoabsorption data at best. All of them are shown to lead to rather similar predictions of the $E1$ strength, both at low energies and for exotic neutron-rich nuclei. Predictions of $\gamma$-ray strength functions and Maxwellian-averaged neutron capture rates for the whole Sn isotopic chain are also discussed and compared with previous theoretical calculations.

About half of the nuclei with $A>60$ observed in nature are formed by the rapid neutron-capture process (r-process) occurring in explosive stellar events \cite{arnould07}. The r-process is expected in environments with high neutron density ($N_n > 10^{20}~{\rm cm^{-3}}$). Successive neutron captures proceed into neutron-rich regions well off the $\beta$-stability valley forming exotic nuclei that cannot be produced and therefore studied in the laboratory. When the temperature or the neutron density required for the r-process are low enough to break the (n,$\gamma$)-($\gamma$,n) equilibrium, the r-abundance distribution depends directly on the neutron capture rates of the so-produced exotic neutron-rich nuclei \cite{go98}. The neutron capture rates are commonly evaluated within the framework of the statistical model of Hauser-Feshbach, although the direct capture contribution play an important role for very exotic nuclei \cite{xu14}. The fundamental assumption of the Hauser-Feshbach model is that the capture go through the intermediary formation of a compound nucleus in thermodynamic equilibrium. In this approach, the Maxwellian-averaged (n,$\gamma$) rate at temperatures of relevance in r-process environments strongly depends on the electromagnetic interaction, i.e the photon de-excitation probability. The well known challenge of understanding the r-process abundances thus requires reliable extrapolations of the $E1$-strength function out towards the neutron-drip line. Large scale calculations of $E1$ $\gamma$-ray strength functions are usually performed on the basis of the phenomenological Lorentzian model \cite{ripl3}. The reliability of the $\gamma$-ray strength predictions can thus be greatly improved by the use of microscopic models. Indeed, provided satisfactory reproduction of available experimental data, the more microscopic the underlying theory, the greater the confidence in the extrapolations out towards the experimentally unreachable regions. Microscopic approaches are rarely used for practical applications. First, the time cost is often prohibitive for large scale calculations. Second, the fine tuning required to reproduce accurately a large experimental data set is very delicate, in addition to be time consuming. A prominent exception is represented by Refs. \cite{Gor02,Gor04} where a complete set of $\gamma$-ray strength functions was derived from mean field plus quasiparticle random phase approximation (QRPA) calculations. In Refs. \cite{Gor02,Gor04}, zero-range Skyrme forces were considered and phenomenological corrections applied to properly describe the splitting of the giant dipole resonance (GDR) in deformed nuclei as well as the damping of the collective motion. The present study aims to go beyond the former approximation providing axially-symmetric-deformed QRPA approach based on Hartree-Fock-Bogoliubov (HFB) calculations using the finite-range Gogny interaction in a fully consistent way. Contrary to the calculations in which radial wave functions are discretized on a mesh, single-particle wave functions are here expanded on an optimized harmonic oscillator (HO) basis. The present approach is specially suited for open-shell nuclei, where pairing correlations are included without any additional parameters. Another asset is an adequate treatment of the deformation at each step of the calculation. More precisely the intrinsic deformation of the nucleus ground state is predicted by the HFB calculations as the minimum of the potential energy surface. Then the QRPA phonons are oscillations around this minimum, spherical or not. This treatment significantly improves the description of the nuclear structure property, hence the $\gamma$-ray strength function predictions. It was first applied to study giant resonances in Si and Mg isotopes \cite{PG08}. Dipole excitations in Ne isotopes and $N=16$ isotones \cite{Mar11} as well as electromagnetic excitations of the heavy deformed $^{238}$U \cite{Per11} have been obtained with an optimized version of the numerical code, opening the way to large scale calculations. This powerful tool has been generalized to treat also charge exchange excitations which are relevant for the $\beta$-decay of experimentally inaccessible nuclei \cite{Martini:2014ura}. In the present paper, we test the predictive power of the aforementioned approach by applying it to a large set of even-even nuclei in order to compare with available experimental data on photoabsorption~\cite{ripl3}, the $E1$ $\gamma$-strength function of odd nuclei being obtained through an interpolation procedure involving neighbouring even-even nuclei. The paper is organized as follows. In Sec.~\ref{sec_mod}, the axially-symmetric-deformed HFB+QRPA formalism is described in its standard form and possible extensions are sketched. In the same spirit as Ref.~\cite{DechargeNPA83}, the impact of the size of the finite HO basis including cut-off effects are analyzed in Sec.~\ref{sec_sens} adding discussion on the choice of the interaction parameter sets. This convergence analysis sets a protocol for large scale calculations whose results are presented in Sect.~\ref{sec_exp}. First, the impact on deformation of the $\gamma$-strength function is illustrated. Second, the comparison with photo absorption data ~\cite{ripl3} is shown. Third, models are introduced to obtain the continuous strength functions starting from the discrete QRPA strength distributions $B(E1)$ using or not microscopic input. The final $E1$ strength functions as well as the corresponding Hauser-Feshbach astrophysical reaction rates are finally estimated for a large set of exotic neutron-rich nuclei and compared with other predictions. Conclusions are drawn in Sect.~\ref{sec_conc}.

\label{sec_conc} Large scale calculations of $E1$ $\gamma$-ray strength functions for even-even nuclei have been undertaken within the axially-symmetric-deformed HFB+QRPA approach in a fully consistent way using the finite-range Gogny interaction. The convergence of the numerical calculation has been analyzed with respect to the size of the basis (number of HO shells) and the 2-$qp$ excitation energy. This analysis allowed to establish practical choices for large scale calculations optimizing both the minimization of the computational cost and the convergence. Predictions obtained for two parameter sets of the Gogny interaction, namely D1S and D1M, have been compared and shown to give rise to a similar global behaviour. Nevertheless, the D1M $E1$ strength is found to be systematically shifted, leading to lower centroid energy and a smaller EWSR in comparison with D1S. The role of the intrinsic deformation has been systematically investigated. The split between $K$=0 and $|K|$=1 total angular momentum projections for deformed nuclei as well as their opposite hierarchy for prolate and oblate shapes has been confirmed in the whole nuclear chart. Large deformations give rise to a double peak structure while small deformations lead to a split too small to be disentangled. We have calculated the $E1$ $\gamma$-ray strength for all even-even nuclei for which photoabsorption data exist. % The comparison between QRPA and measured cross sections revealed a systematic energy shift of about $\sim$2 MeV of the Gogny-QRPA strength with respect to experimental data. This energy shift (of $\sim$2 MeV) induces that the theoretical values of the EWSR obtained for all the nuclei studied in the present work are systematically larger than the experimental ones. Three prescriptions have been proposed to cure this discrepancy. They correspond to a folding procedure of the $B(E1)$ discrete distribution on the basis of a Lorentzian function, which not only shifts the $E1$ strength down by more or less 2~MeV but also widen the distribution, as experimentally observed in photoabsorption cross sections. Such a folding is also needed to produce $S_{E1}$ strength functions as inputs in reaction models (\textit{e.g.} TALYS code). Two of these prescriptions fit the experimental photoabsorption data by adjusting the width and energy shift parameters of the Lorentzian, taking into account the density of dipole 4-$qp$ excitations. This microscopic ingredient allows us to investigate effects beyond the standard QRPA description involving only 2-$qp$ excitations. The three prescriptions are found to give a satisfactory description of experimental data and are shown to provide globally similar $E1$ strength. In particular the neutron capture rates of astrophysical interest do not differ by more than a factor of 2 for the three prescriptions, even far away from the valley of $\beta$-stability. Further improvements of the present approach can be envisioned. First, the interpolation procedure for odd nuclei can be replaced by a fully microscopic QRPA calculation of odd systems. Second, the folding procedure can be improved by including microscopically the particle-phonon coupling. Third, dynamic deformations can be considered when differing from the HFB one (see \textit{e.g.} Fig. 3 of Ref. \cite{Delaroche:2009fa}). The present encouraging results, in parallel to those obtained for nuclear masses \cite{Gor09} and nuclear level densities \cite{Hil12}, will allow us to include such microscopic ingredients (obtained on the basis of one unique Gogny interaction) into cross sections calculations. We believe that working along such a path is a way, in the future, to improve cross section evaluations and predictions on the basis of reliable and accurate microscopic inputs.

1607.08483

1607

1607.03933_arXiv.txt

We report on the analysis of {\it NuSTAR} observations of the Be-transient X-ray pulsar \src during the giant outburst in 2015 and another minor outburst in 2016. We confirm the cyclotron-line energy -- luminosity correlation previously reported in the source and the line energy decrease during the giant outburst. Based on 2016 observations we find that a year later the line energy has increased again essentially reaching the pre-outburst values. We discuss this behaviour and conclude that it is likely caused by a change of the emission region geometry rather than previously suggested accretion-induced decay of the neutron stars magnetic field. At lower luminosities we find for the first time a hint of departure from the anti-correlation of line energy with flux, which we interpret as a transition from super- to sub- critical accretion associated with disappearance of the accretion column. Finally, we confirm and briefly discuss the orbital modulation observed in the outburst light curve of the source.

In binary systems accretion of matter supplied by non-degenerate companion onto a strongly magnetised ($B\sim10^{12}$\,G) rotating neutron star results into pulsed X-ray emission from the vicinity of neutron stars (NSs) magnetic poles. The plasma is channeled to the polar caps by the magnetic field of the neutron star which has also profound effect on the observed X-ray spectra. In particular, the motion of the electrons in strong magnetic field is quantised, which gives rise to the so-called cyclotron resonance scattering features (CRSFs, see \cite{Mushtukov15x} for a recent review). A single (fundamental) or multiple harmonics \citep{Santangelo99} absorption-like features can be observed in X-ray band depending on the conditions in the line forming region. In particular, the energy of the fundamental is related to the magnetic field strength as $E_{\rm cycl}\sim 12\mbox{keV}\,B/10^{12}$\,G. The structure of the emission region in the vicinity of the NS and thus the origin of the CRSF are, however, uncertain. At low accretion rates most of the observed emission likely comes directly from the accretion mounds on the polar caps of the NS where the gravitational energy of the flow is released. However, the observed luminosities of bright pulsars by far exceed the local Eddington limit for kilometre-sized polar caps. At high accretion rates, the plasma must be, therefore, stopped above the NS surface by the radiative pressure and the observed emission has to emerge from the extended ``accretion column'' \citep{Basko,Becker12,2015MNRAS.454.2539M}. Conditions for the transition between the two regimes are determined by the largely unknown geometry of the column, and by the angular- and energy-dependent plasma opacities which makes it extremely hard to make robust theoretical predictions on the transition (critical) luminosity \citep{Mushtukov15,2015MNRAS.454.2539M}. On the other hand, analysis of the luminosity dependence of the observed properties of X-ray pulsars might help to constrain the critical luminosity observationally \citep{Tsygankov06,Staubert07,Klochkov12}. Indeed, in low luminous sources the CRSF energy typically increases with the flux, whereas at higher accretion rates an anti-correlation is observed. As discussed by \cite{Staubert07}, \cite{Becker12}, \cite{Mushtukov15}, and \cite{Mushtukov15pcor}, this behaviour could point on the two accretion regimes corresponding to sub- and super-critical accretion. Observing the transition between the two regimes in a single source would strongly support this interpretation. In this paper we report on the analysis of the CRSF luminosity dependence in the \emph{Be-}transient X-ray pulsar \src during the giant outburst in 2015 and another minor outburst in 2016, and discuss the complex evolution of the line energy throughout the giant outburst and between the two outbursts which, we argue, provides the first evidence for such transition.

Based on the analysis of {\it NuSTAR} observations of \src during the declining part of the 2015 giant outburst we have confirmed the previously known anti-correlation of CRSF energy with luminosity. We also confirm the apparent drop of the CRSF centroid energy during the declining part recently interpreted by \cite{Cusumano16} as the result of the accretion-induced decay of the magnetic field of the NS. We find that line energy decrease is consistent with being time linear throughout the outburst with rate of $\sim0.015$\,keV/d. Furthermore, follow-up \emph{NuSTAR} observations of another outburst in 2016 revealed that the line energy has increased again approximately to values observed before the 2015 outburst which implies a recovery rate of $\sim0.05$\,keV/d. Both timescales imply unprecedentedly fast evolution of the magnetic field of the neutron star if the change of the observed line energy is directly related to field strength as suggested by \cite{Cusumano16}. We argue, however, that evolution of the observed CRSF energy is likely instead associated with a change of the emission region geometry. The later is defined by the magnetosphere size which indeed seems to be different in rising and declining parts of the outburst as sugested by the observed spin evolution of the neutron star. Finally, we find that at luminosities below $\sim10^{37}$\,erg\,s$^{-1}$ the anti-correlation of the CRSF energy with flux reported for higher luminosities seems to break, which we interpret as the first observational evidence for the transition from super- to sub-critical accretion. The transitional luminosity is in agreement with the theoretical predictions and cyclotron line luminosity dependence observed in other sources as shown in Fig.~\ref{fig:lxcrit}. We note, however, that taking into the account complex evolution of line energy throughout the outburst and comparatively low statistical significance of the break, the transition can not be considered to be robustly detected and additional observations are required to confirm our findings.

1607.03933

1607

1607.05724_arXiv.txt

We present a method to model optical images of galaxies using Expectation Maximization Principal Components Analysis (EMPCA). The method relies on the data alone and does not assume any pre-established model or fitting formula. It preserves the statistical properties of the sample, minimizing possible biases. The precision of the reconstructions appears to be suited for photometric, morphological and weak lensing analysis, as well as the realization of mock astronomical images. Here, we put some emphasis on the latter because weak gravitational lensing is entering a new phase in which systematics are becoming the major source of uncertainty. Accurate simulations are necessary to perform a reliable calibration of the ellipticity measurements on which the final bias depends. As a test case, we process $7038$ galaxies observed with the ACS/WFC stacked images of the Hubble eXtreme Deep Field (XDF) and measure the accuracy of the reconstructions in terms of their moments of brightness, which turn out to be comparable to what can be achieved with well-established weak-lensing algorithms.

Optical imaging is without any doubt one of the main tools to investigate galaxies and dark matter through weak and strong gravitational lensing. Because of the large available data sets, it is crucial to extract all information available in noisy data and to simulate images precisely to calibrate the various methods and properly deal with possible biases. There is thus a pressing need to extract clean galaxy images from data. In particular, several studies have shown how all methods used to measure the ellipticity of galaxies require realistic simulations for their calibration \citep{viola11,bartelmann12,refrejer12,2012arXiv1204.5147M,massey13,gorvich16,bruderer16}. This issue is becoming pressing because of the stringent requirements posed by upcoming wide-field surveys such as the ESA Euclid space mission \citep{laureijs09} and the Large Synoptic Survey Telescope \citep{kaiser02}, among others. Galaxy models based on simple analytical recipes, for example based on the Sersic profile \cite{1968adga.book.....S}, have been widely used at this end \citep{heymans06,great08,kitching12}. These models have proven to well suffice for ground-based observations, but more accurate simulations are now necessary to include complex morphologies to account for spiral, irregular and cuspy galaxies. Also in the strong gravitational lensing regime, accurate galaxy models are now needed to investigate the use of substructures of strongly magnified galaxies to better constrain the mass distribution of lenses such as for example galaxy clusters \citep{meneghetti08,zitrin15}. For these reasons, galaxies observed with HST have been modelled with shapelets \citep{refregier03,refregier03b,massey03,massey07} to achieve noise-free images. Even if this approach deals with complex morphologies, artifacts may arise because of the oscillatory behavior of shapelets. Moreover, also smooth galaxies such as for example ellipticals are not very well reconstructed by this approach because of their slope, which is not well compatible with a Gaussian, which the Hermite polynomials in the shapelets are derived from. Also, image cut-outs have been extracted from HST data \citep{rowe14,mandelbaum15}, but these stamps are affected by the instrumental noise limiting their applicability. In this paper, we present a method to retrieve and reconstruct clean galaxy images in a model-independent way, which also preserves the statistical properties of the reconstructed sample. We do this using Expectation Maximization Principal Components Analysis \cite{bailey12}, which we use to derive a set of orthonormal basis functions optimized for the specific data set to be processed. Other studies used standard PCA \citep{Jolliffe86} to model elliptical galaxies \citep{joseph14J}. However, these are galaxies with smooth morphology, and this method cannot deal with weights and missing data. In contrast, the procedure discussed in this work allows us to process astronomical images with masked areas and pixel-dependent variance, and it allows us to introduce regularization terms to be used when deriving the principal components and impose smoothness on the basis. \begin{figure} \centering \includegraphics[width=1.0\hsize]{./fig/image-galaxy-split} \caption{This schematic representation shows how a postage-stamp of a galaxy image is rearranged in form of a vector, $\vec{d}_i$. All rows of pixels in the matrix composing the image are simply concatenated.} \label {fig:subaru-train-cov} \end{figure} \begin{figure*} \centering \includegraphics[width=0.33\hsize]{./fig/magDistribution_str-gal_clash} \hfill \includegraphics[width=0.33\hsize]{./fig/magError_str-gal_clash} \hfill \includegraphics[width=0.33\hsize]{./fig/mag-size_str-gal_clash} \caption {Statistical properties of the sources detected in the simulation and reproducing a Subaru image with 20 minutes of exposure time. A comparison with the real data is provided. From left to right, the panels show the magnitude distribution of the objects, their photometric error as a function of magnitude, and their size versus magnitude relation. } \label{fig:sim-stat} \end{figure*} As a test case, we analyzed $7038$ galaxies extracted from the \cite{rafelski15} catalog with redshift up to $z<4.0$ and maximum magnitude up to $m_{F775W}<30$. The catalog contains all photometric information including the photometric redshifts of the objects. We modeled these galaxies in all 5 optical bands extracted from the Hubble eXtreme Deep Field \citep[XDF hereafter,][]{illingworth13}. We tested the quality of the models by comparing their moment of brightness against those measured with the weak-lensing Shapelens library \citep{2011MNRAS.412.1552M}. Moreover we showed how to use these models to construct realistic simulations of astronomical images. The structure of this paper is a follows: in Sec.~(\ref{sec:empca}), we derive the EMPCA which are then used in Sec.~(\ref{sec:models}) to create the models of the galaxy images. The description of the analysis of the XDF data set in Sec.~(\ref{sec:models}), a simple sky simulation based on our models is presented in Sec.~(\ref{sec:simulations}), and the conclusions are given in Sec.~(\ref{sec:conclusions}).

\label{sec:conclusions} We have described how optical images of galaxies can be fitted with an optimized linear model based on the Expectation Maximization Principal Components Analysis (EMPCA). This method relies on the data alone, avoiding any assumptions regarding the morphology of the objects to be modeled even if they have complex or irregular shapes. As a test case, we have analyzed the galaxies listed in the \cite{rafelski15} catalog which covers the Hubble eXtreme Deep Field (XDF). We selected those objects with magnitude $m_{W755}<30$, far from the field edges and without overlapping artifacts caused by the few stars present across the field. We collected 7038 postage-stamps of noise-free galaxy images with redshift up to $z=4.0$. We have shown how the modeled galaxies well represent the entire collection of galaxies, from small to large and from regular to irregular. Two codes have been implemented to this end: {\it EasySky} to create the simulations and {\it EasyEMPCA} to model the galaxies. The residuals appear uncorrelated except at very sharp features because of the regularization scheme we adopted during the basis construction. To further verify the quality of the reconstructions, we simulated a set of galaxy images, with and without noise, covering the entire spectrum of shapes and luminosities of the objects present in the XDF. We processed the simulations with the same procedure applied to a real data set: we detected the objects with SExtractor, derived the EMPCA basis and fitted the data with the linear model based on this basis. We then measured the brightness moments up to the second order of the model reconstructions and compared them to those of the noise-free simulations. The quality of the reconstructions very well competes with a well-established method to measure galaxy brightness moments such as the iterative adaptive scheme implemented in Shapelens. The procedure discussed in this paper can be used to derive the properties of galaxies such as their fluxes and shapes, or to create reliable simulations of optical images. In this respect, the accuracy of such simulations is gaining importance for the lensing community. For instance, in the strong-lensing regime, they are necessary to understand how substructures in strongly magnified galaxies can be used to access additional information on the lensing mass distribution, such as galaxy clusters. In the weak-lensing regime, all methods to measure the ellipticites of galaxies require precise simulations for their calibration, on which depends the bias of such measurements and all quantities derived from them. The method we discussed in this work appears as a promising solution to create such simulations.

1607.05724

1607

1607.02274_arXiv.txt

In this paper, we present results of a photometric and kinematic study for a sample of 13 edge-on spiral galaxies with pronounced integral-shape warps of their stellar discs. The global structure of the galaxies is analyzed on the basis of the Sloan Digital Sky Survey (SDSS) imaging, in the $g$, $r$ and $i$ passbands. Spectroscopic observations are obtained with the 6-m Special Astrophysical Observatory telescope. In general, galaxies of the sample are typical bright spiral galaxies satisfying the Tully-Fisher relation. Most of the galaxies reside in dense spatial environments and, therefore, tidal encounters are the most probable mechanism for generating their stellar warps. We carried out a detailed analysis of the galaxies and their warps and obtained the following main results: (i) maximum angles of stellar warps in our sample are about 20$^{\rm o}$; (ii) warps start, on average, between 2 and 3 exponential scale lengths of a disc; (iii) stronger warps start closer to the center, weak warps start farther; (iv) warps are asymmetric, with the typical degree of asymmetry of about several degrees (warp angle); (v) massive dark halo is likely to preclude the formation of strong and asymmetric warps.

According to the prevailing point of view, disc galaxies are usually highly thin and flat. This is true, but only to a certain extent. Peripheral parts of galactic discs often exhibit deviations from the united plane and demonstrate global warps. This phenomenon has been revealed in the neutral gas component, through HI observations (e.g. \citealp{san1976}, \citealp{bosma1981}, \citealp{briggs1990}, \citealt{gsk2002}), and in a lesser extent through optical and infrared observations (e.g. \citealp{ss1990}, \citealp{grijs1997}, \citealp{rc1998, rc1999}, \citealp{schd2001}, \citealp{ss2003}, \citealp{ap2006}, \citealp{saha2009}, \citealp{gui2010}). Typically, warps appear in the outskirts of optical discs, and careful investigation of edge-on galaxies have revealed that a significant fraction (e.g. $\approx$50\% -- \citealp{ss1990}, $\approx$40\% -- \citealp{rc1998}, $\approx$50\% -- \citealp{ap2006}) of stellar discs shows integral-shape with typical amplitudes of a few degrees. This high percentage of observed warps would suggest that the majority of galaxies are warped, since the projection effects should hide large fraction of warps, whose line of nodes is perpendicular to the line of sight. Optical warps were also common in the past, with even greater amplitudes at $z \sim 1$ (\citealp{rbcj2002}). Several theoretical mechanisms have been proposed to explain the formation and maintenance of warped discs (e.g. \citealp{bin1992}, \citealp{kgr2001}, \citealp{sell2013} and references therein). Among the proposed scenarios, discrete modes of bending in a self-gravitating disc (\citealp{toomre1983}, \citealp{spcas1988}), misaligned dark halos (\citealp{dubk1995}), galaxy interactions and accretion of satellites (e.g. \citealp{hucar1997}, \citealp{schd2001}, \citealp{kim2014}), direct accretion of intergalactic matter in the outskirts of galaxies (e.g. \citealp{repf2001}, \citealp{vdk2007}, \citealp{ros2010}), extragalactic magnetic fields (\citealp{bat1990}), and others. This large variety of proposed mechanisms and their modifications probably indicates that there is no single mechanism responsible for all observable warps in galaxies. The current situation looks like the largest warps are mostly caused by tidal distortions (\citealp{schd2001}, \citealp{ap2006}), whereas relatively small warps are triggered and supported by a variety of mechanisms. Most previous studies of optical warps have been devoted to receiving a quantitative description of this phenomenon by collecting statistics for their properties and frequency depending on galaxy morphology, while thorough investigations of warped discs for individual galaxies were quite rare. (This is partly explained by the weakness of the phenomenon: warps are usually seen in the periphery of stellar discs, and their amplitudes are often small.) Some representative examples of galaxies with the detailed studies of their optical warps are M~33 \citep{san1980}, NGC~5907 (warp angle $\psi \approx 4^{\rm o}$ -- \citealp{sasaki1987}), Mrk~176 ($\psi \approx 19^{\rm o}$ -- \citealp{resh1989}), NGC~4013, NGC~4565, NGC~6504 (\citealp{flo1991}), MCG~06-30-005 ($\psi \approx 15^{\rm o}$ -- \citealp{kemp1993}), UGC~3697 ($\psi \approx 22^{\rm o}$ -- \citealp{ann2007}). Stellar discs of the Milky Way (\citealp{reed1996}), LMC (\citealp{os2002}), and the Andromeda galaxy (\citealp{inn1982}) were found to be also warped. The main goal of this paper is to perform a detailed photometric and kinematic study of thirteen edge-on spiral galaxies with integral-shaped stellar discs in order to derive observational characteristics of these galaxies and their warps. This information is very important for understanding the mechanisms responsible for the generation and maintenance of these prominent galaxy features. Until now, a detailed analysis of warped galaxies has been carried out for a few objects. In this paper, we are about to make a new contribution in this area. This paper is organized as follows. In the next section we present our sample. In Section~\ref{Preparation}, we introduce our decomposition technique, as well as the preparation and fitting of optical galaxy images. In Section~\ref{spectro}, we describe our own spectroscopic observations. The results of our investigation are presented in Section~\ref{Results}. We summarize our main findings and conclusions in Section~\ref{Conclusions}. Throughout this article, we adopt a standard flat $\Lambda$CDM cosmology with $\Omega_m$=0.3, $\Omega_{\Lambda}$=0.7, $H_0$=70 km\,s$^{-1}$\,Mpc$^{-1}$. \section[]{The sample} \label{s_samples} We selected our sample of galaxies based on the view of their non-flat stellar discs in the Sloan Digital Sky Survey (SDSS, \citealp{2015ApJS..219...12A}). Most of them were visually selected during the creation of the Sloan-based Polar Ring Catalogue \citep[SPRC][]{2011MNRAS.418..244M} and the catalogue of Edge-on disc Galaxies In SDSS \citep[EGIS][]{biz2014}. Our sample contains 13 objects and it is biased to galaxies with strong optical warps. Selected galaxies belong to different environments -- from group members to relatively isolated objects (this information is taken from the NASA/IPAC Extragalactic Database, NED\footnote{http://ned.ipac.caltech.edu/} hereafter). An overview of the basic properties of the sample galaxies is given in Table~\ref{Table1}. Fig.~\ref{SDSS_images} shows SDSS thumbnail images of the galaxies. Some brief notes on each object are given below. {\bf IC~194} (UGC~1542, RFGC~439, PGC~7812) is seen almost exactly edge-on. Optical disc warping is not strong: edge-on disc tips have a barely detectable $S$-shape. This galaxy has no nearby bright companions, but a loose group of galaxies PPS2~136 with the close redshift is located at a projected distance of 170~kpc from IC~194 (\citealp{tb1998}). {\bf 2MFGC~6306} is classified within the framework of the GalaxyZoo project \citep{2008MNRAS.389.1179L, 2011MNRAS.410..166L}, with the weighted fraction of votes out of all responses that this galaxy is edge-on $P=0.89$. Its stellar disc looks asymmetric and warped. 2MFGC-6306 probably forms a group with three other galaxies with close redshifts -- SDSS~J075644.97+440536.7, CGCG 207-007, and 2MASSX~J07563889+4407408. {\bf SPRC-192} is classified as a related to polar-ring galaxies (PRGs) object \citep{2011MNRAS.418..244M}. It is an early-type spiral with the visible dust lane along the disc major axis surrounded by an inclined ring-like structure. Morphologically, it is similar to the well-known polar-ring galaxy NGC~660. Obviously, if this galaxy had a higher inclination, it would look like a strongly warped spiral. The galaxy has no nearby companions of comparable sizes, but its morphology may point to a profound external perturbation or an accretion event in the past. {\bf UGC~4591} (RFGC~1430, PGC~24674) is an edge-on spiral galaxy, a member of the group of three galaxies WBL~193 \citep{1999AJ....118.2014W} (the brightest of them is the spiral galaxy IC~2394). The apparent disc bending starts at the half-radius of the galaxy. {\bf MCG~+06-22-041} (RFGC~1674, PGC~28776) at $z=0.027$ is a thin late-type spiral galaxy without any presence of a bulge. It was classified in the GalaxyZoo as an edge-on galaxy ($P=1.0$). Similar redshifts of nearby objects (elliptical galaxy CGCG~182-048 and 2MASSX~J09574267+3603307, both at $z=0.027$) suggest that MCG~+06-22-041 may be in a group. {\bf NGC~3160} (UGC~5513, PGC~29830) is one of the most interesting objects in our sample. It was classified in the GalaxyZoo as an edge-on galaxy with $P=0.838$. NGC~3160 shows several prominent features such as strong disc warping and a contrast X-shaped structure in the central part of the galaxy. NGC~3160 is a member of the cluster of galaxies NRGb~78, with the brightest elliptical galaxy NGC~3158 at the centre. {\bf UGC~5791} (SPRC-197, PGC~31697) is a blue peculiar galaxy for which S$^4$G (\citealp{s4g}) analysis has been done. It is the nearest galaxy in the sample. The inclination angle for this galaxy is barely measurable because of its peculiar shape. It is a pair member, together with the galaxy UGC~5798 \citep{1979ApJS...40..527P}. UGC~5791 is classified as a PRG-related object \citep{2011MNRAS.418..244M}. {\bf NGC~3753} (UGC~6602, Arp~320, SPRC-203, PGC~36016) is an interacting galaxy, marked by \cite{2011MNRAS.418..244M} as an object related to PRGs. It belongs to the compact group Hickson~57 \citep{Hick1982} and was included in the catalogue of nearby poor clusters of galaxies \citep{1999AJ....118.2014W}. The warped dust disc has apparent edge-on orientation. The stellar disc is significantly warped and has an asymmetric view. The galaxy image gives us a hint that its structure may include a bar or an inner disc overlapping with a dust component. Two bright companions near NGC~3753 are the spiral galaxy NGC~3754 and the elliptical galaxy NGC~3750. {\bf UGC~6882} (SPRC-204, PGC~37372) is another galaxy related to PRGs \citep{2011MNRAS.418..244M}, which is probably close to edge-on orientation according to the classification from the GalaxyZoo ($P=0.575$). It also belongs to the 2MASS selected Flat Galaxy Catalog \citep[2MFGC,][]{2004BSAO...57....5M}. Some small nearby galaxies are seen in the SDSS image, although they do not have optical spectra, and, thus, estimated redshifts. {\bf SDSS~J140639.64+272242.4} and {\bf SDSS~J153538.63+464229.5} are the two most distant and small in angular size galaxies in our sample. Both galaxies have several nearby companions in projection, but, unfortunately, they do not have measured redshifts. {\bf UGC~10716} (RFGC~3242, PGC~59657) is a late-type spiral galaxy viewed almost edge-on ($P=0.933$). It was selected in the RFGC catalogue and in the catalogue of edge-on galaxies created by \cite{2006A&A...445..765K}. The galaxy forms a triplet with UGC~10714 and SDSS~J170726.32+301316.6 (\citealp{ber2006}). {\bf UGC~12253} (RFGC~4028, PGC~70040) is seen in almost exactly edge-on orientation. The sharp dust lane divides the main body of UGC~12253 almost ideally in two halves. The X-pattern at the centre suggests that this galaxy has a boxy/peanut-shaped bulge, a possible evidence of the presence of a bar \citep{2006MNRAS.370..753B}. There are two nearby galaxies with close redshifts (PGC~070044 and SDSS~J225556.61+124701.9) and, therefore, UGC~12253 can be a member of a triplet or a group of galaxies. \begin{table*} \centering \begin{minipage}{150mm} \centering \parbox[t]{150mm} {\caption{Basic properties of the sample galaxies.} \label{Table1}} \begin{tabular}{cccccccccccc} \hline \hline \# & Galaxy & RA & Dec & $D$ (Mpc) & $M$ (mag) & $a$ (\arcsec) & $q$ & $g-r$ & $r-i$ & $T$ & $i$ (deg) \\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9) & (10) & (11) & (12) \\ \hline 1&IC 194&02:03:05&+02:36:51&83.8&-20.88&61.9&0.25&0.75&0.51&Sb & 87.3$\pm$1.2 \tabularnewline 2&2MFGC 6306&07:56:43&+44:05:49&192.0&-21.12&36.4&0.22&0.85&0.52&Sb & 88.0$\pm$1.6 \tabularnewline 3&SPRC 192&08:23:01&+32:00:54&266.5&-21.93&25.7&0.37&0.75&0.46&Sab & 80.0$\pm$5.1 \tabularnewline 4&UGC 4591&08:46:58&+28:14:17&91.5&-20.26&47.4&0.34&0.64&0.34&Scd & 89.7$\pm$0.4 \tabularnewline 5&MCG +06-22-041&09:57:43&+36:04:09&113.3&-18.85&29.9&0.22&0.05&-0.1&Sd & 89.8$\pm$0.6 \tabularnewline 6&NGC 3160&10:13:55&+38:50:34&98.2&-21.40&56.2&0.36&0.85&0.44&Sb & 89.2$\pm$1.7 \tabularnewline 7&UGC 5791&10:39:27&+47:56:50&14.6&-16.68&56.3&0.43&0.31&0.16&Sc & 85.0$\pm$5.0 \tabularnewline 8&NGC 3753&11:37:54&+21:58:53&124.0&-22.29&74.6&0.31&0.89&0.49&Sab& 84.3$\pm$2.7 \tabularnewline 9&UGC 6882&11:54:43&+33:32:12&135.1&-20.98&43.2&0.23&0.74&0.38&Sbc& 80.1$\pm$6.2 \tabularnewline 10&SDSS J140639.64+272242.4&14:06:40&+27:22:42&303.5&-20.89&24.3&0.18&0.64&0.49&---& 88.6$\pm$1.3\tabularnewline 11&SDSS J153538.63+464229.5&15:35:39&+46:42:30&271.4&-20.17&20.7&0.22&0.41&0.18&---& 84.4$\pm$1.9\tabularnewline 12&UGC 10716&17:07:44&+30:19:35&132.5&-20.61&39.9&0.23&0.66&0.41&Sb & 87.9$\pm$1.3 \tabularnewline 13&UGC 12253 &22:56:02&+12:45:59&101.4&-20.61&49.8&0.27&0.84&0.42&Sb& 89.8$\pm$0.5 \tabularnewline \hline\\ \end{tabular} \end{minipage} \parbox[t]{150mm}{ Columns: \\ (1) Designation number in the sample, \\ (2) first name from NED, \\ (3), (4) J2000 coordinates from NED,\\ (5) 3K CMB distance from NED (except for 2MFGC 6306 for which the distance was calculated from the radial velocity taken from the HyperLeda database), \\ (6) absolute magnitude in the $r$ band inside of the isophote of 25.5 mag/arcsec$^2$. Galactic extinction (according to \citealp{2011ApJ...737..103S}) and K-correction (using the NED K-Correction calculator based on \citealp{2012MNRAS.419.1727C}) was taken into account, \\ (7), (8) semi-major axis and galaxy flattening of the 25.5 mag/arcsec$^2$ isophote in the $r$-band, \\ (9), (10) colours corrected for Galactic extinction and K-correction, \\ (11) morphological type from NED. For SDSS~J140639.64+272242.4 and SDSS~J153538.63+464229.5, there is no morphological classification in NED or HyperLeda databases,\\ (12) galaxy disc inclination (see Appendix A in \citealp{2015MNRAS.451.2376M}). For UGC~4591 and MCG~+06-22-041, we applied the first approach from \cite{2015MNRAS.451.2376M} using the 3D disc decomposition available in the \textsc{IMFIT} code. Inclination estimate for UGC~5791 is found on the basis of the apparent axial ratio of the central part of the galaxy. For the remaining galaxies, we estimated the inclination by the orientation of the dust lane (the second method from \citealp{2015MNRAS.451.2376M}). } \end{table*} \begin{figure*} \includegraphics[width=3.5cm, angle=0, clip=]{1_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{2_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{3_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{4_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{5_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{6_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{9_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{10_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{11_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{10_new.eps} \includegraphics[width=3.5cm, angle=0, clip=]{11_new.eps} \includegraphics[width=3.5cm, angle=0, clip=]{7_name.eps} \includegraphics[width=3.5cm, angle=0, clip=]{8_name.eps} \caption{SDSS thumbnail images for the 13 galaxies in our sample with the field of view of 90$\arcsec$ by 90$\arcsec$.} \label{SDSS_images} \end{figure*}

\label{Conclusions} We have performed detailed photometric and kinematic study of the 13 edge-on spiral galaxies with warped stellar discs. The galaxies were selected solely on the basis of their optical morphology and, therefore, they are {\it a priori} with notable integral-shape warp. As it turned out, most of them reside in dense spatial environments -- in pairs, groups, clusters -- and, hence, tidal interaction (current or past) with companions may be possible mechanism for the origin of stellar warps in our sample. Our main conclusions are as follows: (i) The sample galaxies demonstrate wide distribution of the warp angles $\psi$, with maximum values of $\psi \sim 20^{\rm o}$ (Fig.~\ref{hist}a). (ii) On average, stellar warps start at a projected distances of (2--3)\,$h$, i.e. near or just beyond the maximum of the rotation curve of a self-gravitating exponential disc (Table~\ref{Table6}). (iii) Stronger warps have on average a smaller projected starting point (Sect.~\ref{warp_char}). (iv) Warps show notable asymmetry in all the sample galaxies. Typically, asymmetry reaches several degrees in $\psi$, about 50 per cent of the exponential scale length in $R_w$, and about 0.5 mag in $\mu_w$ (Tables~\ref{Table3} and \ref{Table7}). (v) Apparently, as the dark halo becomes more and more massive compared to the stellar disc, it prevents the formation of very strong and asymmetric warps (Fig.~\ref{am}). In order to investigate the formation of stellar warps, \citet{kim2014} recently presented a set of $N$-body simulations of fly-by encounters between galaxies. They have found that fly-bys can excite integral-shape warps in galaxies, and such induced warps can survive for a few billion years. In \citet{kim2014} simulations, the maximum warp angle reaches about 25$^{\rm o}$, and warps are often non-symmetric. These results are quite consistent with our observational data. Significant asymmetry of the projected starting points of warps (Sect.~\ref{warp_char}) could probably arise due to the underlying asymmetry in the dark halo potential (\citealp{saha2009}), which could be induced by the galaxy-galaxy interactions. Our results are based on a small sample of galaxies, and they can only be regarded as preliminary. Also, these results can be biased due to the misidentification of strongly inclined spirals as warped stellar discs. \cite{rc1998} simulated this projection effect and found that no more than $\approx$15\% of integral-shaped warps could actually be spiral arms. Careful study of the sample galaxies showed that, with a certain probability, in 2 of 13 galaxies the global warp can be explained as inclined spiral arms (these galaxies are UGC~6882 and SDSS~J153538.63+464229.5). Thus, the fraction of possible false warps in the sample is in general agreement with the \cite{rc1998} results. We verified the position of the studied characteristics of these two galaxies in our figures and found that they do not bias our conclusions. A further extended studies of warped galaxies in different spatial environments will help us to better understand this common but still puzzling phenomenon.

1607.02274

1607

1607.07451_arXiv.txt

We describe the design and performance of the hardware system at the Bleien Observatory. The system is designed to deliver a map of the Galaxy for studying the foreground contamination of low-redshift (z=0.13--0.43) H$_{\rm I}$ intensity mapping experiments as well as other astronomical Galactic studies. This hardware system is composed of a 7m parabolic dish, a dual-polarization corrugated horn feed, a pseudo correlation receiver, a Fast Fourier Transform spectrometer, and an integrated control system that controls and monitors the progress of the data collection. The main innovative designs in the hardware are (1) the pseudo correlation receiver and the cold reference source within (2) the high dynamic range, high frequency resolution spectrometer and (3) the phase-switch implementation of the system. This is the first time these technologies are used together for a L-band radio telescope to achieve an electronically stable system, which is an essential first step for wide-field cosmological measurements. This work demonstrates the prospects and challenges for future H$_{\rm I}$ intensity mapping experiments. \\

In the coming decades, a wealth of astronomical data in radio wavelength will become available through large survey projects and telescopes such as LOFAR\footnote{\url{ http://www.lofar.org}} \citep{vanHaarlem2013}, GMRT\footnote{\url{http://www.ncra.tifr.res.in/ncra/gmrt}} \citep{Paciga2013}, PAPER\footnote{\url{ http://eor.berkeley.edu}} \citep{Ali2015}, CHIME \citep{Bandura2014}, BINGO \citep{Battye2012, Battye2013}, HERA \citep{Pober2014}, Tianlai \citep{Chen2012}, and SKA\footnote{\url{ http://www.skatelescope.org}} \citep{Mellema2015}. Most of these projects aim at measuring signal from the redshifted 21 cm H$_{\rm I}$ emission line, from either low-redshift large-scale structure \citep[e.g.,][]{Battye2012, Masui2013} or high-redshift Epoch of Reionization \citep[e.g.,][]{Furlanetto2006, Paciga2013}. Many of these surveys will suffer from foreground contamination of our own galaxy, the Milky Way, which is typically several orders-of-magnitude larger than the cosmological signals of interest. As a result, a good understanding for this foreground component is crucial for the interpretation of the cosmological measurements of interest. This need for a deeper understanding of the Milky Way at the relevant H$_{\rm I}$ wavelength is the main driver for this work. To date, only a few of wide-field Galactic maps exist and are available for the community to use to study the H$_{\rm I}$ foreground \citep[see][and reference therein]{deOliveira-Costa2008}. Furthermore, the frequency coverage has been fairly sparse across the different maps. The most commonly used map around 1 GHz has been the so-called ``Haslam map'' at 408 MHz \citep{Haslam1982}. This map has been re-processed several times \citep{Remazeilles2015}, but almost no new data has been taken systematically since then. The main reason for the lack of wide-field maps at these frequencies is that historically, this has not been a wavelength window for cosmological measurements. In addition, the generation of these maps require significant observing time even at moderate resolution. With the next generation of H$_{\rm I}$ cosmological experiments being built, foreground maps at the radio L-band are becoming more important. \begin{figure} \begin{center} \includegraphics[scale=0.32]{figs/overview_schematics.pdf} \caption{Schematic illustration of the entire system for this project. The three grey boxes indicate the three different locations where the instrumentation are -- the front-end unit mounted inside the telescope dish, the telescope tower basement, and the control room located about 60 m away from the telescope. The white squares indicate controller PCs whereas the white rounded squares indicate instruments. All computers in the control room connect to each other as well as the external network. } \label{fig:overview} \end{center} \end{figure} In this work, we investigate the possibility of a simple system that could be used for these foreground maps. This system consists of a dedicated single-dish telescope that scans the sky over several months with very low human intervention during the observation. The hardware system is set up at the 7m telescope at the Bleien Observatory, operating in the frequency range 990--1260 MHz. We emphasise the innovative yet low-cost design that resulted in significant improvement in data quality compared to the more conventional systems. This work continues from \citet{Chang2015} as a series of studies at the Bleien Observatory in preparation for future H$_{\rm I}$ intensity mapping experiments. The paper is organized as follows. In \Sref{sec:system}, we describe our overall system at the Bleien Observatory, both the hardware and the controlling/monitoring mechanism. In \Sref{sec:commission}, we evaluate the performance of the system and the improvement over the previous system with both laboratory measurements and on-sky measurements. Our conclusions are summarized in \Sref{sec:conclusion} and additional tests of the hardware system are presented in the Appendix. A companion paper (Akeret et al. in prep, hereafter A16) describes the analysis software and some early processed data.

\label{sec:conclusion} In this paper we describe the design of an integrated system at the Bleien Observatory for mapping the Galaxy. The system is designed to map the Milky Way in the L-band frequency range 990--1260 MHz with a 7m single-dish telescope. The ultimate science goal of this system is to provide a set of new data for studying the foreground effect of future low-redshift H$_{\rm I}$ intensity mapping cosmology experiments such as BINGO and HIRAX. Specifically, the data from this system will fill the gap between the Haslam map at 408 MHz and the various maps constructed at the 21 cm frequency (1420 MHz). This data will also allow for explorations in different areas of Galactic astronomy. We describe the hardware upgrade of the 7m single-dish telescope at the Bleien Observatory from a more conventional electronic chain to a high-performance, stable system. We also demonstrate that the improvement in the data quality is significant. Several innovative designs in the hardware system result in a very stable system (Allan-time $>$ 5000 s) with high dynamic range (70 dB) and high frequency resolution (50 kHz), which is important for the large-scale Galactic map. These innovations include the pseudo correlation receiver and the cold reference source within the receiver, the FFT-spectrometer and the phase-switch operation of the whole system. Some of these components have been well tested and developed in the microwave instruments for cosmic microwave background (CMB) measurements, but for radio astronomy, these technologies driven by large-scale cosmology are relatively new. The hardware system presented in this work provides an environment to test many of these new ideas for a radio telescope. The hardware system described in this paper together with the software pipeline developed in the companion paper A16 can be fed into the planning of a generic Galactic mapping survey. Future potential upgrades of this system include building a parallel electronic chain for the second polarization, installing ground-shields to reduce ground-pickup and RFI, and making the FFT-spectrometer more programmable so that more powerful functionalities in the spectrometer can be used. One interesting extension is to use the internal correlation function in the FFT-spectrometer to replace the front end correlation receiver \citep{Kooi2004}. This potentially could save hardware cost significantly.

1607.07451

1607

1607.03980_arXiv.txt

HR 8799 harbors arguably the first and best-studied directly-imaged planets. In this brief article, I describe how the HR 8799 planetary system is a benchmark system for studying the atmospheres, orbital properties, dynamical stability, and formation of young superjovian planets. Multi-wavelength photometry and spectroscopy show that HR 8799 bcde appear to have thicker clouds than do field brown dwarfs of similar effective temperatures and exhibit evidence for non-equilibrium carbon chemistry, features that are likely connected to the planets' low surface gravities. Over 17 years of astrometric data constrain the planets' orbits to not be face on but possibly in multiple orbital resonances. At orbital separations of 15--70 au and with masses of $\approx$ 5--7 $M_{\rm J}$, HR 8799 bcde probe the extremes of jovian planet formation by core accretion: medium-resolution spectroscopy may provide clues about these planets' formation conditions. Data from the next generation of 30 m-class telescopes should better constrain the planets' orbits, chemistry, gravity, and formation history.

In March 2008, Christian Marois noticed one (Figure 1, left panel), and then two, faint point sources located at a projected separation of 68 and 38 au from the nearby, dusty A5 star HR 8799. Follow-up observations in July--September 2008 confirmed that these objects were bound companions and added a third at $\rho$ $\sim$ 24 au \citep[][hereafter Ma08]{Marois2008}. Two years later, \citet{Marois2010} announced the discovery of a fourth companion at $\rho$ $\sim$ 15 au: HR 8799 e (Figure 1, right panel). Given the estimated age of the system ($t$ $\sim$ 30 $Myr$, Ma08, Baines et al. 2012), HR 8799 bcde's low luminosities imply masses of $\approx$ 5--7 $M_{J}$, below the deuterium-burning limit ($\approx$ 13 $M_{J}$) nominally separating planets from brown dwarfs. Analyses focused on atmospheric modeling \citep[][hereafter Cu11]{Currie2011a}, dynamical stability (\citealt{Marois2010}; Cu11), and formation (\citealt{Kratter2010}; Cu11) likewise corroborate the conclusion that HR 8799 bcde are bona fide planets, not brown dwarfs. The HR 8799 planetary system resembles a scaled-up version of our outer solar system. The planets orbit in between a warm dust belt ($r_{inner}$ $\approx$ 6-12 au) and a cold Kuiper belt-like structure at $r_{outer}$ $\approx$ 90--145 au \citep{Su2009,Booth2016}. Due to HR 8799's higher luminosity, the planets and dust belt populations receive about as much energy as the solar system's gas/ice giant planets and asteroid belt/Kuiper belt receive from the Sun. \begin{figure*}[htb] \begin{center} \includegraphics[width=0.95\textwidth, trim=0mm 0mm 0mm 0mm,clip]{hr8799_gemkeck.eps} \caption{ (left) The first direct image of an extrasolar planet: detection of HR 8799 b from October 2007 Gemini/NIRI data reduced in March 2008, the first of three planets (HR 8799 bcd) announced by \citet{Marois2008}. (right) Image of HR 8799 bcde from November 2009 Keck/NIRC2 data \citep{Marois2010} depicting the planets' counterclockwise orbital motion. The planets' discoveries were enabled by advances in observing and image processing techniques \citep[e.g.][]{Marois2006,Lafreniere2007}.} \label{hr8799discovery} \end{center} \end{figure*} HR 8799 harbors arguably not just the \textit{first} directly-imaged planets\footnote{While Fomalhaut b was announced on the same day as HR 8799 bcd and claimed to produce variable, accretion-driven emission at 0.6 $\mu m$ and thermal emission at longer wavelengths \citep{Kalas2008}, later work cast doubt on its existence (Janson et al. 2012) and then showed that instead Fomalhaut b is made visible entirely by circumplanetary dust emission \citep{Currie2012b,Galicher2013}. Thus, it is on slightly shakier ground and instead, as noted in \citet{Currie2012b}, is likely a ``planet [of unknown mass] \textit{identified by direct imaging} but not a \textit{directly-imaged} planet." While other planet-mass objects were announced prior to HR 8799 bcd \citep[e.g. 2M 1207 B;][]{Chauvin2004}, their lower mass ratios (compared to the primary) and/or wider separations suggest that they represent the low-mass tail of the substellar mass function. } but among the best studied ones. Photometry and/or low-resolution spectroscopy for HR 8799 bcde span 1--5 $\mu m$ (e.g. Cu11; \citealt{Barman2011a, Galicher2011,Zurlo2016}). HR 8799 bc have 1.4--2.5 $\mu m$ medium-resolution spectroscopy \citep{Konopacky2013,Barman2015}. After the reported discovery of HR 8799 bcd, multiple studies revealed at least one of the HR 8799 planets from data between 1998 and 2007 \citep[e.g.][]{Lafreniere2009,Metchev2009,Soummer2011,Currie2012a}. The planets have been imaged by nearly all other 5--8 m telescopes with adaptive optics systems (e.g. Cu11; \citealt{Currie2014a,Ingraham2014,Oppenheimer2013,Zurlo2016}). This wealth of data makes HR 8799 a benchmark system for studying the atmospheres, orbital properties, dynamical stability, and formation of young superjovian planets. %

1607.03980

1607

1607.06088_arXiv.txt

% We present a panoramic map of the entire Milky Way halo north of $\delta\sim-30\degr$ ($\sim$ 30,000 deg$^2$), constructed by applying the matched-filter technique to the Pan-STARRS1 $3\upi$ Survey dataset. Using single-epoch photometry reaching to $g\sim$22, we are sensitive to stellar substructures with heliocentric distances between 3.5 and $\sim$35~kpc. We recover almost all previously-reported streams in this volume and demonstrate that several of these are significantly more extended than earlier datasets have indicated. In addition, we also report five new candidate stellar streams. One of these features appears significantly broader and more luminous than the others and is likely the remnant of a dwarf galaxy. The other four streams are consistent with a globular cluster origin, and three of these are rather short in projection ($\la10\degr$), suggesting that streams like Ophiuchus may not be that rare. Finally, a significant number of more marginal substructures are also revealed by our analysis; many of these features can also be discerned in matched-filter maps produced by other authors from SDSS data, and hence they are very likely to be genuine. However, the extant $3\upi$ data is currently too shallow to determine their properties or produce convincing CMDs. The global view of the Milky Way provided by Pan-STARRS1 provides further evidence for the important role of both globular cluster disruption and dwarf galaxy accretion in building the Milky Way's stellar halo.

One consequence of the hierarchical galaxy formation process predicted by cold dark matter cosmological models is that a significant fraction of the stellar mass in galaxies has been accreted. In disc galaxies like the Milky Way, stars that formed {\it ex situ} are overall a minority, but dominate the stellar halo \citep[e.g.][]{pil15}. In these outer regions, where dynamical times are extremely long, the accreted material remains coherent for many billions of years \citep[e.g.][]{joh96}. Stellar streams are therefore powerful probes of the formation and evolution of galaxies: in addition to providing direct evidence of past and ongoing accretion and disruption events, the observed properties of these substructures contain a wealth of information on both their progenitors and their host galaxy. For example, the stars from disrupted galaxies and globular clusters approximately follow, and therefore trace, the orbit of their progenitor, which provides an estimate of the mass and morphology of the potential enclosed within the orbit \citep[e.g.][]{kop10}. The apparent width and velocity dispersion of globular cluster streams are strongly affected by density variations along their orbits, and can thus reveal the amount of clumpiness of the dark matter halo \citep[e.g.][]{iba02,nga16}. Finally, \citet{err15} have recently shown that the dark matter profile of dwarf spheroidal galaxies plays an important role in defining the sizes and internal dynamics of their tidal streams. With the advent of wide-field photometric observations and surveys, many streams and substructures have been detected in the Milky Way \citep[see][and references therein; hereafter GC16]{gri16b} and in nearby galaxies \citep[e.g.][]{mal97,sha98,iba01,mar10,iba14,oka15,duc15,crn16}. \defcitealias{gri16b}{GC16} In the Galaxy, most of the known substructures have been discovered by searching for coherent stellar over-densities in the homogeneous, wide-field photometric catalogue provided by the {\it Sloan Digital Sky Survey} \citep[SDSS;][]{yor00}, although several streams have recently been found in other wide-field surveys \citep[e.g.][]{ber14a,kop14,mar14,bal16}. While some streams have clearly originated from the accretion of dwarf galaxies, about three quarters are consistent with globular cluster disruption according to \citetalias{gri16b}. Since several teams have dedicated significant, independent efforts with the goal of detecting new substructures, one could expect that any stream within the detection limit of SDSS would have been found by now. However, like any survey, the SDSS catalogue suffers from artefacts, areas with shallower photometry due to e.g.\ weather conditions, and calibration issues revealing the observation patterns \citep[see e.g.][]{fin16}. Here we present a systematic search for stellar substructures in the whole sky north of $\delta>-30\degr$ by taking advantage of the extensive coverage of the Pan-STARRS1 (PS1) 3$\upi$ Survey. It significantly expands on the previous Milky Way substructure work that was carried out with an earlier data processing version of PS1 \citep{sla13,sla14,ber14b,mor16}. The current processing version reaches to roughly the same depth as the SDSS but covers 30,000 deg$^2$ with homogeneous and well-calibrated photometry. The observational strategy and data reduction procedure are completely different from those of SDSS, thereby allowing a fully independent analysis. We first provide a summary of the substructures recovered in our analysis, including further extensions of known features, then present five new candidate streams, all but one of which lie within the SDSS footprint.

1607.06088

1607

1607.06892.txt

{Direct imaging has led to the discovery of several giant planet and brown dwarf companions. These imaged companions populate a mass, separation and age domain (mass$>1~M_\mathrm{Jup}$, orbits>5~AU, age<1~Gyr) quite distinct from the one occupied by exoplanets discovered by the radial velocity or transit methods. This distinction could pinpoint that different formation mechanisms are at play.} {We aim at investigating correlations between the host star's mass and the presence of wide-orbit giant planets, and at providing new observational constraints on planetary formation models.} {We observed 58 young and nearby M-type dwarfs in $L^\prime$-band with the VLT/NaCo instrument and used ADI algorithms to optimize the sensitivity to planetary-mass companions and to derive the best detection limits. We estimate the probability of detecting a planet as a function of its mass and physical separation around each target. We conduct a Bayesian analysis to determine the frequency of substellar companions orbiting low-mass stars, using a homogenous sub-sample of 54 stars.} {We derive a frequency of $4.4^{+3.2}_{-1.3}\%$ for companions with masses in the range of 2-80~$M_\mathrm{Jup}$, and $2.3^{+2.9}_{-0.7}$\% for planetary mass companions (2-14~$M_\mathrm{Jup}$), at physical separations of 8 to 400~AU for both cases. Comparing our results with a previous survey targeting more massive stars, we find evidence that substellar companions more massive than 1~$M_\mathrm{Jup}$ with a low mass ratio Q with respect to their host star (Q<1\%), are less frequent around low-mass stars. This may represent an observational evidence that the frequency of imaged wide-orbit substellar companions is correlated with stellar mass, corroborating theoretical expectations. On the opposite, we show statistical evidence that intermediate-mass ratio (1\%<Q<5\%) companion with masses >2~$M_\mathrm{Jup}$ might be independent from the mass of the host star. } {} \date{}

After more than two decades of exoplanet discoveries \citep[e.g.][]{Wolszczan.1992, Mayor.1995, Udalski.2002}, more than two thousand planets have been identified, mostly using radial velocity (hereafter RV) and transit methods. Both techniques have detected exoplanets with masses ranging from one Earth mass up to the brown dwarf mass regime ($\sim$13-80~$M_\mathrm{Jup}$), orbiting at close separations from their host star (usually within 5~AU). Direct imaging (DI) has provided data for about 60 large separation ($\ge$10~AU) massive planets and brown dwarfs for more than a decade \footnote{http://exoplanet.eu/} (e.g. 51~Eri~b, \citealt{Macintosh.2015}; HD95086b, \citealt{Rameau.2013.HD95}; GJ504b, \citealt{Kuzuhara.2013}; 2MASS0103(AB)b, \citealt{Delorme.2013}; $\beta$~Pictoris b, \citealt{Lagrange.2010}; HR8799bcde, \citealt{Marois.2008}, \citealt{Marois.2010}; 2M1207b, \citealt{Chauvin.2004}). About a third of the imaged substellar companions (SCs, $<80M_\mathrm{Jup}$) have young M-type dwarf host stars (Table~\ref{Mcompanions}). \\ DI is usually limited to the detection of distant young massive planets because of their more favorable planet-to-host-star contrast and their larger angular separation. Some of the limitations of DI can be mitigated by using thermal imaging (in $L^\prime$-band) where young massive planets emit most of their flux, and by taking advantage of the intrinsically lower flux of low-mass stars, which provides a more favorable star-planet contrast. \begin{table*} \small \begin{center} \caption{Giant planets and brown dwarfs imaged around self-luminous young M-dwarfs (source: http://exoplanet.eu). Ref: $^{(1)}$\citet{Delorme.2013}, $^{(2)}$\citet{Naud.2014}, $^{(3)}$\citet{Artigau.2015}, $^{(4)}$\citet{Kraus.2014}, $^{(5)}$\citet{Itoh.2005}, $^{(6)}$\citet{Todorov.2010},$^{(7)}$\citet{Luhman.2006}, $^{(8)}$\citet{Chauvin.2004}, $^{(9)}$\citet{Gauza.2015}, $^{(10)}$\citet{Burgasser.2010}, $^{(11)}$\citet{Currie.2014}, $^{(12)}$\citet{Bowler.2014}, $^{(13)}$\citet{Luhman.2009}, $^{(14)}$\citet{Wahhaj.2011}, $^{(15)}$\citet{Bejar.2008}, $^{(16)}$\citet{Ireland.2011}, $^{(17)}$\citet{Allers.2005}} \begin{tabular}{c c c c c c} \hline\hline Name$^{ref}$ & SpT & d & Age & Comp. mass & Sep. \\ & & (pc) & (Myr) & (M$_\mathrm{Jup}$) & (AU) \\ [0.5ex] \hline Planetary-mass companions & & & & & \\ [0.5ex] \hline 2M~0103(AB)~b$^{(1)}$ & M5,5 + M6 & $47.2 \pm 3.1$ & $30$ & $13.0 \pm 1.0$ & $84.0$\\ GU~Psc~b$^{(2)}$ & M3 & $48.0 \pm 5.0$ & $100 \pm 30$ & $11.0 \pm 2.0$ & $2000 \pm 200$\\ J0219-3925~B$^{(3)}$ & M6 & $39.4 \pm 2.6$ & $30-40$ & $12-15$ & $156 \pm 10$ \\ FW~Tau~b$^{(4)}$ & M4 & $145.0 \pm 15.0$ & $1.8^{+0.2}_{-1.0}$ & $10.0 \pm 4.0$ & $330 \pm 30$\\ DH~Tau~b$^{(5)}$ & M0,5 & $\sim$ 140 & $1.0 \pm 0.9$ & $11.0_{-0.003}^{+0.01}$ & $330$\\ 2M~044144~b$^{(6)}$ & M8,5 & $140.0$ & $1.0$ & $7.5 \pm 2.5$ & $15.0 \pm 0.6$\\ CHXR~73~b$^{(7)}$ & M3,25 & $\sim$ 1.6 & $2$ & $12.0_{-5.0}^{+8.0}$ & $200$\\ 2M1207~b$^{(8)}$ & M8 & $52.4 \pm 1.1$ & $8 \pm 3$ & $4.0_{-1.0}^{+6.0}$ & $46.0 \pm 5.0$\\ VHS~1256-1257~B$^{(9)}$ & M7.5 & $12.7 \pm 1.0$ & $150-300$ & $11.2^{+9.7}_{-15}$ & $102 \pm 9$\\ Ross~458(AB)~c$^{(10)}$ & M2 & $114.0 \pm 2.0$ & $475 \pm 325$ & $8.5 \pm 2.5$ & $1168.0$\\ ROXs~42B~b$^{(11)}$ & M0 & $135.0$ & $2 \pm 1$ & $10.0 \pm 4.0$ & $140.0 \pm 10.0$\\ \hline Brown dwarfs & & & & & \\ [0.5ex] \hline 1RXS~J034231.8+121622~B$^{(12)}$ & M4V & $23.9 \pm 1.1$ & $60-300$ & $35 \pm 8$ & $19.8 \pm 0.9$\\ FU~Tau~b $^{(13)}$ & M7,25 & $140.0$ & $1.0$ & $15.0$ & $800$\\ CD-35~2722~b$^{(14)}$ & M1V & $21.3 \pm 1.4$ & $100 \pm 50$ & $31.0 \pm 8.0$ & $67.0 \pm 4.0$\\ GJ~3629~B$^{(12)}$ & M3V & $22 \pm 3$ & $30_{-13}^{+30}$ & $46\pm 6$ & $6.5\pm 0.5$\\ 2MASS~J15594729+4403595~B$^{(12)}$ & M2 & $33 \pm 4$ & $120$ & $43\pm 9$ & $187\pm 23$\\ UScoCTIO~108~b$^{(15)}$ & M7 & $145.0 \pm 2.0$ & $11 \pm 2$ & $16.0_{-2.0}^{+3.0}$ & $670.0$\\ GSC~6214-210~b$^{(16)}$ & M1 & $145.0 \pm 14.0$ & $11 \pm 2$ & $17.0 \pm 3.0$ & $320.0$\\ Oph~11~b$^{(17)}$ & M9 & $145.0 \pm 20.0$ & $11 \pm 2$ & $21.0 \pm 3.0$ & $243 \pm 55.0$\\ 1RXS~J235133.3+312720~B$^{(12)}$ & M2 & $50 \pm 10$ & $120$ & $32.0 \pm 6.0$ & $120 \pm 20$\\ %2M 2140+16 b & 21:40:29.0 & +16:25:18 & - & $25.0 \pm 10.0$ & - & $20.0_{ -20.0}^{+80.0}$ & $3.53 \pm 0.15$\\ [1ex] \end{tabular} \label{Mcompanions} \end{center} \end{table*} Both populations of Planetary Mass Companions (hereafter PMCs, $<14M_\mathrm{Jup}$; close-in RV and transit PMCs on one hand, and wide-orbit DI PMCs on the other) challenge our view on planetary formation. Core accretion \citep[CA,][]{Pollack.1996, Alibert.2004} followed by planet migration is the preferred model to explain both the formation of Solar System giant planets and Hot Jupiters found around solar-type stars. Nevertheless, CA does not easily explain the massive gas giants imaged around low-mass stars, especially at large separations, even in the case of pebble accretion \citep{Lambrechts.2012}. Indeed, the mass of the planets are of the same order of magnitude (e.g. 2M0103b, \citealt{Delorme.2013}) or even greater (2M1207b, \citealt{Chauvin.2004}) than the total mass of the protoplanetary disc ($\approx$~10\% of the star's mass). On the other hand, Gravitational Instability \citep[GI,][]{Boss.2011,Cameron.1978} provides an interesting alternative to CA to explain the formation of wide-orbit, massive planets within the disk of a low-mass star. Other initial conditions for the formation of PMCs, such as the position of the ice line and the stellar metallicity, impact their formation mechanisms. For instance, there are indications that planets with masses between 10~$M_\mathrm{Earth}$ and 4~$M_\mathrm{Jup}$ orbiting metal-poor stars are found at larger separations \citep{Adibekyan.2013}. \\ Studying low-mass stars provides an opportunity to test the formation of PMCs at the low end of the stellar mass function. They also allow the study of a wider range of companion to star mass ratios. Different mechanisms might explain the formation of the planets, as a function of their mass ratio. In this paper, we will investigate whether there are different populations of planets orbiting low and high-mass stars. \\ From an observational point of view, \citet{Bonfils.2011} found that the rate of close-in giant planets ($m \sin(i)=100-1000~M_\mathrm{Earth}$) around M-type stars is low ($f<0.01$ for $P=1-10~d$, $f=0.02^{+0.03}_{-0.01}$ for $P=10-100~\mathrm{d}$), while close-in super-Earths are much more numerous ($f=0.35^{+0.45}_{-0.11}$ for $P=10-100d$). Concerning higher-mass stars, \citet{Rameau.2013} find that between 10.8\% and 24.8\% of A to F type stars host at least one planet defined by the parameter intervals [1,13]~$M_\mathrm{Jup}$ and [1,1000]~AU. Recently, \citet{Bowler.2014} showed that few M dwarfs host giant planets closer than for <1000~AU and argue that there is no evidence of a relation between the wide-orbit giant planet (>10~AU) frequency and the stellar mass (for A, FGK and M stars). However, other studies combining different techniques of detection (RV and DI, micro lensing and RV) have shown that GPs are less frequent around M dwarfs \citep[respectively]{Montet.2014,Clanton.2014}. \\ In this paper, we present the results of the M-dwArf Statistical Survey for direct Imaging of massiVe Exoplanets (MASSIVE). Our sample is composed of 58 young and nearby M dwarfs observed in $L^\prime$-band for which we have follow-up data for all our candidate companions. \\ \\ We present the MASSIVE survey in Section~\ref{MASSIVE_survey}. Section~\ref{Frequency} details our Bayesian approach to derive the frequency of planetary companions orbiting low-mass stars. In Section~\ref{Bayesian_Analysis}, we explore the influence of the stellar host mass on the planet occurrence by comparing the planetary companion frequency around low-mass stars in the MASSIVE survey and around a similar VLT/NaCo survey that targeted higher mass, A-F stars \citep{Rameau.2013}. We present our conclusions in Section~\ref{Conclusion}. In all the following Sections, we consider three types of companions: PMCs with masses $<14M_\mathrm{Jup}$, brown dwarfs with masses between $14$ and $80M_\mathrm{Jup}$, and SCs that bring together PMCs and brown dwarfs ($<80M_\mathrm{Jup}$).

\label{Conclusion} We conducted an adaptive optics survey for exoplanets orbiting nearby young low mass stars using $L^\prime$-band observations with NaCo at VLT. Two planetary mass companions were detected and previously reported \citep{Delorme.2013,Chauvin.2004}, but no additional companions were detected. From our significant data set (54 targets), we derived the planetary mass companion (and substellar companion) frequency, defined here by the probability that a star hosts at least one planetary (substellar) mass companion. We used Bayesian statistics with an optimized conjugate prior, and we selected the sub-range in planet mass and semi-major axes where our data provides the best information, so that our statistics does not consider parameter ranges where our survey is not sensitive to companions. Our range of interest corresponds to planets more massive than 2~$M_\mathrm{Jup}$ and a semi-major axis between 8 and 400 AU. Using this range, we determined with a 68\% confidence that substellar companion frequency is $4.4^{+3.2}_{-1.3}$\%, and that planetary mass companion frequency is $2.3^{+2.9}_{-0.7}$\%. Considering all companions regardless of planet-to-star mass ratio, a Monte Carlo comparison between the substellar companion frequencies of the low-mass stars in the MASSIVE survey and a higher-mass A-F type star survey shows that there is a $\sim$~74\% probability that there are two distinct substellar companion populations, one orbiting around the low-mass stars and one orbiting around the higher-mass A-F type stars. We also found a $\sim$~75\% probability considering only low-mass ratios (Q<1\%). These results suggest that the frequency of imaged wide-orbit substellar companions is correlated with the stellar mass, in agreement with theoretical expectations \citep[see e.g.][]{Laughlin.2004}. For intermediate mass ratios (1\%<Q<5\%), the same comparison indicates a statistically significant similarity between the substellar companion populations around the low-mass star and high-mass star samples. We acknowledge that our results showing that substellar companion frequency could be correlated to stellar mass are still moderately significant. We therefore need more observations to confirm this important result.

1607.06892

1607

1607.07540.txt

\ik has discovered hundreds of systems with multiple transiting exoplanets which hold tremendous potential both individually and collectively for understanding the formation and evolution of planetary systems. Many of these systems consist of multiple small planets with periods less than $\sim$50 days known as Systems with Tightly-spaced Inner Planets, or STIPs. One especially intriguing STIP, Kepler-80 (KOI-500), contains five transiting planets: f, d, e, b, and c with periods of 1.0, 3.1, 4.6, 7.1, 9.5 days, respectively. We provide measurements of transit times and a transit timing variation (TTV) dynamical analysis. We find that TTVs cannot reliably detect eccentricities for this system, though mass estimates are not affected. Restricting the eccentricity to a reasonable range, we infer masses for the outer four planets (d, e, b, and c) to be $6.75^{+0.69}_{-0.51}$, $4.13^{+0.81}_{-0.95}$, $6.93^{+1.05}_{-0.70}$, and $6.74^{+1.23}_{-0.86}$ Earth masses, respectively. The similar masses but different radii are consistent with terrestrial compositions for d and e and $\sim$2\% H/He envelopes for b and c. We confirm that the outer four planets are in a rare dynamical configuration with four interconnected three-body resonances that are librating with few degree amplitudes. We present a formation model that can reproduce the observed configuration by starting with a multi-resonant chain and introducing dissipation. Overall, the information-rich Kepler-80 planets provide an important perspective into exoplanetary systems.

\setcounter{footnote}{0} \ik has solidified the existence of a new population of planetary systems that consist of multiple small (1-3 Earth radii), nearly coplanar planets with periods concentrated around 5-50 days \citep[e.g.,][]{2011ApJ...736...19B, 2011ApJS..197....8L} now known as STIPs or Systems with Tightly-spaced Inner Planets (see $\S$\ref{stips}). Though the earliest examples were discovered with radial velocity surveys \citep[e.g.,][]{2006Natur.441..305L,2011arXiv1109.2497M}, \emph{Kepler} has significantly expanded our understanding of this population with its discovery of hundreds of stars with multiple transiting planet candidates \citep[e.g.][]{2011ApJ...736...19B, 2011ApJS..197....8L,2015arXiv151206149C}. Technically, some of the multi-transiting systems are composed of planet candidates, and some of these candidates might be false positives. While $Kepler$'s false positive rate is generally low due to careful candidate vetting, it is clear that candidates in systems with multiple \ik candidates are much more likely to be planets \citep{RH10, 2010ApJ...713L.140L,2011ApJS..197....8L}, especially those with 3 or more candidates \citep{2012ApJ...750..112L,2014ApJ...784...45R,2014ApJ...784...44L}, whose purity is near 99\%. This purity is only one example of the value of multi-transiting systems, with many more aspects discussed in \citet{RH10} and subsequent works. Complementary to studies of the multi-transiting systems as an ensemble are investigations into individual systems to infer the masses of the planets from their mutual gravitational interactions as manifested in deviations from a perfectly periodic sequence of transits. These non-Keplerian motions are characterized by measuring how the times of transits are non-periodic, whence the now-common name of Transit Timing Variations (TTVs). While the value of TTVs was predicted before \ik \citep[e.g.,][]{2005MNRAS.359..567A, 2005Sci...307.1288H}, \ik has measured hundreds of statistically significant TTVs \citep{2013ApJS..208...16M,2015arXiv150400707R}, allowing for precise and numerous mass estimates, particularly of small planets that are difficult to detect with Radial Velocity (RV) measurements \citep[see, e.g., ][]{2014ApJS..210...20M,2014PNAS..11112616F, 2015ARA&A..53..409W,2016MNRAS.tmp...32S}. TTVs in multi-transiting systems are particularly valuable, since the combination of masses and radii can yield multiple density measurements in a single system. In this work, we present such a detailed study for the transiting planets of \kstar (also KOI-500, KIC 4852528, 2MASS J19442701+3958436). \kstar has the historical distinction of being the first system identified with 5-candidates. The outer two candidates in this system were confirmed by observing anti-correlated transit timing variations and were called Kepler-80b and Kepler-80c with periods of 7.1 and 9.5 days, respectively \citep{2013ApJS..208...22X}. The middle two candidates were validated by \citet{2014ApJ...784...44L} and \citet{2014ApJ...784...45R} and were named Kepler-80d and Kepler-80e with periods of 3.1 and 4.6 days respectively. \citet{2016ApJ...822...86M} recently validated the innermost 1.0-day period planet, now Kepler-80f. In this work, we are able to measure the masses of the outer four planets, identify their dynamical relationship, and simulate their formation. We present the observations and data ($\S$\ref{obs and data}) and the inferred stellar properties ($\S$\ref{stellar props}). We then turn to a detailed analysis of the TTV data for \kstar ($\S$\ref{ttv}), including a validation of our fitting procedure and assumptions ($\S$\ref{validation}). With mass estimates, we investigate the physical properties of the planets, including the mass fraction of H/He gas ($\S$\ref{planet props}). We then explore the dynamical configuration of Kepler-80, finding its planets to be in multiple three-body resonances ($\S$\ref{dynamics}). A simulation showing the formation of the system that achieves the observed three-body resonant configuration is presented in $\S$\ref{formation}. Finally, we summarize our conclusions and look forward to future observational and theoretical investigations ($\S$\ref{conclusion}). %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\label{conclusion} \ik has provided us with a wealth of data on the architectures of planetary systems. Herein, we investigated the dynamically intriguing Kepler-80 system (planets f, d, e, b, and c in order of period) and came to several interesting conclusions. A self-consistent dynamical analysis of the system, using TTV fitting under the assumption of restricted eccentricities, inferred masses for the outer four planets (d, e, b, and c) of $6.75^{+0.69}_{-0.51}$, $4.13^{+0.81}_{-0.95}$, $6.93^{+1.05}_{-0.70}$, and $6.74^{+1.23}_{-0.86}$ Earth masses, respectively. The choice to restrict eccentricities to small values resulted from extensive testing of our fitting technique that showed TTV fits using eccentric models infer accurate mass estimates but inaccurate eccentricities and apsidal angles. Further tests showed that we cannot infer two-body resonance angle libration with Kepler-80 TTVs. Although all four planets have very similar masses, planets d and e are terrestrial and planets b and c have $\sim$2\% (by mass) H/He envelopes assuming Earth-like cores. Their orbits are similar and models suggest that photo-evaporation would have removed $\sim$1\% H/He from all four planets. Though simulations suggest the system has been affected by planetary tides, we did not consider the effect of dissipation on the atmospheric history of the planets. It is unusual to have four well-measured densities in the same system and future comparative planetology may constrain the formation and evolution of their atmospheres. Kepler-80 is very interesting dynamically. The system appears to be long-term stable as long as eccentricities are below $\sim$0.2. The outer four planets in Kepler-80 are in a dynamically rare configuration, with multiple three-body resonances librating with only $\sim$3$^{\circ}$ amplitude. This architecture is the natural result of migration simulations, described herein, where the four outer planets were in a resonant chain and a dissipative forces pushed them wide of nominal two-body resonance locations (while retaining two-body resonance angle libration) and deep into three-body resonances. Kepler-80 should thus play an important constraint on the formulation and evolution of STIPS. Many of these conclusions are fruitful starting points for additional study. To assist in the future observational efforts, we extrapolate our restricted eccentricity bootstrap models $\sim$15 years into the future and provide the transit times and estimated uncertainties in Table \ref{table:ttestimates}. Four years of high-precision coverage from \ik has maintained the uncertainty in near-term (e.g., 2016) transit times to about 10 minutes for each planet; a TT measurement more precise than this will be required to significantly improve the model. For transits with a depth of 0.5-1.6 millimagnitudes and a duration of 2 hours on a $V \simeq 15.2$ magnitude star, useful TTV measurements will require space-based observations with large aperture telescopes. Neither TESS \citep{2014SPIE.9143E..20R} nor CHEOPS \citep{2013EPJWC..4703005B} will have sufficient precision. At the estimated time of PLATO observations of the \ik field \citep{2014ExA....38..249R}, the TT uncertainty for the planets will have grown to about 30 minutes, which may be detectable. Note that these are statistical uncertainty estimates that do not include potential sources of systematic error. \begin{deluxetable*}{lcccc} \tabletypesize{\footnotesize} \tablecolumns{5} \tablewidth{0pt} \tablecaption{Transit Time Predictions\label{table:ttestimates}} \tablehead{ \colhead{Planet} & \colhead{Transit No.} & \colhead{Transit Time (BJD)} & \colhead{$\sigma^+$} & \colhead{$\sigma^-$} } \startdata Kepler-80d & 0 & 2455695.130 & 0.003055367 & 0.004165690 \\ Kepler-80d & 1 & 2455698.202 & 0.003098780 & 0.004228180 \\ Kepler-80d & 2 & 2455701.275 & 0.003295472 & 0.004198860 \\ Kepler-80d & 3 & 2455704.347 & 0.003348749 & 0.004229972 \\ Kepler-80d & 4 & 2455707.419 & 0.003356507 & 0.004354759 \\ Kepler-80d & 5 & 2455710.492 & 0.003385405 & 0.004236598 \\ Kepler-80d & 6 & 2455713.564 & 0.003373803 & 0.004251005 \\ Kepler-80d & 7 & 2455716.637 & 0.003396669 & 0.004194434 \\ Kepler-80d & 8 & 2455719.709 & 0.003519851 & 0.003935169 \\ Kepler-80d & 9 & 2455722.782 & 0.003535066 & 0.003922528 \\ Kepler-80d & 10 & 2455725.854 & 0.003600991 & 0.003827943 \\ \enddata \tablecomments{Predictions of future transit times from our integrations through 2025. These predictions were made using our eccentric bootstrapping models. The columns, from left to right, are: the planet's Kepler name, the transit number (where transit 0 indicates the first transit after the epoch of BJD 2454693), the transit time (BJD), the upper uncertainty on the transit time taken to include the 84$^{th}$ percentile ($\sigma^+$), and the lower uncertainty on the transit time taken to include the 16$^{th}$ percentile ($\sigma^-$). Uncertainties are given in units of days and are about 10 minutes in the near term (2016) and grow to about 30 minutes by 2025. Table \ref{table:ttestimates} is published in its entirety in the electronic edition of the \emph{Astrophysical Journal}. A portion is shown here for guidance regarding its form and content.} \end{deluxetable*} A full photodynamical model of Kepler-80 (with stellar parameters updated after GAIA) is a worthwhile endeavor to somewhat improve mass and eccentricity estimates and uncertainties as well as the covariances. For example, we did not use the known durations to constrain the system, which might help to constrain the eccentricities in a less artificial way. Combination with a Bayesian technique would be particularly powerful, as it would be much less susceptible to overfitting, an issue which plagued our inference of eccentricities, apsidal angles, and two-body resonance libration. Additional investigation into the meaning and origin of the three-body resonances might provide interesting constraints of the formation of this system (e.g., damping timescales), which may be broadly applicable to other STIPs. Kepler-80 has proven to be an information-rich multi-transiting system, and we hope that future endeavors will continue to provide insight into this system with implications for the formation and evolution of planetary systems.

1607.07540

1607

1607.04272_arXiv.txt

{We present a new fourth-order, finite-volume hydrodynamics code named {\sc Apsara}. The code employs a high-order, finite-volume method for mapped coordinates with extensions for nonlinear hyperbolic conservation laws. {\sc Apsara} can handle arbitrary structured curvilinear meshes in three spatial dimensions. The code has successfully passed several hydrodynamic test problems, including the advection of a Gaussian density profile and a nonlinear vortex and the propagation of linear acoustic waves. For these test problems, {\sc Apsara} produces fourth-order accurate results in case of smooth grid mappings. The order of accuracy is reduced to first-order when using the nonsmooth circular grid mapping. When applying the high-order method to simulations of low-Mach number flows, for example, the Gresho vortex and the Taylor-Green vortex, we discover that {\sc Apsara} delivers superior results to codes based on the dimensionally split, piecewise parabolic method (PPM) widely used in astrophysics. Hence, {\sc Apsara} is a suitable tool for simulating highly subsonic flows in astrophysics. In the first astrophysical application, we perform implicit large eddy simulations (ILES) of anisotropic turbulence in the context of core collapse supernova (CCSN) and obtain results similar to those previously reported.} \author{A. Wongwathanarat\inst{\ref{riken},\ref{mpa}} \and H. Grimm-Strele\inst{\ref{wien},\ref{mpa}} \thanks{{\it present address:} NUMECA International, Chauss\'{e}e de la Hulpe, 189,Terhulpsesteenweg, B-1170~Brussels, Belgium} \and E. M\"uller\inst{\ref{mpa}}} \authorrunning{A. Wongwathanarat et al.} \institute{RIKEN, Astrophysical Big Bang Laboratory, 2-1 Hirosawa, Wako, Saitama~351-0198, Japan\label{riken} \and Max-Planck-Institut f\"{u}r Astrophysik, Karl-Schwarzschild-Stra\ss e 1, D-85748~Garching, Germany\label{mpa} \and Institute of Mathematics, University of Vienna, Oskar-Morgenstern-Platz 1, A-1090~Vienna, Austria\label{wien}}

In many astrophysical simulations, such as simulations of convection inside a star and stellar explosions, spherical coordinates are often a preferable choice for integrating the respective partial differential equations. However, the spherical polar coordinates possess coordinate singularities at the coordinate origin and along the north and south poles, which prevent an easy and straightforward implementation of numerical methods. These coordinate singularities result in smaller time steps and, hence, special numerical treatments and/or boundary conditions have to be applied. This leads to a degradation of the efficiency of the employed numerical schemes and may introduce numerical artifacts. For example, in recent state-of-the-art, three-dimensional (3D) simulations of core collapse supernovae (CCSN) performed with time-explicit, finite-volume hydrodynamic codes \citep[e.g.,][]{Melsonetal15b, Lentzetal15}, the flow is modeled in spherical symmetry inside a specified sphere representing the inner core of the proto-neutron star, where the symmetry assumption is justified, to alleviate the restriction of the time step due to the Courant–Friedrichs–Lewy (CFL) condition. In addition, a reflecting boundary condition is usually applied at the coordinate origin. \citet{Mueller15} applied the same spherical symmetry assumption, but he also included the effect of proto-neutron star convection by means of a mixing-length theory \citep[e.g.,][]{WilsonMayle88}. \citet{Lentzetal15} avoided time steps that were too small by using nonuniform angular zones in the polar direction and an azimuthal averaging procedure for grid zones near the two poles. Similarly, \citet{Mueller15} circumvented this problem using a mesh coarsening scheme in 30$^\circ$ cones around the poles whereby the short wavelength noise in the azimuthal direction is filtered out. The problem of the severe time step restriction near the singular points at the poles of a sphere, also known as the pole-problem, has received much attention in various fields of research over the past few decades. In the particular case of finite-volume methods on a structured mesh, two popular solutions to the pole-problem are the cubed sphere grid \citep{Ronchietal96} and the Yin-Yang grid \citep{KageyamaSato04}. The cubed sphere grid is based on a projection of the six sides of a cube onto the surface of a sphere and thus consists of six equidistant (in polar and azimuthal direction) identical grid patches. Except for the one point at the middle of each grid patch, the cubed sphere mesh is nonorthogonal. On the other hand, the Yin-Yang grid is formed by combining two identical low-latitude parts of a spherical polar grid. Hence, it is orthogonal everywhere, allowing for an easy extension of existing numerical codes that are based on an orthogonal mesh. Nevertheless, overlapping grids like the Yin-Yang grid have a drawback. Even if the numerical scheme employed on each grid patch is conservative, this does not ensure global conservation in the computational volume unless one applies a flux correction algorithm at the boundaries between the two grid patches \citep[see, e.g.,][]{Pengetal06, Wongwathanaratetal10}. In contrast, the boundaries of each patch of the cubed sphere mesh coincide perfectly with its neigboring patches, \ie, there is no overlap between grid patches. Thus, boundary flux corrections can be implemented in a more straightforward manner than in the case of the Yin-Yang grid. Both the cubed sphere grid and the Yin-Yang grid have been used in various astrophysical applications. For instance, the cubed sphere grid was applied in simulations of accretion flows onto magnetized stars \citep[e.g.,][]{Koldobaetal02, Romanovaetal12} and accretion disks around rotating black holes \citep{Fragileetal09}. The Yin-Yang grid was used, for example, in CCSN simulations \citep[e.g.,][]{Wongwathanaratetal15, Melsonetal15a}, simulations of type Ia supernova remnants \citep{WarrenBlondin13}, and calculations of the solar global convection \citep{Hottaetal14}, solar corona \citep[e.g.,][]{Jiangetal12, Fengetal12}, and coronal mass ejections \citep{Shiotaetal10}. While the cubed sphere grid and the Yin-Yang grid circumvent the pole-problem, the singularity at the coordinate origin remains in both cases. One possible solution is to supplement the cubed sphere mesh or the Yin-Yang mesh with a Cartesian mesh at small radii to cover the central part of the sphere. An example of such a grid arrangement, called the Yin-Yang-Zhong, was recently developed by \citet{HayashiKageyama16}. Using this approach, the numerical code must be able to deal with different coordinate systems on different grid patches. In addition, global conservation is hampered by the overlap of the Cartesian grid patch with the spherical grid patches. On the other hand, the GenASiS code developed by \citet{Cardalletal14} handles both the pole-problem and the singularity at the coordinate origin by employing the adaptive mesh refinement (AMR) approach. Instead of a block-structured AMR framework, GenASiS uses a more flexible cell-by-cell refinement in the Cartesian coordinate system. The cell-by-cell refinement can generate a centrally refined mesh, achieving higher and higher resolutions toward the origin of the sphere, which is a desirable grid property for many astrophysical applications. Nevertheless, a possible disadvantage of the cell-by-cell refinement approach is the cost of a complicated data communication on large-scale machines. \citet{Calhounetal08} proposed grid mappings for circular or spherical domains that work on a logically rectangular mesh. These grid mappings can be used together with the mapped-grid technique in which one formulates the governing equations in an abstract (singularity-free) computational space instead of in a physical space and applies a coordinate transformation between them. A great advantage of the mapped-grid approach is that the computational domain can always be discretized by an equidistant Cartesian mesh. Therefore, any numerical scheme formulated for a Cartesian equidistant grid can be applied in a straightforward manner. The mapped-grid technique is widely used in engineering applications, where complicated geometries are common \citep[see, e.g.,][]{Leveque02}. In the field of astrophysics, the mapped-grid approach is not very well known and has only been used up to now in a few numerical codes {\citep{KifonidisMueller12,Grimmstreleetal14,Miczeketal15}. \citet{Colellaetal11} introduced a new class of high-order (better than second-order), finite-volume methods for mapped coordinates. The methods are based on computing face-averaged fluxes at each face of a control volume using high-order quadrature rules rather than the mid-point rule as in standard second-order accurate schemes. They demonstrated the capability of their approach to solve an elliptic equation and a scalar, linear hyperbolic equation up to fourth-order accuracy. An extension of the method to solve nonlinear hyperbolic equations was presented for Cartesian coordinates by \citet{McCorquodaleColella11} and for mapped coordinates by \citet{Guziketal12}. This extension is nontrivial because it is necessary to perform nonlinear transformations between point and (zone and face) averaged values of the conserved variables and fluxes ensuring fourth-order accuracy. Motivated by the works of \citet{Calhounetal08} and \citet{Colellaetal11}, we have developed a new numerical code for astrophysical applications. This code, which we named {\sc Apsara,} is based on the fourth-order implementation of the high-order, finite-volume method in mapped coordinates by \citet{Colellaetal11}. The code is extended to solve the Euler equations of gas dynamics using a high-resolution, shock-capturing (HRSC) method on arbitrary, structured curvilinear grids in three spatial dimensions as described in \citet{Guziketal12}. This code robustly captures shock waves and discontinuities, while obtaining high-order accurate results in smooth flow regions. On the other hand, simulations of flows on complex grid geometry are possible thanks to the mapped-grid technique. These combined features make the code suitable for simulating a wide range of astrophysical problems. The paper is organized as follows. In Section~\ref{sec:numerics}, we summarize the methods by \citet{Colellaetal11}, \citet{McCorquodaleColella11}, and \citet{Guziketal12} for simulating 3D compressible flows in mapped coordinates. In Section~\ref{sec:mappings}, examples of mapping functions are introduced, which are employed in our test calculations. In Section~\ref{sec:tests}, using a set of hydrodynamic tests, we demonstrate the capability of {\sc Apsara} to compute fourth-order accurate solutions for smooth flows without discontinuities. In addition, we show the ability of {\sc Apsara} to accurately calculate low-Mach number flows. As an astrophysical application, we also performed turbulence simulations in the context of CCSN with {\sc Apsara}. We discuss the results of these simulations in Section~\ref{sec:turbulence}. Finally, we summarize the code properties in Section~\ref{sec:summary}, and mention some plans for future extensions of {\sc Apsara}.

\label{sec:summary} We have developed a new numerical code for simulating astrophysical flows called {\sc Apsara}. The code is based on the fourth-order implementation of a high-order, finite-volume method for mapped coordinates proposed by \citet{Colellaetal11}. An extension of the method for solving nonlinear hyperbolic equations following \citet{McCorquodaleColella11} and \cite{Guziketal12} is applied to solve the Euler equations of Newtonian gas dynamics. {\sc Apsara} comprises great flexibility concerning grid geometry because of the implemented mapped-grid technique, which makes the code suitable for a wide range of applications. By defining a mapping function, which describes the coordinate transformation between physical space and an abstract computational space, the governing equations written in Cartesian coordinates in physical space are transformed into equations of similar form in the computational space. The physical domain of interest can then be discretized as a general structured curvilinear mesh, while the computational space is still discretized by an equidistant Cartesian mesh that allows for easy and efficient implementation of numerical algorithms. We strictly follow the implementation described in \citet{Colellaetal11} for computation of metric terms on faces of control volumes to ensure that {\sc Apsara} preserves the freestream condition. The freestream property is crucial for grid-based codes in curvilinear coordinates because violation of this condition can introduce numerical artifacts as shown, for example, in \citet{Grimmstreleetal14} for the WENO-G finite difference scheme \citep{Nonomuraetal10,Shu03}. The method of \citet{Colellaetal11} can be extended to a higher order of accuracy in time as described in \citet{BuchmuellerHelzel14}. However, RK schemes of order higher than four are considerably more expensive because many more integration stages are needed in this case \citep{Ketcheson08}. Concerning fourth-order accurate RK schemes, the ten-stage algorithm SSPRK(10,4) of \citet{Ketcheson08} has the advantage of possessing the strong stability-preserving (SSP) property, which is not the case for the classical RK4 scheme. In this work, however, we always used the RK4 scheme and the extension to the SSPRK(10,4) scheme is subject of the future development of {\sc Apsara}. We validated our numerical code by simulating a set of hydrodynamic tests problems. In the case of smooth solutions, \ie, for the linear advection of a Gaussian profile, the propagation of a linear acoustic wave, and the advection of a nonlinear vortex, {\sc Apsara} exhibits fourth-order accuracy both on a Cartesian mesh and the sinusoidally deformed mesh considered by \citet{Colellaetal11}. However, our results from the advection of a nonlinear vortex test also revealed that the order of accuracy is severely degraded when using the circular mapping proposed by \citet{Calhounetal08}, \ie, the grid smoothness directly impacts the achievable order of accuracy of the scheme. Our finding agrees with results obtained by \citet{LemoineOu14}, who used a modified version of grid mappings suggested by \citet{Calhounetal08}. \citeauthor{LemoineOu14} also achieved roughly first-order convergence on these nonsmooth grid mappings. Thus, we conclude that the singularity-free mappings for circular and spherical domains of \citet{Calhounetal08} are not suitable for use in conjuction with high-order, finite-volume methods. We quantified how the grid smoothness influences the order of accuracy of the high-order method of \citet{Colellaetal11} using the advection of a nonlinear vortex across a 2D mesh varying the nonequidistant grid spacing in one coordinate direction systematically. We found that the grid expansion ratio, i.e., the ratio of the cell size between two neighboring zones, must be less than $\sim$1.5 in order to preserve fourth-order accuracy. This result can be regarded as a guideline for the grid setup in the case of a logarithmically spaced radial grid that is often employed in astrophysical simulations in spherical geometry. Motivated by the work of \citet{Miczeketal15}, which demonstrated a failure of a Godunov-type scheme employing a Roe flux in the low-Mach number regime, we performed simulations of the Gresho vortex test varying the maximum Mach number of the vortex from $10^{-1}$ to $10^{-4}$. Our results show that the decrease of the kinetic energy of the vortex is independent of the maximum Mach number when computed with {\sc Apsara}. In all our simulations the kinetic energy had decreased by less than $\approx 0.5\%$ at the end of the simulations, \ie, after one complete rotation of the vortex. This contrasts with solutions calculated with {\sc Prometheus}, which uses the dimensionally-split PPM method as do many other codes in astrophysics. The kinetic energy of the vortex decreased more and more rapidly when the maximum Mach number was decreased. In the lowest Mach number case, the Gresho vortex vanished completely and the kinetic energy was reduced by more than half of its initial value because of the numerical dissipation present in {\sc Prometheus}. The superior performance of the high-order method of \citet{McCorquodaleColella11} for highly subsonic flows is further supported by our results obtained for the Taylor-Green vortex problem. We simulated this test for two reference Mach numbers and compared the results computed with {\sc Apsara} and {\sc Prometheus}. For a reference Mach number $0.29$ we obtained similar results for both codes, whereas, for a reference Mach number $10^{-2}$, {\sc Apsara} achieved a larger numerical Reynolds number for the same resolution, implying a lower numerical viscosity than {\sc Prometheus}. To demonstrate {\sc Apsara}'s performance on an astrophysical application, we performed ILES of anisotropic turbulence in a periodic Cartesian domain following the numerical setup used by \citet{Radiceetal15}. The characteristic Mach number of the flow was $0.37$, and the acceleration field driving the turbulence was enhanced along one coordinate direction to mimic the turbulent flow conditions in a CCSN. We obtained very similar results as those reported by \citet{Radiceetal15}. At low resolution the energy spectra are dominated by the bottleneck effect at intermediate scales. We only began to recover the inertial range for a narrow range of wavenumbers at the highest grid resolution of 512 zones per coordinate direction. On the other hand, the effective viscosity as a function of wavenumber displayed fewer variations than that of \citet{Radiceetal15}. For problems to be simulated in a spherical domain, which we are particularly interested in, the singularity-free grid mappings of \citet{Calhounetal08} result in a loss of convergence order. Therefore, we plan to implement the mapped multiblock grid technique into {\sc Apsara} following the strategy of \citet{McCorquodaleetal15}, who extended the high-order, finite-volume method of \citet{Colellaetal11} to include this technique. They described an algorithm for data communication between blocks such that high-order accuracy is maintained. We also plan to implement additional physics into {\sc Apsara}, for example, self-gravity, a nuclear EOS, and a nuclear reaction network, to tackle multiphysics astrophysical problems such as CCSN and stellar convection.

1607.04272

1607

1607.08606_arXiv.txt

Binary-star exoplanetary systems are now known to be common, for both wide and close binaries. However, their orbital evolution is generally unsolvable. Special cases of the $N$-body problem which are in fact completely solvable include dynamical architectures known as central configurations. Here, I utilize recent advances in our knowledge of central configurations to assess the plausibility of linking them to coplanar exoplanetary binary systems. By simply restricting constituent masses to be within stellar or substellar ranges characteristic of planetary systems, I find that (i) this constraint reduces by over 90 per cent the phase space in which central configurations may occur, (ii) both equal-mass and unequal-mass binary stars admit central configurations, (iii) these configurations effectively represent different geometrical extensions of the Sun-Jupiter-Trojan-like architecture, (iv) deviations from these geometries are no greater than ten degrees, and (v) the deviation increases as the substellar masses increase. This study may help restrict future stability analyses to architectures which resemble exoplanetary systems, and might hint at where observers may discover dust, asteroids and/or planets in binary star systems.

The Trojan asteroids hosted by Jupiter and Neptune help constrain the formation of the Solar System \citep{chilit2005,morbidelli2005,lyketal2009} and represent powerful examples of how central configurations -- a topic often limited to the mathematics literature -- apply to a real planetary system. Extrasolar planetary systems potentially provide other opportunities, with their diverse architectures, and our increasing capacity to observe sub-Earth mass solid bodies \citep[e.g.][]{kieetal2014,vanetal2015} and clumps of dust \citep[e.g.][]{rapetal2015,haretal2016}. Central configurations produce exact solutions to the $N$-body problem, which is otherwise generally unsolvable. Hence, characterizing the existence and architectures of all central configurations for all $N$ is a holy grail of celestial mechanics (see \citealt*{moeckel1990} and Chapter 2 of \citealt*{libetal2015}). They can aid in understanding the long-term dynamics of planetary systems, identifying the masses and locations of objects which remain stable, and targeting searches for hidden objects. Strictly, a central configuration is an arrangement of point masses which satisfy \begin{equation} \Lambda \vec{r}_i = \sum_{j \ne i}^{N} \frac{M_j \left(\vec{r}_j - \vec{r}_i \right)}{r_{ij}^3} , \label{orig} \end{equation} \noindent{}where $\vec{r}$ is the barycentric position vector, $M$ is mass, $i,j$ are object indices, $r_{ij} = |\vec{r}_i - \vec{r}_j|$, $N$ is the total number of bodies, and $\Lambda$ is a constant. When $N=2$, all systems are central configurations. For $N=3$, there exists two classes of central configurations: when all three bodies are co-linear \citep{euler1767} and when all three bodies form an equilateral triangle \citep{lagrange1772}. Only some co-linear systems (for correctly chosen masses and distances) are central configurations, whereas all equilateral triangle configurations (for any masses or distances) are central configurations. Co-linear central configurations give rise to the quintic equation which yields Hill stability boundaries \citep{marboz1982}, a now well-used concept in exoplanetary dynamics \citep{davetal2014}. A stability analysis of equilateral triangle architectures reveals the beautiful result that because the Jupiter/Sun mass ratio is less than $(1/2 - \sqrt{23/108}) \approx 0.039$, Jupiter's Trojan asteroids can remain in stable orbits. \begin{figure*} \includegraphics[width=18cm]{CCSetup4.eps} \caption{Six cases of planar four-body architectures containing two stars and one axis of symmetry. Stars are labelled in red and with five-pointed star symbols, and have either unequal masses (left-hand panels) or equal masses (right-hand panels). The green check marks indicate which cases were found to contain central configurations that are potentially applicable to binary star planetary systems. } \label{sumfig} \end{figure*} \begin{figure*} \centerline{ \includegraphics[width=8cm]{Case1_mass1.eps} \ \ \ \includegraphics[width=8cm]{Case1_mass2.eps} } \caption{ The allowed Case \#1 regions that can produce central configurations in binary-star planetary systems. The plots demonstrate that $\alpha$ and $\beta$ are each limited to within a few degrees of $30^{\circ}$. The colour contours indicate how much more restrictive this range is for different object masses: For $M_2/M > 10^2$, both $\alpha$ and $\beta$ are confined to within a tenth of a degree of $30^{\circ}$, whereas for $M_1/M > 10^2$, $\alpha$ and $\beta$ lie along the lower axis of the allowed region. Asteroids, pebbles or dust (with $M_1/M \gg 10^{10}, M_2/M \gg 10^{10}$) effectively lie at the critical point indicated by the red dot. } \label{case1mass} \end{figure*} \begin{figure} \includegraphics[width=8cm]{Case1_distance.eps} \caption{ The distance between both stars in Case \#1. Throughout the region which allows for central configurations to occur, both stars form a near-equilateral triangle shape (with side-length ratios differing by no more than five per cent) with each of the other objects. } \label{case1distance} \end{figure} The $N=4$ case is considerably more complex, but may lend itself well to binary stellar systems which contain dust, minor planets or major planets. A recent significant advancement in the $N=4$ case was provided by \cite{erdczi2016}, who fully characterized all central configurations of four planar bodies, at least two of which contain equal masses, and with an axis of symmetry between the other two bodies. The importance of the result was promoted by \cite{hamilton2016}, who suggested several potential mathematical extensions. Rather than pursue any of those perhaps daunting challenges, I simply wish here to take stock of the results of \cite{erdczi2016}, and determine broadly if and how they may be applicable to exoplanetary systems, in particular those containing two stars. I cover all four-body planar cases with one axis of symmetry and two equal masses about that axis. The two stars are not restricted to lie along this axis, nor restricted to have equal masses. First, in Section 2, I repackage the method for determining these central configurations and provide a straightforward algorithm for the computation. Then, in Section 3 I restrict this class of configurations to masses and distances that correspond to realistic astronomical systems. Doing so allows me to pinpoint, robustly, architectures that might warrant future studies by exoplanet theorists and observers alike. I discuss the implications of this study in Section 4, and then briefly conclude in Section 5.

This study demonstrates that central configurations in binary-star planetary systems may exist, and quantifies the extent of the existence region. Whether objects residing in the immediate vicinity of these configurations are long-term stable and have plausible dynamical origins are different questions. A proper stability analysis of planar four-body central configurations with two equal masses and an axis of symmetry connecting the unequal masses has not yet been carried out, and could represent a significant undertaking. The results here might help focus that effort by quantifying the variations in geometries and mass that result from the most basic planetary system-based restrictions. Regarding the likelihood of a planetary system forming or settling into one of the central configurations discussed here, the chances probably increase as the substellar masses decrease. All these central configurations include objects which are nearly co-orbital or nearly co-linear. No planets as massive as Jupiter have yet to be found in such architectures. The largest known co-orbital objects in the solar system are Janus ($1 \times 10^{18}$ kg) and Epimetheus ($5 \times 10^{17}$ kg), and in extrasolar systems is a minor planet with its disrupted fragments orbiting white dwarf WD 1145+017; the mass of this minor planet is likely to be about $10^{20}$~kg, a tenth of the mass of Ceres \citep{guretal2016,rapetal2016}. Multiple co-orbital masses above $10^{23}$~kg are roughly assumed to become unstable at the 20 per cent level \citep{veretal2016}. Three or more co-linear planetary objects are not known to exist in stable configurations. The fact that the majority of planets in the solar system host at least one Trojan asteroid might indicate that capture into that central configuration is common. If so, the extended Trojan-analogue rhombus configurations from Cases \#1 and \#2 might also be easily populated by asteroids or smaller bodies such as pebbles or dust or gas. In fact, remnant protoplanetary disc features might reside in the locations indicated by the red dots on Figs. \ref{case1mass}-\ref{case2distance}. Finally, I emphasize that these configurations are fully scalable with $y$ and $M$. The binary stars could have wide separations of $10^4$ au or close separations of $10^{-2}$ au. Such close separations are unlikely to host objects in central configurations because any bodies so close to those stars would either be drawn in and destroyed due to tides, or blown away by winds. For particularly wide separations, comparable to where stellar flybys or planet-planet scattering might deposit planets or planetary debris into an exosystem \citep{perkou2012,varetal2012}, capture into a central configuration is a distinct possibility.

1607.08606

1607

1607.05935_arXiv.txt

{Electromagnetic waves arise in many area of physics. Solutions are difficult to find in the general case.}{ In this paper, we numerically integrate Maxwell equations in a 3D spherical polar coordinate system.}{Straightforward finite difference methods would lead to a coordinate singularity along the polar axis. Spectral methods are better suited to deal with such artificial singularities related to the choice of a coordinate system. When the radiating object is rotating like for instance a star, special classes of solutions to Maxwell equations are worthwhile to study such as quasi-stationary regimes. Moreover, in high-energy astrophysics, strong gravitational and magnetic fields are present especially around rotating neutron stars.}{ In order to study such systems, we designed an algorithm to solve the time-dependent Maxwell equations in spherical polar coordinates including general relativity as well as quantum electrodynamical corrections to leading order. As a diagnostic, we compute the spindown luminosity expected from these stars and compare it to the classical i.e. non relativistic and non quantum mechanical results.}{It is shown that quantum electrodynamics leads to an irrelevant change in the spindown luminosity even for magnetic field around the critical value of $\numprint{4.4e9}$~\si{\tesla}. Therefore the braking index remains close to its value for a point dipole in vacuum namely $n=3$. The same conclusion holds for a general-relativistic quantum electrodynamically corrected force-free magnetosphere. }

Nature do not offer us much places in the Universe where to test our current theories of gravity and electromagnetism {\it simultaneously}. However and fortunately strong magnetic and gravitational fields exist inside and around neutron stars. They represent valuable laboratories to check our current theories in the strong field regime. Curvature of space-time is important because of the stellar compactness of about \begin{equation} \label{eq:compacite} \Xi = \frac{\Rs}{R} \approx 0.345 \, \left( \frac{M}{1.4~M_\odot} \right) \, \left( \frac{R}{12 \textrm{ km}} \right)^{-1} \end{equation} where $R$ is the neutron star radius, $M$ its mass, $\Rs=2\,G\,M/c^2$ its Schwarzschild radius, $c$ the speed of light and $G$ the gravitational constant. Moreover, neutron stars are strongly magnetized objects, harbouring fields as high as the critical value of $\BQ\approx\numprint{4.4e9}$~\si{\tesla} or even higher. These regimes of strong gravity and magnetic fields are unreachable on Earth even separately. Since the exact analytical solution for a static dipole in general relativity (GR) found by \cite{Ginzburg1964} and those for multipolar terms in a spherically symmetric vacuum gravitational field by \cite{1972PhRvD...6.1476W}, several authors looked deeper into the effect of rotation with emphasizes to neutron stars. In vacuum, Maxwell equations remain linear even in a background gravitational field. This helped \cite{2001MNRAS.322..723R, 2002MNRAS.331..376Z, 2004MNRAS.352.1161R} to compute the electromagnetic field in the exterior of a slowly rotating neutron star. They gave approximate analytical expressions for the external electromagnetic field close to the neutron star which have later also been reported by \cite{2013MNRAS.433..986P}. \cite{2004MNRAS.348.1388K} extended the previous work by solving numerically the equations for the oblique rotator in vacuum in general relativity. They retrieve \cite{2001MNRAS.322..723R} results close to the surface and the Deutsch solution \citep{1955AnAp...18....1D} for distances larger than the light cylinder~$r\gg\rlight$ where $\rlight=c/\Omega$ and $\Omega$ is the rotation speed of the star. Whereas quantum electrodynamics (QED) effects are known to be relevant for wave propagation in the birefringent vacuum as described by \cite{1971AnPhy..67..599A, 1979PhRvD..19.2868H, 1986ApJ...302..120A, 1988MNRAS.235...51B} or in the review by \cite{1992herm.book.....M, 2006RPPh...69.2631H, 2015SSRv..191...13L}, less attention has been focused so far to the whole picture of the magnetosphere. Let us mention the work of \cite{1997JPhA...30.6475H} who computed corrections to a dipole to first order for any strength of the magnetic field following \cite{1936ZPhy...98..714H} effective Lagrangian. This result has recently been generalized by \cite{2016MNRAS.456.4455P} taking into account the curvature of space-time following the 3+1~formalism developed by \cite{2015MNRAS.451.3581P}. We apply this formalism to a rotating monopole and dipole. Such corrections are relevant for magnetars, those neutron stars with the strongest magnetic fields known in the Universe \citep{2015RPPh...78k6901T}. One of the mystery of the global electrodynamics of pulsar or neutron star magnetosphere is symbolized by the braking index~$n$ relating the braking torque to the rotation rate~$\Omega$ of the star by $\dot\Omega \propto - \Omega^n$ where a dot means derivative with respect to time. For a pure dipole rotating in vacuum, it should be very close to $n=3$, see for instance \cite{1973ApJ...181..161R} for the dipole and the general case of a multipole of order~$\ell$ being $n=2\,\ell+1$ as given in \cite{1991ApJ...373L..69K}, see also \cite{2015MNRAS.450..714P} for the exact expression taking into account the finite size of the star. In the dipole case \cite{2016MNRAS.455.3779P} showed that this result is not affected by the presence of a plasma in the magnetosphere in general relativity. \cite{2015PhRvD..91f3007H} summarize the state of the art in the measurements of pulsar braking indices. They are all less than~3, some of them much smaller, closer to 1 or 1.5 thus definitely ruling out a pure dipolar field in vacuum, force-free (FFE) or MHD regime. See however a very recent outsider reported by \cite{2016ApJ...819L..16A} to have $n=3.15$. Could QED effects account for this discrepancy? \cite{2012EL.....9849001D} claimed that QED can indeed strongly impact on the braking index. Starting from this assertion \cite{2016RAA....16a...9X} proposed to test the hypothesis of superstrong magnetic field in magnetars by inspection of their energy loss that should be dominated by quantum vacuum friction. Even more recently \cite{2016ApJ...823...97C} build on this quantum vacuum friction effect and arrived at the same conclusion within a factor~2. Unfortunately as reported in this work, we do not retrieve their results. Their expression for QED corrections adds a spindown luminosity depending on $\Omega^2$ therefore resembles to radiation from a magnetic monopole very similar to the split monopole solution. Maxwell theory of electromagnetism does not allow radiation from a magnetic monopole in vacuum. Moreover, non linear electrodynamics of a rotating dipole would induce higher multipoles with mode numbers $\ell\geqslant1$ due to non-linearities but never a $\ell=0$ multipole. It is thus very difficult to understand the origin of the quantum vacuum luminosity given by these authors. In this paper, we develop a pseudo-spectral discontinuous Galerkin method in space in the weak formulation to solve Maxwell equations in spherical coordinates using our formalism in general relativity with the effective Euler-Heisenberg QED Lagrangian. The set of equations and the solution techniques are reminded in Section~\ref{sec:GRElectrodynamic}. The algorithm is discussed in depth in Section~\ref{sec:Algorithm}. Results for the dipole in classical flat space-time and with strong field corrections from GR and QED are presented in Section~\ref{sec:ResultatsVide} for vacuum case and in Section~\ref{sec:ResultatsFFE} for FFE case. We conclude about possible extensions of this work in the concluding remarks of Section~\ref{sec:Conclusion}.

\label{sec:Conclusion} Strongly magnetized rotating fields in vacuum such as the one expect in neutron star systems and especially in magnetars are responsible for their electromagnetic activity like pair creation and very high energy emission processes which are effectively observed at Earth. Thus in some sense we get indirect insights into the physics of such strong fields. In this paper we have shown that despite the presence of magnetic field strengths around the critical field, QED corrections would not lead to drastic changes in the global electrodynamics of a neutron star magnetosphere, especially not in the rate of braking through electromagnetic radiation of the large amplitude low frequency electromagnetic wave in vacuum. Filling the magnetosphere with a high density ultra-relativistic pair plasma leading to the force-free regime does not modify this outcome. So far our results have been restricted to fields $B\lesssim\BQ$ because of the first order QED Lagrangian we used. However, for very intense fields $B\gg\BQ$ we do not expect the QED effect to become dominant because asymptotically, the perturbation of the Lagrangian scales as $\ln (B/\BQ)$ and would require unrealistically high fields to become comparable to the unperturbed Lagrangian \citep{LandauLifchitzTome4}. Consequently, our results are fairly robust even in the extreme case of high-B field magnetars.

1607.05935

1607

1607.00384_arXiv.txt

Using deep {\it Hubble Frontier Fields} imaging and slitless spectroscopy from the {\it Grism Lens-amplified Survey from Space}, we study 2200 cluster and 1748 field galaxies at $0.2\leq z\leq0.7$ to determine the impact of environment on galaxy size and structure at stellar masses $\logm>7.8$, an unprecedented limit at these redshifts. Based on simple assumptions---$r_e= f(\Mstel)$---we find no significant differences in half-light radii ($r_e$) between equal-mass cluster or field systems. More complex analyses---$r_e = f(\Mstel, U-V, n, z,\Sigma)$---reveal local density ($\Sigma$) to induce only a 7\%$\pm3\%$ ($95\%$ confidence) reduction in $r_e$ beyond what can be accounted for by $U-V$ color, S\'ersic index ($n$), and redshift ($z$) effects. Almost any size difference between galaxies in high- and low-density regions is thus attributable to their different distributions in properties other than environment. Indeed, we find a clear color--$r_{e}$ correlation in low-mass passive cluster galaxies ($\logm<9.8$) such that bluer systems have larger radii, with the bluest having sizes consistent with equal-mass star-forming galaxies. We take this as evidence that {\it large-$r_{e}$} low-mass passive cluster galaxies are recently acquired systems that have been environmentally quenched without significant structural transformation (e.g., by ram pressure stripping or starvation). Conversely, $\sim20\%$ of {\it small-$r_{e}$} low-mass passive cluster galaxies appear to have been in place since $z\gtrsim3$. Given the consistency of the small-$r_{e}$ galaxies' stellar surface densities (and even colors) with those of systems more than ten times as massive, our findings suggest that clusters mark places where galaxy evolution is accelerated for an ancient base population spanning most masses, with late-time additions quenched by environment-specific mechanisms are mainly restricted to the lowest masses.

\label{sec:intro} In terms of a host of properties---color, star formation activity, structure, morphology---clusters harbor different galaxy populations than average (``field") environments \citep[e.g.,][]{hubble31, dressler80}. The mechanisms that produce these differences has been the subject of intense scrutiny. While evidence of environmental effects have been seen \citep[e.g.,][]{vollmer09, abramson11, mcpartland16, poggianti16}, their roles and relative importance compared to {\it in situ} galaxy evolution remain poorly understood. Indeed, the extent to which clusters are agents that halt galaxy evolution, as opposed to {\it tracers} of regions where it has been accelerated, is still under debate (cf.\ \citealt{peng10} with \citealt{dressler80}, \citealt{thomas05}, \citealt{guglielmo15}, \citealt{abramson16}). One confounding factor is that galaxy-by-galaxy analyses reveal almost no differential environmental effects for systems, e.g., at fixed stellar mass ($\Mstel$) and color \citep{grutzbauch11b}. That is, while galaxy {\it populations} are different in low- and high-density regions, representatives of all parts of parameter space seem to exist everywhere \citep[e.g.,][]{dressler13, dressler16, wu14}.\footnote[9]{Excluding the very most- and very least-massive red objects---dwarf ellipticals and cDs/BCGs---which may never exist in isolation \citep{koester07, geha12}.} This appears to hold even for scaling laws that (seemingly) should contain the signatures of any transformational mechanism, such as the star formation rate--mass and size--mass relations (e.g., \citealt{maltby10, peng10, huertas13, koyama13, allen16}; but see\ \citealt{vulcani10, paccagnella16}, who define environment by spectroscopic membership as opposed to spatial overdensity). However, a key obstacle to many previous investigations has been their relatively high mass limits of $\logm \gtrsim10$. In this regime, a system's self-gravity is strong, perhaps protecting it from environmental influences such as ram pressure stripping or harassment \citep[e.g.,][]{dressler83, moore96, treu03, lin14}. Furthermore, high-mass galaxies might be subject to internal processes---such as feedback from active galactic nuclei, or the suppression of star formation by morphological structures---that act before they enter the cluster, preventing the latter from having any effect at all \citep[e.g.,][]{martig09, hopkins14}. To better constrain the physical processes {\it causally related} to environmental density, targeting the low-mass tail of the galaxy population ($\logm\ll10$) is key. Some studies of the nearest clusters have probed this regime: \citet[][see also \citealt{ferrarese06}]{misgeld11} examined the size--mass relation of Local Group and Coma galaxies at $\logm\gtrsim 6$, and \citet[][]{lisker09} and \citet{toloba15} explored the diversity of dwarf galaxies (dEs, dSphs) in Virgo, with the latter study using kinematical/chemical information to find the evidence of ram pressure stripping-influenced evolution. However, at $z\approx0$, cluster galaxies are so uniformly old that the residual signatures of any transformational mechanisms may be detectable only in the most detailed fossil evidence \citep{mcdermid15}. Shifting focus to $z\sim0.5$ would alleviate this issue by probing an epoch when clusters were still rapidly assembling, and hence more dramatically reshaping their galaxy populations \citep{butcher78}. The recent advent of ultra-deep multi-band imaging and spectroscopy from space enables such studies of low-mass, mid-$z$ galaxies for the first time. In this paper, we use data from the {\it Hubble Frontier Fields} \citep[HFF;][]{lotz16} and {\it Grism Lens-Amplified Survey from Space} \citep[GLASS;][]{schmidt14,treu15} to examine the galaxy populations of clusters and the field at $0.2\leq z\leq0.7$ to a hitherto unexplored mass limit of $\logm>7.8$. We exploit these data to study the dependence of galaxy size on stellar mass and other structural properties as a function of environmental density using an unprecedented sample of over 3900 cluster and field galaxies. We examine these correlations because: (1) processes that depend on environment---e.g., ram pressure stripping, or mergers (rarer in richer systems)---affect galaxy size and structure in significant and well-defined ways \citep[][and many others]{vandokkum10, damjanov11, newman12, nipoti12, patel13}, and (2) the depth and resolution of new \hst\ imaging enables analyses of galaxy structure that current ground-based observations cannot support. This is especially true in the near-infrared, which most directly probes galaxies' stellar mass distributions. By using a new multidimensional approach that holistically examines galaxies in their natural parameter space---spanning color, size, structure, environmental density, and redshift---our analysis provides a new look at both rapid and long-term environmental influences, yielding a ``4D'' view of the galaxy population at unexplored masses and spatial resolutions. We organize our discussion around the central question posed above: how many of the observed differences in cluster/field populations reflect phenomena {\it driven} by clusters versus those {\it traced} by them? Ultimately, our results suggest that, while a cluster-specific process similar to ram pressure stripping is indeed operational now, $\sim20\%$ of present-day passive cluster galaxies with $\logm<10$ must have been ``built into'' the cluster population at very early times. These findings support a scenario in which clusters mark places where galaxy evolution has been accelerated compared to---but not radically divergent from---the cosmic mean, but are now also in a phase of transforming mainly low-mass systems via environmentally specific phenomena. We proceed as follows: in Section \ref{sec:data}, we describe the observations and measurements upon which our analysis is based. In Section \ref{sec:result1}, we explore the canonical size--mass relations of our sample and use these to identify important spatiotemporal trends in the data. In Section \ref{sec:result2}, we adopt a new framework to reinterpret galaxy structural parameters holistically across all environments, performing a multidimensional analysis similar in spirit to the approach that led to the discovery of the fundamental plane \citep{djorgovski87, dressler87} and the more-fundamental plane \citep{bolton07, auger10} of early-type galaxies. We discuss our results in Section \ref{sec:discussion} and summarize in Section \ref{sec:summary}. Details of various parts of our analysis are also provided in Appendices. Magnitudes are quoted in the AB system \citep{oke83, fukugita96}. We assume $\Omega_m=0.3$, $\Omega_\Lambda=0.7$, $H_0=70\,\kms\, {\rm Mpc}^{-1}$, and a \citet{chabrier03} initial mass function (IMF). The catalog for galaxy structural parameters are made available as an electronic table associated with this paper and through the GLASS website.\footnote{\url{http://glass.astro.ucla.edu}}

\begin{enumerate} \item Our target clusters are {\it currently} bright X-ray sources, confirming the presence of dense intra-cluster gas \citep{mantz10}. Especially in the core regions probed by the \hst\ observations, drag from this hot atmosphere can effectively strip gas from infalling galaxies, and also stifle their accretion of new gas for future star formation, the definitions of ram-pressure stripping and starvation. \item Large-size, low-mass systems are most amenable to (ram pressure) stripping as their internal gas supplies are most loosely bound \citep{treu03}. They would also run out of fuel for star formation relatively quickly if starved of external fuel supplies given the generally higher specific star formation rates of their low-mass star-forming galaxy progenitors \citep[e.g.,][]{salim07, whitaker14}. \item The low-mass end of the passive cluster mass function grows over the epochs probed (Figures \ref{fig:smf}, bottom-right), consistent with a scenario where star-forming galaxies are continually being transformed. This is also consistent with previous studies that a single epoch is disfavored for the formation of low-mass passive cluster galaxies \citep[e.g.,][]{roediger11,toloba14}. \item The relatively low S\'ersic indices of the low-mass passive cluster galaxies (Figure \ref{fig:sersic}) point to a ``gentle'' mechanism; one that does not destroy disks/rearrange galaxies' stellar components, which is consistent with starvation, ram pressure stripping, and galaxy harassment \citep{bialas15}. The observed S\'ersic indices of low-mass passive galaxies are slightly higher than those of star-forming ones with similar stellar mass, but this difference is explained by disk fading, where cessation of star formation in the disk leads to more concentrated light profiles (and thus higher S\'ersic indices; \citealt{lackner13}). \end{enumerate} Given the direct evidence for stripping in local analogues \citep[e.g.,][]{cayatte90,abramson11} and at intermediate redshift \citep{vulcani15}, it seems likely that at least this mechanism is {\it currently} operative. {\it However}, to attribute the presence of the smallest-size low-mass passive galaxies---which are {\it also} the oldest (Figure \ref{fig:lowQG})---to the same mechanism(s), one must take into account the evolution of the cluster itself. \subsection{Small-$r_{e}$ Low-mass Passive Cluster Galaxies} \label{ssec:slmass} At $z\sim3$, when many of these $\logm<9.8$ systems seem last to have been in the star-forming population, our target clusters would have been much less massive, and were therefore home to much more tenuous, cooler intra-cluster media. Based on calculations by \citet{trenti08},\footnote{$\Lambda$CDM cosmology based on Wilkinson Microwave Anisotropy Probe (WMAP) year 1 results \citep{spergel03}.} we estimate that our clusters---systems with $\log M_{\rm halo}/\Msun\sim15$ at $z\sim0.5$---had progenitors with $\log M_{\rm halo}/\Msun\sim14$ at $z\sim3$ (also consistent with \citealt{evrard02}). Hence, the question becomes whether or not the environments at those epochs---when the global density and neutral gas fraction was higher---could have supported ram pressure stripping/significantly cut galaxies off from their fuel supplies. If they could, these channels could provide a unified explanation for all low-mass passive cluster galaxies. If they could {\it not}, another channel must have been open. \begin{figure} \begin{center} \includegraphics[width=8.2cm,bb=0 0 288 288]{fig13.pdf} \caption{Extrapolation of the current quenching rate to high lookback times as a function of stellar mass. Each color corresponds to stellar mass bin listed in legend. The derived linear fit slope, $dF_{\rm red}/dt$, for each stellar mass bin is shown near the fitting line in the same color. One stellar mass bin, $9.25<\logm<9.75$ (green), has a negative slope, $dF_{\rm red}/dt=-0.004$, because of a weak statistics, and the value is not shown. } \label{fig:frac_evl} \end{center} \end{figure} To further constrain the strength of the above channels and their ability to produce the smallest low-mass passive cluster galaxies, we can assume all of the evolution of the passive fraction shown in Figure \ref{fig:smf}, bottom-right, is due to the same mechanism(s) and project the effects back in time. We do so by extrapolating a simple linear regression and show the results in Figure \ref{fig:frac_evl}. This is an extreme model---we might expect the recent evolution in passive fractions to be more rapid than its past evolution due to the cluster and cosmic evolutionary effects mentioned above, which should reduce the efficiency of, e.g., stripping, with lookback time. However, it provides something like an upper limit to the quenching that could be driven by such mechanisms, which is what we seek. From this exercise, we obtain $dF_{\rm red}/dt\sim0.1\, {\rm Gyr^{-1}}$ at $\logm<9$. By extrapolating the slope, we find that the passive fraction at these masses would be zero at $2\lesssim z\lesssim4$, depending on the bin. Hence, it could be that {\it all} cluster (core) galaxies with $\logm<9$ were transformed due to gas removal, or whatever other process is active {\it now} in clusters. However, this idea does not hold for even slightly more massive passive cluster galaxies; i.e., those with $9.25\lesssim\logm\lesssim10.25$---there are simply too many of these to have all arisen through the same channel. Our toy calculation suggests that perhaps $30\%$ of these systems were already in place at $z\sim3$, just 2~Gyr after the Big Bang. For these galaxies---and presumably those with yet lower masses, assuming a more-realistic nonlinear $dF_{\rm red}/dt$---other explanations must be sought. We explore one possibility below. \subsubsection{Evidence for Accelerated Evolution for Low-mass Dense Cluster Galaxies} Clues for a formation scenario come from the low-mass passive {\it field} and the high-mass cluster populations. From Figure \ref{fig:uvj}, we see that, at $\logm<9$, there exist passive galaxies in clusters that are systematically redder than those in the field. This implies that the cluster galaxies {\it reached their final mass before their field counterparts}. Turning to the $\logm>9.8$ passive cluster population, Figure \ref{fig:2slopes} shows the size--mass relation of these objects to lie remarkably close to a line of constant stellar surface density, $\Sigma_*=\Mstel/r_{e}^{2}$. This points to common formation time for these systems \citep[][]{franx08, vandenbosch08, stringer14, lilly16, whitaker16, abramson16b}. Given their uniformly ancient stellar populations (based on their colors), the implication is that such high-mass passive cluster galaxies are monolithically old and do not descend from field galaxies transformed over long stretches of time. Assuming this is the case, we can obtain a rough idea of how many low-mass objects formed similarly by calculating the fraction of them that have surface densities similar to the high-mass systems. Given that the 2-$\sigma\approx0.5\,{\rm dex}$ intrinsic spread in sizes at fixed mass corresponds to a factor of 100 difference in surface densities for galaxies at the top and bottom of the size--mass relation, we should expect some $\logm=8$--9 passive objects to have densities comparable to their $\logm=10$--11 counterparts. Indeed, we find $\approx18\%$ of galaxies with $\logm<9.8$ to lie above $\Sigma_*\approx10^{8.3}\, \Msun\,\kpc^{-2}$---the 1-$\sigma$ lower bound to the high-mass sample's surface densities. This fraction is close to that independently obtained by extrapolating $dF_{\rm red}/dt$ above. Combined with the fact that these systems will, by definition, be the smallest low-mass galaxies---and therefore also the {\it reddest} (Figure \ref{fig:lowQG})---this finding strengthens the conclusion that such systems were in place long ago, having formed alongside their massive counterparts. This is consistent with \citet{toloba15}, who find low-mass galaxies with higher stellar velocity dispersions at fixed mass to have lower $r_{e}$ and lie closer to the cluster core in Virgo, and therefore be older than larger-$r_{e}$ systems. Notably, the same density calculation reveals only $6\%$ of low-mass {\it field} passive galaxies to have densities consistent with high-mass passive objects. If we assume, following \citet{geha12}, that all of these dense systems are in fact stripped satellites---i.e., they do not truly arise from the same processes generating high-mass passive galaxies---this estimate can be taken as our measurement uncertainty. As such, combined with the $dF_{\rm red}/dt$ results, we can state the following: regarding the formation of most high-mass and 10\%--30\% of $\logm<9.8$ passive cluster galaxies, clusters mark regions of space where evolution was {\it accelerated} due to a population's residence in a common overdensity. That is, perhaps {\it all} sufficiently dense cluster galaxies---regardless of stellar mass---are consistent with having arisen through a common, prompt, formation channel. This scenario---consistent with that of \citet{dressler80}, \citet{abramson16}, and \citet{kelson16}---is fundamentally different from the environmental quenching of infalling field systems that accounts for the rest of the (mainly low-mass) cluster passive galaxy population at $z\sim0.5$, and presumably today, suggesting a dual (or bimodal) formation scenario for passive cluster galaxies \citep[see also][]{poggianti01,poggianti06}. A consequence of any ``accelerated'' growth is that passive cluster objects would lock-in the smaller sizes of {\it all} star-forming galaxies {\it everywhere} at early epochs, appearing naturally at the bottom of the size distribution for their ``final mass'' and exhibiting globally higher S\'ersic indices, precisely as seen in Figures \ref{fig:sersic} and \ref{fig:lowQG}. This effect would also naturally lead to the color--size anticorrelation revealed by our holistic fit ($\beta_{UV}=-0.29\pm0.01$; Figure \ref{fig:NFIT}), which is therefore actually a reflection of the well-known redshift--size anticorrelation \citep[e.g.,][]{newman12, newman14, vanderwel14, morishita15} that our analysis also reveals ($\beta_{z} = -1.06\pm0.14$). The evidence of the accelerated grown in dense environment is also observed in a higher redshift cluster \citep[e.g.,][though limited to massive galaxies]{papovich12, bassett13}. In sum, our results point to an identifiable ``native'' population of galaxies at all masses that matured rapidly at early times {\it because} it was situated in a region of space that was also collapsing quickly. This is the mechanism that reaches to large distances, causing the population differences between clusters and the field to extend to many virial radii (\citealt{lewis02}; \citealt{treu03}; \citealt{dressler13}). To this is added a frosting of new galaxies at late times driven by something akin to ram pressure stripping or starvation (certainly something gentle), which is especially active at small clusto-centric radii and later epochs. \subsection{Better Tests than Scaling Relations} The HFF provides something close to the best possible imaging data acquirable for objects in the distant universe. As such, it is unclear what more-detailed studies of the size--mass relation will uncover in terms of offsets between mean sizes of populations as a function of environment that cannot be gleaned already. Our approach in Section \ref{sec:result2} points to the fundamental limitations of ever more sophisticated analysis of these kinds of galaxy-integrated, photometric metrics. Instead, our analysis suggests that, if one seeks better knowledge of the detailed physical mechanisms transforming galaxies in clusters (or outside of them), different kinds of data are required. Principally, the addition of {\it spectroscopic} data, and probably in a spatially resolved sense; i.e., deep and wide IFU surveys investigating the star formation and kinematic properties that may differentiate galaxies as a function of mass and environment. For example, using light-weighted stellar ages, $\logm>9.8$ passive galaxies observed in local clusters and the field have been seen to show a similar trend to our $\logm<9.8$ systems, such that the largest-$r_{e}$ galaxies are the youngest \citep[e.g.,][]{valentinuzzi10, poggianti13}. At $z\sim0.5$, we see no significant color--size trend for equal-mass galaxies in this regime, but $U-V$ colors are limited as an age indicator: any passive galaxies with ages $>2$~Gyr would be uniformly red in this index. Hence, deep optical spectroscopy of our sample is needed to determine whether any real evolution at these masses takes place in the $\sim5$ intervening Gyr. Also, the observed scatter in $r_{e}$ could be due to radial stellar migration and not age. Induced by stellar feedback, such migration can cause $r_{e}$ to fluctuate by $\sim2\times$ in just $\sim100$~Myr \citep[e.g.,][]{elbadry16}. We see a clear correlation between color and galaxy size, which suggests such rapid effects are not the principal source of scatter in $r_{e}(\Mstel)$ in the low-mass population, but resolved spectroscopy would constitute a much more stringent physical constraint. The GLASS \citep{jones15, treu15, vulcani15, vulcani16a, vulcani16b}, GASP (Poggianti et al. in preparation), {\sc Sauron} \citep{davies01}, {\sc Atlas3D} \citep{cappellari11}, {\sc 3D-HST} \citep{brammer12,nelson15}, CALIFA \citep{sanchez12}, {\sc MaNGA} \citep{bundy15}, SAMI \citep{allen15}, and KROSS \citep{magdis16} have already started these investigations, and highly resolved analyses (mainly at low-$z$) show both signs of accelerated evolution driven by clusters \citep{mcdermid15}, and stripping \citep{conselice01,conselice03a,nipoti03,janz16}. By studying the sites of star formation (\citealt{wang15,wang16}) and kinematics \citep[e.g., KLASS;][]{mason16} through resolved gas maps at high(er)-$z$ using the next generation of ground- and space-based instruments, we may be able to map the evolving importance of these effects at levels of detail currently available only in the local universe. We derived photometric redshifts, stellar masses, and structural parameters for $>3900$ cluster and field galaxies at $0.2\leq z\leq0.7$ from the HFF and GLASS programs, complete to $\logm=7.8$---an unexplored regime at these redshifts. Using this homogeneous sample: \begin{enumerate} \item We studied the size--mass relations of four ``canonical'' populations---cluster/field, passive/star-forming galaxies---fitting each subsample with a single slope. Though the populations reside in environments of maximally different density, $\sim$ a factor of 1000, we find the relations to be identical within their measurement uncertainties (Figure \ref{fig:slope}). This holds even at the lowest masses where cluster-specific effects would be expected to have the most significant impact on galaxy structure. \item A multivariate analysis---wherein all galaxy classifications are removed and sizes are fit as a function of stellar mass, S\'ersic index, color, redshift {\it and} environment---is consistent with the above results, and quantitatively reveals local density to induce but a $7\%\pm3\%$ reduction in size ($95\%$ confidence) when controlling for these other factors (Figure \ref{fig:NFIT}). Immediate environment therefore appears to have a tiny effect on galaxy size, while stellar mass and color correlate most strongly. \item We studied the trends in $(U-V)$ color in the low-mass passive cluster population ($\logm<9.8$) as a function of offset from the best-fit slope of {\it star-forming} galaxies at the same redshift. We find that smaller-size galaxies are also redder. \item The {\it largest}-size low-mass passive cluster galaxies---which are also the bluest---have sizes and S\'ersic indices similar to those of contemporaneous star-forming galaxies (Figures~\ref{fig:sersic}, \ref{fig:2slopes}). This fact suggests that they are recently acquired systems that have been ``quenched'' by a cluster-specific process that terminates star formation in a non-violent manner/preserves the structure of star-forming galaxies. Given that our clusters all harbor dense intra-cluster gas, the most likely candidate is ram pressure stripping or starvation. \item This explanation holds for the {\it smallest} low-mass passive cluster galaxies only if the progenitors of our clusters were capable of hosting a hot intra-cluster medium at $z\gtrsim3$. If not, the consistent stellar surface densities and colors of these objects with those of their uniformly ancient, more massive ($\logm>9.8$) peers suggest that $10\%$--$30\%$ of these galaxies are ``native,'' having had their evolution accelerated---not terminated---by their presence in a large overdensity at birth. \end{enumerate} Our conclusion is therefore not that environment has no impact on galaxy evolution, but rather that it {\it encodes} the fact that most high-mass and $\sim18\%$ of low-mass passive galaxies found in dense regions at late times had common (accelerated) evolutionary trajectories, with late-time effects playing some role, but only dominant at low masses.

1607.00384

1607

1607.05797_arXiv.txt

Using all the RXTE archival data of Sco X-1 and GX 5-1, which amount to about 1.6 mega seconds in total, we searched for possible occultation events caused by Oort Cloud Objects. The detection efficiency of our searching approach was studied with simulation. Our search is sensitive to object size of about 300 m in the inner Oort Cloud, taking 4000 AU as a representative distance, and of 900 m in the outer Oort Cloud, taking 36000 AU as the representative distance. No occultation events were found in the 1.6 Ms data. We derived upper limits to the number of Oort Cloud Objects, which are about three orders of magnitude higher than the highest theoretical estimates in the literature for the inner Oort Cloud, and about six orders higher for the outer Oort Cloud. Although these upper limits are not constraining enough, they are the first obtained observationally, without making any model assumptions about comet injection. They also provide guidance to such serendipitous occultation event search in the future.

Existence of the Oort Cloud, which was proposed 66 years ago \citep{oort50}, has not yet been directly confirmed, except for the discovery of Sedna \citep{brown04}, 2012 VP$_{113}$ \citep{trujillo14} and three other objects \citep{chen13,brasser15}, which are considered to be inner Oort Coloud Objects. The population properties of Oort Cloud Objects are not yet clear either, such as their total number and size distribution. The knowledge of these properties is important to our understanding of the origin and evolution of the Oort Cloud. Oort Cloud Objects are usually thought to form in the outer planetary region and to be ejected to a distance far away from the sun by giant outer planets (for brief reviews, see, e.g., \citet{dones04a,rickman14,dones15}). Galactic tides, star/giant-molecular-cloud encounters, and planetary perturbations all affect the formation and evolution of the Oort Cloud as well as the injection of Oort Cloud Objects into the inner solar system. With a proper injection model and observed flux of long-period comets, one can in principle infer the number of objects in the Oort Cloud. This has been being the effort made by many authors with various physical considerations. In \citet{duncan87} it was found, by direct numerical integration, that the Oort Cloud has a sharp inner edge at about 3000 AU and the number density of Oort Cloud Objects falls steeply outwards, roughly as $r^{-3.5}$. The Oort Cloud is therefore centrally condensed, with roughly 5 times more objects in the inner Oort Cloud ($a < 20,000$ AU) than in the classical outer Oort Cloud. In a later study, \citet{weissman96} estimated that there are $10^{12}$ objects of diameter larger than 2.3 km in the outer Oort Cloud ($a > 20,000$ AU). The inner Oort Cloud was originally proposed by \citet{hills81}. \citet{kaib09} demonstrated that a significant fraction of long-period comets in fact come from the inner Oort Cloud before their migration into the outer Oort Cloud due to planetary perturbation. It was estimated in \citet{kaib09} that there could be $\sim 10^{12}$ objects in the inner Oort Cloud (3000 AU $<a<$ 20,000 AU). Some estimates yielding smaller numbers of objects in the outer Oort Cloud, about $2\times 10^{11}$ -- $5\times 10^{11}$, were reported in \citet{heisler90,dones04b,brasser13}. Direct observation of objects in the Oort Cloud region is extremely difficult. Their existence, however, in addition to being inferred from observations of long-period comets, may also be explored by occultation events that they cause to distant background stars. Such an approach, looking for serendipitous occultation events from a vast amount of data, has been being employed to study the size distribution of small Kuiper Belt Objects (KBOs) down to sub-kilometer size using optical data \citep{roques06,bickerton08,schlichting12,zhang13,liu15} and even to decameter size using X-ray observations \citep{chang13}. It has also been applied to estimating upper limits to the number of small objects at distances between 100 and 1000 AU with data accumulated by the TAOS project \citep{wang09}. In this paper we report the result of our effort to extend the work of serendipitous KBO occultation search in X-rays to the regime relevant to possible occultation events caused by Oort Cloud Objects. Our conclusion is that in all the RXTE archival data of Sco X-1 and GX 5-1, 1.6 mega seconds in total, no such events were found. The derived upper limits to the number of inner and outer Oort Cloud Objects are about three orders of magnitude higher than the highest theoretical estimates in the literature for the inner, and about six orders higher for the outer, respectively. Although these upper limits are not constraining enough, they are the first obtained observationally, without making any model assumptions about comet injection. They also provide guidance to observations in the future. At the end of this paper, we discuss the possibility of using ASTROSAT \citep{singh14}, Athena (a future ESA X-ray space mission \citep{ayre15}) and LOFT (a proposed mission \citep{zane14}) to conduct such a search to study the size distribution of Oort Cloud Objects and KBOs.

1607.05797

1607

1607.05274_arXiv.txt

Using high-resolution simulations with explicit treatment of stellar feedback physics based on the FIRE (Feedback in Realistic Environments) project, we study how galaxy formation and the interstellar medium (ISM) are affected by magnetic fields, anisotropic Spitzer-Braginskii conduction and viscosity, and sub-grid metal diffusion from unresolved turbulence. We consider controlled simulations of isolated (non-cosmological) galaxies but also a limited set of cosmological ``zoom-in'' simulations. Although simulations have shown significant effects from these physics with weak or absent stellar feedback, the effects are much weaker than those of stellar feedback when the latter is modeled explicitly. The additional physics have no systematic effect on galactic star formation rates (SFRs) . In contrast, removing stellar feedback leads to SFRs being over-predicted by factors of $\sim 10 -100$. Without feedback, neither galactic winds nor volume filling hot-phase gas exist, and discs tend to runaway collapse to ultra-thin scale-heights with unphysically dense clumps congregating at the galactic center. With stellar feedback, a multi-phase, turbulent medium with galactic fountains and winds is established. At currently achievable resolutions and for the investigated halo mass range $10^{10}-10^{13} M_{\odot}$, the additional physics investigated here (MHD, conduction, viscosity, metal diffusion) have only weak ($\sim10\%$-level) effects on regulating SFR and altering the balance of phases, outflows, or the energy in ISM turbulence, consistent with simple equipartition arguments. We conclude that galactic star formation and the ISM are primarily governed by a combination of turbulence, gravitational instabilities, and feedback. We add the caveat that AGN feedback is not included in the present work.

\label{S:intro} Feedback from stars is essential to galaxy evolution. In isolated galaxy simulations without strong stellar feedback, giant molecular clouds (GMCs) experience runaway collapse, resulting in star formation rates (SFRs) orders-of-magnitude higher than observed \citep{2010MNRAS.409.1088B,2011MNRAS.413.2935D,2011IAUS..270..235H,2011ApJ...740...74K,2011ApJ...730...11T,2011MNRAS.417..950H}. This is in direct contradiction with the observed Kennicutt-Schmidt (KS) relation, which shows that the gas consumption time of a galaxy is roughly $\sim 50 -100$ dynamical times \citep{1998ApJ...498..541K,1974ApJ...192L.149Z,1997ApJ...476..166W,1999ARA&A..37..311E,2009ApJS..181..321E}. Cosmological simulations without strong feedback face a similar challenge. The efficiency of cooling causes runaway collapse of gas to high densities within a dynamical time, ultimately forming far too many stars compared to observations (\citealt{1996ApJS..105...19K,1999MNRAS.310.1087S,2000MNRAS.319..168C,2003MNRAS.339..289S,2009MNRAS.396.2332K}, and references therein). \vspace{0.1cm} Recent years have seen great progress in modeling feedback on galaxy scales \citep{2000ApJ...545..728T,2007MNRAS.374.1479G,2009ApJ...695..292C,2012MNRAS.423.2374U,2011MNRAS.417..950H, 2012MNRAS.421.3488H,2012MNRAS.421.3522H,2015arXiv150900853A,2015arXiv151005644H}. In \cite{2011MNRAS.417..950H, 2012MNRAS.421.3488H}, a detailed feedback model including radiation pressure, stellar winds, supernovae and photo-heating was developed and applied to simulations of isolated galaxies. They showed that stellar feedback is sufficient to maintain a self-regulated multi-phase interstellar medium (ISM), with global structure in good agreement with the observations. GMCs survive several dynamical times and only turn a few per cent of their mass into stars, and the galaxy-averaged SFR agrees well with the observed Kennicutt-Schmidt (KS) law. These models were extended with numerical improvements and additional cooling physics, and then applied to cosmological ``zoom in'' simulations in the FIRE (Feedback In Realistic Environments) project\footnote{Project web site: http://fire.northwestern.edu.}. A series of papers, using the identical code and simulation set have demonstrated that these feedback physics successfully reproduce a wide range of observations, including star formation histories of galaxies \citep{2014MNRAS.445..581H}, time variability of star formation \citep[][]{2015arXiv151003869S}, galactic winds \citep[][]{2015MNRAS.454.2691M}, HI content of galaxy halos (\citealt{2015MNRAS.449..987F,2016arXiv160107188F}; Hafen et al., in prep.), and galaxy metallicities \citep[]{2015arXiv150402097M}. Other groups (e.g. \citealt{2013MNRAS.428..129S}, who implemented energy injection from SNe and an approximate treatment of UV radiation pressure, and \citealt[e.g.,][]{ 2015arXiv150900853A}, who included momentum injection from SNe, radiation pressure and stellar winds) have also found that stellar feedback can regulate galaxy SFRs and lead to realistic disc morphologies. However, several potentially important physical processes have not been included in most previous galaxy formation simulations. Magnetic fields have long been suspected to play a role in galaxy evolution because the magnetic pressure reaches equipartition with the thermal and turbulent pressures \citep{1996ARA&A..34..155B,2009ASTRA...5...43B}. Isolated galaxy simulations with magnetic fields -- but using more simplified models for stellar feedback -- have been studied in various contexts and suggest that magnetic fields can provide extra support in dense clouds, thus slowing down star formation \citep{2009ApJ...696...96W,2012MNRAS.422.2152B,2013MNRAS.432..176P}. Turbulent box simulations \citep{2005ApJ...629..849P,2007ApJ...663..183P} also suggest that MRI-driven (magnetorotational instability) turbulence can suppress star formation at large radii in spiral galaxies. In particular, \cite{2015ApJ...815...67K} explicitly demonstrate such suppression from magnetic fields in a simulation of a turbulent box that includes momentum feedback from SNe. Magnetic fields can also be important because of their effects on fluid mixing instabilities, including the Rayleigh-Taylor (RT) and Kelvin-Helmholtz (KH) instabilities \citep{1995ApJ...453..332J, 2015MNRAS.449....2M,2016arXiv160805416A}. These instabilities can potentially affect galaxy evolution through processes including the evolution of supernovae (SN) remnants \citep{1996ApJ...465..800J,1996ApJ...472..245J,1999ApJ...511..774J,2000ApJ...534..915T, 2015ApJ...815...67K}. Another potentially important effect is viscosity, which has been more extensively studied in simulations of galaxy clusters. It has been suggested that viscosity can affect the turbulent motion of the intracluster medium (ICM) or circum-galactic medium (CGM) and affect the KH stability of various structures in the ICM \citep[][]{2007PhR...443....1M}. It has been shown in particular that viscosity may be important for the dynamics of bubbles in the ICM inflated by active galactic nucleus (AGN) feedback or bursts of SNe activity \citep{2005MNRAS.357..242R, 2006MNRAS.371.1025S}. Thermal conduction, which in the presence of magnetic fields is highly anisotropic, affects the stability of plasmas at both galactic and cluster scales \citep{2009ApJ...699..348S,2010ApJ...720..652S, 2012MNRAS.422..704P,2016arXiv160805416A, 2012ApJ...747...86C} and the survival and mixing of multi-phase fluids. Combined with the effect of magnetic fields, conduction may be critical to determine the survival of cool clouds in galactic winds. Turbulent metal diffusion due to small-scale (un-resolvable) eddies may also have important effects. It has been suggested, for example, that unresolved turbulence in galaxy simulations may be important to effectively ``diffuse'' metals in the ISM and intergalactic medium \citep[IGM; e.g.,][]{2010MNRAS.407.1581S}, leading non-linearly to different cooling physics at halo centers and within the dense ISM. While most previous studies considered these physics in isolation, their effects and relative importance may be quite different in a realistic multi-phase ISM shaped by strong stellar feedback processes. Another challenge is that conduction and viscosity in magnetized plasmas are inherently anisotropic. Properly treating this anisotropy requires MHD simulations and is numerically non-trivial; consequently, most previous studies on galactic scales have considered only isotropic conduction and viscosity. However, studies which correctly treat the anisotropy have shown that this anisotropy can produce orders-of-magnitude differences and, in some cases, qualitatively different behavior \citep{2009ApJ...704.1309D,2015ApJ...798...90Z,2009ApJ...699..348S,2010ApJ...720..652S,2012ApJ...747...86C} In this paper, we study the effects of these different microphysics in the presence of explicit models for stellar feedback. While the simulations analyzed here implement the same stellar feedback physics from the FIRE cosmological simulations, we focus primarily on non-cosmological simulations of isolated galaxies, because this allows us to achieve higher spatial and mass resolution, and to have well-controlled experiments with identical galaxy initial conditions. In cosmological runs, on the other hand, the inherently chaotic nature of the problem makes detailed one-to-one comparison of simulations with varied physics more complicated; we do, however, include a limited subset of these experiments. We also make use of a new, more accurate hydrodynamic solver, needed to properly treat MHD and anisotropic diffusion. Overall, we find that at the resolutions currently achievable in isolated galaxy and cosmological simulations, MHD, anisotropic conduction and viscosity, and sub-grid turbulent metal diffusion play a relatively minor role in the regulation of star formation and of the phases and energetics of the ISM \emph{when the dominant effects of stellar feedback are simultaneously modeled}. We caution, however, that despite this result, some of these effects likely have some important and observationally interesting consequences on finer scales, such as for the survival of cool clouds in galactic winds \citep[e.g.,][]{2015MNRAS.449....2M, 2016arXiv160805416A,2016ApJ...822...31B}, and stellar abundance distribution patterns within star clusters or small galaxies. It is also possible that some important effects would only reveal themselves in simulations of much higher resolution than currently possible for galaxy simulations. Furthermore, the interaction of physical processes not included in our simulations with, e.g., magnetic fields is likely to prove important. This is the case in particular for the transport of cosmic rays, which a number of recent studies indicate may be an important form of feedback for galaxy evolution \citep[e.g.,][]{2012MNRAS.423.2374U, 2013ApJ...777L..16B, 2014ApJ...797L..18S, 2016arXiv160204856R,2016arXiv160407399P,2016MNRAS.462.2603P}. The remainder of this paper is organized as follows: in \sref{S:methods}, we describe the initial conditions and the baryonic physics model of our default model. In \sref{s:add_ph}, we summarize the additional physics studied in this paper. In \sref{S:results}, we analyze the effects on the star formation histories, morphologies, phase structures, magnetic and turbulent energies, and outflows of our simulated galaxies. We discuss the reason why the fluid microphysics have minor effects in \sref{s:discussion} and conclude in \sref{S:conclusions}. \vspace{-0.5cm}

\label{S:conclusions} We use simulations with parsec-scale resolution, explicit treatments of stellar feedback identical to those used in the FIRE project, magnetic fields, anisotropic Spitzer-Braginskii conduction and viscosity, and sub-grid turbulent metal diffusion to study how these affect galaxy-scale star formation, the phase structure of the ISM, and the generation of galactic outflows. We consider both isolated (non-cosmological) simulations of a range of galaxy types and fully cosmological zoom-in simulations of a Milky Way-mass halo and a dwarf halo. In all cases, we find the following: \begin{itemize} \item{Stellar feedback plays the dominant role in regulating the SFR. % We find that magnetic fields and additional microphysical diffusion processes change the SFR (and therefore the KS law) by small amount comparing to the effect from stellar feedback in the investigated ma. This is consistent with the models advocated in the aforementioned papers (see the references in \sref{s:discussion.bfield}), in which the SFR and star formation scaling relations are set by self-regulation via feedback, which drives super-sonic turbulence and balances the disc against gravity.} \item{The ISM phase structure and galactic winds is also primarily established by stellar feedback. % Stellar feedback also serves as an extra source of turbulent energy, boosting the rms turbulent velocity by a factor of 2-3. Perhaps surprisingly, however, neither MHD nor the additional diffusion microphysics appear to produce larger than $\sim 10\%$-level systematic effects on these quantities. In fact, in some earlier experiments where we artificially increased the viscosity coefficient $\eta$ by a factor of 100, there were still weak systematic effects. It appears that because the turbulence is super-Alfv{\'e}nic on the scales most important for fragmentation, ISM phase structure and outflow generation (of order the disc scale height), these effects are sub-dominant. A more detailed discussion of why such small effects are seen is provided in \sref{s:discussion}.} \item{The magnetic field energies saturate at $\sim 10\%$ of the turbulent kinetic energies on of order the galactic scale height (\fref{fig:turbulent}). The ratio is smaller still if we include the kinetic energy of small-scale galactic fountains in the ``turbulence'' budget. This is consistent with both observations \citep{,1996ARA&A..34..155B,2002RvMP...74..775W,2008RPPh...71d6901K,2008Natur.454..302B,2008ApJ...676...70K,2012ApJ...757...14J,2012ApJ...761L..11J} and other simulations \citep{2012MNRAS.422.2152B,2013MNRAS.432..176P,2009ApJ...696...96W,2010A&A...523A..72D,2010ApJ...716.1438K,2011MNRAS.415.3189K}. This result partially explains why the magnetic field's effects are sub-dominant on the large scales of order the disc scale height (the scales containing most of the turbulent energy).} \item{A systemic increase of stellar mass and cold gas is observed in CosmoDwarf run with all fluid microphysics included. This may result from conduction dissipating part of the SNe energy making it more difficult to wipe out cold clumps. Our cosmoMW run shows a similar enhancement in late-time cooling from the CGM with all microphysics present. A more detailed discussion of this is provided in \sref{s:discussion}.} \end{itemize} It appears that, at least on galactic scales, in the presence of explicit models for multi-mechanism stellar feedback as well as self-gravity, magnetic fields and additional diffusion microphysics (such as conduction, viscosity, and sub-grid turbulent metal diffusion) are subdominant in the star formation and galaxy formation process at currently achievable resolutions. This general result appears to contradict some earlier claims in the literature. However, to our knowledge, these prior studies have not focused on the combination of large galactic scales (yet with high enough resolution to resolve vertical disc scale heights and the phase structure in discs) and fully explicit models for stellar feedback. For example, it is relatively ``easy'' for magnetic fields to have a large fractional effect in simulations with either no or weak stellar feedback or stellar feedback modeled only in a ``sub-grid'' fashion (so it e.g.\ does not locally alter the gas dynamics but only ejects gas in outflows or adds an effective pressure term). However, the claimed effects in these cases are typically order-unity \citep{2005ApJ...629..849P,2007ApJ...663..183P,2009ApJ...696...96W,2012MNRAS.422.2152B,2013MNRAS.432..176P} and thus still orders of magnitude less than the factor $\sim 100-1000$ changes in the properties we study here that occur when the full model for stellar feedback is introduced. Altogether, our results support the emerging picture wherein galaxy-scale ($\gtrsim 10-100\,$pc) star formation, ISM structure, and outflows are determined primarily by a competition among super-sonic (and super-Alfv{\'e}nic) turbulence, stellar feedback, and self-gravity. The microphysics we study here may certainly be important on smaller scales (e.g. for regulating the structure of turbulent cores as they collapse to form stars) or in the more diffuse CGM and IGM (e.g. the outskirts of galaxy clusters). However, they do not, to leading order, significantly alter the dynamics on the scales we study here. We also caution that certain unresolved processes (e.g. conduction altering mixing and cooling in single SNe blastwaves or cool cloud ``shredding'' in the circumgalactic medium) may have large non-linear effects on the efficiency of feedback or cooling, and these cannot be captured in our simulations. We see tentative evidence of this in our fully cosmological MW-mass simulation, which shows enhanced late-time cooling and a larger gas disc with conduction, viscosity and sub-grid metal diffusion active. Although the magnetic field has little effect on the properties analyzed in our current study, it might for instance provide important pressure support in the violent tidal compression that occurs in galaxy mergers, which could possibly affect the properties of the star clusters formed in merger-induced starbursts. Besides stellar feedback and fluid microphysics, AGN feedback may be an important determinant of galaxies' physical properties, especially for massive galaxies. Moreover, cosmic rays may significantly affecting galaxy evolution, and properly treating cosmic ray transport requires an accurate determination of the magnetic field. Detailed investigations of these processes and their interaction with fluid microphysics in the context of simulations with explicit stellar feedback will be presented in future work. \vspace{-0.7cm}

1607.05274

1607

1607.00498_arXiv.txt

A computed line list for hydrogen peroxide, H$_2{}^{16}$O$_2$, applicable to temperatures up to $T=1250$~K is presented. A semi-empirical high accuracy potential energy surface is constructed and used with an {\it ab initio} dipole moment surface as input TROVE to compute 7.5 million rotational-vibrational states and around 20 billion transitions with associated Einstein-$A$ coefficients for rotational excitations up to $J=85$. The resulting APTY line list is complete for wavenumbers below 6~000 cm$^{-1}$ ($\lambda < 1.67$~$\mu$m) and temperatures up to 1250~K. Room-temperature spectra are compared with laboratory measurements and data currently available in the HITRAN database and literature. Our rms with line positions from the literature is 0.152 \cm\ and our absolute intensities agree better than 10\%. The full line list is available from the CDS database as well as at \url{www.exomol.com}.

Terrestrial hydrogen peroxide exists as a trace molecule in the Earth's atmosphere and contributes to the atmospheres oxidising budget as well as ozone production and water chemistry \citep{74DaXXXX.H2O2,91ChJoTr.H2O2,13AlAbBe.H2O2,14ZiKrxx.H2O2} and its concentration is now being routinely observed \citep{13AlAbBe.H2O2}. Astrophysically there have been multiple detections of \hhoo\ in the atmosphere of Mars \citep{04ClSaMo.H2O2,04EnBeGr.H2O2,12EnGrLe.H2O2,15AoGiKa.H2O2} with seasonal variation, possibly formed by triboelectricity in dust devils and dust storms \citep{12EnGrLe.H2O2} and may well act as an agent in the oxidization of the Martian surface. Hydrogen peroxide has also been detected in the atmosphere of Europa \citep{13HaBrXX.H2O2} in the 3.5$~\mu m$ region. The first detection of interstellar \hhoo\ was made by \citet{11BePaLi.H2O2} and is believed to play an important role in astrophysical water chemistry similar to that on Earth. \citet{12DuPaBe.H2O2} suggest that \hhoo\ is produced on dust-grains via the hydrogenation of grain HO$_{2}$ and released into the gas-phase through surface reactions. On the dust-grain, \hhoo\ acts as an intermediate in the formation of water and aids in the production of other species such as H$_2$CO, CH$_3$OH, and O$_2$. Hydrogen peroxide belongs to the peroxide group of molecules with an HO-OH bond dissociation enthalpy of 17050 \cm \citep{96BaAySc.H2O2} at 0 K. \hhoo\ is an asymmetric prolate rotor molecule and is the simplest molecule that exhibits internal rotation. This torsional motion gives rise to a double minimum potential curve with respect to its internal rotation co-ordinates as well as two alignments of the O-H bonds: \textit{cis} and \textit{trans}. The consequence of this motion means that there are four sub-levels for each torsional excitation which are characterized by their symmetry. This necessitates the use of an additional quantum number, $\tau$, to unambiguously describe its motion. The molecular states can be classified using the $C^{+}_{\rm 2h}$(M) symmetry group which best describes the torsional splitting caused by the \textit{cis} and \textit{trans} tunneling \citep{84HoXXXX.H2O2}. \hhoo\ has six vibrational modes: $\nu_{1}$ and $\nu_{5}$ represent the symmetric and asymmetric O-H stretching respectively, $\nu_{3}$ and $\nu_{6}$ represent the O-H bending modes, $\nu_{2}$ represents the O-O stretch and the $\nu_{4}$ mode represents the torsional excitation with the more common notation of $n$. Experimental studies of ro-vibrational \hhoo\ spectra have mostly probed the torsional motion in the ground \citep{88OlHuYo.H2O2}, the $\nu_{3}$ \citep{92CaFlJo.H2O2} and $\nu_{6}$ \citep{90PeFlCa.H2O2,95PeVaFl.H2O2} vibrational modes. Conversely, the higher-lying O-H stretching modes, $\nu_{1}$ and $\nu_{5}$, are poorly studied using high resolution techniques. The difference between the two stretching bands is about 8 -- 10 \icm\ and torsional splitting from the double minimum of the potential gives rise to doubling \citep{74GiSrXX.H2O2} in the form of 'quasi'-degenerate states \citep{09RaKnWe.H2O2} that are difficult to resolve with a degree of accuracy. \citet{88OlHuYo.H2O2} give an estimate of 3610 - 3618~\icm\ for $\nu_{5}$ and 3601 - 3617~\icm\ for $\nu_{1}$ whilst a Raman study gives a lower value of 3607~\icm\ for the ~$\nu_{1}$~ band-centre \citep{74GiSrXX.H2O2} but determining the accuracy to better than 0.1~\icm\ is difficult. \hhoo\ has been a benchmark system for developing methods aiming to treat large amplitude motion \citep{00LuXXXX.H2O2,02MlXXXX.H2O2,02YuMuXX.H2O2,09CaHaBo.H2O2}. Recent calculations on the ro-vibrational states for \hhoo\ include the {\it ab initio} computation using CCSD(T)-F12 electronic structure calculations of band frequencies accurate to about 4.0 \icm\ by \citet{09RaKnWe.H2O2}, models of the peroxide stretches by \citet{05BaBiXX.H2O2}, a discrete variable representation (DVR) calculation for levels up to 6000 \icm\ by \citet{01ChMaGu.H2O2,03LiGuxx.H2O2} and finally, a potential energy surface (PES) calculations by \citet{98KoCaHa.H2O2} and \citet{99KuRiLu.H2O2}. Calculation which also consider transition intensities are rather rarer but a recent example is provided by \citet{11CaShBo.H2O2}. % The peroxide system was used to benchmark the large amplitude calculations of MULTIMODE \citep{03BoCaHa.methods} up to $J=20$ and showed good agreement against HITRAN line intensities but the PES used had an rms of $\approx$ 20 \icm\ against experimental band centers. However, this PES has been superceded by the higher accuracy \ai\ potential energy surface (PES) of \citet{13MaKoXX.H2O2} which was further modified by \citet{jt553}. This modified PES was used for our room-temperature line list \citep{jt620} and provides the starting point for the refinements performed here. Experimental transitions frequencies and intensities for \hhoo\ are available in the HITRAN 2012 database \citep{jt557} but only for room temperature modelling up to 1800~\cm. This region covers the torsional, O-H bending modes and O-O stretch but misses the O-H stretches in the 3750~\cm~region. Only a few studies deal with absolute intensities of \hhoo\ in the far-infrared \citep{41ZuGiXX.H2O2,96PeFlCA.H2O2,95PeVaFl.H2O2} with only PNNL-IR \citep{PNNL} data providing integrated intensities in the mid-infrared region \citep{09JoSaBu.H2O2}. The thermal decomposition of hydrogen peroxide at 423 K makes it difficult and dangerous to study at higher temperatures. Theoretical line lists can be used to fill in gaps in the experimental data both in terms of wavelength and temperature coverage. The ExoMol project \citep{jt528} aims to produce comprehensive theoretical molecular line lists to aid in studies of the atmospheres of exoplanets, cool stars and other (hot) bodies. A room temperature line list for \hhoo\ was previously computed by us \citep{jt620} using the PES of \citet{jt553} and a new \ai\ dipole moment surface (DMS). It provides about 1 billion transitions at up to 8,000~\icm. However it is limited as the rotational excitation of $J=40$ makes it inadequate for high temperature modelling and the lower energy cut-off means that coverage above 4,000~\icm\ rapidly becomes incomplete. This work aims to build upon this line list by refining the PES towards spectroscopic accuracy and extending the temperature and frequency range for which the resulting line list is applicable.

The frequency and Einstein-$A$ coefficients of almost 20 billion transitions of hydrogen peroxide are computed. These transitions cover wavelengths longer than 1.6 $\mu$m and include all rotational excitations up to \(J=85\), making the line list applicable for temperatures up to 1250 K. The line list gives a room-temperature spectrum in excellent agreement with available experimental data and has good predictive ability for bands and line-positions not available experimentally. The new line list may be accessed via \url{www.exomol.com} or \url{http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/MNRAS/}. The cross-sections of \hhoo\ can be also generated at \url{www.exomol.com} as described by \citet{jt631}.

1607.00498

1607

1607.02440_arXiv.txt

We present the results of spectropolarimetric observations of 12 Herbig Ae/Be objects. Our data have the largest spectropolarimetric wavelength coverage, 4560~$\rm\AA$ to $9480~{\rm \AA}$, published to date. A change in linear polarisation across the H$\alpha$ line, is detected in all objects. Such a line effect reveals the fact that stellar photons are scattered off free electrons that are not distributed in a spherically symmetric volume, suggesting the presence of small disks around these accreting objects. Thanks to the large wavelength coverage, we can report that H$\alpha$ is the spectral line in the optical wavelength range that is most sensitive to revealing deviations from spherical symmetry, and the one most likely to show a line effect across the polarisation in such cases. Few other spectral lines display changes in polarisation across the line. In addition, H$\alpha$ is the only line which shows an effect across its absorption component in some sources. We present a scenario explaining this finding and demonstrate that the detection of the line effect strongly relies on the number of photons scattered into our line of sight. We highlight the special case of R Mon, which is the only object in our sample to show many lines with a polarisation effect, which is much stronger than in all other objects. Given that the object and its nebulosity is spatially resolved, we argue that this is due to scattering of the stellar and emission spectrum off circumstellar dust.

Despite the important role of massive stars in the evolution of the interstellar medium and unlike their counterpart low mass stars that are thought to be formed via magnetospheric accretion \citep{1998ApJ...492..743M}, their accretion mechanism is still open to debate. Observationally, the detection of massive stars is challenging as they reside far away and are deeply embedded in their natal clouds. As an alternative to massive stars, the intermediate mass Herbig Ae/Be (HAeBe) stars are the best candidates to address this issue. They are optically visible and bridge the gap between low and high mass stars. They were first identified by Herbig as having a spectral type A or B with emission lines \citep{1960ApJS....4..337H}. The current view is that HAeBe stars are surrounded by disks through which material continues to accrete onto the star (e.g. \citealt{ilee14}), although the precise scenario is not known (for reviews on the topic, see \citealt{grady15, kraus15,beltran16}). In order to make progress, one needs to study the circumstellar environment through which the material accumulates onto the star via accretion channels. This requires methods capable of probing the matter very close to the star, as this is where the accretion action happens. A key point is the observation whether the ionised region around the star is spherically symmetric at small scales or not. If it is not, then the possibility that a flattened structure is observed lends support to the disk accretion scenario responsible for the formation of such stars. Studying the circumstellar environment at these small scales is possible through linear spectropolarimetry, measuring the scattering of photons off free electrons in a dense, ionised gas. The idea of spectropolarimetry is that in the ionised region, free electrons scatter and polarise continuum photons from the central star. If the geometry is not circular on the sky, for example in the case of flattened disk, a net polarisation can be detected. However, emission photons undergo less scattering as they emerge further away from the central star. If we are to measure the polarisation as a function of wavelength, this difference in scattering will be visible as a change in polarisation across the line, which is often simply referred to as the ``line effect''. The observed de-polarisation across the H$\alpha$ line constituted the first evidence that the class of Be stars were surrounded by disks \citep{1974MNRAS.167P..27C, 1976ApJ...206..182P}. This was only much later confirmed by direct observations (e.g. \citealt{quirrenbach94} based on image reconstruction of interferometric data; \citealt{wheelwright12} using sub-milliarcsecond precision spectro-astrometry). \citet{1999MNRAS.305..166O}, \citet{Vink2002} and \citet{Mottram07} extended the use of the technique to HAeBe objects. \citet{Vink2002} found that 7 out of 12 Herbig Be (HBe) objects they observed have a depolarisation line effect across H$\alpha$, very similar to what was found in Be stars. As the line-effect becomes less pronounced for lower inclinations, and would disappear for a face-on disk, which is circular on the sky, the detection statistic strongly suggests that all HBe stars are surrounded by small disks with sizes of order several stellar radii. In contrast, they found a different line effect for Herbig Ae (HAe) objects, where enhanced {\it polarisation} across the H$\alpha$ was found in 9 out of 11 stars. They proposed that the line itself is intrinsically polarised since part of the emission lines originate from a compact region, where the accretion takes place. \citet{vink05a} found that HAe stars have a similar spectropolarimetric signature as the lower mass T Tauri stars with HBe stars having a different signature. \citet{McLean1979} reported a different line effect, across the absorption component of the emission line which is often called the McLean effect. The general idea is that the absorption blocks the unscattered light from the beam, and photons originally emitted in different directions are scattered into the line of sight, resulting in enhanced polarisation across the absorption. An alternate hypothesis was provided by \citet{kuhn07}, who proposed the polarisation can be caused by selective absorption due to optical pumping of an anisotropic radiation field. In addition, \citet{kuhn11} point out that resonant line scattering, which potentially also produces line polarisation, predicts that the lines within the {Ca \sc ii} near-infrared triplet around 8500~$\rm \AA$ will be differently polarised from each other. For a recent review on the use of linear spectropolarimetry, see \citet{vink15}. In this work we aim to expand the existing spectropolarimetric work that was mostly aimed at H$\alpha$ by observing other emission lines probing different volumes and conditions than H$\alpha$. Here we present a spectropolarimetric study of a sample of 12 HAeBe objects. The new feature of the current study is the broad wavelength range from 4560~$\rm\AA$ to $9480~{\rm \AA}$, covering almost the entire optical spectrum. The paper is organised as follows. In Section 2 we discuss the sample selection criteria, the details of the observations and data reduction. In Section 3 we present the results starting with continuum polarisation and then discussing line spectropolarimetry. The analysis is provided in Section 4. We conclude in Section 5.

\begin{itemize} \item We sample linear line spectropolarimetry in the optical wavelength range, which is much larger than any previous work at similar spectral resolution. \item Changes in the polarisation across the H$\alpha$ emission line are detected in all objects, as an indication of a flattened structure of the circumstellar environment. The line effects vary from depolarisation, line polarisation to the McLean effect. \item Depolarisation and the McLean effect are observed in Be type stars predominantly in early B type stars while line polarisation is observed in Ae type objects. \item The McLean effect is observed only across the absorptive component of H$\alpha$ and the line effect is stronger for a weaker absorption component, while H$\beta$ does not display the effect. We propose a scenario to explain this property. It is based on the fact that the photons from the strong H$\alpha$ line are scattered into the line of sight. As a consequence, the photons in the absorption are more polarised than the emission. A side result is that the selective absorption due to optical pumping as proposed by \citet{kuhn07} is unlikely to be responsible for the polarisation behaviour in these objects. \item We detect a broad depolarisation line effect across {Ca~\sc ii} triplet and {[O \sc i]} in two objects. These lines are emerging further away from the star and H$\alpha$ region in the circumstellar environment. The depolarisation simply implies that the circumstellar environment has an asymmetrical structure in this region. \item The few spectra with calcium triplet lines that show an effect show a similar polarisation profile for all members of the triplet. We confirm a similar observation by \citet{kuhn11} for an evolved, RV Tau star. As these authors explained, resonant line scattering can then not be the cause for the observed line polarisation, as the {Ca~\sc ii} 8662~$\rm \AA$ comes from a different upper level than the other members of the triplet. Similar polarisation due to line scattering would then not be expected. \item Finally, apart from H$\alpha$, few lines in few objects show a line effect in the polarisation. We present the case of R Mon, which displays exceptionally strong line effects in most of its emission lines. In addition, not all lines show the same type of line effect. We explain this by the fact that the object itself is resolved and propose that different line forming regions and geometries are responsible for polarisation properties of the higher hydrogen recombination lines and forbidden lines and H$\alpha$ on the other hand. \end{itemize}

1607.02440

1607

1607.07323_arXiv.txt

Magnetic reconnection is a {plasma} phenomenon where a topological rearrangement of magnetic field lines with opposite polarity results in dissipation of magnetic energy into heat, kinetic energy and particle acceleration. Such a phenomenon is considered as an efficient mechanism for energy release in laboratory and astrophysical plasmas. {An important question is how to make the process fast enough to account for observed explosive energy releases.} The classical model for steady state magnetic reconnection predicts reconnection times scaling as $S^{1/2}$ (where $S$ is the Lundquist number) and yields times scales several order of magnitude larger than the observed ones. Earlier two-dimensional MHD simulations showed that for large Lundquist number the reconnection time becomes independent of $S$ (``fast reconnection'' regime) due to the presence of the secondary tearing instability that takes place for $S \gtrsim 1 \times 10^4$. We report on our 3D MHD simulations of magnetic reconnection in a magnetically confined cylindrical plasma column under either a pressure balanced or a force-free equilibrium and compare the results with 2D simulations of a circular current sheet. We find that the 3D instabilities acting on these configurations result in a fragmentation of the initial current sheet in small filaments, \rev{leading to enhanced dissipation rate that becomes independent of the Lundquist number already at $S \simeq 1\times 10^3$.}

Magnetic reconnection is \ess{a plasma } phenomenon where a rapid rearrangement of magnetic fields of opposite polarity leads to the dissipation of the magnetic energy into heat, plasma kinetic energy and particle acceleration. \ess{In particular}, magnetic reconnection is generally regarded as a mechanism to account for the fast (i.e. much shorter than the dynamical time-scale) and intense variability observed in many astrophysical environments, like active galactic nuclei \citep{Giannios2013} and pulsar wind nebulae \citep{Cerutti2013}. It is also \ess{likely to occur} in space environments like solar flares and coronal mass ejection \citep{Gordovskyy2010a,Gordovskyy2011,Drake2006}. A measure of the conversion of magnetic energy into particle acceleration via magnetic reconnection in Earth's magnetosphere is reported in a recent paper of \cite{Burch2016}. Finally, magnetic reconnection is responsible for sawtooth crashes that prevent the magnetic confinement in laboratory fusion experiments, such as tokamaks \citep{Hastie1997}. The general features of steady state magnetic reconnection are described by the theory of Sweet-Parker \citep{Sweet1958,Parker57}, that \ess{proposed reconnection taking place in current sheets} (\rev{localized} regions of very intense currents where non-ideal effects become important) of length $L$ and thickness $\delta$. In this model the reconnection time scales as $S^{1/2}$ (where $S = LV_{A}/\eta$, is the Lundquist number, $L$ is the characteristic length of the field configuration, $V_{A}$ is the Alfv\'en velocity and $\eta$ is the resistivity). However, considering that the Lundquist number is very large in space, astrophysical and laboratory plasmas \citep[e.g. $S \sim 10^{12}\-- 10^{14}$ in the solar corona and $S \sim 10^6 \-- 10^8$ in tokamaks, see ][]{Loureiro2016}, the above mentioned scaling yields reconnection time-scales that are several order of magnitudes longer than observed. An attempt to solve this problem was suggested by Petschek \citep{Petschek1964}, whose model yields a logarithmic dependence of the reconnection rate on $S$. \ess{Petschek-like configuration and scaling are found in a recent relativistic resistive magnetohydrodynamic (MHD) simulation of \cite{DelZanna2016}. However, this regime was never observed in laboratory experiments. The understanding of this time-scale problem was significantly improved using resistive MHD numerical simulations with large Lundquist number. Two-dimensional simulations \citep[see, e.g., ][]{ Samtaney2009,Loureiro2012,Huang2010,Huang2013} have shown that when $S > S_c\simeq 1 \times 10^4$,} the current sheet is subject to secondary tearing instability \citep{Biskamp1986}, \ess{resulting in the} fragmentation of the current sheet \ess{and formation of } a large number of plasmoids. This leads to \ess{the so called} ``fast reconnection'' regime, where the reconnection rate becomes independent of the resistivity. Recent three-dimensional MHD simulations \citep{Oishi2015} have further shown that 3D instabilities can trigger a ``fast reconnection'' regime even for $S < S_c$. In the present work, we consider both 2D and 3D resistive MHD simulations of magnetically confined cylindrical plasma columns (\es{see Fig. \ref{Plasma_Column}}) featuring a current ring where the azimuthal component of magnetic field changes polarity. \rev{This field configuration was considered in \cite{Romanova1992} and more recently in \cite{McKinney2012}.} We consider two initial equilibria: one in which radial force balance is established by a thermal pressure gradient and one in which the field is force-free. The former can become unstable to pressure-driven instabilities while the latter is prone to the onset of current-driven modes. We then show that the presence of 3D plasma column instabilities results in a fragmentation of the initial current sheet and leads to a ``fast reconnection'' regime also for $S \simeq 1\times 10^3$. This paper is organized as follows. In section \ref{Equations} we summarize the equations of resistive MHD used in the simulations and we present our model setup and initial conditions. In section \ref{Results} we illustrate the results of our simulations. Finally in section \ref{Discussion} we summarize and discuss our findings.

\label{Discussion} We studied magnetic reconnection using three-dimensional resistive MHD simulations of a magnetically confined cylindrical plasma column featuring a circular current sheet. Different equilibrium conditions, including radial pressure balance and a force-free field, have been considered. Results have been compared with 2-dimensional simulation of a circular current sheet. Our 2D simulations generalize previous studies of planar current sheets to the cylindrical case. The main results from these simulations are listed below : \begin{itemize} \item At early stages (phase I), the magnetic dissipation rate in the current ring agrees with the Sweet-Parker scaling of $S^{-0.5}$. \item At later times (phase II) and for values of $S \gtrsim S_c \simeq 1 \times 10^4$, the current sheet is subjected to secondary tearing instability whereby continuous formation of plasmoids is observed. The formation of plasmoids leads to the fragmentation of the initial circular sheet into multiple small-sized current sheets. During this stage, the decay rate increases sharply, and becomes independent of $S$, revealing the transition to a regime of fast reconnection. \item Eventually, the continuous formation and merging of plasmoids results in the random orientation of fragmented current-sheets that closely resemble the turbulent reconnection described by, e.g., \cite{Kowal2009}, \cite{Loureiro2009} and \cite{Takamoto2015}. \item The rate of dissipation of magnetic energy during the fast reconnection regime is \rev{$\sim 0.1 t_A^{-1}$}, consistent with previous numerical results of 2D reconnection. \end{itemize} In the three-dimensional case, our results can be summarized as follows: \begin{itemize} \item Similar to the 2D runs, the magnetic energy (during the initial phase) is dissipated at a rate which is consistent with Sweet-Parker scaling, $S^{-0.5}$. \item At later times the plasma column becomes unstable to either pressure-driven or current-driven instabilities (depending on the initial equilibrium configuration), the onset of which does not depend on the Lundquist number. In runs with same set of parameters ($\beta = 10$ and $P = 0$), the 3D pressure-driven instability starts before the 2D secondary tearing mode. The growth of these instabilities causes the fragmentation of the original current ring into smaller secondary current sheets (see Fig. \ref{Fig:3DPB} and Fig. \ref{Fig:3DFF}). \item \rev{At this time an increased magnetic dissipation is observed (phase II).} The \rev{dissipation rate} becomes independent of $S$ and is of the order of \rev{$(\sim 0.1 \-- 0.5) t_A^{-1}$} (see Fig. \ref{Fig:DissipationPB} and Fig. \ref{Fig:DissipationFF}). \item The \rev{dissipation rate} starts to become independent of $S$ for $S \simeq 10^{-3}$, a threshold value which is an order of magnitude smaller than the one obtained from 2D runs. \end{itemize} \rev{We point out that the dissipation rates reported here result from the interplay between magnetic reconnection and the turbulence induced by the instabilities arising in each configurations. This may lead to energy dissipation rates that are faster than the actual reconnection rate and could explain the differences between our findings ($\gtrsim 0.1 t_A^{-1}$) and the results reported in previous reference studies ($\sim 0.01 t_A^{-1}$). On the other hand, three-dimensional simulations without magnetic shear, that are not expected to develop magnetic reconnection, do not show relevant dissipation. In summary, we find that the 3D instabilities alone dissipate the magnetic energy inefficiently. However, they play a major role in enhancing the rate of magnetic dissipation in presence of reconnection. } We emphasize that the Lundquist numbers for the above 3D simulations lie in the range $10^3\--10^4$ and no formation of secondary tearing instability is observed. The ``fast reconnection'' regime is, therefore, a mere effect of the 3D instabilities. A similar effect was reported in recent 3D simulations by \cite{Oishi2015}, where they attributed the early fast reconnection regime to an unspecified 3D instability. Our detailed analysis obtains consistent results in a different configuration (magnetically confined plasma column) and provides clear evidence that the onset of ``fast reconnection'' is triggered by well-known plasma instabilities (pressure- or current-driven). Our results can be relevant in the context of MHD jets, where these instabilities are likely to operate. Typical astrophysical environments are active galactic nuclei, microquasars and pulsar wind nebulae. Here, magnetic reconnection has been recently invoked as an efficient mechanism to accelerate particles to non-thermal energies \citep{Sironi2014,deGouveia2015} up to PeV energies \citep{Cerutti2013}. Plasma instabilities in jets, therefore, could trigger fast magnetic reconnection episodes \citep{Lyubarsky2012,Giannios2013} that may account for the observed fast variability and non-thermal features in these astrophysical scenarios, like, e.g., the $\gamma$-ray flares from the Crab Nebula \citep{Tavani2011,Striani2011}, or the very rapid variability, $\sim$ 10 min, detected, e.g., in PKS 2155 \citep{Aharonian2007} and PKS 1222 \citep{Aleksic2011}. Our results can, however, be applied only in the reference frame of the jet as no velocity shear has been considered. Besides, a more detailed analysis would require direct investigation of particle acceleration. These issues will be explored in forthcoming studies.

1607.07323

1607

1607.03108.txt

\label{sec:intro} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% According to the Weakly Interacting Massive Particle (WIMP) paradigm, the Dark Matter (DM) component of our Universe is composed of non-relativistic particles whose abundance is set by their annihilations in the early Universe~\cite{Lee:1977ua,Vysotsky:1977pe,Kolb:1990vq}. At high temperatures ($T>m_{DM}$), DM particles are assumed to be in thermal equilibrium with the Standard Model (SM) plasma. As the temperature drops below their mass, their abundance starts to decrease exponentially and number changing annihilation processes like ${\textrm{DM}}\,{\textrm{DM}}\leftrightarrow {\textrm{SM}}\, {\textrm{SM}}$ become inefficient. Eventually, when $T/m_{DM}\lesssim 1/{30}$, the DM comoving number density \emph{freezes out}. The resulting energy density is determined by the DM annihilation cross section, $\left< \sigma v \right>$, \be \label{eq:VanillaOmega} \Omega_{DM}h^2\simeq 0.1\frac{ (20\,{\textrm{TeV}})^{-2}}{\langle\sigma v\rangle}. \ee This result implies that a particle with weak scale mass and electroweak size interactions ($\left< \sigma v \right> \approx (20~\mathrm{TeV})^{-2} \approx 3 \times 10^{-26}~\mathrm{cm}^3/\mathrm{s}$) has a freeze-out abundance that matches the observed relic density: $\Omega_{DM}h^2\approx 0.1$~\cite{Ade:2015xua}. Such a framework has various appealing features. The freeze-out mechanism is insensitive to any initial condition or UV physics due to its thermal nature. If DM mainly annihilates into SM particles, then it must have sizable interactions with the SM, opening up exciting experimental possibilities for its observation. Furthermore, the existence of new weak scale particles is motivated by theoretical considerations such as solving the naturalness problem of the Higgs mass. However, recent experimental results are challenging this picture: DM direct detection experiments have excluded significant WIMP parameter space~\cite{Cushman:2013zza,Akerib:2015rjg,Aprile:2015uzo}, and are getting close to the neutrino background~\cite{Billard:2013qya}. At the same time, collider and precision searches are increasingly constraining the possible existence of new physics around the weak scale. Pending a discovery of a WIMP, it is highly motivated to explore new theories of dark matter and their experimental tests. One possibility is that DM is part of a hidden sector of particles that is very weakly coupled to the SM\@ (see for example Refs.~\cite{Kolb:1985bf,Goldberg:1986nk,Carlson:1992fn,Strassler:2006im,Finkbeiner:2007kk,Pospelov:2007mp,Feng:2008ya,Feng:2008mu,ArkaniHamed:2008qn,Kaplan:2009ag,Cheung:2010gj,Hochberg:2014dra,D'Agnolo:2015koa,cannibal}). At sufficiently high temperatures, interactions within the dark sector will guarantee thermal equilibrium among the dark sector particles. However, if interactions between the dark sector and the SM are sufficiently weak, the dark sector temperature $T_D$ and entropy $s_D$ will in general be different from the SM ones~\cite{Carlson:1992fn,Feng:2008mu,Cheung:2010gj,cannibal}. The fact that the DM is in equilibrium with a thermal bath while its energy density evolves to the observed value implies the same insensitivity to UV dynamics as for standard WIMPs (except for sensitivity to the initial condition that sets the relative temperatures of the SM and dark sector plasmas). \begin{figure}[!!!t] \begin{center} \includegraphics[width=0.4\textwidth]{spectrum.pdf} \end{center} \vspace{-.3cm} \caption{\small { The typical mass spectrum of a gapped hidden sector. $\chi$ plays the role of the DM, which we assume to be stabilized by some symmetry. The various $\phi_i$ are generically unstable. $\phi_{0}$ is the Lightest Dark sector Particle (LDP) and has nonzero mass. The hidden sector will undergo an epoch of cannibalism if the temperature drops below the mass of the LDP, $T_d < m_0$, while number changing interactions are still in equilibrium.}} \label{fig:spectrum} \end{figure} Most works studying hidden sector dark matter assume that the hidden sector contains relativistic particles in thermal equilibrium with DM when its annihilations freeze-out. An alternative possibility is that DM resides in a hidden sector with a mass gap set by the mass of the Lightest Dark sector Particle (LDP), $m_0$. If the hidden sector is sufficiently weakly coupled to the SM, this opens up the possibility that DM is not in thermal contact with radiation when its annihilations decouple. This implies that the hidden sector undergoes an epoch of cannibalism~\cite{Carlson:1992fn}, during which its temperature decreases only logarithmically with the scale factor. This possibility was first studied by Ref.~\cite{Carlson:1992fn}, which considers the possibility that DM freezes out through 3-to-2 annihilations. Recently, some of us proposed that dark matter may reside in a hidden sector with a mass gap and have abundance that follows from 2-to-2 annihilations~\cite{cannibal}. For additional studies that include cannibalism see Refs.~\cite{deLaix:1995vi,Boddy:2014qxa,Yamanaka:2014pva,Garcia:2015loa,Bernal:2015ova,Bernal:2015xba,Kuflik:2015isi,Soni:2016gzf,Forestell:2016qhc}. In this paper, we consider the following framework~\cite{cannibal}: \begin{itemize} \item {\bf Cannibal Dark Matter:} the abundance of DM is set by the freeze-out of 2-to-2 annihilations in a hidden sector that undergoes an epoch of cannibalism that begins before DM annihilations decouple. \end{itemize} Our goal is to map out the different possible cosmologies for Cannibal DM, in order to identify viable models that reproduce the observed relic density. We will find that gapped hidden sectors have a rich phase structure, with multiple novel avenues for DM freeze-out. We assume that dark matter resides in a hidden sector with a mass gap, and that the hidden sector is kinetically decoupled from the SM such that it evolves with an independent temperature. The hidden sector contains a stable particle, $\chi$, that will constitute DM, as well as massive unstable particles $\phi_i$ (see Fig.~\ref{fig:spectrum}). We can identify three relevant timescales in the cosmological evolution of the hidden sector: \begin{itemize} \item $t_f$ (and the corresponding dark temperature $T_f$)--- the time after which the rate of DM number changing processes, $\Gamma_f$, is smaller than the Hubble constant, $H$, \be \Gamma_f \equiv n_\chi^{eq} \left< \sigma_2 v \right> < H, \ee where $\sigma_2$ is the cross section for $\chi \chi \rightarrow \phi \phi$ (which is the leading process that changes DM number density) and $n_\chi^{eq}$ is the equilibrium number density of $\chi$. Here and below the indices on the $\phi_i$ are implied. In general, $n_\chi^{eq} \propto e^{-(m_\chi-\mu_\chi)/T_D}$ includes a chemical potential, $\mu_\chi$, which as we will see below can play an important role. After $t_f$, the comoving DM density is conserved. \item $t_c$ (and the corresponding dark temperature $T_c$)--- the time after which the rate of number changing processes in the hidden sector, $\Gamma_c$, falls below $H$. This timescale is set by the decoupling of the final $n \leftrightarrow m$ process with $n \ne m$. For example, $t_c$ can be set by the decoupling of $\phi \phi \phi \rightarrow \phi \phi$ with cross section $\sigma_3$, \be \Gamma_c \equiv (n_\phi^{eq})^2 \left< \sigma_3 v^2 \right> < H. \ee After the time $t_c$, the total comoving number of particles in the hidden sector is conserved. However, 2-to-2 processes may still be active and can change relative particle abundances. A chemical potential is necessary to enforce the conservation of the total comoving number of hidden sector particles. %%%% \item $t_d$ (and the corresponding dark temperature $T_d$)--- the time at which the DM annihilation products, $\phi$, decay more rapidly than the Hubble time ({\it i.e.} the Universe becomes older than the $\phi$ lifetime), \be \Gamma_\phi > H. \ee As we discuss below, metastability of $\phi_i$ is a necessary condition for $\chi$ to constitute the totality of the observed DM abundance~\cite{cannibal}. If this was not the case, the lightest of the $\phi_i$, being non-relativistic, would dominate the dark sector energy density. \end{itemize} %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% \begin{figure}[!!!t] \begin{center} \includegraphics[width=0.75\textwidth]{PhaseSchematic.pdf} \end{center} \vspace{-.3cm} \caption{\small The three phases of cannibal dark matter: {\it cannibal}, {\it chemical}, and {\it one way}. The phase depends on the ordering of three timescales: $t_f$, $t_c$, and $t_d$ (which are defined in the text). The hidden sector undergoes cannibalism in the red shaded region, which begins when $T_d = m_0$ at time $t_\phi$. Cannibalism ends whichever occurs first: $t_c$ or $t_d$, and either ordering is possible for the cannibal phase. The number of hidden particles is conserved in the blue shaded region, because only 2-to-2 interactions are in equilibrium, implying a nonzero chemical potential. } \label{fig:schematic} \end{figure} %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% \begin{figure}[!!!t] \begin{center} \includegraphics[width=0.5\textwidth]{PhaseDiagram.pdf} \end{center} \vspace{-.3cm} \caption{\small The phase diagram of cannibal dark matter. There are three phases that describe the behavior of the hidden sector when DM annihilations freeze-out: {\it cannibal}, {\it chemical}, and {\it one way}. The phase depends on the ordering of the three timescales: $t_f$, $t_c$, and $t_d$. The relic density takes a different parametric form within each phase, as shown in Eqs.~\ref{scalingcann}, \ref{scalingchem}, and \ref{scalingdecay}. } \label{phases} \end{figure} The key point is that different orderings of these timescales lead to different parametric scalings for the DM relic density as a function of the fundamental parameters of the dark sector. We identify three distinct phases, depending on the order of $t_f$, $t_c$, and $t_d$ (see Fig.~\ref{fig:schematic}). \begin{itemize} \item {\bf Cannibal phase} ($t_f\ll t_c,\, t_d$): this is the scenario proposed in Ref.~\cite{cannibal}. The freeze-out of DM number changing interactions takes place while the hidden sector is undergoing cannibalism. The final relic density is exponentially sensitive to the ratio of the LDP and DM masses, $r = m_0 / m_\chi$. The relic density depends on the $\chi \chi \leftrightarrow \phi \phi$ cross section, $\sigma_2$. Assuming the Universe to be radiation dominated at freeze-out, \be\label{scalingcann} Y_\chi \propto(m_{\chi}M_P\sigma_2)^{-\frac{1-r}{1-2/3\, r}}, \ee where $M_P \approx 1.2 \times 10^{19}~\mathrm{GeV}$ is the Planck mass and $Y_{DM} = n_{DM} / s_{SM}$ is the DM yield (see eq.~\ref{Ycann}). The yield is related to the observed relic density: $\Omega_{DM}/\Omega_{DM}^{\textrm{obs}}=m_{DM}Y_{DM}/(0.4\,{\textrm{eV}})$. %%%% \item {\bf{Chemical phase}} ($t_c\ll t_f\ll t_d$): in this case, at time $t_c$, a chemical potential develops enforcing the conservation of the total number of hidden sector particles. The DM relic density depends on the size of the number changing cross section, $\sigma_3$, through this chemical potential. The DM relic density depends inversely on its annihilation cross section, as for a regular WIMP\@. Assuming the Universe to be radiation dominated at freeze-out (see eq.~\ref{Ychem}), \be\label{scalingchem} Y_{\chi}\propto \frac{(m_{0}^4 M_P \sigma_3)^{1/4}}{m_{\chi} M_P \sigma_2}. \ee Notice the relic abundance is no longer exponentially sensitive to $r$. %%%% \item {\bf{One way phase}} ($t_d\ll t_f$ any value of $t_c>t_\phi$): if the lifetime of the states to which DM is annihilating is shorter than $t_f$, inverse annihilations $\phi \phi \to \chi \chi$ decouple when $\phi$ decays at the time $t_d$. After $\phi$ decays, forward annihilations $\chi \chi \to \phi \phi$ are still active, but inverse annihilations are suppressed. The resulting DM yield depends on the width of $\phi$ (see Eq.~\ref{Omegadecay}), \be\label{scalingdecay} Y_{\chi}\propto \frac{1}{\Gamma_{\phi}^{1/2}M_P^{3/2}\sigma_2}. \ee \end{itemize} %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% \begin{figure}[!!!t] \begin{center} ~~\includegraphics[width=0.45\textwidth]{cc.pdf}~~~~~\includegraphics[width=0.4639\textwidth]{cd.pdf} \end{center} \vspace{-.3cm} \caption{ \small The {\it left} side shows the dark matter relic density as a function of $\sigma_3$ (the cross section of $\phi \phi \phi \rightarrow \phi \phi$), for various values of $\sigma_2$ (the cross section of $\chi \chi \rightarrow \phi \phi$) normalized to the conventional WIMP value $\sigma_0 = 3 \times 10^{-26}~\mathrm{cm}^2/\mathrm{s}$. Moving from smaller to larger $\sigma_3$, the phase transitions from chemical (blue), where $\Omega_\chi$ increases with $\sigma_3$, to cannibal (red), where $\Omega_\chi$ is independent of $\sigma_3$. The {\it right} side shows $\Omega_\chi$ as a function of $\gamma_\phi=\Gamma_\phi/H(m_\phi)$. Moving from smaller to larger $\Gamma_\phi$, the phase transitions from cannibal (red), where $\Omega_\chi$ is independent of $\Gamma_\phi$, to one way (green) where $\Omega_\chi$ decreases with $\Gamma_\chi$. This figure was made using the model of Eqs.~\ref{eq:L} and~\ref{potential}. } \label{transition} \end{figure} %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% These three regimes are displayed in the phase diagram of Fig.~\ref{phases}. Fig.~\ref{transition} shows the behavior of the relic abundance, in terms of the relevant parameters, in the three phases and the transition region between them. In order to draw the curves we use the example model which will be introduced in Section~\ref{sec:GenCo}. Cannibal DM has distinctive phenomenology, which we explore below. Most notably, the DM annihilation rate is generically boosted above the standard value of a thermal WIMP, $\sigma_0 \approx 3 \times 10^{26}~\mathrm{cm}^3/\mathrm{s}$. This is a consequence of the novel parametric form of the DM yield (Eqs.~\ref{scalingcann}, \ref{scalingchem}, and \ref{scalingdecay}) and the requirement that Cannibal DM match the observed DM energy density. The parametric form of the boost to the cross section is summarized in Tab.~\ref{tab:boost} of Section~\ref{sec:GenCo}. The boost implies that Cannibal DM is easier to see through indirect detection than standard WIMPs. Observational constraints on cannibal DM, and future reach, are summarized in Figs.~\ref{figpheno1} and \ref{figpheno2}. Cannibal DM is constrained by Fermi measurements of $\gamma$-rays, Planck measurements of the Cosmic Microwave Background (CMB), and Lyman-$\alpha$ constraints on warm DM\@. However, significant parameter space remains allowed and we find that there is promising reach to discover cannibal DM through future CMB measurements. The rest of this paper is organized as follows. In Section~\ref{sec:Thermo}, we follow the evolution of a decoupled hidden sector as its temperature drops below the mass of the LDP, but when its number changing interactions are still active. In Section~\ref{sec:GenCo}, we introduce the example model that we use to describe the evolution of the DM density during the three phases. In Section~\ref{sec:pheno}, we study possible observational signatures and constraints on the model. Appendix~\ref{boltzmann} contains a complete description of the system of Boltzmann equations that we use to determine the DM energy density. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\label{sec:intro} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% According to the Weakly Interacting Massive Particle (WIMP) paradigm, the Dark Matter (DM) component of our Universe is composed of non-relativistic particles whose abundance is set by their annihilations in the early Universe~\cite{Lee:1977ua,Vysotsky:1977pe,Kolb:1990vq}. At high temperatures ($T>m_{DM}$), DM particles are assumed to be in thermal equilibrium with the Standard Model (SM) plasma. As the temperature drops below their mass, their abundance starts to decrease exponentially and number changing annihilation processes like ${\textrm{DM}}\,{\textrm{DM}}\leftrightarrow {\textrm{SM}}\, {\textrm{SM}}$ become inefficient. Eventually, when $T/m_{DM}\lesssim 1/{30}$, the DM comoving number density \emph{freezes out}. The resulting energy density is determined by the DM annihilation cross section, $\left< \sigma v \right>$, \be \label{eq:VanillaOmega} \Omega_{DM}h^2\simeq 0.1\frac{ (20\,{\textrm{TeV}})^{-2}}{\langle\sigma v\rangle}. \ee This result implies that a particle with weak scale mass and electroweak size interactions ($\left< \sigma v \right> \approx (20~\mathrm{TeV})^{-2} \approx 3 \times 10^{-26}~\mathrm{cm}^3/\mathrm{s}$) has a freeze-out abundance that matches the observed relic density: $\Omega_{DM}h^2\approx 0.1$~\cite{Ade:2015xua}. Such a framework has various appealing features. The freeze-out mechanism is insensitive to any initial condition or UV physics due to its thermal nature. If DM mainly annihilates into SM particles, then it must have sizable interactions with the SM, opening up exciting experimental possibilities for its observation. Furthermore, the existence of new weak scale particles is motivated by theoretical considerations such as solving the naturalness problem of the Higgs mass. However, recent experimental results are challenging this picture: DM direct detection experiments have excluded significant WIMP parameter space~\cite{Cushman:2013zza,Akerib:2015rjg,Aprile:2015uzo}, and are getting close to the neutrino background~\cite{Billard:2013qya}. At the same time, collider and precision searches are increasingly constraining the possible existence of new physics around the weak scale. Pending a discovery of a WIMP, it is highly motivated to explore new theories of dark matter and their experimental tests. One possibility is that DM is part of a hidden sector of particles that is very weakly coupled to the SM\@ (see for example Refs.~\cite{Kolb:1985bf,Goldberg:1986nk,Carlson:1992fn,Strassler:2006im,Finkbeiner:2007kk,Pospelov:2007mp,Feng:2008ya,Feng:2008mu,ArkaniHamed:2008qn,Kaplan:2009ag,Cheung:2010gj,Hochberg:2014dra,D'Agnolo:2015koa,cannibal}). At sufficiently high temperatures, interactions within the dark sector will guarantee thermal equilibrium among the dark sector particles. However, if interactions between the dark sector and the SM are sufficiently weak, the dark sector temperature $T_D$ and entropy $s_D$ will in general be different from the SM ones~\cite{Carlson:1992fn,Feng:2008mu,Cheung:2010gj,cannibal}. The fact that the DM is in equilibrium with a thermal bath while its energy density evolves to the observed value implies the same insensitivity to UV dynamics as for standard WIMPs (except for sensitivity to the initial condition that sets the relative temperatures of the SM and dark sector plasmas). \begin{figure}[!!!t] \begin{center} \includegraphics[width=0.4\textwidth]{spectrum.pdf} \end{center} \vspace{-.3cm} \caption{\small { The typical mass spectrum of a gapped hidden sector. $\chi$ plays the role of the DM, which we assume to be stabilized by some symmetry. The various $\phi_i$ are generically unstable. $\phi_{0}$ is the Lightest Dark sector Particle (LDP) and has nonzero mass. The hidden sector will undergo an epoch of cannibalism if the temperature drops below the mass of the LDP, $T_d < m_0$, while number changing interactions are still in equilibrium.}} \label{fig:spectrum} \end{figure} Most works studying hidden sector dark matter assume that the hidden sector contains relativistic particles in thermal equilibrium with DM when its annihilations freeze-out. An alternative possibility is that DM resides in a hidden sector with a mass gap set by the mass of the Lightest Dark sector Particle (LDP), $m_0$. If the hidden sector is sufficiently weakly coupled to the SM, this opens up the possibility that DM is not in thermal contact with radiation when its annihilations decouple. This implies that the hidden sector undergoes an epoch of cannibalism~\cite{Carlson:1992fn}, during which its temperature decreases only logarithmically with the scale factor. This possibility was first studied by Ref.~\cite{Carlson:1992fn}, which considers the possibility that DM freezes out through 3-to-2 annihilations. Recently, some of us proposed that dark matter may reside in a hidden sector with a mass gap and have abundance that follows from 2-to-2 annihilations~\cite{cannibal}. For additional studies that include cannibalism see Refs.~\cite{deLaix:1995vi,Boddy:2014qxa,Yamanaka:2014pva,Garcia:2015loa,Bernal:2015ova,Bernal:2015xba,Kuflik:2015isi,Soni:2016gzf,Forestell:2016qhc}. In this paper, we consider the following framework~\cite{cannibal}: \begin{itemize} \item {\bf Cannibal Dark Matter:} the abundance of DM is set by the freeze-out of 2-to-2 annihilations in a hidden sector that undergoes an epoch of cannibalism that begins before DM annihilations decouple. \end{itemize} Our goal is to map out the different possible cosmologies for Cannibal DM, in order to identify viable models that reproduce the observed relic density. We will find that gapped hidden sectors have a rich phase structure, with multiple novel avenues for DM freeze-out. We assume that dark matter resides in a hidden sector with a mass gap, and that the hidden sector is kinetically decoupled from the SM such that it evolves with an independent temperature. The hidden sector contains a stable particle, $\chi$, that will constitute DM, as well as massive unstable particles $\phi_i$ (see Fig.~\ref{fig:spectrum}). We can identify three relevant timescales in the cosmological evolution of the hidden sector: \begin{itemize} \item $t_f$ (and the corresponding dark temperature $T_f$)--- the time after which the rate of DM number changing processes, $\Gamma_f$, is smaller than the Hubble constant, $H$, \be \Gamma_f \equiv n_\chi^{eq} \left< \sigma_2 v \right> < H, \ee where $\sigma_2$ is the cross section for $\chi \chi \rightarrow \phi \phi$ (which is the leading process that changes DM number density) and $n_\chi^{eq}$ is the equilibrium number density of $\chi$. Here and below the indices on the $\phi_i$ are implied. In general, $n_\chi^{eq} \propto e^{-(m_\chi-\mu_\chi)/T_D}$ includes a chemical potential, $\mu_\chi$, which as we will see below can play an important role. After $t_f$, the comoving DM density is conserved. \item $t_c$ (and the corresponding dark temperature $T_c$)--- the time after which the rate of number changing processes in the hidden sector, $\Gamma_c$, falls below $H$. This timescale is set by the decoupling of the final $n \leftrightarrow m$ process with $n \ne m$. For example, $t_c$ can be set by the decoupling of $\phi \phi \phi \rightarrow \phi \phi$ with cross section $\sigma_3$, \be \Gamma_c \equiv (n_\phi^{eq})^2 \left< \sigma_3 v^2 \right> < H. \ee After the time $t_c$, the total comoving number of particles in the hidden sector is conserved. However, 2-to-2 processes may still be active and can change relative particle abundances. A chemical potential is necessary to enforce the conservation of the total comoving number of hidden sector particles. %%%% \item $t_d$ (and the corresponding dark temperature $T_d$)--- the time at which the DM annihilation products, $\phi$, decay more rapidly than the Hubble time ({\it i.e.} the Universe becomes older than the $\phi$ lifetime), \be \Gamma_\phi > H. \ee As we discuss below, metastability of $\phi_i$ is a necessary condition for $\chi$ to constitute the totality of the observed DM abundance~\cite{cannibal}. If this was not the case, the lightest of the $\phi_i$, being non-relativistic, would dominate the dark sector energy density. \end{itemize} %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% \begin{figure}[!!!t] \begin{center} \includegraphics[width=0.75\textwidth]{PhaseSchematic.pdf} \end{center} \vspace{-.3cm} \caption{\small The three phases of cannibal dark matter: {\it cannibal}, {\it chemical}, and {\it one way}. The phase depends on the ordering of three timescales: $t_f$, $t_c$, and $t_d$ (which are defined in the text). The hidden sector undergoes cannibalism in the red shaded region, which begins when $T_d = m_0$ at time $t_\phi$. Cannibalism ends whichever occurs first: $t_c$ or $t_d$, and either ordering is possible for the cannibal phase. The number of hidden particles is conserved in the blue shaded region, because only 2-to-2 interactions are in equilibrium, implying a nonzero chemical potential. } \label{fig:schematic} \end{figure} %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% \begin{figure}[!!!t] \begin{center} \includegraphics[width=0.5\textwidth]{PhaseDiagram.pdf} \end{center} \vspace{-.3cm} \caption{\small The phase diagram of cannibal dark matter. There are three phases that describe the behavior of the hidden sector when DM annihilations freeze-out: {\it cannibal}, {\it chemical}, and {\it one way}. The phase depends on the ordering of the three timescales: $t_f$, $t_c$, and $t_d$. The relic density takes a different parametric form within each phase, as shown in Eqs.~\ref{scalingcann}, \ref{scalingchem}, and \ref{scalingdecay}. } \label{phases} \end{figure} The key point is that different orderings of these timescales lead to different parametric scalings for the DM relic density as a function of the fundamental parameters of the dark sector. We identify three distinct phases, depending on the order of $t_f$, $t_c$, and $t_d$ (see Fig.~\ref{fig:schematic}). \begin{itemize} \item {\bf Cannibal phase} ($t_f\ll t_c,\, t_d$): this is the scenario proposed in Ref.~\cite{cannibal}. The freeze-out of DM number changing interactions takes place while the hidden sector is undergoing cannibalism. The final relic density is exponentially sensitive to the ratio of the LDP and DM masses, $r = m_0 / m_\chi$. The relic density depends on the $\chi \chi \leftrightarrow \phi \phi$ cross section, $\sigma_2$. Assuming the Universe to be radiation dominated at freeze-out, \be\label{scalingcann} Y_\chi \propto(m_{\chi}M_P\sigma_2)^{-\frac{1-r}{1-2/3\, r}}, \ee where $M_P \approx 1.2 \times 10^{19}~\mathrm{GeV}$ is the Planck mass and $Y_{DM} = n_{DM} / s_{SM}$ is the DM yield (see eq.~\ref{Ycann}). The yield is related to the observed relic density: $\Omega_{DM}/\Omega_{DM}^{\textrm{obs}}=m_{DM}Y_{DM}/(0.4\,{\textrm{eV}})$. %%%% \item {\bf{Chemical phase}} ($t_c\ll t_f\ll t_d$): in this case, at time $t_c$, a chemical potential develops enforcing the conservation of the total number of hidden sector particles. The DM relic density depends on the size of the number changing cross section, $\sigma_3$, through this chemical potential. The DM relic density depends inversely on its annihilation cross section, as for a regular WIMP\@. Assuming the Universe to be radiation dominated at freeze-out (see eq.~\ref{Ychem}), \be\label{scalingchem} Y_{\chi}\propto \frac{(m_{0}^4 M_P \sigma_3)^{1/4}}{m_{\chi} M_P \sigma_2}. \ee Notice the relic abundance is no longer exponentially sensitive to $r$. %%%% \item {\bf{One way phase}} ($t_d\ll t_f$ any value of $t_c>t_\phi$): if the lifetime of the states to which DM is annihilating is shorter than $t_f$, inverse annihilations $\phi \phi \to \chi \chi$ decouple when $\phi$ decays at the time $t_d$. After $\phi$ decays, forward annihilations $\chi \chi \to \phi \phi$ are still active, but inverse annihilations are suppressed. The resulting DM yield depends on the width of $\phi$ (see Eq.~\ref{Omegadecay}), \be\label{scalingdecay} Y_{\chi}\propto \frac{1}{\Gamma_{\phi}^{1/2}M_P^{3/2}\sigma_2}. \ee \end{itemize} %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% \begin{figure}[!!!t] \begin{center} ~~\includegraphics[width=0.45\textwidth]{cc.pdf}~~~~~\includegraphics[width=0.4639\textwidth]{cd.pdf} \end{center} \vspace{-.3cm} \caption{ \small The {\it left} side shows the dark matter relic density as a function of $\sigma_3$ (the cross section of $\phi \phi \phi \rightarrow \phi \phi$), for various values of $\sigma_2$ (the cross section of $\chi \chi \rightarrow \phi \phi$) normalized to the conventional WIMP value $\sigma_0 = 3 \times 10^{-26}~\mathrm{cm}^2/\mathrm{s}$. Moving from smaller to larger $\sigma_3$, the phase transitions from chemical (blue), where $\Omega_\chi$ increases with $\sigma_3$, to cannibal (red), where $\Omega_\chi$ is independent of $\sigma_3$. The {\it right} side shows $\Omega_\chi$ as a function of $\gamma_\phi=\Gamma_\phi/H(m_\phi)$. Moving from smaller to larger $\Gamma_\phi$, the phase transitions from cannibal (red), where $\Omega_\chi$ is independent of $\Gamma_\phi$, to one way (green) where $\Omega_\chi$ decreases with $\Gamma_\chi$. This figure was made using the model of Eqs.~\ref{eq:L} and~\ref{potential}. } \label{transition} \end{figure} %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% %%%%%%%%%%%%% These three regimes are displayed in the phase diagram of Fig.~\ref{phases}. Fig.~\ref{transition} shows the behavior of the relic abundance, in terms of the relevant parameters, in the three phases and the transition region between them. In order to draw the curves we use the example model which will be introduced in Section~\ref{sec:GenCo}. Cannibal DM has distinctive phenomenology, which we explore below. Most notably, the DM annihilation rate is generically boosted above the standard value of a thermal WIMP, $\sigma_0 \approx 3 \times 10^{26}~\mathrm{cm}^3/\mathrm{s}$. This is a consequence of the novel parametric form of the DM yield (Eqs.~\ref{scalingcann}, \ref{scalingchem}, and \ref{scalingdecay}) and the requirement that Cannibal DM match the observed DM energy density. The parametric form of the boost to the cross section is summarized in Tab.~\ref{tab:boost} of Section~\ref{sec:GenCo}. The boost implies that Cannibal DM is easier to see through indirect detection than standard WIMPs. Observational constraints on cannibal DM, and future reach, are summarized in Figs.~\ref{figpheno1} and \ref{figpheno2}. Cannibal DM is constrained by Fermi measurements of $\gamma$-rays, Planck measurements of the Cosmic Microwave Background (CMB), and Lyman-$\alpha$ constraints on warm DM\@. However, significant parameter space remains allowed and we find that there is promising reach to discover cannibal DM through future CMB measurements. The rest of this paper is organized as follows. In Section~\ref{sec:Thermo}, we follow the evolution of a decoupled hidden sector as its temperature drops below the mass of the LDP, but when its number changing interactions are still active. In Section~\ref{sec:GenCo}, we introduce the example model that we use to describe the evolution of the DM density during the three phases. In Section~\ref{sec:pheno}, we study possible observational signatures and constraints on the model. Appendix~\ref{boltzmann} contains a complete description of the system of Boltzmann equations that we use to determine the DM energy density. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

1607.03108

1607

1607.06325_arXiv.txt

{Cosmological and astrophysical observations lead to the emerging picture of a universe that is spatially flat and presently undertaking an accelerated expansion. The observations supporting this picture come from a range of measurements encompassing estimates of galaxy cluster masses, the Hubble diagram derived from type-Ia supernovae observations, the measurements of Cosmic Microwave Background radiation anisotropies, etc. The present accelerated expansion of the universe can be explained by admitting the existence of a cosmic fluid, with negative pressure. In the simplest scenario this unknown component of the universe, the Dark Energy, is represented by the cosmological constant ($\Lambda$), and accounts for about 70\% of the global energy budget of the universe. The remaining 30\% consists of a small fraction of baryons (4\%) with the rest cold Dark Matter (CDM). The Lambda Cold Dark Matter ($\Lambda$CDM) model, i.e. General Relativity with cosmological constant, is in good agreement with observations. It can be assumed as the first step towards a new standard cosmological model. However, despite the satisfying agreement with observations, the $\Lambda$CDM model presents several lacks of congruence and shortcomings, and therefore theories beyond Einstein’s General Relativity are called for. Many extensions of Einstein's theory of gravity have been studied and proposed with various motivations like the quest for a quantum theory of gravity to extensions of anomalies in observations at the solar system, galactic and cosmological scales. These extensions include adding higher powers of Ricci curvature $R$, coupling the Ricci curvature with scalar fields and generalized functions of $R$. In addition when viewed from the perspective of Supergravity (SUGRA) many of these theories may originate from the same SUGRA theory, but interpreted in different frames. SUGRA therefore serves as a good framework for organizing and generalizing theories of gravity beyond General Relativity. All these theories when applied to inflation (a rapid expansion of early Universe in which primordial gravitational waves might be generated and might still be detectable by the imprint they left or by the ripples that persist today) can have distinct signatures in the Cosmic Microwave Background radiation temperature and polarization anisotropies. We give a review of $\Lambda$CDM cosmology and survey the theories of gravity beyond Einstein’s General Relativity, specially which arise from SUGRA, and study the consequences of these theories in the context of inflation and put bounds on the theories and the parameters therein from the observational experiments like PLANCK, Keck/BICEP etc. The possibility of testing these theories in the near future in CMB observations and new data coming from colliders like the LHC, provides an unique opportunity for constructing verifiable models of particle physics and General Relativity.}

\setcounter{equation}{0} The attempt to find a consistent quantum theory of gravity has motivated attempts to find generalizations of Einsteins General Theory of Relativity. The recent cosmological observations, in fact, suggest that the present Universe is in accelerating phase \cite{accUn}. To explain such a behavior one is forced to introduce the concept of Dark Energy. Moreover, for explaining the rotation curves of Galaxies one needs to introduce Dark Matter. The origin and nature of both Dark Energy (DE) and Dark Matter (DM), that undoubtedly represent a fundamental issue in particle physics, astrophysics and cosmology, is unknown and although many attempts have been done to explain such {\it dark components}, no final conclusions has been reached, and the question till now is completely open. Observations of the cosmic microwave background anisotropy motivate the idea that the universe in the past went through a period of accelerated expansion called inflation. The specific model of inflation which fits the temperature and polarization anisotropies is still not settled. It's not established yet whether inflation requires a fundamental scalar field in a particle physics model or whether the scalar degrees of freedom of the metric in modified gravity theories can give rise to viable inflation. In this review we survey the generailised theories of gravity which can give inflation compatible with observations of the CMB by experiments like Planck \cite{Ade:2015lrj,Planck:2015} and BICEP2/{\it Keck}+Planck~\cite{BKP:2015,bicep2keck2015,Ade:2015xua}. \subsubsection*{\it General features of metric theories of gravity} In order that a theory of gravity might be considered a valid theory, it must fulfill some minimal requirements: $1)$ It must reproduce the Newtonian dynamics in the weak-energy limit, which means that it has to pass the classical Solar System tests which are all experimentally well founded \cite{Will93}. $2)$ It must reproduce Galactic dynamics in relation to the observed baryonic constituents (hence luminous components as stars, or planets, dust and gas), radiation and Newtonian potential which is, by assumption, extrapolated to Galactic scales. $3)$ It must address the problem of large scale structure and the cosmological dynamics. General Relativity (GR) is the simplest theory that satisfies these requirements. It is based on the assumption that space and time are entangled into a single space-time structure, which must reduce to Minkowski's space-time structure in absence of gravitational forces. GR underlies also on Riemann's ideas according to which the Universe is a curved manifold \cite{riemann} and that the distribution of matter affects point by point the local curvature of the space-time structure. GR is strongly based on some assumptions that the Physics of Gravitation has to satisfy: \begin{quote} The {\it Principle of Relativity}: All frames are good frames for describing Physics. This implies that no preferred inertial frame should be chosen a priori (if any exists). \end{quote} \begin{quote} The {\it Principle of Equivalence}: Inertial effects are locally indistinguishable from gravitational effects (i.e. the equivalence between the inertial and the gravitational mass). \end{quote} \begin{quote} The {\it Principle of General Covariance}: Field equations are generally covariant. \end{quote} \begin{quote} The {\it Principle of Causality}: Each point of space-time should admit a universally valid notion of past, present and future. \end{quote} On the basis of the above principles, Einstein postulates that the gravitational forces must be related to the curvature of a metric tensor field $ds^2 = g_{\mu\nu}dx^{\mu}dx^{\nu}$ defined on a four-dimensional space-time manifold (the metric signature is the same of Minkowski's metric, $(-,+,+,+)$) and that space-time is curved, with the curvature locally determined by the distribution of the sources, the latter described by the energy-momentum tensor $T^{(m)}_{\mu\nu}$. An important achievement was to prove that the field equations for a metric tensor $g_{\mu\nu}$ can be obtained by starting from an action linearly depending on Ricci's scalar $R$ (Hilbert-Einstein action \cite{schroedinger}.) \begin{equation}\label{HEaction} S_{HE}=\frac{1}{2\kappa^2}\int d^4 x \sqrt{-g}\, R + S_{matter}\,, \qquad \kappa^2=8\pi G/c^4\,. \end{equation} The choice of Hilbert and Einstein was completely arbitrary, but it was certainly the simplest one. Some years later the GR formulation, it was clarified by Levi--Civita that curvature is not a purely metric notion. It is indeed related to the linear connection, which plays a central role in the definition of parallel transport and covariant derivation \cite{levicivita} (this is, in some sense, the precursor idea of what would be called a "gauge theoretical framework" \cite{gauge}, after the Cartan work \cite{cartan}). As later clarified, the principles of relativity, equivalence and covariance require that the space-time structure has to be determined by either one or both of two fields: a metric $g$ and a linear connection $\Gamma$ (symmetric in order that the theory is torsionless). The metric $g$ fixes the causal structure of space-time (the light cones) as well as some metric relations (clocks and rods); the connection $\Gamma$ fixes the free-fall, {\it i.e.} the locally inertial observers. They have to satisfy a number of compatibility relations, which generally lead to a specific form of the connection in terms of the metric tensor (Levi-Civita connections), but they can be also independent, leading to the Palatini approach of GR. \cite{palatiniorigin}. It is on this basis that the so-called alternative theories of gravitation, or {\it Extended Theories of Gravitation} (ETGs) arise. In fact their starting points is that gravitation is described by either a metric (purely metric theories), or by a linear connection (purely affine theories) or by both fields (metric-affine theories). Here the Lagrangian is a scalar density of the curvature invariants, constructed by mean of $\{g, \Gamma\}$. Attempts to generalize GR along these lines \cite{unification} and investigations about "alternative theories" continued even after 1960 \cite{Will93}. What arises from these studies is that the search for a coherent quantum theory of gravitation or the belief that gravity has to be considered as a sort of low-energy limit of string theories \cite{green}, has renew the idea that there is no reason to follow the simple prescription of Einstein and Hilbert. Other curvature invariants or non-linear functions of them should be also considered, especially in view of the fact that they have to be included in both the semi-classical expansion of a quantum Lagrangian or in the low-energy limit of a string Lagrangian. Not only from a mathematical point of view there is the necessity to generalize GR, but also current astrophysical and cosmological observations suggest that, as already pointed out, Einstein's equations are no longer a good test for gravitation at Galactic, extra-galactic and cosmic scales, unless one does not admit that the matter side of field equations contains some kind of exotic matter-energy which is the {\it dark matter} and {\it dark energy} side of the Universe. One can adopt a different point of view, in the sense that instead of changing the matter side of Einstein field equations to fit the missing matter-energy content of the currently observed Universe (by adding any sort of exotic matter and/or energy), one may change the gravitational side of the equations, admitting corrections coming from non-linearities in the effective Lagrangian. This is a possibility that needs to explored, even if, without a complete theory of gravity, one has to tune up the form of effective theory that is going to study, hence a huge family of allowed Lagrangians can be chosen, trying to fit all possible observational tests, at all scales (Solar, Galactic, extragalactic and cosmic, and so on). \subsection{Shortcomings in the standard cosmological model} We shortly review the shortcomings of the cosmological standard model (the cosmology based on GR) or hot Big Bang. The latter provides a framework for the description the evolution of the Universe. The spacetime evolution is governed by Einstein's field equation (which contains matter content of the Universe) \begin{equation}\label{GRfieldeq} R_{\mu\nu}-\frac{1}{2}g_{\mu\nu}R + \Lambda g_{\mu\nu}= {\kappa}^2 T_{\mu\nu}\,, \end{equation} This equation follows by varying the action (\ref{HEaction}) with respect to the metric $g_{\mu\nu}$, while \[ T_{\mu\nu}=-\frac{2}{\sqrt{-g}}\frac{\delta S_{matter}}{\delta g_{\mu\nu}} \] is the energy-momentum tensor. $\Lambda$ is the cosmological constant (it enters into (\ref{HEaction}) by replacing $R\to R+\Lambda$). There are three important epochs to which the early Universe undergoes: radiation, matter and vacuum (energy) domination eras. They are characterized by an appropriate relation between pressure and energy density of matter, $p=p(\rho)$. Assuming the homogeneity and isotropy of the (flat) cosmological background Friedman-Robertson-Walker (FRW) Universe, the line element is given by\footnote{The homogeneity, isotropy and expanding nature of the space-time is mathematically described by the more general FRW metric given by \be\label{FRWmetrica} ds^{2}=-dt^{2}+a^{2}(t)\left[\frac{dr^{2}}{1-\kappa r^{2}} + r^{2}\left(d\theta^{2}+sin^{2}\theta d\phi^{2}\right)\right], \ee where ($r, \theta, \phi$) are the comoving spatial coordinates. $\kappa$ represents the spatial curvature and can take values $+1, 0, -1$ describing open, flat and closed universe, respectively.} \begin{equation}\label{FRWmetric} ds^2=-dt^2+a^2(t) d{\bf x}^2\,, \end{equation} where $a(t)$ is the scale factor. For a long time, the success of the Big Bang model was based on three cornerstones: $i)$ Hubble expansion, $ii)$ the distribution of relic photons (CMBR), and $iii)$ the light element abundance such as ${}^3 He$, ${}^4 He$, $D$, ${}^7 Li$ (Big Bang Nucleosynthesis (BBN)). With the recent developments of modern Cosmology, our view of the Universe evolution has been completely transformed. For example, the cosmological model present some inconsistencies that can be solved only admitting the existence of the Inflation, i.e. a period of accelerated expansion of the Universe which occurred in the very early phase of the Universe evolution. It is believed that Inflation is responsible for the inhomogeneities in the matter distribution (whose evolution allowed the formation of structures, stars, planets) and the inhomogeneities of the CMB. These perturbations are generated by quantum fluctuations of the scalar field (the inflaton) which drives the Inflation, and they can be scalar, vectorial or tensorial. The challenge of modern cosmology is to verify all predictions of Inflationary scenario. The Inflationary paradigm allows to solve some inconsistence of the standard cosmological model: \begin{itemize} \item Flatness of the Universe - $\Omega=\rho/\rho_c=1$ where $\rho=3H^2m_P^2/8\pi$ is the critical density. Without Inflation, the adjustment of $\Omega$ should be $\sim 10^{-60}$ at the Planck era, and $\sim 10^{-15}$ at the primordial nucleosynthesis. \item The problems of homogeneity, isotropy, and horizon (which created headache in the frameworks of FRW cosmology) are elegantly solved. \item Inflation provides a natural mechanism of generation of {\it small density perturbations} with practically flat spectrum in agreement with observations. \item To solve all problems of standard FRW cosmology, the duration of Inflation must be $N=H t \sim 70-100$. \end{itemize} We will return on the standard cosmological model and Inflation in next Sections.

\setcounter{equation}{0} The idea, that the universe through a period of exponential expansion, called inflation, has proved useful for solving the horizon and flatness problems of standard cosmology and in addition providing an explanation for the scale invariant super-horizon perturbations which are responsible for generating the CMB anisotropies and formation of structures in the universe. A successful theory of inflation requires a flat potential where a scalar field acquires a slow-roll over a sufficiently long period to enable the universe to expand by at least $60$ e-foldings during the period of inflation. There are a wide variety of particle physics models which can provide the slow-roll scalar field 'inflaton' for inflation \cite{Lyth:1998xn,Martin:2013tda}. From the observations of CMB anisotropy spectrum by COBE, WMAP and Planck \cite{Smoot:1998jt,WMAP9,Planck:2015}, it is not yet possible to pin down a specific particle physics model as the one responsible for inflation. Though all of the above experiments gave tighter and tighter constraints on inflationary observables, $e.g.$ power spectrum and spectral index, which allowed several models of inflation to be ruled out but still there is a large degeneracy in inflation models. The 2015 data from Planck observation gives the amplitude, spectral index and tensor-to-scalar ratio as $10^{10}\ln(\Delta_{\mathcal R}^{2}) = 3.089\pm 0.036$, $n_{s}= 0.9666 \pm 0.0062$ at ($68 \%$ CL) and $r_{0.002}<0.11$ at ($95\%$CL), respectively~\cite{Ade:2015lrj}. Also the latest results by BKP collaboration put $r$ at $r_{0.05}<0.07$ at ($95\%$CL)\cite{bicep2keck2015}. Therefore, all those models which can produce the correct amplitude of CMB power spectrum and its spectral tilt along with producing $r<0.07$ are allowed. The future $B-$mode observations are expected to fix $r$ which will allow many existing inflation models to be ruled out. In this review we have presented some recent results of a single-field generalized non-minimal model of inflation, a power law model of inflation and a two-field model of inflation with a non-canonical kinetic term. We have calculated the key inflationary observables : amplitude of the power spectrum of curvature perturbations, spectral index and its running, tensor-to-scalar ratio and amplitude of isocurvature perturbations. We fix the parameters of the models using the measured values of these observables. Also we motivate these models from a fundamental theory called no-scale supergravity. From the analysis of these models we find that the power law model $R+R^{\beta}/M^{2}$ of inflation and therefore all the models of inflation with the Lagrangian $R + \xi \phi^{a} R^{b} + \lambda \phi^{4(1+\gamma)}$ are ruled form the current bound on $r_{0.05}<0.07$ at ($95\%$CL). However for $\beta=2$ which corresponds to Starobinsky model $R+R^{2}/M^{2}$ of inflation produces $r\simeq0.0033$. This prediction of $r\simeq0.0033$ however very small but is currently favored by the observations. Contrary to the large $r$ prediction of power law model, the two field model predicts a range of $r\sim 10^{-1}-10^{-2}$ values close the present bound on $r$. Currently tensor-to-scalar ratio is an extremely important inflationary parameter in view of validating and ruling out models of inflation, and therefore many models that can be classified as extended theories of gravity. The observation of primordial $B$-modes (CMB polarization) will provide the constraint on $r$. The $B$-modes are the signature of inflationary tensor modes, precise observations of which will provide a most distinctive confirmation of occurrence of an inflationary era in the early universe. There are several ongoing experiments for $B-$mode detection and all hope to observe these signals from the inflationary era. There are several experiments, $e.g.$ ground based (Keck/BICEP3, SPT-3G, AdvACT, CLASS, Simons Array), balloons based (Spider,EBEX) and satellites based (CMBPol, LiteBIRD and COrE). These observational experiments will be taking into account the recent Planck data on polarized dust. They aim to probe the tensor-to-scale ratio at the level of $r\sim10^{-3}$ which is a theoretically motivated limit~\cite{Creminelli:2015oda}. High precision measurements of small-scale temperature anisotropies along with the observations of $B-$mode will not only test the inflationary hypothesis but allow to remove a large degeneracy in the models of inflation. \begin{acknowledgement} G.L. thanks the project COST project CA15117 CANTATA. \end{acknowledgement}

1607.06325

1607

1607.08612_arXiv.txt

We present a large survey of giant pulses from the Crab Pulsar as observed with the first station of the Long Wavelength Array. Automated methods for detecting giant pulses at low frequencies where scattering becomes prevalent are also explored. More than 1400 pulses were detected across four frequency bands between 20 - 84 MHz over a seven month period beginning in 2013, with additional followup observations in late 2014 and early 2015. A handful of these pulses were detected simultaneously across all four frequency bands. We examine pulse characteristics, including pulse broadening and power law indices for amplitude distributions. We find that the flux density increases toward shorter wavelengths, consistent with a spectral turnover at 100 MHz. Our observations uniquely span multiple scattering epochs, manifesting as a notable trend in the number of detections per observation. These results are characteristic of the variable interface between the synchrotron nebula and the surrounding interstellar medium.

The average emission profile of the Crab Pulsar exhibits occasional bursts of increased intensity commonly referred to as giant pulses \citep{Staelin1968}. These individual pulses can exceed the average flux density by several orders of magnitude, becoming one of the brightest radio sources in the sky \citep{Jessner2005}. The short duration and high brightness temperatures of these bursts are indicative of non-thermal, coherent emission \citep{Hankins2003}. Nevertheless, the exact mechanisms responsible for these spurious bursts remain elusive. The majority of studies to date have focused on high time resolution radio observations at several GHz, where the effects of dispersion and scattering do not degrade the intrinsic nano to microsecond time resolution of detected pulses. More recently, several campaigns have emerged to study giant pulses and the effects of multi-path propagation at lower frequencies (e.g. \citet{Oronsaye2015}, at 193 MHz; \citet{Karuppusamy2012}, at 110-180 MHz; \citet{Bhat2007}, at 200 MHz). In this paper, we present results from a low frequency survey of giant pulses from the Crab Pulsar as observed with the first station of the Long Wavelength Array (LWA1). In \S 2, we introduce the observations, followed by a discussion of pulse shapes at low frequencies in \S 3. The data reduction scheme and flux density calibration are discussed in \S 4 and \S 5. Finally, in \S 6, we examine pulse characteristics and discuss the implications of our results.

Over 1400 giant pulses from the Crab Pulsar have been detected with the LWA1 radio telescope in 73 hours of observations, compared to 33 pulses detected in 10 hours of observations as in \citet{Ellingson2013}. This corresponds to an approximate increase in the rate of detection by a factor of six. These differences may stem from improved calibration of the LWA1 cable delays since the initial giant pulse investigation. Additionally, results presented in \citet{Ellingson2013} seem to suggest an altogether separate scattering epoch (see \S 6.2). Finally, detection methods presented here differ from those implemented in \citet{Ellingson2013}, in which pulses were initially identified by eye. The use of pulse-matched filtering in the present work likely resulted in detections that would have gone otherwise unnoticed. Our observations uniquely bracket a scattering epoch, given by the anticorrelation between average scattering timescales and the number of observable pulses. These fluctuating timescales represent the variable nature of the surrounding nebula, and in particular, provide an interesting probe of the nebula-ISM interface. A positive spectral index is obtained for those pulses observed concurrently in all four passbands. These results are consistent with a supposed spectral turnover at 100 MHz and indicate that giant pulse detection below 50 MHz becomes increasingly more difficult, relying on the brightest pulses which populate the tail end of the distribution. Continued observations of pulsars with the LWA1 are currently ongoing (see \citet{Stovall2015}). Low frequency studies of pulsars are particularly well-suited for characterizing the effects of multi-path propagation through the interstellar medium. Such studies -- when combined with simultaneous observations at higher frequencies -- will allow for careful analysis of pulse morphologies across a range of frequencies, providing further constraints on the mechanisms responsible for pulsar emission. In particular, simultaneous observations of individual pulses spanning frequencies above and below the 100 MHz spectral turnover will be particularly useful in characterizing the complex nature of giant pulse emission.

1607.08612

1607

1607.08424_arXiv.txt

{Using a photonic reformatter to eliminate the effects of conventional modal noise could greatly improve the stability of a high resolution spectrograph. However the regimes where this advantage becomes clear are not yet defined. Here we will look at where modal noise becomes a problem in conventional high resolution spectroscopy and what impact photonic spectrographs could have. We will theoretically derive achievable radial velocity measurements to compare photonic instruments and conventional ones. We will discuss the theoretical and experimental investigations that will need to be undertaken to optimize and prove the photonic reformatting concept.}

High resolution spectroscopy is one of the pillars of modern astronomy. When applied to stars it allows the derivation of temperature, stellar composition and velocity. It is also is a fundamental technique for detecting exoplanets. However, as astronomers push for greater stability and resolution, the results are limited by new factors. Not only is the light from the star dispersed by larger amounts, meaning only brighter targets can be observed, but new noise limits are set by the instrument and star. Knowing the velocities of stars to the order of cm/s requires large telescopes, this means the spectrographs behind them have to grow in size, becoming more costly to manufacture. This also means controlling them to high precision becomes more difficult. To keep the possibility of mechanical flexure to a minimum, temperatures and pressures are highly controlled (e.g.\ HARPS is maintained at a pressure \textless $10^{-2}$ mbar and a temperature of 17 \degree C, constant within 0.005 \degree C RMS\cite{2003Msngr.114...20M}) and this will only become more stringent in the next generation of instruments. Once the mechanical controls are in place, calibration requires devices like Iodine cells and Thorium Argon lamps, laser frequency combs \cite{steinmetz2008laser} and Fabry-Perot interferometers \cite{wildi2010fabry,schwab2015stabilizing}. Once the above elements are controlled other problems arise, such as the stability of the input fibers to these spectrographs. It has been shown that conventional multimode fibers are subject to modal noise, \cite{goodman1981statistics,lemke2011modal} due to the wave nature of light producing modes within the fiber. If this is not properly removed, the resulting measured barycentre for a given wavelength will be altered (note this is different to incomplete scrambling \cite{mccoy2012optical}), changing the resultant radial velocity measurement. Modern high resolution spectrographs use fiber scramblers, different core geometries and fiber agitators, to attempt to remove modal noise though not always completely. Modal noise also becomes worse at longer wavelengths, where scrambling becomes harder due to the inverse dependence of number of modes with wavelength. The field of Astrophotonics \cite{Bland-Hawthorn2009} provides two potential solutions to the problem of modal noise. The first use a photonic lantern to seperate the modes and use agitation to effectively 'mix' them \cite{haynes2014new}. The second is to reformat the multimode fiber into lots of individual single modes, which can be dealt with separately. This can be done in many ways, \cite{harris2014comparison} such as \ac{PIMMS} \cite{Bland-Hawthorn2010}, TIGER \cite{2012arXiv1208.3006L} the photonic dicer \cite{2015MNRAS.450..428H}. Reformatting into a long slit, such as with the photonic dicer should in practise eliminate any barycentre shift (in one axis), however if the slit is imperfectly manufactured it could cause effects that have yet to be quantified and it is not yet shown whether using single modes may cause other problems (e.g. Ref. \citenum{2015ApJ...814L..22H}). Reformatting, also comes at a price (unless the input from the telescope is diffraction limited). All these extra modes must pass through the spectrograph and be sampled by the detector in order to retain throughput, increasing the number of pixels required for the spectrograph \cite{Harris2012,betters2013beating}. This can lead to uncompetitive sizes and costs in certain types of instruments \cite{2013MNRAS.428.3139H}. Because of this, science cases requiring a small telescope, \ac{AO}, space based instrumentation, a long wavelength and small field of view (few spaxels) \cite{Harris2012} are desirable. The next generation of high resolution spectrographs fit well into this category. This paper will look at three existing high resolution spectrographs and identify if and where modal noise is a limiting factor. It will compare this to models of photonic spectrographs and identify where they can provide a competitive edge.

We have used toy models to investigate where modal noise could be a problem in high resolution spectroscopy. To do this we have compared conventional instrumentation to models of photonic spectrographs for HIRES, CARMENES and GIANO. As with other papers in the field we have concluded that modal noise is most problematic at longer wavelengths. For all directly comparable models we found that photonic spectrographs would achieve lower signal to noise for a given exposure time due to the required number of modes (and hence pixels to sample them). This leads to a lower radial velocity precision when only using at photon statistics to calculate signal to noise (this is inline with previous results). In cases where we reduced the sampling of the slit we found similar or in some cases higher signal to noise for photonic spectrographs. This should be taken with caution as we have assumed that the slit is uniform and no aberrations are introduced by the spectrograph. When the limit due to modal noise is taken into account we found that HIRES and CARMENES should not be hugely affected (provided that the fiber is scrambled, which is the case for both), but that it was the limiting factor for GIANO (again consistent with results). If a photonic spectrograph can be shown to reduce or eliminate modal noise it would be very useful in this case. Our toy model does have limitations and further work is needed. Summarized briefly: \begin{enumerate} \item While initial results are promising it still needs to be shown that photonic reformatters are modal noise free. This will need to be be explored both theoretically and experimentally. \item Even if the photonic reformatter can be shown to be modal noise free, other factors like polarization response (e.g Ref. \citenum{2015ApJ...814L..22H}), may need to be taken into account. Investigations should be conducted to see how the response of the reformatter compares with conventional spectrographs. \item Here we have assumed that each mode along the slit needs to be sampled by two pixels in the spatial direction. There have been suggestions before of reducing this by using cylindrical lenses or rectangular pixels. This should be investigated. \end{enumerate} Finally additional spectrographs can be modelled to see how they respond. To this end the code is freely available. It can be found at https://zenodo.org/record/54765.

1607.08424

1607

1607.03446_arXiv.txt

{Stellar populations are the building blocks of galaxies including the Milky Way. The majority, if not all extragalactic studies are entangled with the use of stellar population models given the unresolved nature of their observation. Extragalactic systems contain multiple stellar populations with complex star formation histories. However, their study is {mainly} based upon the principles of simple stellar populations (SSP). Hence, it is critical to examine the validity of SSP models. } {This work aims to empirically test the validity of SSP models. This is done by comparing SSP models against observations of spatially resolved young stellar population in the determination of its physical properties, i.e. age and metallicity.} {Integral field spectroscopy of a young stellar cluster in the Milky Way, NGC 3603, is used to study the properties of the cluster both as a resolved and unresolved stellar population. {The unresolved stellar population is analysed using the Ha equivalent width as an age indicator, and the ratio of strong emission lines to infer metallicity. In addition, spectral energy distribution (SED) fitting using STARLIGHT, is used to infer these properties from the {integrated spectrum}.} Independently, the resolved stellar population is analysed using the color-magnitude diagram (CMD) for age and metallicity determination. As the SSP model represents the unresolved stellar population, the derived age and metallicity are put to test whether they agree with those derived from resolved stars. } {The age and metallicity estimate of NGC 3603 derived from integrated spectroscopy are confirmed to be within the range of those derived from the CMD of the resolved stellar population, including other estimates found in the literature. The result from this pilot study supports the reliability of SSP models for studying {unresolved} young stellar populations. } {}

Stars together form stellar populations, which are the fundamental building blocks of galaxies in the Universe. A simple stellar population (SSP) is usually defined as a group of stars distributed following an initial mass function (IMF), that was formed from a single cloud of gas with homogeneous chemical composition at the same time in one burst of star formation. As a result, stars within an SSP possess the same age and metallicity. Star clusters have been regarded as nature's closest approximation to the ideal SSP \citep[c.f.][]{bruzual10}. Studies of external galaxies rely heavily on the SSP models. SSP models such as BC03 \citep{bc03}, Starburst99 \citep{leitherer99} and GALEV \citep{kotulla09}, among many others, are widely used as probes of unresolved stellar populations in galaxies. SSP models help interpreting the observed spectral energy distribution (SED) into meaningful physical quantities such as age, metallicity, luminosity, stellar content, and star formation rate. {Correct determination of such properties is important for the study of galaxies itself as well as for other purposes.} {For example}, in the studies of the hosts and environments of extragalactic transient objects such as supernovae (SNe) and gamma-ray bursts (GRBs), the properties of their stellar progenitors are derived from those of the underlying stellar populations \citep[see e.g.][]{galbany16a,galbany16b,galbany14,leloudas15,leloudas11,kruehler15,anderson15,hk13a,hk13b,levesque10}. On that account, correct interpretation of the stellar population associated with the transient is necessary. While the validity of SSP models is of utmost importance, there are questions concerning the consistency between different SSP models and whether those are all well-{tested}. Different SSP models are generated using various different methods, thus the stellar population physical properties derived using one SSP model family may not necessarily be consistent with the result from another. {The main ingredients of SSP models usually consist of isochrones, a stellar spectral library, and an IMF \citep[see][for a review]{conroy13}. The choice of the isochrones and spectral libraries may be different from one model family to the other, and there can be variations at the more detailed levels such as the stellar evolution calculations for the isochrones and the selection of empirical or theoretical stellar spectra for the library, among others. These differences may eventually affect the constructed SSP models.} Figure~\ref{haewcomp} shows the evolution of H$\alpha$ emission line equivalent width (EW) as an age indicator, according to three different SSP models families {(GALEV, Starburst99, and BPASS). These SSP models cover very young age ($\sim$Myr), and are thus suitable for the analysis presented here.} It is apparent that despite the roughly similar behaviour in evolution, the H$\alpha$EW values may differ by a factor of two or more at a given age. {To test the validity of SSP models in the analysis of complex extragalactic systems, a two-step procedure is required. First, the SSP models need to be compared to stellar clusters. Second, more complex systems are analysed using combinations of SSP models with non-instantaneous star formation, as one SSP model is considered too simplistic to represent such systems. In the current work, we aim to perform the first step in such a validity test. This first step serves as the basis for the subsequently more intricate step, which eventually will help in assessing whether the SSP-based techniques of extracting stellar population information in extragalactic systems are robust or not.} Attempts to {test SSP models against stellar clusters} have been undertaken previously, although only for older stellar populations where the massive stars and ionized gas are no longer present \citep[see e.g.][]{renzini88}. \citet{beasley02} collected integrated spectra of 24 Large Magellanic Cloud star clusters and compared the spectroscopically-derived age and metallicity with those from the literature. In general, the age and metallicity derived using Lick/IDS spectral indices show consistency with the literature values, which were mostly derived from the resolved Color-Magnitude Diagram (CMD) or integrated colour for age and Ca-triplet spectroscopy for metallicity. Other efforts along the same line of study have also been done by \citet{koleva08} and \citet{riffel11}, who concluded that SSP models quite satisfactorily reproduce the observed parameters albeit there were minor caveats. {In the context of studying more complex objects with non-instantaneous star formation history,} more recently \citet{ruiz15} presented a comparison of resolved and unresolved stellar populations within a field in the Large Magellanic Cloud in a study of the star formation history in the region. On the other hand, SSP models corresponding to young stellar populations ($\sim$Myr age) are less well-established compared to the older population ($\sim$Gyr) grids {\citep{conroy13}}. {Many SSP models do not include young grids or take into account nebular emissions coming from young stellar populations, while they commonly provide grids for old populations.} These young stellar populations are dynamic objects associated with active star-forming regions where the interplay between the ionizing massive stars and the surrounding gas is intense, and energetic events such as core-collapse SNe and long-GRBs occur {\citep[see e.g.][]{portegies10}}. Even in a galaxy preeminently consisting of old stellar populations, the young populations may dominate the emissions in the ultraviolet/optical regime thus affecting significantly the observed integrated properties {\citep{meurer95,anders03}}. In this paper we report our pilot study {examining the reliability of} young SSP models, by using VLT/MUSE integral field spectroscopy (IFS) of a well-studied young stellar population, NGC 3603. It has been suggested that the cluster is as young as a few Myr \citep[e.g.][]{melnick89,sagar01} and was formed in an instantaneous burst of star formation \citep[][]{kudryavtseva12,fukui14}, {and} thus represents a good approximation for a young SSP. As young stellar populations are still associated with ionized gas, IFS offers the most efficient way to obtain the integrated spectrum of the cluster, including the ionized gas, in addition to the individual spectra of the member stars. {With this technique, it is possible to obtain simultaneously the integrated and resolved information of a stellar population from a single dataset.} This would have taken enormous time to conduct with conventional slit or fiber spectroscopy, and to our knowledge this is the first of such efforts in {testing} young ($\sim$Myr) SSP models. The paper is organized as follows. Following the introduction, the observations and data reduction are described in Section 2, then the analysis of the integrated spectrum for age and metallicity estimates is presented in Section 3. In Section 4 we analyse the properties of the resolved stars, and compare these to the unresolved stellar population from the integrated light. Section 5 provides a final summary. \begin{figure} \centering \includegraphics[width=\hsize]{fig_compssp.eps} \caption{Comparison of H$\alpha$EW age indicator of BPASS {single-star} \citep{eldridge09}, GALEV \citep{kotulla09}, and Starburst99 \citep{leitherer99} {SSPs}, illustrating the differences between various SSP models. BPASS and GALEV H$\alpha$EWs were measured directly from the SSP model spectra (see section \ref{sec:emission} for the measurement method), while in Starburst99 the values are already given in tabulated form. All SSP models assume instantaneous star formation and a standard Salpeter IMF. {As \haew$ $ also depends on the metallicity of the SSP model, two different metallicities are plotted to illustrate this effect. Solid lines indicate solar metallicity ($Z$ = 0.02) models and dashed lines indicate models with $Z$ = 0.004.} } \label{haewcomp} \end{figure}

We have obtained integral field spectroscopy of NGC 3603 using VLT/MUSE, from which an analysis to determine age and metallicity was performed with the aim of {testing the validity of} SSP models. {The integrated spectrum of NGC 3603 was analysed by utilizing the strong-line method, together with SSP model comparisons to both the \haew, and the overall SED. The resulting age and metallicity were then compared to those derived from an independent method utilizing the color-magnitude diagram of resolved stellar population {from this very dataset}, and also to other age and metallicity estimates obtained from the literature.} From the integrated spectrum the age of NGC 3603 was derived as around {4-6} Myr and with metallicity between $Z$ = 0.004 and $Z$ = 0.008. {The uncertainty in the age estimate originates from the different results when comparing the observed \haew$ $ to the different SSP models (BPASS/GALEV/Starburst99). In the case of SED fitting, using different SSP models similarly alters the age estimate although this also results in large variations in metallicity and fraction of contamination from older populations. Overall this illustrates the notion that using different SSP models may yield different results in the interpretation of stellar population properties. } {The 4-6 Myr age estimate from the integrated spectrum} agrees reasonably well with the estimates from photometry ({$\sim$ 1-5 Myr}), and suggests the consistency between the two independent methods. This pathfinder study provides a promising start towards a more systematic and comprehensive investigation in the effort of {testing} SSP models for young stellar populations. In this context, IFS offers a unique way to perform such investigation due to the extended nature of the ionized gas component in young stellar populations. {MUSE as the foremost IFS instrument currently available offers the exciting capability of exploring both the stellar content and nebular component of young star clusters simultaneously.} {For systems located close enough as to have a good spatial resolution, MUSE allows both resolved stellar population and integrated population analyses.} Wide-field IFS observations of young stellar clusters with a significantly larger sample are desired to further advance this study{, opening the possibility of a new field of analysis.} {We encourage the community to undertake efforts such as the one presented here in comparing resolved and unresolved stellar populations. This would help in refining SSP analysis techniques, which is ultimately vital for our understanding of extragalactic systems.}

1607.03446

1607

1607.07445_arXiv.txt

Galaxies fall into two clearly distinct types: `blue-sequence' galaxies that are rapidly forming young stars, and `red-sequence' galaxies in which star formation has almost completely ceased. Most galaxies more massive than $3\times10^{10}\Msol$ follow the red-sequence while less massive central galaxies lie on the blue sequence. We show that these sequences are created by a competition between star formation-driven outflows and gas accretion on to the supermassive black hole at the galaxy's centre. We develop a simple analytic model for this interaction. In galaxies less massive than $3\times10^{10}\Msol$ young stars and supernovae drive a high entropy outflow which is more buoyant than any tenuous corona. The outflow balances the rate of gas inflow, preventing high gas densities building up in the central regions. More massive galaxies, however, are surrounded by an increasingly hot corona. Above a halo mass of $\sim 10^{12}\Msol$, the outflow ceases to be buoyant and star formation is unable to prevent the build up of gas in the central regions. This triggers a strongly non-linear response from the black hole. Its accretion rate rises rapidly, heating the galaxy's corona, disrupting the incoming supply of cool gas and starving the galaxy of the fuel for star formation. The host galaxy makes a transition to the red sequence, and further growth predominantly occurs through galaxy mergers. We show that the analytic model provides a good description of galaxy evolution in the EAGLE hydrodynamic simulations. So long as star formation-driven outflows are present, the transition mass scale is almost independent of subgrid parameter choice.

Galaxies fall into two clearly distinct types: active, `blue-sequence' galaxies that are rapidly forming young stars, and passive `red-sequence' galaxies in which star formation has almost completely ceased. The two sequences are clearly seen when galaxy colours or star formation rates are plotted as a function of galaxy mass (eg., \citet{kauffmann2003, baldry2006}). Low-mass galaxies generally follow the `blue-sequence' with a tight, almost linear, relationship between star formation rate and stellar mass \citep[eg.,][]{brinchmann2004}, while massive galaxies follow the `red-sequence' with almost undetectable levels of star formation \citep[eg.,][]{bower1992}. % At the present day, the transition between the two types occurs at a stellar mass scale of $3\times10^{10}\Msol$. Galaxies less massive than the transition-scale grow through star formation, doubling their stellar mass on a timescale comparable to the age of the Universe; above the transition mass, galaxy growth slows and is driven primarily by galaxy mergers \citep[eg.,][]{delucia2006,parry2009,qu2016,rodriguez2016}. The existence of the transition mass is closely related to the form of the galaxy stellar mass function, creating the exponential break at high masses \citep[eg.,][]{benson2003, peng2010}. The transition mass is sometimes referred to as the `quenching' mass scale. In this paper we will focus on the properties of central galaxies (we will not consider the environmental effects that may suppress star formation in lower mass satellite galaxies, see \citet{trayford2016}) and show that the transition mass scale arises from a competition between star formation driven outflows and black hole accretion. Most of the stars in the Universe today were, however, formed when the Universe was less than half its present age. Recently, deep redshift surveys have been able to convincingly demonstrate the existence of a transition mass at higher redshifts \citep[eg.,][]{peng2010,ilbert2015,darvish2016}. It is useful to illustrate the balance of galaxies on the two sequences by revisiting the analysis of \citet{kauffmann2003}. This is illustrated in Fig.~\ref{fig:sfgrowth_mstar}, where contours show the star formation rate growth timescale of galaxies at $z=1$ using observational data from the COSMOS high-redshift galaxy survey \citep{ilbert2015}. So that passive galaxies appear in the figure, we assign galaxies without detectable star formation a growth timescale of $20\Gyr$ with a scatter of 0.2 dex. The relative contributions of blue-and red-sequence galaxies are then computed using the luminosity functions presented by \citet{ilbert2013}. The figure clearly shows that the division into sequences seen in present-day galaxies was already established 8 billion years ago and that the transition mass scale of $3\times10^{10}\Msol$ has changed little over the intervening time. \begin{figure*} \centering \includegraphics[width=\linewidth]{figures_july20/REFERENCE/bh_plots_sfgrowth_mstar.png} \caption{ The formation timescales of galaxies, $M_*/\dot{M}_*$, as a function of stellar mass. Contours show observational data for galaxies at $z=1$ \citep{ilbert2015}, using the revised absolute star formation rate calibration of \citet{chang2015}. The separation of galaxies into blue (rapidly star forming) and red (passive) sequences is clearly seen. Most low-mass galaxies follow the star forming blue galaxy sequence, doubling their stellar mass every 3 billion years, but more massive galaxies have much longer star formation growth timescales. The horizontal dotted line shows the present day age of the Universe; galaxies with longer star formation timescales are randomly placed above the line. The transition between the sequences occurs at a stellar mass of round $3\times10^{10}\Msol$, similar to the transition mass scale observed in the present-day Universe. We supplement the observational data with model galaxies from the reference EAGLE cosmological simulation (filled circles). The simulated galaxies follow the observed data closely. The points are coloured by the mass of the black hole relative to that of the host halo. Around the transition mass scale, there is considerable scatter in the relative mass of the black hole and in the star formation growth timescale. However, at a fixed galaxy mass, systems with a higher black hole mass tend to have substantially longer growth timescales, implying the existence of an intimate connection between black hole mass and galaxy star formation activity. } \label{fig:sfgrowth_mstar} \end{figure*} A simple way to begin to understand galaxy formation is to view galaxies as equilibrium systems in which the star formation rate must balance the gas inflow rate, either by converting the inflowing gas into stars or, more importantly, by driving an outflow \citep[eg.,][]{white_frenk1991,finlator2008,schaye2010,bouche2010,dave2012}. Such a model broadly explains many aspects of galaxy evolution, such as the almost linear correlation between stellar mass and star formation rate, and the rate of evolution of this sequence. In order to explain the flat faint-end slope of the galaxy mass function, such models require that the mass loading of the outflow depends strongly on galaxy mass. Low-mass galaxies lose most of their mass in the outflow (and form few stars) while galaxies around the transition mass scale consume much of the inflowing mass in star formation. Such models do not generally consider the role of the nuclear supermassive black hole, however. Yet the energy liberated when one solar mass of gas is accreted by a black hole is 10,000 times the supernova energy released by forming the same mass of stars. Observational measurements of energetic outflows and radio jets from galaxy nuclei \citep[eg.,][]{harrison2012, maiolino2012, fabian2003} indeed suggest that black holes play an important role in galaxy formation, most likely by heating the surrounding gas corona, offsetting cooling losses and disrupting the gas inflow \citep[eg.,][]{binney_tabor1995, silk_rees1998, dubois2013}. Such observations have motivated the inclusion of black hole feedback in cosmological galaxy formation models \citep[eg.,][]{dimatteo2005,croton2006,bower2006,sijacki2007,booth2009,dubois2013,sijacki2015,dubois2016}. The success of these models results from including two different modes of black hole feedback or accretion depending on halo mass or Eddington accretion ratio. In the most extreme implementation, the effect of black hole feedback is implemented by switching cooling off in massive haloes \citep{gabor2015}. There are three possible arguments that may be advanced to support a halo mass or accretion rate dependence: (i) (in the absence of disk instabilities) black holes are only able to accrete efficiently if the surrounding gas is hot and pressure supported \citep{croton2006}; (ii) black hole feedback is only effective if the black hole is surrounded by a hot halo that is able to capture the energy of the black hole jet \citep{bower2006}; (iii) only black holes accreting well below the Eddington limit (perhaps in the ADAF regime) are able to generate mechanical outflows \citep{meier1999,nemmen2007}. In the latter case, a halo mass dependence of the feedback efficiency emerges because of the strong dependence of black hole mass on halo mass. In each case, however, the models provide only a qualitative explanation of the transition mass scale. The first case assumes angular momentum prevents accretion of cold gas; the second argument does not explain why black holes do not undergo run-away growth in low-mass haloes; the third links galaxy properties to the highly uncertain AU-scale physics of black hole accretion disks. In contrast to models that include an explicit mass or accretion rate dependence, the MassiveBlackII \citep{khandai2015} and EAGLE \citep{schaye2015} simulations adopt a simpler description in which AGN feedback power is a fixed fraction of the rest mass accretion rate. Despite this simple proportionality, the EAGLE simulations reproduce the observed galaxy transition mass scale well. For example, the star formation growth time scales of EAGLE galaxies, shown by the coloured points in Fig~\ref{fig:sfgrowth_mstar}, match the observation data well (for further discussion, see \citet{furlong2015}). Although a transition mass is not so clearly evident in the MassiveBlackII simulations, the results are broadly consistent with a variation of the EAGLE simulation in which star formation driven outflow are weak (see \S\ref{sec:parameter_variations}). The success of the simple black hole feedback scheme in EAGLE suggests that the transition mass scale emerges from the interaction between star formation feedback and black hole fueling. This has motivated us to explore a scenario in which star formation-driven outflows themselves regulate the density of gas reaching the black hole. This has previously been considered in the context of high resolution simulations of individual high redshift galaxies \citep[eg.,][]{dubois2015, habouzit2016}. In low-mass galaxies, such outflows are efficient, making it difficult for high gas densities to build up in the central regions of the galaxy. The critical and novel component of our model is the recognition that outflows from more massive galaxies are not buoyant. As the mass of the galaxy's dark matter halo increases, a hot corona begins to form \citep[eg.,][]{white_frenk1991, birnboim2003} and the outflow becomes trapped. Star formation can then no longer prevent the gas density increasing in the central regions of the galaxy, triggering a strongly non-linear response from the black hole, heating the corona and disrupting the inflow of fresh fuel for star formation. In this picture, the falling effectiveness of feedback from star formation leads to an increase in accretion on to the black hole. Black hole feedback will take over just as star formation driven feedback fails. The structure of the paper is as follows. In \S\ref{sec:analytic_model}, we develop a simple analytic model that captures the key physics of the problem. We begin by exploring the strongly non-linear growth rates of black holes accreting from a constant density medium. We show that in the Bondi accretion regime black holes initially grow slowly and then abruptly switch to a rapid accretion phase. In the absence of feedback, the black hole grows to infinite mass in a finite time. We go on to consider how the density in the central parts of the galaxy will evolve as the halo mass grows, and develop a model in which the gas density is determined by the buoyancy of the star formation driven outflow. We show that a critical mass scale emerges: below this scale galaxies balance gas inflow with star formation driven outflows, but above this mass the galaxy and its gas corona are regulated by black hole accretion and star formation is strongly suppressed. We use this simple model to explore the growth of black holes in a cosmological context. In \S\ref{sec:validation}, we validate the analytic model by comparing it to galaxies forming in the EAGLE hydrodynamic simulation suite. We show that the approximations of the analytic model are well supported, and explore the effects of varying the sub-grid parameters of the simulations. We show that the galaxy transition mass scale is robust to the choice of subgrid parameters, but that it disappears entirely if star formation driven outflow are absent. We present a summary of the results and a discussion of the model's wider implications in \S\ref{sec:discussion}.

\label{sec:discussion} Galaxy properties show an abrupt change in behaviour around a halo mass of $10^{12}\Msol$. In lower-mass haloes, central galaxies are rapidly star forming, doubling their stellar mass through star formation on a timescale that is shorter than the age of the Universe. In higher-mass haloes, galaxies have much longer star formation growth timescales (Fig.~\ref{fig:sfgrowth_mstar}) and grow primarily through galaxy mergers and the accretion of external stars. This dichotomy is often referred to as the blue and red galaxy sequences, respectively. The change in star formation rate at the transition mass scale leads to a change in slope of the galaxy mass -- halo mass relation, creating a break in the galaxy mass function at a galaxy mass scale of $3\times10^{10}\Msol$. The decline in the importance of on-going star formation and the rise in the importance of growth by mergers may also account for the shift from late to early-type galaxy morphology at this mass scale. Explaining the origin of this mass scale is of fundamental importance to explaining the structure of the observable Universe. In this paper, we have developed a simple analytic model in which black holes and star formation compete to regulate the gas content of a galaxy as it grows by accretion from the cosmic web. The main elements of the model are the non-linear growth of black holes in the Bondi accretion regime and the buoyancy of star formation - driven outflow relative to galaxies hot coronae. In the Bondi regime, black holes grow slowly before abruptly switching to a rapid growth phase (Fig.~\ref{fig:bh_growth_model}). The onset of this phase is highly dependent on the surrounding gas density, and thus on the ability of star formation driven outflows to prevent gas accumulating in the centre of the galaxy. In our model, the effectiveness of stellar feedback depends on the buoyancy of the outflow relative to the corona (in the absence of a corona, the outflow escapes as a rarefraction wave). We show that this leads to a critical halo mass above which star formation driven outflows are unable to prevent the build up of gas. Making simple assumptions about the evolution of the gas density of star forming galaxies, we compare the buoyancy of the star formation driven outflow to that of the halo (Fig.~\ref{fig:halo_density_temperature}) in order to determine the dependence of the central gas density on halo mass and redshift. The model links the build up of the gas density in the central regions of the galaxy, the on-set of rapid black hole growth, the galaxy's eventual transition to the red sequence and the build-up of the hot corona. In order to test the simplified analytic model, we compare the model to the growth of black holes in the EAGLE hydrodynamic simulations. The EAGLE simulations assume a constant efficiency of black hole feedback, but at the same time obtain a remarkable match to the observed properties of galaxies, including the abrupt change in galaxy star formation rates in the transition regime (Fig.~\ref{fig:sfgrowth_mstar}). Although no halo mass dependence of black hole accretion rates is imposed in the simulations, the transition in galaxy properties emerges as we would expect from the analytic model (Fig.~\ref{fig:model}). The greatest value of the hydrodynamic simulations is, however, that we can vary the parameters that control feedback from star formation and the accretion rates of black holes, and thus experiment with their impact on the resulting galaxy and black hole correlations. This enables us to confirm the causal connections implied by the simple analytic model. For a fixed density, varying the implicit accretion disk viscosity, or the seed black hole mass, alters $\kappa$, and hence the onset of the rapid growth phase, $t_{\infty}$. In practice, however, reductions in $\kappa$ are compensated by an increasingly steep relation between the gas density around the black hole and halo mass, and the galaxy transition mass is remarkably insensitive to the parameter choice (Fig.~\ref{fig:bh_model_variations}). Variations in the effective equation of state assumed for star-forming gas or the heating temperature of AGN feedback (Fig.~\ref{fig:more_model_variations}) also have little impact on the transition mass scale. If we reduce or eliminate feedback from supernovae, however, we find that black holes are able to grow effectively in haloes of all masses, as illustrated by the dot-dashed line in Fig. ~\ref{fig:model} and the cyan line in Fig.~\ref{fig:more_model_variations}. This change eliminates the transition mass scale and consequently does not reproduce the properties of observed galaxies. This confirms that, below the transition mass, black hole growth is suppressed by stellar feedback. Galaxies fall into two distinct galaxy types, characterised by the presence or absence of significant on-going star formation. This dichotomy is established in the early universe and driven by a transition in galaxy properties at a mass of $\sim 3\times 10^{10}\Msol$. We have presented a new way of understanding this transition and, thus, the origin of the distinct red and blue galaxy sequences. Rapid growth of black holes is triggered at a halo mass of $\sim 10^{12}\Msol$ by the development of a sufficiently hot diffuse gas corona which confines the star formation-driven outflow by preventing it from rising buoyantly. This simple analytic model is supported by observational data and confirmed in numerical experiments. It makes predictions for the relations between the growth of galaxies and their black holes which will be tested by forthcoming observations. In particular, the model predicts that the black holes present in isolated low-mass galaxies will be small, with most objects having black holes smaller than $10^6\Msol$.

1607.07445

1607

1607.02056_arXiv.txt

{Current optical imaging surveys for cosmology cover large areas of sky. Exploiting the statistical power of these surveys for weak lensing measurements requires shape measurement methods with subpercent systematic errors.}{We introduce a new weak lensing shear measurement algorithm, shear nulling after PSF Gaussianisation (SNAPG), designed to avoid the noise biases that affect most other methods.} {SNAPG operates on images that have been convolved with a kernel that renders the point spread function (PSF) a circular Gaussian, and uses weighted second moments of the sources. The response of such second moments to a shear of the pre-seeing galaxy image can be predicted analytically, allowing us to construct a shear nulling scheme that finds the shear parameters for which the observed galaxies are consistent with an unsheared, isotropically oriented population of sources. The inverse of this nulling shear is then an estimate of the gravitational lensing shear.} {We identify the uncertainty of the estimated centre of each galaxy as the source of noise bias, and incorporate an approximate estimate of the centroid covariance into the scheme. We test the method on extensive suites of simulated galaxies of increasing complexity, and find that it is capable of shear measurements with multiplicative bias below 0.5 percent.}{}

The effect that masses can act as lenses and bend the path of light rays is called gravitational lensing. In the weak lensing regime first considered by \citet{firstwl} we statistically measure the slight distortion of the shapes of background galaxies by foreground lenses, called the shear. The subtle effects of weak gravitational lensing on galaxy shapes are an immensely powerful tool in observational astronomy. Amongst other applications, weak lensing has been an invaluable tool for cosmology through measurements of shear-shear correlations, called cosmic shear, which are connected to the dark matter power spectrum. After its first detection 15 years ago \citep{bacon,vanwaerbeke,wittman00,kaiser_cosmicshear} cosmic shear has been extensively used in cosmological studies (e.g. \citealt{kilbinger,kids450,des_jarvis}). Currently, large ($>$1000 deg$^2$) cosmic shear surveys are ongoing, such as the Kilo Degree Survey \citep{jelte_kids}, the Dark Energy Survey \citep{des}, and Hyper Suprime-Cam \citep{hsc}; more hemisphere-sized missions are planned, such as LSST \citep{lsst}, \textit{WFIRST} \citep{wfirst}, and \textit{Euclid} \citep{euclid}. These surveys will observe unprecedented numbers of galaxies, pushing down statistical errors, and hence requiring percent (for ongoing missions) to subpercent level accuracies (for future missions) on the measured galaxy shapes. In order to conduct weak lensing studies a crucial point is to measure the shapes of faint background galaxies with high accuracy as well as high precision in the face of inevitable noise, finite image resolution, and pixel effects. The first weak lensing techniques used the moments of the galaxy's image to estimate its shape and are known as moment-based methods (e.g. \citealt{kaiser}; \citealt{ksb} (hereafter KSB); \citealt{rhodes00}). These techniques need to use a weighting function with which to cut off the moment integrals so that the moments are not dominated by noise. Having to correct for the effect of the weight function and the PSF convolution are the main challenges for this class of techniques. The widely used KSB method employs an approximate deconvolution scheme, which assumes that the PSF is nearly Gaussian. Newer moment-based methods have improved upon the PSF correction \citep{deimos}, and there have been methods that change the PSF to make the measurement more exact (as explained in \citealt{hirataseljak} and used by \citealt{mandelbaum_regauss,eholics_era,okura_psfsmearing}). An alternative class of techniques relies on models of galaxies which are convolved with a PSF and then fit to the galaxy image and are hence known as model-fitting methods \citep[e.g.][]{kuijken99,lensfit,im3shape}. These techniques have the benefit of an accurate treatment of the PSF, but in return require realistic models of galaxies. The model of a galaxy is usually a parametric model (e.g. a linear combination of Sersic profiles) and if it does not resemble the intrinsic galaxy, the results can be biased \citep{bernstein_mb,voigtbridle}. A similar class of techniques, known as shapelets methods \citep{bernsteinjarvis,refregierbacon_shapelets,kuijken_shapelets}, use a set of basis functions which can, in theory, model any galaxy morphology by invoking ever higher order functions. However, in practice the order has to be truncated as the higher functions are dominated by noise, again leading to an unrepresentative galaxy model. In addition, noise in the galaxy image biases all shape measurement methods due to the non-linear dependence of the galaxy's ellipticity (the usual description of its shape) on the surface brightness \citep{refregier_nb,melchior12,viola14}. In order to quantify these uncertainties and to find ways of calibrating the different techniques, the weak lensing community started shape measurement challenges in which teams competed by using their methods to obtain the most unbiased shear estimate. This started with a general census and benchmark tests in the STEP challenges \citep{step1,step2} and continued with GREAT challenges \citep{great08,great10,great3}, which focused on the understanding of different sources of bias. After the most recent GREAT3 challenge it appears that the development in shape measurement algorithms is slowly reaching the goals set by ongoing cosmic shear surveys. The recent improvement in accuracy was mainly due to the advanced understanding of noise bias. Several authors have introduced correction schemes into their shape measurement methods which are able remove a large portion of the noise bias \citep{great3}. An alternative route is to avoid biased shear estimators by using estimators with a linear response to the pixel data instead of traditional non-linear variables, such as the ellipticity. Several authors have used the second moments of the galaxy's image brightness to estimate the shear \citep{zhangkomatsu,bernsteinarmstrong14,viola14}. Recently, \citet{bernsteinarmstrong16} have reported that their Bayesian method based on moments is able to reach subpercent accuracy even with low signal-to-noise ($S/N$) galaxies. However, the drawback of any Bayesian analysis is the requirement of accurate priors, for which external deep observations would be required. This requirement also means that there is no shear estimate for single galaxies, as then knowledge of the intrinsic galaxy profile would be needed, but only a shear estimate for an ensemble of galaxies. In this paper we propose a novel shape measurement method which may help to reach the ambitious goals of future cosmic shear experiments. Shear nulling after PSF Gaussianisation (SNAPG) is a moment-based method based on a circular Gaussian PSF and weight function, and requires the images to be preprocessed with a PSF Gaussianisation routine. For such galaxies we have an analytic relation between the moments of the galaxy and the shear. Shearing a population of galaxies introduces anisotropy to their ellipticity distribution. Using the analytic expressions, SNAPG reintroduces isotropy to this population by applying a nulling shear to the weighted second moments. The inverse of the nulling shear is then the shear estimate. Such a nulling technique was first advocated by \citet{bernsteinjarvis}. We propose an analytic correction to mitigate the bias due to centroid errors \citep{bernsteinjarvis}, which is directly computed from the galaxy image. SNAPG is similar to the \citet{bernsteinarmstrong16} method, but instead of a Bayesian framework it uses a nulling technique extract the shear from the second moments of a population of galaxies. It does not require a prior on the intrinsic moments of the galaxy population, but instead relies on the more general requirement that galaxy ellipticities are isotropic. Our novel method thus only produces a shear value for an ensemble of galaxies, but has the benefit that no auxiliary data is needed. In Sect. \ref{sec:theory} we introduce the SNAPG concept and a correction for the bias due to centroid errors. Section \ref{sec:sims} describes the image simulations we use to test SNAPG, and in Sect. \ref{sec:tests} and Sect. \ref{sec:great08} we present the results of the test runs. This is followed by a detailed discussion in Sect. \ref{sec:discussion}, and a summary in Sect. \ref{sec:summary}.

\label{sec:discussion} \subsection{SNAPG formalism} We have introduced the SNAPG formalism and tested its performance as a shear measurement method. For galaxy images convolved with a round Gaussian PSF the effect of shear on weighted second moments of image brightness can be analytically calculated. This analytical treatment is used to create a pipeline which finds the gravitational lensing shear by nulling the polarisations for an ensemble of galaxies. This procedure thus finds an estimate for the shear experienced by the galaxies. On test images with high $S/N$ galaxies convolved with a circular Gaussian PSF, the method obtained shear estimates deviating from the input shears by only parts per thousand. \subsection{Noise bias} Like most shape measurement methods, SNAPG suffers a noise bias when applied to images of galaxies with low $S/N$. However, by using only linear combinations of second moments instead of ratios of moments such as the polarisation or ellipticity, much of the noise bias can be avoided. This strategy allows SNAPG to obtain only a percent level bias in images with a $S/N\approx10$. Noise in the data introduces errors in the centroid estimates, which in turn biases the shear estimates. We compute an analytic treatment to correct the centroids and show that it can significantly improve the performance of SNAPG for low $S/N$ galaxies. Remaining biases after correction for $S/N\approx10$ are in the range of less than one percent. The residual biases increase with decreasing $S/N$, which indicates that the centroid error correction does not account for the full effect of noise bias. We traced their origin to the correlation between the centroid errors and pixel noise in the second moments. By removing the correlation, we can greatly decrease the measured bias, and also correct for the remaining bias with our centroid bias correction to subpercent accuracy. For multi-band surveys a possible solution is to use different filters for the estimates of the centroid and the measurement of the moments. In this way, the correlation between the centroid and the image is removed and without this correlation SNAPG can produce almost unbiased results. The impact on the bias of such a scheme will have to be investigated as galaxy colours and colour gradients may become an issue. In addition, this introduces a correlation with the photometric redshift estimate, which might pose a problem for cosmic shear measurements. \subsection{Galaxy resolution} The shape of a galaxy similar in size to the PSF is heavily distorted by the PSF, making it difficult to estimate the intrinsic shape. However, the analytic treatment of the PSF in the SNAPG formalism ensures that shear estimation is possible even for barely resolved galaxies. For galaxies 0.84 times smaller than the PSF, the shear was retrieved to similar accuracy, as were resolved galaxies for $S/N\approx$7-10 galaxies. By being able to measure unresolved galaxies reliably, SNAPG is able to use the large population of faint small galaxies to boost statistical power. \subsection{PSF Gaussianisation} We have run SNAPG on the images of the `LowNoise\textunderscore Known' and `RealNoise\textunderscore Known' branches of the GREAT08 challenge. To make them suitable for SNAPG, the GREAT08 images were first passed through the PSF Gaussianisation. We find a slight overestimation of the shear for the LNK images with $S/N=200$, of the order of 1-2\%. The PSF Gaussianisation introduces a correlation in the noise, which is analytically corrected for. SNAPG can retrieve the shear from the RNK images with $S/N=20$ to an accuracy that is similar to that for the high $S/N$ images. Further tests revealed that this percent level bias is probably inherent to the PSF Gaussianisation routine that we have used. For a variety of PSF profiles the multiplicative shear bias remained constant around 2\%. The PSF Gaussianisation appears to be the limiting factor for SNAPG to obtain subpercent shear bias and detailed investigation into this routine is necessary before SNAPG can be reliably applied to observations. We can compare the performance of SNAPG to the performance of the other methods tested in the GREAT08 challenge. This will only provide an indication as we did not run our pipeline on all datasets in the challenge and shear measurement methods have evolved since. However, a comparison to figures C3 and C4 in \citet{great08} shows that the 1-2\% bias SNAPG has obtained is at least competitive with other shear measurement methods. A more quantitative comparison to other (recent) shape measurement methods will require testing on image simulations which incorporate realistic observational features. However, optimistically the performance we find for SNAPG is sufficient to meet the requirements of the largest cosmic shear survey to date \citep{kids450} without any calibration being required. \subsection{Shear precision} \label{sec:precision} So far we have been concerned only with the accuracy of SNAPG, but an equally valid demand is high precision. To estimate the scatter in the shear estimate we use the simulated images of galaxies observed with the Hubble Space Telescope (HST) included in the GalSim software. These galaxies were observed as part of the COSMOS survey \citep{koekemoer07} and we used galaxies between magnitudes 20 and 24.5, similar to the depths of the Kilo Degree Survey and the Dark Energy Survey. These galaxies were rescaled to a pixel size of 0.214 arcseconds and convolved with a circular Gaussian PSF. We find that the scatter in the shear estimate for this set of galaxies is roughly $0.45/ \sqrt{N_{gal}}$, where $N_{gal}$ is the number of galaxies in the image. Thus the scatter in the SNAPG shear estimate for a fully realistic ensemble of galaxies is worse than an ellipticity based estimate; roughly 2-3 times more galaxies are needed by SNAPG for the same precision. This result is more optimistic than the increase by a factor of 10 found by \citet{viola14} in their analysis of a shear estimator based on Stokes parameters. Our use of a weight function reduces the variation of the moments, thereby shrinking the scatter in the Stokes parameter. In our tests we used identical weights for all sources, which naturally downweighs large, bright galaxies, which would otherwise dominate the ensemble average of second moments. Ideally, in order to optimise the $S/N$ of the individual moment measurement, the size of the weight function should match the observed size of the galaxy. However, fitting weight functions to individual galaxies is in itself a noisy process that may lead to a bias. We therefore advocate using the same weight function size for all galaxies (since most will be only partially resolved, it is not difficult to find a size that is nearly optimal for most of the galaxies by picking a small multiple of the PSF size; see also Eq. \ref{eq:wfsize}). A possible improvement is to assign each galaxy a weight to reduce the variance in the shear estimate. We find that for our sample of HST galaxies weighting by the inverse of the true flux can reduce the required number of galaxies by a factor of $\sim$4. This would bring the precision of SNAPG close to the precision of shear estimates based on galaxy ellipticities. In practice, estimating this weight factor from the galaxy fluxes measured in other images (e.g. adjacent photometric bands in a multi-band survey) uncorrelated with the lensing images will avoid introducing noise bias. \subsection{Variable shear} Observational weak lensing deals with varying shear fields, for instance in cosmic shear measurements or when measuring the mass of groups or clusters of galaxies. The traditional method is then to average the shear estimate for individual galaxies to obtain the lensing signal. This is not possible with SNAPG as it does not produce a shear estimate per source. In addition, SNAPG requires a large number of galaxies to obtain a precise shear estimate and satisfy the condition that the intrinsic ellipticities average to zero. Instead of nulling a single shear value for an ensemble of galaxies, we therefore advocate nulling a parametrised model shear field for that ensemble. For example, to measure a galaxy-galaxy lensing signal, the model would include parameters that describe the average shear profile of galaxies and their scaling with pertinent galaxy properties. The model parameters would then be varied until the average shear in a number of annular bins around the lensing galaxies is nulled, analogous to a traditional tangential shear stacking analysis. As another example, for cosmic shear measurements, the amplitudes of independent Fourier modes in the shear field could be nulled. \\ Developing this procedure will be left to the future. We have presented a new moment-based method that attempts to combine the best aspects of earlier approaches to the problem of high-accuracy, precise shear measurement from galaxy images. Moment-based methods generally approximate the deconvolution of the PSF, but do not require any information beyond the data and generally run very fast. Model fitting methods perform exact forward modelling, including convolution with the PSF, but are expensive to run because they need to search through a large parameter space, and may suffer model bias. The shear nulling after PSF Gaussianisation or SNAPG technique deals analytically with the PSF deconvolution and as a moment-based method only requires a few measurements on the data. In addition, SNAPG incorporates a correction scheme to mitigate the effects of noise bias, a major hurdle to all shape measurement techniques. Idealised test images show that SNAPG can retrieve shears to percent level accuracy for galaxies with low signal-to-noise, even if they are smaller in size than the PSF. The main issue limiting this technique is the correlation between the noisy estimate of the centroid and the pixel noise, which may be mitigated by incorporating further data about the sources, such as images from neighbouring bands in a multi-wavelength survey. In such a set-up, SNAPG can obtain shear estimates to subpercent accuracy for galaxies with a Gaussian PSF. Application to real data requires PSF Gaussianisation and if this routine is imperfect it can introduce percent level biases. This level of accuracy is comparable to what is required of the shape measurement algorithms used for ongoing surveys. As such, we expect SNAPG to be a useful asset for current and future weak lensing experiments. \subsection{Acknowledgements} We would like to thank Massimo Viola, Henk Hoekstra, and Peter Schneider for useful discussions. We also thank the anonymous referee, whose suggestions helped to improve this paper. RH acknowledges support from the European Research Council FP7 grant number 279396. AB was supported for this research partly through a stipend from the International Max Planck Research School (IMPRS) for Astronomy and Astrophysics at the Universities of Bonn and Cologne and through funding from the Transregional Collaborative Research Centre `The dark Universe' (TR 33) of the DFG. KK acknowledges support from an Alexander von Humboldt Foundation research award.

1607.02056

1607

1607.00265_arXiv.txt

We present the results of simultaneous spectroscopic and photometric observations of the cataclysmic variable star (hereafter CVs) V455 Andromedae, belonging to the WZ Sge sub-class, in quiescence. Using the spectroscopic data we computed time-resolved Doppler tomograms of the system demonstrating its behavior at different orbital phases. In the tomograms one can see the periodic brightening of different regions within one orbital cycle. We interpret this brightening as being due to the interaction of four phase-locked shock waves in the disk with a specific internal precessing density wave that develops inside the disk, because of the tidal influence of the secondary star, and this density wave propagates up to the disk's outer regions. The shock waves, located in the outer regions of the disk, are: two arms of the tidal shock; the "hot line", a shock occurring in the region of the interaction between the stream from the $L_{1}$ point and the disk; and a bow-shock, occurring ahead of the disk due to its orbital supersonic motion in the circumbinary gas. When the outer part of the density wave in its precessional motion reaches a shock wave the local density grows, which amplifies the shock (by increasing $\rho V^{2}/2$). This results in an additional energy release in the shock and can be observed as a brightening. Analysis of the tomographic results and the photometric data shows that two main sources contribute to the light curves of the system: the radiation of the "hot line" and the bow-shock gives us two major orbital humps, located approximately at the orbital phases $\phi=0.25$ and $\phi=0.75$; the amplification of the four shock waves may give us up to four "superhumps" shifting over the light curve with the precessional period. These two effects, when overlapping, change the shape of the light curves, shift the hump maxima, and they sometimes produce more than two humps in the light curve. We should emphasize that when saying "superhumps" we imply an effect that is observed in quiescent light curves of WZ Sge stars, as opposed to ''classical'' superhumps usually observed in outbursts.

WZ Sge stars (sub-class of SU UMa stars) are evolved close binary systems comprised of late type secondary components, filling their Roche lobes, and white dwarfs as primary components. To date, there have been about 60-100 representatives of this class discovered (see, e.g., \citet{katy2014, goly2014}). Specific orbital periods of these objects are near the period minimum (usually, $\approx$80 min.), and mass ratios are extremely low ($q=M_{2}/M_{1}<0.1$). In some works (see, e.g., \citet{araj} and references inside) the authors suppose that some representatives of this sub-class can even harbor brown dwarfs, no longer undergoing fusion reactions, as the secondary. One of the most noticeable observational features of the WZ Sge stars is their very powerful and rare (recurrent periods up to decades) super-outbursts when a star increases its brightness by over $6^{m}$. "Normal" outbursts in these systems are very rare or even absent (\citet{hara}). The main photometric observational features of these objects on short time scales (one orbital period) are the so-called double-humped light curves (see, e.g., \citet{araj, katy2009}), demonstrated in quiescence, i.e. one observes two pronounced humps within a single orbital period of a WZ Sge type system. There is still no commonly accepted model about how WZ Sge orbital light curves appear. Some authors (see, e.g., \citet{avil}, \citet{tovm}) suppose that the two humps may be caused by radiation from two arms of a tidal shock or density waves, formed in the disk under the action of tidal resonances. The shape and position of the humps in this model are determined by the viewing aspect of each arm. Another model proposes the ricochet of the gas stream, issuing from the inner Lagrangian point, and the formation of two bright regions in the disk, one "classical" hot spot and another on the opposite side of the disk (see \citet{wolf, Silv12}). In the frame of this model two bright spots give two humps in the light curve. However the proposed models cannot explain a number of effects, observed in WZ Sge stars. One of these effects is the varying number of the humps observed within one orbital period. For example, \citet{kitsi} report the results of long-term monitoring of the WZ Sge system, the sub-class prototype, and mention that in different observational epochs they see from two to three pronounced humps. Studying the system SDSS J080434.20+510349.2, related to WZ Sge stars, \citet{pavl2009} reports four (!) humps. Sometimes the reported number of the humps may be reduced to only one per orbital period (\citet{araj, katy2009}). This also means that it would be more correct to call the light curves of WZ Sge stars multi-humped instead of double-humped. In addition, the humps may shift over binary phases from one observational epoch to another and change their shape. The observational facts, mentioned above, contradict the proposed models of two tidal arms or two bright spots. First of all, no one of them can explain the varying number of the humps. Then, the results of the numerical simulations (see, e.g. \citet{smh_1, smis, BiBo04, Bi05, KuBi01}), conducted using grid-based methods, and the results of Doppler tomography (see e.g. \citet{shh97}) predict that the arms of the tidal shock in the disks of CVs are phase-locked. Indeed, if these spiral shocks had not been phase-locked we wouldn't have seen them in Doppler tomograms as separate bright regions, since, by definition, any flow element moving in the reference frame, co-rotating with the binary, should be smeared into a ring in the resulting tomogram. Therefore, if we suppose that the humps occur due to the viewing aspects of these features, we should always see them at the same phases. Recently (\citet{ko2015}) proposed a model, based on the results of 3D gas dynamic simulations of the V455 Andromedae system, that can explain almost all the properties of double-humped light curves including the varying number of the humps and their shift over orbital phases. Within this model we interpret the observed effects by the interaction of four phase-locked shock waves in the disk with a specific internal precessing density wave that develops inside the disk due to the tidal influence of the secondary star, and propagates up to its outer regions. We discuss the model in detail below. In this work we aim to find observational evidence that our model is correct. The structure of the paper is as follows. In Section 2 we describe our gas-dynamical model. In Section 3 we describe our spectroscopic and photometric observations. Section 4 is focused on analysis of observational Doppler tomograms, calculated using the obtained profiles of the $H_{\alpha}$ emission line. In Section 5 we compare the observational and synthetic Doppler tomograms and analyze the photometric data. In the section "Conclusions" we summarize the results of our study.

Recently \cite{ko2015} proposed a model, based on the results of 3D gas dynamic simulations of the V455 Andromedae system, that can explain almost all the properties of WZ Sge stars' light curves. The model predicts the periodic amplification of four major shock waves in the accretion disk (hot line, two arms of the tidal shock, and bow-shock ) by a specific internal precessional density wave. This periodic amplification can cause up to four humps in the light curve of the system. To examine the model, we conducted simultaneous spectroscopic and photometric observations of the cataclysmic variable star V455 Andromedae (WZ Sge subclass). Within 6 days, from September 22 to September 27, 2014, at the Kourovka Observatory of the Ural Federal University (Russia), we obtained 5 light curves of the system (except the night of September 26 when the weather conditions were bad). The spectroscopic observations, conducted at the Terskol Observatory (Terskol branch of The Institute of Astronomy of the Russian Academy of Sciences), resulted in obtaining 60 profiles of the $H_{\alpha}$ line on September 22, 2014. The other nights, unfortunately, were unsuccessful for spectroscopy because of poor weather conditions. We used the $H_{\alpha}$ line profiles for Doppler tomography of the V455 And system to calculate 10 Doppler tomograms of the system that demonstrate its 10 subsequent states within one orbital period. Besides, using the results of 3D gas-dynamical modeling of V455 And, described by \cite{ko2015}, we simulated line profiles and calculated 10 corresponding synthetic Doppler tomograms. The profiles were simulated under the assumption that the disk is optically thin. By analyzing the simulated and observed data we have found the following properties of the quiescent accretion disk in V455 And as they were on September 22, 2014: \begin{itemize} \item In the simulated disk four major shock waves occur. These are: two arms of the tidal shock wave; the hot line, a shock wave occurring due to the interaction of the circum-disk halo and the gas stream from the inner Lagrangian point; and the bow-shock caused by the super-sonic motion of the accretor and disk in the gas of the circum-binary envelope. We can also observe a one-armed spiral density wave, starting in the vicinity of the accretor and propagating to the outer regions of the disk. \item Under the tidal action of the secondary component the density wave retrogradely precesses with an angular velocity that differs from the orbital angular velocity of the system by only 1-2\%. In course of disk's rotation each of the four shocks passes through the outer part of the density wave. This results in the periodical local density increase in the region of each shock wave. Since the energy, released in shock waves as radiation, depends on the kinetic energy of gas $\rho V^{2}/2$, by increasing the density $\rho$ we increase the energy release in the shock wave, which may be observed as increased brightness or a hump in the light curve. \item Observed and simulated trailed spectra demonstrate similar structure and behavior. This enables us to suppose that the shape of the observed line profiles is determined by the gas-dynamical processes that we consider in our model. \item In both the synthetic and observational Doppler tomograms we see the periodical brightening of four regions. The regions correspond to the four major shock waves in the disk, which is in a good agreement with the proposed model. \item In the light curve of September 22 we observe that primary minima are shallower than the secondary minima. This may be explained by the fact that the hot line was amplified by the density wave when the system was observed at the orbital phase $\phi=0.0$. The resulting additional brightness overlaps with the primary eclipse, which makes the primary minimum shallower. \item Analysis of the light curves shows that the light curves of V455 Andromedae consist of two main components. One component is determined by the viewing aspects of the two strongest shocks in the disk, the hot line, and the bow-shock. The system's brightness variations in this component are modulated with the orbital period. The second component of the light curve contains up to four superhumps occurring because of the interaction of the spiral density wave with four shocks in the disk, with the most powerful superhumps produced by the amplification of the hot line and bow-shock. These superhumps shift over the orbital phases with the period of apsidal motion of the spiral density wave. The overlapping of the main humps and the superhumps results in the shift of the humps' maxima and sometimes, may even change the number of observed humps. \end{itemize} The obtained observational results allow us to suppose that our model of light curve formation is correct at least for V455 And. It can explain the varying number of the humps in the light curves of WZ Sge stars, since it implies the amplification of four shock waves, potentially giving four superhumps. This effect cannot be explained in the frame of models proposed by, e.g. \citet{avil}, or \citet{wolf, Silv12}, since they imply only two sources contributing to the light curves (two-armed tidal shock by \citet{avil}, or two hot spots by \citet{wolf, Silv12}). Our model can also explain the shift of the hump maxima, observed in different nights, by the difference between the precession angular velocity and orbital angular velocity of the system. At the next orbital cycle, the retrogradely precessing density wave interacts with the certain phase-locked shock wave a little earlier, which causes the corresponding superhump to occur at an earlier orbital phase. This may give negative superhumps. The mentioned alternative interpretations can explain neither positive nor negative superhumps since the sources of radiation they consider are phase-locked (tidal shock or hot spot). Thus, we should observe their contribution to the light curves always at the same orbital phases, which would eliminate periodicities, except the orbital variations. The results, discussed in this paper, should, in a sense, be considered as preliminary, though giving much new information on the flow structure in WZ Sge-type stars. Further, we plan on obtaining better observational material. In particular, we, of course, need more successful nights of simultaneous spectroscopic and photometric observations to perform a more detailed analysis. Also, we need to test our model with the other representatives of WZ Sge subclass. In particular, we need to clarify the question about the positive superhumps, if they are observed in quiescent light curves of WZ Sge-type stars, since within the proposed model, implying the retrograde disk precession, we can obtain only negative superhumps. The archive, containing animated .gif files, demonstrating the evolution of the simulated disk, "dynamic" synthetic and observational tomograms of V455 Andromedae, is available via DOI 10.13140/RG.2.1.1107.8884 on the ResearchGate (www.reaserchgate.net) web page.

1607.00265

1607

1607.00924_arXiv.txt

{Research on supernovae (SNe) over the past decade has confirmed that there is a distinct class of events which are much more luminous (by $\sim2$ mag) than canonical core-collapse SNe (CCSNe). These events with visual peak magnitudes $\lesssim-21$ are called superluminous SNe (SLSNe). The mechanism for powering the light curves of SLSNe is still not well understood. The proposed scenarios are circumstellar interaction, the emergence of a magnetar after core collapse, or disruption of a massive star through pair production.} {There are a few intermediate events which have luminosities between these two classes. They are important for constraining the nature of the progenitors of these two different populations as well as their environments and powering mechanisms. Here we study one such object, SN 2012aa.} {We observed and analysed the evolution of the luminous Type Ic \sn. The event was discovered by the Lick Observatory Supernova Search in an anonymous ($z\approx0.08$) galaxy. The optical photometric and spectroscopic follow-up observations were conducted over a time span of about 120 days.} {With an absolute $V$-band peak of $\sim-20$ mag, the SN is an intermediate-luminosity transient between regular SNe~Ibc and SLSNe. \sn\ also exhibits an unusual secondary bump after the maximum in its light curve. For \sn, we interpret this as a manifestation of SN-shock interaction with the circumstellar medium (CSM). If we would assume a $^{56}$Ni-powered ejecta, the quasi-bolometric light curve requires roughly 1.3 \msun\ of $^{56}$Ni and an ejected mass of $\sim 14$ \msun. This would also imply a high kinetic energy of the explosion, $\sim5.4\times10^{51}$ ergs. On the other hand, the unusually broad light curve along with the secondary peak indicate the possibility of interaction with CSM. The third alternative is the presence of a central engine releasing spin energy that eventually powers the light curve over a long time. The host of \sn\ is a star-forming Sa/Sb/Sbc galaxy.} {Although the spectral properties of \sn\ and its velocity evolution are comparable to those of normal SNe~Ibc, its broad light curve along with a large peak luminosity distinguish it from canonical CCSNe, suggesting the event to be an intermediate-luminosity transient between CCSNe and SLSNe at least in terms of peak luminosity. Comparing to other SNe, we argue that \sn\ belongs to a subclass where CSM interaction plays a significant role in powering the SN, at least during the initial stages of evolution. }

\label{intro} The study of SLSNe \citep{2012Sci...337..927G} has emerged from the development of untargeted transient surveys such as the Texas Supernova Search \citep{2006PhDT........13Q}, the Palomar Transient Factory (PTF; \citealt{2009PASP..121.1334R}), the Catalina Real-time Transient Survey (CRTS; \citealt{2009ApJ...696..870D}), and the Panoramic Survey Telescope \& Rapid Response System (Pan-STARRS; \citealt{2004AN....325..636H}). SLSNe are much more luminous than normal core-collapse SNe (CCSNe; \citealt{1997ARA&A..35..309F}). It was with the discovery of objects such as SNe 2005ap and SCP-06F6 \citep{2007ApJ...668L..99Q, 2011Natur.474..487Q}, as well as SN 2007bi \citep{2009Natur.462..624G} that events with peak luminosities $\gtrsim7\times10^{43}$erg\,s$^{-1}$ ($\lesssim-21$ absolute mag, over 2 mag brighter than the bulk CCSN population) became known. SLSNe have been classified into three groups: SLSN-I, SLSN-II and SLSN-R \citep{2012Sci...337..927G}. The SLSN-II are hydrogen (H) rich (e.g., SN 2006gy, \citealt{2007ApJ...666.1116S}; CSS100217, \citealt{2011ApJ...735..106D}; CSS121015, \citealt{2014MNRAS.441..289B}), while the others are H-poor. The SLSN-R (e.g., SN 2007bi, \citealt{2009Natur.462..624G}) have post-maximum decline rates consistent with the $^{56}$Co $\rightarrow$ $^{56}$Fe radioactive decay, whereas SLSNe-I (e.g., SNe 2010gx, \citealt{2010ApJ...724L..16P}; SCP-06F6) have steeper declines. H-poor events mostly exhibit SN~Ic-like spectral evolution \citep{2010ApJ...724L..16P, 2011Natur.474..487Q, 2013ApJ...770..128I}. However, the explosion and emission-powering mechanisms of these transients are still disputed. It is also not clear whether CCSNe and SLSNe are originated from similar or completely different progenitor channels, or if there is a link between these kinds of explosions. In particular, the SLSNe-R were initially thought to be powered by radioactive decay, but later the energy release of a spin-down magnetar was proposed \citep{2013ApJ...770..128I}, as was the CSM-interaction scenario \citep{2015ApJ...814..108Y, 2015arXiv151000834S}. In this present work, although we use the notation of the initial classification scheme, by SLSNe-R we simply mean those Type Ic SLSNe which show shallow (or comparable to $^{56}$Co decay) decline after maximum brightness, while the events which decline faster than $^{56}$Co decay are designated as SLSNe-I. The basic mechanisms governing the emission of stripped-envelope CCSNe are relatively well known. At early times ($\lesssim5$\,d after explosion), SNe~Ibc are powered by a cooling shock \cite{2013ApJ...769...67P, 2015A&A...574A..60T}. Thereafter, radioactive heating ($^{56}$Ni $\rightarrow$ $^{56}$Co) powers the emission from the ejecta. Beyond the peak luminosity, the photosphere cools and eventually the ejecta become optically thin. By $\sim100$\,d post explosion the luminosity starts to decrease linearly (in mag) with time, and it is believed that $^{56}$Co is the main source powering the light curve during these epochs (\citealt{1982ApJ...253..785A}, and references therein). \begin{figure} \centering \includegraphics[width=8.5cm]{2012aa_id.pdf} \caption{Identification chart of SN 2012aa. The image is about $10'$ on a side, taken in the $R_c$ band with the 1.04\,m Sampurnanand telescope at ARIES, Nainital. The local standard stars are numbered. North is up and east is to the left. The location of the target is marked ``T'' --- as observed with the 8\,m Gemini South telescope (bottom-left inset) on 2012 April 26, +88\,d after discovery. The SN is resolved; its position and the host center are marked. The field was also observed with the 3\,m TNG (top-left inset) on 2013 April 9, long after the SN had faded away. The SN location is not resolved in the TNG image.} \label{fig:snid} \end{figure} Three different models have been suggested for the explosion mechanism of SLSNe. A pair-instability supernova (PISN) produces a large mass of radioactive $^{56}$Ni, a source of enormous optical luminosity \citep{1967PhRvL..18..379B, 1967ApJ...148..803R}. Alternatively, a similarly large luminosity can be produced by interaction between the SN ejecta and a dense circumstellar medium (CSM) surrounding the progenitor, and the radiation is released when the strong shocks convert the bulk kinetic energy of the SN ejecta to thermal energy \citep{2007ApJ...671L..17S, 2011ApJ...729L...6C, 2012ApJ...757..178G}. The third alternative is to invoke an engine-driven explosion: a protomagnetar is a compact remnant that radiates enormous power during its spin-down process \citep{2010ApJ...717..245K, 2010ApJ...719L.204W}. Although both magnetar engines and CSM interaction can reproduce the range of luminosities and shapes of SLSN light curves, there are very few SNe that show a PISN-like $^{56}$Co $\rightarrow$ $^{56}$Fe decay rate after maximum light. Only $\sim10$\% of SLSNe show a slowly fading tail \citep{2013Natur.502..346N, 2015MNRAS.448.1206M}, but do not exactly follow the radioactive decay rate of cobalt. SLSNe 2006gy and 2007bi could be explained as PISN producing $\sim5$--10 \msun\ of $^{56}$Ni \citep{2007AIPC..937..412N, 2009Natur.462..624G}, but there are also counterarguments because the rise times of some similar events are shorter than the PISN model predictions (\citealt{2011ApJ...734..102K, 2013MNRAS.428.1020M, 2013Natur.502..346N}, and references therein). Beyond the general evolutionary picture of Type I SNe (hereafter we exclude SNe~Ia), many events exhibit different kinds of peculiarities. One of these, observed in SN 2005bf, is the appearance of a short-lived bump before the principal light-curve peak \citep{2005ApJ...631L.125A}. Two hypotheses have been proposed to explain this early short-duration bump: either a burst of radiation from a blob containing radioactive $^{56}$Ni that was ejected asymmetrically \citep{2007ApJ...666.1069M}, or a polar explosion powered by a relativistic jet \citep{2006ApJ...641.1039F}. Although a double-peaked light curve was observed in only a single SN~Ic, there are at least a handful ($\sim 8$) Type I SLSNe that exhibit a short-duration bump before the principal broad peak \citep{2016MNRAS.457L..79N}. Most of them have been described using the magnetar model. For SN-LSQ14bdq it was proposed that the intense radiation in optical and ultraviolet wavelengths was due to the emergence of the shock driven by a high-pressure bubble produced by the central engine, a proto-magnetar\citep{2016ApJ...821...36K}. Since H-poor SLSNe and normal Type Ic events show similar spectral features beyond maximum light, it is worth to explore whether there is any link between them. In fact the gap between SNe and SLSNe has not been well explored. The Type Ib SN 2012au showed spectral signatures which mark the transition between H-poor CCSNe and SLSNe, although its peak magnitude was comparable to those of H-poor (stripped envelope) CCSNe \citep{2013ApJ...770L..38M}. Recently, a few transients such as PTF09ge, PTF09axc, PTF09djl, PTF10iam, PTF10nuj, and PTF11glr with peak absolute magnitudes between $-19$ and $-21$ were discovered \citep{2014ApJ...793...38A}. Except PTF09ge, all had H-dominated spectra, and except PTF10iam they showed spectral behavior like that of tidal disruption events (TDEs; \citealt{1988Natur.333..523R}). However, note that TDEs may be even more luminous than SLSNe --- for example, the ROTSE collaboration found an event (nicknamed Dougie) that had a peak magnitude of $\sim-22.5$ \citep{2015ApJ...798...12V}. Unlike CCSNe or SLSNe, TDEs are always found at the center of their host galaxies, and their spectral features are normally dominated by H/He emission lines, thus differing from Type Ic SNe. \citet{2016ApJ...819...35A} also found four SNe having rapid rise times ($\sim10$\,d) and peak magnitudes of $\sim-20$ along with an \ha\ signature in their spectra (see Sect. \ref{res:otherpossibility}). In this work, we discuss the evolution of a SN~Ic that shows properties intermediate between those of SNe~Ic and H-poor SLSNe. \sn\ (= PSN J14523348-0331540) was discovered on 2012 January 29.56 UT (JD 2,455,956.04) in a relatively distant anonymous galaxy (redshift $z = 0.083$; see Sect. \ref{DistExt}) during the Lick Observatory Supernova Search (LOSS; \citealt{2001ASPC..246..121F}) with the Katzman Automatic Imaging Telescope (KAIT). An image of the field of \sn\ is shown in Fig.~\ref{fig:snid}. Basic properties of the host galaxy are given in Table~\ref{tab:propgal}. The event was observed at $\sim17.7$ mag (unfiltered) at discovery, and was further monitored in the subsequent nights and then spectroscopically classified as a SN~Ic \citep{2012CBET.3015....1C} on February 2 (UT dates are used throughout this paper). It was not observed in the X-ray in radio domains. However, independent optical photometry along with spectroscopic follow-up observations were obtained using various telescopes around the globe. We find (Sect. \ref{phot}) that the event was actually discovered near its peak, but it was caught in its rising phase by the CRTS survey \citep{2009ApJ...696..870D}. It shows some spectroscopic resemblance to H-poor SLSNe, although photometrically the peak is about 1\,mag fainter than that of typical SLSNe. More remarkably, \sn\ shows an unusual secondary bump after a broad primary peak in its light curve (Sect. \ref{phot}). \begin{table} \centering \caption{\label{tab:propgal}Properties of SN 2012aa and its host galaxy} \vskip 0.2cm \scriptsize \begin{tabular}{llc} \hline\hline Parameters&Value&Ref\\ \hline {\bf Anonymous Host:}&&\\ Type& Sa/Sb/Sbc&Sect. \ref{Host}\\ \\ Position & $\alpha_{\rm J2000} = 14^{\rm h} 52^{\rm m} 33\fs55$&Sect. \ref{obs}\\ & $\delta_{\rm J2000} = -03\degr 31\arcmin 53\farcs82$& \\ \\ Abs. magnitude& $M_{B}=-18.13$ mag&Sect. \ref{Host}\\ \\ Redshift& $z=0.0830\pm0.0005$ &Sect. \ref{DistExt}\\ \\ \\ Distance& $D=380.2$\,Mpc&Sect. \ref{DistExt}\\ \\ Distance modulus& $\mu \approx 37.9$\,mag&Sect. \ref{DistExt}\\ \\ Projected dimension& $\sim 6$\arcsec &Sect. \ref{Host}\\ \\ Metallicity of the host& Z$_{\rm host}\approx0.92\pm0.34$\,\zsun\,&Sect. \ref{Host}\\ \\ {\bf SN 2012aa:}&&\\ Position & $\alpha_{\rm J2000} = 14^{\rm h} 52^{\rm m} 33\fs56$&Sect. \ref{obs}\\ & $\delta_{\rm J2000} = -03\degr 31\arcmin 55\farcs45$& \\ \\ Separation from center&$\sim3$ kpc&Sect. \ref{Host}\\ \\ Discovery date (UT)& 29.56 January 2012&Sect. \ref{intro}\\ & (JD 2,455,956.04)& \\ \\ Epoch of $V$ maximum: &JD 2,455,952&Sect. \ref{phot}\\ \\ Total reddening toward SN: & \ebv\,$=0.12$\,mag&Sect. \ref{DistExt}\\ \hline \end{tabular} \end{table} Here we present an optical photometric and low-resolution spectroscopic investigation of \sn. We also make comparisons among \sn, SLSNe, and CCSNe. The paper is organized as follows. Section~\ref{obs} presents a brief description of the observations of \sn\ along with a description of the data-reduction procedure. The distance and extinction are estimated in Sect.~\ref{DistExt}. In Sect.~\ref{phot}, we study the light-curve evolution, while the evolution of the colours and of the quasi-bolometric luminosity are discussed in Sect. \ref{ColBol}. The spectroscopic evolution is presented in Sect. \ref{spec}, and the properties of the host galaxy in Sect. \ref{Host}. Section~\ref{Progenitor} explores different physical parameters of the explosion, and a comparative study of this event with SLSNe and CCSNe is given in Sect. \ref{res:comp}. Finally, Sect.~\ref{concl} summarizes our conclusions.

\label{concl} The nature of the progenitors and the explosion mechanism of SLSNe is still unclear. Whether there is a difference in the explosion mechanisms of SLSNe and canonical SNe~Ibc is also unknown. In this context, the study of luminous events that fall between CCSNe and SLSNe in various respects is important. Here we have carried out optical photometric and spectroscopic observations of \sn\ over an period of 120\,d. In this investigation, we proposed that \sn\ is a Type~Ic SN that occurred in a dense medium, which made the event more luminous owing to SN-CSM interaction and resulting in a photometric evolution similar to that of H-poor SLSNe. Different aspects of this study include the following. {The peak bolometric luminosity of \sn\ is $\sim1.6\times10^{43}$\,ergs ($M_{\rm Bol}\approx-20$\,mag), which is less than the typical peak luminosity of SLSNe~I but larger than that of SNe~Ic. The spectroscopic properties of \sn\ are similar to those of normal SNe~Ic. However, its rise timescale ($\tau_{\rm ris}=34$\,d) and decline timescale ($\tau_{\rm dec}=57$\,d) are more consistent with those of H-poor SLSNe.} {The event shows a secondary bump in all optical bands between +40\,d and +55\,d after maximum brightness. It is more pronounced in the $R_c$ and $I_c$ bands than in the $B$ and $V$ bands.} {Beyond +47\,d after maximum brightness, we noticed an emission peak in the spectra at rest wavelength 6500\,\AA. There are two possibilities: either this is \Oia\ $\lambda\lambda$6300, 6364 or \ha\ in emission. In the first scenario, the line-emitting region is moving away with a projected velocity of $\sim9500$\,\kms. In the second scenario the emission is \ha\ blueshifted by 3000\,\kms; since it appeared at late epochs, it is probably part of the CSM and was powrwed by interaction. The identification of \ha\ is consistent with other line profiles. On the other hand, comparison of spectra of \sn\ with those of other SNe~Ic and SLSNe suggests that this particular feature is more likely \Oia\ rather than \ha.} {Assuming that only $^{56}$Ni decay is responsible for powering the SN, from the quasi-bolometric light curve we require roughly 1.3\,\msun\ of $^{56}$Ni ejected in this explosion. The ejected mass is $\sim 14$\,\msun, implying a kinetic energy of $\sim5.4\times10^{51}$\,erg. On the other hand, if the entire explosion is CSM-interaction dominated, a similar explosion can be produced under the presence of CSM with a mass of $\sim5$--10\,\msun.} {We have also explored the possibility of the emergence of a magnetar. Assuming that the velocity of the ejecta near peak luminosity is $\sim12,000$\,\kms, the kinetic energy of the explosion is $\sim10^{51}$\,ergs, and the rise time is 30--35\,d, we find that the quasi-bolometric light curve of \sn\ can be fitted with a magnetar having a magnetic field of $(6.5\pm0.3)\times10^{14}$\,G, an initial spin period of $7.1\pm0.6$\,ms, and an explosion ejected mass of $16.0\pm1.7$\,\msun.} {With a limited SN sample consisting of PTF10hgi, SN 2011ke, SN 2012aa, and SN 2012il, we found that there is a potential subset of SLSNe~I showing a radioactivity-powered tail at relatively early times ($\sim 60$\,d after maximum brightness). This subset also has a narrower peak (i.e., smaller diffusion time) in comparison with other SLSN~I/SLSN-R events.} {The $z\approx0.08$ host of \sn\ is a star-forming (Sa/Sb/Sbc) galaxy. The overall metallicity ($Z_{\rm host}\approx0.92\pm0.34$\,\zsun) of the host is comparable to the solar metallicity and also to the metallicities of typical nearby spiral galaxies hosting SNe~IIP and normal SNe~Ibc.}

1607.00924

1607

1607.00115_arXiv.txt

{We present a star catalogue extracted from the Lunar-based Ultraviolet Telescope (LUT) survey program. LUT's observable sky area is a circular belt around the Moon's north pole, and the survey program covers a preferred area for about 2400\,$\deg^2$ which includes a region of the Galactic plane. The data is processed with an automatic pipeline which copes with stray light contamination, artificial sources, cosmic rays, flat field calibration, photometry and so on. In the first release version, the catalogue provides high confidence sources which have been cross-identified with Tycho-2 catalogue. All the sources have signal-to-noise ratio larger than 5, and the corresponding magnitude limit is typically 14.4\,mag, which can be deeper as $\sim$16\,mag if the stray light contamination is in the lowest level. A total number of 86,467 stars are recorded in the catalogue. The full catalogue in electronic form is available on line.

\label{sec:Intro} Lunar-based Ultraviolet Telescope (LUT) is the first robotic telescope deployed on the moon's surface, and is loaded inside the lander of China's Chang'e-3 lunar exploration program \citep{Ip2014RAA}. LUT and the lander are located at 44.12$\degr$N and 19.52$\degr$W on a basin of the moon named Sea Of Showers. LUT is an imaging telescope working at a characteristic near-ultraviolet (NUV) band. Since the successful launch and landing of Chang'e-3 in December 2013, LUT had finished the task of its one-year's mission phase and continued to work stably for another one year. The stability of its performance has been verified by an 18-month magnitude zero point (zp) calibration work. The photometric calibration gives $zp=17.53\pm0.05$\,mag, which is highly consistent with the results of the first 6-months of $zp=17.52\pm0.07$\,mag \citep{WangJ2015zp}. One of LUT's main scientific objectives is to perform a sky survey for an area about 2400\,$\deg^2$ \citep{Caoli2011}. Tens of star catalogues in NUV bands have already been published, which are contributed by Galaxy Evolution Explorer (GALEX), Hubble Space Telescope (HST), etc. GALEX has an all sky survey project named All-Sky Imaging survey (AIS), whose detection limit is $\sim$21\,mag at an NUV band. GALEX avoids the Galactic plane during the prime mission phase because of its high-countrate safety limits. Its latest survey covers regions in the Galactic plane, but these data are not reachable yet in the public archive \citep{BianchiGALEX2014}. LUT survey covers a part of the low Galactic latitude region within its available sky area, so it would be helpful for future researches for this region. Further more, the survey observation strategy of LUT enables the telescope to revisit some sky areas for more than 10 times, so it is potentially possible to find variable stars through further data mining. The LUT survey data are processed with an automatic pipeline, which is inherited from LUT's pointing observation data processing pipeline \citep[see][]{Meng2015ApSS}. The main different parts from the pointing program are: (1) the survey data processing has to clean off all the cosmic rays, so it has to perform image stacking; (2) a series of processes have to be carried out to clean off artificial sources that have arisen from stray light residuals. The first release of LUT survey data product, as described in this paper, is a star catalogue covering the whole LUT's available sky area, and is cross-identified with an optical star catalogue Tycho-2 \citep{Tycho2000}. Tycho-2 is the reference catalogue for LUT's astrometry, and its positional and photometric data are very helpful for LUT to remove artificial sources. After a quick look of the LUT instrument in Sect.~\ref{sec:Inst}, the survey observation strategy and its footprints are described in Sect.~\ref{sec:obs}. The details of the survey data processing pipeline are described in Sect.~\ref{sec:Pipeline}, and the performances in various aspects are also shown. Compilation of the catalogue and some statistical results are presented in Sect.~\ref{sec:catalogue}. Discussion on Galactic extinction and aperture correction is given in Sect.~\ref{sec:Discuss}.

\label{sec:Discuss} \subsection{Extinction} \label{subsec:extinct} The photometric results in the catalogue are not corrected by Galactic extinction. The reasons are: (1) the sources are in principle stars in the Milky Way and extinction correction for Galactic sources is complicated. At least, we need to have the information of their distances, but actually they are lacking; (2) the uncertainty of the correction should be significant. First, even if we had their distance information, the dust/gas amounts they had went through are hard to estimate; second, for example, a certain method of galactic dust reddening estimation may have an uncertainty of about 16\% \citep{Schlegel1998}, which may be much larger than the uncertainty of the photometry results; (3) the extinction correction may induce errors to photometric results, and these are systematic errors which come from the correction factors and should be treated differently from random errors. \subsection{Aperture Correction} \label{subsec:Apertcor} The catalogue doesn't include aperture corrections for photometry results. We give the correction factor separately in this section as a choice that is subject to the users. The reason is as follows. Because of the image combination, the form of the brightness profile of stars may vary in a considerable range, depending on the precision of the image alignment, profiles variation between single-exposures, field distortion, etc. Therefore, the aperture effect of the survey data photometry is a complicated problem. The aperture correction method for survey data follows the pointing observation \citep[see][Sect. 4]{Meng2015ApSS}. We measure the ``curve of growth'' for 26 bright stars on different survey images and their physical positions were varying. The scatter of the growth curves of survey program is obviously larger than that of pointing program, and their profiles are different (see the median stacked profiles shown in Figure~\ref{fig:surveygrowth}). Therefore, the uncertainty of aperture correction factor for survey program is larger than that for pointing program. Derived from the growth curve, the aperture correction factor for 2$\times$FWHM aperture radius is $\Delta m_{r_{\rm 2}}=-0.029\pm0.011$, if the aperture correction is performed with \begin{equation} m_{{r_{\rm 2}},{\rm cor}} = m_{r_{\rm 2}} + \Delta m_{r_{\rm 2}} \end{equation} where $m_{r_{\rm 2}}$ and $m_{{r_{\rm 2}},{\rm cor}}$ are the magnitudes before and after aperture correction, respectively. \begin{figure}[!htbp] \centering \includegraphics[width=0.8\columnwidth]{msRAA-2016-0054-R1fig8.eps} \caption{The curve of growth and its dispersions of LUT combined images of the survey data (blue circles and red error bars). The blue solid line is a Voigt model fitting to the curve of growth. The red dashed line representing the curve of growth of LUT single-exposure images, as has described by \citet{Meng2015ApSS}, is shown here for comparison.} \label{fig:surveygrowth} \end{figure} A star catalogue obtained from the observation data of LUT survey program is presented here, which has 86,467 entries of stars at NUV band. The sky coverage of the catalogue is about 2400\,$\deg^2$, a circular belt around the Moon's north pole, and a part of it has low Galactic latitude of $b<15\degr$. An automatic pipeline is developed to process the data, coping with stray light contamination and thereby false sources, cosmic rays, flat field calibrations, photometries, etc. In this first released version, the catalogue provides high confidence sources that have been cross-identified with Tycho-2 catalogue. The SNR is constrained to be $\geq 5$, and the corresponding detection limit of the LUT survey is about 14.4\,mag. Some statistical properties are given here. The full catalogue in electronic form is available as catalog I/335 from CDS.

1607.00115

1607

1607.00323_arXiv.txt

Either bulk rotation or local turbulence is widely invoked to drive fragmentation in collapsing cores so as to produce multiple star systems. Even when the two mechanisms predict different manners in which the stellar spins and orbits are aligned, subsequent internal or external interactions can drive multiple systems towards or away from alignment thus masking their formation process. Here, we demonstrate that the geometrical and dynamical relationship between the binary system and its surrounding bulk envelope provide the crucial distinction between fragmentation models. We find that the circumstellar disks of the binary protostellar system L1551\,IRS\,5 are closely parallel not just with each other but also with their surrounding flattened envelope. Measurements of the relative proper motion of the binary components spanning nearly 30\,yr indicate an orbital motion in the same sense as the envelope rotation. Eliminating orbital solutions whereby the circumstellar disks would be tidally truncated to sizes smaller than are observed, the remaining solutions favor a circular or low-eccentricity orbit tilted by up to $\sim$25\degr\ from the circumstellar disks. Turbulence-driven fragmentation can generate local angular momentum to produce a coplanar binary system, but which bears no particular relationship with its surrounding envelope. Instead, the observed properties conform with predictions for rotationally-driven fragmentation. If the fragments were produced at different heights or on opposite sides of the midplane in the flattened central region of a rotating core, the resulting protostars would then exhibit circumstellar disks parallel with the surrounding envelope but tilted from the orbital plane as is observed.

While there is a generally accepted framework for how single stars form, the manner in which binary and high-order multiple stars form remains contentious. Yet, the formation of such systems, rather than of single stars, constitutes the primary mode by which the majority of stars having masses comparable to and higher than that of the Sun form. At the present time, the debate on how binary (hereafter, used generically to also include higher-order multiple) stars form are focussed on two primary issues: (i) what mechanism(s) create the seeds for a binary system, and when are these seeds produced during the collapse of a gravitationally-bound condensation (core); and (ii) what factors determine the growth of the individual protostellar components to reach their final, often different, masses. Knowledge of the ingredients involved in defining the individual masses of binary components is critical for addressing another outstanding question in star formation, the origin of the stellar initial mass function. Historically, a number of pathways have been proposed for the formation of binary stars: fission, capture, and fragmentation \citep[e.g., review by][]{Bodenheimer2000, Tohline2002}. Numerical simulations show that rotational instabilities in self-gravitating spheroids do not lead to fission -- splitting into two (equal) pieces. Instead, the spheroid develops a central bar and spiral arms, the latter of which transport both angular momentum and mass outwards to stabilize the structure against fission \citep[e.g.,][]{Durisen1986, Williams1988}. Capture as a result of a close encounter between single protostars is energetically unfavourable, even when large and massive circumstellar disks are invoked to help absorb the impact energy \citep{Clarke1991, Clarke1993}. Encounters typically result in flybys, with each flyby partially disrupting the disk and reducing its ability to absorb the impact energy of the next encounter. Recently, gravitational instabilities in circumstellar disks have been invoked to form relatively low-mass binary companions \citep[e.g.,][]{Adams1989}. Given the large range in mass ratios exhibited by binary systems, however, this process can constitute, at best, a minor pathway for the formation of binary stars. In part therefore by a process of elimination, fragmentation -- the internal break-up of a core into two or more pieces -- has emerged as the leading contender for how the majority of binary stars form \citep[e.g., review by][]{Goodwin2007}. Two different mechanisms have been proposed to drive fragmentation: (i) bulk (large-scale ordered) rotation; and (ii) {local (small-scale)} turbulence. As described in more detail in $\S\ref{Fragmentation Models}$, depending on the actual circumstances involved, these two mechanisms can predict very different geometries and dynamics for the resulting binary system: i.e., alignment between circumstellar disks and/or spin axes of the binary components, as well as alignment between circumstellar disks and orbital plane or between the spin and orbital axes. Comparisons between binary properties and model predictions for their formation, however, are complicated by possible internal or external interactions during or after the protostellar phase. Depending on the nature of the interaction, the binary system can be driven either towards or away from alignment, altering its original geometry and dynamics {thus masking its formation process}. As we shall demonstrate in this manuscript, a more promising approach to distinguish between different fragmentation models is to study binary systems in the process of formation. Still embedded in their parental cores, the different proposed drivers of fragmentation make very different predictions for the geometrical and dynamical relationship between the {resulting} protostellar binaries and their surrounding envelopes (see $\S\ref{Fragmentation Models}$). In this manuscript, incorporating new data having a significantly improved angular resolution and sensitivity, we accurately deduce the alignment between the circumstellar disks, as well as between the circumstellar disks and surrounding envelope, of a binary protostellar system, the Class\,I object L1551\,IRS\,5. Furthermore, we thoroughly explore best-fit orbits to measurements of relative proper motion for the binary protostars, thus deducing the alignment between the orbital plane and the circumstellar disks. From the geometrical and dynamical relationship between the binary system and its parental core, % we reason that the physical properties of L1551\,IRS\,5 reflect the manner in which it formed rather than being imposed by subsequent interactions. As a consequence, its physical properties can be directly used to differentiate between different mechanisms that drove the fragmentation of its parental core. The work reported in this manuscript lays the foundation for theoretical simulations on how the binary components in L1551\,IRS\,5 are interacting with and accreting from their envelope, motivating planned observations of this system with the Atacama Large Millimeter and Submillimeter Array (ALMA). We recently conducted such a theoretical simulation for the Class\,I object L1551\,NE, assuming a circular coplanar orbit for the binary protostars as well as a coplanar circumbinary disk. We found the model predictions to be in good agreement with the observed structure and kinematics of its circumbinary disk as measured with ALMA \citep{Takakuwa2014}. This manuscript is arranged as follows. In $\S\ref{Observations}$, we describe our most recent observation of the circumstellar disks of the binary protostars in L1551\,IRS\,5, surpassing in both sensitivity and angular resolution previous observations. The results are presented in $\S\ref{Results}$, along with an explanation of how we separated the ionized jets from the circumstellar dust disks. In $\S\ref{Physical Parameters Circumstellar Disks}$, we describe how we determine the inclinations, relative alignments, and sizes of the binary circumstellar disks. In $\S\ref{Physical Parameters Envelope}$, we describe how we determine the geometry of their surrounding envelope, how this geometry relates to that of the circumstellar disks, and place constraints on a central cavity in the envelope. In $\S\ref{Orbit}$, we present all currently available measurements for the relative proper motion of the binary protostars, acceptable best-fit orbits to these measurements, and constraints placed on the orbits by the sizes of the circumstellar disks and tentative upper limits on the size of any central cavity. In $\S\ref{Discussion}$, we bring together all the aforementioned properties of L1551\,IRS5 to assemble a coherent picture for how this system formed. We also extrapolate into the future the likely evolution of L1551\,IRS\,5, and how its projected properties on the main sequence will compare with those of other binary systems having similar masses and orbital separations. In $\S\ref{Summary}$, we provide a thorough summary of our results, analyses, and interpretation. Throughout this manuscript, we assume a distance to L1551\,IRS\,5 of 140\,pc.

\label{Discussion} \subsection{Observational Constraints on Fragmentation Models}\label{Constraints on Fragmentation Models} We begin by summarizing the constraints imposed on models for the formation of L1551\,IRS\,5 as derived in $\S\ref{Physical Parameters Circumstellar Disks}$--$\S\ref{Orbit}$. From the measured geometry of its circumstellar disks and its surrounding envelope: \begin{itemize} \item the circumstellar disks of the binary protostars are very closely parallel \item the circumstellar disks also are very closely parallel to their surrounding flattened envelope \end{itemize} \noindent Based on the measured relative proper motion of the binary protostars and published measurements of the kinematics of its surrounding envelope: \begin{itemize} \item the orbital motion is in the same direction as the rotational motion of the surrounding envelope \end{itemize} \noindent {The close geometrical and dynamical relationship between the circumstellar disks and surrounding flattened envelope suggests that the angular momentum of material that gave rise to the protostellar system and that comprising its parental core are aligned.} \noindent Finally, although orbits spanning a wide range of eccentricities provide acceptable fits to the available measurements of the relative proper motion for the binary protostars in L1551\,IRS5, theoretical predictions for limits in the sizes of the circumstellar disks as imposed by tidal truncation rule out the vast majority of best-fit orbits having moderate eccentricites ($0.4 \lesssim e \lesssim 0.6$) and all best fit orbits having high eccentricities ($e \gtrsim 0.8$). For the remaining best-fit orbits: \begin{itemize} \item in the majority of cases, there is a systematic tilt (by up to typically $\sim$25\degr) between the circumstellar disks and the binary orbital plane; nevertheless, there is a clear preference for a relatively close alignment if not coplanarity \end{itemize} \subsection{Predictions of Fragmentation Models}\label{Fragmentation Models} As explained in $\S\ref{Introduction}$, current models invoke either {local (small-scale)} turbulence in or the bulk rotation of cores to drive fragmentation. In cores that have little or no bulk rotation, turbulence % introduces velocity and density inhomogeneities that can seed and drive the growth of multiple density perturbations to become self gravitating \citep[e.g.,][]{Bate2002, Bate2003, Bate2005, DelgadoDonate2004a, DelgadoDonate2004b, Goodwin2004a, Goodwin2004b, Goodwin2006}. Thereafter, dynamical interactions between the fragments can play an important role in determining the final properties of the binary system (e.g., number and masses of component stars, together with their orbital separation and eccentricity). Multiple fragments produced in different turbulent cells are predicted to exhibit random orientations between the circumstellar disks or spin axes of the binary components, and no particular relationship between the circumstellar disks and orbital plane or the spin and orbital axes. If multiple fragments are produced in a common region where turbulence conspires to create local angular momentum, however, the binary system thus assembled can exhibit quite well aligned circumstellar disks and orbital plane or spin and orbital axes. Rather than through turbulence, the large-scale ordered rotation of the core also can drive dynamical instabilities to induce fragmentation during collapse. In such models, conservation of angular momentum forces cores to become increasingly flattened as they collapse. As a result, a disequilibrium disk-like (i.e., flattened and rotating) structure forms at the center of the core. The central region of the core can become especially flattened if magnetic fields are invoked to direct infalling matter onto the mid-plane of the disk-like structure; the resulting structures closely resemble, at least morphologically, rotationally-supported disks, and are therefore referred to as pseudodisks \citep{Galli1993a, Galli1993b}. By introducing an initial density or velocity perturbation, the large-scale ordered rotation of the core can drive dynamical instabilities in the form of a spiral, bar, or ring in its central flattened region \citep{Matsumoto2003, Cha2003, Machida2008}. The pattern of the dynamical instability does not depend on the nature of the initial velocity or density perturbations. Furthermore, because the perturbations introduced can either promote or hinder fragmentation, in the latter case through the effective removal of angular momentum by gravitational torques from the resulting dynamical instability, increasing the amplitude of the initial perturbation does not necessarily increase the likelihood of fragmentation. Fragments form in localised regions of the resulting dynamical instabilities that are gravitationally unstable (according to the Toomre criterion) and have masses exceeding the local Jeans mass. Theoretical considerations suggest that, when driven by the large-scale ordered rotation of the core, fragmentation most likely occurs during the adiabatic phase (when a first hydrostatic core forms and grows) or the protostellar phase \citep[see arguments in][]{Machida2008}. Binaries that form through the rotation fragmentation of disk-like structures should naturally exhibit a close alignment between the spin axes of the component stars. Because fragments can be produced at different heights from and perhaps even on opposite sides of the mid-plane, however, the circumstellar disks and orbital planes of the resulting binary system need not be closely aligned but can span a range of angles. The mass ratio and orbital parameters (i.e., orbital separation and eccentricity) of such binary systems depend not only on the initial mass ratios and orbits of their parental fragments, but also how these fragments interact with each other (especially in systems initially comprising three or more fragments) as well as with the surrounding material from which they accrete. % \subsection{Formation of L1551\,IRS\,5} As explained above, both turbulent- and rotationally-driven fragmentation can produce binary systems in which the circumstellar disks are parallel with each other, as well as quite closely parallel if not coplanar with the binary orbit, satisfying some of the observed physical properties of L1551\,IRS\,5. The crucial distinction between the two fragmentation models, however, is the physical relationship between the binary system and its surrounding bulk envelope, constituting the remnant parental core. In models that invoke turbulent fragmentation, even in the situation where multiple fragments are produced in a common region where turbulence conspires to create local angular momentum (and thus produce an aligned binary system), the binary system need not orbit in the same direction as the rotation of the bulk envelope, let alone have circumstellar disks aligned with any flattening of the bulk envelope. On the other hand, in models that invoke rotational fragmentation, the binary system will of course orbit in the same direction as the rotation of its surrounding envelope, and its circumstellar disks aligned with the surrounding flattened envelope. The observed physical properties of L1551\,IRS\,5 therefore directly point to its formation through rotational fragmentation. In support of this argument, both the measured rotational velocity and the high degree of flattening exhibited by the envelope around L1551\,IRS\,5 suggest that its parental core possessed considerable angular momentum. As mentioned above, even when formed by rotational fragmentation, the circumstellar disks of the resulting binary protostars, although parallel with each other, need not be {accurately} aligned with the orbital plane. \citet{Bateetal2000} argue that tidal interactions can quite rapidly align the circumstellar disks of binary protostars with their orbital planes, such that low-mass binary protostellar systems with orbital separations $\lesssim 100$\,AU should have coplanar disks and orbital planes. If this process had indeed driven the circumstellar disks of L1551\,IRS\,5 towards close alignment with the binary orbital plane, the same process would have led to a misalignment between the circumstellar disks/orbital plane and the surrounding flattened envelope. Consider the situation where the two protostars formed at different heights from, and perhaps even on opposite sides of, the midplane of a flattened core, as shown in the schematic diagram of Figure\,\ref{Fragmentation}($a$). At birth, the circumstellar disks of the two protostars are therefore closely parallel with their surrounding flattened envelope, but not coplanar with each other or the orbital plane. If the circumstellar disks are then brought into alignment with the orbital plane by tidal interactions, the mutual plane of the binary system now becomes misaligned with the surrounding flattened envelope, as shown in Figure\,\ref{Fragmentation}($b$). By contrast, we find that the circumstellar disks of L1551\,IRS\,5 are very closely parallel to their surrounding flattened envelope, but probably misaligned by up to typically $\sim$25\degr\ from the orbital plane, a geometry more akin to that shown in Figure\,\ref{Fragmentation}($a$). The probable misalignment between the circumstellar disks/flattened envelope and the binary orbital plane may therefore reflect the formation of the binary fragments in different planes sharing the same rotational axis. By imposing the condition that the inferred break radii of the circumstellar disks cannot be larger than their theoretically predicted tidally-truncated radii, only a very narrow range of orbital periods, if any, are permitted for the best-fit orbits having eccentricities $e \gtrsim 0.4$, as explained in $\S\ref{Constraints Circumstellar Disks}$. Thus, most probably, L1551\,IRS\,5 has a circular or low-eccentricity orbit. Early in the protostellar and perhaps even first core phase when fragmentation occurs, the fragments have a very low mass compared with their surrounding envelope (parental core). In the situation where, at apastron, the orbital velocity of the fragments is smaller than the rotational velocity of their surrounding envelope (in the immediate vicinity), tidal interactions between the envelope and the fragments transfer both angular momentum and energy from the envelope to the binary. This action results in an increase in the orbital velocity and hence separation during the subsequent passage through periastron, and consequently a reduction in the orbital eccentricity. Similarly, at periastron, if the orbital velocity of the fragments is larger than the rotational velocity of their surrounding envelope (in the immediate vicinity), tidal interactions between the fragments and the envelope transfer both angular momentum and energy from the binary to the envelope. This action results in a decrease in the orbital velocity and hence separation during the subsequent passage through apastron, and once again a reduction in the orbital eccentricity. Thus, the probable low orbital eccentricity of L1551\,IRS\,5 may have been imprinted at birth. \subsection{Future Evolution} The total mass of L1551\,IRS\,5 is at least $\sim$$0.5 {\rm \ M_\sun}$ (based on the best-fit orbits), compared with a mass for its envelope of about $0.1 {\rm \ M_\sun}$ \citep{Momose1998}. Thus, during its remaining protostellar phase, L1551\,IRS\,5 is not likely to accrete more than $\sim$20\% of its present total mass and perhaps much less. Thus, both the total mass and the mass ratio of the binary system should be largely preserved as it evolves onto the pre-main-sequence and then the main-sequence. Whether or not the orbital parameters of L1551\,IRS\,5 will evolve significantly during its remaining protostellar phase is not so clear. In the situation where the binary system dominates in mass over its circumbinary disk (so that, even at apastron, the binary orbits faster than the circumbinary disk rotates), as must be the case here, theoretical simulations show that tidal interactions between the binary system and its (coplanar) circumbinary disk result in a transfer of both angular momentum and energy from the binary to its circumbinary disk \citep{Artymowicz1991, Bate2000}. The effect is to drive up the orbital eccentricity, opposite to the situation described above where the envelope dominates in mass over the binary fragments. By preserving the apastron distance, $a(1+e)$, the effect also is to decrease the orbital semimajor axis. Thus, L1551\,IRS\,5 will likely evolve into a solar-mass binary system that has roughly equal companion masses, quite closely coplanar circumstellar disks and orbital planes (especially if tidal interactions between one protostar and the circumstellar disk of the other protostar bring their circumstellar disks further into alignment with the orbital plane), and hence presumably also closely aligned stellar spin and orbital axes. Unless its orbital parameters are dramatically altered during its subsequent evolution, L1551\,IRS\,5 will retain a likely circular or low eccentricity orbit with a semimajor axis of many tens to perhaps a few hundred AU, and a corresponding orbital period of many hundreds to perhaps a few thousand years. The physical properties that we project L1551\,IRS\,5 will have in its final state closely resemble those observed for binary systems on the pre-main-sequence and main-sequence having similar masses and orbital separations \citep[e.g., see review by][]{Duchene2013}. Such systems do not show a preferred mass ratio or orbital eccentricity, but do exhibit preferentially aligned circumstellar disks on the pre-main-sequence and preferentially aligned stellar spin and orbital axes on the main-sequence. For example, the circumstellar disks of T\,Tauri star binaries having projected orbital separations spanning the range $\sim$100--1000\,AU are preferentially aligned to within a few tens of degrees \citep[e.g.,][]{Wolf2001, Jensen2004, Monin2006}. Main-sequence solar-type {binaries} having orbital separations smaller than about 30--40\,AU (close to the median orbital separation) have approximately aligned spin and orbital axes, whereas those with larger orbital separations commonly show misaligned spin and orbital axes \citep{Hale1994}. Many of the systems that are closely aligned may have formed in a manner similar to L1551\,IRS\,5, rather than having being brought into alignment by internal or external interactions.

1607.00323

1607

1607.07735_arXiv.txt

We report on simultaneous X-ray and radio observations of the mode-switching pulsar \psr\ obtained with the {\it XMM-Newton} satellite and the LOFAR, LWA and Arecibo radio telescopes in November 2014. We confirm the synchronous X-ray/radio switching between a radio-bright (B) and a radio-quiet (Q) mode, in which the X-ray flux is a factor $\sim2.4$ higher than in the B-mode. We discovered X-ray pulsations, with pulsed fraction of 38$\pm$5\% (0.5-2 keV), during the B-mode, and confirm their presence in Q-mode, where the pulsed fraction increases with energy from $\sim$20\% up to $\sim$65\% at 2 keV. We found marginal evidence for an increase in the X-ray pulsed fraction during B-mode on a timescale of hours. The Q-mode X-ray spectrum requires a fit with a two-component model (either a power-law plus blackbody or the sum of two blackbodies), while the B-mode spectrum is well fit by a single blackbody (a single power-law is rejected). With a maximum likelihood analysis, we found that in Q-mode the pulsed emission has a thermal blackbody spectrum with temperature $\sim3.4\times10^6$ K and the unpulsed emission is a power-law with photon index $\sim$2.5, while during B-mode both the pulsed and unpulsed emission can be fit by either a blackbody or a power law with similar values of temperature and photon index. A {\it Chandra} image shows no evidence for diffuse X-ray emission. These results support a scenario in which both unpulsed non-thermal emission, likely of magnetospheric origin, and pulsed thermal emission from a small polar cap ($\sim$1500 m$^2$) with a strong non-dipolar magnetic field ($\sim10^{14}$ G), are present during both radio modes and vary in intensity in a correlated way. This is broadly consistent with the predictions of the partially screened gap model and does not necessarily imply global magnetospheric rearrangements to explain the mode switching.

Shortly after their discovery, radio pulsars were convincingly interpreted as rapidly rotating neutron stars with a very strong magnetic field, whose rotation and magnetic axes may differ. Radio pulsar profiles were early found to be very stable in time, so it was surprising to find that some were not, even assuming several stable forms, along with nulling, drifting and other pulse-sequence phenomena \citep[e.g.][]{ran86}. We now know that radio pulsar emission displays a wide range of variations on almost all intensity and time scales, from sparse bursts or nulling, to multi-decade fluctuations. Remarkably, some objects exhibit mode changes (or switches): transitions between otherwise stable states with distinct pulse shapes, flux densities, polarization properties, and sometimes different slow-down rates \citep{kra06,lyn10}. The study of these sources is very important, as they provide glimpses into the dynamics of the neutron star magnetosphere and the poorly understood physics of the pulsar radio emission \citep[e.g.,][]{sob15}. Here we concentrate on \psr\ which, being the prototypical mode-switching radio pulsar, has been extensively studied in the radio band and is a key target to investigate in more details the high-energy variability of pulsars. Its timing parameters ($P=1.1$ s, $\pdot$=3.5$\times$10$^{-15}$ s s$^{-1}$) imply, under the usual assumptions, a characteristic age of $\tau$=$P/(2\pdot$) = 5 Myr, a dipolar surface magnetic field $B$=4$\times$10$^{12}$ G, and a rotational energy loss rate $\dot{E}_{rot}$ = 10$^{32}$ erg s$^{-1}$. Detailed modelling of the radio pulse profiles and polarization indicate that \src\ is a nearly aligned rotator (angle between the rotation and magnetic axis $\sim$15$^{\circ}$) seen nearly pole-on \citep{des01}. Its distance, based on the dispersion measure, and the Galactic electron density distribution of \citet{cor02}, is $\sim$630 pc. In the radio band, \src\ alternates between two different states: when it is in the so-called B (bright) mode, the radio emission displays a regular pattern of drifting subpulses, while it is chaotic, and on average fainter, when in the Q (quiescent) mode \citep{sul84,ran06}. The phenomenon of drifting subpulses, observed in many radio pulsars, is believed to originate from a system of sub-beams of radio emission rotating around the magnetic axis \citep{rud75}. The existence of two modes of emission in \src\ indicates that such a structure is subject to some instability of unclear origin. Two short \xmm\ observations, carried out in 2003, showed that \src\ is a faint X-ray source, with a 0.5-8 keV flux of $\sim5\times10^{-15}$ erg cm$^{-2}$ s$^{-1}$ \citep{zha05}. For a distance of 630 pc (which we adopt in the whole paper), this corresponds to a luminosity L$_X\sim2\times10^{29}$ erg s$^{-1}$, and implies an X-ray efficiency in line with that of other rotation-powered pulsars of comparable characteristic age \citep{pos12}. A deeper study of the X-ray properties of \psr\ was performed by \citet{her13}, who used five \xmm\ pointings, supplemented by simultaneous radio observations with LOFAR and the GMRT. This first multi-wavelength campaign was carried out in 2011 November-December and provided a useful exposure of about 100 ks. Most importantly, using the mode-change times derived from the radio observations, it was possible to analyze separately the X-ray data of the Q and B modes. Quite surprisingly, it was discovered that the X-ray properties in the two modes are different. The X-ray flux is larger by more than a factor two during the Q-mode (when the radio flux is lower by roughly a factor of two at low frequencies; \citet{sul84}). X-ray pulsations at the rotation period of 1.1 s were detected for the first time, but only during the Q-mode. The Q-mode X-ray spectrum was well fit by the sum of a blackbody and a power law, with single component models clearly rejected. The spectrum of the fainter B-mode was less constrained and could be described equally well by either a power law or a blackbody. From the analysis of the pulsed spectrum in the Q-mode and the timing properties, \citet{her13} concluded that, during the B-mode, \src\ emits only an unpulsed non-thermal component and that the higher flux in Q-mode is caused by the addition of a thermal component with a 100\% pulsed fraction. This interpretation challenges the geometry of \psr\ derived from the radio observations, which predicts a smaller modulation of the thermal emission observed from the hot polar cap. According to \citet{sto14}, a strongly modulated thermal component can be obtained with beamed emission from a magnetic atmosphere or with an offset dipole geometry. \citet{her13} proposed instead an interpretation based on time-dependent scattering or absorption in the magnetosphere. This requires some global and rapid rearrangement of the pulsar magnetosphere to explain the different X-ray properties in the two modes. A reanalysis of the 2011 \xmm\ observations was carried out by \citet{mer13}, who concluded that the data are also consistent with the possibility that a constant, or slightly modulated thermal emission is present in both modes and the flux increase in the Q-mode is caused by the appearence of a pulsed non-thermal component. To study the remarkable correlated X-ray/radio variability of \psr\ in more detail, and possibly distinguish between the different interpretations, we obtained new X-ray observations in November 2014, within an \xmm\ Large Program with simultaneous radio monitoring provided by the Low-Frequency Array (LOFAR), Long Wavelength Array (LWA), and Arecibo radio telescopes. To assess the possible contribution of a pulsar wind nebula to the unpulsed non-thermal X-ray emission seen in \psr , we also obtained the first high spatial resolution X-ray image of this pulsar using the {\it Chandra X-ray Observatory}. In this paper we concentrate on the results from the new X-ray data. Further studies of the radio observations, as well as a joint analysis of the whole \xmm\ data set will be presented in future works. All the errors are at 1 $\sigma$, unless specified differently. \begin{figure*} \includegraphics[angle=-90,width=20cm]{B0943_Session5_X-ray_Radio_LCs.pdf} \vspace{-2.5cm} \caption{\label{fig_radio_2014} Radio and X-ray lightcurves for Session 5. {\it Top panel}: \xmm\ EPIC pn counts from \psr\ in the 0.2-10 keV band, binned in 3670-s intervals and corrected for the exposure reduction due to the exclusion of high-background periods. The radio-identified B/Q-modes are shown by the shaded regions. The dashed lines show the average B/Q-mode count levels determined from the entire 2014 data set. {\it Bottom panel}: Radio pulse profiles, each 300\,s, zoomed-in on a narrow range around the main pulse peak. The sudden transition in profile shape and noise properties at MJD = 56984.35 is when observing coverage shifted from LOFAR (observing at $\sim 150$\,MHz) to LWA (observing at $\sim 60$\,MHz). Note that the observed brightness of PSR~B0943+10 is modulated by both intrinsic effects (the mode switching) as well as the effective sensitivity of the telescope due to the changing source elevation. The start/stop times of B/Q-modes are indicated by the green/red ticks at the top/bottom of the panel. The vertical blue bars indicate {\it XMM-Newton} observation start/stop times. } \end{figure*}

Thanks to the long duration of the \xmm\ Large Program with simultaneous LOFAR, LWA and Arecibo radio monitoring carried out in 2014, we obtained several new results on the X-ray emission of the prototypal bimodal radio pulsar \src . Though the cause for the X-ray variability correlated with the radio modes remains unknown, we could explore in much more detail some of the scenarios that were proposed to explain the remarkable X-ray spectral and timing properties of this pulsar. In particular, the discovery of pulsations in the B-mode and the failure of a single power-law to fit the B-mode spectrum rule out simple interpretations which tried to explain the difference between the two modes as a change in a single pulsed component, present only during the Q-mode \citep{her13,mer13}. We showed that the situation is indeed more complex, and propose a consistent picture in which pulsed thermal and unpulsed non-thermal emission are simultaneously present in both radio modes and vary in a correlated way. Such a correlation is not surprising in the framework of the Vacuum Gap models discussed above, since the pairs are produced relatively close to the star surface and the accelerated particles are responsible both for the non-thermal emission in the magnetosphere and for the heating of the polar regions through return currents. A quantitative assessment of the relation between thermal and non-thermal emission by comparing different neutron stars is complicated by the presence of other factors, e.g. orientation, magnetic field strength, age, which introduce a variance in the observed properties. In the case of \psr\ we have instead the unique possibility to examine the relation between these components without such complications. On the other hand, Space Charge Limited Flow models with particles leaving freely the star surface \citep[e.g.,][]{aro79,zha00}, predict the pair-creation front much higher in the magnetosphere. Since in this case only a few percent of the return particles might reach the surface, one would not naturally expect a correlation beteween the thermal and non-thermal emission. Finally, we found some evidence for an evolution of the timing properties during the B-mode, which, if confirmed, can provide another important diagnostic to study the correlation between radio and X-ray properties in \psr\ and shed light on the physical processes responsible for the mode switching behavior. \\

1607.07735

1607

1607.03116_arXiv.txt

{Mergers of two stellar origin black holes are a prime source of gravitational waves and are under intensive investigation. One crucial ingredient in their modeling has been neglected: pair-instability pulsation supernovae with associated severe mass loss may suppress the formation of massive black holes, decreasing black hole merger rates for the highest black hole masses. } {We demonstrate the effects of pair-instability pulsation supernovae on merger rate and mass using populations of double black hole binaries formed through the isolated binary classical evolution channel. } {The mass loss from pair-instability pulsation supernova is estimated based on existing hydrodynamical calculations. This mass loss is incorporated into the {\tt StarTrack} population synthesis code. {\tt StarTrack} is used to generate double black hole populations with and without pair-instability pulsation supernova mass loss. } {The mass loss associated with pair-instability pulsation supernovae limits the Population I/II stellar-origin black hole mass to $50\msun$, in tension with earlier predictions that the maximum black hole mass could be as high as $100\msun$. In our model, neutron stars form with mass $1$--$2\msun$, then we encounter the first mass gap at $2$--$5\msun$ with an absence of compact objects due to rapid supernova explosions, followed by the formation of black holes with mass $5$--$50\msun$, with a second mass gap at $50$--$135\msun$ created by pair-instability pulsation supernovae and by pair-instability supernovae. Finally, black holes having masses above $135\msun$ may {\em potentially} form to arbitrarily high mass limited only by the extent of the initial mass function and the strength of stellar winds. Suppression of double black hole merger rates by pair-instability pulsation supernovae is negligible for our evolutionary channel. Our standard evolutionary model with the inclusion of pair-instability pulsation supernovae and pair-instability supernovae is fully consistent with the LIGO observations of black hole mergers: GW150914, GW151226, and LVT151012. The LIGO results are inconsistent with high ($\gtrsim 400$ km s$^{-1}$) BH natal kicks. We predict the detection of several, and up to as many as $\sim 60$, BH-BH mergers with a total mass of $10$--$150\msun$ (most likely range: $20$--$80\msun$) in the forthcoming $\sim 60$ effective days of the LIGO O2 observations, assuming the detectors reach the \emph{optimistic}\/ target O2 sensitivity. }

In September 2015 the upgraded Laser Interferometer Gravitational-wave Observatory (LIGO) began observations with a sensitivity to the merger of two neutron stars to an average distance of $d_{\rm nsns} \approx 70$ Mpc. In this first upgraded science run (O1; \citet{LigoO1a,LigoO1b}) LIGO made two firm detections of black hole-black hole (BH-BH) mergers with component masses of $36\msun$ and $29\msun$ (GW150914), and $14\msun$ and $8\msun$ (GW151226), and also reported a candidate BH-BH merger with masses of $23\msun$ and $13\msun$ (LVT151012). These discoveries verified earlier predictions that: \emph{(i)} the first detection would happen when LIGO sensitivity reaches $d_{\rm nsns}=50$--$100$ Mpc, that \emph{(ii)} BH-BH mergers would dominate the gravitational-wave signal, and that \emph{(iii)} the merging black holes would be substantially more massive than typical $10\msun$ Galactic BHs \citep{Belczynski2010a,Dominik2015,Belczynski2016a}. BH-BH mergers have been proposed as potential gravitational wave sources since the 1980s \citep{Bond1984,Thorne1987,Schutz1989} and have been studied since the 1990s \citep{Tutukov1993,Lipunov1997,Flanagan1998}. More recently, a number of groups have provided evolutionary models leading to potential BH-BH formation in a typical Galactic environment with high metallicity stars \citep{Brown2001,Nelemans2001,Belczynski2002,Voss2003,Postnov2006}. Subsequently it was shown that BH-BH merger formation in the Galactic environment, with its high metallicity, leads to suppression of BH-BH merger formation \citep{Belczynski2007,Mennekens2014}; a finding that was not encountered in previous studies. Only later was it noted that low metallicity stars will dominate the formation of massive BHs \citep{Belczynski2010b} and BH-BH mergers in general \citep{Belczynski2010a}. Full scale predictions that took into account low metallicity stars were performed in advance of the 2015 LIGO detections \citep{Dominik2012,Dominik2013,Dominik2015,Rodriguez2015,Marchant2016, Belczynski2016a,Mandel2016}. Additionally, Population III stars have also been considered as possible venues for BH-BH formation and gravitational wave detection for over thirty years \citep{Bond1984,Belczynski2004,Kinugawa2014}. Finally, prior to the first BH-BH merger detections, Population I/II very massive stars ($>150\msun$; \citep{Crowther2010}) were also introduced into predictions of BH-BH merger rates \citep{Belczynski2014,Marchant2016}. Since the first detection of GW150914, several investigations have examined how double black hole binaries could have been produced from the evolution of massive stars, whether from classical isolated evolution in low-metallicity environments \citep{Belczynski2016b,Eldridge2016}; via the aid of rapid rotation and hence homogeneous chemical evolution \citep{deMink2016, Woosley2016}; via Population III stars \citep{Hartwig2016,Inayoshi2016,Dvorkin2016}; or from dynamical formation in interacting environments \citep{Mapelli2016,Rodriguez2016}. Other more exotic scenarios have been introduced and discussed in the context of GW150914; dark matter primordial BH-BH formation \citep{Sasaki2016,Eroshenko2016}, formation of a BH-BH merger from a divided core of a massive rapidly rotating single star \citep{Loeb2016}, or formation of BH-BH mergers with disks around BHs formed from fallback material in weak supernova explosions \citep{Perna2016}. In this study we consider the effects of Pair-instability Pulsation Supernovae (PPSN) and Pair-instability Supernovae (PSN) on BH-BH mergers. PPSN are associated with severe mass loss \citep{Heger2002,Woosley2007} that may significantly reduce BH mass and thus detectability of BH-BH mergers. PSN are expected to completely disrupt massive stars with no resulting BH formation \citep{Bond1984,Fryer2001,Chatzopoulos2012a} and thus suppress formation of BH-BH mergers. While PSN are taken into account in some of the predictions for BH-BH merger formation (e.g., \cite{Marchant2016,Mandel2016,Spera2016}), PPSN and associated mass loss have thus far been ignored in studies of BH-BH formation (e.g., \cite{Dominik2015, Rodriguez2015,Belczynski2016a,Marchant2016,Mandel2016,Rodriguez2016,deMink2016, Belczynski2016b,Eldridge2016}) with the exception of recent work by \cite{Woosley2016}. We quantify the effect of PPSN and PSN on BH-BH mergers in our isolated classical binary evolution channel. In brief, these processes introduce a maximum mass of stellar-origin black holes which differs from previous projections; compare to, e.g., the review in \cite{AstroPaper}.

We have incorporated pair-instability pulsation supernovae and pair-instability supernovae into predictions of double compact object merger rates and masses in context of near future LIGO observations. We find that; \begin{enumerate} \item The mass of Population I/II stellar-origin black holes is limited to $50\msun$ by severe mass loss imposed by pair-instability pulsation supernovae (see Fig.~\ref{bhmass}). This may be contrasted with earlier predictions that the maximum mass of black holes can reach $80$--$130\msun$ in the evolution of Population I/II stars with modest initial masses: $M_{\rm zams}<150\msun$~\citep{Zampieri2009,Mapelli2009,Belczynski2010b, Spera2015,Spera2016}. This conclusion applies to black holes formed below the second mass gap (no compact objects in the mass range: $50$--$135\msun$; see Sec.~\ref{ex1}) imposed by pair-instability pulsation supernovae and pair-instability supernovae. If stars reach high enough mass to avoid disruption by pair-instability supernovae (i.e., if they can form helium cores above $135\msun$) then black holes with mass above $135\msun$ may form. Such massive black hole formation would require any combination of: very high star mass $>200$--$300\msun$ (whether it is initial mass or mass of a stellar merger), or very low metallicity (e.g., Population III stars), or very rapid rotation (homogeneous evolution). \item We show that the introduction of pair-instability pulsation supernovae and the associated mass loss does not affect our predictions for detection of NS-NS, BH-NS and BH-BH mergers during the LIGO O2 observational run. In particular, our isolated binary classical evolution channel produces a similar number of detections for the O2 run whether or not pair-instability pulsation supernovae and pair-instability supernovae are included; $\sim 60$ BH-BH merger detections with a total redshifted mass in the range $10$--$150\msun$. Detections of BH-NS and NS-NS mergers originating from our classical isolated binary evolution model are not very likely in O2 (see Tab.~\ref{tab1}). We also note that the detection rates may be significantly smaller if pessimistic assumptions on binary evolution are adopted (i.e., model M3). To demonstrate this we have allowed for high black hole and neutron star natal kicks to obtain: only $\sim 2$ BH-BH merger detections in the entire O2 run. Since this model is just below the current LIGO empirical BH-BH merger rate estimate, it may serve as a lower limit on the number of predicted detections during O2. However, note that we use the optimistic target O2 sensitivity in all our predictions. \item The detection of very massive BH-BH mergers ($M_{\rm tot,z}>150$--$200\msun$; see Fig.~\ref{final}) could distinguish between models with and without pair-instability pulsation supernovae and pair-instability supernovae. However, our results argue that such a detection is unlikely. A detection of any binary with a BH mass $M_{\rm BH}>50\msun$ will rule out our adopted model for mass loss by pair-instability pulsation supernovae. Such an observation would require reconsideration of physics currently believed to be driving pair-instability pulsation supernovae and pair-instability supernovae (see Sec.~\ref{ex1}). An alternate solution for the detection of a massive BH ($M_{\rm BH}>50\msun$) is that the massive BH was formed through dynamical interactions. Any dynamical interaction that increases BH mass (either merger of two BHs, or rapid accretion onto a BH in tidal disruption event) can potentially accomplish this. For example, the merger of two lighter BHs (first burst of gravitational waves) may form a massive BH. This massive BH can then undergo a dynamical capture/exchange in a dense stellar environment (e.g., in a globular cluster) placing it in a new, massive binary. This binary generates the second BH-BH merger, in which one BH is very massive. In this scenario, the capture/exchange rate may be limited by the first merger natal kick that could potentially remove the massive BH from a cluster environment (Giersz et al. 2015). There is so far no published probability/rate estimate for such a specific scenario. \end{enumerate}

1607.03116

1607

1607.08518_arXiv.txt

We present the identification of 634 variable stars in the Milky Way dSph satellite Sculptor based on archival ground-based optical observations spanning $\sim$24 years and covering $\sim$ 2.5 deg$^2$. We employed the same methodologies as the ``Homogeneous Photometry'' series published by Stetson. In particular, we have identified and characterized one of the largest (536) RR~Lyrae samples so far in a Milky Way dSph satellite. We have also detected four Anomalous Cepheids, 23 SX~Phoenicis stars, five eclipsing binaries, three field variable stars, three \textit{peculiar} variable stars located above the horizontal branch -- near to the locus of BL~Herculis -- that we are unable to classify properly. Additionally we identify 37 Long Period Variables plus 23 probable variable stars, for which the current data do not allow us to determine the period. We report positions and finding charts for all the variable stars, and basic properties (period, amplitude, mean magnitude) and light curves for 574 of them. We discuss the properties of the RR Lyrae stars in the Bailey diagram, which supports the coexistence of subpopulations with different chemical compositions. We estimate the mean mass of Anomalous Cepheids ($\sim$1.5M$_{\sun}$) and SX~Phoenicis stars ($\sim$1M$_{\sun}$). We discuss in detail the nature of the former. The connections between the properties of the different families of variable stars are discussed in the context of the star formation history of the Sculptor dSph galaxy.

\label{sec:introduction} Pulsating variable stars are powerful tools to investigate the evolution of their host galaxy, as they trace the age and the metallicity of the parent population. Most importantly, the coexistence of different types of variable stars provides, thanks to their pulsational properties, independent constraints not only on the star formation history and the chemical evolution, but also on the distance of the system. Indeed, because pulsations occur at specific evolutionary stages that depend on the stellar mass, variable stars trace the spatial distribution of stellar populations of given ages. Therefore they can be used as markers of spatial trends across the galaxy under examination \citep[e.g.,][]{Gallart2004}. Moreover, even the range of pulsational properties among individual stars of a particular type can trace some differences in the age and metallicity of the corresponding population \citep[e.g.,][]{Bernard2008,MartinezVazquez2015}. This paper focuses on the variable-star content of the Local Group dwarf spheroidal (dSph) Sculptor. Sculptor is one of the ``classical'' Milky Way dSph satellites. After the Magellanic Clouds, it was the first to be discovered along with Fornax \citep{Shapley1938}. Sculptor's stellar content has been investigated in a large number of papers, using different techniques. Large scale and/or deep photometric surveys provided colour-magnitude diagrams (CMDs) showing an extended horizontal branch \citep[HB,][]{Majewski1999,Hurley-Keller1999,Harbeck2001}, and a wide colour spread of the red giant branch first mentioned by \citet{DaCosta1984}. While it is well established that Sculptor is composed of a predominantly old population \citep{Monkiewicz1999,deBoer2011}, it clearly presents some age spread \citep{Tolstoy2004,deBoer2012}. The chemical enrichment history of Sculptor has been investigated through spectroscopy of its RGB \citep{Tolstoy2004,Kirby2009,Starkenburg2013,Skuladottir2015} and HB stars \citep{Clementini2005}, revealing a large range in metallicity, of the order of 1 dex. In \citet{MartinezVazquez2015} (hereafter \citetalias{MartinezVazquez2015}), based on the pulsational properties of RR Lyrae (RRL) stars, we showed that a similar metallicity spread ($\sim$ 0.8 dex) was already in place at an early epoch (>10 Gyr), imprinted in the parent population that we observe today as RRL stars. The first investigation of the variable-star content in Sculptor dates back to the work by \citet{Baade1939} and \citet{ Thackeray1950}. However it was not until \citet{vanAgt1978} that a conspicuous population of 602 candidate variable stars was discovered and periods for 64 of them were provided. The most complete catalogue of variable stars (in terms of providing pulsational properties) in Sculptor is that of \citet{Kaluzny1995}. They investigated the central region of the Sculptor dSph ($\sim$ 15$\arcmin \times$15$\arcmin$) as a side-program of the OGLE I project. They identified 231 variable stars that were classified as 226 RRL, 3 Anomalous Cepheids (AC), and 2 long period variable (LPV) stars. Their properties are consistent with a metal-poor population ([Fe/H] $<$ --1.7). A spectroscopic follow-up made by \citet{Clementini2005} confirmed this result through low resolution (R$\approx$800) spectroscopy of 107 variables using the $\Delta$S method. In particular, they found that the metallicity peaks at [Fe/H] $\sim$--1.8. In \citetalias{MartinezVazquez2015} we reported on the detection of a large metallicity spread and spatial gradients within the population of Sculptor's RRL star population. In this work we present the full catalog of variable stars detected in this galaxy employing the same methodologies as the ``Homogeneous Photometry'' series \citep{Stetson1998,Stetson2000,Stetson2003, Stetson2005a,Stetson2005b,Stetson2014}. In \S~\ref{sec:photometry} we present the extensive data set of 4,404 images used in this analysis. In \S~\ref{sec:cmd_var} we discuss the variable-star detection and classification. We later discuss in detail different families of variable stars: RRL stars (\S~\ref{sec:rrl}), AC (\S \ref{sec:acep}), SX Phoenicis (SX~Phe,\S~\ref{sec:sxpho}), and other groups (peculiar, binaries, long period and probable variable stars, \S~\ref{sec:others}). In \S~\ref{sec:discussion} we discuss the properties of the old populations of Sculptor, analysing its RRL stars in detail. A summary of our conclusions (\S~\ref{sec:conclusions}) closes the paper. We highlight that in the online version of the paper we provide full details on all the variable stars discussed: time series photometry, light curves, mean photometric and pulsational properties and finding charts.

1607.08518

1607

1607.06396_arXiv.txt

In this paper, we perform a full next-to-leading order (NLO) QCD calculation of neutralino scattering on protons or neutrons in the MSSM. We match the results of the NLO QCD calculation to the scalar and axial-vector operators in the effective field theory approach. These govern the spin-independent and spin-dependent detection rates, respectively. The calculations have been performed for general bino, wino and higgsino decompositions of neutralino dark matter and required a novel tensor reduction method of loop integrals with vanishing relative velocities and Gram determinants. Numerically, the NLO QCD effects are shown to be of at least of similar size and sometimes larger than the currently estimated nuclear uncertainties. We also demonstrate the interplay of the direct detection rate with the relic density when consistently analyzed with the program \DMNLO.

1607.06396

1607

1607.08668_arXiv.txt

{} {Previous literature suggests that the densest structures in the interstellar medium form through colliding flows, but patent evidence of this process is still missing. Recent literature proposes using SiO line emission to trace low-velocity shocks associated with cloud formation through collision. In this paper we investigate the bright and extended SiO(2--1) emission observed along the $\sim$5\,pc-long W43-MM1 ridge to determine its origin.} {We used high angular resolution images of the SiO(2--1) and HCN(1--0) emission lines obtained with the IRAM plateau de Bure (PdBI) interferometer and combined with data from the IRAM 30\,m radiotelescope. These data were complemented by a \textit{Herschel} column density map of the region. We performed spectral analysis of SiO and HCN emission line profiles to identify protostellar outflows and spatially disentangle two velocity components associated with low- and high-velocity shocks. Then, we compared the low-velocity shock component to a dedicated grid of one-dimensional (1D) radiative shock models. } {We find that the SiO emission originates from a mixture of high-velocity shocks caused by bipolar outflows and low-velocity shocks. Using SiO and HCN emission lines, we extract seven bipolar outflows associated with massive dense cores previously identified within the W43-MM1 mini-starburst cluster. Comparing observations with dedicated Paris-Durham shock models constrains the velocity of the low-velocity shock component from 7 to 12\,km.s$^{-1}$. } {The SiO arising from low-velocity shocks spreads along the complete length of the ridge. Its contribution represents at least 45\% and up to 100\% of the total SiO emission depending on the area considered. The low-velocity component of SiO is most likely associated with the ridge formation through colliding flows or cloud-cloud collision.}

High-mass stars (OB-type, $\geq$8\,M$_\odot$) are known to form in massive dense cores \citep[MDCs, diameter of $\sim$\,0.1\,pc and density $>$\,10$^5$\,cm$^{-3}$; see, e.g.,][]{motte07}. But the formation of MDCs, as well as the physical process by which high-mass stars form within MDCs, are still badly constrained. The formation of high-mass stars is explained in two ways: quasi-static and dynamical scenarios. The quasi-static scenario describes a monolithic collapse of the MDC that is supported by supersonic turbulent pressure \citep[e.g.,][]{mckee2002}. However, the dynamical scenarios suggest either a coalescence of low/intermediate-mass protostars \citep[e.g.,][]{bonnell-bate02} or a formation of high-mass stars through the interplay of colliding flows associated with the cloud formation \citep[e.g.,][]{bonnell-bate02, heitsch08}. In this latest dynamical picture, cloud formation generates colliding flows that funnel mass from large potentials to small scales \citep[e.g.,][]{vazquezsemadeni05, hartmann12, smith13}. While the gas flows are colliding, we expect low-velocity shocks at the colliding interfaces. At the MDC scale ($\sim$0.1\,pc) dynamical signatures, such as gravitational streams and shearing motions, have been found using N$_2$H$^+$ and H$^{13}$CO$^+$ lines \citep[see, e.g.,][]{csengeri11a,csengeri11b,henshaw14}. At that scale, low-velocity shocks have also been reported using CH$_3$CN and SiO lines \citep[e.g.,][]{csengeri11b, duarte-cabral14}. The \textit{Herschel} key program HOBYS \cite[see][]{motte10,motte12} demonstrated that MDCs, and later high-mass stars, tend to form in high-density filaments with typical sizes of $\sim$1--10\,pc, above 10$^{23}$\,cm$^{-2}$ in column density. These structures are called ridges \citep{hill11, quang11-w43, martin12}. On the scale of ridges, gas inflow, global collapse, and velocity gradients have been reported by several groups \citep[e.g.,][]{schneider10,quang13,peretto13,tackenberg14}. However, all the studies suggesting that SiO emissions could trace low-velocity shocks associated with the collision of gas inflows \citep{jimenez10, quang11-w43, quang13, sanhueza13} were impaired by their limited sensitivity and/or angular resolution. All of these studies reported the SiO emission line profile as the association of two Gaussian components at approximately the same central velocity. While the broad Gaussian component was attributed to high-velocity shock linked to protostellar outflows, the narrow Gaussian component was equally attributable to low-velocity shock of either cloud-cloud collision or less powerful outflows beyond the angular resolution performances. \begin{table*}[htbp!] \caption{Main observational parameters} \label{t:pariram} \begin{center} \begin{tabular}{c|cc|cc} \hline \hline Parameter & \multicolumn{2}{c}{IRAM PdBI} & \multicolumn{2}{c}{IRAM 30\,m} \\ & SiO (2-1) & HCN (1-0) & SiO (2-1) & HCN (1-0) \\ \hline Frequency (GHz) & 86.846 & 88.631 & 86.846 & 88.631 \\ Bandwidth (MHz) & 40 & 80 & -- & -- \\ Spec. Res. (km.s$^{-1}$) & 0.27 & 1.06 & 0.67 & 0.67 \\ Primary beam & 59\arcsec & 59\arcsec & 29.9\arcsec & 29.9\arcsec \\ Synthesized beam & $4\farcs96\times3\farcs14$ & $5\farcs02\times3\farcs08$ & - & - \\ $1\sigma$ rms & 0.20 Jy/Beam.km/s & 0.49 Jy/Beam.km/s & 0.02 K.km/s & 0.04 K.km/s \\ System temperature (K) & 120 & 120 & - & - \\ \hline \end{tabular} \end{center} \end{table*} We were able to disentangle these broad and narrow Gaussian components in the most extreme of the ridges: \object{W43-MM1}. This ridge lies in the W43 molecular complex \citep{quang11-w43,carlhoff13, motte14}. A distance from the sun of $5.5\pm0.4$\,kpc was inferred for this complex by parallax measurements \citep{zhang14}. Given their uncertainties, and since the measurements were based on only four sources located at the periphery of the complex, we adopted a round 6\,kpc distance, which is consistent with \cite{louvet14}. \cite{louvet14} modeled \object{W43-MM1} using a 3.9~pc$\times$2~pc$\times$2~pc ellipsoid with a total mass of $\sim$2$\times 10^4$~M$_\odot$ and an average density of $\sim$4.3$\times 10^4$ cm$^{-3}$. This ridge hosts a protocluster of 12 MDCs of mean size of $\sim$0.07\,pc and masses ranging from 20\,M$_\odot$ to 2000\,M$_\odot$ \citep{louvet14,sridharan14}. It undergoes a remarkable burst of (high-mass) star formation with an instantaneous star formation rate of $\sim$6000\,M$_{\odot}$.Myr$^{-1}$ \citep{louvet14}. As part of the W43-HERO IRAM large program, \cite{quang13} discovered a bright and extended SiO(2--1) emission with $N_{\rm SiO}\sim 6\times10^{13}$~cm$^{-2}$ over a $\sim$43~pc$^2$ area along and around \object{W43-MM1}. They interpreted this SiO emission as arising from a low-velocity cloud collision the ridge experienced during its formation, but their observations lacked the resolution to rule out an origin from a protocluster. In this paper, we show that SiO emission can be an excellent tracer of colliding flows. Making use of the interferometric and single-dish data cubes described in Sect.~\ref{s:obs}, we investigate the high-density parts of \object{W43-MM1}, look for outflows driven by protostars forming in the \object{W43-MM1} ridge, and quantify the SiO intensity unambiguously associated with low-velocity shocks in Sect.~\ref{s:analysis}. Section~\ref{s:model} confronts the integrated SiO emission values with dedicated shock models. Finally, Sect.~\ref{s:conclusion} gives our conclusions.

\label{s:conclusion} We mapped the high-mass star-forming ridge \object{W43-MM1} at high angular resolution with the IRAM Plateau de Bure Interferometer. We obtained SiO(2--1) and HCN(1--0) emission lines that we complemented with single-dish IRAM 30\,m radiotelescope data (see Fig.\ref{f:compa}). We used both transitions to look for outflows associated with star formation in \object{W43-MM1}. Our main results and conclusions may be summarized as follows: \begin{itemize} \item[$\bullet$] At high angular resolution, when observed with the IRAM PdBI interferometer, the SiO(2-1) emission appears very extended: $\sim$5\,pc in projected distance. \item[$\bullet$] We retrieve evidence of seven outflows that were all detected both in SiO(2-1) and HCN(1-0) line transitions. Six of these outflows are associated with massive dense cores previously identified by \cite{louvet14}, pointing to the protostellar nature of these objects. The other outflow is associated with a lower mass object, the source \emph{G} identified by \citet{sridharan14}. \item[$\bullet$] From the line profile analysis of the SiO(2-1) emission on the entire filament, we attribute half of the emission to high-velocity shocks linked with the stellar formation. The second half of the line emission appears to be very homogeneous along the \object{W43-MM1} filament with a mean FWHM of $\sim$6\,km.s$^{-1}$. \item[$\bullet$] We characterize the line emission profile by defining eight independent areas along the filament and discover that four of the positions are only composed of one narrow component line profile (width of $\sim$6\,km.s$^{-1}$), one is composed of one single broad component line profile (source N12 that displays an outflow), and three are composed of both broad\, and \,narrow component line profiles. In the latter case, the narrow component is generally similar to those observed at the narrow-only positions, suggesting a common origin for this narrow component. \item[$\bullet$] We run dedicated Paris-Durham shock models to confront the narrow component of the SiO(2-1) emission profile with SiO emission from low-velocity shocks. From this analysis, we show that the SiO emission intensity observed for the narrow component can be reproduced with shocks at low velocities in the range 7\,km.s$^{-1}$ to 12\,km.s$^{-1}$. These model-constrained velocities are in perfect agreement with the observational constraints of the SiO line width, which suggest velocities of shock between 4\,km.s$^{-1}$ to 14\,km.s$^{-1}$. \end{itemize}

1607.08668

1607

1607.01780_arXiv.txt

\sourcelong\ is a transient neutron star low-mass X-ray binary that exhibited a bright accretion outburst in 2015. We present \nustar, \swift, and \chan\ observations obtained around the peak brightness of this outburst. The source was in a soft X-ray spectral state and displayed an X-ray luminosity of $L_{\mathrm{X}}$$\simeq$$(2-3)\times10^{37}~\dist~\lum$ (0.5--10 keV). The \nustar\ data reveal a broad Fe-K emission line that we model as relativistically broadened reflection to constrain the accretion geometry. We found that the accretion disk is viewed at an inclination of $i$$\simeq$$27^{\circ}$--$35^{\circ}$ and extended close to the neutron star, down to $R_{\mathrm{in}}$$\simeq$5--7.5 gravitational radii ($\simeq$11--17~km). This inner disk radius suggests that the neutron star magnetic field strength is $B$$\lesssim$$2\times10^8$~G. We find a narrow absorption line in the \chan/HEG data at an energy of $\simeq$7.64~keV with a significance of $\simeq$4.8$\sigma$. This feature could correspond to blue-shifted Fe\,{\sc xxvi} and arise from an accretion disk wind, which would imply an outflow velocity of $v_{\mathrm{out}}$$\simeq$$0.086c$ ($\simeq$25\,800$~\kms$). However, this would be extreme for an X-ray binary and it is unclear if a disk wind should be visible at the low inclination angle that we infer from our reflection analysis. Finally, we discuss how the X-ray and optical properties of \sourcelong\ are consistent with a relatively small ($P_{\mathrm{orb}}$$\lesssim$3~hr) binary orbit.

\label{sec:introduction} Low-mass X-ray binaries (LMXBs) consist of a neutron star or a black hole plus a companion star that is less massive than the compact primary. The companion typically overflows its Roche-lobe, thereby feeding gas to an accretion disk that surrounds the neutron star or black hole. In transient systems, outbursts of active accretion are interleaved with long periods of quiescence when the accretion rate onto the compact primary is strongly reduced. LMXBs are excellent laboratories to investigate how accretion proceeds under the influence of strong gravity over a wide range of accretion rates. X-ray spectroscopy is a powerful tool to study accretion processes in LMXBs; both the X-ray continuum spectral shape and discrete emission/absorption features give insight into the accretion geometry. Broadly speaking, two main X-ray spectral states can be distinguished during the active phases of LMXBs \citep[see e.g.,][for a detailed overview]{homan2005_specstates,remillard2006}. During a ``soft state'', thermal emission from the accretion disk is prominently detected in the X-ray spectrum. However, during a ``hard state'', the disk emission is much weaker and a strong non-thermal, hard X-ray emission component is seen that likely arises from a hot plasma in the inner accretion flow. Hard X-rays can illuminate the accretion disk and produce a reflection spectrum of emission lines, the most prominent being Fe-K at 6.4--6.97~keV \citep[e.g.,][]{george1991,matt1991}. The shape of such lines depends on the chemical abundances and ionization state of the accretion disk, but also on its inner radial extent since both relativistic and dynamical effects act to broaden reflected emission lines. Disk reflection studies suggest that in some neutron star LMXBs the stellar magnetic field is strong enough to truncate the inner accretion disk \citep[][]{degenaar2014_groj1744,pintore2016,king2016}, whereas in others it can extend all the way down to the innermost stable circular orbit (ISCO) or the neutron star surface \citep[e.g.,][]{bhattacharyya2007,dai2009,papitto2009,cackett2010_iron,miller2013_serx1,degenaar2015_4u1608,disalvo2015}. High-resolution X-ray spectroscopy of LMXBs can reveal narrow, blue-shifted absorption features arising from a highly ionized wind that is blown off the accretion disk \citep[e.g.,][]{brandt2000,ueda2004,miller2006,miller2016_wind,diaztrigo2007,bozzo2016_wind}. These winds are preferentially detected during soft X-ray spectral states \citep[e.g.,][]{miller2006_winds,miller2008,neilsen2009,ponti2012_winds}. Typical outflow velocities are $\simeq$400--3000$~\kms$ ($\simeq$0.001--0.01$c$) and the properties are very similar in black hole and neutron star LMXBs \citep[e.g.,][]{diaztrigo2015}. Powerful disk winds may carry away a significant amount of mass, similar to or even exceeding the amount of mass that is accreted \citep[e.g.,][]{lee2002,neilsen2011,king2012,king2015,ponti2012_winds,degenaar2014_groj1744}. Winds are therefore important in considering the energy budget of the accretion process, and may also lead to instabilities in the accretion flow \citep[e.g.,][]{begelman1983b,shields1986}. \vspace{-0.2cm} \subsection{\sourcelong\ (\source)}\label{subsec:source} \sourcelong, hereafter referred to as \source, was detected as a faint X-ray source with \rosat\ in 1990 \citep[][]{voges1999}. It remained unclassified until 2012, when \inte\ detected a thermonuclear X-ray burst that identified \source\ as an accreting neutron star in, most likely, an LMXB. Assuming that the peak luminosity of the X-ray burst did not exceed the empirically-determined Eddington limit \citep[$L_{\mathrm{Edd}}$$=$$3.8\times10^{38}~\lum$;][]{kuulkers2003}, places the source at a distance of $D$$\lesssim$5.8~kpc \citep[][]{chenevez2012}. No persistent emission could be detected with \inte\ when the X-ray burst was seen in 2012, implying an accretion luminosity of $L_{\mathrm{X}}$$\lesssim 4 \times 10^{35}~\lum$ at the time (3--10~keV). Indeed, follow-up \swift\ observations detected the source at $L_{\mathrm{X}}$$\simeq$$10^{33}-10^{34}~\lum$ (0.5--10~keV), suggesting that it had exhibited a faint accretion outburst that was missed by all-sky monitors \citep[][]{chenevez2012,kaur2012_1804}. In 2015 January, an outburst from \source\ was seen by \swift/BAT and \maxi\ monitoring observations \citep[][]{krimm2015,negoro2015}. Figure~\ref{fig:bat} shows the publicly available \maxi\ \citep[2--20~keV;][]{maxi2009}\footnote{http://maxi.riken.jp/top/index.php?cid=1\&jname=J1804-343} and \swift/BAT \citep[15--50 keV;][]{krimm2013}\footnote{http://swift.gsfc.nasa.gov/results/transients/weak/1RXSJ180408.9-342058/} daily monitoring light curves of this outburst. The source remained at relatively constant flux in a hard X-ray spectral state \citep[][]{ludlam2016} for $\simeq$2~months, during which radio emission from a jet was detected \citep[][]{deller2015_1804}. Around April 3, however, the source transitioned to a soft X-ray spectral state \citep[][dotted line in Figure~\ref{fig:bat}]{degenaar2015_1804}. The activity ceased $\simeq$2~months later, i.e., the total outburst duration was $\simeq$4.5~months (see Figure~\ref{fig:bat}). In this work we report on X-ray spectral observations of \source\ obtained with \nustar\ \citep[][]{harrison2013_nustar}, \swift\ \citep[][]{gehrels2004}, and \chan\ \citep[][]{weisskopf2000}. These data were taken during the soft X-ray spectral state near the peak brightness of the 2015 outburst (see Table~\ref{tab:obs} and Figure~\ref{fig:bat}). \begin{figure} \begin{center} \includegraphics[width=8.5cm]{figure1.eps} \end{center} \caption[]{Combined \swift/BAT (red, 15--50 keV) and \maxi\ (black, 2--20 keV) X-ray light curves of the 2015 outburst of \source\ (2-day bins). The hard to soft state transition and the estimated return to quiescence are indicated by the vertical dotted lines. Times of the \chan, \nustar\ and \swift/XRT observations discussed in this work are indicated by the arrows. } \label{fig:bat} \end{figure}

\label{sec:discuss} We report on \nustar, \swift, and \chan\ observations obtained during a $\simeq$11-day window around the peak brightness of the 2015 outburst of the neutron star LMXB \source. The broad-band 0.7--35 keV \nustar/\swift\ spectral data can be described by a continuum model composed of a disk black body and thermal Comptonization. The source was in a soft spectral state during our observations and we measured a 0.5--10 keV luminosity of $L_{0.5-10}$$\simeq$$(2-3) \times 10^{37}~\dist~\lum$. Assuming a bolometric correction factor of 2--3 \citep[e.g.,][]{zand07}, this corresponds to $\simeq$10\%--25\% of the Eddington limit \citep[$L_{\mathrm{Edd}}$$=$$3.8\times10^{38}~\lum$;][]{kuulkers2003}. Superimposed on the continuum, we detect a broad Fe-K emission line that likely arises from hard X-rays reflecting off the accretion disk. Modeling this feature as relativistically blurred reflection allows us to constrain the accretion geometry in the soft state. We find a disk inclination of $i$$=$29$^{\circ}$--35$^{\circ}$ (1$\sigma$ confidence level), consistent with results obtained from reflection analysis in the hard state of the same \source\ outburst \citep[$i$$=$18$^{\circ}$--29$^{\circ}$ at 90\% confidence;][]{ludlam2016}. Our analysis suggests that the inner edge of the accretion disk extended inwards to $R_{\mathrm{in}}$$\simeq$5--7.5$~\rg$ ($\simeq$11--17~km). Bearing in mind the caveats stated in Section~\ref{subsec:refl}, we note that a similar inner disk radius is implied by the normalization of the disk black body component in our spectral fits ($R_{\mathrm{in}}$$\simeq$9.5--11~km). These values are consistent with the expected radius of a neutron star \citep[e.g.,][for a recent overview]{lattimer2014}, which suggests that the accretion disk may have been truncated by the stellar surface (and hence that the ISCO lies within the neutron star). The inner disk may have also been truncated by the magnetic field of the neutron star, rather than its surface. We can thus use our measured inner disk radius to estimate an upper limit for the magnetic field strength of the neutron star. To this end we use equation (1) of \citet{cackett2009_iron}, which was adapted from the formulation of \citet{ibragimov2009}. Extrapolating our \nustar/\swift\ spectral fit to 0.01--100 keV suggests a bolometric flux of $F_{\mathrm{bol}}$$\simeq$$1.2\times10^{-8}~\flux$. We assume $D$$=$5.8~kpc, $M$$=$$1.5~\Msun$, $R$$=$10~km and make similar assumptions regarding geometry and the accretion efficiency as in \cite{cackett2009_iron}; an anisotropy correction factor $f_{\mathrm{ang}}$$=$1, a geometry coefficient $k_{\mathrm{A}}$$=$0.5, and an accretion efficiency $\eta$$=$0.1. The constraint that $R_{\mathrm{in}}$$\lesssim$$7.5~Rg$, then yields $B$$\lesssim$$2\times10^{8}$~G for \source. \subsection{On the possible detection of a disk wind}\label{subsec:wind} We detect a narrow absorption line at $\simeq$7.64~keV in our \chan/HEG data that is $\simeq$4.8$\sigma$ significant after accounting for the number of trials. Other narrow absorption features were seen in the \chan/HEG data and two of the \swift/XRT soft-state observations at energies of $\simeq$7--8 keV, but at lower significance ($\simeq$2--4$\sigma$ after accounting for trials). Blue-shifted, narrow absorption lines have been seen in several black hole and neutron star LMXBs during their soft states, and are interpreted as outflowing disk winds \citep[e.g.,][]{brandt2000,lee2002,ueda2004,miller2006,miller2008,miller2011,miller2016_wind,diaztrigo2007,neilsen2009,neilsen2011,king2012,king2015,ponti2012_winds,degenaar2014_groj1744,bozzo2016_wind}. Since \source\ was in a soft state during the observations analyzed in this work, it is plausible that a disk wind was present. Assuming that the 7.64~keV line in our \chan\ data is both real and due to a disk wind, our photo-ionization modeling suggests that the most likely identification is blue-shifted Fe\,{\sc xxvi}, which would imply an outflow velocity of $\simeq$$0.086c$$\simeq$$ 25\,800~\kms$. However, we note that a solid identification is not trivial with just a single line. Moreover, there are several concerns about interpreting this feature as being due to a fast, outflowing disk wind. Firstly, observational evidence suggests that disk winds are concentrated in the accretion disk plane \citep[e.g.,][]{miller2006_winds,miller2006,miller2015_winds,king2012,ponti2012_winds,degenaar2014}. Indeed, simulations of (thermally-driven) disk winds suggest that at near-polar angles, the low density and high ionization of the gas may prevent its detection \citep[e.g.,][]{sim2008,higginbottom2015}. However, our spectral analysis suggests that \source\ is seen at a relatively low inclination of $i$$\simeq$$30^{\circ}$ \citep[see also][]{ludlam2016}. It is therefore not clear whether a disk wind, if present, would be observable. Nevertheless, a signature of a wind was recently reported for the neutron star LMXB GX 340+0, which has a similarly low inclination of $i$$\simeq$$35^{\circ}$ \citep[][]{miller2016_wind}. It was proposed that in this source the wind may be radiatively or magnetically driven \citep[see e.g.,][for a discussion on wind driving mechanisms]{diaztrigo2015}. Indeed, simulations of magnetically-driven winds in LMXBs suggest that these may be observable at lower inclination angles \citep[e.g.,][]{chakravorty2015}. Secondly, the implied outflow velocity of $\simeq$$25\,800~\kms$ ($\simeq$0.086$c$) would be extreme for an LMXB. Such fast outflows are not uncommon in supermassive black holes \citep[e.g.,][for a recent review]{tombesi2016}, but wind velocities tend to be more modest in LMXBs. The highest velocities claimed for LMXBs so far are $\simeq$$9\,000-14\,000~\kms$ \citep[$\simeq$0.03--0.04$c$;][]{king2012,degenaar2014_groj1744,miller2016_wind}, but more typical values are $\simeq$$400-3\,000~\kms$ \citep[$\simeq$0.001--0.01$c$; e.g.,][for an overview]{diaztrigo2015}. Wind velocities in LMXBs are thus commonly a factor $>$10 lower than implied for \source, and even the most extreme cases so far have a factor $>$3 lower outflow velocities than we find. Although there is reason to be skeptical about the detection of the $\simeq$7.64~keV line, we briefly speculate on its possible implication assuming that it is both real and that its interpretation as blue-shifted Fe\,{\sc xxvi} is correct. Observations seem to suggests that winds and jets generally do not co-exists, and leave room for the possibility that winds and jets perhaps evolve into one another \citep[e.g.,][]{miller2006_winds,neilsen2009,king2012,ponti2012_winds,diaztrigo2013_nature}. If so, winds may accelerate and become denser away from the disk plane when moving toward lower luminosity. This could perhaps account for detecting a fast outflow from \source\ in our \chan\ observations, as it was accreting at a relatively low luminosity of $L_{0.5-10}$$\simeq$$2 \times 10^{37}~\dist~\lum$ at the time. Within this context it is interesting to note that \source\ was detected at radio wavelengths on 2015 April 12, just two days before our \chan\ observation, albeit with a much lower flux density than during the preceding hard spectral state (Gusinskaia et al., in preparation). It is also possible that the neutron star magnetic field acts as a propeller and ejects material. MHD simulations of \citet{romanova2009} reveal both a jet-like and a wind-like outflow for an active propeller. In these simulations the wind component is a thin conical shell with a half-opening angle of $\simeq$$30^{\circ}$--$40^{\circ}$ and a velocity up to $\simeq$0.1$c$. Such an outflow may be compatible with our X-ray observations of \source\ and the quasi-simultaneous weak radio detection. In particular, the putative X-ray absorption lines are very narrow and it would require a thin ejected shell to be internally consistent with the radial velocity implied by the blue-shift. A propeller operates when the magnetospheric radius is larger than the co-rotation radius (i.e., the radius at which the Keplerian frequency of the disk is equal to the spin frequency of the neutron star). If the inner disk in \source\ were to be truncated by the magnetosphere and this radius ($R_{\mathrm{in}}$$\lesssim$17~km) were to be larger than the co-rotation radius, then the neutron star would have to be spinning at $\lesssim$0.4~ms. However, no neutron stars with sub-millisecond spin periods are known to date \citep[e.g.,][]{patruno2012_gravwav}. In summary, there is an absorption feature significantly detected at $\simeq$7.64~keV in our \chan/HEG data that can plausibly be identified as blue-shifted Fe\,{\sc xxvi} and arise from a disk wind. However, this would require that winds in neutron star LMXBs are observable at low inclination angles of $i$$\simeq$30$^{\circ}$, and are able to reach an outflow velocity as high as $v_{\mathrm{out}}$$\simeq$$25\,800~\kms$. \subsection{On the size of the binary orbit}\label{subsec:ucxb} \citet{baglio2016} reported that an optical spectrum obtained in late April, after \source\ had transitioned to the soft X-ray spectral state, was featureless apart from a possible He{\sc ii} emission line at 4686~\AA. Based on the lack of H-emission lines, which are usually seen in the optical spectra of LMXBs, the authors proposed that the donor star must be H poor and thus that \source\ is a candidate ultra-compact X-ray binary (UCXB). UCXBs are X-ray binaries in which the donor is an evolved star in a tight $P_{\mathrm{orb}}$$\lesssim$90~min orbit. \citet{baglio2016} estimated the time-averaged mass-accretion rate of \source\ based on \rxte/ASM and \maxi\ monitoring, and concluded that this is broadly consistent with the source having a He-rich donor in a $\simeq$40~min orbit, based on a comparison with the disk instability model (\citet{lasota2008}; see also the evolutionary tracks of \citet{vanhaaften2012}). We note that the lack of H-emission lines during single-epoch spectroscopy does not necessarily has to imply an UCXB nature. There are a few (black hole) LMXBs that did not show H-emission lines in their optical outburst spectra at some epochs, but have measured orbital periods of $P_{\mathrm{orb}}$$\gtrsim$90~min: e.g., Swift J1357.2--0933 \citep[][]{torres2011}, and Swift J1753.5--0127, \citep[][]{jonker2008_faint}. Nevertheless, their orbits are relatively small: $P_{\mathrm{orb}}$$=$2.8 and 2.85 or 3.2~hr, respectively \citep[][]{zurita2008,corralsantana2013,neustroev2014}. This may suggests a possible link between a short orbital period and the absence of H-emission lines. Indeed, there are several additional arguments that would support a relatively short orbital period for \source. Firstly, the full-width half maximum (FWHM) of the tentative He{\sc ii} 4686-\AA\ line identified by \citet{baglio2016} is quite low. From their figure~1, we estimate a FWHM of $\simeq$15~\AA\ and hence an equivalent resolved velocity of $\simeq$960$~\kms$. If a similar FWHM-$K_2$ correlation as recently found by \citet{casares2015} would apply for He and during outburst, this would suggest a low velocity amplitude of the companion star of $K_2$$=$$0.233\times$FWHM$\simeq$224$~\kms$. For a short orbital period of 1--3 hr and a very small mass ratio (applicable for a short period system), the mass function then is $f$$\simeq$0.05--0.14~$\Msun$. For a 1.5~$\Msun$ neutron star, this implies a low inclination of $i$$\simeq$19$^{\circ}$--27$^{\circ}$, which is remarkably similar to that obtained from modeling the Fe-K line in the X-ray spectra \citep[$i$$\simeq$18$^{\circ}$--35$^{\circ}$; this work and][]{ludlam2016}. Such a low inclination could account for the lack of orbital signatures in the optical data \citep[][]{baglio2016}, as ellipsoidal modulations would be small. Secondly, the orbital period of LMXBs correlates with their absolute visual magnitude and X-ray luminosity (\cite{vanparadijs94}, see also \cite{russell2006}). We measured an X-ray luminosity of $L_{0.7-35}$$\simeq$$2\times10^{37}~\lum$ on 2015 April 14, whereas \citet{baglio2016} reported an apparent visual magnitude of $\simeq$17~mag on 2015 April 24 (see their table 2). For $D$$=$5.8~kpc, we can then roughly estimate an orbital period of $P_{\mathrm{orb}}$$=$2.7~hr using the empirical relation of \citet{vanparadijs94}. The X-ray spectral shape may also provide clues about the size of the binary. \citet{sidoli2001} analyzed the broad-band (0.1--100 keV) spectra of a sample of Galactic globular cluster LMXBs, fitting these with a disk black body and a Comptonizing component. These authors found that only in the (candidate) UCXBs, the obtained inner disk temperatures were similar to that of the seed photons of the Comptonized emission, and the obtained radii (several km) broadly consistent with that expected for the location of the inner disk. For the non-UCXBs, the inner disk temperatures were much higher ($\simeq$2--3~keV), and their radii un-physically small ($\lesssim$1~km). This seems to suggest that only in UCXBs, the region providing the seed photons of the Comptonized emission is small and consistent with the inner accretion disk \citep[][]{sidoli2001}. The tentative classification scheme of \citet{sidoli2001} gained credence through testing for an UCXB nature by alternative methods such as the relation of \citet{vanparadijs94} discussed above, and the composition of the accreted matter inferred from thermonuclear X-ray bursts \citep[see][]{verbunt2006}. Its validity has also been strengthened by a number of additional sources \citep[e.g.,][]{gierlinski2005,falanga2005,fiocchi2008}. However, \citet{engel2012} identified a 2.15-hr orbital period for XB~1832--330 in the Galactic globular cluster NGC 6652, which was originally put forward as an UCXB based on the spectral-shape classification \citep[][]{parmar2001}. It was therefore proposed by \citet{stacey2012}, that the \citet{sidoli2001} scheme may not strictly apply to UCXBs (i.e., with $P_{\mathrm{orb}}$$\lesssim$90~min), but can still distinguish between short and long orbital period systems. Against this background, we note that our fits of the 0.7--35 keV \swift/\nustar\ spectrum of \source\ yield an inner disk temperature and seed photon temperature of the Comptonized component that are consistent within their 1$\sigma$ errors; $kT_{\mathrm{in}}$$=$$1.15\pm0.02$~keV and $kT_{\mathrm{s}}$$=$$1.27\pm0.45$~keV, respectively. Furthermore, from the continuum spectral shape we obtain an inner disk radius of $\simeq$9.5--11~km if we boldly apply a color-correction factor of 1.7 (see Section~\ref{subsec:refl} for caveats), i.e., consistent with inner disk radius measured from the reflection component. The \citet{sidoli2001} classification scheme would thus favor a small disk size, hence small orbital period, for \source. Finally, the maximum luminosity that can be reached during outburst should scale with the orbital period of the binary \citep[e.g.,][]{lasota01}, which is borne out by observations \citep[e.g.,][]{wu2010}. For the estimated bolometric peak X-ray luminosity of \source\ ($L_{\mathrm{bol}}$$\simeq$$5\times10^{37}~\lum$), the empirical relation of \citet{wu2010} yields an orbital period of $\simeq$1.4~hr. Disk instability theory predicts an orbital period of $\simeq$2.4~hr for an irradiated He disk \citep[][]{lasota2008_aa}. A He-dominated disk would be consistent the lack of H-features and presence of He{\sc ii} emission in the optical spectra \citep[][]{baglio2016}. We note that if \source\ harbors an evolved companion, the detection of Fe-K reflection in both the soft and the hard X-ray spectral state might argue in favor of a He donor rather than a C/O or O/Ne/Mg white dwarf donor, since a high C/O abundance could screen Fe and hence suppress the Fe-K line in C/O rich systems (\cite{koliopanos2013}, \cite{koliopanos2014}, but see \cite{madej2014}). In conclusion, the lack of H and the narrowness of the tentative He line in the optical spectrum, the X-ray and optical luminosity, the broad-band X-ray spectral shape, and the peak X-ray luminosity are all consistent with \source\ being a neutron star LMXB with a relatively small orbit of $P_{\mathrm{orb}}$$\lesssim$3~hr.

1607.01780

1607

1607.04670_arXiv.txt

NGG 4660, in the Virgo cluster, is a well-studied elliptical galaxy which has a strong disk component (D/T about 0.2-0.3). The central regions including the disk component have stellar populations with ages about 12-13 Gyr from SAURON studies. However we report the discovery of a long narrow tidal filament associated with the galaxy in deep co-added Schmidt plate images and deep CCD frames, implying that the galaxy has undergone a tidal interaction and merger within the last few Gyr. The relative narrowness of the filament implies a wet merger with at least one spiral galaxy involved, but the current state of the system has little evidence for this. However a 2-component photometric fit using GALFIT shows much bluer B-V colours for the disk component than for the elliptical component, which may represent a residual trace of enhanced star formation in the disk caused by the interaction 1-2 Gyr ago. There are brighter concentrations within the filament which resemble Tidal Dwarf Galaxies, although they are at least 40 times fainter. These may represent faint, evolved versions of these galaxies. A previously detected stripped satellite galaxy south of the nucleus is seen in our residual image and may imply that the filament is a tidal stream produced by perigalactic passages of this satellite.

NGC 4660 is an E5 elliptical galaxy with a strong incrusted disk component, and has frequently been used as a `classic' example of this type of galaxy. As such it has appeared frequently in the literature. It is located in the Virgo cluster, and has been identified as a member of the nearest compact group satisfying the Hickson criterion \citep{mam89}, along with other Virgo galaxies like M59 and M60, but \citet{mam08} concluded from surface brightness fluctuation measurements that M59 and NGC 4660 form a pair $\sim 1 - 2$ Mpc closer to us than the other three galaxies of the potential compact group. Hence we adopt a distance of 18 Mpc for this galaxy, giving a scale of 87 pc/arcsec and 5.2 kpc/arcmin. Generally it has been assumed that the disk in NGC 4660 is primordial and that the galaxy has not undergone a significant interaction or merger, but here we report the discovery of a long, curved filament apparently emanating from NGC 4660, in the direction of M59, which implies that it has experienced some kind of major event with a dynamical timescale of a few $\times 10^{8}$ years. The filament was discovered on a co-added array of scanned Kodak Technical Pan films of the SE region of the Virgo cluster taken with the UK Schmidt Telescope \citep{kat98, kat01} which have previously been used to identify or confirm other filamentary structures, e.g. the filamentary `ellipse' of IC 3481 \citep{per09}, the filamentary connections between NGC 4410 A/B, C and D \citep{per08} etc. Although such Schmidt plate/film studies have become almost obselete in modern astronomy, this discovery of the present, previously unsuspected, filament of NGC 4660 shows that they could still be put to productive uses. Here we report the discovery of the filament using the Schmidt data, and describe subsequent multiband CCD observations with the 2.1m telescope at the San Pedro M\'artir Observatory, which provide colour information for this galaxy and further imaging of the filament. In Section 2 we provide a summary of the previous work on the galaxy NGC 4660, while in Section 3 we describe the observational data and their processing. Results of the photometry of NGC 4660 and of the filament are given in Section 4, and in Section 5 we discuss the age and the formation of the filament and give general conclusions. \section {Previous work on NGC 4660} NGC 4660 (VCC 2000) has been studied in various ways, and has formed part of many samples of early-type galaxies over the years, on account of its proximity to us in the Virgo cluster and its incrusted disk component. It is located at RA = 12:44:32.0 Dec. = +11:11:26 (J2000), has a recessional velocity of 1083 km/s, and redshift of 0.003612 (NASA/IPAC Extragalactic Database). Although its classification in the SIMBAD database is E5, the Virgo Cluster Catalogue \citep{bin85} gives a classification of E3/S0--1(3), which emphasises its transitionary nature. There have been a few previous photometric and structural studies of NGC 4660, though three of them have used the same observational dataset. We compare our photometric results with theirs in Section 4. \citet{benet88} carried out CCD photometry and isophote analysis on NGC 4660 in the V, R and I bands. They found an effective radius $r_{e}$ between 9--9.8 arcsec and an isophote twist of $5^{\circ}$. Peak ellipticity was $\sim 0.5$ while $a_{4}/a$ (used as the 4th cosine coefficient in this case) rises to 0.04 at 30 arcsec. \citet{rix90} performed photometric modelling of galaxies with 2 components (spheroidal $r^{1/4}$ law bulge plus exponential disk). They applied their techniques to \citet{benet88}'s data on NGC 4660 and found they achieved better fits with a significant disk component included ($L_{disk}$ $\sim 50$\% $L_{bulge}$). \citet{ben88} described NGC 4660 as a rapidly-rotating elliptical while \citet{rix90} conclude that it is the disk rotation which is the major contributor to the detected velocity. \citet{sco95}, again using the same photometric data of \citet{benet88} for NGC 4660, find disk signatures in the ellipticity and $a_{4}/a$ profiles (again used as the 4th cosine coefficient). The ellipticity is $\sim 0.2$ at the centre and peaks at 0.5 at 10--30 arcsec, while $a_{4}/a$ is $< 0.01$ in the centre, rising to 0.02--0.04 between 10--40 arcsec. The velocity dispersion $\sigma$ is $\sim 200$ km s$^{-1}$ at $< 5$ arcsec (0.43 kpc) and 100 km s$^{-1}$ at 30 arcsec (2.6 kpc), while the disk rotation velocity is 180 km s$^{-1}$ at 10--25 arcsec (0.87--2.17 kpc) and $\sim 100$ km s$^{-1}$ at 0--10 arcsec ( 0--0.87 kpc). They find evidence of an isophotal twist, with PAs in the range $90^{\circ}$ to $100^{\circ}$, between 10--20 arcsec (0.87--1.74 kpc), taking this as evidence of a disk with inclined bar extending to 12 arcsec (1.04 kpc) from the centre. Its surface brightness profile is similar to early-type barred galaxies \citep{com93}. For the disk, $v/{\sigma}$ is very high, $\sim 3.3$, and D/B within the bar is 0.5. \citet{sco95} suggest that the disk may be unstable to bar formation, and the bar may have a long evolution time. The dataset of \citet{sco95} comprises mainly disky elliptical galaxies i.e. E galaxies with `pointed' isophotes (of which, NGC 4660 has the highest D/B ratio, 0.28). These galaxies are found to lie on the same correlation between central surface brightness and disk scale length defined by lenticular and spiral galaxies. This implies that disky E's are not inclined S0's but form `transition' objects between E's and S0's. The disk profile is not generally exponential. In NGC 4660 the disk and bulge rotate in the same direction and have similar surface brightness profiles, suggesting that they formed at about the same time out of the same material. This galaxy was observed by \citet{fer06} with HST/ACS as part of a sample of 100 early-type galaxies in Virgo. These data deal with the central regions of the galaxy at higher spatial resolution, and so the present data are complementary to this dataset. Their image shows the presence of a possible stripped satellite galaxy, 2.5 arcsec across and 4.5 arcsec south of the centre, described as blue, `resembling a nucleus with two spiral arms'. They fit the luminosity profile of NGC 4660 with one component with S\'ersic indices of 4.0 in the $g$ band and $4.5$ in the $z$ band and an average $g-z$ colour of 1.51. Structural parameters such as ellipticity, PA, $a_{3}, a_{4}, b_{3}, b_{4}$ are reported, giving results quite similar to previous references. $b_{4}$ is used here as the 4th cosine coeffcient and provides a clear separation of the boxy bulge and the disk. The $g-z$ colour profile shows a blueward tendency from 0-3 -- 20 arcsec then rises slightly to 50 arcsec. NGC 4660 forms part of the SAURON sample, a survey of 72 early-type galaxies using the integral-field spectrograph SAURON at the William Herschel Telescope. Stellar population studies of this sample give us ages and metallicities of the dominant stellar populations in central regions of this galaxy \citep{bac01, zee02}. \citet{kun10} carried out a stellar population analysis of absorption line strength maps of 48 of these galaxies including NGC 4660. They find only old stellar populations in this galaxy, $12.2 ^{+ 1.2} _{-0.6}$ Gyr at $R_{e}/8$ (1.4 arcsec, 1.3 kpc) and $13.4 ^{+1.3} _{-1.2}$ Gyr at $R_{e}$ (11.5 arcsec, 1.0 kpc). The metallicity [Fe/H] varies from $0.15 \pm 0.02$ at $R_{e}/8$ to $0.11 \pm 0.02$ at $R_{e}$, while the $\alpha$-element composition varies slightly, [$\alpha$/Fe] is $0.24 \pm 0.04$ at $R_{e}/8$ and $0.29 \pm 0.05$ at $R_{e}$. Although the [$\alpha$/Fe] variation is within the errors, it is typical of the other galaxies in the sample to have depressed [$\alpha$/Fe] values in the centre corresponding to higher metallicities. The velocity dispersion $\sigma$ is 221 km s$^{-1}$ at $R_{e}/8$ and 181 km s$^{-1}$ at $R_{e}$. All galaxies which are fast rotators in the sample and which have flattened components with disk-like kinematics are found to have different stellar populations in these flattened components. Those with young populations ($ < 3$ Gyr) frequently have circumnuclear disks and rings which are still forming stars. Meanwhile NGC 4660 belongs to the other extreme, in which the structure with disk-like kinematics has slightly older ages and lower [${\alpha}$/Fe] ratios in which the `secondary' star-formation event which formed the disk-like structure is presumably nearly as old as the `elliptical' component of the galaxy. We compare these results with populations indicated by our photometric colours in Section 5. Another interesting line of work carried out concerning NGC 4660 is its possible membership of a Compact Group (CG). \citet{mam89} considered it, alongside M59, M60, NGC 4638 and NGC 4647, as a possible CG involving members of the Virgo cluster, as it was found to satisfy Hickson (1982)'s criteria according to new magnitude measurements. However, \citet{mam08} concluded from surface brightness fluctuation measurements of individual early-type galaxies that NGC 4660 may form a pair with M59, at a distance of $\sim 2$ Mpc closer to us than most of the rest of the Virgo cluster. Although the various statistical calibrations differ, M59 and NGC 4660 as a pair are at least 440 kpc closer to us than M60 and NGC 4638. This line-of-sight depth would be too great for a CG \citep{mam08}. Our new data do not provide anything new on this interesting possibility. In the following section we consider the new observational data and their processing.

Tidal tails and other fine low surface brightness features around galaxies can only be formed in gravitational interactions between galaxies (usually major mergers with proportion of masses less than 4) over a timescale of a few $\times 10^8$ years \citep{tom72}. Therefore the detection of the filament near to NGC 4660 appears to indicate a past gravitational interaction for this galaxy, for which there was no previous evidence. Properties of tidal tails can be used to date the merger event which formed them. Tidal tails may be long-lived after a gravitational encounter. NGC 4660 may have had an encounter, or a `dry' interaction (not merger), with another early-type galaxy. The nearest candidate galaxy for such an interaction with NGC 4660 is LEDA 42878 (VCC 1991) which can be seen at about 6 arcmin { West-SouthWest of NGC 4660 (seen in Figs.\ 1 and 9), at 12:44:09.2 +11:10:32 (J2000), is classified as a dE in the SIMBAD database and identified as a nucleated dwarf galaxy by \citet{bin85}. This has a radial velocity of 1681 km s$^{-1}$ \citep{san11}, approximately 600 km s$^{-1}$ more than that of NGC 4660, which would make any significant interaction between them unlikely. There is also the dE 1 arcmin NE of the centre of NGC 4660, VCC 2002, but this appears to be too small to produce such a long filament in an interaction with NGC 4660. Such prominent tidal tails are more likely to be formed in `wet mergers' (in which at least one progenitor is a gas-rich spiral). There are no H I observations which could demonstrate whether the filament observed near NGC 4660 is gas-rich. The presence of only one tail may indicate that only one of the progenitors of NGC 4660 was a gas-rich spiral, The period of time for which a tail may be observed is important in studies of galactic history and evolution, especially in the case of NGC 4660 where it was the only direct evidence of the galaxy having experienced a tidal interaction. Tails are formed on the dynamical timescale of $\sim 10^{8}$ years and the existence of a bright tail does not imply that the interaction which formed it was recent, bright tails may survive for a few Gyr. The galaxy evolution numerical simulations of \citet{pei10} show early-type/spiral mergers producing peaks in Star Formation Rate (SFR) for $\sim 2$ Gyr after closest approach, while shells are seen at $\sim 1.5$ Gyr after closest approach. \citet{con09} predicts maximum merger timescales of $1.1 \pm 0.3$ Gyr, without taking into account the detailed history of the fallback of tidal features. \citet{hib95} modelled the spatial morphology and velocity structure of the famous merger product NGC 7252, suggesting that the merger occurred 0.6 Gyr ago and 80\% of the mass would fall back to the merger remnant in 2.5 Gyr. So detection of tidal debris would get difficult after about 3 Gyr, while if the merger were older there would be time for more minor mergers to occur, disrupting the filament \citep{duc11}. Most of the tail material would and does remain bound -- there are currently velocity reversals along the tails indicating material which has already reached the turnaround point in its orbit and has started falling to smaller radii. 20\% of current tail particles will not fall back to $< 5 R_{e}$ in a Hubble time. A filament as relatively narrow as the one in NGC 4660 would require at least one of the parent galaxies to be a spiral galaxy, though the other may have been an early-type dynamically hot galaxy, which would only have produced plume-like debris which may disappear more rapidly. Alternatively, if the other galaxy was indeed another spiral galaxy, the other tail has evaporated or was destroyed in a subsequent minor merger interaction with the general gravitational potential well of the cluster/group environment. So in Virgo one may expect tails to survive for maybe less than 2 Gyr. The presence of a spiral galaxy would make this a gas-rich (`wet') merger. The two peaks in surface brightness in the North of the CCD field (Figure 8) may correspond to Tidal Dwarf Galaxies (TDGs) which are gravitationally bound small systems of stars and gas formed in major mergers (mass ratio $> 1:4$), see \citet{mir92}. TDGs are generally thought to form from gas clouds pulled out of galaxies during mergers and so do not represent previously existing stellar systems. Here we only have data in red filters for these TDGs, which are often found to have a bluish colour \citep{duc11}. Studies of the colours and spectra of the candidate TDGs in NGC 4660 would prove interesting, but are made difficult by their faintness. However the candidate old tidal dwarfs studied by \citet{duc14} have typical sizes of 0.8--2.3 kpc and absolute magnitudes corresponding to --13.5 to --17.5 in $R$, while their central surface brightnesses are 23.5--26.5 $R$ mag arcsec$^{-2}$. Younger dwarfs in the same article have scale lengths of 2--6 kpc, absolute $R$ magnitudes of --14 to --17.5, and central surface brightnesses ranging from 20.5 to 25 $R$ mag arcsec$^{-2}$. While our candidate TDGs in NGC 4660 are at the low end of the surface brightness range for old dwarfs they are much smaller objects, with luminosities at least 40 times less than previously identified TDGs, and only a few pixels in size, so they inevitably have small scale lengths given the detection threshold. The implied masses will also roughly be 2 orders of magnitude less, so while the objects in \citet{duc14} are of the order of $10^8 M_{\odot}$, the present objects may be only about $10^6 M_{\odot}$ in mass. \citet{duc11} consider that the TDGs in the eastern tidal tail of NGC 5557 may be at least 2 Gyr old and they are still 1.3 mag bluer than the rest of this tail, while \citet{duc14} give a spectroscopic age of 4 Gyr for one of these TDGs. If the filament in NGC 4660 is several Gyr old then these could represent faint, evolved, even older TDGs, however they are so much smaller that they seem to represent extreme lower-luminosity cases of the class of TDGs or a different class of objects. Evidence for long lived TDGs, or detection of many fainter but similar objects, would be significant in terms of studies of numbers of dwarf satellite galaxies. The `stripped satellite' detected by \citet{fer06} may provide an alternative origin for the filament, formed by stars tidally stripped from this satellite galaxy during the closest parts of its orbit to NGC 4660, making the filament a tidal stream, comparable to the tidal stream associated with the Sagittarius dwarf galaxy satellite of the Milky Way \citep{iba94}. This would account for the presence of only one detectable filament. \citet{sta15} note that detection of extragalactic tidal streams from the ground is complicated by their low surface brightnesses, below 28 mag arcsec$^{-2}$ and this seems just about compatible with the surface brightness of the filament of NGC 4660 away from the possible TDGs. The colour maps of Figure 4a and 4b do imply some complicated structure in the inner disk area, with gradients from east to west, on the South side of the disk. Studies of tidal streams are now commonly made with the 3.6$\mu$m and 4.5$\mu$m filters of the {\it Spitzer Space Telescope} on its warm mission \citep{sta15}. It would be interesting to map the area around NGC 4660 in these filters. While the present filament is relatively bright, this tidal stream hypothesis does provide a possible explanation for the origin of the filament. The ages of stellar populations in the galaxy given by the SAURON data \citep{kun10} are 12--14 Gyr, with metallicities of 1.3--1.4 times solar (with marginal evidence for lower central metallicities) though this only covers the galaxy interior to and at $R_e$ (11.5 arcseconds). The $(B-V)$ colour profile varies from $\sim 0.9$ in the disk region to $\sim 1.15$ in the exterior regions, which is a bigger colour change than implied by the marginal metallicity gradient in the SAURON data (according to the Tables in \citet{bre94} the change in metallicity will only raise the $B-V$ colour by a few hundredths of a magnitude). However, the GALFIT results do show a rather bluer colour for the disk ($B-V \approx 0.7$ compared with $\approx 1.0$ for the much larger and brighter S\'ersic component), and the disk is progressively less prominent in the 2-component fit from B to I. The models of \citet{bre94} give $B-V$ colours of 1.03--1.13 for a 12.5 Gyr population with metallicity 1.0--2.5 times solar, slightly redder than the S\'ersic component here but compatible with the colours in the inner colour profile. Meanwhile the disk component colours can be reproduced with models of age 1-2 Gyr for solar metallicity and 1 Gyr for 2.5 times solar metallicity. So this may imply relatively young stellar populations in the disk, masked by the fact that the elliptical `bulge' component dominates at all radii, and because the SAURON data use circular apertures, and now revealed by the ability of GALFIT to separate the light from the two components. This may therefore be evidence for star formation produced by a merger 1-2 Gyr ago and now detectable through a bluer colour in the stellar population of the disk component separated by GALFIT. In conclusion, we have detected a tidal filament in the vicinity of NGC 4660 using the enhanced Schmidt film material, and detected brightness peaks of the filament in subsequent deep CCD data. This is not a galaxy expected to have undergone a tidal interaction/merger recently, and the SAURON data indicate only old populations, and a disk coeval with, or only slightly younger than, the elliptical component of the galaxy. However the two-component fit shows that the disk component has bluer colours, indicating that the disk may have been formed in a merger 1-2 Gyr ago, or that this interaction caused enhanced star formation in a pre-existing disk. There are brighter concentrations within the filament which resemble Tidal Dwarf Galaxies, although they are at least 40 times fainter. These may represent faint, evolved TDGs or may be another class of object. A previously detected stripped satellite south of the nucleus it detected in our residual image and may imply that the filament is really a tidal stream produced by perigalactic passages of this satellite. It would be interesting to carry out deep H I observations in the region of the filament, and its brightness peaks, to determine its gas content, and to obtain {\it Spitzer} images of the filament, mapping the old stellar popualtions to fainter levels.

1607.04670

1607

1607.01055_arXiv.txt

The Galactic blue supergiant SBW1 with its circumstellar ring nebula represents the best known analog of the progenitor of SN~1987A. High-resolution imaging has shown H$\alpha$ and IR structures arising in an ionized flow that partly fills the ring's interior. To constrain the influence of the stellar wind on this structure, we obtained an ultraviolet (UV) spectrum of the central star of SBW1 with the {\it Hubble Space Telescope} ({\it HST}) Cosmic Origins Spectrograph (COS). The UV spectrum shows none of the typical wind signatures, indicating a very low mass-loss rate. Radiative transfer models suggest an extremely low rate below 10$^{-10}$ $M_{\odot}$ yr$^{-1}$, although we find that cooling timescales probably become comparable to or longer than the flow time below 10$^{-8}$ $M_{\odot}$ yr$^{-1}$. We therefore adopt this latter value as a conservative upper limit. For the central star, the model yields $T_{\rm eff}$=21,000$\pm$1000~K, $L\simeq$5$\times$10$^4$ $L_{\odot}$, and roughly Solar composition except for enhanced N abundance. SBW1's very low mass-loss rate may hinder the wind's ability to shape the surrounding nebula. The very low mass-loss rate also impairs the wind's ability to shed angular momentum; the spin-down timescale for magnetic breaking is more than 500 times longer than the age of the ring. This, combined with the star's slow rotation rate, constrain merger scenarios to form ring nebulae. The mass-loss rate is at least 10 times lower than expected from mass-loss recipes, without any account of clumping. The physical explanation for why SBW1's wind is so weak presents an interesting mystery.

SN~1987A was the nearest supernova (SN) in modern times. Two surprising observations associated with SN~1987A (see review by \citealt{arnett+89}) were the identification of a blue supergiant (BSG) progenitor in pre-explosion images \citep{walborn89,rousseau78,arnett87,arnett+89} and its very unusual triple-ring circumstellar nebula \citep{burrows95,crotts95}. These two are intimately related, since the geometry of the nebula bears the imprint of mass loss shaped by binary interaction and/or rapid rotation as the star evolved to its blue pre-SN state. The dynamical age of the nebula is only $\sim$20,000 yr \citep{meaburn95,ch00}, so the nebular structures trace recent pre-SN mass loss on a time scale much shorter than core He burning and longer than C burning. The total mass of the ring is uncertain (due to an uncertain neutral fraction), but may be 0.1-1 $M_{\odot}$ \citep{fransson15}. An important unresolved question is whether or not close binary evolution was key in determining the progenitor's BSG state (mass transfer or merger, mass loss, etc.). The ring nebula may therefore provide important clues to how and why the progenitor came to be a BSG. Understanding the origin of the observed triple-ring structure has been difficult, however. Several early models showed that a faster BSG wind expanding into a previous slower red supergiant (RSG) wind with an equatorial density enhancement could yield an equatorial ring and bipolar structure \citep{lm91,bl93,ma95,collins99}. However, the pair of polar rings around SN~1987A really are empty rings, rather than limb-brightened polar lobes or filled caps, and their origin has not been satisfactorally explained. It is difficult to understand the origin of the equatorial density enhancement in a RSG wind without invoking a binary \citep{collins99}. In subsequent studies, two different types of models have been proposed as plausible ways to form the nebula. A scenario involving a binary merger as a RSG and subsequent blueward evolution was proposed \citep{mp07,mp09}, and was suggested to account for the observed nebular structure. However, this specific merger model predicts filled polar caps and relatively empty mid-latitudes in the nebula, whereas the observed nebula has no polar caps and may have some emission in the side walls of an hourglass structure. Moreover, the model requires that the merger product should be rotating very rapidly, which seems to be at odds with Galactic analogs (see below). A somewhat different model involves rotating single-star evolution, where a massive star spins up as it contracts on a post-RSG blue loop, nearly reaching critical rotation and ejecting a ring \citep{chita08}. Then a bipolar wind from the rotating BSG expands into the RSG wind and ring, forming transient structures that may resemble the rings of SN~1987A \citep{chita08}. However, this model also requires a rapidly rotating BSG, inconsistent with observations of Galactic analogs. In either case the strength of the BSG wind is a critical ingredient. Radio observations of SN~1987A (or rather, the lack of radio emission at early times) suggested that for the first 1000-1500 days after explosion, the blast wave was expanding relatively unimpeded through a very low-density wind \citep{ss93}. After about 1500 days, however, the radio emission brightened and the expansion speed slowed to only 3500 km s$^{-1}$ \citep{gaensler97,gaensler00,zanardo13}. In order to reach the angular scale of the resolved radio emission when it turned on, the blast wave must have been expanding at about 35,000 km s$^{-1}$ for those first 1500 days \citep{ss93}. \citet{cd95} suggested that after the initial free expansion through a rarefied wind with $\dot{M}$=7.5$\times$10$^{-8} M_{\odot}$ yr$^{-1}$ or less, the blast wave slammed into a dense H~{\sc ii} region that partly filled the ring's interior. In this model, the dense H~{\sc ii} region was created by photoionization of the RSG wind by the BSG's UV radiation. This collision slowed the forward shock's expansion and caused the radio emission to brighten dramatically \citep{cd95}. Since the progenitor star Sk$-$69$^{\circ}$202 is now dead, it is hard to improve our understanding of the connection between the star and its nebula. For this reason, nearby analogs of SN~1987A's progenitor -- where the BSG has not yet exploded -- become interesting and valuable. There are currently three well-studied analogs known in our Galaxy: (1) Sher~25 \citep{brandner97,smartt02}, which is significantly more luminous and has partial polar caps instead of rings, (2) HD~168625 \citep{smith07}, which is a more luminous LBV candidate that does have polar rings, and (3) SBW1 \citep{sbw} (discussed below).\footnote{Another possible member of the group is MN18, which is a similar BSG with a ring-like bipolar nebula \citep{gvaramadze15}, although this object has not yet been studied in as much detail as the others.} An interesting recent result places important constraints on the formation of these rings: \citet{taylor14} monitored the central stars of all three Galactic analogs with high-resolution spectroscopy and did not find any radial velocity variations consistent with close massive binaries (the value of sin $i$ is presumed from the resolved equatorial rings; if a binary system's orbit were significantly misaligned with these rings, then binary interation would not help to explain them). Perhaps even more interesing, \citet{taylor14} found that all three BSG central stars have relatively {\it slow rotation speeds}. For SBW1, the rotation speed is only about 40 km s$^{-1}$. Such slow rotation may be quite problematic for some merger models that would predict a rapidly rotating BSG post-merger product -- especially if the stellar wind in the BSG phase is very weak. We will return to this issue later in the paper. Of these three Galactic analogs, SBW1 is currently the best known analog to SN~1987A in terms of stellar properties and nebular structure. The ring around SBW1 was discovered by \citet{sbw} during a survey of the Carina Nebula, but two factors suggest that it is actually located several kpc {\it behind} the Carina Nebula and is seen there in projection. First, its positive radial velocity compared to expectations for Galactic rotation in that direction suggest that it is on the far side of the Sagittarius-Carina arm and outside the Solar circle. Second, its apparent magnitude and color only match its spectral type and luminosity class (B1.5 Iab; \citealt{sbw}) if it is at a much larger distance than the Carina Nebula. At a distance of 6-7 kpc, the stellar luminosity (0.5--1$\times$10$^5$ $L_{\odot}$) as well as the size of the ring nebula ($r \simeq 0.2$ pc) make SBW1 a close match for SN~1987A. Detailed analysis of the SBW1 nebula has provided interesting clues that may alter our ideas about the formation of the nebula around SN~1987A. {\it HST} images of SBW1 \citep{smith13} show a pattern of clumps around the ring and a radial extent that closely resemble the spacing and size scale of spots in the ring of SN~1987A. The interior of the ring is filled with diffuse H$\alpha$ emission, although the ring would probably appear much brighter relative to the interior if it were flash ionized by a SN \citep{smith13}. A very interesting result is that high-resolution ground-based infrared (IR) images show that the interior of the ring is also partly filled with diffuse emission from warm dust. Since BSGs don't produce dust in their winds, this requires that the dust inside the ring was entrained from the ring itself. \citet{smith13} proposed that this structure arises because the inner surface of the dense and neutral equatorial ring is ionized by the central star, and that this triggers a dusty ionized photoevaporative flow that fills the interior of the ring. The ionized gas expands into the ring until it collides with the stellar wind; entrained dust piles up at this interface, producing the observed peaks of thermal-IR emission inside the ring \citep{smith13}. This simple ionized flow is able to dramatically influence the observed structure and dynamics of the nebula because the ring's expansion is slow (10-20 km s$^{-1}$), comparable to the sound speed of the ionized gas. This directly imaged structure around SBW1 appears to validate the picture of an H~{\sc ii} region inside the ring of SN~1987A proposed by \citet{cd95}, which was deduced from the time evolution of the SN's radio emission. In this scenario, the main requirement is that the star ejected a thin, dense ring about 10$^4$ yr ago. While this might have occurred in post-RSG evolution to the blue (perhaps with a merger), a ring might also be ejected in a brief mass-transfer episode of a close binary or in an eruptive mass loss event from a rotating star (e.g., \citealt{st07}). A previous RSG phase is not necessarily required \citep{sbw}. \begin{figure*} \includegraphics[width=\textwidth]{fig1.eps} \caption{\label{modeloptical} Comparison between the observed optical spectrum of SBW1 in the range $3760-4690$~\AA~(black line) and the best fit CMFGEN model (red line). The strongest spectral features are identified. The broad feature at $\lambda\simeq$ 4430 \AA \ in the observed spectrum is due to absorption by a known diffuse interstellar band. There is a slight error in the wavelength solution at the blue end of the spectrum (note that Balmer lines are shifted slightly blueward compared to the model), due to poor signal to noise in the arc spectrum used for calibration. When we inferred \geff from this spectrum, we compared to individual lines with an appropriate shift for each.} \end{figure*} \begin{figure*} \includegraphics[width=6in]{fig2.eps} \caption{\label{modeluv} Similar to Fig.~\ref{modeloptical}, but in the ultraviolet region between $1370-1750$~\AA. The strongest spectral lines are identified, while the remaining features are mostly due to \ion{Fe}{iii}, \ion{Fe}{iv}, or \ion{Ni}{iii} lines. This figure also includes a plane-parallel (p.p.) CMFGEN model in the UV (top panel, in blue) to illustrate the atmospheric spectrum with no wind.} \end{figure*} \begin{figure*} \includegraphics[width=5.8in]{fig3.eps} \caption{\label{uvcomparehd} Comparison between the ultraviolet spectrum of SBW 1 (B1.5~Iab; black) and HD~13854 (B1~Iab; red). HD~13854 was scaled to roughly match the continuum flux of SBW1. Note the stronger \ion{Si}{iv} and \ion{C}{iv} resonance lines in the latter because of the higher mass-loss rate. } \end{figure*} In order to test whether this proposed scenario actually works for the specific case of SBW1, we need to know the mass-loss rate of the wind from the central star, because this was an assumed parameter in our previous analysis \citep{smith13}. For this reason, we proposed to obtain UV spectra of the central star. Optical spectra are useful for constraining atmosphere/wind models as well, but the UV resonance lines are usually the most sensitive probes of the wind density and speed. Our new observations are discussed in \S 2, our analysis of the data including a comparison with radiative transfer models is presented in \S 3, and a discussion of the results and implications is given in \S 4.

\begin{figure*} \includegraphics[width=5.5in]{fig4.eps} \caption{\label{halpha} . Detail of the H$\alpha$ line profile in an echelle spectrum of the star obtained with EMMI (from \citealt{sbw}). The emission at line center is due to nebular emission from the ring (this is clear in resolved long-slit spectra; see \citealt{sbw}), but the line wings trace the photospheric absorption profile due to pressure broadening. The rotation rate is only $\sim$40 km s$^{-1}$ \citep{taylor14}, so rotation does not alter the broad line wings. To this high-resolution spectrum, we compare a CMFGEN model with the same parameters as discussed above, except that we explore different values of the effective gravity. From this comparison, it is evident that log $g_{\rm eff}$=2.5 (cgs units) is too low and 4.0 is too high, but log $g_{\rm eff}$ values of 2.7-3.0 provide a good match to the data, with 3.0 being somewhat preferred in some wavelength ranges.} \end{figure*} \begin{figure*} \includegraphics[width=5.5in]{fig5.eps} \caption{\label{hrd} HR Diagram with representative single-star evolution tracks from \citet{brott11}. We denote the location of the progenitor stars of SN~1987A \citep{maund+04} and the companion of SN~1993J \citep{fox+14}, as compared to SBW1. For SBW1, the two black filled dots show the luminosity indicated by our CMFGEN model for assumed distances of $D_7$=7 kpc and $\sqrt{2} \times D_7$=9.9 kpc, while the red and black text give the implied present-day stellar masses for log $g_{\rm eff}$=2.7 and 3, respectively. The implication is that an assumed distance of 7 kpc is too small, because $g_{\rm eff}$ gives stellar masses that are far below the mass one would infer from comparing the luminosity to evolutionary tracks. A slightly larger distance of 9.9 kpc, on the other hand, gives double the luminosity and stellar mass, and is in much better agreement with the luminosity for evolutionary models with that mass if we adopt log $g_{\rm eff}$=3. } \end{figure*} \subsection{Central star properties} Our CMFGEN analysis confirms many of the physical paramaters that had been inferred previously from photometry and spectral type. The value of $\teff$ = 21,000~K that we derive from the CMFGEN model is the same as assumed previously from the B1.5~Iab spectral type and spectral energy distribution \citep{sbw,smith13}. The luminosity derived previously from the SED was uncertain (0.5-1)$\times$10$^5$ $\lsun$. Our CMFGEN model gives a somewhat lower value of 2.5$\times$10$^4$ $\lsun$ $D_7^2$, where $D_7$ is the distance relative to our assumed value of 7 kpc. As noted earlier, however, the true luminosity would be raised to (5-6.5)$\times$10$^4$ $L_{\odot}$ if we adopted a larger value of $R_V$, as may be appropriate. Below, we also find that the likely distance is larger than 7 kpc. Interestingly, with the weak wind of SBW1, we can use the effective gravity $g_{\rm eff}$ from the model to place constraints on both the true luminosity and present-day mass if we assume that the spectroscopically derived mass $M_{\rm spec}$ is comparable to the evolutionary mass $M_{\rm evo}$. \citet{lk14} have discussed this in detail, and concluded that $M_{\rm spec}$ is usually a reliable representation of the true stellar mass as long as the star is not close to the Eddington limit (i.e. for moderately massive and intermediate-mass stars). Since SBW1 has no detectable signatures of a wind, we surmize that it is nowhere near its Eddington limit. Our model derived from a comparison to low-resolution spectra gave log $g_{\rm eff}$=2.6 (Table 1). However, by comparing CMFGEN models to higher resolution spectra of Balmer lines, as shown in Figure~\ref{halpha}, we favor a somewhat higher value of 2.7-3.0 for log $g_{\rm eff}$. This difference arizes because the lower-resolution spectra are compromised by nebular emission that affects the line profiles, whereas the stellar Balmer line wing shapes are resolved in the echelle spectrum. From the definition of $g_{\rm eff} = GM/R^2$ we can write \begin{displaymath} M_{\rm spec} = \frac{g_{\rm eff} \, R_*^2}{G} = \frac{g_{\rm eff} \, L_*}{4 \pi \sigma \, G \, T_{\rm eff}^4} \end{displaymath} \noindent where $M_{\rm spec}$ is the present-day spectroscopically derived stellar mass, $R_*$ and $L_*$ are the star's photospheric radius and bolometric luminosity, respectively, $\sigma$ is the Stefan-Boltzmann constant, and G is the gravitational constant. We note that the uncertainty is dominated by errors in $g_{\rm eff}$ rather than errors in $T_{\rm eff}$. Inserting fiducial values of $T_{\rm eff}$ = 21,000 K, $g_3 = (g_{\rm eff} / 10^{3} {\rm cm \, s^{-2}})$ and $L_* = 2.5 \times 10^4 L_{\odot}$ ($D_7$)$^2$, where $D_7$ is the adopted distance relative to 7 kpc, we then have \begin{displaymath} M_{\rm spec} \, \approx \, 7 \, M_{\odot} \, ( g_3 \, D_7^2) \end{displaymath} \noindent for the {\it present-day} stellar mass as indicated by the effective gravity. We can then attempt to constrain the actual luminosity and mass of SBW1 by seeing what combinations of $D$, $L_*$, $g_{\rm eff}$, and $M_{\rm spec}$ give values consistent with evolutionary models. Figure~\ref{hrd} shows a Hertzsprung-Russel (HR) diagram comparing the inferred luminosity of SBW1 to single-star evolutionary tracks \citep{brott11}, for reference. This comparison shows that the lower luminosity for an assumed distance of 7 kpc, combined with the $g_{\rm eff}$ indicated by the spectrum, gives a stellar mass of $\sim$7 $M_{\odot}$ --- but this is much lower than one expects in this region of the HR Diagram. For a slightly larger distance of 9.9 kpc, however, the luminosity and $M{\rm spec}$ rise by a factor of two, and importantly, are then in very good agreement with the expected luminosity for a 14-15 $M_{\odot}$ evolved star. In fact, the $T_{\rm eff}$ and $L_*$ we derive would agree very well with the hook in the evolutionary track for a 14 $M_{\odot}$ star that occurs after core H exhaustion, which would seem to make sense with the blue supergiant spectral type. Of course, the comparison to single-star evolutionary tracks may be misleading if SBW1 is the result of binary evolution that may alter its $L$/$M$ ratio; one could argue that it would be appropriate to compare SBW1's spectroscopic mass to models for a merger product or mass gainer. Such a comparison might favor a slightly different combination of $D$, $g_{\rm eff}$, $L_*$, and $M_*$, especially if these values evolve in the $\sim$10$^4$ yr after a merger or mass transfer episode. We could, of course, also achieve a luminosity of 5$\times$10$^4$ $L_{\odot}$ with a smaller distance (8 kpc, say) and a slightly higher $R_V$ value. The luminosity is unlikely to be much lower than 2.5$\times$10$^4$ $\lsun$, however, due to the fact that the optical spectrum has a supergiant luminosity class, and that the implied mass from $g_{\rm eff}$ would be too low for the corresponding $L$. Based on this comparison, we therefore favor values of log $g_{\rm eff}$=3.0 cm s$^{-2}$ and $L_*$ = 5$\times$10$^4$ $L_{\odot}$ for SBW1. Altogether, these parameters make the central star of SBW1 only a little hotter than Sk$-$69$^{\circ}$202, and about 50\% of its bolometric luminosity. It has an effective initial mass of around 14~$\msun$, as compared to 18~$\msun$ for Sk$-$69$^{\circ}$202 \citep{arnett89,arnett+89}. From $L \propto R^2 T_{\rm eff}^4$, the implied stellar radius is of order 15-20 $\rsun$, and so the surface escape velocity is about the same as for the Sun or slightly lower. The chemical abundances we derive from the photospheric spectrum show basically Solar composition, except for an enhanced N abundance that is elevated by a factor of 3 or 8 compared to Solar N/O or N/C ratios, respectively. This, too, is quite similar to the enhanced N abundances inferred from the emission-line spectrum of the ring around SN~1987A, and is indicative of significant CNO processing present at the star's surface. Will the central star of SBW1 be the next Galactic SN? The dynamical age of the nebula is about 10$^4$ yr, similar to SN~1987A, so perhaps it is a good candidate. Of course, the uncertainty of such a clock is huge. Aside from the nebular age and an analogy to SN~1987A, we have little from which to infer the time until the impending core collapse. \begin{figure*} \includegraphics[width=4.5in]{fig6.eps} \caption{\label{flow}Plot of the flow timescale (black solid) compared to the cooling timescale in a CMFGEN model as a function of radius for some representative assumed values of mass-loss rate and shock velocity $U_0$. If the wind outflow speed is 300-500 km s$^{-1}$, then the typical speed of internal shocks in the wind is probably 100--200 km$^{-1}$ or less, and very likely less than 300 km s$^{-1}$. Thus, we see that for expected shock speeds, the wind cooling timescale becomes comparable to the flow timescale in the inner wind when the mass-loss rate drops below 10$^{-8}$ $M_{\odot}$ yr$^{-1}$. As such, it is possible that UV diagnostics become less reliable at such low wind densities due to increased ionization.} \end{figure*} \subsection{Stellar wind properties} The most significant observational result from our COS spectrum is the lack of any strong wind features in the spectrum, which is very unusual for a blue supergiant. Consequently, the most interesting result from our quantitative CMFGEN analysis is the astonishingly low derived mass-loss rate of SBW1. Comparing our lowest mass-loss rate CMFGEN models (which still show some evidence of wind emission) to the observations, which show none, implies and upper limit of $\dot{M} < 10^{-10}$ \ $\msun$ yr$^{-1}$ for an assumed terminal wind speed of 300 $\kms$. However, this assumes that the UV resonance lines are modeled correctly at such low wind densities. A cautionary remark relates to the so-called ``weak-wind problem'' (see \citealt{smith14} for a review), where UV-diagnostics of late-type O dwarfs yield mass-loss rates that are 100 times lower than expected from the Vink et al.\ recipe and from H$\alpha$ diagnostics. The cautionary comment is that an independent method of deriving the mass-loss rate based on the structure of a bow shock around $\zeta$~Oph gives a mass-loss rate estimate that is $\sim$10 times higher than UV diagnostics, but still an order of magnitude lower than expected from standard mass-loss prescriptions \citep{gvaramadze12}. Thus, there are some indications that the weak-wind problem for late O dwarfs is perhaps not as severe as indicated by UV estimates. Thus, when mass-loss rates are low, CMFGEN and similar models might underestimate the mass-loss rate somewhat based on UV diagnostics. This may be caused by inefficient cooling at low wind densities, so that shocks within the wind keep the ionization level higher than expected \citep{bouret15,puebla16}. Does some version of this weak wind problem translate to BSG winds such as the case of SBW1? Figure~\ref{flow} shows how the flow timescale in the wind compares to the cooling timescale for some representative assumed values of the mass-loss rate and the collision speed $U_0$ of internal shocks in the wind. The temperature and ionization balance of the wind depends on heating by shocks within the flow, and cooling, which depends on the density. If the cooling timescale becomes long at low densities, the wind may expand before it can cool, and so the ionization in the inner wind may be higher than in a CMFGEN model. Typical wind speeds for an early B supergiant would be around 500 km s$^{-1}$, and we would expect shocks within the clumpy outflowing wind to be some fraction of that -- perhaps 100-200 km s$^{-1}$ and almost certainly less than 300 km s$^{-1}$ (a caveat is that we can't be certain about the value of the wind speed, since we don't actually detect any wind absorption; thus, it remains possible that SBW1's wind might be faster than typical winds for B1.5 supergiants, which might raise the allowed mass-loss rate). Therefore, in Figure~\ref{flow} we should expect SBW1's wind to reside somewhere above the green dashed line and below the blue dash-dotted line (for $U_0$=100 and 300 km s$^{-1}$, respectively). For an intermediate shock speed around 200 km s$^{-1}$, for example, a mass-loss rate below 10$^{-8}$ $M_{\odot}$ yr$^{-1}$ would make the cooling timescale and the flow timescale about the same in the inner wind (a few stellar radii). This means that below 10$^{-8}$ $\msun$ yr$^{-1}$, the bulk of the wind might remain hotter than in the CMFGEN model, and would be harder to observe in the typical UV diagnostic lines that we are referrring to. Thus, if the mass-loss rate drops much below 10$^{-8}$ $M_{\odot}$ yr$^{-1}$, we cannot be confident that CMFGEN is properly treating the relevant physics, whereas above this, we should begin to see some evidence of a wind in the UV lines. We adopt 10$^{-8}$ $M_{\odot}$ yr$^{-1}$ as a fairly conservative upper limit to the mass-loss rate of SBW1, rather than 10$^{-10}$ $M_{\odot}$ yr$^{-1}$, due to this uncertain treatment of the cooling and ionization balance in CMFGEN at such low mass-loss rates. Our hypothesis that the wind of SBW1 has too low a density to cool --- and therefore remains hot --- can be tested. This hypothesis would predict detectable X-ray, EUV, and possibly FUV emission signatures from the wind, which may be verified with future observations. An observational determination of $L_X/L_{\rm Bol}$ with future X-ray observations would thus help provide a direct constraint on the mass-loss rate of SBW1 and the amount of shock heating within the wind. At the stellar temperatures appropriate for SBW1's spectral type, CMFGEN does not predict any N~{\sc v} or O~{\sc vi} features, but CMFGEN's reatment of the wind is not appropriate if the wind remains hot as we suspect. If low density inhibits cooling, we can make a qualitative prediction that N~{\sc v} and O~{\sc vi} may be observed (see, e.g., \citealt{bouret15,puebla16,zsargo08}), but a more detailed model beyond CMFGEN's current capabilities would be needed to derive a specific line strength for a quantitative mass-loss rate. Even this revised upper limit to the mass-loss rate of 10$^{-8}$ is much lower than one would expect for this star. For example, from the mass-loss prescriptions given by \citet{vink01}, we would expect $\dot{M}$ = 1.2$\times$10$^{-7}$ $\msun$ yr$^{-1}$ for $L = 5 \times 10^4 \lsun$, for line-driven winds at the appropriate $\teff$ of SBW1 (note that 21,000 K places this on the cool side of the bistability jump). Our observationally derived upper limit to the mass-loss rate for SBW1 is more than 10 times lower than this expected value, even with no reduction in the observed value to correct for clumping. Why is the wind of SBW1 so weak as compared to expectations, and as compared to other observed BSGs? The solution to this puzzle may hold important clues related to the origin of the ring nebula and the star's evolutionary history, and perhaps also for extragalactic SNe that appear similar to SN~1987A. For SN~1987A, the physical properties of the pre-SN stellar wind were uncertain, but some considerations also pointed to an anomalously low mass-loss rate compared to other BSGs. On the one hand, models derived from interpreting the early radio observations in the context of free-free self absorption of the SN radio emission by the freely expanding wind yielded a relatively high mass-loss rate of order 3.5--6 $\times$ 10$^{-6}$ $M_{\odot}$ yr$^{-1}$ with $v_w$=550 km s$^{-1}$ \citep{cf87,lf91,cd95}. On the other hand, hydrodynamic interacting-winds models used to explain the formation of the nebula required much weaker BSG winds in order to reproduce the slow expansion speed of the equatorial ring \citep{bl93,ma95}. To keep the ring expanding at the slow observed value of $\sim$10 km s$^{-1}$, these models would require upper limits to the mass-loss rate and wind speed of $\dot{M} < 3 \times 10^{-7}$ $M_{\odot}$ yr$^{-1}$ and $v_w < 300$ km s$^{-1}$. \citet{bl93} suggested that this discrepancy might be explained if the star's mass-loss rate increased in the last decades or century leading up to the moment of explosion, but \citet{cd95} suggested that synchrotron self-absorption, rather than free-free self absorption by the wind, might explain the early radio observations. Further indication that the progenitor star's mass-loss rate was low compared to normal BSGs came from the rebrightening in the radio at $\sim$1500 days after the SN \citep{ss92,ss93,gaensler97,gaensler00,manchester02,zanardo13}. This rebrightening was attributed to the collision between the fast SN ejecta and an H~{\sc ii} region from the photoionized RSG wind, as noted in the Introduction \citep{cd95}. Assuming that the interior region was filled with a relatively low-density freely expanding BSG wind, \citet{cd95} showed that this collision could occur at the observationally inferred radius of $\sim$0.1 pc with a model that adopted $\dot{M} = 7.5 \times 10^{-8}$ $M_{\odot}$ yr$^{-1}$ and $v_w = 450$ km s$^{-1}$. Later models refined this value, in some cases including constraints from the evolution of X-ray emission, to even lower values of around 5$\times$10$^{-9}$ $M_{\odot}$ yr$^{-1}$ or less \citep{vikram07,dewey}. This is a very low mass-loss rate for a BSG star of $\sim$10$^5$ $L_{\odot}$. According to the standard recipie for hot star mass-loss rates usually used in evolutionary codes \citep{vink01}, a star with log($L/L_{\odot}$)=5, $T_{\rm eff}$=21,000~K, and $M$=18 $M_{\odot}$ should have a mass-loss rate of 4.8$\times$10$^{-7}$ $M_{\odot}$ yr$^{-1}$ at LMC metallicity. The mass-loss rate inferred for the progenitor of SN~1987A based on the expansion of the blast wave is at least 6 and as much as 100 times lower than this expected value. This appears very similar to the case of SBW1 outlined above. Thus, both the progenitor of SN~1987A and SBW1 seem to share the pecularity that they have BSG winds that are extremely weak compared to the expected wind strength for their stellar parameters. This is not the case for the other two well-studied Galactic analogs with ring nebulae; both Sher~25 and HD~168625 have strong H$\alpha$ wind emission and have mass-loss rates that are normal (Sher~25) or strong (HD~168625) compared to other BSGs \citep{smartt02,nota96}. Although accounting for clumping has been argued to require a reduction to mass-loss rate recipies by factors of 3-5 (see \citealt{smith14} for a review), the deficits for SN~1987A and SBW1 are greater than this (and again, we did not include a clumping correction for SBW1). In models that aim to explain the formation of SN~1987A's triple ring nebula with a merger \citep{mp07,mp09}, BSG mass-loss rates of (1-2) $\times$ 10$^{-7}$ $M_{\odot}$ yr$^{-1}$ are adopted to shape the ring nebula (e.g., Table 4 in \citealt{mp09}). Recently, \citet{orlando15} adopted this same value for the mass-loss rate in their simulations of the SN interaction with the CSM, although they did not explore the impact of other assumed values for the mass-loss rate. These are higher than the observationally inferred values for SN~1987A (from the time history of radio emission, as noted above) and for SBW1. It is therefore unclear if interacting stellar winds can provide a viable physical explanation for the shaping of the nebulae around SN~1987A and SBW1. The issue of pressure balance is discussed more below. \subsection{Implications for the nebula and the pre-SN evolution of SN~1987A} Previous studies have discussed the formation of bipolar and ring nebulae, like the ones around SN~1987A and SBW1, in the context of interacting winds where a fast BSG wind expands into a slower and asymmetric RSG wind with an equatorial density enhancement (see the Introduction). However, a somewhat different scenario was discussed wherein a fast BSG interacts with an H~{\sc ii} region or photoevaporative flow for the specific cases of SN~1987A \citep{cd95} and SBW1 \citep{smith13} based on the inferred density structure inside the ring, which is inconsistent with a simple interacting winds scenario. In this section, we discuss how the extreme weakness of the BSG wind from SBW1 requires further modification to the story. In this scenario, the location of the shock between the BSG wind and the ionized photoevaporative flow is determined by the mass-loss rate and speed of the wind from the central star, balanced by the pressure of the photoevaporative flow. Specifically, ram pressure of the stellar wind $\rho \, v^2$ is balanced by the thermal pressure of the ionized gas inside the ring nebula. The photoevaporation rate of the ring that is the source of gas and dust in the H~{\sc ii} region depends on geometry and the ionizing photon flux of the star, $Q_H$, which is given in Table 1. However, in this case we can avoid the uncertainty introduced by the detailed geometry of the ring (clump size, ring height, whether gas in the walls of an hourglass contributes, etc.) because spectral observations of the nebula (the H$\alpha$ emission measure and the [S~{\sc ii}] $\lambda\lambda$6717,6731 line instensity ratio in the spatially resolved diffuse interior of the ring) directly constrain the density of the ionized flow filling the inside the ring to be roughly 300-500 cm$^{-3}$ \citep{smith13}. Thus, the ionized gas pressure is directly constrained observationally, and so the pressure there is known regardless of the geometry that creates it. While the pressure within this H~{\sc ii} region is roughly uniform, the ram pressure of the wind drops with radius from the star if we assume a steady BSG wind ($R^{-2}$ density profile). Then $R$ is the radius where the two balance, given by \begin{displaymath} R \ = \ 0.05 \ \Big{(} \dot{M}_{-7} \ V_{300} \Big{)}^{1/2} \Big{(}\frac{n_e}{500 \ {\rm cm}^{-3}} \Big{)}^{-1/2} \ {\rm pc}, \end{displaymath} \noindent where we have assumed $T = 10^4$\,K in the H~{\sc ii} region, $\dot{M}_{-7}$ is the BSG wind mass-loss rate in units of $10^{-7}$\,M$_{\odot}$\,yr$^{-1}$, and $V_{300}$ is the wind speed in units of 300\,km\,s$^{-1}$. We assumed a value for $V_{\rm BSG}$ of 300\,km\,s$^{-1}$, as above. These fiducial values are similar to the values adopted for the progenitor of SN~1987A by \citet{cd95}. With these values, the stand-off shock will be at $R \approx 0.05$\,pc from the BSG. This is about 25\% of the radius of the ring (note that both SN~1987A and SBW1 have the same ring radius of $\sim$0.2 pc). In the case of SBW1, 25\% of the ring radius roughly matches the location of the observed inner peaks of hot dust and enhanced H$\alpha$ emission in images, which is why we chose these fiducial values. Since the innermost dust near the shock front will be the hottest and brightest because it is radiatively heated by the star, we argued \citep{smith13} that this physical scenario may give a plausible explanation for the structures inside the ring. We subsequently proposed to obtain UV spectra to directly constrain SBW1's mass-loss rate in order to test this picture. We were therefore surprised to find a mass-loss rate for SBW1 that is at least an order of magnitude lower than the fiducial value above. With $\mdot <$ 10$^{-8}$ $\msun$ yr$^{-1}$ (a conservative upper limit), the radius of the stand-off shock between the BSG wind and the ionized photoevaporative flow should be much smaller, roughly $<$0.015 pc or only about 5--10\% of the ring's radius. Essentially, the BSG wind is so weak that it would be overwhelmed by the gas pressue of the photionized photoevaporative flow. Colliding winds may therefore have difficulty explaining the pile-up of dust at the location of the observed IR peaks in images \citep{smith13}. A renewed investigation of this problem using hydrodynamic simulations is warranted. What, then, causes the peaks of dust emission at $\sim$25\% of the radius of the ring (at $R$$\approx$0.05 pc from the star)? As noted in our previous paper \citep{smith13}, the observed dust temperature estimated from the SED is only about 190~K (and the expected equilibrium temperature is even lower at that radius) so 0.05 pc cannot mark the dust vaporization radius. Something else must hold back the dust and prevent it from flowing closer to the star. A possibility is that direct stellar radiation pressure on dust grains helps keep them at bay, and that collisions couple this radiation pressure on dust to the gas. Indeed, the magnitude of the radiation pressure $L / (4 \pi R^2 c)$ inside the ring, for our derived stellar parameters of SBW1, is comparable to or greater than the inferred ionized gas pressure for $T = 10^4$ K and $n_e$=500 cm$^{-3}$, suggesting that direct radiation pressure on dust should affect the structure and dynamics of the interior of the ring. So far, radiation pressure has not been included in simulations aiming to explain the origin and shaping of BSG rings like the ones around SN~1987A and SBW1. However, the weakness of the observed wind from SBW1 reported here (as well as the inferred weakness of the wind of SN~1987A's progenitor) suggest that this should be undertaken. Examining the hydrodynamics including radiation pressure is beyond the scope of this paper, but we note that the problem is reminiscent of recent studies of the dynamics and structure of dusty H~{\sc ii} regions, where radiation pressure on dust is also found to be important \citep{km09,draine11,kim+16}. The relative influence of radiation pressure is even stronger in the case of SBW1 due to its extremely weak stellar wind for its luminosity. Another possibility, which is difficult to rule out, is that the inner dust peaks arise from a past eruptive mass ejection akin to LBV eruptions \citep{smith11}. While the BSG wind cannot form dust in its steady wind, it could potentially form dust in an episodic ejection of a dense shell (see, e.g., \citealt{kochanek11}). This dust shell might then expand until it is stopped by the pressure of the photoevaporative flow, leaving a cavity in its wake to be filled by the very weak BSG wind. In this case it would be the momentum of the (hypothetical) eruptive mass ejection rather than the ram pressure of the BSG wind that would set the location of the inner dust peaks. This scenario is admittedly somewhat {\it ad hoc}, but there is precedent for it. Sequential episodic ejections of rings have been inferred based on direct proper motions of the ring nebula around the massive binary RY Scuti \citep{smith+11}, for example. For a somewhat different type of system, hydrodynamic simulations of nova eruptions inside a slow, equatorially concentrated CSM produced by RLOF can yield a similar torus structure with inner density peaks \citep{booth16}. \subsection{Spindown} The very weak wind of SBW1 has an important consequence regarding the star's rotational evolution (e.g., \citealt{meynet11}). Such a low mass-loss rate will impair the star's ability to shed angular momentum via its wind. SBW1 currently has a rather slow rotation rate, with an equatorial rotation speed of only about 40 km s$^{-1}$ \citep{taylor14}, which is only about 5\% of its critical rotation speed. The current slow rotation rate coupled with the currently observed very low wind mass-loss rate presents a puzzle in connection with the observed ring nebula. As noted in the Introduction, most scenarios to explain the existence of ring nebulae like the ones around SN~1987A and SBW1 invoke either (1) mass transfer through RLOF in an interacting binary (which would spin up the mass gainer and then shed mass through the outer Lagrange point), (2) the merger of a close binary system resulting in a rapid rotator that excretes a disk, ring, or torus in the merger, or (3) post-RSG contraction to a BSG, spinning the star up to a rapidly rotating star that sheds an equatorial disk. All of these include a star that is rotating at or close to critical rotation when the ring is ejected. In the case of SBW1, the ring is only about 10$^4$ yr old \citep{sbw,smith13}. The puzzle, then, is how a star can go from (presumably) nearly critical rotation (several 10$^2$ $\kms$) to being such a slow rotator (only 40 $\kms$) in such a short time if its wind is very weak. In this time, the star would shed a tiny fraction ($\sim$10$^{-5}$ or less) of its total mass. Magnetic breaking would be the key mechanism to spin down the star, and indeed, it has been suggested that a stellar merger event - which might eject a ring - might also lead to very strong stellar magnetic fields \citep{schneider16}. However, one expects the loss of angular momentum via magnetic breaking to be directly proportional to the mass-loss rate, which in the case of SBW1 is exceedingly low. Even for massive stars with very strong (a few to several kG) fields and stronger winds, the spin-down timescale is a few to several Myr \citep{uddoula09}, not 10$^4$ yr. Indeed, using a parameterized estimate for the spin-down time from Equation 25 in \citet{uddoula09}, and adopting a generous 3 kG magnetic field, the parameters we estimate for SBW1 would suggest a spin-down timescale of $>$6 Myr. It is therefore difficult to understand how the star could have slowed its rotation rate during the age of the nebula of only 10$^4$ yr unless the mass-loss rate was much higher in the past. Ways out of this puzzle may require some different ideas. Observationally, at least, the gas and dust that partly fills the interior of the ring \citep{smith13} could be interpreted as evidence for a previous high $\mdot$ \ phase, with a slow, dense, dusty wind that followed a merger and the ring's ejection. Perhaps a highly time variable wind or eruption needs to be invoked to help resolve this issue. Alternatively, perhaps a merger scenario different from proposed models is in order. For example, a merger of two blue stars (rather than a RSG) may lead to an envelope that is out of thermal equilibrium, as rotational energy is used to heat the merger product's envelope. The subsequent inflation of that envelope might allow a merger product to have a slow surface rotation rate at such a young age after a merger event. It is difficult to see how very rapid rotation can be avoided in a scenario wherein a merger occurs as a RSG, and then the merger product contracts to the blue while also maintaining a very low mass-loss rate. The low BSG wind mass-loss rate that we derive here is therefore an important constraint for models that aim to explain the origin and shaping of such ring nebulae with a merger event. SBW1 may be an interesting target for spectropolarimetry to investigate the possibility of a strong magnetic field, although this may be complicated by large interstellar polarization. It is interesting to note that some models predict that magnetic massive stars can avoid the RSG phase altogether, staying blue and exploding as BSGs \citep{petermann15}. How this BSG star can avoid driving a much stronger wind with its current luminosity remains puzzling.

1607.01055

1607

1607.04994_arXiv.txt

Observations of pulsars across the radio spectrum are revealing a dependence of the characteristic scattering time ($\tau$) on frequency, which is more complex than the simple power law with a theoretically predicted power law index. In this paper we investigate these effects using simulated pulsar data at frequencies below 300 MHz. We investigate different scattering mechanisms, namely isotropic and anisotropic scattering, by thin screens along the line of sight, and the particular frequency dependent impact on pulsar profiles and scattering time scales of each. We also consider how the screen shape, location and offset along the line of sight lead to specific observable effects. We evaluate how well forward fitting techniques perform in determining $\tau$. We investigate the systematic errors in $\tau$ associated with the use of an incorrect fitting method and with the determination of an off-pulse baseline. Our simulations provide examples of average pulse profiles at various frequencies. Using these we compute spectra of $\tau$ and mean flux for different scattering setups. We identify setups that lead to deviations from the simple theoretical picture. This work provides a framework for interpretation of upcoming low frequency data, both in terms of modelling the interstellar medium and understanding intrinsic emission properties of pulsars.

Observed radio pulses from pulsars pass through the ionised interstellar medium (ISM) before terrestrial detections are made. As the waves propagate through the ISM, inhomogeneities in the electron densities of dense, ionised regions cause the radio waves to scatter. This scattering is observed through, amongst other effects, the frequency dependent temporal broadening of the received radio pulses. The temporal broadening typically leads to characteristic exponential \textit{tails} in averaged pulsar profiles (e.g. \citealt{Lohmer2001}). Understanding the frequency dependence of the temporal broadening allows us to investigate the properties of the scattering media. Compensating for scattering effects can be used to infer the intrinsic radio emission characteristics of pulsars. The scattering tails of the broadened profiles are routinely modelled with exponential functions that take the form $\tau^{-1}e^{-t/\tau}$, where $\tau$ is referred to as the \textit{characteristic scattering time}. This functional form, as described in detail in Sec. \ref{sec:scatsetup}, is based on the assumption of a single, thin scattering screen extending infinitely transverse to the line of sight to the pulsar. The screen is invoked as an approximation of the combined scattering that the pulse signal undergoes in the extended ISM along the line of sight \citep{Scheuer1968, Cronyn1970,Williamson1972, Williamson1973}. The characteristic scattering time $\tau$ depends on the properties of the scattering screen as well as its location along the line of sight (see eq. \ref{eq:tau}) and it is maximised for a midway screen. Apart from approximating the overall scattering geometry to a single scatterer, historic studies have also investigated distinct electron density distribution models for the screen. These include an isotropic Gaussian electron column density wavenumber spectrum, which leads to a temporal broadening dependence on frequency, $\tau \propto \nu^{-4}$, \citep{LeeJokipii1976, Lang1971}. Alternatively power law wavenumber spectra, characterised by an inner and outer cut-off length scale, are adopted to model the plasma inhomogeneities. Most often a Kolmogorov power spectrum, which describes the turbulence of a neutral gas, is preferred \citep{Rickett1977, RCB1984, Rickett1990}. A Kolmogorov spectrum will have a frequency dependence of $\tau \propto \nu^{-4.4}$, for wavenumbers lying between the inner and the outer cut-off length scales\footnote{Note that for power-law spectra, where $\tau \propto \nu^{-\alpha}$, the fluctuation spectral index $\beta$ is often quoted instead of the spectral index $\alpha$. In the 2D case, for $\alpha \geq 4$, they are related through: $\alpha = 2\beta/(\beta -2)$. A Kolmogorov spectrum with $\alpha = 4.4$ has $\beta = 11/3.$ Also note that the symbols $\alpha$ and $\beta$ are sometimes used interchangeably.}. The scale at which electron density inhomogeneities lead to strong diffractive scintillation and consequently the observed pulse broadening analysed in this paper, is set by the field coherence scale, $s_{0} = 1/(k \theta)$, with $k$ the radio wavenumber \citep{Rickett1990}. For strong diffractive scintillation $s_{0} \ll L_{F}$, where $L_{F} = (D/k)^{1/2}$, is the Fresnel scale (e.g. \citealt{SmithThompson1988}. Refer to Fig. \ref{GeoGraph} for descriptions of $\theta$ and $D$.). Estimations of $\tau$ values from pulsar profile observations are done by either deconvolution of the folded pulse profile with a broadening function such as a one-sided exponential function through e.g. the $CLEAN$-based algorithms \citep{Bhat2003CLEAN}, or by forward modelling which convolves a template unscattered pulse profile (or alternatively an observed high frequency pulse profile) with a parametrised broadening function and compares the model to the observed profile (e.g. \citealt{Lohmer2001,Lohmer2004}). It should be noted that using a high frequency observation as a template will, for pulsars that exhibit intrinsic frequency evolution, introduce biases to measurements of scattering parameters. A deconvolution method does not suffer from this effect. Alternatively $\tau$ can be estimated by means of the \textit{scintillation bandwidth}, $\delta f$ \citep{LyneRickett1968}. The scattering process imparts frequency dependent phase changes on the passing radio waves. An interference pattern will emerge for phase changes of the order $2\pi \delta f \tau \approx 1$ radian. A measurement of $\delta f$, over which the interference effects are observed in the dynamic spectrum, can therefore lead to an approximation of $\tau$ \citep{Cordes1986}. Using the ISM transfer function implied by the thin screen model, results in fits of $\tau$ in observed data that follow the expected spectral dependence \citep{Armstrong1995}. However, deviations from the theoretical spectral dependences have also been proposed and observed - both with spectral index $\alpha$ (where $\tau \propto \nu^{-\alpha}$) smaller than predicted \citep{Bhat2004, Lohmer2001, Lewandowski2015} and larger than predicted \citep{Rickett1990, LambertRickett1999,Tuntsov2012}. Different models will lead to different values for $\alpha$, but the critical minimum for strong diffractive scattering, is $\alpha = 4$ \citep{RNB1986}. Anisotropic scatterers, i.e. media for which the distribution of scattering angles exhibit directionality have been considered as causes for steeper spectra \citep{Tuntsov2012}. They were first invoked to explain organised patterns in the dynamic spectra of pulsar observations \citep{Gupta1994}. \citet{Stinebring2001} found that complex patterns in dynamic spectra are often related to parabolic arcs in the secondary (power) spectra of pulsar observations, which can be explained by anisotropic scattering \citep{Walker2004}. There has also been some successes in associating anisotropy with observations of Intra Day Variability (IDV) \citep{Bignalletal2003, Tuntsov2012}. Extreme Scattering Events (ESE's) as observed by \citet{Fiedler1987} and analysed by e.g. \citet{Walker2006} have motivated works studying the impact of alternative scatterers, such as self-gravitating, AU-sized ionised gas clouds. A one dimensional Kolmogorov power spectrum in electron density distribution is often used to model extreme anisotropy. On the contrary, weaker dependences on frequency have been attributed to screens which have a limited physical size rather than being infinite transverse to the line of sight as most models assume \citep{CordesLazio2001}. Beyond the edges of such a screen, no radio waves coming from the pulsar will be bent back into the line of sight of the observer, leading to no further scatter broadening with increasing wavelength. Such configurations also lead to a loss in observable flux. A physical inner scale of the scattering screen, that is smaller than the diffractive scattering scale, will restrict diffraction to a maximum scattering angle. This too will lead to less power observed at large angles, or alternatively flatter $\tau$ spectra with lower power law indices, at long wavelengths. In the case of an inner scale however, flux is preserved \citep{RickettJohnston2009}. Until recently, characterising the frequency dependence of temporal scatter broadening was done using pulsar observations taken at different receivers and telescopes at different epochs. With the construction of telescopes such as the Low-Frequency Array (LOFAR; \citealt{vanHaarlem2013}) and the Murchison Widefield Array (MWA; \citealt{Tingay2012}), we have the advantage of studying pulsars at low frequencies, where scattering is most pronounced, using broad bands (see \citealt{Stappers2011} for an outline of LOFAR's pulsar observing modes). This allows accurate measurements of $\tau$ spectra and simultaneous studies of the flux spectrum and profile evolution of the pulsar. Analysis of broad band data will therefore test our current understanding of pulsar scattering. Recent works showing average pulse profiles in the LOFAR bands (30 - 90 MHz and 110 - 190 MHz) confirm the expectation of scatter broadened pulse shapes at these frequencies \citep{Bilous2015, Pilia2015}. In this paper we provide a framework based on simulations, for interpreting observations such as the ones presented in the LOFAR works. This can be used to gain insight on the size, location and nature of scattering screens. The paper is split into sections as follows. In Sec.~\ref{sec:simdataandtech} we describe the general scattering setup, the simulated data and the techniques by which we fit the data to retrieve $\tau$ spectra. The mathematical formalism for examples of the broadening functions for isotropic and anisotropic scattering is also presented in this section. In Sec.~\ref{sec:results} we present the results of our experiments. These include example profile shapes as well as $\tau$ and flux spectra for infinite and finite screens respectively. We end with a discussion of our results and their relevance to broad band observations in Sec.~\ref{sec:discussion}, followed by a short conclusion.

We have showed that our train+DC method is an efficient forward fitting technique for scattered profiles. Our investigations show flatter $\tau$ spectra (than what is theoretically expected) when using an isotropic model to fit profiles shaped by an anisotropic scattering mechanism. Flatter $\tau$ spectra can also be the result of low frequency profile shape changes caused by a truncated scattering screen. Over a large enough sampled frequency range, such a $\tau$ spectrum will exhibit a break. Analysing the impact of scattering on flux spectra, has revealed turnovers at low frequencies due to baseline effects. Similarly truncated screens can cause turnovers in flux spectra. These turnovers are a function of the screen size and its scattering strength. We anticipate new data from instruments such as LOFAR and the MWA in the frequency range studied here. This paper will serve as a framework to compare real data to the simulated results obtained here. Our analysis of scattering imprints is by no means complete, as we have only investigated simple profile shapes and Gaussian scattering mechanisms. However, our results indicate where attention should be paid in terms of creating accurate fitting methods as well as describing possible sources of deviations to theoretically predicted results.

1607.04994

1607

1607.03861_arXiv.txt

Detection of the Epoch of Reionization HI signal requires a precise understanding of the intervening galaxies and AGN, both for instrumental calibration and foreground removal. We present a catalogue of \Total extragalactic sources at 182~MHz detected in the RA=0 field of the Murchison Widefield Array Epoch of Reionization observation programme. Motivated by unprecedented requirements for precision and reliability we develop new methods for source finding and selection. We apply machine learning methods to self-consistently classify the relative reliability of 9490 source candidates. A subset of \Cand are selected based on reliability class and signal-to-noise ratio criteria. These are statistically cross-matched to four other radio surveys using both position and flux density information. We find 7369 sources to have confident matches, including 90 partially resolved sources that split into a total of 192 sub-components. An additional \NewSrcs unmatched sources are included as new radio detections. The catalogue sources have a median spectral index of $-$0.85. Spectral flattening is seen toward lower frequencies with a median of $-$0.71 predicted at 182~MHz. The astrometric error is $7$" compared to a 2.3' beam FWHM. The resulting catalogue covers $\sim$1400~deg$^2$ and is complete to approximately 80~mJy within half beam power. This provides the most reliable discrete source sky model available to date in the MWA EoR0 field for precision foreground subtraction.

\label{sec:Intro} Observations of the Epoch of Reionization (EoR) 21~cm neutral Hydrogen (HI) signal are one of the key science goals of the Murchison Widefield Array \citep[MWA;][]{Lonsdale2009,Tingay2013,Bowman2013}. Astrophysical foreground sources are estimated to be 4--5 orders of magnitude brighter than this signal, presenting a major obstacle and motivating a careful dedicated survey to be used for both calibration and foreground power removal within the EoR analysis. During the commissioning phase of the MWA and early development of the EoR analysis pipeline, the Molonglo Reference Catalogue (MRC; \citealt{Large1991}) was used to model the foregrounds. The MRC is complete to 1~Jy at 408~MHz or about 2~Jy at 182~MHz with an assumed average spectral index of $-$0.8. This is not only much shallower than desired, but large errors are introduced by this naive flux density extrapolation. The MRC catalogue was later replaced by the MWA Commissioning Survey \citep{Hurley-Walker2014}, giving us a deeper and frequency-specific sky model that greatly improved calibration and foreground power subtraction. Unfortunately the MWACS does not cover the northernmost $5^{\circ}$ of the EoR fields, and as an early product of a partially built instrument it contains large flux density and astrometric uncertainties. While the GaLactic and Extragalactic All-sky MWA survey \citep[GLEAM;][Hurley-Walker {\it in prep.}]{Wayth2015} was still underway, an extragalactic survey tuned to the requirements of EoR foreground subtraction in the MWA EoR0 (RA=0, Dec=-27) field was initiated. Processing MWA data is different in a number of ways from traditional interferometric radio data processing. Traditional radio analysis recipes are well tuned for arrays with narrow fields of view, stable beams, sparse instantaneous $uv$ coverage, and limited source confusion. MWA data break these assumptions. The MWA primary beam is $\sim$1400~deg$^2$ and changes with time as the field drifts through, the instantaneous $uv$ coverage is excellent, and the continuum confusion limit is reached very quickly. Background sky regions reach a signal-to-noise ratio (SNR) $\sim$~10 in a 112~sec integration. For these reasons most MWA analyses have deconvolved sources on times scales of a few minutes or less to minimize time-dependent beam effects and leverage the snapshot $uv$ coverage. Many different deconvolution algorithms may be chosen, though typically a pixel based algorithm such as CLEAN is used (i.e. source components are located at pixel centres and may take either positive or negative values). Going from radio deconvolution products to a source catalogue is often performed by fitting sources in a restored image. For the MWA this has meant restoring the individual snapshot observations, mapping to celestial coordinates to remove widefield distortions, and then co-adding or mosaicing (e.g. \citealt{Wayth2015}). Combining snapshots through co-adding increases the SNR of real sources and drives down the amplitude of false side lobe sources as the array beam (PSF) rotates. While this reduces contamination, it removes information that is valuable for determining reliability, that is whether or not a source is in fact true and real. True sources should be detected consistently across all observations, while noise and side lobes will vary in time. This information is lost if source finding is not performed prior to image stacking. Reliability is a primary concern for the EoR foreground model as it will be used for both calibration and subtraction \citep{Barry2016}. Various source extractors have been developed that isolate flux density peaks in an image, fit an assumed PSF or morphological shape, and measure the integrated flux density and background rms (e.g. \citealt{Hancock2012}). This approach does well but is not without its limitations. \citet{Hopkins2015} demonstrate clear variations in the performance of eleven different source finders in terms of completeness, reliability, and measurement accuracy. Blended or extended sources were particularly troublesome despite the fact that all sources in the test images were artificially positioned at pixel-centers. In this work we take a novel approach to source finding that does not require the use of restored images or an assumed source shape. Individual snapshots are deconvolved using the Fast Holographic Deconvolution \citep[][Sullivan et al., {\it in prep.}]{Sullivan2012} software package. FHD deconvolution is similar to CLEAN in that sources are iteratively removed using point-like source components with a CLEAN gain. However, FHD differs in using a full direction-dependent PSF, centroiding each source component (not fixed at pixel centers), and using only positive components. For a bright unresolved point source, FHD will create a set of positive source components, each with a floating-point precision position, tightly scattered about the actual source position ($\ll$ PSF FWHM). Diffuse sources will be modeled by a cloud of source components approximating the extended flux density distribution. The FHD source components trace the sub-resolution flux density distribution of sources remarkably well. We can therefore identify and extract sources simply by spatially clustering the components regardless of shape. This does not require a restored image, nor the assumption of a PSF or morphological model. Machine-learning classification methods are then used to self-consistently assess source reliability. This process for source finding, measurement, and classification has been termed KATALOGSS (KDD Astrometry, Trueness, and Apparent Luminosity of Galaxies in Snapshot Surveys; hereafter abbreviated to KGS). We use the KGS results in combination with a typical signal-to-noise detection threshold and cross-matching to maximize the overall completeness and reliability of the final catalogue. This paper presents these methods and the resulting source survey of the MWA EoR0 field. The observations and pre-processing are described in \S\ref{sec:Data}; the process of source finding and association across snapshots in \S\ref{sec:SourceFinding}; and the reliability classification is detailed in \S\ref{sec:Classification}. In \S\ref{sec_XMatch} we introduce the Positional Update and Matching Algorithm (PUMA) used to cross-match the catalogue to other radio surveys and fit for the power-law spectral index. The final catalogue is described in \S\ref{sec:catalogue} along with measures of the astrometric accuracy, spectral index distribution, and completeness. We further identify \NewSrcs new sources previously undetected at radio frequencies and discuss potential associations in \S\ref{sec:newsources}.

\label{sec:Conclusions} We have presented a catalogue of \Total extragalactic radio sources in the MWA EoR RA=0 field at 182~MHz. This survey was motivated by the EoR analysis and the need for an accurate foreground model. The foreground catalogue is used for the purposes of calibration and subtraction, and is the predominant systematic hurdle to making an EoR detection. A catalogue of high precision, reliability, and completeness at low frequencies in the southern sky is required. To this aim, new methods were tested for deconvolution and source finding. An in-depth analysis confirmed source reliability and excluded contamination from noise and side lobe sources. Seventy-five consecutive snapshot observations were processed, covering 2.5 hours approximately centered on zenith. These were independently deconvolved using FHD, resulting in an array of centroided positive components. Source finding was done by spatially clustering the deconvolved components into source candidates for each observation. This approach was chosen to reduce errors inherent to the process of producing a restored image and fitting simplified morphological shapes to sources in the image plane. By source finding independently for each snapshot we retain information on the detection frequency, a valuable diagnostic of source reliability. We identified 9490 unique source candidates detected in at least two snapshots. Radio surveys are afflicted by contamination from both noise and side lobes. This is especially troublesome outside of the typical half-power cutoff of the primary beam response. We wished to push this boundary to 5\% beam response while maximizing both completeness and reliability. Given the basic knowledge of how true sources, noise, and side lobes are expected to behave in the data, we used machine learning methods to categorize the source candidates into 10 "reliability" classes based on their observed properties. This gave us a more informative indication of confidence than SNR or detection frequency alone, while remaining entirely self-consistent. For the purposes of this work, we selected only \Cand robustly detected above $5\sigma$ or reliability class $R_{\rm class}<7$. These were probabilistically cross-matched to overlapping radio surveys from 74~MHz to 1400~MHz using both positional and broad-band spectral information. Outliers, complex sources, and unmatched sources were flagged and individually investigated. The reliability classification and postage stamp images further aided the decision to include, modify, or exclude a source. Individually inspecting outliers for consistency resulted in the serendipitous identification of variable, peaked or steep spectrum, as well as morphologically complex sources. Among sources unmatched to another catalogue, \NewSrcs are included in the final catalogue. Some of these detections can be attributed to the low surface-brightness sensitivity of the MWA. Others may be ultra steep spectrum HzRG candidates. We identify several possible associations with galaxy clusters or groups. The final catalogue contains \Total sources. Compared to NVSS or SUMSS the median absolute offset is only 7" compared to a 2.3' beam FWHM. No significant evidence is found for flux scale bias in the catalogue. The median broad-band spectral index is found to be $-$0.85 but is dependent on the catalogues matched for each source. A trend in median spectral index is observed as a function of frequency, possibly indicating spectral flattening among a sub-population of sources. The median spectral index at 182~MHz is estimated at $-$0.71 using a second order polynomial fit to sources detected in 3 or more surveys. This catalogue provides the most reliable discrete source sky model available to date in the MWA EoR0 field for foreground subtraction. The impact of the improved foreground model on the EoR power spectrum will be presented in Carroll et al. ({\it in prep.}). Based on the lessons learned through the creation of this catalogue, particularly the value of repeated snapshots in assessing source reliability, we are currently taking new observations to extend the EoR foreground survey across the southern sky.

1607.03861

1607

1607.06443_arXiv.txt

{Ultra-compact dwarf galaxies (UCDs) share many properties with globular clusters (GCs) and are found in similar environments. Here, a large sample of UCDs and GCs around NGC~1399, the central giant elliptical of the Fornax galaxy cluster, is used to infer their formation history and also to shed light on the formation of NGC~1399 itself. We assumed that all GCs and UCDs in our sample are the result of star cluster (SC) formation processes and used them as tracers of past star formation activities. After correcting our GC/UCD sample for mass loss, we interpreted their overall mass function to be a superposition of SC populations that formed coevally during different formation epochs. The SC masses of each population were distributed according to the embedded cluster mass function (ECMF), a pure power law with the slope $-\beta$. Each ECMF was characterized by a stellar upper mass limit, $M_{\mathrm{max}}$, which depended on the star formation rate (SFR). We decomposed the observed GC/UCD mass function into individual SC populations and converted $M_{\mathrm{max}}$ of each SC population to an SFR. The overall distribution of SFRs reveals under which conditions the GC/UCD sample around NGC~1399 formed. Considering the constraints set by the age of the GCs/UCDs and the present stellar mass of NGC~1399, we found that the formation of the GCs/UCDs can be well explained within our framework with values for $\beta$ below 2.3. This finding agrees very well with the observation in young SCs where $\beta \approx 2.0$ is usually found. Even though we took into account that some of the most massive objects might not be genuine SCs and applied different corrections for the mass loss, we found that these considerations do not influence much the outcome. We derived the peak SFRs to be between approximately 300 and 3000 $M_{\odot}\mathrm{yr}^{-1}$, which matches the SFRs observed in massive high-redshift sub-millimeter galaxies and an SFR estimate inferred from NGC~1399 based on the so-called downsizing picture, meaning that more massive galaxies must have formed over shorter periods of time. Our findings give rise to the interpretation that NGC~1399 and its GC/UCD system formed in a relatively short, intense starburst early on.}

\label{intro} The Fornax cluster around the giant elliptical \object{NGC~1399} possesses a very rich system of globular clusters (GCs). \citet{gregg09} estimated $11\,100 \pm 2\,400$ GCs in total within a radius of 0$^{\circ} \!$.9 around \object{NGC~1399}, corresponding to a radius of 320~kpc at a distance of 19~Mpc based on a distance modulus of $(m-M) = 31.4 \pm 0.2$ (\citealt{dirsch03, dirsch04}, see also \citealt{ferrarese00, blakeslee09}). Within $15'$, which corresponds to roughly 83~kpc, \citet{dirsch03} estimated a smaller number of $6\,450 \pm 700$ GCs. Of particular interest is the upper end of the GC mass function, which is dominated by ultra-compact dwarf galaxies (UCDs). The term UCD \citep{phillipps01} is somewhat misleading since these objects are not necessarily dwarf galaxies: They show similarities with nuclei of dwarf galaxies, but they also share many properties with GCs, for which reason two main formation scenarios for UCDs in the Fornax cluster are discussed in the recent literature \citep[e.g.,][]{evstigneeva08, chilingarian08, chilingarian11, mieske12, francis12}: \begin{enumerate}[(a)] \item UCDs are dynamically evolved nucleated dwarf galaxies, from which outer stellar components were removed while orbiting in the center of the Fornax galaxy cluster and suffering from its strong tidal gravitational field \citep[threshing scenario; e.g.,][]{bekki01, bekki03, drinkwater03, goerdt08, thomas08, pfeffer13, pfeffer14, pfeffer16}. \item UCDs are the brightest GCs of the rich \object{NGC~1399} globular cluster system and thus they are the result of star cluster formation processes \citep[e.g.,][]{mieske02, mieske12}. Moreover, it has been proposed that the very massive UCDs could also form in the so-called merged star cluster scenario, where massive complexes of star clusters merge and thereby form a super star cluster \citep{kroupa98, fellhauer&kroupa02, mieske06, bruens11, bruens12}. \end{enumerate} Although almost two decades of research have passed since the discovery of UCDs \citep{minniti98, hilker99, drinkwater00}, their nature is still puzzling \citep[e.g.,][]{phillipps13}: Various investigations found evidence for both formation scenarios, but could not confirm either of the hypotheses with certainty \citep[e.g.,][]{mieske04, evstigneeva08, wittmann16}. However, there is growing evidence that UCDs are rather a "mixed bag of objects" \citep{hilker09book} than a distinct type of object \citep[e.g.,][]{chilingarian11}. In the Fornax galaxy cluster, GCs have effective radii from smaller than 1 and up to 10~pc with an average of around 3~pc \citep{masters10, puzia14}. Their masses range from $10^4~M_{\odot}$ up to lower than $10^7~M_{\odot}$ \citep[][their Fig.~15]{puzia14}. UCDs, on the other hand, have some overlap with GCs, but also extend the parameter space to larger sizes and masses: their effective radii range from a few pc to up to 100~pc \citep[e.g.,][]{drinkwater03, evstigneeva07, evstigneeva08, hilker07, mieske08}, while their masses lie between $10^6~M_{\odot}$ and lower than $10^8~M_{\odot}$ \citep[e.g.,][]{drinkwater03, evstigneeva07, hilker07, mieske08, frank11}, bridging the region between classical GCs and compact elliptical galaxies. In the literature, several arguments support that most of the Fornax UCDs are very massive GCs: \begin{itemize} \item The luminosities of GCs and UCDs are distributed smoothly and their luminosity functions do not show any bimodality \citep{mieske02, mieske04}. Furthermore, the UCDs in Fornax are consistent with being drawn from the bright tail of the GC luminosity function. However, a small excess at the bright end is statistically possible \citep{mieske04, gregg09, mieske12}. \item GCs and UCDs form a continuum in the luminosity-size plane \citep{mieske06}. \item UCDs exhibit the full range of (high and low) metallicities as observed for GCs \citep{francis12}. \item The spread of age and metallicity of the UCDs is consistent with that observed for GCs \citep{francis12}. \item Most of the UCDs have super-solar $\alpha$-element abundances, implying short formation times, similar to those of GCs \citep{francis12}. \end{itemize} However, even if the UCDs were not genuine GCs, several findings suggest that they are at least the result of a star cluster (SC) formation process: \begin{itemize} \item \citet{kissler-patig06} placed young massive clusters (YMCs) with masses higher than $10^7 M_{\odot}$ on three different scaling relations and found their positions to be similar to those of the UCDs, in particular for the most massive YMCs. \item UCDs have metallicities close to but slightly below those of YMCs of comparable masses \citep{mieske06}. \item Fitting a simple stellar population model to the spectra of UCDs reveals that UCDs are in agreement with a pure stellar content \citep{hilker07, chilingarian11} so that no DM component is needed in UCDs within their 1--3 half-mass radii \citep{hilker07}. \citet{chilingarian11} found almost all UCDs to be compatible with no DM inside. More recently, \citet[their Table~3]{mieske13} found that only the most massive UCDs require an additional mass component to compensate the elevated $M/L$ ratio, which they suggested might be massive black holes, while the lower-mass UCDs do not need any form of additional, non-luminous matter \citep[see also][]{dabringhausen09, dabringhausen10, dabringhausen12}. \end{itemize} Following \citet{mieske12}, who found that most UCDs are compatible with being formed in the same way as GCs, we assume that it is justified to treat UCDs, like GCs, as (very) massive SCs and assume that they formed in SC formation processes. We aim to determine the conditions under which this occurred and what it indicates about the assembly history of NGC~1399. However, as mentioned before, even if the majority of the UCDs are compatible with being giant ancient SCs, it is very likely that some of the most massive UCDs did not form in an SC formation process. For instance, some of the UCDs exhibit extended surface brightness profiles or even tidal features \citep{richtler05, voggel16} or appear asymmetric or elongated \citep{wittmann16}. In addition, \citet{voggel16} detected for the first time the tendency of GCs to cluster around UCDs. In the Virgo galaxy cluster, a fraction of the very massive UCDs are thought to be of galaxy origin \citep[e.g.,][]{strader13, seth14, liu15a, liu15b, norris15}, while a fraction of the faintest (and thus lowest mass) UCDs might instead be related to compact SCs \citep{brodie11}. Thus, in two additional scenarios we examine how the distribution of required SFRs changes if some of the most massive UCDs in Fornax are excluded from the GC/UCD mass distribution.

\label{sect_concl} We combined a spectroscopic and a photometric sample of GCs/UCDs around the giant elliptical NGC~1399 in the Fornax galaxy cluster to derive their overall mass distribution (Sect.~\ref{sect_sample}) and corrected it for stellar and dynamical evolution (Sect.~\ref{sect_corrections}) to obtain their mass function at the time of formation. Then, this "natal" mass function was decomposed into individual SC populations distributed according to the ECMF (Sect.~\ref{sect_decomposition}, schematic plot in Fig.~\ref{fig_decompo_sketch}). The upper mass limit of each population, $M_{\mathrm{max}}$, was converted into an SFR according to the SFR-$M_{\mathrm{max}}$ relation \citep{weidner04} as was fit in \citet{schulz15} (Sect.~\ref{sect_sfr.distr}). The resulting SFR distributions (Fig.~\ref{fig_sfr}) reveal which SFRs are required to form the entire GCs/UCDs system around NGC~1399. When restoring the natal GC/UCD mass function, we assumed different lifetimes of the SCs, parameterized by $t_4$, to account for the variable strength of the tidal field depending on the distance to the center of NGC~1399 (Sect.~\ref{sect_corrections}). Moreover, we used three different approaches regarding the treatment of the natal GC/UCD mass function. First, in the standard approach we assumed that all GCs/UCDs were formed in an SC formation process and are therefore ancient SCs. Second, since at least the most massive UCD in our sample, UCD3, shows hints of being a merged SC or the nucleus of a stripped dwarf galaxy, we excluded this object from our sample. Third, \citet{pfeffer14} derived a possible distribution of dwarf galaxies whose envelopes were stripped away. We assumed that all objects in our GC/UCD mass function compatible with their distribution are stripped nuclei so that only the remaining GCs/UCDs need to be replicated by our method. All these modified samples were then treated as described above, meaning that they were decomposed into individual SC populations from which a distribution of SFRs was deduced based on $M_{\mathrm{max}}$ of each population. Although we made different assumptions regarding the lifetime of the GCs/UCDs and modified the original GC/UCD sample, the outcome was mainly determined by the index $\beta$ of the underlying ECMF, while the influence of the parameter $t_4$ or the modified sample was rather of second order. We extracted from our analysis for each combination of parameters the distribution of SFRs required to build up the observed GC/UCD sample (Fig.~\ref{fig_sfr}), the time this takes, $t_{\mathrm{SCFE,tot}}$, and the total stellar mass, $M_{\mathrm{tot}}$, formed during that time (Table~\ref{tab_statistics}). We found that \begin{itemize} \item the results for $t_{\mathrm{SCFE,tot}}$ and $M_{\mathrm{tot}}$ are well within the constraints set by the age of the GCs/UCDs and the total stellar mass of NGC~1399 for $\beta \lesssim 2.2$, \item the favored values for $\beta$ nicely fit the observations of young SCs where usually $\beta \approx 2.0$ is measured, and \item the peak SFRs derived by us agree well with the range of SFRs observed in massive high-redshift SMGs and also with an estimate based on downsizing. \end{itemize} As discussed in Sect.~\ref{sect_discussion}, our assumption for the lower mass limit of SCs, $M_{\mathrm{min}} = 5~M_{\odot}$, might be an underestimate, given that high SFRs as derived here might suppress the formation of very low-mass SCs. Moreover, the length of one SC formation epoch, $\delta t$, that we used for large $\beta$ is higher than observed in star-forming regions. Increasing $M_{\mathrm{min}}$ and, for large $\beta,$ decreasing $\delta t$ would lead to a shorter total time for SC formation, $t_{\mathrm{SCFE,tot}}$, and to a lower total stellar mass, $M_{\mathrm{tot}}$, so that even higher $\beta$ could still be in agreement with the observational constraints in NGC~1399. Additionally, we regard the standard approach, where all GCs/UCDs are assumed to be genuine SCs, as not very well justified because UCD3, the most massive object in our sample, is clearly not a typical GC. However, except for the higher peak SFRs, caused by UCD3, the outcome from the standard approach does not differ much from the other two approaches. In conclusion, NGC~1399 might have originated from an intense starburst similar to those observed in massive SMGs in the distant Universe. During that starburst, in particular the most massive GCs/UCDs were formed along with many lower mass GCs within a few Gyr. The dissolution and tidal disruption of a part of the GCs/UCDs probably fed the build-up of NGC~1399 and its halo, while a part of the GCs/UCDs was able to survive, allowing us to observe them today. However, here, the GC/UCD sample was analyzed in its entirety without differentiating the red and blue subsamples, that is, the metal-rich and metal-poor GCs/UCDs. In a future paper we will reapply the method described here to the two subsamples to see whether and by how much the formation conditions of the red and blue GCs/UCDs differ.

1607.06443

1607

1607.06958_arXiv.txt

The Peccei-Quinn mechanism presents a neat solution to the strong CP problem. As a by-product, it provides an ideal dark matter candidate, ``the axion", albeit with a tiny mass. Axions therefore can act as dark radiation if excited with large momenta after the end of inflation. Nevertheless, the recent measurement of relativistic degrees of freedom from cosmic microwave background radiation strictly constrains the abundance of such extra relativistic species. We show that ultra-relativistic axions can be abundantly produced if the Peccei-Quinn field was initially displaced from the minimum of the potential. This in lieu places an interesting constraint on the axion dark matter window with large decay constant which is expected to be probed by future experiments. Moreover, an upper bound on the reheating temperature can be placed, which further constrains the thermal history of our Universe.

\label{sec:intro} The strong CP problem is one of the outstanding puzzle of particle physics today. It is a well-known fact that quantum chromodynamics (QCD) allows a CP violating term of the form $ \theta (g_s^2/32 \pi^2) G^{b \mu \nu} \widetilde{G}^b_{\mu \nu}$, where $\theta$ is a constant parameter~\cite{'tHooft:1976up}. The stringent bound on the electric dipole moment of neutron implies that $\vert \theta \vert < 0.7 \times 10^{-11}$~\cite{Crewther:1979pi}. Such a small value for $\theta$ is quite unnatural. This problem can be elegantly solved by the Peccei-Quinn (PQ) mechanism in which a global U(1)$_{\rm PQ}$ symmetry with a chiral anomaly is introduced and the CP violating $\theta$-term can be dynamically relaxed to zero~\cite{Peccei:1977ur}. The corresponding Goldstone boson, the axion~\cite{Weinberg:1977ma} remains massless at the classical level, but acquires a periodic potential and consequently a mass inversely proportional to the PQ symmetry breaking scale, $f_{\rm PQ}$, due to non-perturbative QCD effect~\cite{'tHooft:1976up}. Even though the original PQ proposal with $f_{\rm PQ}$ at the electroweak (EW) scale was soon ruled out by several experiments~\cite{Cheng:1987gp}, other variants of the PQ mechanism circumvent the problem by creating a hierarchy between the PQ breaking scale and the EW one via the introduction of a new Standard Model (SM) complex singlet field whose vacuum expectation value (VEV) breaks the PQ symmetry~\cite{Kim:1979if,Dine:1981rt}. The scale of PQ symmetry breaking is subjected to several observational constraints. For instance, the observation of the supernova SN1987A, white dwarfs and the globular clusters set a lower bound of $(2 \textendash 4) \times 10^{8}$~GeV on the scale of PQ symmetry breaking ~\footnote{More precisely, on the axion decay constant, $f_a = f_{\rm PQ}/N_{\rm DW}$, where $N_{\rm DW}$ is the number of domain walls; $N_{\rm DW} \geq 1$ for KSVZ models~\cite{Kim:1979if}, and $N_{\rm DW} = 6$ for DFSZ models~\cite{Dine:1981rt}. We elaborate on these models in Appendix-A and B. } (see e.g.~\cite{Raffelt:1996wa,Raffelt:2006cw} and references therein). With such a high PQ symmetry breaking scale, axion can be a good dark matter (DM) candidate~\cite{Preskill:1982cy}, whose couplings to other fields are suppressed by powers of $f_{\rm PQ}$. In fact, the energy stored in the coherent oscillations of axions today can make the entirety of the observed DM abundance for $f_{\rm PQ} \sim 7 \times 10^{11} \ {\rm GeV} \ N_{\rm DW} \langle \theta_i^2 \rangle^{-0.84} $~\cite{Turner:1985si} where $ \langle \theta_{i}^2 \rangle$ is the axion misalignment angle at beginning of the axion oscillation phase. On the other hand, many puzzles of early Universe cosmology can be solved by an early epoch of accelerated expansion, ``inflation'' (for a review, see~\cite{Mazumdar:2010sa}). Inflation is also responsible for seeding the primordial perturbations for cosmic microwave background (CMB) and large scale structures. During inflation, if there exists any light field, such as moduli, whose mass is below ${\cal O}(\mathcal{H}_{\rm inf})$, they obtain vacuum induced quantum fluctuations of ${\cal O}({\cal H}_{\rm inf}/2 \pi)$~\cite{Starobinsky:1986fx}, where $\mathcal{H}_{\rm inf}$ denotes the Hubble parameter during inflation. In this case, such a light moduli can obtain large VEV, i.e. ${\cal O}(M_{\rm P})$, where $M_{\rm P} \simeq 2.43 \times 10^{18}$~GeV is the reduced Planck mass. Typically, the moduli behaves like a condensate within our Hubble patch~\cite{Enqvist:2003gh}, and begins its coherent oscillations when the Hubble expansion rate of the Universe drops to the mass of the moduli. Similarly, if the PQ field is light compared to the Hubble expansion rate during inflation, then the PQ field can also be displaced from its minimum, which is determined by $f_{\rm PQ}$~\cite{Starobinsky:1986fx,Linde:1991km}, and consequently after the end of inflation the PQ field will start coherent oscillations when its mass would exceed the time dependent Hubble scale. The PQ field can also be displaced away from $f_{\rm PQ}$ during inflation if it is coupled to the inflaton field~\cite{Linde:1991km}, and later starts oscillating once inflaton begins its own coherent oscllations around the minimum of its potential. If the initial VEV of PQ field during inflation is displaced by $\gg f_{\rm PQ}$, the initial phase of oscillation takes place around the origin. This can lead to the restoration of the PQ symmetry and formation of dangerous topological defects~\cite{Zel'dovich:1975,Sikivie:1982qv}~\footnote{For $N_{\rm DW} = 1$ (i.e. for KSVZ-like models with only one extra heavy quark species), these defects are unstable and decay to cold axions leading to an upper bound on the PQ breaking scale, $f_{\rm PQ} \lesssim (4.6 \textendash 7.2) \times 10^{10} \,{\rm GeV} ( \Omega_a/\Omega_{\rm CDM} )^{0.84}$, where $\Omega_a$ denotes the cold axions abundance and $\Omega_{\rm CDM}$ is the observed abundance of cold dark matter (CDM)~\cite{Hiramatsu:2012gg,Kawasaki:2014sqa}. On the other hand, when $N_{\rm DW} > 1$ (i.e. for DFSZ-like models or KSVZ-like models with several extra heavy quark species), the topological defects are stable and dominate the energy density of the Universe ruling out such scenario unless one fine-tunes a bias term that explicitly breaks the shift symmetry, and in this case $f_{\rm PQ}$ is constrained to be less than $\mathcal{O}(10^{10})\,{\rm GeV}$ in order to avoid the overproduction of axions~\cite{Kawasaki:2014sqa,Hiramatsu:2012sc}. }. The non-thermal restoration of the PQ symmetry can be avoided if the amplitude of the PQ field at the beginning of the oscillation phase is less than $\lesssim 10^4 f_{\rm PQ}$~\cite{Kawasaki:2013iha} or, if there is a coupling between the PQ field and the total energy density of the inflaton, and the oscillation of the PQ field is driven by a higher order term in the potential~\cite{Harigaya:2015hha,Kearney:2016vqw}. Note that similar constraints would follow, if we had considered a moduli field instead of an inflaton field. Once the amplitude of the PQ field drops below $f_{\rm PQ}$, the oscillation of the field continues around its minimum at $f_{\rm PQ}$. In such a case, there will be no non-perturbative production of QCD axion during the second phase of the oscillation~\cite{Mazumdar:2015pta}, but it can still lead to dangerous consequences from a perturbative decay of the PQ field. In a wide range of parameter space, the coherent oscillation of the PQ field can decay dominantly into ultra-relativistic axions. If this decay occurs at sufficiently late times, the resultant axions will not thermalise with the plasma keeping their initial abundance and momenta. Such hot axions will act as dark radiation, increasing the effective number of relativistic degrees of freedom (dof), i.e. $N_{\rm eff}$. The value of $N_{\rm eff}$ is constrained by the observation of CMB~\cite{Ade:2015xua}, allowing us to put constraints on the PQ parameter space. The constraints on the extra relativistic species induced by heavy decaying particles are extensively discussed in the literature. These include the discussion in the context of the supersymmetric axion models~\cite{Kawasaki:2007mk}, and in the context of the heavy moduli decay in string cosmology~\cite{Cicoli:2012aq,Cicoli:2010ha}, where reheating SM degrees of freedom remains a challenge. Instead of considering these scenarios, where the mass of decaying particles is generically controlled by the supersymmetry breaking effects, here we focus on the standard non-supersymmetric PQ mechanism, in which the decaying particle is identified as the radial component of the PQ field. The mass of the radial field is determined by the self coupling constant and the PQ symmetry breaking scale. We show that an upper bound on the reheating temperature can be placed, which is relevant to axion DM with a large decay constant. This paper is structured as follows. In section~\ref{sec:setup}, we review the dynamics of the PQ symmetry breaking followed by a brief review of axion thermalisation and thermal production in section~\ref{sec:th-prod}. In section~\ref{sec:nonth-prod}, we discuss the non-thermal production of ultra-relativistic axions from the coherent oscillation of the radial component of the PQ field. We discuss the different constraints on the axion parameter space in section~\ref{sec:constraints}. Finally, we conclude our discussion in section~\ref{sec:conclusion}.

\label{sec:conclusion} The PQ mechanism presents an elegant solution to the strong CP problem, and the angular field, the axion, can be a good DM candidate due to its largely suppressed coupling to all SM particles. However, axions being very light can also act as dark radiation if they are produced with large momenta at sufficiently late times. We showed that this can happen if the radial part of the PQ field was displaced from $f_{\rm PQ}$ due to an initial condition or a direct coupling to the inflaton/moduli field. The perturbative decay which we have discussed here happens when the PQ field oscillates coherently, during which most of the co-moving energy density stored in these coherent oscillations gets transferred into light axions. The energy density of axions which contribute to the radiation energy density is constrained from number of observations listed above. The bound is mostly relevant to larger values of the axion decay constant. Axion DM with a large decay constant is expected to be probed by future experimental studies such as CASPEr~\cite{Budker:2013hfa}. Since it is impossible to realize such a large PQ scale in the post-inflationary PQ symmetry breaking scenario~\cite{Kawasaki:1995vt}, one should seriously consider the pre-inflationary PQ symmetry breaking scenario if axion DM were to be found in such experiments. We have seen that the cosmological evolution of the PQ field is quite non-trivial in such a scenario, and the overproduction of ultra-relativistic axions leads to an upper bound on the reheating temperature, which further constrains the thermal history of the Universe. Furthermore, one can also expect many $\phi$ fields to oscillate simultaneously, either arising from inflation~\cite{Liddle:1998jc}, or due to multi-moduli fields, whose effects can be discussed by following similar arguments developed in this paper.

1607.06958

1607

1607.05288_arXiv.txt

{We present a systematic study of the intensity mapping (IM) technique using updated models for the different emission lines from galaxies. We identify which ones are more promising for cosmological studies of the post reionization epoch. We consider the emission of $\lya$, $\ha$, H$\beta$, optical and infrared oxygen lines, nitrogen lines, CII and the CO rotational lines. We show that $\lya$, $\ha$, OII, CII and the lowest rotational CO lines are the best candidates to be used as IM probes. These lines form a complementary set of probes of the galaxies emission spectra. We then use reasonable experimental setups from current, planned or proposed experiments to assess the detectability of the power spectrum of each emission line. Intensity mapping of $\lya$ emission from $z=2$ to 3 will be possible in the near future with HETDEX, while far-infrared lines require new dedicated experiments. We also show that the proposed SPHEREx satellite can use OII and $\ha$ IM to study the large-scale distribution of matter in intermediate redshifts of 1 to 4. We find that submilimeter experiments with bolometers can have similar performances at intermediate redshifts using CII and CO(3-2).}

Measurements of the 3-dimensional (3D) large-scale structure of the Universe across cosmic time promise to bring exquisite constraints on cosmology, from the nature of dark energy or the mass of neutrinos, to primordial non-Gaussianity and tests of General Relativity (GR) on large scales. Most of these surveys are based on imaging a large number of galaxies at optical or near-infrared wavelengths, combined with redshift information to provide 3D positions of the galaxies [BOSS (SDSS-III) \citep{2009astro2010S.314S}, DES \citep{Flaugher:2004vg}, eBOSS \citep{Dawson:2015wdb}, DESI \citep{Levi:2013gra}, 4MOST \citep{deJong:2012nj}, LSST \citep{Abate:2012za,Bacon:2015dqe}, WFIRST \citep{Spergel:2015sza}, and the Euclid satellite \citep{Laureijs:2011gra,Amendola:2012ys}]. These observations at optical and infrared wavelengths will be limited to galaxy samples between $z=0.3$ to 2, and in some cases, redshifts will be determined only by the photometric data. Instead of counting galaxies, the intensity mapping (IM) technique uses the total observed intensity from any given pixel. For a reasonably large 3D pixel (with a given angular and frequency/redshift resolution), also referred to as \emph{voxel}, we expect it to contain several galaxies. The intensity in each pixel will thus be the integrated emission from all these galaxies. This should then provide a higher signal-to-noise compared to standard galaxy ``threshold" surveys. Moreover, since most cosmological applications rely on probing large scales, the use of these large pixels will not affect the cosmological constraints. By not needing galaxy detections, the requirements on the telescope/survey will be much less demanding. However, since we are no longer relying on ``clean" galaxy counts, we need to be much more careful with other contaminants of the observed intensity. In order to have redshift information, the measured intensity should originate from specific emission lines. The underlying idea is that the amplitude of this intensity will be related to the number of galaxies in the 3D pixel emitting the target line. The fluctuations in the intensity across the map should then be proportional to the underlying dark matter fluctuations. Several lines can in principle be used for such surveys. In particular, a significant focus has been given in recent years to the HI 21cm line \citep{Battye:2004re, 2013MNRAS.434L..46S, Chang:2007xk, 2008PhRvL.100p1301L,% Bagla:2009jy,% 2012A&A...540A.129A} and how well it can perform cosmological measurements \citep[see e.g.][]{Camera:2013kpa, Bull:2014rha}. The 21cm signal is also being used to study the Epoch of Reionization at higher redshifts, with several experiments running or in development (see e.g. \citealt{2013ExA....36..235M} for a review). As this line will be observed at radio frequencies, telescopes will naturally have lower resolutions, making this line an obvious application for intensity mapping. These telescopes also usually have large fields of view which allows them to cover large areas of the sky more quickly. Moreover, given the low frequencies, the HI line has very little background and foreground contamination from other lines, which is an advantage when comparing to intensity mapping with other lines as we will discuss later. Still, all lines have significant Galactic foregrounds that need to be removed or avoided. Several experiments have been proposed so far \citep{chime, 2013MNRAS.434.1239B, 2012IJMPS..12..256C, Bigot-Sazy:2015tot,Newburgh:2016mwi}, and some observations have already been made \citep{2011AN....332..637K}, including a detection in cross-correlation \citep{Chang:2010jp}. Moreover, a large HI intensity mapping survey has been proposed for SKA1-MID \citep{Santos:2015bsa}. Although intensity mapping with the HI line holds great promises for cosmology, it is still relevant to ask if we can use other lines for such purposes. An obvious reason is that the 21cm line is quite weak. Although galaxy surveys with other lines have become routine, with detections made up to very high redshifts, the highest redshift detection for HI is $z=0.376$ \citep{2016ApJ...824L...1F}. Some studies have already been done for other lines, in particular for the Epoch of Reionization, such as Ly$\a$ \citep{Silva:2012mtb}, CII \citep{Gong:2011mf}, CO \citep{Gong:2011ts}, H$_2$ \citep{Gong:2012iz} and others \citep{Visbal:2010rz}. At lower redshifts, \citet{Pullen:2013dir} have considered Ly$\a$ to probe the underlying dark matter power spectrum in a fiducial experiment, although they assume a very strong signal coming from the intergalactic medium (IGM). In fact, $\lya$ IM has already been detected in cross-correlation with quasars \citep{Croft:2015nna}. \citet{Uzgil:2014pga} considered the 3D power spectrum of CII and other far-infrared (FIR) emission lines. Similarly, \citet{Pullen:2012su} used the large scale structure distribution of matter to study CO emission while \citet{Breysse:2014uia} studied the possibility of using the first rotational transition of CO to do a survey at $z\sim3$. This paper presents a systematic study of all lines (besides HI) that can in principle be used for intensity mapping with reasonable experimental setups. It also compares them to determine which are the optimal lines to target for intensity mapping. It uses the latest observational and simulation results to predict the strength of the intensity and the expected bias towards dark matter. It then discusses the feasibility of surveys with these lines and which ones are more appropriate for cosmological applications. This is particularly relevant as the combination of different lines can bring exquisite constraints on large scale effects, such as primordial non-Gaussinanity, using the multi-tracer technique \citep{Seljak:2008xr, 2015PhRvD..92f3525A, 2015ApJ...812L..22F}. Moreover, the cross-correlation can be particularly useful in dealing with foregrounds/backgrounds as well as systematics. Finally, even if such intensity mapping surveys are not extremely competitive for cosmology, the simple detection of such a signal and its power spectrum will bring invaluable information about the astrophysical processes involved in the production of such lines and the clustering of the corresponding galaxies. We start with a review of the basic calculations for the average observed intensity of an emission line in Section \ref{sec:reviewIM}. In Section \ref{sec:lineemission} we model line emission in terms of the star formation rate of a given halo and estimate the signal for different lines. We then compare lines in Section \ref{sec:line_comp} so that we have an indication of which surveys to target in Section \ref{sec:Surveys}. The two following sections are focused on briefly assessing other sources of emission/contamination and address how to deal with such issues. We conclude in section \ref{sec:discussion}.

\label{sec:discussion} IM studies open new windows to probe not only the distribution of matter in the universe, but also to study global properties of galaxies. Particularly IM is useful to probe the BAO scale independently from galaxy surveys. The evolution of the halo mass function is also an interesting cosmological goal although it will be difficult to disentangle it from the evolution of astrophysical properties of galaxies. On the astrophysical side IM of metal lines can be used to probe the poorly known chemical evolution of the gas in the universe. In table \ref{tab:expsum} we summarize the experimental and sample details for each line we consider for IM. In the last two columns we present the signal-to-noise ratio for each survey for a fixed $k$-range. Although higher SNRs can be attained going to smaller scales, we caped at $k=0.3$ Mpc$^{-1}$ so we do not need to worry about issues with non-linear scales. Also, below these scales the voxel of IM may be too small for the assumption that astrophysical fluctuations between galaxies average out. From the last column of table \ref{tab:expsum} we see that using HETDEX as a $\lya$ IM survey would be the best performing survey out of the ones considered in this paper. Although $\ha$ IM with SPHEREx is possible, its SNR is not as good as for $\lya$ or the FIR lines with a TIME-like experiment. The proposed OII survey will have the worst performance of the IM experiments considered, given the smaller physical volume to be observed. This is due to the relatively low redshifts of the observational frequency window where OII is the dominant emission line. CII with a TIME-like experiment also has a high SNR but it is still lower than CO(3-2). This is mainly due to the fact that for a fixed survey time a CO(3-2) IM survey can sample a bigger area of the sky. In fact, a CO(3-2) IM survey can do as well as $\lya$ with HETDEX. Since TIME-like experiments are in the realm of possibilities we also performed the calculations for a pilot experiment. For both CII and CO(3-2) we computed the SNR for a smaller 50h survey covering 2$\deg^2$ of the sky. These can be used as proof of concept despite the lower SNR. Looking at table \ref{tab:expsum} one can say that $\lya$ and CO(3-2) would be the best lines to perform IM. This is not necessarily true due to contamination issues. As said before, different lines are sensitive to different mass ranges of the HMF, to the chemical evolution of the Universe and to the underlying astrophysical processes within galaxies. Hence, we would like to stress that all lines should be seen as complementary, although $\lya$ is the best starting point. This work is intended to probe the new windows opened by line IM for studying our cosmology. On top of looking at the distribution of matter in the Universe, one can study the astrophysical properties of galaxies. We started by modelling the luminosity of a line in terms of halo mass. We proposed a general prescription based on previous theoretical and observational studies. Although one expects scatter in these relations on small scales it should average out in big enough voxels. Under our models we find figures \ref{fig:nuI:comp_opt}, \ref{fig:nuI:comp_fir} and \ref{fig:nuI:comp_cos} which indicate which are the best candidates for line IM. Unsurprisingly, these are $\lya$, OII, $\ha$, CII and the lowest rotation lines of CO. We then follow by estimating power spectrum error bars for each line with reasonable experimental settings for a fixed redshift range. We find that one can measure the power spectrum of these lines, assuming a cleaned signal. % Unlike 21cm emission of HI, UV, optical and infrared lines have stronger line contaminants that need to be considered. We discussed foreground contamination and ways to extract the target signal from observational maps in Section \ref{sec:Foregrounds} and found that there are promising ways to successfully clean many of these maps. The literature already has a wide discussion on such methods either for the EoR \citep{Visbal:2010rz,Lidz:2011dx,Silva:2012mtb,Gong:2013xda,Silva:2014ira,Breysse:2015baa} or for the late universe \citep{Croft:2015nna,Lidz:2016lub} which we will explore in the future. This work considers the 3D power spectrum P(k), as it is conventional in galaxy surveys. Nonetheless, one should point out that IM allows for tomographic studies using the angular power spectra $C_\ell$. There are several advantages on doing so. First of all P(k) is gauge dependent and suffers from projection effects, which does not happen when using $C_\ell$. Furthermore, the contributions from lensing and the so called GR corrections to number count fluctuations \citep{Yoo:2010ni,Challinor:2011bk,Bonvin:2011bg} and temperature/intensity fluctuations \citep{Hall:2012wd} are easily included in the angular power but not in the 3D P(k). We note that in order to successfully do science with line intensity maps it will be necessary to perform further detailed studies. Galaxy emission studies will have to be done for foreground/interloping lines. Also, continuum contamination estimates should be made accounting for dust absorption and/or propagation in the IGM. This is needed to properly access the feasibility of all contamination cleaning techniques. As prescribed in section \ref{sec:lineemission} one should use cosmological simulations of the 3D distribution of halos and attribute line luminosities and galactic continuum emission to each on top of a continuous extragalactic continuum. Then, by redshifting all the emission to the observers' frame, one should assess how well one can disentangle the signal, accounting for masking rates of interloping lines, the frequency smoothness of the continuum contamination and cross-correlations. % These are general guidelines since the ideal method depends on the target line. We conclude by emphasizing that IM, with lines other than HI, shows great potential to measure the large-scale distribution of matter in the Universe. This is especially true since independent sets of lines can measure the BAO scale around $z=2$, a further test of the expansion of the universe at intermediate redshifts. One should also note that these lines acquire an important role in the context of the multi-tracer technique \citep{Seljak:2008xr,McDonald:2008sh} to beat cosmic variance. As \citet{2015ApJ...812L..22F} showed, the improvement in measuring large scale effects is greatly increased as the ratio of the two biases deviates from $\sim1$. Hence, these lines can give a better bias ratio for a particular paring of DM tracers. \subsection*

1607.05288

1607

1607.07628_arXiv.txt

We construct foreground simulations comprising spatially correlated extragalactic and diffuse Galactic emission components and calculate the `intrinsic' (instrument-free) two-dimensional spatial power spectrum and the cylindrically and spherically averaged three-dimensional $k$-space power spectra of the Epoch of Reionization (EoR) and our foreground simulations using a Bayesian power spectral estimation framework. This leads us to identify a model dependent region of optimal signal estimation for our foreground and EoR models, within which the spatial power in the EoR signal relative to foregrounds is maximised. We identify a target field dependent region, in $k$-space, of intrinsic foreground power spectral contamination at low $k_{\perp}$ and $k_{\parallel}$ and a transition to a relatively foreground-free intrinsic EoR window in the complement to this region. The contaminated region of $k$-space demonstrates that simultaneous estimation of the EoR and foregrounds is important for obtaining statistically robust estimates of the EoR power spectrum; biased results will be obtained from methodologies that ignore their covariance. Using simulated observations with frequency dependent $uv$-coverage and primary beam, with the former derived for HERA in 37-antenna and 331-antenna configuration, we recover instrumental power spectra consistent with their intrinsic counterparts. We discuss the implications of these results for optimal strategies for unbiased estimation of the EoR power spectrum.

\label{Sec:Introduction} The redshifted 21-cm hyperfine line emission from the neutral hydrogen that pervades the intergalactic medium (IGM) at high redshifts ($z \gtrsim 6$) provides a unique probe of the Cosmic Dawn and the Epoch of Reionization (EoR) when the first stars, galaxies and quasars formed, effecting a transition of the IGM from a neutral to ionized state. The redshift dependent hydrogen spin-temperature distribution during reionization (see e.g. \citealt{1958PIRE...46..240F, 1959ApJ...129..525F, 1990MNRAS.247..510S}) encodes a wealth of information which, if extracted, can be used to constrain cosmological parameters (e.g. \citealt{2006ApJ...653..815M, 2008PhRvD..78b3529M, 2009MNRAS.394.1667F, 2014ApJ...782...66P, 2015aska.confE..12P}), probe directly the initial stages of structure formation and deduce the nature of the first ionizing sources (e.g. \citealt{2012MNRAS.424..762D, 2013MNRAS.431..621M, 2014MNRAS.439.3262M, 2015MNRAS.449.4246G}). Over the last few years, the first generation of low-frequency interferometers designed to detect the highly redshifted 21-cm signal from the EoR have begun taking measurements. These include: the Giant Metrewave Radio Telescope (GMRT; \citealt{2013MNRAS.433..639P})\footnote{http://www.gmrt.ncra.tifr.res.in}, the LOw Frequency ARray (LOFAR; \citealt{2013A&A...556A...2V})\footnote{http://www.lofar.org/}, the Murchison Widefield Array (MWA; \citealt{2013PASA...30....7T})\footnote{http://www.mwatelescope.org/} and the Donald C. Backer Precision Array for Probing the Epoch of Reionization (PAPER; \citealt{2010AJ....139.1468P})\footnote{http://eor.berkeley.edu/}. Over the next few years, the second generation of instruments with larger collecting areas will begin construction and operation. These include the Hydrogen Epoch of Reionization Array (HERA; \citealt{2016arXiv160607473D})\footnote{http://reionization.org/} and the Square Kilometre Array (SKA; \citealt{2013ExA....36..235M})\footnote{https://www.skatelescope.org/}. Statistical information about the expansion of ionized bubbles and the astrophysical sources that produced them is encoded in the 21-cm temperature power spectrum of the EoR. For isotropic emission, the spatial power spectrum can be calculated through two-dimensional spatial averaging over circular annuli as a function of spatial scale. For the redshifted 21-cm line emission, observation frequency maps to line of sight distance. Thus, for a sufficiently slowly evolving 21-cm signal averaging of the power spectrum can be extended to three dimensions for narrow frequency bands within which the evolution of the signal is minimal \citep{2012MNRAS.424.1877D, 2014MNRAS.442.1491D}. Performing three-dimensional averaging of this form using spherical shells or cylindrical annuli defines the spherical and cylindrical power spectrum respectively. As a result of the significant gains in signal to noise, facilitated by averaging of the data, detection of the power spectrum of the EoR signal and disentangling the cosmological and astrophysical information that it encodes is the primary goal of current EoR experiments (see e.g. \citealt{2015MNRAS.449.4246G}). Theoretical models predict the 21-cm temperature signal from the EoR to be of the order of $10~\mr{mK}$ (e.g. \citealt{2011MNRAS.411..955M}). The largest challenge faced by both current and next generation instruments is the detection of this signal in the presence of astrophysical foreground contaminants that are up to 5 orders of magnitude brighter in intensity \citep{1999A&A...345..380S}. In order to detect the faint EoR power spectrum and disentangle it from the foregrounds in an unbiased manner, accurate characterisation of each in power spectral space can provide valuable insight. In this paper we develop EoR and foreground simulations and present the results of applying a Bayesian methodology to infer the `intrinsic' (instrument-free) two-dimensional spatial power spectrum, and cylindrically and spherically averaged three-dimensional $k$-space power spectra of these simulations. We also estimate the spherically averaged three-dimensional $k$-space power spectra from the simulated interferometric observation of the foregrounds. While the foregrounds do not possess the frequency-to-redshift mapping valid for the 21-cm emission, their $k$-space power spectra nevertheless describe foreground power at the spatial and spectral scales of interest for measuring the EoR. Their power spectra, therefore, define a fundamental measurement of contamination of the EoR power spectrum by foreground emission. In \autoref{Sec:ForegroundRemoval} we summarize common foreground avoidance / removal strategies used to date. In \autoref{Sec:PowerSpectralModelandAssumptions} we summarize the method of Bayesian power spectral estimation as presented in Lentati et al. 2016 (in prep.) with particular focus on the method of direct sampling from the spherical power spectrum coefficients of the EoR which we make use of in what follows. In Sections \ref{Sec:CosmologicalSignal}, \ref{Sec:GalacticForegroundEmission} and \ref{Sec:ExtragalacticForegroundEmission} we develop our EoR, Galactic and extragalactic emission simulations respectively. In \autoref{Sec:SimulatedHERAObservations} we describe our instrumental simulation modelled on HERA in 37 and 331-antenna configurations. We use a frequency dependent Gaussian approximation to the HERA primary beam and include realistic frequency dependent $uv$-coverage of the simulated observations. In \autoref{Sec:Analysis} we analyse the power spectra of our foreground and 21-cm simulations. Using the spatial power spectrum we calculate a model dependent region of optimal signal estimation within which the ratio of EoR to foreground power is maximised. We estimate the cylindrically and spherically averaged intrinsic power spectra of each of our simulation components and analyse contamination of the EoR power spectrum by foregrounds as a function of position in $k$-space. Finally, we derive the corresponding instrumental power spectra and show their detected coefficients are fully consistent with the intrinsic power spectra. We discuss the implications of these results for optimal strategies for unbiased estimation of the EoR power spectrum and in the context of current experimental approaches to EoR power spectral estimation. We offer some concluding remarks and discuss future work in \autoref{Sec:Conclusions}.

\label{Sec:Conclusions} We have argued that independent estimation of the level of foreground and EoR power as a function of position in $k$-space provides a means of assessing EoR power spectrum contamination. We have developed foreground simulations comprising spatially correlated extragalactic and diffuse Galactic emission components to investigate this claim. In constructing the foreground simulations we include: accurate modelling of large scale Galactic structure (relative to the simulation region), Galactic diffuse synchrotron emission (GDSE) simulations for a range of levels of brightness temperature -- spectral index correlation, and a physically motivated spectral absorption model for compact extragalactic sources. We have used the Bayesian power spectral estimation framework detailed in Lentati et al. 2016 (in prep.) to perform unbiased estimation of the `intrinsic' (instrument-free) two-dimensional spatial power spectrum, $P_{uv}$, the three-dimensional spherically averaged power spectrum, $P(k)$, and the maximum likelihood cylindrically averaged power spectrum, $P(k_{\perp}, k_{\parallel})$, of each of the foregrounds and of the EoR. Unbiased estimation is achieved by performing a joint estimate for the power spectrum on scales fulfilling the Nyquist criterion, given the data sampling rate, with models for both large and small scale power. Our large spatial and spectral scale power models are given by a sub-harmonic grid sampled on a set of 10 log-uniformly spaced spatial scales between the size of the image and 10 times the size of the image and a set of quadratics respectively. Power from emission components at higher frequencies than the Nyquist sampling rate of the data manifests itself as an additional source of noise that we model in the data covariance matrix. We additionally assume that the redshifted 21-cm signal is spatially isotropic and, over the redshift interval under consideration, homogeneous (assuming the power spectrum of the 21-cm signal is approximately constant) and uncorrelated between spatial scales. As such, we impose a spherical prior on the three-dimensional $k$-space power spectrum. Comparison of $P_{uv}$ calculated for our EoR and foreground simulations has lead us to identify a model dependent region of optimal signal estimation within which the ratio of the spatial power in the EoR signal to foregrounds is maximised. This assumes no uncertainty on the flux-density distribution resulting from sources with flux-densities in excess of five times the confusion noise limit for HERA in 37-antenna configuration. The region of optimal signal estimation extends over the spatial resolution range $35~\lambda \le \abs{\mathbfit{u}} \le 55~\lambda$ at $126~\mr{MHz}$, with lower and upper bounds corresponding to the optimal spatial resolutions for estimation of the EoR signal in a cold region out of the plane of the Galaxy and a region of intense GDSE emission in the Galactic plane respectively. We discuss the application of this procedure to observational data via calculation of the spatial power spectrum of the foregrounds in a specific target field. We find that for a fixed EoR neutral fraction, it is possible to calculate well defined limits on the region of optimal signal estimation, dominated by the structure of the spatial power spectrum. We consider the $\zeta$, $R_\mr{mfp}$, $T^\mr{Feed}_\mr{vir}$ parametrisation of the EoR signal (see e.g. \citealt{2015MNRAS.449.4246G}), with $\zeta$ the galactic ionizing efficiency, $R_{\rm mfp}$ the mean free path of ionizing photons in the IGM and $T^{\rm Feed}_{\rm vir}$ the minimum virial temperature of star-forming haloes, and find that the limits on the region of optimal signal estimation can be defined to within uncertainties of $15\%$, for our foreground simulations and EoR power spectra across an observationally motivated astrophysical parameter range and with neutral fractions, $x_{H}$, in the range $0.45$ to $0.55$. For a spatially separable three-dimensional foreground power spectrum (which is a good approximation across the $\sim 8~\mr{MHz}$ bandwidth considered here), this provides a method for estimating the optimal baseline lengths (those on which intrinsic contamination of the EoR power spectrum by foreground emission is minimal), on which to perform our three-dimensional $k$-space Bayesian power spectral estimation of the EoR signal, subject to a choice of EoR neutral fraction, target field and instrument model. When performing this calculation, and in general when estimating the power spectrum of the EoR from observational data, ideally, one would include prior information on the resolved extragalactic point source population in the power spectral likelihood. The amplitude of the signal would then be estimated given this prior. This allows the available information on sources to be accounted for when estimating the power spectrum in a statistically robust manner. In future work we will investigate the level to which including realistic uncertainties on resolved source parameters, as well as including priors on unresolved sources (for example, from external source catalogues constructed using high resolution instruments), impacts the power spectral constraints that can be obtained. Extending our analysis to the three-dimensional power spectrum, we have calculated the intrinsic spherically averaged $k$-space power spectrum of our EoR simulation and of each foreground simulations in the spatial resolution range $1.6\times10^{-3}~h\mr{Mpc^{-1}} < k_{\perp} < 3.3\times10^{-2}~h\mr{Mpc^{-1}}$, broadly matching the region of maximum spatial sensitivity of current generation 21-cm experiments. For our GDSE simulation we find that power on the scales of interest for EoR power spectral estimation varies as a function of position in the Galaxy and with the level of brightness temperature -- spectral index correlation of the emission. In all cases, we find GDSE is the dominant foreground component and that contaminating power in our GDSE simulations exceeds the power in the EoR emission by up to several orders of magnitude on large spatial scales (for $k \lesssim 0.15~h\mr{Mpc^{-1}}$ in cold regions, and $k \lesssim 0.3~h\mr{Mpc^{-1}}$ for emission in the Galactic plane). The fact that the intrinsic power spectra of the foregrounds are significant across a range of spatial scales of interest for estimating the EoR power spectrum (specifically, at low-$k$) demonstrates that without a priori knowledge of the complexity and covariance between the foregrounds and the EoR signal, foreground subtraction prior to power spectral estimation is liable to bias the power spectral estimates, as well as resulting in incorrect uncertainties. This highlights the importance of joint estimation of a model for the foregrounds and the EoR power spectrum if unbiased estimates of the EoR power spectrum are to be obtained. We have further investigated the distribution of contamination of the EoR power spectrum by foregrounds as a function of position in $k$-space by calculating the maximum likelihood intrinsic cylindrically averaged power spectra corresponding to the intrinsic spherically averaged power spectra of a reference region. For this region we select a relatively cold patch of sky, lying out of the plane of the Galaxy (centred on $\mr{RA}=0\fdg0,\ \mr{Dec}=-30\fdg0$), as being typical of preferred regions for estimation of the EoR signal from observational data. We find that foreground power is dominated by GDSE across ($k_{\perp}$, $k_{\parallel}$)-space in the spatial resolution range probed. In each of the foreground components we observe a rapid drop-off in power as a function of $k_{\parallel}$ which reflects the relative smoothness of their spectra. In the free--free and GDSE cylindrical power spectra, steep power law structure is apparent spatially. This structure is well described by the spatial power spectra of those foregrounds. We additionally calculate the fractional contamination of the total recovered power by foreground emission as a function of $k_{\perp}$ and $k_{\parallel}$. We find that the power law drop-off of the dominant foreground in this regime, GDSE, as a function of both $k_{\perp}$ and $k_{\parallel}$ produces significant contamination of the total power in ($k_{\perp}$, $k_{\parallel}$)-space evident at low $k$. Outside of this region, we find a relatively foreground-free intrinsic EoR window in which the 21-cm signal dominates the intrinsic power spectrum. We propose a method for estimating the intrinsic foreground contamination as a function of position in $k$-space and its complement, the intrinsic EoR window, from observational data in a target field. By estimating the intrinsic power spectrum in observations at sufficiently high frequencies for the reionization signal to be negligible ($\nu \gtrsim 200~\mr{MHz}$), the intrinsic foreground power spectrum can be independently estimated. In the simplest case, the mean spectral index of the foreground components can be used to extrapolate from the measured high frequency intrinsic foreground contamination to the frequencies of interest for estimating the redshifted 21-cm power spectrum. We note that, by supplementing the spherically symmetric prior on the power spectrum with physically motivated priors on the power spectra of the foregrounds, better separation of the foreground and EoR $k$-space power spectral estimates will be possible. With appropriate extrapolation, the measured foreground power spectral contamination obtained at high frequencies can provide such a foreground prior at the frequencies of interest. Extrapolation of the high frequency intrinsic foreground power spectrum to the frequencies of interest in combination with numerical modelling of the EoR power spectrum can also enable the estimation of an intrinsic EoR window within which the ratio of EoR to foreground power in $k$-space is maximal. A balance between sensitivity and predicted power spectral bias can then allow a preferred region of $k$-space within which to estimate the EoR power spectrum to be determined. Finally, we estimate the power spectrum from simulated interferometric observations incorporating frequency dependent $uv$-coverage and primary beam. We consider two test cases: firstly, in the low signal-to-noise regime, using $uv$-sampling relevant to HERA in 37-antenna configuration and secondly, in the high signal-to-noise regime, using the $uv$-sampling of HERA in 331-antenna configuration for a restricted set of baselines with lengths in the HERA 37 range (which enables us to take advantage of the increased sensitivity of the larger instrument in this resolution range without additional computational cost). In both cases power spectral detections are fully consistent with power predicted by the intrinsic power spectra of the emission components calculated with a filled $uv$-plane and in the absence of instrumental effects.

1607.07628

1607

1607.00839_arXiv.txt

We present the first three-dimensional gas-dynamical simulations of the grazing envelope evolution (GEE) \textred{{{{{{of stars}}}}}}, with the goal of exploring the basic flow properties and the role of jets at the onset of the GEE. In the simulated runs, a secondary main-sequence star grazes the envelope of the primary asymptotic giant branch (AGB) star. The orbit is circular at the radius of the AGB primary star on its equator. We inject two opposite jets perpendicular to the equatorial plane from the location of the secondary star, and follow the evolution for several orbital periods. We explore the flow pattern by which the jets eject the outskirts of the AGB envelope. After one orbit the jets start to interact with gas ejected in previous orbits and inflate hot low-density bubbles.

\label{sec:intro} Numerical hydrodynamical simulations of the common envelope evolution (CEE) have been performed for thirty years, e.g., the 2D simulations conducted by \cite{BodenheimerTaam1984}, followed by the 3D simulations conducted by \cite{LivioSoker1988} and then by \cite{RasioLivio1996}. More sophisticated simulations have followed (e.g., \citealt{SandquistTaam1998, Lombardi2006, DeMarcoetal2011, Passy2011, Passyetal2012, RickerTaam2012, Ohlmannetal2016, Iaconietal2016, Staffetal2016, IvanovaNandez2016, Kuruwitaetal2016, NandezIvanova2016}). Two main goals are behind many of these simulations. The first goal is to determine the final orbital separation between the core of the giant star and the more compact companion (secondary) star. The second goal is to understand the manner by which the envelope is removed. Although in many cases less than the energy released by the in-spiraling binary system is required to unbind the CE (e.g., \citealt{DeMarcoetal2011, NordhausSpiegel2013}), the removal of the CE is not free from problems (e.g., \citealt{DeMarcoetal2011, Passy2011, Passyetal2012, RickerTaam2012, Soker2013, Ohlmannetal2016}). To overcome the envelope-removal obstacle, extra energy sources have been proposed, such as recombination energy (e.g., \citealt{Nandezetal2015} for a recent paper), and jets launched by the secondary star (e.g., \citealt{Soker2004}). Also, \cite{Soker2004} suggested that in some cases, both for stellar and sub-stellar companions, the energy source might be the giant luminosity itself (nuclear burning in the core), under the assumption that fast rotating envelopes efficiently form dust and have high mass-loss rates. We adopt the view that in many cases jets launched by the companion can facilitate CE removal (e.g., \citealt{Soker2014}). This holds for main-sequence (MS) stars as well \citep{SchreierSoker2016}, as they can accrete mass at a high rate \citep{Shiberetal2015}. \cite{ArmitageLivio2000} and \cite{Chevalier2012} studied CE ejection by jets launched from a neutron star companion, but they did not consider jets to be a general CE ejection process. When the jets become efficient in envelope removal, such that they remove the entire envelope outside the orbit of the companion, the system does not enter a CE phase. Instead, the system experiences the grazing envelope evolution (GEE; \citealt{SabachSoker2015, Soker2015, Soker2016a, Soker2016b}). Our goal is to explore some basic properties of the flow at the start of the GEE. The numerical set up is described in section \ref{sec:numerics}, and the results are presented in section \ref{sec:results}. We summarize in section \ref{sec:summary}.

\end{table} \begin{figure*} \centering \includegraphics[trim= 1cm 0.4cm 0.2cm 0.3cm,clip=true,width=0.9\textwidth]{4panel.eps} % \caption{The hydrodynamic properties of Run 1 ($\dot{M}_{\rm{jet}}=10^{-3} \msyr$; $P=198.4 \; \rm{days} $) after four orbital periods. The companion is at $(x,y,z)=(0,1~ \rm{AU},0)$ on every panel. The plane $x-y$ is taken to be the equatorial plane. The upper-left panel shows the density and velocity in the $x-z$ plane at $y=1~ \rm{AU}$, namely, goes through the secondary location in perpendicular to the orbital plane. The lower-left panel shows density and velocity in the orbital plane. The lower-right panel shows the density and velocity in the $z-y$ plane, a plane cutting the centers of the two stars. The upper-right panel shows the temperature in the same plane. } \label{fig:4panel} \end{figure*} With \textit{MASS\_SCALAR} option in \texttt{FLASH}, we use two `tracers' to follow the material originated from the AGB envelope and the gas injected by the jets. The tracers indicate the fraction of mass originate from the AGB star and from the jets in each cell. Fig. \ref{fig:tracers} presents a 3D map of contour surfaces with $99.5$ per cent of AGB gas (red) and of $50$ per cent of jet gas (blue), at $t = 119~{\rm days}$ for Run 1. The morphology of the jets can clearly be seen in the figure. The jets are injected from the secondary star in two opposite directions perpendicular to the orbital plane (they also have an azimuthal velocity component due to the orbital motion). The jets are immediately diverted outwards by the dense AGB gas. \begin{figure} \includegraphics[trim= 0.0cm 0.5cm 0.0cm 1.0cm,clip=true,width=0.95\columnwidth]{stmt.eps} % \caption{3D map of contour surfaces with $99.5$ per cent of AGB gas (red) and of $50$ per cent of jet gas (blue), after one orbital period for Run 1. Axes run from $-7\times 10^{13} \cm$ to $7\times 10^{13} \cm$. The orbital plane is along the line of sight and the secondary star is on the left side of the AGB surface.} \label{fig:tracers} \end{figure} In Fig. \ref{fig:three-dimensional_density} we present a 3D map of density contour surfaces. Arrows represent the flow velocity. This figure clearly emphasizes the complicated flow structure. The figure shows separate color schemes and vector properties for velocity above and below $400~\rm{km~s^{-1}}$. The high velocity gas comes from the jets and the AGB layers it collided with. It can be clearly seen that the jets are deflected towards lower density regions. The slow velocity gas shows both rotation as a result of the secondary orbit, and an outflow as a result of the jets. \begin{figure} \includegraphics[trim= 0.0cm 0.0cm 1.2cm 0.0cm,clip=true,width=0.95\columnwidth]{dens3d.eps} % \caption{ The flow structure after four orbital periods in Run 1. Axes run from $-7\times 10^{13} \cm$ to $7\times 10^{13} \cm$. The AGB is at the center and the secondary is marked by the ``$+$'' sign. The secondary is moving to the left in the figure. The red-green-blue colors show gas with velocities larger than $400~\rm{km~s^{-1}}$, that originated in the jets and got deflected. The gray colors show slower gas and at higher densities, corresponding to gas originated from the AGB envelope. The blue-scale velocity arrows are limited to a minimal velocity of $400~\rm{km~s^{-1}}$ and their size and color are proportional to the module of the velocity. The black arrows are of uniform size and represent velocities smaller than $400~\rm{km~s^{-1}}$. } \label{fig:three-dimensional_density} \end{figure} We measured the mass that left the system through a sphere near the edge of the grid. Fig. \ref{fig:logMout} shows this ejected mass as a function of time for the four simulations. In Runs 3 and 4 the secondary star starts its circular orbit inside the envelope, and larger amount of mass is ejected from the system during the time of the simulations. Higher values of $\dot{M}_{\rm jet}$ also yielded higher value of ejected mass. Both these results are according to expectations. In late times of $t \ga 300$~days, the time dependence is close to being exponential with similar rate for all four simulations. Most of the mass injected in the jets leaves the grid. The ejected mass is up to $0.1 \rmModot$ at $t=800$~days, which is $\approx10 \textrm{--} 50$ times the mass injected in the jets. In Run 1 we find that $94$ per cent of the mass that left the grid has a positive energy, namely it is unbound. In Run 3 we find that $88$ per cent is left unbound. \begin{figure} \includegraphics[trim= 0.0cm 0.6cm 0.0cm 0.4cm,clip=true,width=0.99\columnwidth]{ejmass.eps} \caption{The ejected mass as function of time in our four simulations. The ejected mass is the mass that crosses out through a spherical shell of radius $4 \AU$. } \label{fig:logMout} \end{figure} In Fig. \ref{fig:angle} we present the mass that left the grid over the first 8 orbital periods per unit solid angle as a function of the angle from the pole ($\theta=0^\circ$ is the polar direction and $\theta =90^\circ$ is the equatorial plane). The graph represents the two sides of the equatorial plane as they are symmetric. We also show the value of {{{{$\sqrt{<v^2>} \equiv \sqrt{ 2 dE_k(\theta) / dM(\theta)}$}}}}, where $d E_k(\theta)$ and $d M(\theta)$ are the kinetic energy and mass, respectively, that left the grid through a circular surface from $\theta$ to $\theta + d \theta$. The quantities were calculated from $t=0$ to $t=4.4 \yr$, about 8 orbital periods. \begin{figure} \includegraphics[trim= 0.0cm 0.6cm 0.0cm 0.4cm,clip=true,width=0.99\columnwidth]{mv_theta.eps} \caption{The mass lost from the grid per unit solid angle (green line) and $\sqrt{<v^2>}$ (dashed blue line) as a function of the polar angle $\theta$, for the first $4.4 \yr$ of the simulation of Run 1.} \label{fig:angle} \end{figure} From Fig. \ref{fig:angle} we learn the following. (1) There is a concentration of outflow in the equatorial plane (see also Fig. \ref{fig:three-dimensional_density}). This equatorial outflow is achieved without any gravitational field of the secondary star and without any rotation of the primary AGB star. The equatorial outflow is a result of the action of the jets. (2) Along $\simeq 35^\circ$--$45^\circ$ from the polar directions there is a fast relatively massive outflow. This is mainly due to the material that originated in the jets and is deflected by the AGB envelope (see also Figs. \ref{fig:tracers} and \ref{fig:three-dimensional_density}). The outflow geometry and its role in shaping the descendent nebula is a subject of a forthcoming paper. We conducted preliminary three-dimensional hydrodynamical simulations of the onset of the GEE. The GEE is based on jets launched from the secondary MS star that efficiently remove mass from the envelope of the primary AGB star, to the extent it prevents the formation of a common envelope. In the GEE the secondary star grazes the surface of the giant star, as in Run 1 and Run 2 listed in Table \ref{tab:simulationsummary}. For comparison we also simulated two cases where the secondary star starts rotating inside the giant envelope (Run 3 and Run 4). As expected, the flow structure is complicated, e.g., the outflow strongly depends on latitude and longitude, and when integrated over time around the orbital axis there are two latitudes with peaks in the mass outflow rate and variations in the outflow velocity (Figs. \ref{fig:tracers}, \ref{fig:three-dimensional_density}, and \ref{fig:angle}). The two opposite jets inflate two hot bubbles, one at each side of the equatorial plane (temperature panel in Fig. \ref{fig:4panel}). Due to the orbital motion these bubbles stretch in the azimuthal direction and expand outwards. The jets and the inflated bubbles lead to a concentrated equatorial outflow, despite that neither the gravity of the secondary star nor AGB rotation were included in the simulation. We also find that if jets are indeed launched according to the expectation from the mass accretion rate (eqs. \ref{eq:dotmaccrete} and \ref{eq:dotmjet}), they can efficiently remove gas from the AGB envelope as the secondary star moves on a grazing orbit. We find that even if the secondary star starts to launch jets only after it penetrated to the giant envelope, the jets can nonetheless eject a large amount of mass from the region of the envelope outside the secondary orbit and near the equatorial plane. The main result of the present study is that the GEE has a merit. Our results also hint that jets might help in removing the envelope in a full common envelope evolution, i.e., when the secondary is deep inside the envelope rather than grazing the surface. In a future paper we will improve our modelling by introducing a spiraling-in orbit that fully conserves the losses of energy and will follow the secondary for longer times as it dives into the AGB. Jets play an important role in the common envelope phase, especially during the onset of the common envelope. Long lasting imprints of the jets, both in the GEE and in the common envelope evolution, might be a complicated morphology of the descendant nebula, whether a planetary nebula or a nebula around a massive star.

1607.00839

1607

1607.00008_arXiv.txt

We apply two Bayesian hierarchical inference schemes to infer shear power spectra, shear maps and cosmological parameters from the CFHTLenS weak lensing survey --- the first application of this method to data. In the first approach, we sample the joint posterior distribution of the shear maps and power spectra by Gibbs sampling, with minimal model assumptions. In the second approach, we sample the joint posterior of the shear maps and cosmological parameters, providing a new, accurate and principled approach to cosmological parameter inference from cosmic shear data. As a first demonstration on data we perform a 2-bin tomographic analysis to constrain cosmological parameters and investigate the possibility of photometric redshift bias in the CFHTLenS data. Under the baseline $\Lambda$CDM model we constrain $S_8 = \sigma_8(\Omega_\mathrm{m}/0.3)^{0.5} = 0.67 ^{\scriptscriptstyle+ 0.03 }_{\scriptscriptstyle- 0.03 }$ $(68\%)$, consistent with previous CFHTLenS analyses but in tension with \emph{Planck} \citep{Planck2015XIII}. Adding neutrino mass as a free parameter we are able to constrain $\sum m_\nu < 4.6\mathrm{eV}$ (95\%) using CFHTLenS data alone. Including a linear redshift dependent photo-$z$ bias $\Delta z = p_2(z - p_1)$, we find $p_1=-0.25 ^{\scriptscriptstyle+ 0.53 }_{\scriptscriptstyle- 0.60 }$ and $p_2 = -0.15 ^{\scriptscriptstyle+ 0.17 }_{\scriptscriptstyle- 0.15 }$, and tension with \emph{Planck} is only alleviated under very conservative prior assumptions. Neither the non-minimal neutrino mass or photo-$z$ bias models are significantly preferred by the CFHTLenS (2-bin tomography) data.

Light from distant galaxies is continuously deflected by the gravitational potential fluctuations of large-scale structure on its way to us, resulting in a coherent distortion of observed galaxy images across the sky --- weak gravitational lensing. This weak lensing effect is a function of both the geometry of the Universe (through the distance-redshift relation) and the growth of potential fluctuations along the line-of-sight, making it a tremendously rich cosmological probe; the statistics of the weak lensing fields are sensitive to the initial conditions of the potential fluctuations, the relative abundance of baryonic and dark matter (through baryon acoustic oscillations), the linear and non-linear growth of structure, the mass and hierarchy of neutrinos (\eg, \citealp{Jimenez2010}), dark energy and gravity on large scales (see, \eg, \citealp{Weinberg2013} for a review). The goal of cosmic shear analyses is to extract cosmological inferences from the statistics of the observed weak lensing shear field --- the distortion of observed galaxy shapes measured across the sky and in redshift. In \citet{Alsing2016} we developed a Bayesian hierarchical modelling (BHM) approach to infer the cosmic shear power spectrum (and thus cosmological parameters) from a catalogue of measured galaxy shapes and redshifts, building on previous work on cosmic microwave background (CMB) power spectrum inference \citep{Wandelt2004, Jewell2004, Eriksen2004, ODwyer2004, Chu2005, Larson2007, Eriksen2007} and large scale-structure analysis methods \citep{Jasche2010, Jasche2012, Jasche2013, Jasche2015}. The Bayesian hierarchical approach has a number of desirable features and advantages over traditional estimator-based methods: In contrast to frequentist estimators whose likelihoods need calibrating against large numbers of forward simulations (introducing assumptions and uncertainties that are often hard to propagate), the Bayesian approach explores the posterior distribution of the parameters of interest directly with clearly stated (and minimal) model assumptions, without the need for calibration. The Bayesian approach is exact and optimal, up to our ability to model the cosmic shear statistics (and systematics). Masks and complicated survey geometry are dealt with exactly and cleanly, in contrast to, \eg, pseudo-$C_\ell$ estimators that must carefully correct for mixing of $E$- and $B$-modes and physical (angular) scales arising from the mask inversion, which can be difficult in practice. The BHM approach can be readily extended to include models of non-Gaussian fields, exploiting more of the information-content of the weak lensing fields than is possible through $n$-point statistic estimators \citep{Jasche2013, Leclercq2015, Carron2012}. More generally, the BHM approach can also be extended to incorporate more of the weak lensing inference pipeline (\eg, shape measurement, PSF modelling etc), formally marginalising over nuisance parameters and systematics in a principled way and ultimately leading to more robust science at the end of the day (see \citet{Alsing2016, Schneider2015} for a discussion of the global hierarchical modelling approach to weak lensing). The joint map-power spectrum inference approach proceeds in two distinct steps: In step one we sample the joint posterior of the shear map and power spectrum\footnote{We use ``power spectrum" as a shorthand for the full set of $EE$, $BB$ and $EB$ auto- and cross-power spectra for all tomographic bins.}, using, \eg, the Gibbs sampling approach described in \citet{Alsing2016}. In step two, we construct a smooth posterior density from the power spectrum samples and proceed to infer cosmological parameters by Markov Chain Monte Carlo (MCMC) sampling the power spectrum posterior as a function of the cosmological parameters (under our model of interest). The posterior distribution of the power spectrum is a valuable intermediate product; cosmological parameter inference can be performed for a large number of cosmological (and systematics) models directly from the power spectrum posterior \emph{a posteriori}, without loss of information and without having to re-analyse the entire data-set (since the initial power spectrum inference was independent of cosmology, assuming only statistical isotropy of the lensing fields). However, the need to estimate the continuous posterior density from a set of posterior samples may come with practical challenges; this \emph{density estimation} step may introduce uncertainties at some level, and accurate density estimation may be challenging for analyses with a large number of tomographic bins (for example, with $10$ tomographic bins we may need to estimate the joint distribution of the $55$ tomographic cross-power spectra from a set of MCMC samples --- a challenging density estimation task). In this paper, we develop a second Bayesian approach to cosmic shear inference, whereby we jointly sample the shear maps and cosmological parameters, rather than the maps and power spectra. By going straight to cosmological parameters and bypassing the explicit power spectrum inference step, we circumvent the need to transform posterior samples into a continuous posterior density and hence avoid the prickly (high-dimensional) density estimation issues altogether. There are other advantages, too: by parametrising the power spectrum with a handful of cosmological parameters, the number of interesting parameters has been reduced from a few thousand power spectrum coefficients to typically $\lsim 10$ cosmological parameters -- this reduction in the parameter space will inevitably improve the sampling efficiency. Map-cosmology inference also extends more naturally to incorporate models for non-Gaussian shear where the power spectrum no longer fully specifies the lensing statistics. These benefits come at the cost of having to re-analyse the full shear maps for every cosmological model of interest, whereas the power spectrum posterior obtained in a cosmology-independent way represented a highly compressed intermediate product that could be used for (fast) \emph{a posteriori} cosmological parameter inference without the need to revisit the full data-set. However, we argue that nonetheless the map-cosmology inference scheme is a (comfortably) computationally feasible approach for current and future surveys. In this paper we apply the map-power spectrum and map-cosmology sampling schemes to infer power spectra, shear maps and cosmological parameters from the Canada-France-Hawaii Telescope (CFHTLenS) weak lensing survey - the first application to data. The structure of the paper is as follows: In \S \ref{sec:data} we describe the CFHTLenS data and compression of the full galaxy catalogue into pixelized (noisy) tomographic shear maps. In \S \ref{sec:formalism} we review the tomographic weak lensing formalism, and in \S \ref{sec:bayes_schemes} we describe the map-power spectrum and map-cosmology Gibbs sampling schemes. In \S \ref{sec:implementation} we outline the cosmological and systematics models considered in this analysis. In \S \ref{sec:results} we demonstrate the Bayesian inference schemes on simulations before presenting the inferred $E$- and $B$- mode power spectra, tomographic shear maps and cosmological parameters from the CFHTLenS cosmic shear data. We discuss computation costs and prospects for future surveys in \S \ref{sec:cost} and conclude in \S \ref{sec:conclusions}.

\label{sec:conclusions} We have applied Bayesian map-power spectrum and map-cosmology hierarchical inference schemes for extracting cosmic shear power spectra, shear maps and cosmological parameters from CFHTLenS -- the first application to data. Under the baseline flat-$\Lambda$CDM model, we obtain cosmological parameter constraints consistent with previous CFHTLenS analyses, and in-line with previous studies our inferred $(\sigma_8, \Omega_\mathrm{m})$ constraints are in tension with \emph{Planck} 2015 results at the $2\sigma$ level (under the baseline $\Lambda$CDM model). Extending the baseline model to include massive neutrinos, we are able to constrain the total neutrino mass to $\sum m_\nu < 4.6\mathrm{eV}$ (95\%) from CFHTLenS data alone. The inclusion of neutrino mass as an extra degree-of-freedom does little to alleviate the tension between CFHTLenS and Planck, and the more flexible model is not preferred over $\Lambda$CDM by the CFHTLenS data. Including the possibility of a linear redshift-dependent photo-$z$ bias $\Delta z = p_2(z - p_1)$ we find the CFHTLenS data prefer $p_1=-0.25 ^{\scriptscriptstyle+ 0.53 }_{\scriptscriptstyle- 0.60 }$ and $p_2 = -0.15 ^{\scriptscriptstyle+ 0.17 }_{\scriptscriptstyle- 0.15 }$, although both are consistent with zero. Including $p_1$ and $p_2$ as additional parameters under a broad flat prior completely alleviates tension between CFHTLenS and Planck in our analysis, due to a significant increase in the error bars from the extra degrees-of-freedom. Imposing an informative prior on $(p_1, p_2)$ from the CFHTLenS-BOSS cross-correlation analysis of \citet{Choi2015}, the impact of the photo-$z$ bias on the cosmological constraints is more modest and the tension with Planck remains. The CFHTLenS data alone do not prefer the more flexible photo-$z$ bias model over the baseline $\Lambda$CDM under either the flat or informative prior. As well as demonstrating the Bayesian inference schemes on current data, we have argued that for future large-area surveys such as \emph{Euclid} and \emph{LSST} the map-cosmology sampling scheme is of comparable computational cost to traditional estimator-likelihood sampling methods, and the map-power spectrum sampling scheme is also a computationally practical approach for obtaining cosmology independent power spectrum inference for future surveys. Both the map-power spectrum and map-cosmology inference schemes implemented in this work assume Gaussian lensing fields. Whilst this is appropriate (and optimal) on large scales, on smaller scales the lensing fields are well known to be non-Gaussian. This is both an opportunity and a curse: if we are to assume Gaussianity, we must rigorously validate the algorithms against $N$-body/non-Gaussian shear simulations to test for any model biases that could arise from the Gaussian assumption. On the other hand, the Bayesian hierarchical approach can be readily extended to include models for the non-Gaussian shear, allowing us to extract information beyond the two-point statistics and exploiting the full information content of the cosmological fields, leading to tighter constraints on cosmology and better science at the end of the day. We will explore the limits of the Gaussian approximation for weak lensing analyses and develop extended hierarchical models for non-Gaussian shear inference in future work.

1607.00008

1607

1607.04906_arXiv.txt

The Fermi bubbles are two lobes filled with non-thermal particles that emit gamma rays, extend $\approx$10 kpc vertically from the Galactic center, and formed from either nuclear star formation or accretion activity on Sgr A*. Simulations predict a range of shock strengths as the bubbles expand into the surrounding hot gas halo distribution ($T_{halo} \approx 2 \times 10^6$ K), but with significant uncertainties in the energetics, age, and thermal gas structure. The bubbles should contain thermal gas with temperatures between $10^6$ and $10^8$ K, with potential X-ray signatures. In this work, we constrain the bubbles' thermal gas structure by modeling the \ion{O}{7} and \ion{O}{8} emission line strengths from archival \textit{XMM-Newton} and \textit{Suzaku} data. Our emission model includes a hot thermal volume-filled bubble component cospatial with the gamma-ray region, and a shell of compressed material. We find that a bubble/shell model with $n \approx 1 \times 10^{-3}$ cm$^{-3}$ and with log($T$) $\approx$ 6.60-6.70 is consistent with the observed line intensities. In the framework of a continuous Galactic outflow, we infer a bubble expansion rate, age, and energy injection rate of $490_{-77}^{+230}$ km s$^{-1}$, $4.3_{-1.4}^{+0.8}$ Myr, and $2.3_{-0.9}^{+5.1} \times 10^{42}$ erg s$^{-1}$. These estimates are consistent with the bubbles forming from a Sgr A* accretion event rather than from nuclear star formation.

\label{section.introduction_chap_fb} The Fermi bubbles are important Galactic structures that were recently discovered by the \textit{Fermi Gamma-ray Space Telescope} \citep{su_etal10}. The bubbles are two diffuse lobes of material extending $\sim$50$\arcdeg$ above and below the Galactic plane ($\approx$10 kpc at the Galactic center). Their surface brightness shows little variation on the sky, their gamma-ray spectrum follows a power law with $dN/dE \propto E^{-2}$ between $\approx$1 and 200 GeV, and they have a counterpart in microwaves, known as the $Wilkinson\ Microwave\ Anisotropy\ Probe$ ($WMAP$) haze \citep{dobler_finkbeiner08, dobler_etal10, su_etal10, ackermann_etal14}. It is still unclear what produces the gamma rays, but all plausible mechanisms imply that energetic cosmic-ray particles exist within the bubbles. This inference combined with the bubbles' size and location on the sky suggests that they are affiliated with a massive energy injection event near the Galactic center. The bubbles' morphology is similar to wind-blown bubbles observed in other galaxies, indicating that they formed from either a period of enhanced nuclear star formation or a Sgr A* outburst event (see \citealt{veilleux_etal05} for a review). Star formation can drive outflows through a combination of stellar winds from young stars and multiple type-II supernova explosions \citep[e.g., ][]{leitherer_etal99}, while black hole accretion episodes can produce energetic jets or winds that inflate a cavity with thermal and non-thermal particles \citep[e.g., ][]{mcnamara_nulsen07,yuan_narayan14}. Both of these scenarios are critical events in galaxy evolution, as they both can deposit significant amounts of energy into the rest of the galaxy on scales $\gtrsim$10 kpc (see \citet{mcnamara_nulsen07} for a review). However, the details of these ``feedback'' effects (mass displacement, energy transport, etc.) are poorly understood since we observe them in external galaxies. The Fermi bubbles are a unique laboratory for understanding these processes since we can spatially resolve the bubbles across multiple wavebands. A popular strategy to probe these effects and bubbles' origins has been the use of magnetohydrodynamic (MHD) simulations to reproduce the bubbles’ global morphology and non-thermal properties. Simulations produce cosmic rays either from a black hole accretion event \citep{zubovas_etal11,guo_mathews12a,guo_mathews12b,yang_etal12,yang_etal13,zubovas_nayakshin12,mou_etal14,mou_etal15}, from nuclear star formation activity \citep{crocker12,crocker_etal14,crocker_etal15,sarkar_etal15,ruszkowski_etal16}, or in-situ as the bubbles evolve \citep{cheng_etal11,cheng_etal15b,mertsch_sarkar11,fujita_etal14,lacki14,sasaki_etal15}, and compare the non-thermal emission to the bubbles' gamma-ray emission. All of these origin scenarios can reproduce the bubbles' morphology, but they imply significantly different input energetics and timescales required to inflate the bubbles ($\dot{E} \gtrsim 10^{41}$ erg s$^{-1}$, $t \lesssim 5$ Myr for black hole accretion compared to $\dot{E} \lesssim 5 \times 10^{40}$ erg s$^{-1}$, $t \gtrsim 50$ Myr for star formation). This variation in the feedback rate is a significant uncertainty in how the bubbles impact the Galaxy, but there are additional factors that can constrain the characteristic bubble energetics. Constraining the bubbles' thermal gas distribution is a promising avenue to solve this problem, since the characteristic densities and temperatures should be significantly different depending on the bubble energetics. In the framework of expanding galactic outflows and shocks \citep[e.g., ][]{veilleux_etal05}, a higher energy input rate leads to a higher plasma temperature and a larger expansion rate for a fixed bubble size and ambient density. Thus, the plasma temperature at the interface between the bubbles and surrounding medium encodes information on the bubbles' shock strength, expansion properties, and overall energy input rate. A generic prediction from simulations and observations of galactic outflows is that the bubbles are overpressurized and hotter than the surrounding medium ($\gtrsim 2 \times 10^6$ K), implying that the bubbles' thermal gas should have signatures at soft X-ray energies. Indeed, the bubbles appear to be bounded by X-ray emission seen in the \textit{ROSAT} 1.5 keV band \citep{bh_cohen03, su_etal10}; however, these observations do not constrain the bubbles' intrinsic thermal gas structure since the broad-band images are a weak temperature diagnostic. Spectral observations with current X-ray telescopes are a much better temperature diagnostic for this type of environment. Initial efforts to observe the bubbles in soft X-rays with \textit{Suzaku} and \textit{Swift} and constrain their temperature and shock strength were carried out by several groups \citep{kataoka_etal13,kataoka_etal15,tahara_etal15}. \citet{kataoka_etal15} extracted soft X-ray background (SXRB) spectra in the 0.5--2.0 keV band for 97 sight lines that pass through the Fermi bubbles, and fit the spectra with thermal plasma models. They consistently measured plasma temperatures of $kT$ = 0.3 keV for these sight lines, which is systematically higher than the characteristic temperature measured in sight lines away from the Galactic center \citep[$kT \approx$0.2 keV; ][]{hs13}. From this temperature ratio, they inferred a shock Mach number of $\mathcal{M} \approx $0.3 keV / 0.2 keV = 1.5, and corresponding expansion rate of $\approx$300 km s$^{-1}$. This is a valuable attempt to constrain these quantities, but the analysis assumes that the Fermi bubble plasma dominates the hotter spectral component. In practice, there are other known emission sources that contribute to the SXRB spectrum, and accounting for this emission can change the inferred thermal gas temperature. The Milky Way hosts a hot gas distribution with $T \approx 2 \times 10^6$ K extending on scales $\gtrsim$10 kpc based on shadowing experiments from \textit{ROSAT} all-sky data \citep{snowden_etal97, kuntz_snowden00}. This plasma is believed to dominate any SXRB spectrum, with \ovii and \oviii being the characteristic observed line transitions \citep[e.g., ][]{mccammon_etal02, yoshino_etal09, hs12}. The structure of this extended plasma distribution has been debated in the literature, but numerous studies on but numerous studies on both absorption and emission line strengths indicate that the plasma is spherical and extends to at least $r \sim 50$ kpc \citep{fang_etal06, bregman_ld07, gupta_etal12, fang_etal13, miller_bregman13, miller_bregman15}. In particular, Miller \& Bregman (\citeyear{miller_bregman15}, defined as \citetalias{miller_bregman15} henceforth) modeled a set of 648 \oviii emission line intensities from Henley \& Shelton (\citeyear{hs12}, defined as \citetalias{hs12} henceforth), and found that a hot gas density profile with $n \propto r^{-3/2}$ extending to the virial radius reproduces the observed emission line intensities. These modeling studies have placed useful constraints on the Galactic-scale hot gas distribution, but also highlight the fact that this extended plasma is likely the dominant emission source in all 0.5--2.0 keV band spectra. In this study, we expand the analysis Kataoka et al. (\citeyear{kataoka_etal15}, defined as \citetalias{kataoka_etal15} henceforth) analysis by modeling the combined emission from the Fermi bubbles and hot gas halo present in \oviii emission line measurements. We modify the Galactic-scale hot gas models from \citetalias{miller_bregman15} to include a geometry, density, and temperature structure for the Fermi bubbles. Given a set of model parameters, we predict the contribution to the emission made by the Fermi bubbles and hot gas halo along any sight line. This results in a more careful comparison between the Fermi bubbles' emission and the total observed emission in any SXRB measurement. The \oviii observations used in our analysis consist of published \textit{XMM-Newton} measurements from \citetalias{hs12}, and a new \textit{Suzaku} data set produced for this work. The \textit{XMM-Newton} data are mostly the same measurements used in \citetalias{miller_bregman15}, but we now include data near the Fermi bubbles. We supplement these data with archival \textit{Suzaku} measurements of SXRB spectra, which more than doubles the number of emission line measurements projected near the bubbles. These data are processed in a similar way to the \textit{XMM-Newton} data reduction outlined in \citetalias{hs12}, resulting in a uniformly processed data set of emission line intensities from the SXRB. Following the methodology from \citetalias{miller_bregman15}, we constrain the Fermi bubbles' density and temperature structure by finding the parametric model that is most consistent with the observed emission line intensities. We measure the characteristic bubble temperature from analyzing the distribution of observed \oviii/\ovii line ratios near the bubbles, and the characteristic bubble density from explicitly modeling the \oviii emission line intensities. We infer a similar bubble shock strength compared to the \citetalias{kataoka_etal15} analysis, and discuss the systematic differences between our approaches and results and theirs. We also estimate the bubbles' age and energy input rate, and these with the possible formation mechanisms discussed above. The rest of the paper is outlined as follows. Section~\ref{section.data_chap_fb} discusses how we compiled our emission line sample, including an overview of the \textit{XMM-Newton} data set and the \textit{Suzaku} data processing. Section~\ref{section.model_chap_fb} definesdefines our parametric density and temperature model and discusses our line intensity calculation. Section~\ref{section.results_chap_fb} discusses our model fitting routine and results. Section~\ref{section.discussion_chap_fb} discusses our constraints on the Fermi bubbles in the context of galactic outflows, previous X-ray studies, and simulations. Section~\ref{section.conclusions_chap_fb} summarizes our results.

1607.04906

1607

1607.08405_arXiv.txt

{We present the results of a comprehensive numerical simulation of the environment around three exoplanet-host stars (HD\,1237, HD\,22049, and HD\,147513). Our simulations consider one of the latest models currently used for space weather studies in the Heliosphere, with turbulent Alfv\'en wave dissipation as the source of coronal heating and stellar wind acceleration. Large-scale magnetic field maps, recovered with two implementations of the tomographic technique of Zeeman-Doppler imaging, serve to drive steady-state solutions in each system. This paper contains the description of the stellar wind and inner astrosphere, while the coronal structure was previously discussed in \citetads{2016A&A...588A..28A}. The analysis includes the magneto-hydrodynamical properties of the stellar wind, the associated mass and angular momentum loss rates, as well as the topology of the astrospheric current sheet in each system. A systematic comparison among the considered cases is performed, including two reference solar simulations covering activity minimum and maximum. For HD\,1237, we investigate the interactions between the structure of the developed stellar wind, and a possible magnetosphere around the Jupiter-mass planet in this system. We find that the process of particle injection into the planetary atmosphere is dominated by the density distribution rather than velocity profile of the stellar wind. In this context, we predict a maximum exoplanetary radio emission of 12 mJy at 40 MHz in this system, assuming the crossing of a high-density streamer during periastron passage. Furthermore, in combination with the analysis performed in \citetads{2016A&A...588A..28A}, we obtain for the first time a fully simulated mass loss-activity relation, which is compared and discussed in the context of the relation based on astrospheric detections proposed by \citetads{2005ApJ...628L.143W}. Finally, we provide a characterisation of the global 3D properties of the stellar wind of these systems, at the inner edges of their habitable zones.}

\noindent As well as driving stellar activity cycles, magnetic fields strongly influence different aspects of the stellar structure and evolution. It is known that they play a major role in the coronal heating processes in the Sun and other late type stars (\citeads{2015RSPTA.37340269D}; \citeads{2015RSPTA.37340259T}), as well as in the generation of persistent stellar winds and astrospheres (see \citeads{2004LRSP....1....2W}). These stellar winds are crucial to understand the evolution of rotation and magnetic activity in cool stars on the early main sequence. G to K type stars tend to rotate rapidly on the Zero Age Main Sequence (ZAMS); braking torques exerted by winds cause them to spin down, losing most of their angular momentum within the first 500 Myr (\citeads{2016A&A...587A.105A}; \citeads{2010ApJ...721..675B}). Strong winds from the young Sun have been used to explain both the stripping of the Martian atmosphere (\citeads{2013oepa.book.....L}; \citeads{2009AsBio...9...55T}), and address the ``faint young Sun paradox''. This paradox is that terrestrial geological records indicate that water existed in liquid form very early in the history of Earth and Mars, despite the young Sun having only 70\% of its current luminosity. Solar wind sputtering is a leading candidate to explain the loss of Mars' once-thick atmosphere, because Mars is not protected by a strong magnetosphere, unlike Earth (see \citeads{2007SSRv..129..245L}). However, recent work presented by \citetads{2014ApJ...781L..33W} argues that while the young Sun was more magnetically active, it does not necessarily follow that it would have hosted stronger winds. This last result comes from the close relation between the winds in Sun-like stars and their surrounding astrospheres. In the case of the Sun, the solar wind creates a comet-like bubble (the heliosphere) that extends far past the orbits of the planets, and interacts with the local interstellar medium (LISM)\footnote[2]{This classical shape of the heliosphere has been recently revisited in various observational and numerical works, pointing towards a far more complex description including magnetized jets (see \citeads{2013ApJ...771...77M}; \citeads{2015ApJ...800L..28O}; \citeads{2015ApJ...808L..44D}).}. The heliosphere is populated by hot hydrogen atoms created through charge exchange between the ionized gas in the solar wind and the cold LISM hydrogen. Hot hydrogen builds up particularly in the region between the termination shock and the heliopause. This is the region which the Voyager mission may recently have crossed (\citeads{2013Sci...341.1489G}), although this is still a matter of debate (see \citeads{2014ApJ...789...41F}; \citeads{2015ApJ...806L..27G}). This hydrogen wall is detected as extra H\,I Lyman-$\alpha$ absorption in the UV spectra of cool stars. Stronger winds result in a larger astrosphere and increased absorption (\citeads{2014ASTRP...1...43L}). By measuring the column densities and velocities of this extra absorption it is possible to derive the only available observationally-driven estimates of mass loss rates in cool stars (see \citeads{2015ASSL..411...19W}). However, these estimates strongly depend on the assumed characteristics and topology of the LISM (see \citeads{2014ASTRP...1...43L}), for which there is still no complete agreement in the literature (e.g., \citeads{2009ApJ...696.1517K}; \citeads{2014A&A...567A..58G}; \citeads{2015ApJ...812..125R}). On the theoretical and modelling side, recent studies have provided different frameworks for the stellar wind origin, behaviour, and influence in the angular momentum evolution of late-type stars. Among the 1D and 2D models, a non-comprehensive list includes semi-empirical approaches for thermally-driven winds, within a hydro- (e.g. Johnstone et al. \citeyearads{2015A&A...577A..27J}, \citeyearads{2015A&A...577A..28J}) or magneto-hydrodynamic (MHD) regime (e.g. Matt et al. \citeyearads{2008ApJ...678.1109M}, \citeyearads{2012ApJ...754L..26M}; R\'eville et al. \citeyearads{2015ApJ...798..116R}, \citeyearads{2015ApJ...814...99R}), physically-motivated descriptions involving scaling relations for the stellar magnetic fields, rotation periods, convective properties, and X-ray fluxes (e.g. \citeads{2012ApJ...746...43R}; \citeads{2016MNRAS.458.1548B}), and semi-analytic and numerical formulations based on Alfv\'en wave MHD turbulence (e.g. \citeads{2011ApJ...741...54C}; \citeads{2013PASJ...65...98S}). While providing reasonable agreement in the rotational evolution of late-type stars at different stages (e.g. \citeads{2013A&A...556A..36G}, \citeads{2015ApJ...799L..23M}), such approaches are very generic and cannot capture the specifics of the stellar wind of a given system. The same is true for the complex interplay between the magnetic field topology, coronal structure, and the stellar wind. These elements are fundamental for a better understanding of the environment around planet-hosting stars, including the relative influence of the wind and the high-energy emission on the exoplanetary conditions and habitability (see \citeads{2003ApJ...598L.121L}; \citeads{2013oepa.book.....L}; \citeads{2014ApJ...795..132S}; \citeads{2014RSPTA.37230084F}). Such detailed descriptions are crucial for the current and future perspectives in the area of exoplanetary characterisation from the ground and space (see \citeads{2014Natur.513..358P}; \citeads{2014Natur.513..353H}). In this context, we presented in \citetads{2016A&A...588A..28A} the initial results of a detailed 3D numerical study aimed at simulating the environment around planet-hosting stars. This previous article described the developed coronal structure and high-energy environment on three exoplanet-hosts, namely HD\,22049 (K2V), HD\,1237 (G8V), and HD\,147513 (G5V). The basic stellar and planetary (orbital) properties in these systems are listed in Table \ref{table_0}. We employed one of the latest physics-based models, compatible with recent satellite solar observations (see \citeads{2007Sci...318.1574D}; \citeads{2011Natur.475..477M}), and currently used for space weather forecast in the solar system (see \citeads{2012JCoPh.231..870T}). This model considers a data-driven approach where surface magnetic field distributions (i.e. \textit{magnetograms}) are decomposed in a high-degree spherical harmonics expansion and implemented as an initial condition, following the methodology presented in \citeads{2011ApJ...732..102T}. Then the simulation evolves self-consistently, calculating coronal heating and stellar wind acceleration based on Alfv\'en wave turbulence dissipation (\citeads{2013ApJ...764...23S}; \citeads{2014ApJ...782...81V}). \begin{table} \caption{Basic observational properties of the considered systems.} \label{table_0} \centering {\small \begin{tabular}{l c c c c c c c} \hline\hline & & & & & & \\[-9pt] Star ID & $M_{*}$ & $R_*$ & $P_{\rm rot}$ & $M_{\rm p}\sin i$ & $a$ & $e$ \\ & [$M_{\odot}$] & [$R_{\odot}$] & [days] & [$M_{\jupiter}$] & [AU] &\\ \hline & & & & & & \\[-8pt] HD 1237\,\tablefootmark{a} & 0.86 & 1.00 & 7.00 & 3.37 & 0.49 & 0.51\\ HD 22049\,\tablefootmark{b,$\dagger$} & 0.74 & 0.86 & 11.68 & 1.05 & 3.38 & 0.25\\ HD 147513\,\tablefootmark{c} & 0.98 & 1.07 & 10.00 & 1.21 & 1.32 & 0.26\\ \hline \end{tabular}} \tablefoot{References for each system: \tablefoottext{a}{\citetads{2001A&A...375..205N}; \citetads{2010ApJ...720.1290G}; \citetads{2015A&A...582A..38A}.} \tablefoottext{b}{\citetads{1993ApJ...412..797D}; \citetads{1996ApJ...466..384D}; \citetads{2005ApJS..159..141V}; \citetads{2006ApJ...646..505B}.} \tablefoottext{c}{\cite{2004A&A...415..391M}, \citetads{2007ApJS..168..297T}; \citeads{2016A&A...585A..77H}.} \\\tablefoottext{$\dagger$}{The listed orbital parameters in \citetads{2016A&A...588A..28A} for the exoplanet in this system were taken from the discovery paper (\citeads{2000ApJ...544L.145H}).}} \end{table} For other stars this information can be nowadays retrieved (to some extent) from high-resolution spectropolarimetric observations and the technique of Zeeman-Doppler imaging (ZDI, \citeads{1989A&A...225..456S}; \citeads{1991A&A...250..463B}; \citeads{1997A&A...326.1135D}; \citeads{2000MNRAS.318..961H}; \citeads{2002A&A...381..736P}). For the stellar systems of interest, ZDI maps were previously recovered by \citetads{2014A&A...569A..79J}, \citetads{2015A&A...582A..38A}, and \citetads{2016A&A...585A..77H}. As discussed in detail in \citetads{2016A&A...588A..28A}, we considered two different implementations of this mapping technique (i.e. ZDI and SH-ZDI). The first one considers an image reconstruction with independent pixels for the radial, meridional, and azimuthal components, yielding an unconstrained distribution of the magnetic field (\citeads{1991A&A...250..463B}; \citeads{1997A&A...326.1135D}). The second one follows the methodology presented in \citetads{2001MNRAS.322..681H} and \citetads{2006MNRAS.370..629D}, where the vector field is reconstructed in terms of a spherical harmonics decomposition. The main difference with the previous approach is the possibility to impose physical and geometrical constraints to the final reconstructed topology (i.e. purely potential or toroidal fields, symmetry or antisymmetry). In this way, the SH-ZDI implementation permits the completion of the map in the un-observed hemisphere (due to the inclination of the star). As explained in \citetads{2016A&A...588A..28A}, the SH-ZDI maps were completed preserving the same level of fit to the spectropolarimetric observations as the one obtained with the standard ZDI procedure. This last consideration is fundamental due to the fact that the field strengths in the final map depend on the degree of fit to the observations. This ZDI\,/\,SH-ZDI comparison was performed in order to evaluate the effect of these observational constraints in our simulations. Additionally, the analysis included a comprehensive evaluation procedure of two solar simulations against satellite data, covering activity minimum (CR\,1922) and maximum (CR\,1962), that were spatially-filtered to a similar level of resolution as the ZDI maps. Finally, the results of the ZDI-driven stellar simulations were consistently compared with the solar cases, and with various observational estimates of their coronal conditions. This paper provides continuity to this previous work, to include the stellar wind properties and inner astrospheric structure. A description of the numerical set up, boundary conditions, and general characteristics in each simulation domain is provided in Sect. \ref{sec_model}. Section \ref{sec_results} contains the results for each system, including the reference solar cases (i.e., activity minimum and maximum). We discuss our results in the context of previous observational and numerical studies in Sect. \ref{sec_Analysis}, and the main conclusions of our work are summarised in Sect. \ref{sec_Summary}.

\label{sec_Summary} \noindent We carried out elaborated simulations of the stellar wind and inner astrospheric structure of three planet-hosting stars (HD\,22049, HD\,1237, and HD\,147513), using the Space Weather Modelling Framework (SWMF, T\'oth et al. \citeyearads{2005JGRA..11012226T}, \citeyearads{2012JCoPh.231..870T}). This paper complements the study presented in \citetads{2016A&A...588A..28A}, which contains the results of the coronal structure modelling of these systems. Steady-state solutions were obtained for two coupled simulation domains, ranging from $1 - 30$ $R_{*}$ (SC domain) and from $25 - 215$ $R_{*}$ (IH domain). Large-scale magnetic field maps of these stars, recovered with Zeeman-Doppler imaging, serve to drive the solutions inside the SC domain, which are coupled self-consistently for a combined solution in the IH domain. A summary of our results and main conclusions is provided below: \begin{itemize} \item [-] Following \citetads{2016A&A...588A..28A}, simulations driven by two sets of similar large-scale magnetic field distributions (i.e. ZDI and SH-ZDI) were compared. It is worth noting that both sets of magnetic field maps provided equivalently good fits to the observations and showed substantial similarities in the overall structure of the stellar wind. However, several differences in the magneto-hydrodynamic properties of the solutions were found, including $\sim$10\,--\,30\% denser and $\sim$25\,--\,35\% colder stellar winds in the SH-ZDI solutions compared to the ZDI cases. In addition, the SH-ZDI simulations led to larger values in the average Alfv\'en surface size (by a factor of $\sim$1.5), the mass loss rate $\dot{M}$ (by a factor of $\sim\,$3\,--\,4), and the angular momentum loss rate $\dot{J}$ (by roughly one order of magnitude). Therefore, the values listed in Table \ref{table_1} should be actually interpreted as predicted ranges from this ZDI-driven model. These variations arise as a consequence of the available magnetic energy to heat the corona and accelerate the wind, which in turn, relates to the different field strengths and map completeness provided by the ZDI and SH-ZDI reconstructions. This strongly differs from previous studies where older implementations of the numerical code used here are considered, and where the completeness in the driving magnetic field distributions yield no significant changes in the wind structure (e.g. \citeads{2012MNRAS.423.3285V}; see Sect. \ref{sec_alfven}). \item[-] The results from two different solar simulations, covering activity minimum (CR\,1922) and maximum (CR\,1962), were also considered. We showed that this numerical framework properly recovers the expected structure of the solar wind, including thermodynamical properties (e.g. density, temperature), mass loss and angular momentum loss rates ($\dot{M}_{\odot}$ and $\dot{J}_{\odot}$, respectively; see Table \ref{table_1}), and global topology during each activity state (Figs. \ref{fig_1} and \ref{fig_5}). However, the solar maximum simulation showed an over-enhanced plasma density at 1 AU (Sect. \ref{sec_wind-HZ}, Table \ref{table_3}), as a consequence of an overestimated mass loss rate (by $\sim$\,40\%)\footnote[2]{This corresponds to a very rough estimate, as it relies on the single spatial point measurements and location of Voyager II as reference (\url{http://voyager.jpl.nasa.gov/mission/weekly-reports/index.htm})}. This is interpreted as the result of a considerable fraction of missing mixed polarity regions in the driving magnetogram, which was artificially degraded for a more consistent comparison with the stellar cases (see \citeads{2016A&A...588A..28A}). \item[-] In general, the stellar wind solutions showed a clear relation with the driving magnetic field distribution, and the developed coronal structure in each case. For HD\,22049 (Figs. \ref{fig_2} and \ref{fig_6}) and HD\,1237 (Figs. \ref{fig_4}, \ref{fig_8} and \ref{fig_9}) various fast-wind regions appeared self-consistently in the simulations, nearly perpendicular to the astrospheric current sheet structure (defined by $B_{\rm r}$ = 0), and with a spatial correspondence with the dominant features in their lower corona (e.g. coronal holes, see \citeads{2016A&A...588A..28A}). The radial wind velocity in these regions reached up to $\sim$\,1100 km s$^{-1}$ in the ZDI simulations, dropping to $\sim$\,700 km s$^{-1}$ in the SH-ZDI cases (see Sect. \ref{sec_resultsIH}). \item[-] On the other hand, the simulation of HD\,147513 yielded a much more complex wind solution (Figs. \ref{fig_3} and \ref{fig_7}), compared to what could have been expected from the simple magnetic field distribution driving the simulation (\citeads{2016A&A...588A..28A}). A highly warped astrospheric current sheet was obtained in this case, over which a dominant slow-wind component ($u_{\rm r}$\,$\simeq$\,500 km s$^{-1}$) was developed. While these results could be affected by the comparatively low-resolution of the SH-ZDI map driving the simulation (see \citeads{2016A&A...585A..77H}), this example indicates that numerical descriptions based on first order extrapolations of surface magnetic field properties alone, cannot provide a complete picture of the wind complexity in a given system (e.g. \citeads{2012ApJ...754L..26M}). \item[-] For HD\,1237 we investigated in detail the wind environment and the conditions experienced by the exoplanet of this system (Sect. \ref{sec_HD1237b}). For this purpose we additionally coupled the Global Magnetosphere (GM) module of the SWMF, to the developed wind structure inside the IH domain. For each simulation (i.e. ZDI and SH-ZDI cases), two representative spatial locations were considered (Figs. \ref{fig_8} and \ref{fig_9}). These included a low-density, fast wind stream, and a high-density, slow wind region. This analysis showed that the density structure of the stellar wind dominates, over the wind velocity, the process of particle injection into the planetary atmosphere (see Table \ref{table_2}). This is a consequence of the large density gradients obtained in the wind solutions (i.e., dense streamers with increments up to 4 orders of magnitude in $n$), compared to the relatively narrow range of resulting radial wind speeds (with variations only up to a factor of 2 in $u_{\rm r}$ for the entire 3D domain). \item[-] Following \citeads{2005MNRAS.356.1053S}, we additionally calculated the amount of exoplanetary radio emission from the wind-magnetospheric interaction in this system. We obtained a maximum radio flux on Earth of $F_{\oplus}^{R} \simeq 12$ mJy at 40 MHz, associated with a high-density streamer crossing during periastron passage (at 0.25 AU, \citeads{2001A&A...375..205N}). This value is reduced by an order of magnitude during the orbital motion of the planet (approximately at 2/3 of a right-hand oriented orbit with respect to periastron, see Fig. \ref{fig_10}). Our maximum emission prediction is lower by a factor of $\sim$\,2 compared to the estimates of \citetads{2005MNRAS.356.1053S}, which were based on the mass loss-activity relation of \citeads{2002ApJ...574..412W}, and the assumption of a spherically symmetric wind. Given the system's low declination, SKA is the only facility which could robustly detect and analyse this emission. \item[-] From our simulations, and applying the methodology explained in \citetads{2014ApJ...783...55C} and \citetads{2015ApJ...813...40G}, we calculated the mass loss rate, $\dot{M}$, and angular momentum loss rate, $\dot{J}$, in these systems. We obtained absolute $\dot{M}$ values, ranging from approximately 1\,$\dot{M}_{\odot}$ up to $\sim$\,7\,$\dot{M}_{\odot}$, and $\dot{J}$ within a broader range of $\sim$\,1--\,60 times the solar prediction (Table \ref{table_1}). In combination with the results for the coronal structure (\citeads{2016A&A...588A..28A}), we constructed, for the first time, a fully simulated mass loss-activity relation, expressed as $\dot{M}_{\rm (sim)} \propto F_{\rm X\,(sim)}^{\rm \,\gamma}$ with $\gamma = 0.79^{\,+\,0.19}_{\,-\,0.15}$. A thoughtful discussion is presented in Sect. \ref{sec_B-Mloss}, comparing this result with the observational relation of \citetads{2005ApJ...628L.143W} (e.g. $\dot{M} \propto F^{1.34 \pm 0.18}_{\rm X}$), exploring various possibilities that could explain the discrepancy in these relations. \item[-] Finally, by exploiting the 3D capabilities of our simulations we characterised the stellar wind structure at the inner edge of the Habitable Zone (HZ) of these systems (Sect. \ref{sec_wind-HZ}). The optimistic and conservative limits of this boundary, provided in the Habitable Zone Gallery (HZG, \citeads{2012PASP..124..323K}), were considered. We included a 10$^{\circ}$ range in orbital inclination (e.g. Fig. \ref{fig_12}), in order to provide more realistic stellar wind parameters (allowing possible off-the-equator variations), and to capture the region where the majority of exoplanets have been found so far (\citeads{2014PASP..126..827H}). The results of this characterisation are presented in Table \ref{table_3}, and consider all the magneto-hydrodynamic properties of the stellar wind in these systems. Using the solar simulations, reference calculations at the locations of Mercury, Mars, and the Earth are also provided. These results will be used in a future study to perform a dynamical parametrisation of the inner edge of the HZ in these and other systems, accounting for the effects due to the stellar wind and the high-energy environment of the host star. \end{itemize}

1607.08405

1607

1607.08543_arXiv.txt

\vskip12pt In solar and solar-terrestrial physics the phenomenon of hysteresis has long been known [1, 2, 3]. This effect manifests itself (for different levels of solar activity) in an ambiguous dependence of the radiation of the solar photosphere, chromosphere and corona, as well as indices characterizing the state of Earth's ionosphere. We have chosen to study the effect of the hysteresis in solar activity versus off the index $F_{10.7}$ (radio flux on the wave 10.7 cm -- 2800 MHz). The aim of this work is to study the variations of indices of activity during three solar cycles: (1) the changes in the values of the activity indices (AI - Activity Indices) depending on the solar activity level (i.e. versus the $F_{10.7}$ index), and (2) the analysis of the ambiguous dependence of the relationship of indices of activity versus $F_{10.7}$ on the rise and decline phases of the solar cycle. For the monthly values of 7 activity indices AI we have determined the coefficients of the quadratic regression for the correlations with an $F_{10.7}$ index (AI $\leftrightarrow $ $F_{10.7}$). Also we determined the coefficients of a quadratic regression for 5 daily values of indices of solar and ionospheric activity (AI $\leftrightarrow $ $F_{10.7}$). Thus a series of observations of the indices of the SSN, TSI and $F_{L \alpha}$ were analyzed for monthly and daily data. We also analyzed the hysteresis of the solar-type stars (with a stable cycles of activity similar to the sun's cycles) between the radiation fluxes from their photospheres and chromospheres (in the lines H and K CaII) from the Mount Wilson HK-project and Lowell observatory data [4,5]. In this study we used the data of observations of $F_{10.7}$, SSN and other indices of solar and ionospheric activity from the archives of the NOAA National Geophysical Data Center and of the Solar-Geophysical Data Reports [6, 7].

The effect of hysteresis is typical not only for pairs of activity indices in 11-yr solar cycles, but for the stars of the HK- project with detected the stable cycles of activity similar to the Sun. Hysteresis is a real delay in the onset of the maximum and decline phase of solar and stellar activity and is an important key in the search for physical processes responsible for changing radiation at different wavelengths. The effect of hysteresis is apparently a common property of astronomical systems, which are characterized by different manifestations of the cyclic activity associated with the time evolution of magnetic fields. The hysteresis of foF2 is exist due to the ionospheric response to variations in solar activity, in particular, we can see the hysteresis in EUV radiation and hysteresis of the geomagnetic activity index AP. All this hysteresis effects is due to the hysteresis of the solar Dynamo. For pairs of indices of solar activity and solar-type stars activity the effect of hysteresis appears in different ways. The curves of a hysteresis on the rise and decline cycle's phases vary from one cycle to another.

1607.08543

1607

1607.08069_arXiv.txt

{The 70-month Swift/BAT catalogue provides a sensitive view of the extragalactic X-ray sky at hard energies ($>$10 keV) containing about 800 active galactic nuclei (AGN). We explore its content in heavily obscured, Compton-thick AGN by combining the BAT (14-195 keV) with the lower energy XRT (0.3-10 keV) data. We apply a Bayesian methodology using Markov chains to estimate the exact probability distribution of the column density for each source. We find 53 possible Compton-thick sources (probability range 3 -- 100\%) translating to a $\sim$7\% fraction of the AGN in our sample. We derive the first parametric luminosity function of Compton-thick AGN. The unabsorbed luminosity function can be represented by a double power law with a break at $L_{\star} \sim 2 \times 10^{42}$ $\rm ergs~s^{-1}$ in the 20-40 keV band. The Compton-thick AGN contribute $\sim$17\% of the total AGN emissivity. We derive an accurate Compton-thick number count distribution taking into account the exact probability of a source being Compton-thick and the flux uncertainties. This number count distribution is critical for the calibration of the X-ray background synthesis models, i.e. for constraining the {intrinsic} fraction of Compton-thick AGN. We find that the number counts distribution in the 14-195 keV band agrees well with our models which adopt a low intrinsic fraction of Compton-thick AGN ($\sim12\%$) among the total AGN population and a reflected emission of $\sim5\%$. In the extreme case of zero reflection, the number counts can be modelled with a fraction of at most 30\% Compton-thick AGN of the total AGN population and no reflection. Moreover, we compare our X-ray background synthesis models with the number counts in the softer 2-10 keV band. This band is more sensitive to the reflected component and thus helps us to break the degeneracy between the fraction of Compton-thick AGN and the reflection emission. The number counts in the 2-10 keV band are well above the models which assume a 30\% Compton-thick AGN fraction and zero reflection, while they are in better agreement with models assuming 12\% Compton-thick fraction and 5\% reflection. The only viable alternative for models invoking a high number of Compton-thick AGN is to assume evolution in their number with redshift. For example, in the zero reflection model the intrinsic fraction of Compton-thick AGN should rise from 30\% at redshift z$\sim0$ to about 50\% at a redshift of z=1.1. }

X-ray surveys provide the most efficient way to detect AGN (see \citealt{brandt2015} for a recent review). The 4 Ms Chandra Deep Field-South Survey (CDFS) catalog uncovered a surface density of 20,000 AGN/$\rm deg^{2}$ \citep{xue2011}, a number which is expected to increase significantly with the additional 3Ms observations to be released within this year. In comparison, optical surveys which detect the most luminous AGN (QSOs) yield surface densities of a few hundred AGN per square degree \citep{ross2013}. The huge contrast in the efficiency between X-ray and optical surveys lies in the fact that X-ray surveys detect the most highly obscured and low luminosity AGN. The deficit of AGN in optical surveys could only partially be recovered using either variability \citep{villforth10} or emission line ratio diagnostics \citep{bongiorno10}. On the other hand, infrared selection techniques, although not affected by obscuration \citep{Stern2012, Donley2012, Mateos2013, Assef2013}, can miss a significant fraction of the less luminous AGN because of contamination by the host galaxy. In conclusion, it is only the X-ray surveys that reliably track the history of accretion into supermassive black holes (SMBH) \citep{ueda2014, miyaji2015, aird2015a, aird2015b, ranalli2015}. Even the extremely efficient X-ray surveys performed by {\it XMM-Newton} and {\it Chandra} in the 0.3-10 keV band face difficulties when they encounter the most heavily obscured AGN, i.e. those with column densities above $10^{24}$ $\rm cm^{-2}$. These are the Compton-thick AGN where the attenuation of X-rays is due to Compton scattering on electrons rather than photoelectric absorption, which is the major attenuation mechanism at lower column densities. The deep {\it Chandra} and {\it XMM-Newton} surveys found a number of Compton-thick AGN at moderate to high redshift \citep{comastri2011, georgantopoulos2013, brightman2014, lanzuisi2015}. Harder X-ray ($>$10 keV) surveys, which are much less prone to obscuration, can yield the least biased samples of Compton-thick AGN compared to any other wavelegth. The {\it Swift}-BAT (Burst Alert Telescope; \citealt{barthelmy2000}) all-sky survey detected a number of heavily obscured AGN at bright fluxes, $\rm f_{14-195keV}\sim 10^{-11}$ $\rm erg~cm^{-2}~s^{-1}$ \citep{burlon2011, ajello2012, ricci2015} arising from 5-7\% of the BAT AGN population. The BAT cannot probe much deeper fluxes because it is a coded-mask detector and thus its spatial resolution is limited. The recently launched {\it NuSTAR} mission is carrying the first telescope operating at energies above 10 keV and therefore it can reach a flux limit two orders of magnitude deeper than {\it Swift}-BAT before it encounters the confusion limit at about $\rm f_{8-24keV}\sim10^{-14} erg~cm^{-2}~s^{-1}$. The {\it NuSTAR} surveys of the COSMOS and the e-CDFS surveys (\citealt{civano2015} and \citealt{mullaney2015}, respectively) could yield the first examples of Compton-thick AGN at faint fluxes. However, so far only a few bona fide Compton-thick sources have been detected by {\it NuSTAR} owing to its small field of view. Larger numbers will become available when a large number of serendipitous sources have been accumulated. Despite the scarcity of Compton-thick AGN even in the hard X-ray band, there are two arguments that support the necessity for a large number of these sources.The first argument is the comparison of the X-ray luminosity function with the number density of SMBH in the local Universe first proposed by \citet{soltan1982}. This suggests that a fraction of the SMBH density found in the local Universe cannot be explained by the X-ray luminosity function \citep{merloni2007, ueda2014, comastri2015}. An explanation for this disagreement is that the accretion is heavily obscured. The second argument has to do with the spectrum of the integrated X-ray light in the Universe, the X-ray background. The X-ray background is mainly due to the X-ray emission from SMBH, but unlike the luminosity function, which is derived from the observed sources, it incorporates the emission from heavily obscured AGN most of which are too faint to be detected even in the deepest X-ray surveys. A number of models have been developed to reconstruct the spectrum of the X-ray background \citep{comastri1995, gilli2007, treister2009, ballantyne2011, akylas2012, ueda2014}. All these models require a substantial number of Compton-thick AGN to reproduce the peak of the spectrum between 20 and 30 keV \citep{marshall1980, gruber1999, revnivtsev2003, frontera2007, Ajello2008_xrb, moretti2009, turler2010}. However, the exact number is still unconstrained with the various models predicting a fraction of Compton-thick AGN between 10 and 35\% of the total AGN population. The most recent X-ray background synthesis models \citep{treister2009, akylas2012} use the number density of Compton-thick AGN found in the local Universe by {\it Swift}/BAT as a calibration. It is therefore important to determine this number precisely. In this paper, we make use of the 70-month {\it Swift}-BAT catalogue in combination with the {\it Swift}-XRT, X-ray Telescope (\citealt{burrows2005}) to estimate accurate absorbing column densities for all AGN detected in the local Universe in the 14-195 keV energy band. Parallel to our work, \citet{ricci2015} used exactly the same sample to search for Compton-thick AGN. The present work extends their analysis as we make use of Bayesian statistics to estimate the probability distribution of a source being Compton thick. In addition, using the above Bayesian approach we derive the accurate number count distribution comparing with our X-ray background synthesis models. This comparison derives the {intrinsic} number of Compton-thick AGN beyond the flux limit of the BAT survey. Finally, we derive the first luminosity function of Compton-thick AGN in the local Universe.

We explore the X-ray spectral properties of AGN selected from the 70-month {\it Swift}-BAT all-sky survey in the 14-195 keV band to constrain the number of Compton-thick sources in the local universe. We combine the BAT with the XRT data (0.3-10keV) at softer energies adopting a Bayesian approach to fit the data using Markov chains. This allows us to consider all sources as potential Compton-thick candidates at a certain level of probability. The probability ranges from 0.03 for marginally Compton-thick sources to 1 for the bona fide Compton-thick cases. The important characteristic of this approach is that intermediate sources, i.e. sources whose column densities lie on the Compton-thick boundary, are assigned a certain weight based on a solid statistical basis. Based on our analysis, 53 sources in the {\it Swift}-BAT catalogue present a non-zero probability of being Compton-thick corresponding to 40 `effective' Compton-thick sources. These sources represent $\sim$7\% of the sample in reasonable agreement with the figures quoted in \citet{ricci2015} and \citet{burlon2011}. We use the same approach to derive the Compton-thick luminosity function in the 20-40 keV band. This can be represented by a double power law with a break luminosity at $L_\star \approx1.4 \times 10^{42}$ $\rm erg~s^{-1}$. The Compton-thick AGN contribute 17\% of the total AGN emissivity in the 20-40 keV band where the X-ray background energy density peaks. We compare this logN-logS with our X-ray background synthesis models \citep{akylas2012}. The main aim of this comparison is to constrain the {intrinsic} fraction of Compton-thick AGN. In all X-ray background synthesis models, there is a close dependence of the fraction of Compton-thick AGN on the amount of reflected emission close to the nucleus. Assuming 5$\%$ reflected emission, we find that the Compton-thick fraction is $\sim$15\% of the obscured AGN population (or 12\% of the total AGN population). Alternatively, a 30\% Compton-thick AGN fraction (with no reflected emission) provides an equally good fit to the 14-195 keV number counts. This can be considered as the upper limit on the fraction of Compton-thick AGN. In addition, we compare the above models with the number count distribution in the 2-10 keV band as this band is more sensitive to the amount of reflected emission. Therefore, this comparison could help us to break the degeneracy between the amount of reflected emission and the fraction of Compton-thick AGN. We compare our model with the {\it XMM-Newton} COSMOS field results by \citet{lanzuisi2015}. A 12\% Compton-thick fraction (among the total AGN population) with 5\% reflection provides a good fit to the data, while the 30\% Compton-thick fraction model falls well below the data. Instead, a model with a 50\% Compton-thick AGN fraction would be in agreement with the 2-10keV number counts. An alternative possibility is that there is evolution in the number of Compton-thick AGN between z$\sim$0 and z$\sim$1.1 (the average redshift) of the COSMOS Compton-thick AGN. Such a strong evolution of the number of Compton-thick AGN is along the lines of the luminosity function models of \citet{ueda2014}. Most X-ray background synthesis models involve Compton-thick AGN with intrinsic luminosities of the order $\rm L_{2-10keV})>10^{42}$ $\rm erg~s^{-1}$. However, it is likely that there is a large number of Compton-thick AGN which are too faint and remain undetected even in the deepest {\it Chandra} surveys. This is the often called ``bottom of the barrel'' of Compton-thick AGN. For example, \citet{risaliti1999} found that optically [OIII] selected Compton-thick AGN form at least 50\% of the obscured AGN population. These AGN may not contribute significantly to the spectrum of the X-ray background owing to their faint luminosities. However, these AGN could form a substantial fraction of the black hole mass density in the Universe \citep{comastri2015}.

1607.08069

1607

1607.01181_arXiv.txt

{ The central regions of galaxies show the presence of super massive black holes and/or very dense stellar clusters. Both objects seem to follow similar host-galaxy correlations, suggesting that they are members of the same family of Compact Massive Objects. We investigate here a huge data collection of Compact Massive Objects properties to correlate them with absolute magnitude, velocity dispersion and mass of their host galaxies.

Various studies suggest that massive galaxies, both elliptical and spiral, harbor a SuperMassive Black Hole (SMBH) in their centers, with masses between $10^6 - 10^9$ M$_\odot$. Galaxies across the entire Hubble sequence also show the presence of massive and compact stellar clusters referred to Nuclear Star Clusters (NSCs). In elliptical galaxies, the NSCs are also referred to as Resolved Stellar Nuclei (RSN). Despite their different morphologies, some galaxies of the Local Group present both a NSC and a SMBH, and in such case, the SMBH is surrounded by the NSC. Those objects follow a similar host-galaxy correlation, suggesting that they are members of the same family of Compact Massive Objects (CMOs). CMOs constitute a interplay between either SMBH or a compact stellar structure (NSCs, nuclear stellar Disks or resolved stellar nuclei). Studies in the literature also showed that the NSC mass versus the host galaxy velocity dispersion ($\sigma$) relation is roughly the same observed for SMBHs. Graham (2012) claimed, instead, that the $M_{NSC}-\sigma$ relation is shallower for NSCs ($M_{NSC} \propto \sigma^{1.5}$) than for SMBHs. Here we investigate these scaling correlations for a set of NSC and SMBH data wider than what already studied in the literature

1607.01181

1607

1607.06768_arXiv.txt

We present the rationale for and the observational description of ASPECS: The ALMA SPECtroscopic Survey in the {\it Hubble} Ultra--Deep Field (UDF), the cosmological deep field that has the deepest multi--wavelength data available. Our overarching goal is to obtain an unbiased census of molecular gas and dust continuum emission in high--redshift (z$>$0.5) galaxies. The $\sim$1$'$ region covered within the UDF was chosen to overlap with the deepest available imaging from {\it HST}. Our ALMA observations consist of full frequency scans in band~3 (84--115\,GHz) and band~6 (212--272\,GHz) at approximately uniform line sensitivity ($L'_{\rm CO}\sim$2\,$\times$10$^{9}$\,\Kkmspc), and continuum noise levels of 3.8\,$\mu$Jy\,beam$^{-1}$ and 12.7\,$\mu$Jy\,beam$^{-1}$, respectively. The molecular surveys cover the different rotational transitions of the CO molecule, leading to essentially full redshift coverage. The \Cii{} emission line is also covered at redshifts $6.0<z<8.0$. We present a customized algorithm to identify line candidates in the molecular line scans, and quantify our ability to recover artificial sources from our data. Based on whether multiple CO lines are detected, and whether optical spectroscopic redshifts as well as optical counterparts exist, we constrain the most likely line identification. We report 10 (11) CO line candidates in the 3\,mm (1\,mm) band, and our statistical analysis shows that $<$4 of these (in each band) are likely spurious. Less than 1/3 of the total CO flux in the low--J CO line candidates are from sources that are not associated with an optical/NIR counterpart. We also present continuum maps of both the band~3 and band~6 observations. The data presented here form the basis of a number of dedicated studies that are presented in subsequent papers.

Characterizing the molecular gas content of distant galaxies is essential in order to understand the evolution of the cosmic star formation rate density \citep{madau14}, and the build--up of stellar mass \citep{bell03} throughout cosmic time \citep{carilli13}. A unique way to fully characterize the molecular gas content in galaxies in the early universe is through spectral line scans in well--studied cosmological deep fields. In comparison to targeted observations of individual galaxies, spectral scans have the advantage that molecular gas reservoirs can be characterized without pre--selection through other information (e.g., stellar mass, star--formation rate). Such spectral line scans can also potentially reveal the presence of gas--rich `dark' galaxies, i.e., galaxies that are invisible in the optical wavebands, and that would not be selected as targets to search for molecular gas emission \citep[e.g.,][]{walter12}. In a sense, spectral line scans follow the spirit of the original {\it HST} deep fields \citep[e.g.,][]{williams96,beckwith06}, as essentially no prior knowledge/selection based on galaxy properties enters the choice of field. As the main constituent of the molecular gas in galaxies, molecular hydrogen (H$_2$), is too weak to be detected, the next most abundant tracer is typically used to measure the molecular gas content: $^{12}$CO (hereafter: CO). Although this molecule is 10$^4$ times less abundant, the line can be detected in various environments. As a consequence, this molecule has been used at low and high redshift to measure gas masses and kinematics. The CO line emission is observed in various rotational transitions in galaxies \citep[e.g.,][]{carilli13}. The rotational ground--state (J=1--0) of CO is at 115.271\,GHz, and the higher rotational states (J$>$1) are approximately equally spaced by that frequency\footnote{In reality, the spacing changes slightly as the dipole moment changes for the higher transitions as a result of centrifugal forces.}. The amount of high--J emission depends on the {\em a priori} unknown excitation of the molecular gas. Nevertheless, full frequency scans in the lowest frequency ALMA bands cover CO emission at essentially all redshifts (see Fig.~\ref{fig_z_range}). We here present the rationale for and the observational description of ASPECS: The ALMA SPECtroscopic Survey in the {\it Hubble} Ultra--Deep Field (UDF). This paper is structured as follows: Sec.~2 summarizes our field choices, as well as the observations and data reduction. In Sec.~3 we describe our methodology to identify line candidates in our data cubes, and present the continuum maps of both the band~3 and band~6 observations. In Sec.~4 we compare our findings to simple expectations based on previous multi--wavelength analysis of the galaxies in the field. We present our summary in Sec.~5. A number of accompanying papers build on the data presented in this paper (hereafter: {\it Paper~I}). In {\it Paper~II} (Aravena et al.~2016a) we analyse the continuum information (mostly based on the band~6 observations); in {\it Paper III} (Decarli et al.~2016a) we discuss the implications for CO luminosity functions and the redshift evolution of the cosmic molecular gas density; in {\it Paper~IV} (Decarli et al.~2016b) we examine the properties of those galaxies in the UDF that show bright CO emission; in {\it Paper~V} (Aravena et al.~2016b) we search for \Cii{} emitters; in {\it Paper~VI} (Bouwens et al.~2016) we investigate where high--redshift galaxies from ASPECS lie in relation to known IRX--$\beta$ and IRX--stellar mass relationships, and finally, in {\it Paper~VII} (Carilli et al.\ 2016) we describe implications on intensity mapping experiments. Throughout the paper we assume a standard cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\rm m}=0.3$ and $\Omega_{\Lambda}=0.7$, broadly in agreement with the most recent {\em Planck} measurements \citep{planck15}. Where required, we refer to the AB photometric system (Oke \& Gunn 1973) for the magnitude definitions and to Chabrier (2003) for the stellar initial mass function. \begin{figure} \includegraphics[width=0.99\columnwidth]{z_range.eps}\\ \caption{CO and \Cii{} redshift coverage of our molecular line scans at 1mm and 3mm. See Table \ref{tab_z_range} for the exact redshift ranges of each transition. The 1mm+3mm synergy provides continuous CO redshift coverage at virtually any redshift, with only a tiny gap at $0.6309<z<0.6950$. The \Cii{} emission line is covered in the redshift range 6$<$z$<$8 and is discussed in {\it Paper~V}} \label{fig_z_range} \end{figure} \begin{figure*} \includegraphics[width=0.99\columnwidth]{fig_freq3_coverage.eps} \includegraphics[width=0.99\columnwidth]{fig_freq1_coverage.eps}\\ \caption{RMS noise as a function of frequency in the 3mm ({\em left}) and 1mm ({\em right}) scans. At 3mm, each channel is 19.5\,MHz wide (five of the native channels), corresponding to 70\,\kms{} at 84 GHz, and 51\,\kms{} at 115 GHz. The original frequency settings (A-H) are labeled in the bottom panel, together with the frequency blocks (a-k) used in the data reduction. At 1mm, the channels are 31.3\.MHz wide (four of the native channels), corresponding to 44\,\kms{} at 212\,GHz, and to 34\,\kms{} at 272\,GHz. To first order, we reach uniform sensitivity as a function of frequency in both bands. The increase in noise towards high frequencies ($>$113\,GHz) in band 3 is due to the atmosphere (O$_2$).} \label{fig_noise} \end{figure*} \begin{table} \caption{\rm Lines and corresponding redshift ranges covered in the molecular line scans. For the 3mm data, comoving volume and volume--weighted average redshifts are computed within the primary beam, accounting for its frequency dependence. For the 1mm data, the area is fixed ($3700$ arcsec$^{2}$, as set by the size of the final mosaic).} \label{tab_z_range} \begin{center} \begin{tabular}{cccccc} \hline Transition & $\nu_0$ & $z_{\rm min}$ & $z_{\rm max}$ & $\langle z \rangle$ & Volume \\ & [GHz] & & & & [Mpc$^3$] \\ (1) & (2) & (3) & (4) & (5) & (6) \\ \hline \multicolumn{6}{c}{band 3: 3mm (84.176--114.928 GHz)}\\ CO(1-0) & 115.271 & 0.0030 & 0.3694 & 0.2801 & 89 \\ CO(2-1) & 230.538 & 1.0059 & 1.7387 & 1.4277 & 1920 \\ CO(3-2) & 345.796 & 2.0088 & 3.1080 & 2.6129 & 3363 \\ CO(4-3) & 461.041 & 3.0115 & 4.4771 & 3.8030 & 4149 \\ CO(5-4) & 576.268 & 4.0142 & 5.8460 & 4.9933 & 4571 \\ CO(6-5) & 691.473 & 5.0166 & 7.2146 & 6.1843 & 4809 \\ CO(7-6) & 806.652 & 6.0188 & 8.5829 & 7.3750 & 4935 \\ \hline \Ci{}$_{1-0}$ & 492.161 & 3.2823 & 4.8468 & 4.1242 & 4287 \\ \Ci{}$_{2-1}$ & 809.342 & 6.0422 & 8.6148 & 7.4031 & 4936 \\ \hline \multicolumn{6}{c}{band 6: 1mm (212.032--272.001 GHz)}\\ CO(2-1) & 230.538 & 0.0000 & 0.0873 & 0.0656 & 1.4 \\ CO(3-2) & 345.796 & 0.2713 & 0.6309 & 0.4858 & 314 \\ CO(4-3) & 461.041 & 0.6950 & 1.1744 & 0.9543 & 1028 \\ CO(5-4) & 576.268 & 1.1186 & 1.7178 & 1.4297 & 1759 \\ CO(6-5) & 691.473 & 1.5422 & 2.2612 & 1.9078 & 2376 \\ CO(7-6) & 806.652 & 1.9656 & 2.8044 & 2.3859 & 2864 \\ \hline \Ci{}$_{1-0}$ & 492.161 & 0.8094 & 1.3212 & 1.0828 & 1233 \\ \Ci{}$_{2-1}$ & 809.342 & 1.9755 & 2.8171 & 2.3973 & 2875 \\ \Cii{}$_{3/2-1/2}$ &1900.548 & 5.9873 & 7.9635 & 6.9408 & 4431 \\ \hline \end{tabular} \end{center} \end{table}

We present the rationale for and the observational description of ASPECS, our complete band~3 and band~6 spectral line scan with ALMA of the {\it Hubble} Ultra--Deep Field (UDF). This field was chosen because it has the deepest multi--wavelength data available, it will remain a key cosmological deep field in the future (in particular in the era of JWST) and is easily observable by ALMA. We discuss our survey design of the full frequency scans in band~3 (84--115\,GHz) and band~6 (212--272\,GHz) and report the relevant parameters of our final dataset. Critically, ALMA allows us to reach approximately uniform depth (line sensitivity: $\sim L'_{\rm CO}\sim 2 \times 10^{9}$\,\Kkmspc) across a broad range of redshifts. The spectral line scans cover the different rotational transitions of the CO molecule at different redshifts, leading to essentially full redshift coverage. We present a customized algorithm to identify line candidates in our data. This algorithm takes varying linewidths of the possible emission lines into account. We assess the {\rm fidelity} of our line search by comparing the number of positive candidates to the respective number of negative candidates, the latter being unphysical. We also calculate the {\em completeness} of our search, by quantifying our ability to recover artificial sources in our data. We present CO spectra and {\it HST} postage stamps of the most signficant detections. Based on whether multiple CO lines are detected, and whether optical spectroscopic (either slit or grism) redshifts as well as optical/NIR counterparts exist, we give constraints on the most likely line identification of our candidates. Out of the 10 line candidates (3mm band) reported in our search (Tab.~\ref{tab_lines}), we expect $<$4 candidates to be spurious, given our statistical analysis. There are a number of line candidates at positions where no optical/NIR counterpart is present. The total CO flux of these candidates is less than 33\% of the total flux of all candidates, i.e. candidate sources without counterparts only contribute a small fraction of the total measured flux in the targeted field. We also present continuum maps of both the band~3 and band~6 observations. The observed flux distribution of the line candidates is in general agreement with the empirical expectations by \citet{dacunha13} based on SED modeling of the optical/NIR emission of galaxies in the UDF. The data presented in this paper ({\it Paper~I}) form the basis of a number of dedicated studies presented in subsequent papers: \noindent $\bullet$ In {\it Paper~II} (Aravena et al.\ 2016a) we present 1.2\,mm continuum number counts, dust properties of individual galaxies, and demonstrate that our observations recover the cosmic infrared background at the wavelengths considered. \noindent $\bullet$ In {\it Paper~III} (Decarli et al.\ 2016a) we discuss the implications for CO luminosity functions and the resulting constraints on the gas density history of the Universe. Based on our data we show that there is a sharp decrease (by a factor of $\sim$5) in the cosmic molecular gas density from redshift $\sim$ 3 to 0. \noindent $\bullet$ In {\it Paper~IV} (Decarli et al.\ 2016b) we examine the properties of those galaxies in the UDF that show bright CO emission, and discuss these also in the context of the bright optical galaxies that are not detected in CO. \noindent $\bullet$ In {\it Paper~V} (Aravena et al.\ 2016b) we search for \Cii{} emitters in previously reported Lyman--break galaxies at 6$<$z$<$8. \noindent $\bullet$ In {\it Paper VI} (Bouwens et al.\ 2016) we investigate where high--redshift galaxies from ASPECS lie in relation to known IRX--$\beta$ and IRX--stellar mass relationships, concluding that less dust continuum emission is detected in z$>$2.5 than expected (unless high dust temperatures, T$\sim$50\,K, are assumed). \noindent $\bullet$ Finally, in {\it Paper VII} (Carilli et al.\ 2016) we discuss implications on CO intensity mapping experiments, and contributions towards the emission from the cosmic microwave background. The data presented here demonstrate the unique power of ALMA spectral scans in well--studied cosmological deep fields. The current size of the survey is admittedly small, limited by the amount of time available in ALMA `early science'. More substantial spectral scan surveys with ALMA of the full UDF (and beyond) will become feasible once ALMA is fully operational.

1607.06768

1607

1607.06242_arXiv.txt

{Information on the origin of internetwork magnetic field is hidden at the smallest spatial scales.} {We try to retrieve the power spectra with certainty to the highest spatial frequencies allowed by current instrumentation.} {For this, we use 2D inversion code that were able to recover information up to the instrumental diffraction limit.} {The retrieved power spectra have shallow slopes extending further down to much smaller scales than found before. They seem not to show any power law. The observed slopes at subgranular scales agree with those obtained from recent local dynamo simulations. Small differences are found for vertical component of kinetic energy that suggest that observations suffer from an instrumental effect that is not taken into account.} {Local dynamo simulations quantitatively reproduce the observed magnetic energy power spectra on the scales of granulation down to the resolution limit of Hinode/SP, within the error bars inflicted by the method used and the instrumental effects replicated.}

Recent high-resolution observations, together with state-of-the-art MHD simulations reinforced the hypothesis, first suggested by \cite{Petrovay1993}, that the small-scale dynamo is the dominant source of magnetic field in internetwork regions. Only indirect support was offered so far. It was shown that a small-scale dynamo can be efficiently sustained in strongly stratified, compressible, and non-helical surface convection without enforced recirculation \citep{vogler07,PietarilaGraham:etal:2010}. The overall structure of the magnetic field resulting from such an action seems to agree with observations, but fall short by about a factor of 2-3 in the field strength \citep{manfred2008,Danilovic:etal:2010}. Small scale dynamo solutions that are in agreement with the field strength implied by observations require a setup that mimics a deep magnetized convection zone and account for an upward directed Poynting flux in upflow regions at the bottom boundary \citep{Rempel2014}. Additional support for a small scale dynamo comes from observations where it has been shown that the mean unsigned magnetic flux in the internetwork is not dependent on the solar cycle or location on the solar disc \citep{trujillo04,David2013}. Furthermore, the net flux imbalance in internetwork is not correlated to the surrounding network \citep{lites2011}. Demonstrating small-scale dynamo action is difficult. If the dynamo were operating in the kinematic regime, it would be sufficient to demonstrate that the magnetic power spectrum peaks at scales smaller than the kinetic power spectrum \citep{Abramenko2011}. However, in the saturated state, such as we find in most astrophysical contexts, the power in both the magnetic field and velocity moves to larger wavelengths and the peak in the magnetic power spectrum no longer needs to be at smaller wavelengths than the peak in the kinetic energy power spectrum \citep{Moll}. Nonetheless, the power spectra remain an important diagnostic, but retrieving them is in itself a challenging task. One is not only limited by instrumental effects, but also by methods and diagnostics used to obtain information on velocity and magnetic fields. Recent studies \citep[][and references therein]{Abramenko2001,Abramenko2011,Katsukawa2012,Stenflo2012} showed that the shallow magnetic energy spectrum tends to extend towards higher wavenumbers as the spatial resolution of the instrumentation improves, but the slope remains steep at the subgranular scales. The largest slope value of approximately $-1$ was fitted by \cite{Katsukawa2012}, who accounted for the Modulation Transfer Function (MTF) of the instrument. In their study, a simple deconvolution was applied, not directly to the observables, but to the derivatives - proxies for line-of-sight (LOS) velocity and magnetic flux densities. In this paper, we perform a power spectral analysis on the same Hinode/SP \citep{Lites:etal:2001,Kosugi:etal:2007} data, but to retrieve the magnetic field and the LOS velocity, we use 2D inversions \citep{Michiel:2012}. As demonstrated by \cite{Michiel:2012}, the advantage of this code over the simple deconvolution is using the information contained in the full observed spectral range, simultaneously. In this way, the 2D inversions are not only able to retrieve the information up to the instrumental diffraction limit, but also to minimize the influence of noise on the retrieved highest spatial frequencies. In \cite{Danilovic:etal:2015}, we tested the code on three different simulations that give the same level of spectropolarimetric signals as the quiet Sun observations. We showed that the inversion code behaves well when a certain combination of node positions is chosen. We demonstrated that in this case, the code is able to retrieve the overall distributions of the field strength and inclination. In this paper we concentrate on kinetic and magnetic power spectra. Again, we use the comprehensive MHD simulations to estimate how well and in what case can we recover the power spectra correctly. We also quantitatively compare our with results obtained by \cite{Katsukawa2012}. \begin{figure} \includegraphics[angle=90,width=0.95\linewidth,trim= 0cm 0cm 1.5cm 0cm,clip=true]{blapp_60G_16km_0702277_338500.eps} \includegraphics[angle=90,width=0.95\linewidth,trim= 0cm 0cm 1.5cm 0cm,clip=true]{btapp_60G_16km_0702277_338500.eps} \caption{Probability density functions (PDFs) for the longitudinal (upper panel) and transverse (lower panel) apparent magnetic flux density. PDFs from MHD simulations reduced to Hinode resolution (Sim1:black plus signs and Sim2: blue plus signs) are compared with the PDFs obtained from observations (red solid line). \label{fig:obs_pdf}} \end{figure} \begin{figure} \includegraphics[angle=90,width=0.95\linewidth,trim= 3.5cm 0cm 1cm 0cm,clip=true]{power60_vel_paper_wc_lt2_norm_brewer.eps} \includegraphics[angle=90,width=0.95\linewidth,trim= 3.5cm 0cm 1cm 0cm,clip=true]{power60_bl_inv_paper_wc_lt2_norm_brewer.eps} \includegraphics[angle=90,width=0.95\linewidth,trim= 0cm 0cm 1cm 0cm,clip=true]{power60_b_inv_paper_wc_lt2_norm_brewer.eps} \caption{Sim~2: Power spectra of the vertical component of kinetic (top panel) and magnetic energy (middle panel) and power spectra of the total magnetic energy (bottom panel). Blue lines mark the original spectra before any spatial smearing at different heights ($\log \tau =0$ solid and $-2.0$ dashed). Purple lines show the result of 2D inversions at the same optical depths. Black lines are spectra of the corresponding parameters obtained from \textit{Solarsoft} routines (see the text) before (dashed line) and after (solid line) spatial smearing. Vertical line marks the resolution limit of Hinode/SP. \label{fig:sim_power}} \end{figure} \begin{figure} \includegraphics[angle=90,width=0.95\linewidth,trim= 3.5cm 0cm 1cm 0cm,clip=true]{power_vel_sim_psfs_norm_brewer.eps} \includegraphics[angle=90,width=0.95\linewidth,trim= 3.5cm 0cm 1cm 0cm,clip=true]{power_blos_sim_psfs_norm_brewer.eps} \includegraphics[angle=90,width=0.95\linewidth,trim= 0cm 0cm 1cm 0cm,clip=true]{power_b_sim_psfs_norm_brewer.eps} \caption{Sim~2: Dependence of the inversion results on the PSF used. Power spectra of the vertical component of kinetic (top panel) and magnetic energy (middle panel) and power spectra of the total magnetic energy (bottom panel). Purple lines mark the spectra inverted with the correct PSF. Solid and dashed lines correspond to different heights, $\log \tau =0$ and $-2.0$, respectively. Blue and green lines show the result when the defocus accounted for is 4 defocus steps less or more, respectively, than the one used in degradation. \label{fig:sim_power_psfs}} \end{figure} \begin{figure} \includegraphics[angle=90,width=0.95\linewidth,trim= 3.5cm 0cm 1cm 0cm,clip=true]{power_vlos_338500_evol_norm_brewer.eps} \includegraphics[angle=90,width=0.95\linewidth,trim= 3.5cm 0cm 1cm 0cm,clip=true]{power_blos_338500_evol_norm_brewer.eps} \includegraphics[angle=90,width=0.95\linewidth,trim= 0cm 0cm 1cm 0cm,clip=true]{power_b_338500_evol_norm_brewer.eps} \caption{Sim~1: How temporal averaging influences the power spectra of the vertical component of kinetic (top panel) and magnetic energy (middle panel) and power spectra of the total magnetic energy (bottom panel). Different colors mark the power spectra obtained after integration over 0~s, 20~s, 2~min and 9~min as denoted in the legend. \label{fig:sim_power_evol}} \end{figure} \begin{figure} \includegraphics[angle=90,width=0.95\linewidth,trim= 3.5cm 0cm 1cm 0cm,clip=true]{power_vlos_10032007_wc_sim_psf32_norm_brewer.eps} \includegraphics[angle=90,width=0.95\linewidth,trim= 3.5cm 0cm 1cm 0cm,clip=true]{power_blos_10032007_wc_sim_psf32_norm_brewer.eps} \includegraphics[angle=90,width=0.95\linewidth,trim= 0cm 0cm 1cm 0cm,clip=true]{power_b_10032007_wc_sim_psf32_norm_brewer.eps} \caption{Observations: Power spectra of the vertical component of kinetic (top panel) and magnetic energy (middle panel) and power spectra of the total magnetic energy (bottom panel). Green lines mark the results of 2D inversions at different heights ($\log \tau =0$ solid and $-2.0$ dashed). Black lines are spectra of the corresponding parameters obtained from \textit{Solarsoft} routines (see the text). Purple lines mark the curves obtained from Sim~2 (purple curves in Fig.~\ref{fig:sim_power}). \label{fig:obs_power}} \end{figure}

Applying a 2D inversion code to spectra synthesized from MHD simulated snapshots showed that all atmospheric parameters can be retrieved reliably up to the diffraction limit of the telescope when all instrumental effects are taken into account properly. The power spectra are recovered to much smaller spatial scales than with any other method used before, without being affected significantly by the limited spatial resolution. Although we cannot claim that any of the spectra follow a power-law, we find much gentle slopes at subgranular scales than previous studies. The observed magnetic power spectra follow closely the power spectra obtained from the most recent local dynamo runs, however, a mismatch of the observed and simulated kinetic power spectra was still observed. The inherent sensitivity of this quantity to the instrumental properties suggests that perhaps some inaccuracies in the instrumental properties still remain. Although state-of-the-art simulations show that the effect of small scale dynamo action has its peak at the scales comparable to the resolution limit of Hinode, looking at smaller scales is of course still desirable, as it is having an optical system whose properties are well known.

1607.06242

1607

1607.00244_arXiv.txt

We present a class of left-right symmetric models where Dirac as well as Majorana mass terms of neutrinos can arise at one-loop level in a scotogenic fashion: with dark matter particles going inside the loop. We show the possibility of naturally light right handed neutrinos that can have interesting implications at neutrinoless double beta decay experiments as well as cosmology. Apart from a stable dark matter candidate stabilised by a remnant $Z_2$ symmetry, one can also have a long lived keV sterile neutrino dark matter in these models. This class of models can have very different collider signatures compared to the conventional left-right models.

Left-right symmetric model (LRSM) \cite{lrsm, lrsmpot} has been one of the most widely studied beyond standard model (BSM) scenarios in the last few decades. The model not only explains the sub-eV neutrino mass \cite{PDG, T2K, chooz, daya, reno, minos} naturally through seesaw mechanisms \cite{ti, tii0, tii, tiii} but also gives rise to low energy parity violations though spontaneous breaking of a parity preserving symmetry at high scale that can also be incorporated within grand unified theory groups like $SO(10)$ naturally. The minimal version of this model give rise to tiny neutrino masses through a combination of type I \cite{ti} and type II \cite{tii0, tii} seesaw mechanisms. Due to the structure of type I seesaw term, one usually has heavy right handed neutrinos in such a model, in order to generate sub-eV neutrino mass from type I seesaw. In TeV scale LRSM, one can have right handed neutrino mass $M_R \approx 1$ TeV (say) which can give sub-eV light neutrinos from type I seesaw if the Dirac neutrino Yukawa couplings should be fine tuned to $Y_{\nu} \leq 10^{-5.5}$. Similar fine-tuning is also involved in the type II seesaw term as we discuss in the next section. Such fine-tunings become more severe if we wish to keep one or two of the right handed neutrinos at very light scale, say between eV to keV. Such light sterile neutrinos could have very interesting implications for low energy neutrino oscillation experiments \cite{LSND1,miniboone,react,gall1,gall2}, neutrinoless double beta decay $(0\nu \beta \beta)$ \cite{kamland_zen, KamLAND-Zen:2016pfg,GERDA}, dark matter \cite{whitekev} and cosmology in general \cite{wmap9, Planck15}. The LRSM also received significant attention after the recent hints from the large hadron collider (LHC) about the possibility of TeV scale new physics: the CMS $eejj$ excess \cite{CMSeejj}, the ATLAS diboson excess \cite{diboson} and more recently, the 750 GeV diphoton excess \cite{atlasconf,lhcrun2a,CMS:2015dxe}. The first two excesses around 2 TeV can be explained within a version of LRSM where the discrete left-right symmetry (to be introduced below) gets broken at a high energy scale \cite{Chang:1983fu} whereas the gauge symmetry of LRSM remains unbroken all the way down to the TeV scale. The recent 750 GeV diphoton excess can however, be explained within LRSM framework only when it is extended with additional vector like fermions \cite{LR750GeV1, LR750GeV2, LR750GeV3, LR750GeV4} or with fields having high $SU(2)$ dimensions \cite{LR750GeV5}. The vector like fermions, as discussed in these works, can serve three purposes: (i) to explain the large diphoton cross section at 750 GeV, (ii) to provide masses to all the fermions of the standard model (SM) through a universal seesaw mechanism \cite{VLQlr, univSeesawLR} and (ii) to assist in gauge coupling unification \cite{LR750GeV3, LR750GeV4}. This model can also have interesting implications for cosmology \cite{gulr} and $0\nu \beta \beta$ \cite{lr0nu2beta}. In the present work, we study a class of left-right symmetric models to see the possibility of having light sterile neutrinos at a scale much below the scale of left-right symmetry. Typically, the right handed neutrinos in LRSM have masses around the scale of left-right symmetry. Thus, in TeV scale LRSM, one usually have right handed neutrinos in the GeV to TeV range. If we fine tune the Yukawa couplings to be as small as electron Yukawa, then we can get right handed neutrinos of the order of tens of MeV. We therefore study some non-minimal versions of LRSM where one can have light right handed neutrinos in the eV-keV regime even with order one Yukawa couplings. Starting with the LRSM having universal seesaw for all fermions, we explore a few other possibilities: (i) light neutrinos as Dirac fermions, (ii) both left and right handed neutrinos acquire Majorana mass terms through radiative type II seesaw and (iii) left handed neutrinos acquire masses through radiative type I seesaw. In all of these possibilities, we consider additional discrete symmetries to stabilise the particles going inside the loops so that the lightest neutral particle among them can be a cold dark matter (CDM) candidate. This follows the basic idea of scotogenic models first proposed by \cite{m06}. Apart from the stable cold dark matter candidate, these models can also have a long lived keV scale right handed neutrino that can be a warm dark matter (WDM) candidate. Such keV scale sterile neutrino dark matter within minimal LRSM was discussed by \cite{wdmlr1, wdmlr2} and a radiative neutrino mass model with both cold and warm dark matter components have been proposed recently by \cite{abm16}. We also discuss the possible implications of these scenarios in $0\nu \beta \beta$, collider experiments and cosmology. In all the models we discuss in this work, there exists additional vector like fermions which are singlets under the left-right gauge symmetry. Since they are singlets, their bare mass terms can be written in the Lagrangian and the symmetry of the theory does not restrict their masses to the TeV scale. In a generic theory, one expects these bare mass terms to be close to the grand unified theory (GUT) scale or the Planck scale. This is in contrast with the minimal LRSM where all the fermions acquire masses as a result of spontaneous gauge symmetry breaking. One can however introduce a discrete $Z_2$ symmetry forbidding the bare mass terms of vector like fermions in a way that was shown by the authors of \cite{LR750GeV1}. Another singlet scalar with appropriate $Z_2$ charge can be introduced in such a way to give rise to a Yukawa term for the vector like fermions. This singlet scalar can then acquire a non-zero vacuum expectation value (vev) and explain the masses of the vector like fermions in a dynamical manner. Also, the fermion content of minimal LRSM can be accommodated within the spinor representation $\textbf{16}$ of $SO(10)$ GUT models which is not possible in the present class of models. Although the present class of models can give rise to gauge coupling unification \cite{LR750GeV3, LR750GeV4}, the GUT group should be bigger than the minimal $SO(10)$ to accommodate the vector like fermions. This article is organised as follows. In section \ref{sec1}, we briefly discuss the minimal LRSM and then discuss the LRSM with universal seesaw in section \ref{sec2}. In section \ref{sec3} we discuss different possible versions of scotogenic LRSM where neutrinos acquire masses at one loop level with dark matter particles going inside the loops. In section \ref{sec4}, we discuss the possibilities of light sterile or right handed neutrinos in these models followed by their implications in $0\nu \beta \beta$ and colliders in section \ref{sec5} and \ref{sec6} respectively. We then comment upon cosmological implications of these scenarios in section \ref{sec7} and finally conclude in section \ref{sec8}.

\label{sec8} We have presented a class of left-right symmetric models without the conventional Higgs bidoublet that allows the possibility of having light sterile neutrinos along with a cold dark matter candidate stabilised by additional discrete symmetries. The charged leptons acquire masses through a universal seesaw mechanism due to the presence of additional vector like fermions. The neutrinos can either acquire masses through the same universal seesaw or through one loop radiative corrections, if suitable particles are added with non-trivial transformations under additional symmetries. The same additional symmetry can also stabilise one of the particles going inside the loop in neutrino mass diagram, resulting in a stable dark matter candidate. Depending on the additional particles and symmetries, the light neutrinos can be either Dirac or Majorana. In case of Dirac light neutrinos, the dark sector is composed of only cold dark matter component. However, in case of Majorana light neutrinos, the one or more of the right handed neutrinos can have mass in the eV-keV range without any unnatural fine tuning of the Yukawa couplings. A keV sterile neutrino can be a warm dark matter candidate giving rise to a mixed dark matter scenario along with the cold component. Such light right handed neutrinos can have interesting implications for neutrinoless double beta decay and collider experiments, constraining the right handed gauge boson masses. Cosmology constraints on the number of relativistic degrees of freedom also constraints the right handed gauge boson masses to be heavier than a few TeV so that the right handed neutrinos get decoupled before the QCD phase transition temperature. A detailed calculation of the relic abundance of such mixed dark matter scenario is worth investigating and we leave it for a future work. We also briefly comment upon the possible ways to get rid of domain walls that are formed due to the spontaneous breaking of discrete symmetries in the models.

1607.00244

1607

1607.05461_arXiv.txt

The fairly recent detection of a variety of anions in the Interstellar Molecular Clouds have underlined the importance of realistically modeling the processes governing their abundance. To pursue this task, our earlier calculations for the radiative electron attachment (REA) rates for C$_4$H$^-$, C$_6$H$^-$, and C$_8$H$^-$ are employed in the present work, within a broad network of other concurrent reactions, to generate the corresponding column density ratios of anion/neutral (A/N) relative abundances . The latter are then compared with those obtained in recent years from observational measurements. The calculations involved the time-dependent solutions of a large network of chemical processes over an extended time interval and included a series of runs in which the values of REA rates were repeatedly scaled over several orders of magnitude. Macroscopic parameters for the Clouds' modeling were also varied to cover a broad range of physical environments. It was found that, within the range and quality of the processes included in the present network,and selected from state-of-the-art astrophysical databases, the REA values required to match the $observed$ A/N ratios needed to be reduced by orders of magnitude for C$_4$H$^-$ case, while the same rates for C$_6$H$^-$ and C$_8$H$^-$ only needed to be scaled by much smaller factors. The results suggest that the generally proposed formation of interstellar anions by REA mechanism is overestimated by current models for the C$_4$H$^-$ case, for which is likely to be an inefficient path to formation . This path is thus providing a rather marginal contribution to the observed abundances of C$_4$H$^-$, the latter being more likely to originate from other chemical processes in the network, as we discuss in some detail in the present work. Possible physical reasons for the much smaller differences against observations found instead for the values of the (A/N) ratios in two other, longer members of the series are put forward and analyzed within the evolutionary modeling discussed in the present work.

\label{intro} Molecular negative ions in the Interstellar Medium (ISM) were first discovered as an unidentified series of lines within a radio astronomical survey of the evolved carbon star IRC+10216 by \cite{Kawaguchi1995} and were later assigned by \cite{McCarthy2006}, who confirmed the presence of C$_6$H$^-$ in TMC-1 on the basis of laboratory rotational spectroscopy. The observation was followed by the detection of both C$_6$H$^-$ and C$_4$H$^-$ in the protostellar core L1527 by \cite{Sakai2007} and by \cite{Agundez2008}. The following year \cite{Gupta2009} carried out an extensive dedicated survey of C$_6$H$^-$ in twenty four different molecular sources: they observed that anion only in two star-forming clouds and confirmed the anion-to-neutral (A/N) ratio of the order of a few percent, in keeping with the earlier measurements. Those detections had been limited to the Taurus Molecular Cloud complex until the more recent discoveries of the presence of the C$_6$H$^-$ anion in the Lupus, Cepheus and Auriga star-forming regions \cite{Sakai2010,Cordiner2011} which therefore established the widespread presence of molecular anions in different corners of the ISM. The presence of similar molecular species, the cyanopolyynes, has also been observed in the same star-forming cloud IRC-10216 via the detection of C$_3$N$^-$ by \cite{Thaddeus2008}. The possible chemical importance of molecular anions in the ISM had been put forward over the years even well before their recent detections. The work in \cite{Dalgarno1973}, in fact, first discussed this possibility suggesting the chief mechanism of formation to be the Radiative Electron Attachment (REA) path: M + e$^-\rightarrow$ M$^-$ + $h\nu$. Attempts at detection using radioastronomy were discussed by \cite{Sarre1980}, while \cite{Herbst1981} further argued that large interstellar molecules, including the polyynic chains, could efficiently undergo the anionic stabilization process via the REA process. Evolutionary models were later put forward, after the anions' observations of above, by \cite{Millar2007} and by \cite{Harada2008}, although marked discrepancies were found between the estimated A/N ratios for C$_4$H$^-$ and C$_6$H$^-$ and those suggested by the above models \cite{Agundez2008}. It has also been demonstrated later on by \cite{Cordiner2012} that anion measurements have the potential to offer additional information on the molecular properties of interstellar clouds because of the marked reactivity of these ionic species, the latter being very sensitive to the relative abundances of electrons and of C, H, N atoms in the environment. It therefore becomes important to be able to model the evolutions of the anionic abundances in different molecular clouds in order to find out which are the most important molecular mechanisms that are likely to preside over their chemical formation . In the present work we shall therefore revisit the A/N ratios recently observed experimentally in the dense interstellar medium \cite{Cordiner2013} by using the existing estimated rates of REA processes which have been obtained by calculations. We shall employ such rates for generating the A/N ratios and the latter quantities will be compared with those from astronomical observations, a comparison that will help us to evaluate what is the expected importance of REA anions' formation vis \'{a} vis other chemical routes currently present in the existing databases and routinely used within chemical evolutionary models. In our study we shall refer to a database that we here label as KIDA\footnote{\url{http://kida.obs.u-bordeaux1.fr/}}, \cite{Wakelam2015}. The following two sections briefly review the existing computational models for the REA anions' formation and further discuss our use of them in the evolutionary calculations we have employed to produce the A/N ratios over large time intervals. Such calculations involve an extended network of chemical processes that have to be included to realistically describe the environment of a dark cloud core. Our final results and our present conclusions are given in Section~\ref{sect:conclusions}.

\begin{figure} \centering \includegraphics[width=0.5\textwidth]{ratioH2.eps} \caption{Computed A/N ratios for the C$_{6}$H$^{-}$ anion following the evolutionary model described in the main text. The curves are referring to different times in the evolutionary modelling, i.e. $t_{end}=10^5$~s (squares), $t_{end}=10^6$~s (circles), and $t_{end}=10^7$~s (diamonds), and initial conditions, i.e. from \cite{Cordiner2012} (C12, dashed) and \cite{Wakelam2008} (WH08, solid). The figure reports the changes in the A/N ratio as a function of the total density over the same range discussed by \cite{Cordiner2013}.} \label{fig01} \end{figure} In the survey on A/N ratios reported by \cite{Cordiner2013} it was also found by their simulations that, in the case of the C$_{6}$H$^{-}$ anion, there was a positive correlation between the value of that ratio and the total number density. Their Fig.~5 indicated, in fact, that the ratio varied between $\sim$2.5\% and $\sim$4.4\% when the hydrogen molecule's density was varied by about one order of magnitude. Similar calculations were carried out by us using the present model with the chemical set-up previously described (see Fig.~\ref{fig01}): the general behaviour is seen to depend on the time evolution ($t_{end}$) and on the initial conditions employed. In the case of \cite{Cordiner2012}, given by the initial conditions (C12), we find a negative dependence of the A/N ratio with the total density that, as expected, converges to a constant value for long evolutions ($t_{end}=10^8$~yr). In the case of \cite{Wakelam2008}, given by the initial conditions (WH08), we note that the positive trend is found for $t_{end}=10^5$~yr, but this dependence is cancelled at later evolutionary stages ($t_{end}\geq10^6$~yr). The difference with the results found by \cite{Cordiner2012} is due to the different set-up employed by us and the more recent rates and chemical network values used within our modelling. We note here, in fact, that the initial conditions play a key role in determining the comparison with the observed quantities, especially when more metals are present (WH08 case). The effects due to the presence of the dust are not taken into account but, as mentioned in the previous section, they could modify the behaviour found in all models. We have carried out a further, and more extensive, test on the dependence of the A/N ratios on the total density (i.e. H$_2$ density) and evolution time for C$_{n}$H$^{-}$ ($n=4,6,8$). In particular, we have run simulations based on three different choices of such densities (see Fig.~\ref{fig02}): the low-density values typical of a quiescent cloud ($10^4$~cm$^{-3}$, upper panels), the middle-range densities relevant for the star-forming regions ($10^5$~cm$^{-3}$, middle panels), and the high-density values associated with protostars (lower panels, $10^6$~cm$^{-3}$). Analogously, in Fig.~\ref{fig03} we report the absolute evolutions of neutral and corresponding anion species for different densities and initial conditions that show a behaviour similar to the one found by \cite{Cordiner2012}. \begin{figure*} \centering \includegraphics[width=1.1\textwidth]{multiAll.eps} \caption{Computed A/N ratios for C$_{n}$H$^{-}$ ($n=4,6,8$), obtained by following the evolutionary model described in the main text. In each of the panels the ratios have been artificially scaled over six orders of magnitude from top to bottom, see as a reference the labels with $\log(n_{tot}/{\rm cm}^{-3})$ in the second panel, top row, where the powers of ten used in the scaling for all panels are explicitly reported.The scalings have been obtained by varying of the same factors the REA rates employed in the calculations. The A/N experimental values (with errors) are given by the grey-shadowed areas in each panel. Note that observed values are valid for TMC-1~CP, that has density of $10^{4}$~cm$^{-3}$ and an age of $10^5$~yr. The choice of the varying total densities indicated in each panel are explained in the main text when discussing this figure. The experimental values with their errors are from \cite{Cordiner2013} and \cite{Brunken2007}. Initial conditions are from \cite{Cordiner2012}. The second panel, top row, is also discussed in Fig.~\ref{fig:cfrC6H}.} \label{fig02} \end{figure*} When we examine Fig.~\ref{fig02}, we see there that the ratio for the C$_{4}$H$^{-}$ matches the experimental ratios only when the REA value, initially taken from the calculations of the CSGG or OH models and included in all the simulations of the cloud shown in the three panels of that figure, is reduced by nearly six orders of magnitude. This indicates that, within the quality of our chemical network, the REA path to the formation of that anion must make only negligible contributions to the overall rates of its formation. We further see that when the macroscopic system has reached later stages in the time of its evolution, sometime even later than the estimated cloud age, the scaling can be reduced to only about three orders of magnitude to match experimental values. Furthermore, the analysis of the evolutionary models indicates that at low-densities it is chiefly the reaction \mbox{O + C$_5$H$^-\rightarrow$ CO + C$_4$H$^-$} that supports the anions formation while there are no significant contributions to its formation coming from the REA process. It is important to note , especially at later times in the cloud's evolution, that the results of our simulations reported in the center- and right-columns of panels of Fig.~\ref{fig02} for the behaviour for the other two members of the present series of anions indicate a significant role of the REA mechanism ( from the CSGG model which we have included in our study) to match the experimental values given by the shadowed areas in all the panels. In fact, the larger anions show that the efficiency of the REA channel to their formations produced by the present CSGG theory is matching the experimental observations: they are in fact reproduced by using less than an order of magnitude reduction of the theoretical values. \noindent It is also interesting to note that the theoretical values we have employed are orders of magnitude smaller than the Langevin-type estimates usually included in the databases: for the C$_{4}$H$^{-}$ case the corresponding value reported is $1.10\times10^{-8}\left(T_{gas}/300~{\rm K}\right)^{-0.5}$. Likewise, for the cases of C$_{6}$H$^{-}$ and of C$_{8}$H$^{-}$ the Langevin estimates are also much larger than the ones provided by the CSGG theoretical modelling: $6\times10^{-8}\left(T_{gas}/300~{\rm K}\right)^{-0.5}$. It therefore follows that the values produced by the CSGG model for the REA rates for the longer chain anions are much more realistic than the Langevin rates.They further show a temperature dependence, feature which is absent in the latter approximation \cite{Carelli2013}. If we now compare the evolutions of the absolute abundances, in Fig.~\ref{fig03}, for both the neutral compounds and their anions, we note that for ages of the cloud larger than $10^6$~yr the amount of both the neutral C$_{n}$H and of their anions reaches very low concentrations with respect to the total gas density, especially for the initial conditions labelled~C12. This suggests that in the employed models the error carried by the A/N ratios after long evolutionary times could become larger than expected, since the uniformly small values of their absolute quantities naturally have a substantial effects on the errors of their ratios. It then follows that the role played by obtaining increasingly more accurate REA rates becomes also crucial in controlling the quality of their A/N ratios. Another indication of the importance of including the most realistic REA rates is provided, for the case of the C$_{6}$H$^{-}$ anion's A/N ratio, by the evolutionary comparison reported in Fig.~\ref{fig:cfrC6H}, where the low density case associated with the TMC-1~CP conditions of a quiescent cloud \cite{Cordiner2013} are also employed in our calculations. We note here that, when we use the initial conditions employed by \cite{Cordiner2012} (C12),while including the REA rates from the CSGG model \cite{Carelli2013}, the results labelled (Car13) match the observational data for ages that are compatible with that of the TMC-1~CP cloud, i.e. $\sim10^5$~yr. On the other hand, the A/N ratio given by the present modelling when using the REA rate value given by the KIDA database is larger than the observed one for nearly the whole evolution time. As mentioned in the previous cases, adding more metals to the initial conditions (i.e. using the WH08 prescription) affects the global behaviour of the A/N ratio, as clearly indicated by the upper curves of the present Fig.~\ref{fig:cfrC6H}. \begin{figure*} \centering \includegraphics[width=.9\textwidth]{cfrInitial.eps} \caption{Computed time evolutions for C$_{n}$H$^{-}$ ($n=4,6,8$) and their anions, following the evolutionary modeling described in the main text. In each panel we report the computed evolutions using different choices for the initial conditions: from \cite{Wakelam2008}, labelled as WH08, and from \cite{Cordiner2012}, labelled as C12. The labels in the third panel, top row, indicate the various values for the total cloud density and are omitted in the other panels because they are the same .} \label{fig03} \end{figure*} The following general considerations could be made by looking at all the figures' data presented thus far,also keeping in mind the effects from the most efficient chemical reactions already mentioned earlier: \begin{enumerate}[label={(\roman*)}] \item all the figures consistently show that a reduction in size of the REA rates within the modelling causes an almost linear reduction of the A/N ratios to be compared with observations, in spite of the complexity of the chemical links which are active within each of the networks that we have used in the evolutionary modeling. \item When the REA rates are reduced by more than five-six orders of magnitude, the ``humps'' shown by the time evolution curves, around time evolution values between 10$^{5}$~yr and 10$^{6}$~yr, were found to be linked to the following chemical formation channels, an effect occurring independently of the initial conditions: \begin{equation}\label{eqn:reaOC} \mO + \mC_{n+1}\mH^{-} \rightarrow \mC\mO + \mC_{n}\mH^{-}\,, \end{equation} for $n=4,6$ at low densities, and $n=4,6,8$ at high densities, while \begin{equation} \mH + \mC_{n+2}\mH^{-} \rightarrow \mC_2\mH + \mC_{n}\mH^{-}\,, \end{equation} for $n=8$ at low densities. These reactions turn out to be the main anion formation channels at the times when the REA processes become negligible. Thus, for this reason they represent the lower limits of the anion formation efficiency. Note that both at low and high densities, when REAs are still significant processes in our modelling, the attachment efficiency is always larger than reaction~(\ref{eqn:reaOC}), except for $n=4$, where reaction~(\ref{eqn:reaOC}) dominates at low densities. \item In the database employed here the reactions: \mbox{O + C$_{n}$H$^{-}$} and \mbox{O + C$_{n}$H} have different efficiencies in yielding their products, so that when the REA anion formation rates are very small the chemical kinetics of the anions is faster than the analogous chemistry for neutrals. As expected, when the REA rates are fully operative, or reduced by less than five order of magnitude, the A/N ratio increases because the REAs reduce the efficiency of the two chemical processes mentioned above. It is important to note here that these reaction rates are still given by the databases as the simple Langevin-type reactions \cite{Harada2008}, and therefore more accurate calculations, or more modern experiments, on such reactions would help to better clarify this issue. \item The uniform drops of the absolute densities shown in our model by the data in Fig.~\ref{fig:cfrC6H} for C$_{n}$H$^{-}$ after time values around 10$^{5}$~yr are due to the following contributing reactions: \begin{equation} \mO + \mC_{n}\mH^{-} \rightarrow \mC\mO + \mC_{n-1}\mH^{-}\,, \end{equation} this being true when using the C12 initial conditions, while for the different choices of WH08 we found that \begin{equation} \mC_{n}\mH^{-} + {\rm Mg}^{+} \rightarrow \mC_{n}\mH + {\rm Mg} \end{equation} \begin{equation} \mC_{n}\mH^{-} + {\rm Na}^{+} \rightarrow \mC_{n}\mH + {\rm Na} \end{equation} \begin{equation} \mC_{n}\mH^{-} + {\rm Fe}^{+} \rightarrow \mC_{n}\mH + {\rm Fe} \end{equation} become an important chemical paths to anion destruction, since more metals are present in the selected cloud's conditions. This result is similar to what we have recently found on CN$^{-}$ formation \cite{Satta2015}, where we have shown that when the electron-driven processes become slow and ineffective in forming anions, then other chemical reactions take over and become dominant with respect to the REA mechanism. Furthermore, this inefficiency in the case of the CN$^{-}$ formation was also shown by the accurate calculations on the REA path for the same molecule discussed in \cite{Douguet2015}. \item When looking at the C$_{4}$H$^{-}$/C$_{4}$H evolution over time, we see that to match the A/N experimental value \cite{Cordiner2013} requires an electron attachment REA rate which should be $\leq10^{-12}$~cm$^{3}$~s$^{-1}$. This value, at 10~K, is at least four orders of magnitude smaller than that estimated by the existing calculations with the CSGG modelling \cite{Carelli2013}, which we have employed here, and even smaller than the larger value from the HO modelling \cite{Herbst2008}, which was also around 10$^{-8}$~cm$^{3}$s$^{-1}$. On the other hand, the more recent results from the DFRDOK model \cite{Douguet2015} include an approximate form of non-adiabatic coupling between the impinging electrons and the molecular vibrations \cite{Douguet2015}.They found for this system an REA rate of formation around $10^{-16}$~cm$^{3}$~s$^{-1}$ at 30~K. In other words, both the latest modelling for the REA process, and the present modelling that examines several scalings of that rate, suggest that the REA path to C$_{4}$H$^{-}$ formation in the cores of cold molecular clouds should be a more inefficient path than previously indicated. \item In the case of the C$_{6}$H$^{-}$/C$_{6}$H ratio, the recent experiments by \cite{Cordiner2013} indicate an A/N mean value over all their sources of 3.10\%. From the present numerical experiments of Fig.~\ref{fig02}, we see that to match their estimates corresponds to employing in our modelling an REA rate of $5\times10^{-9}$~cm$^{3}$~s$^{-1}$, which is now only about one order of magnitude lower than the previously calculated values obtained by the CSGG dynamical modelling \cite{Carelli2013,Herbst2008}. Given the uncertainties that we know to exist among several of the chemical rates reported in our employed database (KIDA), one can consider the discrepancy between estimated and computed REA formation rates for this member of the series to be within the expected error bars of the present runs, where the agreement with observations is bested around $10^{5}$~yr. \item The results of Fig.~\ref{fig02} for the A/N ratio involving the longest member of the present series, the C$_{8}$H$^{-}$/C$_{8}$H ratio, further indicate that the experimental values between 3.8\% and 5\% is accurately given when using an REA attachment rate, at 10~K, of about $10^{-8}$~cm$^{3}$~s$^{-1}$. The CSGG computational estimate used in our present study , at 10~K, is \mbox{$3\times10^{-7}$~cm$^{3}$~s$^{-1}$}, i.e. only about one order of magnitude larger. This indicates once more that the quantum dynamical models of reference \cite{Carelli2013} become increasingly more reliable for the larger members of the series, while overestimating REA rates for the smallest members. \end{enumerate} \begin{figure} \centering \includegraphics[width=.47\textwidth]{cfr_C6H.eps} \caption{Computed time evolutions of the A/N ratio for C$_6$H, using different initial conditions, i.e. \cite{Cordiner2012} (C12, lines with squares) and \cite{Wakelam2008} (WH08, lines without squares), and also employing different electron attachment (REA) rate coefficients. The latter are from \cite{Carelli2013} (Car13, dashed lines) and from the KIDA database (KIDA, solid lines). The observational values from \cite{Cordiner2013} for TMC-1~CP are also reported (dash-dotted line) in the figure, together with their errors (grey-shadowed area).} \label{fig:cfrC6H} \end{figure} In all cases, the large A/N ratio's drops in value as time increases aredue to the different chemical efficacy of the reactions \mbox{O + C$_{n}$H$^{-}$} and \mbox{O + C$_{n}$H}, where the latter generally exhibits smaller rate coefficients. The slower chemical processes reduce the abundance of the neutrals less efficiently than that of the anions, hence causing the latter to disappear in the dark clouds more rapidly than the neutrals. This behaviour induces a reduction of their A/N ratios as the evolutionary times of the clouds increase. Analogously, when the A/N ratios increase, the presence of the REA processes offsets the effects of the two reactions with atomic oxygen. They become more significant in controlling anions' formation paths. Our present evolutionary study indicates that the attachment rates existing in the current literature are uniformly larger than those which can match the experimental A/N ratios at 10~K within rather sophisticated chemical networks, as those we have used to describe the molecular cloud's evolutions in time. Finally, we should also note here that we have tested the robustness of our evolutionary model by running it at at least one higher temperature, i.e. at 30~K, and found essentially no changes in all the parameters produced by the model. Our earlier scattering calculations using the CSGG model \cite{Carelli2013}, which we have employed here within a fairly large chemical network of reactions, had already suggested that the computed REA values using that model were to be considered as upper bounds to the true rates, since no estimates for the autodetachment channels were included explicitely. The present scaling experiments on the A/N evolution in dark molecular clouds confirm that, for the case of the smallest member of the series examined, the actual REA rates have to be smaller than those suggested by calculations: for C$_{4}$H$^{-}$) the rates required to match experimental A/N ratios should be about four orders of magnitude smaller. On the other hand, as the length of the carbon chain increases, e.g. for C$_{6}$H$^{-}$ and C$_{8}$H$^{-}$, the CSGG model calculations indicate that, to obtain realistic A/N ratios within the evolutionary model of a dark molecular clouds, our calculated REA rates are only slightly larger: given the general uncertainties for many of the chemical rates contained in the chosen database, the calculated electron attachment rates are closer to reality for the latter molecules than for the case of the C$_{4}$H$^{-}$ anion. As a possible explanation for these differences, it is worth mentioning at this point that C$_{4}$H is the only partner molecule in the series which does not possess a critical value for its dipole moment \cite{Carelli2013} and therefore would not have the dipole-driven metastable states which can contribute to a more efficient REA mechanism, as we have explained earlier in the Introduction. Since the REA rates are expected to increase with molecular complexity (e.g. for larger polyynes), it seems reasonable to surmise that rates obtained using the CSGG model for the longer C-bearing chains of these molecular systems \cite{Carelli2013} would be likely to match even better future possible observations of A/N ratios in such longer chains.

1607.05461

1607

1607.00628_arXiv.txt

{Pairs of extrasolar giant planets in a mean motion commensurability are common with 2:1 resonance occurring most frequently. Disc-planet interaction provides a mechanism for their origin. However, the time scale on which this could operate in particular cases is unclear. We perform 2D and 3D numerical simulations of pairs of giant planets in a protoplanetary disc as they form and maintain a mean motion commensurability. We consider systems with current parameters similar to those of HD 155358, 24 Sextantis and HD 60532, and disc models of varying mass, decreasing mass corresponding to increasing age. % For the lowest mass discs, systems with planets in the Jovian mass range migrate inwards maintaining a 2:1 commensurability. Systems with the inner planet currently at around 1 au from the central star could have originated at a few au and migrated inwards on a time scale comparable to protoplanetary disc lifetimes. Systems of larger mass planets such as HD 60532 attain 3:1 resonance as observed. For a given mass accretion rate, results are insensitive to the disc model for the range of viscosity prescriptions adopted, there being good agreement between 2D and 3D simulations. % However, in a higher mass disc a pair of Jovian mass planets passes through 2:1 resonance before attaining a temporary phase lasting a few thousand orbits in an unstable 5:3 resonance prior to undergoing a scattering. Thus finding systems in this commensurability is unlikely.}

\label{intro} Pairs of extrasolar giant planets in a mean motion commensurability are a common occurrence. It has been estimated that sixth of multi-planet systems detected by the radial velocity technique are in or close to a 2:1 commensurability (Wright et al. 2011). Parameters for some cases of interest that are considered in this paper are shown in table \ref{table1} (for additional examples see e.g. Emelyanenko 2012). In addition there are two known systems in 3:2 resonance { (Correia et al. 2009, Rein et al. 2012, Robertson et al. 2012a)} and one in a 4:3 resonance { (Johnson et al. 2011, Rein et al. 2012)} consistent with the systems in 2:1 resonance being the most commonly observed commensurability. The existence of these resonant systems indicates that dissipative mechanisms that result in changes to planet semi- major axes that produce related changes to period ratios in planetary systems have operated. This is because the probability of forming resonant configurations in situ is expected to be small (e.g. Beauge et al. 2012). Disc-planet interaction can produce the required evolution of the semi-major axes. This may result in convergent migration leading to the formation of a commensurability (see Baruteau et al. 2014 and references therein). Accordingly understanding the observed configuration of such systems has the potential for either revealing how disc-planet interactions may have operated or for ruling them out. Previous { numerical} studies of commensurabilities forming and evolving as a result of disc-planet interactions have focused on systems such as GJ 876, HD 45364 and HD 6805 { interacting with disc modelled with either constant kinematic viscosity, or with the $\alpha$-viscosity parameter of Shakura \& Sunyaev (1973) taken to be constant} (for a review see Baruteau et al. 2014 and references therein). In this paper we extend such studies, considering systems with { orbital} parameters similar to those of HD 155358 { (Robertson et al. 2012b)}, 24 Sextantis {(Johnson et al. 2011)}, HD 60532 { (Laskar \& Correia 2009)} { as well as HD 6805 (Tifonov et al. 2014) each of} which have the inner planet with semi-major axis in the range $0.5-1.4$ au. We perform 2D and 3D numerical simulations of pairs of giant planets interacting with a protoplanetary disc that attain a mean motion commensurability for up to $2.5 \times 10^4$ orbital periods of the inner planet. We investigate whether such systems could have originated at larger radii beyond the ice line and then migrated inwards, the commensurability possibly being formed in the same neighbourhood. \begin{table} \begin{tabular}{c c c c c c} \hline\hline System &$ M_* $&$M_1 $ & $M_2$ &$ a_1 $ &$a_2$\\ HD 155358&$0.92$&$0.85\pm0.05$&$0.82\pm0.07 $&$0.64$&$1.02$\\ 24 Sextantis&$1.54$& $1.99\pm 0.4$& $ 0.86\pm 0.4$ &$1.33$&$2.08$\\ HD 6805&$1.7$& $2.5\pm 0.2$&$3.3\pm0.2$&$1.27$&$1.93$\\%e=0.13,0.1\pm0.05,0.06 HD 60532&1.44 &$3.15$&$7.46$&$0.77$&$1.58$\\ \hline \end{tabular} \caption{Properties of the HD 155358, 24 Sextantis, HD 6805 and HD 60532 systems. The first three either are or possibly in 2:1 resonance while the fourth is in a 3:1 resonance. The first column identifies the system, the second column gives the mass of the central star in solar masses. The third and fourth columns give the masses of the planets in Jupiter masses and the fifth and sixth columns give their semi-major axes in au.} \label{table1} \end{table} In order to study the role of the nature of the underlying disc model we consider models with an inner MRI active region producing a significant effective viscosity and an outer inactive region for which a significant effective viscosity may occur only in the upper layers of the disc (Gammie 1996) as well as models with a uniform $\alpha$-viscosity prescription throughout. We also consider models with different surface density scaling corresponding to varying the total disc mass or equivalently the steady state accretion rate. In this way the disc-planet interaction at different stages of the life of the protoplanetary disc can be studied with lower mass discs corresponding to later stages (e.g. Calvet et al. 2004). We find that when low mass disc models are considered, systems with planets in the Jovian mass range maintain a 2:1 commensurability while undergoing inward type II migration. This is found to be at a rate such that formation at a few au from the central star and migration to their current locations on a time scale comparable to the expected protoplanetary disc lifetime is possible in principle. We find that there is a relative insensitivity of results to the disc model employed and find good agreement between 2D and 3D simulations. Planets containing larger masses such as the HD 60532 system which is observed to be in 3:1 resonance {(Laskar \& Correia 2009)} are found to attain this resonance in low mass low viscosity discs. We find that for systems with planets in the Jovian mass range, increasing the disc mass results in the formation of an unstable 5:3 resonance. This instability results in the rapid destruction of the commensurability implying that the occurrence of such systems should be less common. The plan of this paper is as follows. We give the basic equations and coordinate system used in Section \ref{Beq}. In Section \ref{Numsims} we outline the numerical methods and computational domains adopted going on to describe aspects of the physical set up and disc models used in Sections \ref{discmod} and \ref{layered}. We then indicate how results might be scaled to different radii and summarise important aspects of type II migration in Sections \ref{scaling} and \ref{typeiimig}. We go on to describe our numerical results in Section \ref{Numres}, {beginning with a comparison with previous results for two migrating planets presented in section \ref{CompPreviousResults}. Finally, we discuss our conclusions in Section \ref{Disc}.}

\label{Disc} \label{sec:disc} In this paper we have performed 2D and 3D simulations of pairs of giant planets that have attained a mean motion resonance in a protoplanetary disc. We considered disc models both with an inner active region and an outer inactive region with lower effective viscosity as well as disc models incorporating only one of these regions. Different magnitudes for the viscosity in both regions were considered. This was found to have only minor effects on the results as long as the surface density was scaled such as to maintain the same steady state accretion rate -- or equivalently, outward directed angular momentum flux, and this scaling did not result in the surface density becoming so small that the planet mass dominated its local neighbourhood. Disc models with a range of masses corresponding to a range of accretion rates were considered, the smallest corresponding to the late stages of the protoplanetary disc life time. Simulations were run for up to 20,000 initial orbits of the inner planet. \subsubsection*{Maintenance of a 2:1 commensurability} When the mass ratio for both planets was $10^{-3}$, a 2:1 commensurability was maintained for small enough disc masses. For our standard case with a low mass disc and a corresponding steady state accretion rate of $6\times 10^{-10}\,M_{\odot}\, \mbox{yr}^{-1},$ the inward migration time $-r/{\dot r} \sim 1.5\times 10^5$ inner planet initial orbital periods is a characteristic viscous time scale and so corresponds to standard type II migration. Noting that observed parameters are somewhat uncertain, { those for} this model may approximately correspond to those for HD 155358 and 24 Sextantis (see table \ref{table1}) if the inner planet semi-major axis is respectively taken to be 0.64 au and 1.33 au { respectively.} The inward migration times are then respectively $\sim 0.77\times 10^5\, \mbox{yr}$ and $ 2.3\times 10^5\, \mbox{yr}$ which are relatively short compared to a characteristic protoplanetary disc life time. { For illustrative purposes, let us assume} the planets have migrated in resonance for a time, $t,$ and use the scaling procedure given in section \ref{scaling} to estimate the initial radius of the inner planet. { For} either of the two examples { we obtain} $r = (t/(1.5\times 10^5\, \mbox{yr}))^{2/3}\, \mbox{au} \sim 3.5\, \mbox{au}$ for $t= 10^6\, \mbox{yr}.$ We note that this is beyond the ice line, with location estimated at about 2.7 au from a Solar mass star (see e.g. Martin \& Livio 2013). But we emphasise that the scaling used restricts the disc model such that $\Sigma \propto r^{-2}.$ While this might be relaxed to some extent the surface density cannot exceed this projection by a large amount because a commensurability with period ratio closer to unity would be formed (see below). However, we note that the quoted eccentricity for the inner planet in HD 155358 is $0.17\pm 0.03$ (Robertson et al. 2012b). The mean value is exceeded for the standard run at $4000$ orbits after going into resonance with the implication that the planets could not have been in resonance longer than this time. If this is the case the above discussion would have to be modified to allow the planets to migrate independently from larger radii before converging on to resonance close to their final locations. This is likely to need to be considered for different possible exterior disc models and in addition the planets may have built up their masses as they went (see e.g. Tadeu dos Santos et al. 2015). These considerations are beyond the scope of this paper. Nonetheless, because the migration rates for single planets and the resonantly coupled planets are in general similar, the estimated starting radii would also be similar for disc models that are similar to those we considered. But note that the attained eccentricities depend on the eccentricity damping rates which depend on the details of the disc model (see Crida et al. 2008). For example, we found that for the same amount of relative resonant migration, the entirely inactive disc model led to smaller eccentricities while the 3D layered model led to larger eccentricities. Thus it is important to note that there is uncertainty as to how long the planets could have been in resonance. In the same context we comment that migration in the completely active disc model was slower by a factor $\sim 1.6$ compared to the standard case on account of its lower mass, that being determined so as to maintain the same steady state accretion rate as in the standard case. { Furthermore the potential importance a residual inner gaseous disc for damping the eccentricity of the inner planet and so preventing the eccentricities of both planets from continuing to increase in the later stages of the orbital evolution has been stressed by Crida et al. (2008). In addition Murray et al. (2002) indicate that a residual disc of planetesimals could produce a similar effect.} We undertook 3D simulations that incorporated consideration of the vertical structure of the disc. Both models that adopted a viscosity that was independent of $\theta$ and layered models for which a viscosity was only applied in the upper portion of the $\theta$ domain were considered. The orbital evolution that was obtained was found to be in good agreement with that obtained from corresponding 2D simulations. One effect seen in the 3D simulations that cannot be recovered from the 2D simulations is the vertical flow towards the mid-plane in the interior neighbourhood of the planet. However, because this occurs in the gap region where the density is very low, this does not lead to significant departures from the 2D results for the orbital evolution. In order to consider a system resembling the HD 6805 system, we performed a simulation identical to the standard one except that the mass of the outer planet was increased by a factor of two. This behaved like the standard case with maintenance of a 2:1 commensurability and an inward migration rate $-r/{\dot r}\sim 10^5$ inner planet initial orbital periods. Using the same scaling argument as above, the inner planet can be estimated to start at a radius being $\sim 5.5\, \mbox{au}$ if resonant migration is assumed for $t=10^6\, \mbox{yr}.$ \subsubsection*{Increasing planet mass and the formation of a 3:1 resonance } We have also studied a case with planet masses chosen to correspond to the HD 60532 system which has the larger planet mass ratios $2.2 \times 10^{-3}$ and $5.2 \times 10^{-3}$ (see table \ref{table1}). These planets are observed to be in a 3:1 resonance (Laskar \& Correia, 2009). In our simulation, the planets attained a 3:1 resonant configuration with $-r/{\dot r}\sim 6\times 10^4$ initial inner orbits. { For the purposes of an illustrative discussion, if we assume that the scaling to larger radii discussed in Section \ref{scaling} applies, we find that if the system arrived in its present location after having undergone inward migration in resonance for $10^6\, \mbox{yr}$, it should have started with the inner planet at an orbital radius of $\sim 7.4$ au. For a shorter evolution time of $4\times 10^5\, \mbox{yr}$ the corresponding starting location shifts to 4 au. However, note that as the orbital configuration obtained in the simulation is like that observed, the two planets may have only spent a relatively small time in resonance, comparable to our simulation run time of $\sim 6000$ orbits. In that case the planets could have migrated independently, starting at initial radii that did not differ by a large factor on account of the single planet migration rates being comparable. If this is the case, detection of the system in resonance would be unlikely. On the other hand only one system of this kind is currently known.} \subsubsection*{Effect of increasing disc mass} When the surface density or equivalently the steady state accretion rate was increased, the character of the migration of the resonantly coupled pairs of Jupiter mass planets changed. When it was increased by a factor of 5 the planets are found to enter a 5:3 resonance which became unstable after about 5000 orbits leading to a planet-planet scattering as was found by Lee et al. (2009) who adopted an N body approach. The unstable character of the 5:3 resonance makes the observed occurrence of such a resonances less likely than a 2:1 resonance. If pairs of planets formed at a few au and then migrated to their present locations with the inner planet being at around 1 au (such as for the HD 155358 and 24 Sextantis systems) while maintaining a 2:1 commensurability for a characteristic time comparable to the disc life time, the disc should have a low mass as might occur during the later stages of a protoplanatary disc lifetime. { We have found that a disc with significantly larger mass produce an unstable 5:3 resonance resulting in its observed occurrence being less likely. Although 3:2 resonances may be produced in other situations (eg. Rein et al. 2010 and see the end of Section \ref{IncSig} above) characteristic evolution times are again short.} { Finally consideration of systems containing a pair of giant planets with the innermost one being significantly more massive, for which the mechanism outlined by Masset \& Snelgrove (2001) may operate more efficiently should be undertaken but is beyond the scope of this paper. A recently discovered system of this kind is HD 204313 (Robertson et al. 2012a). This contains an inner planet of mass $3.55\, M_J$ with semi-major axis $3.04\, \mbox{au}$ and an outer planet of mass $1.68\, M_J$ with semi-major axis $3.93\, \mbox{au}$, the pair being in or close to a 3:2 commensurability. Accordingly, this will be the focus of a future study.}

1607.00628

1607

1607.03906_arXiv.txt

GJ 1132 b is a nearby Earth-sized exoplanet transiting an M dwarf, and is amongst the most highly characterizable small exoplanets currently known. In this paper we study the interaction of a magma ocean with a water-rich atmosphere on GJ 1132b and determine that it must have begun with more than 5 wt\% initial water in order to still retain a water-based atmosphere. We also determine the amount of \ce{O2} that can build up in the atmosphere as a result of hydrogen dissociation and loss. We find that the magma ocean absorbs at most $\sim$ 10\% of the \ce{O2} produced, whereas more than 90\% is lost to space through hydrodynamic drag. The most common outcome for GJ 1132 b from our simulations is a tenuous atmosphere dominated by \ce{O2}, although for very large initial water abundances atmospheres with several thousands of bars of \ce{O2} are possible. A substantial steam envelope would indicate either the existence of an earlier \ce{H2} envelope or low XUV flux over the system's lifetime. A steam atmosphere would also imply the continued existence of a magma ocean on GJ 1132 b. Further modeling is needed to study the evolution of \ce{CO2} or \ce{N2}-rich atmospheres on GJ 1132 b.

\label{sec:intro} With the success of the Kepler and K-2 missions and ground-based follow-up efforts of the brightest targets, significant strides have been made in understanding the size and density distribution of planets around other stars \citep[e.g.][]{Burke2015,Dressing2015b}. Planets with radii less than ~1.5 to 1.6 Earth radii and masses less than about 7 Earth masses are universally consistent with a rocky, Earth-like composition \citep{Rogers2015,Weiss2014}. However, most of these likely rocky planets have been found at very close orbital periods and are therefore significantly hotter than the Earth. Some of these planets receive orders of magnitude more stellar insolation than the Earth, and their atmospheres will be sculpted and altered by interactions with the stellar insolation, particularly the high energy extreme ultra-violet (XUV, 1-120 nm) radiation. Therefore models of atmospheric loss and evolution for close-in planets are timely. There has been substantial work done on atmospheric loss from planets in the solar system, particularly Venus \citep[e.g.,][]{Walker1981,Kasting1983,Zahnle1986,Chassefiere1996,Kulikov2006,Lichtenegger2010,Erkaev2013,Hamano13}. Several recent studies extend this type of modeling to atmospheric loss on habitable zone exoplanets with \ce{H2O}-rich atmospheres \citep{Wordsworth2013,Wordsworth2014,Ramirez2014,Tian2015,Luger2015}. \citet{Bolmont2016} have modeled water loss from the recently discovered TRAPPIST-1 system of planets around an ultracool dwarf star. Others have also studied whether or not close-in rocky exoplanets could be the residual core remnants of gas giant planets stripped of massive \ce{H2} atmospheres \citep[e.g.,][]{Lammer2009,Lopez2013,Luger2015b,OwenMohanty2016}. Many of the solar system studies have noted that preferential loss of H from steam atmospheres may lead to build up of \ce{O2} in a planet's atmosphere \citep[e.g.,][and references therein]{Kasting1995}. This is particularly a problem for Venus, where minimal \ce{O2} is observed, despite an assumed massive early loss of atmospheric water. \citet{Luger2015} applied this type of model to rocky exoplanets in the habitable zones of M and K dwarf stars, where \ce{O2} may be a biosignature mimic. In the present paper, we also study atmospheric loss and oxygen build up, but we extend previous models by including an interior model that allows for uptake of \ce{O2} by the planet's mantle. Our interior model includes both a magma ocean stage, as well as parameterized solid state convection with passive outgassing following solidification. This model is based on magma ocean thermal evolution models long used to study the Solar System terrestrial objects \citep[e.g.,][]{AM85,ET03,Lebrun13,Hamano13}. In comparison, few exoplanet models consider the solid body except as a lower boundary condition for the atmosphere. The present model is an improvement on these treatments and is the first fully coupled model of atmosphere-interior exchange of oxygen. We focus on GJ 1132b, a planet only slightly larger than the Earth (M$_{p}$ = 1.62 M$_{\oplus}$, R$_{P}$ = 1.16 R$_{\oplus}$), which was recently discovered by the MEarth ground-based transiting planet survey \citep{Berta15}. GJ 1132 is a nearby M3.5 dwarf (0.181 M$_{\odot}$) located only 12 parsecs away. The planet GJ 1132 b has an orbit of 1.6 days and at 0.0153 AU, receives $\sim$ 19 times more stellar insolation than the Earth and 10 times more than Venus. With a large relative transit depth, GJ 1132 b will be amenable to near-term follow-up both from large ground based telescopes, as well as orbiting observatories like HST and JWST. It is our goal to determine if the planet could have sustained a water or \ce{O2} rich atmosphere over its lifetime. We focus on O and H in order to be able to thoroughly explore the parameter space in a timely manner. Future models may wish to include a more detailed chemistry incorporating carbon and nitrogen-bearing species. The magma ocean stage on close-in rocky exoplanets may be extremely long-lived. Observations of these objects may present a means to test magma ocean models which are also used to study processes occuring during Solar System accretion. As such, observations of GJ 1132b and other planets like it may help us improve models for our own Solar System, in particular, models for water and \ce{O2} loss on Venus. This paper is organized as follows. Section \ref{sec:escape} discusses our atmospheric escape model and line-by-line climate model. Section \ref{sec:model} describes the planetary interior model and the coupling to the atmospheric model. Section \ref{sec:results} presents results from the coupled model, including the amount of water lost from the planet, the final \ce{O2} abundance in the atmosphere, and the mantle compositon. In section \ref{sec:discussion} we discuss some of the limitations of the model. Finally, in Section \ref{sec:predictions}, we give predictions for the atmospheric composition of GJ 1132 b.

\label{sec:discussion} \subsection{Sensitivity of loss rate to atmospheric composition}\label{sec:sensitivity} In our nominal models, we assume that atmospheric loss is energy-limited, where the loss rates are dependent on the O and H molar concentrations. We assume that energy-limited escape driven by hydrodynamic loss of H occurs until the \ce{O2} and \ce{H2O} total atmospheric pressures are equal. After this cross-over point, we assume that H must diffuse through the O background gas, at which point the hydrodynamic loss halts and O no longer escapes. However, the transition composition is uncertain because H should diffuse more readily into the upper atmosphere than O. We explore the sensitivity of our final results to this transition point in Figure \ref{fig:crossover}. Here we show results for both XUV flux models with $\chi_d = 1$ for a constant FeO abundance of 8 wt\% as a function of initial water abundance for different transition points ($X_H$ = 0.4(nominal), 0.1, and 0.001). For XUV model A, the final \ce{O2} pressure is insensitive to the transition point up to $\sim$1 wt\% of \ce{H2O}. At higher water abundances, the final \ce{O2} pressure is reduced by several orders of magnitude as the transition point drops, except at the very highest water abundance where the magma ocean persists. For XUV model B, the transition point has a strong effect on the \ce{O2} abundance for initial water abundances less than $\sim$10 wt\%. Reducing the transition abundance results in more tenuous \ce{O2} atmospheres, since more of the O can escape. \begin{figure} \begin{center} \includegraphics[scale=0.65]{loss-crossover2b.pdf} \end{center} \caption{Sensitivity of the final \ce{O2} pressure to the transition point between energy-limited loss of H (with hydrodynamic drag of O) to diffusive-loss of H through an O background gas (with no loss of O). The transition composition is given in terms of the molar abundance of H in the atmosphere, calculated from the total pressures of \ce{O2} and \ce{H2O} ($X_H = 0.4(nominal), 0.1, 0.001$). Blue lines are for XUV model A, pink lines are XUV model B. } \label{fig:crossover} \end{figure} \subsection{Loss of an earlier \ce{H2} envelope} It is possible that GJ 1132b began with an envelope dominated by \ce{H2} gas, rather than \ce{H2O}. As discussed in Section \ref{sec:primordial}, a significant mass of \ce{H2} can be lost from the planet, up to 15\% of the planet's mass over 10 Gyr. Interaction of an \ce{H2} atmosphere with mantle FeO might result in reduction of mantle FeO to Fe metal through a reaction such as: \begin{center} \ce{H2(g) + FeO <=> H2O(g) + Fe} \end{center} The forward reaction is thermodynamically unfavorable and has been shown to go nearly to completion in the reverse direction (oxidation of metal) at all temperatures and pressures on the present and early Earth \citep{Fukai84,Kuramoto1996}. In fact, these experimental studies of the iron-water reaction at high pressure have shown that hydrogen liberated from water can be sequestered into a \ce{FeH_x} metallic phase via a reaction such as: \begin{center} \ce{H2O(g) + 2Fe -> FeO + FeH_x}. \end{center} However we expect that relative to the total duration of the magma oceans, the presence of metal within the magma ocean was relatively short-lived. We therefore consider that the primary effect of an initial \ce{H2}-envelope would be to prolong the magma ocean lifetimes and reduce the loss of water and \ce{O2} from those calculated here. \subsection{Effect of \ce{CO2}} \label{sec:CO2} \ce{CO2} is a common atmospheric component that is often included in magma ocean models \citep{ET08,Lebrun13}, due to both its large abundance and its contribution to greenhouse warming. We do not consider it here in order to minimize the complexity of the model, but we will qualitatively discuss its possible effect on the evolution of GJ 1132b. The solubility of \ce{CO2} in silicate melt is much lower than that of \ce{H2O}, but it is much more soluble in metal alloy. Therefore, numerous papers on Earth-based magma ocean models have noted that \ce{CO2} will be concentrated in either the atmosphere or the core. \citet{Hirschmann2012} argues based on alloy/melt partition coefficients that a magma ocean that equilibrates with only 1 wt\% of alloy would lose at least 60\% of its total carbon to the core. However, we noted above that the presence of metal within the magma ocean was likely relatively short-lived, so unless carbon is removed during core formation, it seems likely that there should be substantial carbon remaining in the magma ocean and atmosphere. Solubility of \ce{CO2} in the mantle depends on the temperature, pressure, melt composition, and oxygen fugacity: as GJ 1132 b becomes more oxidized, \ce{CO2} should become more soluble in the melt. However, solubility relationships indicate that it is unlikely that more than about 20 - 30\% of the \ce{CO2} could be dissolved in the magma ocean \citep{Holloway1998}. \citet{Hirschmann2012} also suggests the possibility of diamond precipitation in the mid to lower mantle or a magma ocean carbon pump to lower atmospheric \ce{CO2} abundances. This would sequester carbon in the mantle where it would be available for later outgassing during a post-magma ocean state, much like water in our efficient degassing scenarios. However, while this is a possibility, it would require detailed additional modeling to evaluate. \ce{CO2} in the atmosphere will prolong the magma ocean lifetimes by additional greenhouse warming, which may enhance atmospheric loss of both water vapor and \ce{O2}. \citet{Tian2009} showed that in highly irradiated super-Earth atmospheres dissociation of \ce{CO2} can lead to both carbon and oxygen loss, with carbon escaping more rapidly due to its lower atomic weight. \citet{Wordsworth13} showed that water loss from \ce{CO2}-rich atmospheres can still be substantial, especially for planets that receive more insolation than the present day Earth, such as GJ 1132 b. While \ce{CO2} is effective at cooling the upper atmosphere, which can hinder loss in more temperate planets, a back-of-the-envelope calculation suggests that the degree of cooling from the \ce{CO_2} 15 $\mu$m non-LTE emission would still be far lower than the XUV flux received by GJ 1132 b, at least for the first Gyr. Therefore, cooling of the upper atmosphere would likely be insufficient to hinder the escape of \ce{O_2} and \ce{CO2}. Therefore, tenuous \ce{O2} atmospheres are the most likely scenarios for GJ 1132 b after loss of an \ce{H2} envelope. Non-thermal effects provide additional loss avenues as discussed below. \subsection{Non-thermal loss mechanisms} Considering non-thermal mechanisms for atmospheric escape from GJ1132b, it is safe to assume no planetary magnetic field, in analogy to Venus and as a conservative choice. Also, while GJ1132b is closer than Venus to its star in terms of bolometric irradiation and stellar wind flux, it is significantly more massive. As a result, charge exchange (e.g., $H + H^{+*} \rightarrow H^{+} + H^{*}$) and ion escape could increase the loss of hydrogen, and dissociative recombination (e.g., $O_{2}^{+} + e^{-} \rightarrow O^{*} + O^{*}$) could increase the loss of oxygen. The latter mechanism releases only $0.6 \times 10^{-18}$ J per atom, which is not enough to permit escape given the high mass of GJ1132b. The hydrogen non-thermal escape might be insignificant, by analogy to Venus (e.g. Pierrehumbert 2010), but a dedicated study is warranted. Similarly, the stellar wind of M dwarfs like GJ1132 is expected to be too tenuous to lead to significant atmospheric erosion, but no firm conclusion is possible without detailed modeling \citep{Kislyakova2013,Kulikov2006}. \subsection{Mantle convection after soldification} In order to have efficient degassing during the post-magma ocean state, the mantle of GJ 1132 b must continue to convect. However, progressive oxidation of the mantle by liberated O should lead to lower density materials at the top of the mantle, as shown in Figure \ref{fig:Fe2O3profile}. This may prevent overturn of mantle and delay the onset of solid state mantle convection, which would lead to reduced outgassing efficiencies. However, \citet{ET03} calculated the mineralogy of a solidifying magma ocean (without atmospheric oxidation), and found that the cumulate pile of the solidified magma ocean is unstable due to the partitioning of FeO into later (near surface) crystal phases. Although oxidation of FeO to \ce{FeO_{1.5}} may change the exact mineral condensation sequence, the additional oxygen should not be sufficient to counteract the density effect of enhanced FeO abundance in the upper mantle. For low FeO abundances, the density instability may be insufficient to cause mantle overturn, in which case GJ 1132 b may become stuck in a stagnant lid regime. This would mimic the low degassing efficiency model, which we have shown is only important in the case of the low XUV model B. Inefficient degassing reduces the final \ce{O2} abundance by about an order of magnitude in pressure.

1607.03906

1607

1607.01304_arXiv.txt

We report dynamical mass measurements for the components of the previously known double-lined spectroscopic subdwarfs G 006-026 B and C using the Fine Guidance Sensors (FGS) on the {\it Hubble Space Telescope}. To build the empirical mass-luminosity relation for low metallicity subdwarfs, we collect four other subdwarf systems with dynamical masses that we compare to theoretical models for various metallicities on the mass-luminosity relation. For most stars, they fall in the regions where the models predict to be low metallicity. This effort highlights the scarcity of dynamical masses for subdwarfs and that much work remains to be done to improve the mass errors and metallicity measurements of low mass subdwarfs in our Galaxy.

\label{sec:intro} Subdwarfs are Galactic fossils that presumably comprise the bulk of the halo, and are crucial touchstones of the star formation and metal enrichment histories of the Milky Way. The local paucity of subdwarfs and their intrinsic faintness make them difficult to characterize, unlike their disk counterparts. For example, there are currently only three confirmed subdwarf systems within 10 parsecs: $\mu$ Cas AB, CF UMa, and Kapteyn's Star \citep{Monteiro2006}. Many subdwarfs are detected via high proper motion star surveys (e.g.~\citealt{Ryan1991}, \citealt{Carney1994}, \citealt{Gizis1997}, \citealt{Jao2005, Jao2008}, \citealt{Lepine2007}, and \citealt{Savcheva2014}), yielding important information about their kinematics, chemical compositions, and space densities. However, their masses --- the single most important parameter for a star --- are still mostly unknown, but would be of great use in studies of the Galactic halo, globular clusters, and other old and/or low metallicity systems. There are currently only a half dozen subdwarf systems (here defined as stars having [m/H] or [Fe/H]$\leq$$-$0.5) with dynamical mass measurements. \cite{McCarthy1993} used infrared speckle interferometry and \cite{Drummond1995} used Adaptive Optics to measure the masses of $\mu$ Cas AB, which has [Fe/H]$=-$0.71 \citep{Karaali2003} and is one of the nearest subdwarf systems. \cite{Soderhjelm1999} reported masses for GJ 704 AB (99 Her), which has [Fe/H]$=-$0.58 (See Table~\ref{tbl:GJ704mH}). Recently, \cite{Horch2015} reported total masses for two subdwarf binaries (HIP 85209AB and HIP 95575AB) using the Differential Speckle Survey Instrument (DSSI) to resolve these two doubled-lined spectroscopic binaries \citep{Goldberg2002, Halbwachs2012}. We use the mass-ratio from SB2 results and total masses from \cite{Horch2015} to get individual masses. All of these results are summarized in Table~\ref{tbl:mass}. We note that HIP 81023AB and HIP 103987AB both have metallicity less than $-$0.5, but have preliminary orbital results from \cite{Horch2015}. We did not include them in our analysis. \cite{Ren2013} combined both single-lined spectroscopic data, photocentric orbital data from the Hipparcos Intermediate Astrometric Data and an empirical mass-luminosity relation to estimate masses for four additional subdwarf systems with [Fe/H]$\leq$$-$0.5 (HIP 705, HIP 39893, HIP 55022, HIP 59750 and HIP 73440). However, these masses are not measured dynamically, as they rely on an empirical mass-luminosity relation from stars with solar metallicity, so these stars are not included in our discussion. Finally, \cite{Soderhjelm1999} reported the total system mass and mass ratio of GJ 60AB with [Fe/H]$=-$0.65 \citep{Holmberg2009}. A few years later, \cite{Watson2001} reported that the system could harbor four components rather than two, with close binaries A-C and B-D, where B-D is an eclipsing binary with a period of 0.45 days. Because of the complexity of this system, we conclude that no reliable masses for individual components have been measured, so do not include this system in the discussion. Nearby galactic globular clusers (GCs) are also homes of low metallicities stars. Recently, the Clusters AgeS Experiment (CASE) project \citep[and references therein]{Kaluzny2015} has detected and measured many detached eclipsing binaries in $\omega$ Cen, 47 Tuc, M4, M55 and NGC6362. These binaries are comprised of either massive or evolved stars, and their metallicities are much lower than nearby field subdwarfs, so they are not discussed further in this paper either. Because the local known subdwarf density is far less than that for dwarfs, coupled with the fact that the multiplicity of cool subdwarfs appears to be lower than that of main sequence stars \citep{Jao2009, Lodieu2009}, it has been difficult to identify appropriate binary systems to target for dynamical mass determinations. This is particularly true for low mass M-type subdwarfs with masses less than 0.6 $M_{\odot}$. Consequently, the number of dynamical masses measured for low metallicity subdwarfs is much less than for their counterparts, the main sequence dwarfs (\citealt{Henry1999, Delfosse2000, Torres2010, Benedict2016}). Building an empirical mass-luminosity relation for low metallicity stars has proven difficult, and rigorous testing of low metallicity theoretical models remains to be done. In this manuscript, we present new dynamical masses of $\sim$0.47 and 0.44 $M_{\odot}$ for the low metallicity subdwarfs, G 006-026 B and C, adding two subdwarf masses to establish the empirical mass-luminosity relation. We also recalculate the masses for GJ 704 A and B by reassessing the resolved measurements to date. G 006-026 BC ia a double-lined spectroscopic binary with period of about 302 days \citep{Goldberg2002}. \cite{Carney1994} reported that the primary star in the system, the G-type star G 006-026 A, has a metallicity of $[m/H]= -0.88$. Later, \citet{Casagrande2011} and \cite{Holmberg2009} use photometry to determine the primary's $[Fe/H]=-0.52$ and $-0.60$, respectively. We assume that the B and C components have the same metallicity as the primary, and that they therefore meet our criterion for subdwarfs as stars having [Fe/H]$\leq$$-$0.5. Five of the seven orbital elements have been determined: period ($P$), time of the periastron ($T$), semi-major axis ($a$), eccentricity ($e$) and position of periastron ($\omega$). The binary's orbital inclination ($i$) and longitude of the ascending of node ($\Omega$) remain unknown from spectroscopic data alone. Because of its relatively large $a\sin i$ value among the 34 SB2 systems in \citet{Goldberg2002}, we expected the Fine Guidance Sensor on the Hubble Space Telescope to be able to resolve the system at appropriate orbital phases. By accurately measuring the angular separation of the two components at only a few epochs, we can reliably determine all of the orbital elements and thus the individual masses of the B and C components. \begin{deluxetable}{lccccccccc} \rotate \tablewidth{0pt} \tabletypesize{\small} \tablecaption{Metallicities and Dynamical Masses of Subdwarfs\label{tbl:mass}} \tablehead{ \colhead{Name} & \colhead{$\pi$} & \colhead{[m/H]} & \colhead{[Fe/H]} & \colhead{Ref.} & \colhead{Mass (A)} & \colhead{$M_{V}$} & \colhead{Mass (B)} & \colhead{$M_{V}$} & \colhead{Ref.} \\ \colhead{} & \colhead{mas} & \colhead{} & \colhead{} & \colhead{} & \colhead{$M_{\odot}$} & \colhead{} & \colhead{$M_{\odot}$} & \colhead{} & \colhead{} } \startdata HIP 85209 AB & 19.76$\pm$0.82 &($-$0.75)$^b$& & 2 & 0.84$\pm$0.09 & 5.42 & 0.74$\pm$0.10 & 5.41 & 1$^a$ \\ HIP 95575 AB & 39.98$\pm$0.73 & $-$0.80 & & 3 & 0.69$\pm$0.10 & 6.80 & 0.58$\pm$0.09 & 7.03 & 1$^a$ \\ $\mu$ Cas AB &132.67$\pm$0.74 & &$-$0.71 & 9 & 0.742$\pm$0.059 & 5.78 & 0.173$\pm$0.011 & 11.6 & 4 \\ GJ 704 AB & 64.30$\pm$0.68 & &$-$0.58 & 6 & 0.89$\pm$0.03$^c$ & 4.16 & 0.51$\pm$ 0.03$^c$ & 7.46 & 5 \\ G 006-026 BC & 25.67$\pm$0.04 & $-$0.88 & & 8 & 0.474$\pm$0.053 &10.34 & 0.436$\pm$0.049 & 10.64 & 7 \\ \enddata \tablecomments{All parallaxes other than $\mu$ Cas AB and GJ704 AB are from \cite{Leeuwen2007}. The weighted mean parallax of $\mu$ Cas AB is from \cite{Leeuwen2007} and \cite{YPC}. The weighted mean parallax of GJ704AB is from \cite{YPC} and \cite{Soderhjelm1999}. $^a$\citet{Horch2015} did not report individual mass errors, but they have total mass ($M_{A}+M_{B}$) errors. We adopt these total mass errors to get errors in percentage of masses and apply them to each component. $^b$The metallicity of HIP 85209 is discussed in $\S$6. $^c$The masses of GJ 704 AB are discussed in $\S$\ref{sec:GJ704AB} and details are given in Table~\ref{tbl:GJ704}} \tablerefs{ (1) \cite{Horch2015}, (2) \cite{Latham1992}, (3) \cite{Holmberg2009}, (4) \cite{Drummond1995}, (5) \cite{Soderhjelm1999}, (6) mean value of several measurements given in Table 5, (7) this work, (8) \cite{Carney1994}, (9) \cite{Karaali2003} } \end{deluxetable}

\label{sec:MLR} We plot the 10 low-metallicity subdwarfs discussed here as solid points on and empirical mass-luminosity relation shown in Figure~\ref{fig:mlr}. The masses come directly from Tables 1, 3 and 4, where in addition to the subdwarfs, main sequence stars with masses from \cite{Delfosse2000}, \cite{Torres2010}, and \cite{Benedict2016} are plotted in Figure~\ref{fig:mlr}. These main-sequence stars generally fall along the 1 Gyr theoretical isochrone from \cite{Baraffe2015}, shown with a solid line. The ages of low metallicity stars are primarily measured using nearby GCs. Recently, \cite{Bono2010} and \cite{Correnti2016} acquired deep optical and near-IR photometry from the ground and HST to extend the detected GC populations down to early M dwarfs. Many of these clusters/cool dwarfs have ages estimated to be around $\sim$11 Gyrs. In addition, \cite{Monteiro2006} measured the first ages for two cool field K subdwarfs via their white dwarf companions' cooling curves, and found that these likely thick disk subdwarfs have ages of 6--9 Gyr. However, until we can link the origins of these field subdwarfs to nearby GCs, we conservatively adopt 6 Gyr as the age for the all nearby subdwarfs. We sketch 6 Gyr isochrones for stars with [m/H] = $-$0.5, $-$1.0, and $-$2.0 from BT-Settl isochrones \citep{Baraffe2015} in Figure~\ref{fig:mlr} to permit comparisons to the empirical points for the subwarfs. According to the models, low metallicity subdwarfs appear brighter in $M_{V}$ than dwarfs at a given mass, or alternately, are less massive for a given $M_{V}$. Systems like $\mu$ Cas AB and GJ 704 AB are clearly seen above the mass-luminosity relation of main sequence stars. The locations of $\mu$ Cas A and B, which have [m/H] = $-$0.71, are well-matched to the model predictions for stars with [m/H] = $-$1.0 at an age of 6 Gyr. The metallicity of GJ 704 AB appears much lower ([m/H]$<-$1.0) than what has been measured [Fe/H]=$-$0.58, but the two subdwarfs certainly appear to be of lower metallicity than main sequence stars. We note that \citet{Kennedy2012} used {\it Herschel's} Photodetector and Array Camera and Spectrometer (PACS) at 100 and 160 $\mu$m to reveal a rare and convincing circumbinary polar-ring debris disk around GJ 704 AB. This makes this system very unusual because this disk has somehow been maintained or created around a presumably old subdwarf binary. $M_V$ values used in Figure~\ref{fig:mlr} for GJ 704 A and B do not include any adjustment for obscuration by the disk, which is likely to be minimal. If such an adjustment is required, the points would move to brighter $M_V$, and the locations of the points would be towards even lower metallicity. The remaining subdwarfs systems generally fall along the low metallicity curves. Because of the large mass errors of HIP 85209 AB and HIP 95575 AB measured by \cite{Horch2015}, it is difficult to make a direct comparison between observations and models. \cite{Latham1992} reported a metallicity for HIP 85209 AB (HD 157948 AB) of [m/H]=$-$0.75, whereas \cite{Goldberg2002} reported [m/H]=$-$0.5. \cite{Horch2006} measured the dynamical masses of these subdwarfs using FGS, and determined via comparisons to the theoretical Yale-Yonsei models that the system should have metallicity close to $-$0.5. We note that the mass errors for HIP 85209 A and B are relatively large, so these two stars do not place strict constraints on the models. Finally, both components in the new system described in this paper, G 006-026 B and C, merge with points for main sequence stars with solar metallicity in the MLR. This presents a conflict with the low metallicity measurements discussed in Section~\ref{sec:intro} --- the points should lie near the [m/H] = $-$0.5 or $-$1.0 lines in Figure 7. For example, \cite{Carney1994} used high S/N echelle spectrograph to determine the metallicity of the primary in the system G6-26A, which is also a SB2 binary with a period of 8.7 days \citep{Goldberg2002}. They discussed the challenges of measuring the metallicities of SB2 binaries in detail, but there is no flag of the system regarding quality of their result of [m/H] = $-$0.88. \citet{Casagrande2011} and \cite{Holmberg2009} used photometry of the unresolved primary to determine [Fe/H] = $-$0.52 and $-$0.60, respectively, thereby providing results that are match Carney's, again indicating that the system is metal-deficient. In order to move the masses of B and C components to the low metallicity region, B component's mass would need to be reduced to $\sim$0.3 M$_{\odot}$ (62.7\% of 0.478 M$_{\odot}$) to fall on the $-0.5$ metallicity curve on Figure~\ref{fig:mlr}. In Kepler's Third Law, such a mass decrease requires that the semi-major axis drop by 15\%. There would then be a corresponding reduction in the projected separation ($\rho$) measured using FGS. Figure~\ref{fig:fakeS} shows an example of changing the projected separation along the X-axis by 15\% while comparing to our best fitted S-curve. Note the poor fit to the FGS data, compared to the much better correct fit. In addition, because our orbital parallax matches that from Hipparcos, it is difficult to change the semi-major axis of this system from our derived value of 10.8 mas. Given that it is hard to dispute the three independent metallicty measurements or shrink the separation between B and C, we are left with no easy answer to the conundrum that that two points do not fall on the MLR where they should. Given the overall trend that subdwarf systems are elevated in the mass-luminosity relation and match with the model, but there are several confounding factors that challenge observational programs focusing on mapping the mass-luminosity relation for subdwarfs, including (1) metallicity errors may be large because all available subdwarf binaries have separations less than 1\arcsec, so the acquired spectra used to determine metallicities are for two objects, not one, with consequent complications when measuring linewidths (G 006-026 AB is the only system to have an early type primary, which can be used to determine metallicity independently.), (2) mass errors are large because subdwarfs are rare and further away than comparable main sequence systems in the solar neighborhood, so parallax errors are larger, (3) accurate $\Delta V$ measurements are not available for a few systems and conversions from other filters are imperfect (although overall, this is not likely to be a major problem), (4) not all of the systems are of the same age, so comparison to a single isochrone for 6 Gyr-old systems may not be appropriate. This work shows that building an empirical mass-luminosity relation for subdwarfs has just begun. Only two subdwarf systems, $\mu Cas$ AB and GJ 704 AB, are clearly off the metallicity curves for main-sequence stars and have low mass errors, so there remains significant work to do in understanding the astrophysics that drives the locations of low metallicity stars on the mass-luminosity relation. We can make progress in our understanding of these Galactic fossils, as well as the star formation history and mass distribution of stars in the Milky Way, by discovering more subdwarf binaries for which high-quality metallicities, luminosities, and masses can be measured. \begin{figure} \centering \includegraphics[scale=0.8]{fig6.ps} \caption{Relation $\Delta V_{Tycho}$ vs $\Delta m_{692}$ from 51 binaries in Washington Double Star Catalog. The solid straight line is a fit ($\Delta V_{Tycho}=-0.35+0.95\times\Delta m_{692}$) for these points.} \label{fig:692totychov} \end{figure} \begin{figure} \centering \includegraphics[scale=0.7, angle=90]{fig7.ps} \caption{The mass-luminosity relation for subdwarfs (filled symbols) and dwarfs (open circles). Names are given for subdwarfs. Open circles are from \cite{Delfosse2000}, \cite{Torres2010}, and \cite{Benedict2016}. Subdwarfs from \cite{Horch2015} are filled triangles, with a different sizes used simply to match the two stars in a system. Filled circles represent $\mu$ Cas AB and filled boxes represent the two systems discussed in detail in this paper, GJ 704 AB and G 006-026 BC. The thick solid line represents the main-sequence isochrone for age 1 Gyr from BT-settl \citep{Baraffe2015}. The three other lines represent the predicted locations of stars with metallicities of [m/H]=$-$0.5, $-$1.0 and $-$2.0 at an age of 6 Gyr from bottom to top, taken from \cite{Baraffe2015}} \label{fig:mlr} \end{figure} \begin{figure} \centering \includegraphics[scale=0.7]{fig8.ps} \caption{Displayed is the best fitting model (blue line) to the 2013.09 observation of G 006-26 BC (red line) along the FGS X-axis, with an angular separation of -21.0 mas. Also displayed is a model with an angular separation = -17.85 mas (black line), which would have resulted in the B component mass of $\sim$0.3 M$_{\odot}$ on the -0.5 theoretical metallicity line of the mass luminosity relation. Clearly this is inconsistent with the observation.} \label{fig:fakeS} \end{figure}

1607.01304

1607

1607.06462_arXiv.txt

Microlensing provides a unique tool to break the stellar to dark matter degeneracy in the inner Milky Way. We combine N-body dynamical models fitted to the Milky Way's Boxy/Peanut bulge with exponential disk models outside this, and compute the microlensing properties. Considering the range of models consistent with the revised MOA-II data, we find low dark matter fractions in the inner Galaxy: at the peak of their stellar rotation curve a fraction \onesigbounds~of the circular velocity is baryonic (at $1\sigma$, \twosigmaxfrac~at $2\sigma$). These results are in agreement with constraints from the EROS-II microlensing survey of brighter resolved stars, where we find \onesigboundseros~at $1\sigma$. Our fiducial model of a disk with scale length 2.6\kpc, and a bulge with a low dark matter fraction of 12\%, agrees with both the revised MOA-II and EROS-II microlensing data. The required baryonic fractions, and the resultant low contribution from dark matter, are consistent with the NFW profiles produced by dissipationless cosmological simulations in Milky Way mass galaxies. They are also consistent with recent prescriptions for the mild adiabatic contraction of Milky Way mass haloes without the need for strong feedback, but there is some tension with recent measurements of the local dark matter density. Microlensing optical depths from the larger OGLE-III sample could improve these constraints further when available.

The concept of microlensing, whereby a background star is gravitationally lensed by a foreground object, was conceived by \citet{Einstein:36} however it was not until \citet{Paczynski:86} that modern microlensing experiments were born. One of the primary observational microlensing measurements is the optical depth. The optical depth is defined as the fraction of stars whose projection lies within the Einstein radius of a lens, and it is these stars who see a significant ($>3/\sqrt{5}$) brightening due to microlensing. For a star at distance \ds the optical depth $\tau$ is given by \citep[\eg][]{Kiraga:94} \begin{equation} \tau(\ds) = \frac{4\pi G }{c^2} \int_0^\ds \rhol(\dl) \left(\frac{1}{\dl} - \frac{1}{\ds} \right) \dl^2 \, d\dl \label{eq:tau} \end{equation} where $\rhol(\dl)$ is the density of lenses. That the optical depth depends only on the density of lenses, and not on their distribution of masses or velocities, makes it theoretically very elegant. An original motivation for microlensing studies was to explore the possibility that the Galactic dark halo could be composed of Massive Astrophysical Compact Halo Objects (MACHOs) \citep{Paczynski:86}. The first results from the MACHO survey measuring the optical depth towards the LMC suggested, based on $\sim 10$ events, that 20\% of the Galactic halo could be composed of MACHOs with average mass $\sim 0.4\msun$ \citep{Alcock:00,Bennett:05}. In contrast the EROS survey found that less than 8\% of the halo could be composed of MACHOs with mass $\sim 0.4\msun$ \citep{Tisserand:07}. The OGLE-III survey was much larger in both duration and coverage and found stringent limits of the contribution of MACHOs: a fraction less than $7\%$ for sub-solar lenses \citep{Wyrzykowski:11}. Hence, the optical depth observed towards the LMC has turned out to be too small for the halo to contain more than a small fraction of dark objects able to microlense, most likely because dark matter is composed of particles with low mass, so that they have an Einstein radius far too small to produce a significant brightening \citep[\ie $\lesssim 10^{-7}\msun$, ][]{Paczynski:86}. This instead makes microlensing of Milky Way bulge stars a unique tool for unravelling the structure of galaxies because only stellar matter is able to produce microlensing events. Typically the mass profiles of Galaxies can be constrained dynamically, but this may be distributed between baryonic and dark matter \citep[\eg][]{Courteau:14}. In contrast we are able to use bulge microlensing to break the stellar to dark matter degeneracy in the inner Milky way ($R \lesssim 5\kpc$) in a manner that is generally not possible in external galaxies without additional assumptions such as on the stellar mass-to-light ratio \citep{Iocco:11}. Microlensing towards the bulge was first considered by \citet{Paczynski:91} and \citet{Griest:91bulge}. Initial predictions were based on bulge models fitted to COBE data \citep{Han:95,Bissantz:97,Evans:02,Bissantz:02}. However the measured optical depths were significantly higher than these models predicted. OGLE measured $\tau=(3.3\pm1.2)\times 10^{-6}$ \citep{Udalski:94} and MACHO $\tau =(3.9\substack{+1.8 \\ -1.2})\times 10^{-6}$ \citep{Alcock:97}. These early optical depths were higher even than the theoretical bounds: no model could be constructed that reproduced $\tau$ without overshooting the rotation curve \citep{Binney:00}. Later estimates with improved treatment of blending through difference image analysis (DIA) \citep{Alcock:00bulge} and focusing on brighter resolved giants \citep{Popowski:05,Hamadache:06} produced reduced optical depths, so that at least some possible Galactic models were consistent with the data \citep[\eg][]{Wood:05}, and actually brought the optical depth into line with the earlier predictions. Pioneering studies from OGLE \citep{Udalski:94} and MACHO \citep{Alcock:97} were based on a handful of microlensing events. Since then several thousand microlensing events have been detected by recent and ongoing microlensing surveys such as EROS-2 \citep{Afonso:03,Hamadache:06}, MOA-2 \citep{Sumi:11,Sumi:13,Sumi:16}, MACHO \citep{Popowski:05}, WeCAPP \citep{Lee:15}, OGLE-III \citep{Wyrzykowski:15}, and OGLE-IV \citep{Udalski:15}. These larger samples, the convergence between measurements of recent optical depths for bright resolved stars with fainter unresolved stars \citep{Sumi:16}, and the new made-to-measure models of the bulge \citep{Portail:15}, prompts us to revisit of the constraints provided by bulge microlensing data on Galactic models. We use recent microlensing data to constrain the amount of stellar matter towards the Galactic bulge, and consequently the fractional contributions of stellar and dark matter in the inner Galaxy. The data is primarily taken from the MOA-II survey \citep{Sumi:13} which was recently updated by \citet{Sumi:16} with improved estimates of the effective number of monitored stars. We use the data from \citet{Sumi:16} throughout, referring to it as the revised MOA-II data. We additionally consider and cross check our results with data from the EROS-II survey \citep{Hamadache:06}. \begin{table*} \caption{Summary description of the model Galaxy used to predict microlensing quantities.} \label{tab:model} \begin{tabular}{llll} \hline & Property & Fiducial & Variations \\ \hline Bulge Model & N-body model from \citet{Portail:15} & M90 \ & M80, M85 \\ \hline \multirow{4}{*}{Disk Model} & Disk scale length, $R_d$ & 2.6\kpc & Range: $1.9-3.4\kpc$\\ & Solar neighbourhood disk scale height, $H_\odot$ & 0.3\kpc & --- \\ & Inner disk scale height , $H_{4.5}$ & 0.18\kpc & 0.3\kpc \\ & Local stellar surface density, $\Sigma_\odot$ & $38\, \msun\,\kpc^{-2}$ & --- \\ \hline \multirow{4}{*}{Stellar Population} & Source stellar population & 10Gyr, Baade's Window MDF from \citet{Zoccali:08} & --- \\ & Isochrones & $\alpha$-enhanced BASTI \citep{Pietrinferni:04} & --- \\ & IMF & \citet{Kroupa:01} & \citet{salpeter:55} \\ & Remnant masses & Prescription from \citet{Maraston:98} & --- \\ \hline \end{tabular} \end{table*} This paper proceeds as follows: In \autoref{sec:models} we describe our models of the distribution of stellar matter in the Milky Way. In \autoref{sec:microtheory} we describe how we calculate microlensing properties such as the optical depth and timescale distribution which can then be subsequently compared to the data in sections \ref{sec:data} and \ref{sec:eros}. We compute the resultant rotation curves consistent with data in \autoref{sec:rotcurve}, and briefly investigate the timescale distribution of events in \autoref{sec:tedistdiscuss}. We discuss the consequences of the results in \autoref{sec:discuss}, in particular the limits these place on the halo contribution, and place our Milky Way constraints in the context of external galaxies. We conclude in \autoref{sec:conc}.

Although the focus of this work are the constraints that surveys of microlensing in the bulge place on Galactic structure, which is best revealed by the optical depth, we briefly consider the timescale distribution in this section. While the optical depth has not yet been computed from the OGLE-III survey, the timescale distribution has by \citet{Wyrzykowski:15}. In the upper panel of \autoref{fig:tevssurvey} we show a comparison of this with the timescale distribution measured by MOA-II. \begin{figure} \includegraphics[width=\linewidth]{tevssurvey} \caption{In the above panel we show the efficiency corrected timescale distribution of the OGLE-III survey (blue) compared to the MOA-II sample (red). We also show the best fitting log-normal distributions to the surveys as the curves. In the lower panel we compare the efficiency corrected OGLE-III timescale distribution to our fiducial model with different IMFs: A \citet{salpeter:55} IMF (short dashed line), a \citet{Zoccali:00} IMF (long dashes), a \citet{Kroupa:01} IMF (dash-dot) and the log-normal IMF of \citet{Calamida:15} (dotted). \label{fig:tevssurvey}} \end{figure} The two samples agree in their timescale distribution remarkably well. Fitting a log-normal distribution to each we find the mean log timescales are $<\log \te>=(1.275\pm0.008)$ and $(1.25\pm0.03)$ for the OGLE-III survey and MOA-II all source sample respectively, in statistical agreement. Likewise the standard deviation of the log timescales are $\sigma(\log \te)=(0.409\pm0.006)$ and $(0.41\pm0.02)$, also in agreement. Note that a log-normal distribution does not capture the short timescale events present in the MOA-II sample. These are presumed to arise from microlensing by planetary mass objects \citep{Sumi:11}. It is also not expected to capture the asymptotic distribution of short or long timescale events \citep{Mao:96} but, given the current data, is still a useful parameterisation. This is particularly true since it provides the first two moments of the log timescale distribution. The log timescale distribution is a convolution of the two PDFs in \autoref{eq:teconvobv}, resulting from the dynamical model, and the lens mass distribution. Because of this the means and variances add (and indeed all cumulants) to give the mean and variance of the log timescale distribution. The mean log timescale and its variance are therefore straightforward to interpret and consider the effects of other mass distributions or dynamical models. In the lower panel of \autoref{fig:tevssurvey} we show the timescale distribution predicted from our model for four different IMFs: those from \citet{salpeter:55} (a single power law with slope $\alpha=2.3$), a \citet{Zoccali:00} IMF ($\alpha=1.3$ for $M<1\msun$), a \citet{Kroupa:01} IMF, and log-normal IMF fitted for the bulge by \citet{Calamida:15} (log-normal with $M_c=0.25\msun$, $\sigma=0.5$). All IMFs use lower and upper mass limits of $0.01\msun$ and $100\msun$ respectively, aside from the \citet{salpeter:55} IMF which uses a lower limit of $0.1\msun$. For all we use the remnant prescriptions of \citet{Maraston:98} above the turnoff of a $10\Gyr$ population as described in \autoref{sec:pop}. Given that we have not tailored the model to the timescale distribution it is reassuring that the \citet{Kroupa:01} and particularly the \citet{Calamida:15} IMFs together with the dynamical model reproduce the timescale distribution fairly well. There is a discrepancy at long timescales which the model under predicts. As a result, with the \citet{Calamida:15} IMF the model predicts $<\log \te>=1.21$ slightly less than the observed $<\log \te>=(1.275 \pm 0.008)$. The most uncertain region of the mass function is the brown dwarf region. That the short timescale distribution agrees with the OGLE-III data for the \citet{Calamida:15} log-normal IMF suggests that our model requires a low number of brown dwarfs to be present in the bulge compared to rising power law IMFs like the \citet{Kroupa:01} IMF, similar to local estimates \citep[\eg][]{Andersen:08}. More detailed comparison is beyond the scope of this work and would require consideration of variations of the kinematical model, particularly outside the central region where the N-body model was constrained by \citet{Portail:15}. We are presently constructing of made-to-measure models of the entire inner Galaxy, including the inner disk, and we will model the timescale distribution and place more robust constraints on the IMF when that is complete. We emphasise that the results outside this section are based on the optical depth, which is insensitive to the timescale distribution. This is demonstrated in \autoref{sec:minoraxis} by the insensitivity of the optical depth to the much shorter timescales produced by the \citet{salpeter:55} IMF. \subsection{A high baryonic fraction in the inner Milky Way} The quantity best constrained by the microlensing optical depth is the stellar mass per unit solid angle towards the bulge and inner Galaxy. We find that there is a degeneracy in where that mass is placed between the bulge and the foreground disk: models with lower stellar fractions in the bulge require more prominent foreground disks and vice versa. The scale length of the Milky Way disk is highly uncertain. Earlier estimates using optical wavelengths tended to favour larger scale lengths \citep[\eg 3.5\kpc by][]{Bahcall:80}, while more recent measurements in the NIR and large photometric parallax surveys have favoured lower measurements typically in the range 2-3\kpc ( \eg \citealt{Binney:97,Bissantz:02,Robin:03,Juric:08,Bovy:13}; see \citealt{Ortwin:araa} for a fuller discussion). Despite this uncertainty, in all the models which reproduce the revised MOA-II optical depths, the baryonic contribution to the rotation curve at its peak is high: \onesigbounds~at $1\sigma$ and $\twosigmaxfrac$ at $2\sigma$. These are consistent with the constraints from EROS-II of \onesigboundseros. As discussed in \autoref{sec:rotcurve}, we consider maximal disks to have $\fv > 0.85$ \citep{Courteau:14}. These levels of baryonic contribution therefore place the Milky Way on the boundary between maximal disk fits and sub-maximal disks . The DiskMass survey has measured the vertical velocity dispersions of nearly face on disk galaxies. Together with the statistically determined scale height from a sample of side on galaxies, this directly measures the disk mass under the assumption of a locally isothermal disk. They find that their sample is generally sub-maximal with a fractional baryonic contribution at $2.2R_d$ of $0.57\pm0.07$ \citep{Martinsson:13}. However the measured light weighted stellar kinematics will be biased towards younger stars as opposed to old dynamically relaxed populations. Measuring the vertical kinematics of the Milky Way locally suggests that this correction could be a factor of $\sim 2$ \citep{Aniyan:16}. In addition the resultant IR mass-to-light of the disk mass survey is also a factor of $\sim 2$ smaller than that estimated through other methods \citep{McGaugh:15}. The derived maximal disk in this work would therefore not necessarily make the Milky Way unusual in the context of external disk galaxies. In addition in external barred galaxies a further possibility exists to break the baryonic to dark matter degeneracy. The stellar distribution can be determined from the light, while the non-circular gas motions constrain the shape of the effective potential, and therefore the more spherical dark matter, in the bar. Using this method several studies have found maximal or near maximal disks in the inner regions of barred galaxies \citep{Weiner:01,Weiner:04,Perez:04,Sanchez:08}. \subsection{Consequences for the Milky Way's dark halo} \begin{figure} \includegraphics[width=\linewidth]{rotadiabaticgrey} \caption{We show the range of baryonic rotation curves consistent with the revised MOA-II data at the $1\sigma$ level together with a selection of possible dark matter haloes. The baryonic rotation curves are shown as the cross hatched area, the dark matter profiles as the black lines, and the resultant range of profiles of the total rotation curve in grey. The top left panel shows an NFW profile with concentration $c_{200}=9.0$. The upper right an NFW halo, also with $c_{200}=9.0$, adiabatically contracted through the prescription \citet{Blumenthal:86}. The lower left panel shows the contraction suggested by \citet{Gnedin:04} for a smaller but still cosmologically possible initial $c_{200}=7.0$, and the lower right the contraction fitted by \citet{Abadi:10} for $c_{200}=9.0$. \label{fig:rotadiabatic} } \end{figure} The high baryonic contribution to the rotation curve in the inner Galaxy seen in \autoref{fig:rotadiabatic}, and the resultant high levels of disk maximality, are not however inconsistent with the current understanding of the $\Lambda$CDM paradigm. An NFW dark matter halo with a Milky Way halo mass $M_{\rm 200}=1.1 \times 10^{12} M_\odot$ and radius $R_{\rm 200}=270\kpc$ \citep{Ortwin:araa} and a concentration of $c_{\rm 200}=9.0$ motivated by cosmological simulations (\citealt{Correa:15c} with Planck 2015 cosmology \citealt{PlanckCosmo:15}) gives the profile shown in the top left panel of \autoref{fig:rotadiabatic}. Adding this to the range of baryonic contributions allowed by the revised MOA-II data at $1\sigma$ gives rotation curves that remain below the total rotation curve. \begin{figure*} \includegraphics[width=0.7\linewidth]{fullmaps} \caption{The microlensing optical depth of the fiducial model across the entire Milky Way bulge (top left), compared to the same variants considered earlier in the paper (\eg in \autoref{fig:minor_sumicomp}). We show the effect of reducing the disk scale length to $R_d=2\kpc$ from $R_d=2.6\kpc$ in the top right, using a bulge model with a higher dark matter fraction and lower stellar mass in the bulge (M80) in the lower left, and a model with a constant scale height of $H_{4.5}=0.3\kpc$ in the lower right. All maps are the average optical depth over stars with $14<I_0<19$. \label{fig:fullmaps}} \end{figure*} Adiabatically contracted versions of this halo using the prescription of \citet{Blumenthal:86} are generally not consistent with the revised MOA-II data at the $1\sigma$ level. In this prescription the adiabatic invariant is $r M(<r)$ which results in rotation curves that overshoot that observed in the inner Galaxy (top right panel \autoref{fig:rotadiabatic}). This is true even for lower concentrations as small as $c_{\rm 200}=7.0$, which is the expected cosmological halo-to-halo scatter at $1\sigma$ level \citep{Maccio:08}. It is also true for the contraction detirmined by \citet{Piffl:15} since it results in more contracted haloes than \citet{Blumenthal:86}. \citet{Binney:15} similarly concluded that halos with this level of contraction, in conjunction with the local dark matter density, could not both be consistent with the Milky Way's rotation curve and earlier microlensing optical depths. However other milder contraction prescriptions are consistent with the revised MOA-II data at the $1\sigma$ level. The prescription of \citet{Gnedin:04} is consistent only if the initial halo has lower concentrations ($c_{\rm 200}=7.0$ is shown in the lower left panel of \autoref{fig:rotadiabatic}). The prescription of \citet{Abadi:10} results in less contracted halos, and as a result is consistent with an initial $c_{\rm 200}=9.0$ (lower right panel of \autoref{fig:rotadiabatic}). All haloes were contracted using a baryonic exponential scale length of $2.6\kpc$ and contraction of the \citet{Gnedin:04} and \citet{Abadi:10} halos was implemented through the $\Gamma$ prescription of \citet{Dutton:07,Dutton:11} with $\Gamma=0.8$ and $0.4$ respectively. All contraction prescriptions can be consistent with the range of rotation curves allowed by the MOA-II data at the $2\sigma$ level, and the EROS-II data. It has been shown that strong feedback could alter the central profiles of Milky Way sized haloes \citep{Maccio:11}. The results here suggest that it is not \emph{required} in the centre of Milky Way sized galaxy haloes in a similar way to dwarf galaxies. The microlensing results however do not rule it out, for example the feedback in \citet{Chan:15} for $10^{12} M_\odot$ mass haloes results in central profiles similar to the pre-contracted NFW haloes which would still be consistent. We would expect our results to also be consistent with other recent simulations \citep[\eg][]{Marinacci:14,Schaller:15} which generally produce only mildly contracted haloes. We conclude that the levels of baryonic contribution towards the inner Galaxy, \fv, are high, but not inconsistent with the estimates of contribution of CDM haloes in the inner region of disk galaxies, provided the level of baryonic contraction is not too extreme. There is however possible tension between the low baryonic fractions in the inner galaxy measured here, with recent estimates of the local dark matter density. \citet{Piffl:14} measured the local dark matter density to be $0.0126 q^{-0.89} \msun \pc^{-3} \pm 15\%$ in agreement with $0.014 \msun \pc^{-3}$ measured by \citet{Bienayme:14}. These more recent measurements utilising the exquisite RAVE data are higher than most over that past decade \citep[see][and references therein]{Read:14}. The haloes shown in \autoref{fig:rotadiabatic} however have local dark matter densities in the range $0.005-0.008 \msun pc^{-3}$. The NFW profile considered by \citet{Piffl:14} would be consistent with the range of models considered here and would reproduce also local dark matter density. However this model had an uncomfortably high concentration of $c=20$, significantly higher than expected in dark matter only simulations for Milky Way sized halos. Instead it seems more natural that a less concentrated halo experienced a mild degree of contraction due to growth of the baryonic disk, and this increased the dark matter density at $\sim R_0$. Some tension remains however between the high baryonic fraction in the inner Galaxy measured here, the local dark matter density, and even mild adiabatic contraction prescriptions. Contracting a halo with $c_{200}=15.$, towards the upper end of the cosmological values, with the prescription of \citet{Abadi:10} results in rotation curves just consistent with the microlensing results here. This however gives a solar dark matter density of $0.01 \msun pc^{-3}$, still somewhat lower than the estimates with RAVE. We therefore encourage direct comparison between the high baryonic fraction in the inner Galaxy required by the microlensing here, the local dark matter density, and recent simulations of Milky Way like galaxies in the $\Lambda$CDM paradigm such as \citet{Marinacci:14}. \subsection{Microlensing as a tool for Galactic structure} The measured optical depths close to the plane are particularly important in driving the need for high stellar contributions in the inner Galaxy and low dark matter fractions, because the majority of the stellar mass lies here. It is therefore important to confirm the optical depth measurements in this region. We anticipate that the constraints in this work could be verified and improved on by the larger sample of 3560 events, and the better detection efficiency of longer events, provided by the OGLE-III sample \citep{Wyrzykowski:15}. Of these 2047 lie inside $|b|<3\deg$ and 546 inside $|b|<2\deg$, therefore the optical depths measured by MOA-II in this region can be verified. The optical depth and event rates require detailed field-by-field efficiency calculations that are not yet complete, but we encourage their computation. Microlensing in the Milky Way is a unique tool to break the dark-matter to baryonic matter degeneracy, however since a large fraction of the mass of the disk lies at these low latitudes measurements are essential in this challenging region. A further important tool to this goal will be microlensing surveys in the NIR. To help guide the design and assess the impact of future surveys we provide in \autoref{fig:fullmaps} maps of the optical depth across the entire Galactic bulge. Recently \citet{Awiphan:16} also modelled the MOA-II data within the Besan\c on model, finding optical depths smaller than the observed data, particularly close to the Galactic plane. Some of this discrepancy was resolved by the revised MOA-II data, but the optical depths are remain lower than those observed. The dynamical models used in this work have higher bulge mass than those in \citet{Awiphan:16} which resolves the discrepancy. We note that while the bulge dynamical models do not include a nuclear bulge component inside $1\deg$ \citep{Launhardt:02}, since this is significant only inside $|b|<1\deg$ it does not effect the MOA-II measurements which lie outside this. The optical depth is fundamentally a weighted mean of the density of lenses along the line-of-sight towards each source star (see \eg \autoref{eq:tau}). Because of this the microlensing data of the bulge constrains the stellar mass distribution towards centre of the Galaxy. The weighting of this density measurement is most sensitive to events halfway between the source and lens, and because of our position in the Galaxy is weighted to be sensitive to the stellar density at $R \sim 4\kpc$. Since this is likely to be near the position of the peak in the stellar rotation it makes the microlensing optical depth an attractive method to measure the disk maximality. A disadvantage is that since there is effectively only one line-of-sight (that towards the bulge) assumptions on the form of the density are needed to convert the optical depth to a stellar density constraint. In this work we assume the disk is axisymmetric, however this assumption is likely to be violated to some extent, and certainly inside the radius of the Galactic bar of $(5.0\pm0.2)\kpc$ \citep{Wegg:15}. Our fiducial model has a disk mass between 2.2\kpc and 5\kpc of $2.1\times10^{10}\msun$ while the mass of the bar outside the bulge derived by \citet{Wegg:15} is $1\times10^{10}\msun$. As a result we expect a significant fraction of the mass in this range may be part of the bar and therefore non-axisymmetric. Fortunately our line-of-sight to the bulge lies at an angle $(27\pm 2)\deg$ to the bar \citep{Wegg:13}, and as a result we are neither looking along or perpendicular to the bar. In addition the microlensing optical depth could be sensitive to the presence of a ring or trailing structures at the end of the bar, since this would lie almost halfway to the bulge, where the microlensing is most sensitive to the density of lenses. We note that microlensing models of the bulge, such as that presented here, are important not just for the constraints they provide on the galaxy. Planet detection through microlensing has become as important tool to probe distant and low mass planets, and this is an important component of the upcoming EUCLID and WFIRST missions \citep{Beaulieu:11}. The expected yields and targeting requires microlensing models such as those described in this work.

1607.06462

1607

1607.01242_arXiv.txt

Sulphur trioxide (SO$_3$) is a trace species in the atmospheres of the Earth and Venus, as well as well as being an industrial product and an environmental pollutant. A variational line list for $^{32}$S$^{16}$O$_{3}$, named UYT2, is presented containing 21 billion vibration-rotation transitions. UYT2 can be used to model infrared spectra of SO$_3$ at wavelengths longwards of 2 $\mu$m ($\nu < 5000$ cm$^{-1}$) for temperatures up to 800 K. Infrared absorption cross sections are also recorded at 300 and 500 C are used to validate the UYT2 line list. The intensities in UYT2 are scaled to match the measured cross sections. The line list is made available in electronic form as supplementary data to this article and at \url{www.exomol.com}.

SO$_3$ known to exist naturally in the Earth's atmosphere; its main source being volcanic emissions and hot springs \citep{05MiKrGrAn.SO3} but it also plays role in the formation of acid rain. The oxidisation of SO$_2$ to SO$_3$ in the atmosphere, followed by subsequent rapid reaction with water vapour results in the production of sulphuric acid (H$_2$SO$_4$) \citep{ 85CaLaKo.SO3} with many adverse environmental effects \citep{11VaGoHa.SO3,94KoJaWo.SO3,04SrMiEr.SO3}. So SO$_3$ is a natural product whose concentration in the atmosphere is significantly enhanced by human activity, particularly as a byproduct of industrialisation. SO$_3$ is observed in the products of combustion processes \citep{04SrMiEr.SO3,14HiMexx.SO3} and selective catalytic reduction (SCR) units, where the presence of both is undesirable within flue gas chambers in large quantities, as well as other industrial exhausts \citep{05RaHeSoOa.SO3,12FlVaAnBr.SO3}. The control of these outputs is therefore of great importance. The spectroscopic study of sulphur oxides can also provide insight into the history of the Earth's atmosphere \citep{13WhXiHu.SO2}. All this means that observation of SO$_3$ spectra and hence concentrations provide a useful tool for understanding geological processes and controlling polution. Sulphur oxide chemistry has been observed in a variety of astrophysical settings. Within the solar system, SO$_3$ is a constituent of the atmosphere of Venus \citep{83CrReRa.SO3,10ZhLiMoBe.SO3,13ZhLiMa.SO3}. Although SO$_3$ has yet to be observed outside our solar system, it needs to be considered alongside other sulphur oxides namely, sulphur monoxide (SO) and SO$_2$, which are well-known in several astronomical environments \citep{90NaEsSk.SO2,92PeBexx.SO,03MaMaMa.SO2,05MaMaMa.SO2,12BeMoBe.SO2,06ViLoFe.SO2,13BeMuMe.SO2,13AdEdZi.SO2,15KhViMu.SO2}. SO$_3$ chemistry has been considered in a number of enviroments including giant planets, brown dwarfs, and dwarf stars \citep{06ViLoFe.SO2}. Unlike SO and SO$_2$, SO$_3$ is a symmetric species with no permanent dipole moment making it hard to detect in the interstellar medium. In practice, the identification of SO$_3$ in the infrared is hindered by the presence of interfering SO$_2$ where both species are found simultaneously; a number of their spectral features overlap, particularly the $\nu_3$ bands of both molecules in the 1300 - 1400 \cm\ (7.4 $\mu$m) region. From this point of view SO$_2$ can also be seen as a spectral `weed' with respect to the detection of SO$_3$. An understanding of the spectroscopic behaviour of both of these molecules within the same spectral window is therefore required to be able to correctly identify each species independently. In this context we note that a number of line lists are available for SO$_2$ isotopologues \citep{14HuScLe.SO2,16HuScLe.SO2,jt635}; of particular relevence is the recent hot ExoAmes line list of \citet{jt635}. The experimental spectroscopic studies of SO$_3$ have significant gaps, notably the absence of any measurement of absolute line intensities in the infrared. This may be attributed to its vigorous chemical reactivity which make measurements difficult. SO$_3$ is a symmetric planar molecule with equilibrium S-O bond lengths of 1.41732 \AA\ and inter-bond angles of 120$^{\circ}$ \citep{89OrEsMa.SO3}, described by \Dh(M) symmetry. The $\nu_1$, $\nu_2$, $\nu_3$ and $\nu_4$ fundamental frequencies are attributed to the totally symmetric stretch at 1064.9 \cm \citep{02BaChMa.SO3}, the symmetric bend at 497.5 \cm, and the asymmetric stretching and bending modes at 1391.5 and 530.1 \cm, respectively \citep{03ShBlSa.SO3}. The infrared and coherent anti-stokes vibration-rotation spectra of a number of isotopologues of SO$_3$ have been extensively investigated in a series of papers by Maki and co-workers \citep{73KaMaDo.SO3,89OrEsMa.SO3,01ChVuMa.SO3,01MaBlSa.SO3,02BaChMa.SO3, 03ShBlSa.SO3,04MaBlSa.SO3}, reassessing and confirming fundamental constants and frequencies. 18 bands were analysed based on an empirical fitting to effective Hamiltonian models, yielding rovibrational constants and energy levels assigned by appropriate vibrational and rotational quantum numbers. Some temperature-dependent infrared cross sections are also available from laboratory studies \citep{13GrFaNi.SO2,14GrFaCl.SO3}, and we present new measurements in this work. Unlike all other measurements of SO$_3$ spectra, these cross sections are absolute. However, assigned spectra represented by line lists allow for the modelling of both absorption and emission spectra in different environments. The ``forbidden'' rotational spectrum, for which centrifugal distortions can induce transitions, was investigated for the first time by \citet{91MeSuDr.SO3} using microwave Fourier-transform spectroscopy. Assignments for 25 transitions were made, as well as the determination of a number of rotational constants, including the only direct measurement of the $C_0$ rotational constant. The work was analysed and extended theoretically \citep{jt580}. There have been a few studies on the ultraviolet spectrum of SO$_3$ by \citet{36FaGoxx.SO3}, and \citet{81LeBrRi.SO3}, both between 220 and 270 nm where overlap with SO$_2$ is small. \citet{97BuMcxx.SO3} reported cross sections for the 195 to 330 nm range for the purposes of photolysis rate calculations of SO$_3$. All measurements were taken at room-temperature, and neither reported assignments for any of the bands, which exist as weak, diffuse vibrational band structures superimposed on a continuous background. As such, the rovibronic behaviour of SO$_3$ is much less well understood than for SO$_2$. Prior to our studies, there was limited theoretical work on SO$_3$. Early work on anharmonic force constants \citep{73DoHoMi.SO3,92FlRoTa} for \sothree\ led to first accurate, fully \ai\ anharmonic quartic force field computed by \citet{99Maxxxx.SO3}. There have been no theoretical studies into the UV spectrum of SO$_3$. As for the experimental studies for SO$_3$, none of this work provided transition intensities. Our preliminary study for this project \citep{jt554}, which produced the \ai, room-temperature UYT line list, therefore provides the first absolute transition intensities for \sothree. These were used in the 2012 release of the HITRAN database \citep{jt557} to scale the relative experimental measurements allowing \sothree\ to be included in the database for the first time. As discussed below, the present work suggest that these intensites may need to be reconsidered. The present study on SO$_3$ was performed as part of the ExoMol project. ExoMol aims to provide comprehensive line lists of molecular transitions important for understanding hot atmospheres of exoplanets and other bodies \citep{jt528}. Besides the ExoAmes SO$_2$ line list mentioned above, ExoMol has produced very extensive line lists for a number of polyatomic species including methane (CH$_4$) \citep{jt564}, phosphine (PH$_3$) \citep{jt592}, formaldehyde (H$_2$CO) \citep{jt597}, hydrogen peroxide (HOOH) \citep{jt620} and nitric acid (HNO$_3$) \citep{jt614}. These line lists, all of which contain about 10 billion distinct vibration-rotation transitions, required the adoption of special computational procedures to make their calculation tractable. The UYT2 SO$_3$ line list presented here is the largest computed so far with 21 billion lines; as SO$_3$ is a system comprising four heavy atoms this meant considering rotation states up to $J = 130$ as part of these calculations. These calculations therefore required further enhancement of our computational methods which are described below. The lack of measured SO$_3$ spectra at temperatures above 300~K on an absolute scale is clearly a problem for validating our calculations. Here we present infrared absorption cross sections for SO$_3$ measured at a range of temperatures up to 500~C. The next section describes our theoretical procedures; our experiments are described in section 3. Section 4 presents the UYT2 line list. Section 5 compares the UYT2 line list with our measurements with a particular emphasis on intensity comparisons. The final section gives details on how to access the line list and our conclusions.

The UYT2 line list contains 21 billion transitions, and a total of 18 million energy levels below 10~000\cm. This provides an improvement upon the initial room-temperature line list, UYT, in terms of both line positions and temperature coverage. Table \ref{rmsBands_hot} provides a measure of the improvement introduced by the PES refinement present in the UYT2 line list. The total RMS deviation for the bands included in the potential adjustment is 1.35 \cm, compared to 3.23 \cm\ for the unrefined PES of UYT. The majority of simulated line positions across these bands is improved by an order of magnitude. It is difficult to ascertain the overall quality of the \ai\ DMS used in the production of line intensities. However comparing with newly available cross section data at two different temperatures heavily suggests that the DMS used in the calculation of intensities is slightly overestimated, causing an apparently constant shift in all intensity values. The evidence suggesting this temperature- and band-independent scaling factor is certainly not conclusive, and one may wish to take care in which scaling factor to use for each band. In particular, bands for which there are no experimental intensity data available can not be considered to be truly represented well in UYT2 and the lack of exhaustive absolute intensity knowledge for SO$_3$ limits our ability to effectively correct for the disagreements observed. Nevertheless it is hoped that the scaling factor improves the \ai\ intensity values produced in the UYT2 line list. Further work, probably starting with a systematic \ai\ study of the type recently performed for H$_2$S by \citet{jt607}, is required in order to fully investigate the source of this discrepancy. An experimental determination of individual line intensities would also be extremely helpful. The increased size of the basis set, the computation of rovibrational energies up to $J$ = 130, and the increased spectral range of line strength calculations allows for UYT2 to be used in the simulation of spectra between 0 $ < \nu \leq$ 5000 \cm, with approximately 90\%\ completion at $T$ = 773.15 K (500 C), however calculated cross sections for the region beyond 4500 \cm\ should be treated cautiously, and will have to be further investigated. Given that this is the largest data set of its kind for \sothree, it is recommended that UYT2 be used in the production of cross sections at room-temperature, and up to $T$ = 773.15 K, for both astronomical and other applications. The UYT2 line list contains 21 billion transitions. This makes it use in radiative transport modelling computationally challenging. Work on an even larger methane line list \citep{jtCH4} suggests that it should be possible to split the list into a temperature-dependent but pressure-independent, background cross section which is used to augment a hugely reduced list of stong line whose profiles are treated in detail. This idea will be explored further in the future.

1607.01242

1607

1607.03779_arXiv.txt

Understanding the interaction of primordial gravitational waves (GWs) with the Cosmic Microwave Background (CMB) plasma is important for observational cosmology. In this article, we provide an analysis of an effect apparently overlooked as yet. We consider a single free electric charge and suppose that it can be agitated by primordial GWs propagating through the CMB plasma, resulting in periodic, regular motion along particular directions. Light reflected by the charge will be partially polarized, and this will imprint a characteristic pattern on the CMB. We study this effect by considering a simple model in which anisotropic incident electromagnetic (EM) radiation is rescattered by a charge sitting in spacetime perturbed by GWs and becomes polarized. As the charge is driven to move along particular directions, we calculate its dipole moment to determine the leading-order rescattered EM radiation. The Stokes parameters of the rescattered radiation exhibit a net linear polarization. We investigate how this polarization effect can be schematically represented out of the Stokes parameters. We work out the representations of gradient modes (E-modes) and curl modes (B-modes) to produce polarization maps. Although the polarization effect results from GWs, we find that its representations, the E- and B-modes, do not practically reflect the GW properties such as strain amplitude, frequency and polarization states.

} Ever since its experimental detection \cite{penzias1965} and subsequent interpretation as a relic of the earliest epoch of the universe \cite% {dicke1965}, studies of the Cosmic Microwave Background (CMB) have been crucial for constraining cosmological models. Present-day space-based measurements of the angular power spectrum of the CMB temperature distribution on the sky are able to constrain the parameters included in the $\Lambda $CDM model of cosmology \cite{bahcall1999} with statistical uncertainties down to a few per cent \cite{planck2014}. Of special importance is the fact that state-of-the-art cosmological observations can, at least in principle, place constraints on the large family of models of cosmic inflation \cite{martin2014} which is a cornerstone of hot big-bang cosmology. A key prediction made by models of inflation is the occurrence of primordial gravitational waves (GWs). Even though not observable directly (currently), they should be detectable indirectly via a characteristic linear polarization of the CMB \cite{polnarev1985}. CMB polarization can be decomposed into contributions from gradient modes (E-modes) and curl modes (B-modes), with the latter ones being excited by either tensor perturbations, i.e., propagation of primordial GWs through the plasma emitting the CMB, or gravitational lensing by foreground matter \cite{zaldarriaga1997, kamionkowski1997}. Polarimetric observations \cite{trippe2014} of the CMB at radio frequencies around 100~GHz have been numerous, with different classes of observations aimed at different polarization modes (E or B) and angular scales. A first observation of E-mode polarization was achieved by the Degree Angular Scale Interferometer in 2002 \cite{kovac2002}. B-mode polarization due to gravitational lensing was first detected by the South Pole Telescope in 2013 \cite{hanson2013} and has since been observed by POLARBEAR \citep{polarbear2014, polarbear2015}, ACTPol \citep{act2014}, and BICEP1 \citep{bicep12014}. A detection of B-mode polarization caused by primordial GWs was claimed by the BICEP2 collaboration in 2014 \cite{bicep2014} but was shown to be likely caused by interstellar dust soon thereafter \cite{flauger2014}. The aforementioned studies aim at comparisons of theoretical polarization -- specifically B-mode -- patterns to observational ones in order to constrain cosmological quantities. In a nutshell, the polarization patterns studied so far originate as follows: GWs propagating through the CMB plasma cause local quadrupole anisotropies in the radiation field; this in turn causes characteristic polarization patterns in the light scattered at CMB plasma electrons. However, at least a priori one should also expect polarization immediately from interactions between GWs and individual electric charges. An electric charge located in the path of propagation of a GW will be forced into an oscillatory motion by the wave. This should lead to characteristic linear polarization of light scattered by the charge. We herewith present what appears to be the first quantitative analysis of this effect. We work out the representations of gradient modes (E-modes) and curl modes (B-modes) to produce maps of the polarization patterns resulting from scattering at a single electric charge. Throughout this work, we adopt the negative metric signature $(-,+,+,+)$, and the Minkowski metric is given by $\eta _{\mu \nu}=\mathrm{diag}\left(-1,1,1,1\right) $.

We present an analysis of an apparently as yet overlooked effect in the interaction between (primordial) GWs and individual electric charges. We computed the Stokes parameters for the EM radiation rescattered by a charge that is forced into oscillatory motion by GWs propagating through the CMB plasma: as shown by Eqs. (\ref{I01}) -- (\ref{V01}) and (\ref{i0}) -- (\ref{v0}) in Section \ref{Stokes}, the Stokes parameters exhibit a net linear polarization -- as expected. The GWs \emph{open polarization channels} for the rescattered radiation by setting a charge in motion along particular directions. In other words, the GWs provide for polarization in the photon-electron scattering by agitating a charge along particular directions. Taking the $\mathcal{O}\left( h\right) $ terms into consideration, the resulting contribution to our Stokes parameters will be as shown in Eqs. (\ref{Is}) -- (\ref{Vs}), which is $\mathcal{O}\left( h^{2}\right) $. In our analysis, due to the assumption of linearized gravity in Section II A, $\mathcal{O} \left(h^{2}\right)$ would be considered as very small, and therefore the terms of this order may be disregarded for the final expressions of the Stokes parameters as in Eqs. (\ref{I01}) -- (\ref{V01}). However, until we are able to estimate how large or small the strain $h$ is in the CMB regime [unlike its estimates for present time observations as in Figure 3], it is not clear whether the negligence of $\mathcal{O}\left( h^{2}\right) $ will make any significant difference (or not) in our final E-mode/B-mode maps. Our analysis is presently limited to modeling single-electron scattering; accordingly, we are not yet able to provide a reliable estimate of the amount of polarized flux to be expected. Evidently, reliable amplitude estimates will be crucial for attempts to observe the effect we describe in this study; current B-mode polarization studies have sensitivities of about $0.3\,\mu{\rm K}\sqrt{s}$ (in CMB temperature units, using BICEP2 as benchmark; \cite{bicep2014}). If our measurement were indeed sensitive enough to tell the values of the strain $h$ and the frequency $\omega $ of the GWs in the CMB regime, then we should rather express our Stokes parameters using Eqs. (\ref{Is}) -- (\ref{Vs}). This way, we would be able to see the properties of the GWs directly from our E-mode/B-mode polarization maps through the Stokes parameters. In particular, the second parameter $Q$ would take different values, depending on the GW polarization states, $+$ and $\times $, while the first parameter $I$ would be the same regardless of the states. This implies that the linear polarization of our CMB radiation would be quantified differently, depending on the GW polarization states. That is, $Q_{+}=Q_{\mathrm{o}}+\mathcal{O}\left( h^{2}\right) $ and $Q_{\times }=Q_{\mathrm{o}}-\mathcal{O}\left( h^{2}\right) $ while $I_{+}=I_{\times }=I_{\mathrm{o}}+\mathcal{O}\left( h^{2}\right) $, and therefore we find the ratios, $Q_{+}/I_{+}\neq Q_{\times }/I_{\times }$. This means that the GW-induced polarization should show this distinctive feature, which will distinguish itself from other degenerate cases, where polarization could possibly have been induced from non-GW sources. In order to analyze the CMB polarization on the sky, we have redefined the Stokes parameters on a sphere such that their representations can be expressed as spin-weighted spherical harmonics; namely, spin $\pm 2$ quantities. The representations yield the two polarization patterns, \textquotedblleft electric\textquotedblright\ E-mode and \textquotedblleft magnetic\textquotedblright\ B-mode patterns \cite{Hu}. They are constructed by means of Eqs. (\ref{amp}), (\ref{angE}) and (\ref{angB}), and graphically represented in Figure \ref{fig4} (see \textbf{Graphic representations 1}) in Section \ref{sphere}. While these patterns represent the local signature from scattering over the sphere, in real-world observations, the polarization patterns on the sky are not simply the local signature from scattering but are modulated by plane wave fluctuations on the last scattering surface \cite{Hu}. Thus we have modified the representations of the Stokes parameters such that they contain the properties of the plane wave projected into the spherical sky. The modified representations are spin $0$ quantities, and they also yield the E-mode and\ B-mode polarization patterns \cite{zaldarriaga1997, kamionkowski1997, Kim}. They are constructed by means of Eqs. (\ref{amp1}), (\ref{angE1}) and (\ref% {angB1}), and graphically represented in Figures \ref{fig5}, \ref{fig6} and % \ref{fig7} (see \textbf{Graphic representations 2}) in Section \ref{plane}; with different conditions of the plane wave projection into the spherical sky. Overall, GWs can generate the gradient- and curl-like E- and B-modes as expected. Even though, notable differences between \textbf{Graphic representations 1 }and\textbf{\ 2} are observed due to the plane wave modulation; parameterized by $\rho $, which defines the plane wave projection into the spherical sky. With small to medium $\rho $, which corresponds to the plane wave modulation at \textit{long} to \textit{medium} wavelengths or at \textit{small} to \textit{medium} comoving distances to the last scattering surface, the E-modes and B-modes show broken symmetric and anti-symmetric patterns, respectively in \textbf{Graphic representations 2}. However, with large $\rho $, which corresponds to the plane wave modulation at \textit{short} wavelengths or at \textit{large} comoving distances to the last scattering surface, the E-modes and B-modes take proper symmetric and anti-symmetric patterns. In fact, in the limit $\rho \gg 1$, the E-modes and\ B-modes in \textbf{Graphic representations 2} scale to their counterparts in \textbf{Graphic representations 1}, apart from the factor $\rho ^{2}\sin ^{2}\theta $ in the amplitude \cite{Chiang}. In both \textbf{Graphic representations 1 }and\textbf{\ 2}, however, the periodicity of the patterns in $\phi $ for the E-modes and B-modes remains the same, i.e. $\pi /2$. In this work we have restricted our attention to the case of a single charged particle under the influence of GWs passing through the CMB. We have investigated how the EM radiation is rescattered by the charge being agitated by the GWs and how the rescattered radiation becomes polarized in the CMB. More realistic astronomical scenarios, however, would involve multiple charged particles interacting with each other. At the same time, as we assume that our universe was expanding during the epoch of reionization, the interaction between the particles would be significantly affected by the scale factor of the expanding universe. Therefore, to handle a situation with the interaction properly, the motion of the particles should be described in the spacetime geometry which is prescribed by the perturbed Friedmann-Robertson-Walker metric rather than by the perturbed flat spacetime metric. Unlike the case of a single particle, dynamic evolution of the multiple particles in this setup should be treated statistically, e.g. using the Boltzmann equation \cite{zaldarriaga1997, kamionkowski1997, Hu2}. Likewise, astrophysically realistic CMB polarization maps -- which we do not provide here -- result from convolution of the patterns caused by single-particle interaction with the density distribution of the CMB plasma. Furthermore, we do not yet present power spectra in this study. Following \cite {Hu2}, we can determine the $\ell$-mode expressions for the power spectra out of Eqs. (\ref{QpmU3}) and (\ref{Qt}) -- (\ref{Bt}) by means of the Clebsch-Gordan relation. With these we will be able to estimate the E-mode/B-mode contributions in a quantitative manner as shown by Fig. 11 (b) in \cite{Hu}. We plan to address these open issues in a follow-up study.

1607.03779