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1607.01066_arXiv.txt

We consider dynamo action driven by three-dimensional rotating anelastic convection in a spherical shell. Motivated by the behaviour of the solar dynamo, we examine the interaction of hydromagnetic modes with different symmetries and demonstrate how complicated interactions between convection, differential rotation and magnetic fields may lead to modulation of the basic cycle. For some parameters, Type~1 modulation occurs by the transfer of energy between modes of different symmetries with little change in the overall amplitude; for other parameters, the modulation is of Type~2 where the amplitude is significantly affected (leading to grand minima in activity) without significant changes in symmetry. Most importantly we identify the presence of `supermodulation' in the solutions where the activity switches chaotically between Type~1 and Type~2 modulation; this is believed to be an important process in solar activity.

The origin of magnetic activity in stellar interiors is a fundamental problem of magnetohydrodynamics. The global solar magnetic field oscillates with a mean period of twenty-two years (leading to an eleven-year activity cycle) and is believed to be generated via a dynamo acting (at least in part) deep within the Sun. The Sun's magnetic field is largely dipolar; i.e. the mean azimuthal field that leads to the formation of active regions is generally antisymmetric about the equator. However, when this field is weak at the end of a cycle, it takes on a more mixed character, with a quadrupole component that becomes significant \citep{sokoloff1994}. Furthermore, direct observations and proxy data demonstrate that the amplitude of the solar cycle is modulated on longer time scales. There is indeed a period of reduced activity between 1645 and 1715 ---~the Maunder minimum~--- when the occurrence of sunspots was much reduced \citep{eddy1976,usoskin2015}. Analysis of the abundances of the cosmogenic isotopes \ce{^{10}Be} in polar ice and \ce{^{14}C} in tree rings reveals 27~grand minima in the past 11\,000~yr, separated by aperiodic intervals of approximately 200~yr \citep{usoskin2013,cracken2013}. A key observation for our understanding of the processes leading to modulation is that as the Sun emerged from the Maunder minimum, sunspots were largely restricted to the southern hemisphere, showing that the magnetic field emerged with a mixed character with both dipole and quadrupole components \citep{sokoloff1994}. Moreover, between 1750 and 1775, the solar magnetic field took on a more quadrupolar character, with sunspots appearing at the equator \citep{arlt2009}. There is now evidence from the cosmogenic isotope records that the Sun switches on a long time scale between strong modulation with clusters of deep grand minima and weaker modulation, which can be associated with symmetry breaking. This `supermodulation' is an example of chaotic (though deterministic) modulational effects \citep{weisstobias2016}. Evidence for modulation in other stars arises from the long-term monitoring of the CaII H+K flux of solar-type stars started by Wilson in 1968. The so-called Mount Wilson Observatory survey provides a panel of different stellar activities, in which \SI{60}{\percent} of stars exhibit periodic variations, and \SI{25}{\percent} show irregular or aperiodic variability \citep{baliunas1998,olah2009}. Evidence for changes of symmetry in young rapidly rotating stars is also now beginning to emerge \citep{hlrk2016}. Stellar magnetic fields are thought to be maintained against ohmic dissipation by dynamo action through the flow of an electrically conducting fluid \citep{moffattHK}. Although it is known that systematic activity can be generated through the interaction of turbulent flows with rotation, shear and magnetic fields, no satisfactory, self-consistent nonlinear model of dynamo action is currently available \citep{jonesetal2010,char2014}. Direct numerical simulations aimed at understanding these interactions are restricted to parameters well away from those pertaining to stellar interiors (with Reynolds numbers and magnetic Reynolds numbers (\Rm{}) orders of magnitudes smaller than would be realistic). For this reason, much attention has been focused on mean-field models of dynamo action \citep{krause1980}. In this paradigm, only the large-scale flows and magnetic fields are modelled, with small-scale interactions being parameterised via transport coefficients such as the $\alpha$-tensor and the turbulent diffusivity. Although there are many issues with the mean-field formalism ---~primary among these is whether mean fields can ever be seen at high \Rm{} or whether the solution is dominated by the fluctuations~--- these models are of use in describing the dynamics of mean fields once they have been generated. In particular, the mean-field equations naturally respect the symmetries of the underlying rotating spherical system \citep{knobloch1994}, and capture the nonlinear interactions between magnetic modes of different symmetries and the underlying large-scale velocity field that is driving the dynamo. Mean-field dynamo models have demonstrated that modulation of the basic cycle may occur through stochastic fluctuations in the underlying transport coefficients \citep{schmitt1996,choudhuri2012,hazra2014} or more naturally via nonlinear interactions inherent in the dynamo equations leading to chaotic (though deterministic) modulation \citep{pipin1999,bushbymason2004}. The type of modulation can be classified according to the key nonlinear interactions that are primarily responsible \citep{tobias2002}. In the first (Type~1 modulation) magnetic modes of different symmetry (e.g. dipole and quadrupole modes) interact to produce modulation of the basic cycle, with significant changes in the symmetry (parity) of solutions. This behaviour is similar to that seen in the sunspot record over the past 300~years. In the second (imaginatively termed Type~2 modulation) a magnetic mode with a given symmetry undergoes modulation via interaction with a large-scale velocity field; here changes in the amplitude of the basic cycle occur with no significant changes in the symmetry of solutions. Recently, \citet{weisstobias2016} have argued from analysis of cosmogenic isotope records that both of these modulational mechanisms have been at play in the solar dynamo, leading to `supermodulation' on long time scales. The precise modulational effects important in the system are sometimes model-dependent. For this reason, progress can also be made by considering low-order systems based on symmetry considerations \citep{knobloch1996,knobloch1998,weiss2011}. These models demonstrate that the dynamics found in the ad hoc mean-field models is robust and may be expected in simulations of the full three-dimensional dynamo system. Symmetry arguments have also proved useful in explaining the dynamics of dynamo experiments and the geodynamo where the first bifurcation is stationary \citep{petrelis2009}. In this paper, we present the results of three-dimensional numerical solutions of dynamos driven by anelastic convection in a spherical shell. We do not attempt to model solar convection directly, as this is well beyond the scope of modern-day computations. Rather, we focus on the symmetries and nonlinear interactions that lead to modulation in dynamos, and provide examples of the basic types of modulation and of supermodulation. These results are important for our understanding of magnetic field generation via dynamo action, not only in late-type stars, but also in other astrophysical objects such as planets.

In this paper we have examined the hydromagnetic interactions between dynamo modes generated by rotating anelastic convection in a spherical shell. Motivated by direct and indirect observations of solar magnetic activity, our primary aim was to investigate the interactions between modes with different equatorial symmetries. Mathematically these dynamos display a dynamical behaviour reminiscent of the results obtained with (axisymmetric) mean-field models or low-order systems, with the caveat for the comparison being that the dynamo solutions presented here are dominated by a non-axisymmetric ($m=1$) mode. Hemispheric dynamos of the type reported by \citet{groteBusse2000}, and studied in more detail by \citet{gallet2009}, have also been found. The present study demonstrates that this hemispheric configuration is also pertinent to the understanding of the dynamics of oscillatory dynamos, and thus could be relevant to explaining the hemispheric magnetic configuration that has been observed on the Sun at the end of the Maunder minimum \citep{sokoloff1994, beer1998, knobloch1998}. We stress again that all current direct numerical simulations of convective dynamos ---~including those here~--- are far away from what one can imagine as a ``realistic'' parameter regime. There is, therefore, the question of the robustness of these results. Of course, increasing $Ra$ for fixed Ekman number should lead to more disordered states, gradually breaking all symmetries. What happens after this is a matter of conjecture/debate. It is possible to argue that for very high $Ra$ symmetry is re-established on average in the turbulent state and then similar symmetry-breaking interactions may occur in the averaged equations. Support for this comes from the finding of such interactions in mean-field models, which (despite all their drawbacks) retain the symmetry properties of the underlying system. We note that symmetry arguments are therefore very powerful and we expect similar behaviour to be observed in Boussinesq and indeed fully compressible models. Our primary result is that we have demonstrated that the interactions between such modes can lead naturally to a pattern of supermodulation \citep{arltweiss2014,weisstobias2016} where the system alternates between modulation with little change of symmetry (with clusters of deep minima) and modulation that involves significant changes in symmetry. We believe that this is the first demonstration of such an interaction between the two types of modulation leading to supermodulation in the full partial differential equations for convective dynamos.\\ This study was granted access to the HPC resources of MesoPSL financed by the R\'{e}gion \^{I}le-de-France and the project Equip@Meso (reference ANR-10-EQPX-29-01) of the program Investissements d'Avenir supervised by the Agence Nationale pour la Recherche. Numerical simulations were also carried out at the TGCC computing center (GENCI project x2013046698). R.~Raynaud thanks E.~Dormy, C.~Gissinger, L.~Petitdemange and F.~P\'{e}tr\'{e}lis for various discussions. The authors thank N.~O.~Weiss for helpful comments.

1607.01066

1607

1607.00898_arXiv.txt

{In Liu (2015), we propose selecting binary active galactic nuclei (AGNs) candidates using the centroid shift of the images, which is induced by the non-synchronous variations of the two nuclei. In this paper, a known binary AGN (SDSS J233635.75-010733.7) is employed to verify the ability of this method. Using 162 exposures in the $R$ band of \textit{Palomar Transient Factory} (PTF), an excess of dispersion in the positional distribution of the binary AGN is detected, though the two nuclei cannot be resolved in the images of PTF. We also propose a new method to compare the position of the binary AGN in PTF $g$ and $R$ band and find the difference is highly significant even only with 20 exposures. This new method is efficient for two nuclei with different spectral energy distributions, e.g., type I + type II AGN or off-set AGN. Large-scale surveys, e.g., the Panoramic Survey Telescope and Rapid Response System and the Large Synoptic Survey Telescope, are expected to discover a large sample of binary AGN candidates with these methods.} {}

Supermassive black holes (SMBHs) are discovered in the majority of massive galaxies (Kormendy \& Richstone 1995). They should have grown through mergers and gas accretion (Volonteri et al. 2003; Di Matteo et al. 2008; Kormendy \& Ho 2013). Thus, binary active galactic nuclei (AGNs) \footnote{In this paper, binary AGNs generally stand for any double sources discovered by imaging or spectroscopy. We do not strictly distinguish binary AGNs from dual AGNs or AGN pairs. } are expected to be common, since galaxy mergers can trigger the activity of SMBHs (Begelman et al. 1980; Hernquist 1989; Kauffmann \& Haehnelt 2000; Hopkins et al. 2008). Although hundreds of binary AGNs are found with separations of tens of kiloparsec (kpc), only about ten binary AGNs are known with separations $\sim$10 kpc (Junkkarinen et al. 2001; Komossa et al. 2003; Hennawi et al. 2006; Fu et al. 2011; Koss et al. 2011; Mazzarella et al. 2012; Liu et al. 2013; M{\"u}ller-S{\'a}nchez et al. 2015). The occurrence rate of kpc-scale binary AGNs ($\sim$1\%) is lower than the expected value by a factor of ten if each major merger can induce a binary AGN (Yu et al. 2011). Non-simultaneous activity and the gas content of galaxies are invoked to reconcile this apparent discrepancy (Foreman et al. 2009; Yu et al. 2011; Van Wassenhove et al. 2012). On the observational aspect, a large and complete sample is important for the assessment of this discrepancy. Serendipitously discovered binary AGNs are rare and not sufficient to build a statistically meaningful sample. A systematic method is required for identifying kpc-scale binary AGN candidates. Double-peaked narrow lines (e.g., double-peaked [O III] lines) have been utilized to select kpc-scale binary AGNs (Wang et al. 2009; Liu et al. 2010; Smith et al. 2010; Shen et al. 2011). However, double-peaked narrow lines are usually produced by the gas dynamics (Fu et al. 2012; Blecha et al. 2013; M{\"u}ller-S{\'a}nchez et al. 2015). Moreover, the line shift of a binary AGN in a face-on orbit is small and undetectable. Therefore, an efficient method is needed to well complement spectroscopy methods. In Liu (2015, hereafter Paper I), we show that the imaging centroid of a binary AGN will shift due to the non-synchronous variation of the two nuclei. Thus, such binary AGNs can be revealed by multi-epoch observations, even if the separation is smaller than the angular resolution. This method utilizes the violent variation of AGNs; thus, it is more suitable for two type I AGNs or blazars with strong variability and its efficiency is low for type I + type II binaries. Therefore, we still need an efficient method for the pairs including type II AGN or a galactic nucleus without an AGN. The continuum of type I AGN is bluer than type II AGN or a galactic nucleus without an AGN, which is another distinguishable property of AGNs that can be used to uncover binary AGNs. If the spectral energy distributions (SEDs) of two nuclei are significantly different and two or more filters are available, the images of the blue and red filter are more dominated by the bluer and redder nucleus respectively. As a result, the centroid of the binary in the blue (red) filter should shift towards the bluer (redder) nucleus compared with the centroid in the red (blue) filter. This new method should be valid for two nuclei with different SEDs, e.g., type I + type II AGN, type I AGN+ broad absorption line (BAL) quasar, and type I AGN + galaxy (offset AGN). It is also appropriate to two type II nuclei if the colors of the two host galaxies are remarkably different. As we mentioned in Paper I, the fast variation in jet or the reflection from a cloud near a single AGN can mimic another ``nuclei'' in the image and thus contaminates the candidates selected by the centroid shift. Follow-up high resolution imaging and spectroscopy are still required to confirm the candidates. Therefore, before hunting for new binary AGNs, we would like to select a known binary AGN to verify the two methods. In Sect. 2, we describe the target and data selection. In Sect. 3, the procedure for calculation of the centroid in one band is explained. In Sect. 4, the result of the centroid shift between two bands is shown. In Sect. 5, we discuss the implications of the results and present our conclusions.

Using a confirmed binary AGN, we have tested the validity of the two centroid methods. Both of the two methods have revealed the signal of two nuclei. The significance of the method using one filter is relatively low (p-value=0.0075), since the variation of the BAL nucleus is not as strong as the type I nucleus. For the method using two filters, the distributions of the residual in $R$ and $g$ band are significantly different (p-value=$1.6\times10^{-39}$) and consistent with the expectation from the SEDs of the two quasars. Actually, if we randomly select 20 exposures of $R$ and $g$ band respectively, the significance of the difference between the distributions of $R$ and $g$ band is still at $10^{-9}-10^{-6}$ level. These results indicate the success of the centroid method in uncovering the known binary AGN and its potential power in discovering new binary AGNs. During the merging process of galaxies, the enhanced inflow may induce strong star formation and significant obscuration. Thus type I+type II or type I+BAL cases should be more common than the type I+type I case (Sanders et al. 1988; Treister et al. 2010). The so-called ``offset" AGNs selected by the multi-filter method are also thought to be the important tracer of the merging process and valuable for the understanding on the fueling of AGNs (Comerford \& Greene2014; Steinborn et al. 2015). However, the multi-filter method is not helpful for two type I or type II AGNs with similar SEDs. For the multi-filter method, the relative astrometric error between two bands is crucial. Pier et al. (2003) calibrated the astrometry of SDSS images and found that the standard deviation of the distribution of position differences between $r$ filter and other filters is about $0.025''$, which is better than the astrometric error of $r$ filter against the USNO CCD Astrograph Catalog by a factor of 2. Our method only needs the relative astrometry in arcmin scale and thus the accuracy should be much better. Actually, the distributions of $\tilde{x}^r$ and $\tilde{y}^r$ of $R$ and $g$ band are well consistent with each other. {A binary system such as SDSS J233635.75-010733.7 is hard to discover by optical spectroscopy due to its high redshift ($z=1.285$). The normally employed lines, e.g., [O III] and H$\beta$ lines, are already shifted into the infrared band. A large scale and deep infrared spectroscopy survey on AGNs is still not available at present. However, the centroid methods proposed here are not limited by the redshift in this content if the astrometric error is small enough to identify the centroid shift. The multi-filter method does not require strong variations. Thus, even the exposures in the same night are helpful for identifying the relative shift between different bands. A large sample of binary AGNs should be quickly established by large-scale surveys, e.g., the Panoramic Survey Telescope and Rapid Response System and the Large Synoptic Survey Telescope.

1607.00898

1607

1607.07476_arXiv.txt

We have analyzed images from the VST ATLAS survey to identify candidate gravitationally lensed quasar systems in a sample of WISE sources with $W1 - W2 > 0.7$. Results from followup spectroscopy with the Baade 6.5 m telscope are presented for eight systems. One of these is a quadruply lensed quasar, and two are doubly lensed systems. Two are projected superpositions of two quasars at different redshifts. In one system two quasars, though at the same redshift, have very different emission line profiles, and constitute a physical binary. In two systems the component spectra are consistent with the lensing hypothesis, after allowing for micro-lensing. But as no lensing galaxy is detected in these two, we classify them as {\it lensless twins}. More extensive observations are needed to establish whether they are in fact lensed quasars or physical binaries.

Doubly and quadruply lensed quasar systems are valuable for widely disparate purposes. Treu and Marshall (2016) present a current survey of the use of time delay measurements for cosmography. The micro-lensing of lensed quasars can be used to determine sizes for the emitting regions of quasars (Rauch and Blandford 1991; Agol and Krolik 1999; Pooley et al, 2007; Blackburne et al 2011) and to measure the dark matter fraction in lensing galaxies (Schechter and Wambsganss 2004, Pooley et al 2012, Jimenez Vicente et al 2015). For each of these efforts the accuracy achieved is limited by the relatively small number of lensed systems. The most productive lensed quasar discovery program to date has been the Sloan Digital Sky Survey Quasar Lens Search, henceforth SQLS (Inada et al. 2012), which yielded a statistical sample of 26 lensed quasar systems brighter than a limiting magnitude $i_{lim} =19.1$ over 8000 square degrees in the redshift range $0.6 < z < 2.2$. An additional 36 systems that did not satisfy all of the selection criteria were also catalogued. Of 62 systems {\it in toto}, 40 were newly identified. The ATLAS survey, carried out with VLT Survey Telecope (Shanks et al. 2015), promises to yield comparable if not greater numbers of lensed quasar systems. Its $ugriz$ limiting magnitudes are nearly identical to those of SDSS. While its $ugriz$ photometry covers only 4500 square degrees, the typical ATLAS seeing is 3/4 that of SDSS (Shanks et al. 2015), permitting the discovery of quasar pairs with smaller separations. We have undertaken a search for lensed quasars in the ATLAS survey, and report here the first newly discovered lensed quasars: a quadruple system and two doubly lensed quasars. We also report two systems that have two nearly identical quasar spectra but for which no lensing galaxy has been detected. In \S2 we outline our method for identifying candidate lensed quasar systems, using the WISE catalog and ATLAS $ugriz$ cutouts to choose candidates for spectroscopic and direct imaging followup. The method will be described in greater detail in a forthcoming paper. In \S3 we describe direct imaging and spectroscopic followup observations of seven cadidate systems obtained with IMACS on the Baade 6.5 m telescope of the Magellan Observatory. In \S4 we analyze these observations. In \S5 we present simple models for the the newly discovered quadruple system, WISE 2344-3056 and for one of the doubles, WISE 2304-2214. We discuss our method and results in \S6 and summarize our findings in \S7.

\subsection{The method} We have searched for gravitationally lensed quasars by analyzing VST-ATLAS image cutouts of red WISE sources. Those that could be consistently split into two nearly pointlike objects with roughly constant flux ratios across multiple filters and quasar-like colors in those filters (specifically $u - g < -0.5$) were selected as candidates for followup spectroscopy and imaging. Three new lenses, two pairs of lensless twins, one binary quasar and two projected quasar pairs were found. Owing to bad weather during 2015 and 2016, only a fraction of the candidates (albeit some of the best) were observed. The list of candidates is likely to increase considerably with a) completion of the ATLAS survey, b) application of the method to fainter WISE sources and c) incorporation of a more sophisticated scheme for gauging whether the $ugriz$ colors of a particular pair of objects are quasar-like. The method could be extended to splitting sources into triples, with a primary goal of more readily identifying quadruply lensed quasars. We encountered an unanticipated bottleneck in the speed with which the OmegaCam Science Archive can produce cutouts -- something on the order of one cutout per second, far more than one might think necessary for retrieving 2500 pixels. The servers for the DES and KiDS surveys are not qualitatively faster. This casts a pall on programs that might require $10^7$ or $10^8$ cutouts. While it might be difficult to retool existing archives to speed up the process, we imagine that future systems will deal with such programs more efficiently. Our method (modified for use with $grizY$ photometry) is one of several now being used to search for lenses in the Dark Energy Survey (Agnello et al 2015; Ostrovski et al 2017), which ought to produce at least as many lensed quasars as ATLAS. As no one method will be perfect, comparison of the results will shed light on their relative strengths and weaknesses. \subsection{Lensless Twins} The VST-ATLAS survey has better seeing than SDSS, allowing, in principle, the discovery of less widely separated lensed quasars. But the galaxies that produce close pairs are fainter than those that produce wide pairs and are more crowded by the lensed quasar images, making it more difficult to identify them. We have argued that it is premature to classify either of our ``lensless twin'' quasars as binary quasars -- fraternal twins -- despite the fact that no lensing galaxy has as yet been identified. In Figure~8 we show the two systems, with the component spectra shifted to overlap. For WISE 0326-3122 the agreement is nearly perfect, while for WISE 1051-1142 the differences are consistent with what one would expect for a micro-lensed system. We note that WISE 2329-1258 was at first classified as a lensless twin system, but in the time since the original submission of this paper has been reclassified as lensed quasar with the detection of a lensing galaxy by other investigators. With sufficiently deep direct images in sufficiently good seeing one can set upper limits on the lensing galaxy that rule out the lensing hypothesis, but only with extensive modeling of lensing scenarios. If higher signal-to-noise spectra, or spectra at other wavelengths were to show significant differences in the line profiles, they would again rule out the lensing hypothesis. Confirmation of the lensing hypothesis might come either from identification of the lensing galaxy or from correlated variations in the fluxes, as might be obtained from synoptic observations with the LSST. \bigskip

1607.07476

1607

1607.02779_arXiv.txt

Understanding properties of the first sources in the Universe using the redshifted \HI ~21-cm signal is one of the major aims of present and upcoming low-frequency experiments. We investigate the possibility of imaging the redshifted 21-cm pattern around the first sources during the cosmic dawn using the SKA1-low. We model the \HI ~21-cm image maps, appropriate for the SKA1-low, around the first sources consisting of stars and X-ray sources within galaxies. In addition to the system noise, we account also for the astrophysical foregrounds by adding them to the signal maps. We find that after subtracting the foregrounds using a polynomial fit and suppressing the noise by smoothing the maps over $10\arcmin - 30\arcmin$ angular scale, the isolated sources at $z \sim 15$ are detectable with $\sim 4 - 9 \, \sigma$ confidence level in 2000 h of observation with the SKA1-low. Although the 21-cm profiles around the sources get altered because of the Gaussian smoothing, the images can still be used to extract some of the source properties. We account for overlaps in the patterns of the individual sources by generating realistic \HI ~21-cm maps of the cosmic dawn that are based on $N$-body simulations and a one-dimensional radiative transfer code. We find that these sources should be detectable in the SKA1-low images at $z = 15$ with an SNR of $\sim 14 (4)$ in 2000 (200) h of observations. One possible observational strategy thus could be to observe multiple fields for shorter observation times, identify fields with SNR $\gtrsim 3$ and observe these fields for much longer duration. Such observations are expected to be useful in constraining the parameters related to the first sources.

\label{intro} Detection of the first sources of radiation in the universe which appeared during the ``cosmic dawn'' is at the forefront of modern observational astronomy. It is believed that these sources formed within the dark matter haloes sometime around redshifts $z \sim 15 - 20$ \citep{2007ApJ...665..899W, 2010ApJ...716..510G, 2011ApJ...731...54P, wise2012}. Observing these first sources will not only reveal their unknown properties but also help us in understanding their influence on the formation and evolution of astrophysical objects during later epochs. In recent times a large number of galaxies have been detected at redshift $z \gtrsim 6$ using the broad-band colour \citep{Ellis13, Bouwens15} and the narrow-band $\lya$ emission \citep[e.g.,][]{Ouchi10, Hu10, Kashikawa11}. In addition, a significant number of bright quasars have been detected at high redshifts through various surveys \citep{Fan06b, Venemans15}. New space missions in the near future, e.g., the James Webb Space Telescope ($JWST$)\footnote{http://jwst.nasa.gov}, are expected to detect the most faint sources at even higher redshifts. In addition to the above, 21-cm radiation from the neutral hydrogen (\HI) in the intergalactic medium (IGM) can also be used as a probe to detect the very early sources. Motivated by this fact, many of the present low-frequency radio telescopes like the Low Frequency Array (LOFAR)\footnote{http://www.lofar.org/} \citep{van13}, the Precision Array for Probing the Epoch of Reionization (PAPER)\footnote{http://eor.berkeley.edu/} \citep{parsons13}, the Murchison Widefield Array (MWA)\footnote{http://www.mwatelescope.org/} \citep{bowman13, tingay13}, the Giant Metrewave Radio Telescope (GMRT)\footnote{http://www.gmrt.tifr.res.in}\citep{ghosh12, paciga13} etc have dedicated a large amount of their observing resources to detect the signal from the epoch of reionization (EoR). While most of these telescopes are still not able to probe the very early stages of the EoR as they lack the very low-frequency detectors, the future radio telescope like the Square Kilometre Array (SKA)\footnote{http://www.skatelescope.org/} is expected to detect the signal even from the cosmic dawn. While the first generation telescopes are expected to detect the signal from the EoR statistically (e.g., in terms of the rms, power spectrum, skewness etc), the highly sensitive SKA1-low should be able to image the signal from cosmic \HI ~\citep{2015aska.confE..10M, 2015aska.confE..15W}. Recently, many studies have been done using analytical calculations \citep[e.g.,][]{furlanetto04, 2014MNRAS.442.1470P}, semi-numerical simulations \citep{zahn2007, mesinger07, santos08, Thom09, choudhury09, ghara15a, ghara15b}, and full numerical simulations involving radiative transfer \citep{Iliev2006, mellema06, McQuinn2007, shin2008, baek09} to understand the behaviour of the redshifted 21-cm signal from the cosmic dawn and EoR for different source models. Though most of these studies have concentrated in detecting the signal using statistical quantities, it will be interesting to study the detectability using imaging techniques. Some recent attempts have been made to understand the detection possibility of large ionized bubbles with LOFAR, MWA, GMRT \citep{kanan2007MNRAS.382..809D, 2008MNRAS.386.1683G, 2008MNRAS.391.1900D, 2009MNRAS.399L.132D, 2011MNRAS.413.1409M, datta2012a, Datta2012b}. In addition, \citet{2012MNRAS.425.2964Z} show that the redshifted 21-cm signal from the EoR can be detected in low-resolution images with LOFAR. Studies have also been done in the same context to detect the signal in post-reionization epochs with SKA1 \citep{2014JCAP...09..050V}. Our earlier work \citet[][hereafter paper I]{ghara15c} investigated the detectability of very early sources like metal-free Population III (PopIII) stars, galaxies containing Population II (PopII) stars, mini-QSOs and high-mass X-ray binaries (HMXBs) in the presence of system noise and astrophysical foregrounds using a visibility based techniques. The study showed that the SKA1-low should be able to detect the signal from the sources like the PopII stars, mini-QSOs and HMXBs with $\sim 9-\sigma$ confidence by integrating the visibilities signal over all baselines and frequency channels within $\sim 1000$ hours of observation time. Once the signal from the cosmic dawn is detected, the challenge would be to interpret it and understand the properties of the first sources and the surrounding IGM. One probably needs to use some sophisticated parameter estimation method like the Markov chain Monte Carlo (MCMC) to extract the relevant information. However, before getting involved in the complexities of the parameter estimation methods, one needs to set up appropriate observational strategies to detect the signal. Detection of the signal from the cosmic dawn is itself very challenging as it is very weak compared to the system noise and the astrophysical foregrounds. In general, one has to integrate the signal over a large observing time to reduce the noise and also use some efficient foreground subtraction method to recover the signal given that the foregrounds are 4-5 orders of magnitude stronger. In this paper, we explore, in detail, the detection of the early sources during the cosmic dawn in \HI ~21-cm images in the presence of system noise and the foregrounds. Our analysis is based on realistic simulations of the signal, system noise, and the relevant astrophysical foregrounds, followed by predictions related to the detectability of the early sources using the SKA1-low. These predictions would be quite useful to plan for observational strategies for detecting the sources in 21-cm observations. The paper is organized in the following way. In section \ref{simu}, we describe the simulations used in this work. In particular, we describe the model for the sources used in this study in section \ref{source_rt}, while simulating the baseline distribution of the SKA1-low is described in section \ref{ska_base_dist}. The methods to simulate the signal maps, system noise maps and foregrounds maps are described in section \ref{sig_map}, \ref{noi_map} and \ref{FG_map} respectively. The main results of the paper are given in section \ref{res} before we conclude in section \ref{conc}. We choose the Cosmological parameters $\Omegam=0.32$, $\Omega_\Lambda=0.68$, $\OmegaB=0.049$, $h=0.67$, $n_{\rm s}=0.96$, and $\sigma_8=0.83$, which are consistent with the recent $Planck$ mission results \citep{Planck2013}.

\label{conc} We have investigated the detectability of the first sources during the cosmic dawn using imaging techniques through future radio observations with the SKA1-low. Detecting the 21-cm signature of these sources is expected to reveal, at least to some extent, their properties and also the physical state of the surrounding IGM. However, their detection would be significantly challenging because the signal is much too weak compared to the system noise and the astrophysical foregrounds. Our fiducial source model consists of stars within a galaxy along with a mini-QSO type X-ray source. The model for the sources can be parametrized by several unknown parameters, e.g., the stellar mass ($M_{\star}$), the escape fraction of the UV photons ($f_{\rm esc}$), the ratio of X-ray and UV luminosities ($f_X$), the X-ray spectral index ($\alpha$), the age of the source ($t_{\rm age}$), and the redshift of observation ($z$). In addition, we also need to specify the overdensity of the surrounding IGM ($1+\delta$), assuming it to be uniform. The fiducial values of these parameters are taken to be $M_{\star}=10^7 ~\MSUN, f_{\rm esc}=0.1, f_X=0.05, \alpha =1.5, t_{\rm age}=20$ Myr, $z=15$ and $1 + \delta = 1$. We have considered a fiducial observation using the present antenna configuration of the SKA1-low. Assuming that we observe a region at declination $\delta = -30^{\circ}$, we have used the baseline distribution to obtain the ``dirty'' map. We have added the system noise as well as the astrophysical foregrounds (Galactic synchrotron and extragalactic point sources) to the images. Our main aim is to explore whether the images can be used for detecting the signal from the first sources and if one can extract the properties of these sources from the maps. Our main findings are listed below. \begin{itemize} \item If we assume the target source to be isolated, then in the situation where foregrounds can be perfectly subtracted out, it is possible to achieve a signal to noise ratio (SNR) $\sim 11$ for the fiducial source at a redshift of 15 for 2000 h of observation and a frequency resolution of 100 kHz. This SNR is achieved by smoothing the images with a Gaussian filter of size 30\arcmin which helps in reducing the rms of the noise considerably. In general, the SNR increases with increasing width of the Gaussian filter. \item It is not possible to detect the signal in any reasonable observational time without smoothing the maps. Unfortunately, this smoothing alters the intrinsic brightness temperature profile around the sources which in turn makes it difficult to reliably extract their properties from the maps. We find that it is still possible to constrain the parameters $M_\star$, $f_{\rm esc}$ and $1+\delta$, while it will be difficult to extract any information on $f_X$, $\alpha$ and $t_{\rm age}$ from the smoothed $\TB$ profiles. \item Although the expected brightness temperature profiles around different types of sources are different, smoothing the maps makes it difficult to distinguish between these sources. In particular, we find that the smoothed profiles of the different X-ray sources, e.g., mini-QSOs and HMXBs, are similar to the case where there are no X-rays from the galaxy. \item The cosmological 21-cm signal is largely contaminated by the astrophysical foregrounds. In order to account for these, we model the Galactic synchrotron emission and extragalactic point sources \citep{Choudhuri2014MNRAS.445.4351C} and add them to our maps. We then use a third order polynomial fitting method to subtract the foregrounds. We are able to achieve an SNR $\sim 9$ for the fiducial source model which is only $\sim 20\%$ worse than the foreground-free scenario. \item Since the first galaxies are not expected to form in complete isolation, we generate more realistic signal maps from the output of a $N$-body simulation and using a one-dimensional radiative transfer code \citep{ghara15a}. The reionization history is calibrated to recent Planck measurements of the electron scattering optical depth \citep{2015arXiv150201589P}. We apply the same smoothing and foreground removal technique on these maps as discussed above. The SNR of the map at the redshift 15, after the foregrounds subtraction and smoothing with the fiducial filter, is $\sim$ 14 (4) for 2000 (200) h of observations. The corresponding SNR value is 10 (3) at redshift 16. This suggests that a possible observation strategy for the SKA1-low could be to observe multiple fields for small observation time like 200 h. If one is able to detect a $3-\sigma$ signal in any of these fields (after smoothing with filters of widths $\sim 30\arcmin$), then one can perform a deeper observation of $\sim 2000$ h and possibly constrain properties of the first sources along with the surrounding IGM. \end{itemize} Finally, we discuss some of the aspects of the study which need to be addressed in more details. Although we have modelled the foregrounds in a fairly detailed manner, they can be more complex in the actual case. One probably needs to devise more sophisticated methods to disentangle the signal in that case. Our analysis ignores various other complications, e.g., those arising from instabilities in the ionosphere, calibration of the signal, man-made interference, and instrumental systematics. One possible extension of the present work could be to consider all these complexities and develop a complete pipeline to prepare mock data sets for analysis. On the modelling aspect, one needs to work out the signal in different reionization scenarios accounting for the uncertainties in the galaxy formation processes at high redshifts. This could include studying the effects of, e.g., the small mass sources of ionization and heating leading to a relatively early overlap of $\lya$ regions \citep{ghara15a}, alternate reionization scenarios driven by quasars \citep{2015ApJ...813L...8M, 2016MNRAS.457.4051K, 2016arXiv160602719M}. It is also possible to improve the methods used for detecting the signal. In this paper, we have mainly concentrated on the possibility of imaging the 21-cm pattern of the first sources which can be useful, particularly for visual identification, in a situation when the patterns around different sources overlap with each other. However, it is possible that a more efficient search can be performed in the visibility space where the noise is uncorrelated \citep{ghara15c}. In addition, the smoothing filters used in this work have been constructed assuming that we do not have any prior idea of the signal. One could also explore devising more sophisticated filters (e.g., matched filters) which account for the nature of the signal to make a more efficient detection.

1607.02779

1607

1607.03495.txt

Our Galaxy hosts the annihilation of a few $\times 10^{43}$ low-energy positrons every second. % Radioactive isotopes capable of supplying such positrons are synthesised in stars, stellar remnants, and supernovae. % For decades, however, there has been no positive identification of a main stellar positron source leading to suggestions that many positrons originate from exotic sources like the Galaxy's central super-massive black hole or dark matter annihilation. %, but such sources would not explain the recently-detected positron signal from the extended Galactic disk. % Here we show that a single type of transient source, deriving from stellar populations of age 3-6 Gyr and yielding $\sim 0.03 \msun$ of the positron emitter $^{44}$Ti, can simultaneously explain the strength and morphology of the Galactic positron annihilation signal and the solar system abundance of the $^{44}$Ti decay product $^{44}$Ca. % This transient is likely the merger of two low-mass white dwarfs, observed in external galaxies as the sub-luminous, thermonuclear supernova known as SN1991bg-like.

1607.03495

1607

1607.05535_arXiv.txt

We have designed and developed, from scratch, a global circulation model named \texttt{THOR} that solves the three-dimensional non-hydrostatic Euler equations. Our general approach lifts the commonly used assumptions of a shallow atmosphere and hydrostatic equilibrium. We solve the ``pole problem'' (where converging meridians on a sphere lead to increasingly smaller time steps near the poles) by implementing an icosahedral grid. Irregularities in the grid, which lead to grid imprinting, are smoothed using the ``spring dynamics'' technique. We validate our implementation of spring dynamics by examining calculations of the divergence and gradient of test functions. To prevent the computational time step from being bottlenecked by having to resolve sound waves, we implement a split-explicit method together with a horizontally explicit and vertically implicit integration. We validate our global circulation model by reproducing the Earth and also the hot Jupiter-like benchmark tests. \texttt{THOR} was designed to run on Graphics Processing Units (GPUs), which allows for physics modules (radiative transfer, clouds, chemistry) to be added in the future, and is part of the open-source Exoclimes Simulation Platform (ESP; www.exoclime.org).

\label{sec:intro} The observational exploration of planetary atmospheres has shown that planets present a large diversity of climates and atmospheric circulations. This diversity raises the need to build versatile numerical tools capable of interpreting the observational data and help unveil the main mechanisms driving the atmospheric climate and dynamics. These tools have to be self-consistent and based on theory that does not compromise the accuracy of the results under particular planetary conditions. For this purpose, our goal is to develop the first robust Global Circulation Model (GCM)\footnote{Also termed ``general circulation models''.} capable of simulating a vast range of planetary conditions, which will work as a virtual ``planetary lab''. Global Circulation Models (GCMs) solve the complex physical and dynamical equations that include a representation of the evolution of the resolved fluid flow and various idealizations for radiative transfer, dry or moist convection and for heat and momentum's turbulent surface fluxes. These models are powerful tools to simulate self-consistently, for example, the dynamical heat transport in the atmosphere and represent 3D temperature maps of the atmosphere essential to interpret observational data. They have been important to the study of the atmospheric circulation and climate of Earth, of the Solar System planets and more recently of extrasolar planets. In general, these planetary atmospheric simulations have been explored by changing the general parameters that characterize the planets (e.g. rotation rate, distance to the star, planetary radius, atmospheric mass and composition) and adapting the physics package to deal with a wider range of atmospheric conditions (sometimes radically different from Earth). Some of the most successful simulations of the Solar System planetary atmospheres are: Venus (\citealt{2010Lebonnois} and \citealt{2016Mendonca}), Mars (\citealt{1999Forget}), Titan (\citealt{1995Hourdin}), and the giant gas planets like Jupiter, Saturn, Neptune and Uranus (\citealt{1998Dowling}, \citealt{2004Yamazaki} and \citealt{2009Schneider}). Good reviews on recent advances on Earth GCMs can be found in \cite{2007Randall} and \cite{2013Dowling}. The dynamical part of these atmospheric models usually remains the same, however, that the core has been modified in \cite{2010Lebonnois} and \cite{2016Mendonca} to include, for example, variations of the specific heat with temperature for the Venus atmosphere. This numerical exploration of planetary atmospheres also expands to the extrasolar planets, as for example, the works on tidally locked hot Jupiters (e.g., \citealt{2002Showman}; \citealt{2009Showman}; \citealt{2011Heng}; \citealt{2015Kataria}). These models obtained a robust equatorial jet feature in the simulations, which seems to be a consistent phenomenon with the observational shift of the maximum flux in the secondary eclipse seen in e.g., \cite{2009Knutson}. A good review on the models that have been used for extrasolar planets can be found in \cite{2015Heng}. The majority of the 3D atmospheric models for planetary studies used as their basis dynamical cores that were developed to do Earth climate studies, with the exception of a model called EPIC (\citealt{1998Dowling}). EPIC model is the only 3D climate model that we are aware of, that has been developed from ground-up with the main goal of exploring different planet atmospheres. Despite some success in planetary studies, the vast diversity of planetary characteristics observed raises questions about the flexibility of current atmospheric models to represent accurately the atmospheric physics of those planets. Planetary climate models adapted from Earth climate studies have usually included approximations that are Earth-centric. Some of the most commonly-used approximations are the shallow atmosphere and the hydrostatic approximation (described in section \ref{sec:eqs_thor}). In the new platform that we present in this work, we want to avoid making any ad hoc assumption that can compromise the physics of the problem. The amount of detailed information about extrasolar planetary atmospheres is still very limited and not enough to validate 3D climate models. Due to these limitations the modeling of extrasolar planets has to be formulated from first principles. The perfect simulation would be based on universal physical schemes with no artificial forcing that could reproduce all the observational data or help in preparing the observations (e.g., by making predictions). Those ideal ``planetary labs'' still do not exist and we are still in the important process of assessing the robustness of our models for different atmospheric conditions. A careful exploration of the parameter space needed for this process can teach us important lessons on how the atmospheric circulation and climate work. In addition, it continues to be poorly understood how new atmospheric circulations are driven for a large range of astronomical and planetary bulk parameters. Numerically, we also need to improve our knowledge on the balance between physical and numerical sources and sinks of quantities such as angular momentum and total energy for a large diversity of planetary conditions, which can be crucial to assess the accuracy and robustness of the simulations. The dynamical core is the part of the atmospheric models that solves the resolved dynamical fluid equations including thermodynamics and mass conservation. The new and flexible dynamical model \texttt{THOR} is part of the Exoclimes Simulation Platform (ESP), and the code is intended to be freely open-source (see more information in www.exoclime.org). The ESP is divided into two cores, the core that solves the physics equations such as radiative transfer and convection, and the core that solves the fluid equations. The physics core used in this work is very simplified, and more sophisticated schemes are being developed to be included in the physics core. Currently, two physical schemes are being implemented: \texttt{HELIOS} and \texttt{VULCAN}. \texttt{HELIOS} is an efficient and flexible radiative transfer code suitable for 3D climate models (\citealt{2015Malik}). This code will represent the radiative emission/absorption/scattering in the atmosphere due to gas molecules and clouds in the ESP simulations. The scheme uses as input k-distribution tables, which are produced by the open-source program called \texttt{HELIOS-K} (\citealt{2015Grimm}). Other physics modules such as \texttt{VULCAN} will represent the atmospheric chemistry, including thermochemistry and photochemistry (\citealt{2015Tsai}). In addition to the dynamical core presented here, based on a horizontally explicit and vertically implicit (HE-VI) type time integrator, we are also working on a fully explicit Riemann solver (\citealt{2016Grosheintz}). The dynamical core \texttt{THOR} was planned to be flexible, simple to use, and suitable for parallel computations which will allow us to run efficient high-resolution simulations necessary to resolve important instabilities and to improve numerical accuracy. The main goal when developing \texttt{THOR} was to provide a solid dynamical core capable of simulating and helping us to have a better understanding on the atmospheric dynamics and climate of a large diversity of planets. In this work, we present the first version of the 3D atmospheric model \texttt{THOR}. In the next section, we will start describing the spherical grid used to discretize the model domain. In this section, we also describe the necessary modifications to the icosahedral to improve the numerical accuracy of the divergence and gradient operators. In section \ref{sec:eqs_thor}, we present the equations solved by \texttt{THOR} and the methods used for the time and space integrations. We also discuss the main approximations often used in climate models and their validity. The numerical diffusion applied to the prognostic variables (these are variables predicted by the integration of the equations) in \texttt{THOR} is presented and discussed in section \ref{sec:num_diss}. In section \ref{sec:ProfX}, we include the description of the simple physics representations for radiation, convection and boundary layer friction, which were coupled with the new dynamical core. In section \ref{sec:gpu}, we briefly describe the Graphics Processing Unit (GPU) implementation of the new code. In section \ref{sec:simu}, the new model is validated against other models for two distinct atmospheric conditions. Finally, in section \ref{sec:conclusion}, we present the general conclusions.

\label{sec:conclusion} We have developed a new global circulation model for a large diversity of atmospheric conditions. This new platform has passed two important benchmark tests with very different atmospheric conditions: for Earth and for a hot Jupiter-like planet. The grid implemented is able to reach a second-order accuracy and the grid imprinting noise is efficiently removed from the simulations. In this, work we have described the structure of the new complex dynamical model that uses a numerical solver based on an explicit-implicit formulation. This formulation allows for a very good performance since we use a split-explicit time step method coupled with an implicit integration of the vertical momentum, which is numerically constrained by the propagation of the fast sound waves across the fine vertical resolution. The model does not use any of the traditional approximations often used in Earth climate studies that can limit the model flexibility to explore other planetary conditions. This new platform solves the deep non-hydrostatic Euler equations, and more physical representations can be easily added to the original set of equations. The two experiments which were successfully done represent a small sample in the general diversity of planets. However, \texttt{THOR} showed a very good performance in the simulation of these two very distinct atmospheric conditions, which allow us to be in an advantageous position to continue exploring the diversity of planetary parameters that are expected to characterize the atmospheres of solar system and extrasolar planets. In summary, the main advantages of using our new platform against other recent planetary models are: \begin{itemize} \item The resolved atmospheric fluid flow is completely represented and no approximations are used that could compromise the physics of the problem; \item The model uses, for the first time in exoplanetary studies, an icosahedral grid that solves the pole problem; \item The interface is user friendly and can be easily adapted to a multitude of atmospheric conditions. \end{itemize} We have developed a solid basis for a virtual atmospheric laboratory. This model has also been developed keeping the user interface easy to use since we aim to make this code free and open-source to the community (more details about \texttt{THOR} in www.exoclime.net). \newpage

1607.05535

1607

1607.03512_arXiv.txt

We present the discovery of a hot-Jupiter transiting the $V=9.23$\,mag main-sequence A-star KELT-17 (BD+14 1881). KELT-17b is a $1.31_{-0.29}^{+0.28}\,M_\mathrm{J}$, $1.645_{-0.055}^{+0.060}\,R_\mathrm{J}$ hot-Jupiter in a 3.08 day period orbit misaligned at $-115.9\pm4.1$ deg to the rotation axis of the star. The planet is confirmed via both the detection of the radial velocity orbit, and the Doppler tomographic detection of the shadow of the planet during two transits. The nature of the spin-orbit misaligned transit geometry allows us to place a constraint on the level of differential rotation in the host star; we find that KELT-17 is consistent with both rigid-body rotation and solar differential rotation rates ($\alpha < 0.30$ at $2\sigma$ significance). KELT-17 is only the fourth A-star with a confirmed transiting planet, and with a mass of $1.635_{-0.061}^{+0.066}\,M_\odot$, effective temperature of $7454\pm49$\,K, and projected rotational velocity $v\sin I_* = 44.2_{-1.3}^{+1.5}\,\mathrm{km\,s}^{-1}$; it is amongst the most massive, hottest, and most rapidly rotating of known planet hosts.

The properties of planets orbiting high mass stars provide an important piece of the planet formation puzzle. The occurrence rate of giant planets increases with stellar mass \citep[e.g.][]{Johnson:2007,Johnson:2010,Crepp:2011,Jones:2014,Jones:2016}, at least until $\sim 2\, M_\odot$ \citep{Reffert:2015}. Observations of protoplanetary disks also show a correlation between the mass of the host star and the surface density and mass of the protoplanetary disk \citep[e.g.][]{Muzerolle:2003,Natta:2006}, as well as the disk accretion rate \citep{Manara:2016}. As such, the conditions around young, massive stars are more conducive to the formation of giant planet embryos \citep[e.g.][]{Liu:2016}, and may even lead to more massive planets being formed \citep{Mordasini:2012}. Despite the apparent ease of giant planet formation around massive stars, only three transiting planets have been confirmed around A stars to date: WASP-33b \citep{Collier:2010b,Johnson:2015}, KOI-13b \citep{Szabo:2011,Shporer:2011,Johnson:2014}, and HAT-P-57b \citep{Hartman:2015}. Confirming planets around these stars is difficult via traditional techniques: in addition to the low mass and radius ratios of these systems (therefore low radial velocity amplitudes and transit depths), main sequence A-stars have rapid rotation rates and few metal spectral lines, inhibiting precise radial velocity measurements that are typically necessary for the planet confirmation. One successful approach is to perform radial velocity searches for planets around `retired A-stars' \citep[e.g.][]{Johnson:2010,Wittenmyer:2011} -- giants and sub-giants with masses $>1.6\,M_\odot$ that have spun-down over their post-main-sequence evolution, and allow precise radial velocity measurements to be made. These surveys have revealed some intriguing trends, such as the apparent lack of high eccentricity warm-Jupiters around sub-giants \citep{Jones:2014}. Transiting planets offer an unique set of opportunities, such as the characterization of planet radius, orbital orientation, and atmospheric properties, that are not available to planets detected by radial velocities only. A sample of well characterized planets around massive stars is necessary to understand the mass-dependence of planet properties. The Kilodegree Extremely Little Telescope (KELT) \citep{Pepper:2007} is designed to target planets orbiting stars with brightnesses of $8 < V_\mathrm{mag} < 10$: systems around bright stars that are conducive to follow-up characterization. As discussed in \citet{Bieryla:2015}, a direct result of this KELT sample selection is that 55\% of KELT-North targets are hotter than 6250\,K, with masses $\gtrsim 1.3\,M_\odot$ and median rotational velocities of $\gtrsim 20$\,km\,s$^{-1}$ \cite[inferred from the Kepler sample of stellar rotational velocities in][]{Nielsen:2013}. A similar stellar sample will also be surveyed by the TESS full frame dataset \citep{Ricker:2014}. Strategies for confirming planets around massive stars from the KELT survey are therefore directly transferable to future planet candidates from TESS. Transiting planets around rapidly rotating, high mass stars can be confirmed via Doppler tomography. During a transit, the planet occults parts of the rotating stellar disk, thereby distorting the observed spectral line profile of the star. For relatively slowly rotating host stars, this results in a net shift in the apparent velocity of the host star -- the Rossiter-McLaughlin effect \citep{Rossiter:1924,McLaughlin:1924}. In the cases where the rotational broadening of the star is significantly higher than other broadening factors, the shadow of the planet can be resolved as an intrusion in the rotationally broadened line profile of the star, yielding a Doppler tomographic detection of the planet. In addition to the three planets around A-stars confirmed via Doppler tomography, detections have been achieved for 9 more planets: WASP-3b \citep{Miller:2010}, WASP-38b \citep{Brown:2012}, CoRoT-11b \citep{Gandolfi:2012}, HAT-P-2b and Kepler-25c \citep{Albrecht:2013}, KOI-13b \citep{Johnson:2014}, KOI-12b \citep{Bourrier:2015}, KELT-7b and HAT-P-56b \citep{Bieryla:2015,Zhou:2016}. The depth and width of the spectroscopic shadow of the planet is directly correlated with the planet-star radius ratio. In the cases where the depth agrees with that from the photometric transit, we can rule out blend scenarios often associated with transiting planet candidates. This is especially useful in eliminating the scenarios of background eclipsing binaries, where a Doppler tomographic observation will yield no planet detection. Subsequent out-of-transit radial velocities, at the km\,s$^{-1}$ level, are then taken to constrain the nature of the orbiting companion. In this paper, we report the discovery of KELT-17b, a hot-Jupiter transiting a rapidly rotating $(v\sin I_* = 44\,\mathrm{km\,s}^{-1})$ A-star. KELT-17b was discovered in the equatorial field jointly surveyed by the KELT-North \citep{Pepper:2007} and KELT-South \citep{Pepper:2013} observatories. The discovery involves a series of photometric follow-up observations that confirmed and characterized the transit light curve, and spectroscopic monitoring that constrained the mass of the system. Finally, blend false positive scenarios were ruled out by two in-transit spectroscopic observations that confirmed the Doppler tomographic signal induced by the transiting planet.

We report the discovery of KELT-17b, a hot-Jupiter around an A-star discovered by the joint KELT-North and KELT-South survey. KELT-17b is only the fourth hot-Jupiter found transiting an A-star, after WASP-33b, KOI-13b, and HAT-P-57b. The host star is also amongst the most rapidly rotating known transit-planet-hosting stars, with a $v\sin I_*$ of $44.2_{-1.3}^{+1.5}\,\mathrm{km\,s}^{-1}$, only WASP-33 \citep{Collier:2010b}, KELT-1 \citep{Siverd:2012}, KOI-12 \citep{Bourrier:2015}, KOI-13 \citep{Szabo:2011}, KELT-7 \citep{Bieryla:2015}, and HAT-P-57 \citep{Hartman:2015} exhibit faster spin rates. KELT-17, with a mass of $1.635_{-0.061}^{+0.066}\, M_\odot$, is amongst the most massive (hottest) 3\% (0.5\%) of known planet hosts\footnote{NASA exoplanet archive \url{http://exoplanetarchive.ipac.caltech.edu/}}. Blend scenarios for KELT-17b are eliminated by the detection of the Doppler tomographic signal, from which we also measured a projected spin-orbit angle of $\lambda = -115.9\pm4.1^\circ$ for the system. With a mass of $1.31_{-0.29}^{+0.28}\,M_J$, and radius of $1.525_{-0.060}^{+0.065}\,R_J$, KELT-17b is inflated compared to standard evolution models. It receives an incident flux of $4\times 10^9\,\mathrm{erg\,s}^{-1}\mathrm{cm}^{-2}$, larger than the empirical threshold of $2\times10^8\,\mathrm{erg\,s}^{-1}\mathrm{cm}^{-2}$ where radius inflation is observed for the hot-Jupiter sample \citep{Demory:2011}. KELT-17b is one of 26 known transiting hot-Jupiters around a host star hotter than 6250\,K, of which 70\% are spin-orbit misaligned ($|\lambda| > 10^\circ$). In fact, all four hot-Jupiters around A-stars are in severely misaligned orientations\footnote{Multiple $\lambda$ solutions are allowed for HAT-P-57b \citep{Hartman:2015}}. The spin-orbit synchronization timescale for the KELT-17 system is $\sim 10^{11}$ yrs \citep[using Equation 3,][]{Hansen:2012}, so the current system misalignment is unlikely to have been modified by star-planet tidal interactions, and will be stable for the duration of the main-sequence lifetime of KELT-17. KELT-17b is super-synchronous: the host star has a maximum spin period of 1.7 days, while the planet orbital period is $\sim 3.08$ days, as with a number of other systems around F-A stars (CoRoT-3b, CoRoT-11b, HAT-P-56b, KELT-7b, KOI-13b, WASP-7b, WASP-8b, WASP-33b, WASP-38b). In contrast, no Kepler candidates, which are largely around cooler host stars, are found in super-synchronized orbits \citep{Walkowicz:2013}. While the synchronization timescale is long, the orbit circularization timescale should be only $10^7$ years, assuming $Q_\mathrm{planet} = 10^5$, and adopting the circularization timescale from \citet{Goldreich:1966}, so we expect the orbit of the planet to be circular in the present day. The spin-orbit misaligned orientation of KELT-17b means the planet crosses a wide-range of stellar latitudes during the transit, which allowed us to constrain the latitudinal differential rotation of the star. As a result, we find KELT-17 to be consistent with both rigid body rotation and solar-levels of differential rotation ($\alpha<0.30$ at $2\sigma$). An equivalent technique was applied to the transits of HD 189733b \citep{Cegla:2016}, a significantly more difficult case given the near-aligned geometry of the planet $(\lambda = -0.4\pm 0.2^\circ)$ and the low rotation rate of the star $(v\sin I_* = 3.25\pm0.02\,\mathrm{km\,s}^{-1})$. Nevertheless, they were able to rule out rigid-body rotation for the host star. Future Doppler tomographic follow-up of KELT-17 can further refine the differential rotation constraints on the star, and search for nodal precession of the planet's orbit \citep[e.g.][]{Johnson:2015}.

1607.03512

1607

1607.06930_arXiv.txt

In this paper we present a combined photometric, spectroscopic and orbital period study of three early-type eclipsing binary systems: \astrobj{XZ~Aql}, \astrobj{UX~Her}, and \astrobj{AT~Peg}. As a result, we have derived the absolute parameters of their components and, on that basis, we discuss their evolutionary states. Furthermore, we compare their parameters with those of other binary systems and with the theoretical models. An analysis of all available up-to-date times of minima indicated that all three systems studied here show cyclic orbital changes; their origin is discussed in detail. Finally, we performed a frequency analysis for possible pulsational behavior and as a result we suggest that \astrobj{XZ~Aql} hosts a $\delta$ Scuti component.

Classification of eclipsing binaries is performed according to the physical and evolutionary characteristics of their components, in addition to the shapes of their light curves. The degree to which their inner critical equipotential surfaces (Roche lobes) have been filled is a critical parameter for their classification, which helps in the understanding of their physical nature. Semi-detached binaries constitute an important class of objects, with one component filling its Roche lobe. The shape of a typical semi-detached binary light curve is an Algol-type light variation. These objects are important in understanding this stage of evolution when the mass transfer starts to take place and alters the evolution of the components as single stars. Whether or not one or both of the components are in contact with their Roche lobes or very close to filling them is very important in understanding the evolution of interacting binary systems. In order to achieve this goal, light and radial velocity observations are analyzed to determine their absolute parameters (temperatures, masses, radii, and surface potentials in particular). If the mass transfer has started in these systems at some point during their evolution, it also manifests itself as orbital period changes because a transfer of mass would alter a system's moment of inertia.\\ In this study, we present the results of light curve and period change analyses of three Algol-type binary systems: \astrobj{XZ~Aql}, \astrobj{UX~Her}, and \astrobj{AT~Peg}. We derive their absolute parameters from the analysis of their light and radial velocity curves that we obtained at two different observatories. We analyze the differences between the observed (O) and the calculated (C) eclipse timings, occuring due to the changes in their orbital period, by constructing O-C diagrams. We also perform a frequency analysis to investigate the potential pulsation signals in the data of the studied systems. Finally, we discuss the evolutionary states of the components of these systems on Hertzsprung-Russell, Mass-Luminosity and Mass-Radius diagrams (hereafter HRD, MLD, and MRD, respectively). Such thorough analyses for these systems were being performed for the first time within this study. \section[]{Systems} \subsection{\astrobj{XZ Aql}} \astrobj{XZ Aql} (HD 193740, BD-07$^{\circ}$ 5271, GSC 5174-108, SAO 144345) is an Algol type eclipsing binary. It was discovered by \citet{cannon22}. The first detailed description of the light curve (without a plot) was presented by \citet{witkowski25}. \citet{erleksova59} was the first (and, so far the only one) to present the graphical light curve and give the first discussion of the O-C diagram. She proposed two alternative models of the period variation: first, as an abrupt change between JD 2430000 and 2433000, and second as an increase at a constant rate. \citet{pokzla76} discussed the same subject, using a larger number of eclipse timings. \citet{woodforbes63} and \citet{samolyk96} noted a quadratic term in the ephemeris. A detailed discussion of shape of O-C diagram was given by \cite{soydugan06}. They modeled the variation with a sinusoidal superimposed on a secular parabolic change. They attributed the sinusoid variation to the light time effect caused by an unseen third body, and the secular term to a conservative mass transfer from the less massive component to the more massive one, which had already been noted by previous studies. The orbital period of the hypothetical third body was P$_{3}$ = 36.7$\pm$0.6 yr. The spectral type of XZ Aql is A2, found by \citet{cannon22} when the star was discovered. The mass ratio based on radial velocity observations has not been published until now, nor has any light curve solution. \subsection{\astrobj{UX Her}} \astrobj{UX Her} (HD 163175, BD+16$^{\circ}$ 3311, HIP 87643, SAO 103195, GSC 01557-01268) system was discovered by Cannon \citep{pickering08}, who found its spectral type as B9 or A. \citet{zinner13} confirmed the discovery and determined the correct period of the system. Later, \citet{tsesevich44}, \citet{kaho52}, \citet{kurzemnietse52}, \citet{ashbrook52}, \citet{tsesevich54}, \citet{koch62}, revised the elements. \citet{gordonkron63} and \citet{gordonkron65} published the first light curve solution based on the spectroscopic orbit determined by \citet{sanford37} and their own photometric observations. They proposed that the secondary was an evolved, low-mass, late-type star. \citet{hilletal75} estimated the primary component's spectral type as A0V - A3V for different orbital phases. Following the light curve studies of \citet{cester79}, \citet{mardirossian80} and \citet{giurmar81}, \citet{lazaro97} analyzed the first infrared light curves of the system in J, H and K bands, together with published B and V band light curves of \citet{gordonkron65}. They determined the parameters of the system and found that none of the components of the binary filled their Roche lobes. Although \astrobj{UX Her} is a short-period system, they assigned it to a category of slightly detached systems, most of which, as they pointed out, were long-period systems. \citet{gojko06} computed the mass ratio (q = m$_{2}$ / m$_{1}$ = 0.248) as the result of the q-search method. When combined with the results of their B and V band light curve analysis, this mass ratio value suggests a semi-detached configuration of \astrobj{UX Her}.\\ \citet{kurzemnietse52} noted for the first time that the period was variable. \citet{tremko04} first published a period-variation study of the system, which excluded a mass transfer between the components as the cause of the observed variations in the orbital period of the system, since neither of the stars filled their Roche lobes according to their interpretation. They proposed that a low mass (0.3 $M_\odot$) unseen companion, bound to the system, was causing the period to change periodically. \subsection{\astrobj{AT Peg}} Cannon \citep{canpick24} published the first spectroscopic observation of \astrobj{AT~Peg} (HD 210892, BD+07$^{\circ}$ 4824, SAO 127380, GSC 1137-185), and determined its spectral type as A0. The variability of the system was first announced by \citet{schneller31}, who identified it as an Algol-type binary. \citet{guthnick31}, \citet{rugemer34}, \citet{lass35}, \citet{criswal63}, \citet{woodforbes63}, \citet{oburka64a}, \citet{oburka64b}, \citet{oburka65} and \citet{criswal65} published their photometric observations and revised the light elements. \citet{hillbarnes72} published the orbital elements of the system based on the first detailed spectroscopic observations and determined the spectral type as A7 V. They found the mass ratio ($m_{1}$ / $m_{2}$) to be 2.4 and the orbital inclination 76$^{\circ}$.7. \citet{gulmenetal93} found that AT Peg was a semi-detached binary with a later type subgiant secondary component filling its Roche lobe. \citet{maxtedetal94} obtained spectra of the system and determined a spectroscopic mass ratio of 2.115 ($m_{1}$ / $m_{2}$), smaller than the one determined by \citet{hillbarnes72} from photographic plates. They determined absolute parameters of the system using the photometric observations of \citet{criswal63}. They also classified the spectral type of the primary component (A4 V) from a combined spectrum of their data, which was consistent with the temperature estimates obatined from Str\"omgren photometry by \citet{hilditch75}. The configuration of the system as a semi-detached eclipsing binary has been widely accepted since this study and that of \citet{giuricin81}.\\ \citet{savedoff51} was the first to notice that the orbital period was variable. Although the secular change in the eclipse timing variations has been noted by recent studies \citep{margrave81,guduretal87,liakos12a, hanna12}, \citet{liakos12a} noticed also the discrepancy between the configuration and the direction of the mass transfer. They suggested either mass loss via stellar winds or system angular momentum loss via magnetic braking as possible explanations. Periodic variations in the O-C diagram have been also noted and attributed to unseen third bodies with parameters differing from one study to another \citep{borhege95,borhege96,liakos12a}, or to magnetic activity with single \citep{sarna97}, or to multiple cycles \citep{hanna12}. However there is no firm evidence for strong magnetic activity other than enhanced X-ray emission \citep{hanna12}.

In this study we have performed thorough analyses of the light curves and period variations of three close binary systems. Analysis of our light curves was based on combined photometric and spectroscopic data. The mass ratios of the studied systems were determined from our own spectroscopic observations. We made use of the Wilson-Devinney algorithm to derive the best fits to new, multicolor light curves, and based on obtained results we calculated the absolute parameters of components. The complementary information about the systems has been derived by analyzing their period behaviour. \subsection{\astrobj{XZ~Aql}} The light curve of {\astrobj{XZ~Aql} has equal levels of light maxima, and we were able to obtain its solution that required neither a spot nora third light. This is a high inclination (i=86 deg), semi-detached binary with the less massive secondary filling its Roche lobe, while the primary component is well inside its Roche lobe. The frequency analysis performed after having removed the binarity effects from its light curves hints that this A-type primary is a $\delta$ Sct type pulsator. Due to a significant temperature difference between components, the contribution of the secondary star to the total system light is small: only 2\% in the $B$ filter and about 13\% in the $I$ filter. \\ From the O-C analysis we found a parabolic change in the eclipse timings combined with a periodic variation that could be due to an unseen third body as previously suggested by \citet{soydugan06}. We made trial computations with a third light, however, its resulting intensity was negligible, reaching only 1\% in the $I$ filter. Assuming that this third body is in coplanar orbit with the binary and that it is an MS star, its luminosity contribution to the total light would only be 0.2\%. We found that the orbit of this third companion would be stable according to Harrington's criterion \citep{harrington77}. We have found a parabolic relation in the residuals that may be due to mass transfer from the less massive component to the more massive one. When the corresponding orbital period changes are removed, one might argue that there is a possibility of a second periodicity in the O-C points, but the number of data points is not sufficient to prove it. Alternatively, the periodic variation of the orbital period could also be explained with the quadrupole moment variation of the secondary.\\ Our results from the orbital period variation analysis also support the finding that the secondary component fills its Roche lobe and transfers mass to the more massive primary conservatively. The positions of its component on the HRD, MLD, and MRD based on the absolute parameters computed for each of the components as a result of the light curve modeling support this finding. Both the computed masses and the effective temperatures from our fit are qualitatively consistent with evolutionary tracks, which also point to a main-sequence primary half way between the ZAMS and the TAMS with $\sim$ 2.42 M$_{\odot}$, and an evolved red giant secondary, somewhat more massive ($\sim$ 0.65 M$_{\odot}$) than our computed value ($\sim$ 0.45 M$_{\odot}$). This discrepancy can be explained by the mass transfer from the secondary to the primary star, the time scale of which should be less than 10 Myr assuming a constant rate for the mass transfer. The fact that the main sequence primary has not gained all the mass the secondary transfers can only be explained by a non-conservative mass loss. \subsection{\astrobj{UX~Her}} Our light curve modelling for \astrobj{UX~Her} used the new value for the spectroscopic mass ratio resulting from our DAO data q$_{sp}$ = 0.184. It is significantly smaller than the previous determination from photometry alone as a result of the q-search by \citep{gojko06}. We obtained a semi-detached configuration for this system with the less massive component filling its Roche lobe while the primary being well within it. Therefore, mainly due to the lower q value, the derived secondary mass is also smaller than that previously found \citep{lazaro97,gojko06}, both considered this system as a detached one based on either photometrically or spectroscopically determined mass ratios from photographic plates. The evolutionary state of \astrobj{UX~Her} components deduced from their positions on the HRD, MLD and MRD diagrams point to the interpretation that the primary star is still on the main sequence, while the secondary is the more evolved star.\\ We did not find a significant quadratic component to the period variation, therefore we conclude there is no evidence of mass transfer. The cyclic period variations can more plausibly be explained by a third body of mass M$_{3}$ = 0.87 M$_{\odot}$, orbiting the center of mass on a significantly eccentric orbit (e = 0.41). \citet{tremko04} found a smaller mass, less than half of the value that we found (0.30 M$_{\odot}$). As a test, we also performed computations adding a third light as a free parameter. It turned out that contribution of this hypothetical tertiary to the total light is negligible: it reached about 1\% only in the $I$ filter so therefore we present the solution without l$_3$ as the final one. If we assume that this body has a coplanar orbit with that of the eclipsing binary and that it is an M5 star, its luminosity contribution will be less than 1\% (0.75\% according to our calculations), which is below our detection limits with the photometry. Previously \citet{selam07} also found cyclic changes in the orbital period in \astrobj{UX~Her}. They attributed the variation to magnetic activity, which would also explain the observed asymmetries in the light curve. However, the latest CCD data have a very low scatter around our best fit. More scatter would be expected in the case of magnetic activity because it would complicate the measurement of the eclipse times from asymmetric minimum profiles due to the surface spots therefore causing more errors. Further observations of the system covering a longer time base will be needed to make sure that the period and the amplitude of the variation change from one cycle to another, which would be the case in the presence of strong magnetic activity. Otherwise, the variation could be argued to cause from the gravitational pull of an unseen third body. In additon, the temperature of the primary is too high (8770 K) to expect cool surface spots. Though it would be somewhat reasonable to locate spots on the cooler secondary, our non-spotted solution satisfactorily fits the observations. Therefore, we give here only the third body solution because we do not have further evidence (e.g. variations in magnetic activity indicators) supporting the magnetic activity argument. Moreover, the dynamical stability test according to Harrington's criterion \citep{harrington77} points to a stable orbit for the suggested third body, thus supporting the second hypothesis. \subsection{\astrobj{AT~Peg}} Light curve modelling for \astrobj{AT~Peg} proceeded without difficulty for both radial velocity data sets (our own and that of \citet{maxtedetal94}). A convergence was quickly found for a model that required neither a third light nor a spot. Theoretical light curves fit observed ones very well with just very small discrepancies visible around the primary minimum but only in $B$ and $I$ filters. We arrived at the semi-detached solution with the less massive secondary filling the Roche lobe and the primary well within it. Absolute parameters and their errors differ only within a few percent between the two models obtained by using two different radial velocity data sets. \\ The O-C analysis shows a parabolic trend. A periodic relation for the residuals can also be asserted. The quadratic term cannot be interpreted as being due to conservative mass transfer between components as it would require the more massive star to be the mass loser, in contradiction with the results from light curve modeling. This discrepancy can be explained by non-conservative mass loss from the system or losing of angular momentum via magnetic breaking. Such systems with orbital period decrease may be the progenitors of contact binaries \citep{bradstreet94,qian2000}, because the fill-out factor of the primary may increase while the orbital period decreases, eventually causing the primary to fill its Roche lobe. We do not have evidence for strong magnetic activity such as light curve modulations and asymmetries in the light curves due to stellar spots. The spectral window of our spectral observations is also limited to 5000-5250 \AA ~region, where we have not observed a signature of activity or stellar wind in the resolution our spectra were taken. The spectral type of the primary (A4) would not support this argument either, assuming it is a normal star without any kind of pecularity. The positions of the components relative to the evolutionary tracks on the HRD indicate that the primary might have lost some mass. The evolutionary model value with no mass transfer is $\sim$ 1.90 M$_{\odot}$ for the primary component, while the computed mass from the light curve analysis is 2.22 M$_{\odot}$ and the mass of the secondary is consistent in both models ($\sim$ 1.08 M$_{\odot}$).\\ The residuals from the quadratic fit may be alleged to follow a periodic behavior, which can be attributed to an unseen body (M$_{3,min}$ = 0.68 M$_{\odot}$) gravitionally bound to the system, orbiting its center of mass once in every $\sim$33 years on a a significantly eccentric (e = 0.53) orbit. The relatively low mass of the tertiary can explain its negligible contribution to the total light. Additional bodies have been plausibly argued to extract momentum from the binary, causing the orbit to shrink and hence the orbital period to decrease \citep{yangwei09}. This could be at least a part of the reason that causes a period decrease although the direction of the mass transfer is towards the more massive primary.

1607.06930

1607

1607.03038_arXiv.txt

94\,Ceti is a triple star system with a circumprimary gas giant planet and far-infrared excess. Such excesses around main sequence stars are likely due to debris discs, and are considered as signposts of planetary systems and, therefore, provide important insights into the configuration and evolution of the planetary system. Consequently, in order to learn more about the 94\,Ceti system, we aim to precisely model the dust emission to fit its observed SED and to simulate its orbital dynamics. We interpret our APEX bolometric observations and complement them with archived \textit{Spitzer} and \textit{Herschel} bolometric data to explore the stellar excess and to map out background sources in the fields. Dynamical simulations and 3D radiative transfer calculations were used to constrain the debris disc configurations and model the dust emission. The best fit dust disc model for 94\,Ceti implies a circumbinary disc around the secondary pair, limited by dynamics to radii smaller than 40\,AU and with a grain size power-law distribution of $\sim a^{-3.5}$. This model exhibits a dust-to-star luminosity ratio of $4.6\pm 0.4\,\times 10^{-6}$. The system is dynamically stable and N-body symplectic simulations results are consistent with semi-analytical equations that describe orbits in binary systems. In the observations we also find tentative evidence of a circumtertiary ring that could be edge-on.

Approximately half of the stars belong to binary or multiple systems \citep{duquennoy1991,lada2006,eggleton2008,raghavan2010,duchene2013}. Stellar companions may have a large impact on planetary formation processes. In particular, they are found to truncate circumstellar discs and put a limit on the extent of material available for planetary formation \citep{artymowicz1994,jangcondell2008,andrews2010,jangcondell2015}. They are also expected to stir and increase the eccentricities and relative velocities of planetesimals and planetary embryos that might have formed around the primary \citep{quintana2007,thebault2009,rafikov2015}. Consequently, multiple star systems generate significant collisional activity which may shorten the lifetime of a disc and make it difficult to form planets. Observations tend to support this, since it was found that protoplanetary discs are half as common in young (a few Myr) binary stars as compared to single stars \citep{cieza2009,kraus2012}. In addition, these discs are much less massive compared to those found around single stars \citep[by a factor $\sim 5-25$, ][]{harris2012} and were found to be more short-lived as well \citep{duchene2010}. However, debris discs have been found around old binary stars, first with \textit{Spitzer} \citep{trilling2007} and then with \textit{Herschel} \citep{rodriguez2015}. They were detected through their host star infrared excess emission, signature of the presence of micron-sized dust grains. These grains are expected to acquire unstable orbits due to both gravitational and non-gravitational perturbations. The latter scenario includes radiation pressure and Poynting-Robertson drag that effectively clear dust clouds of small and large grains ($<1\,$\um\ and between 1\,\um\ to 1\,mm, respectively). As such, dust in circumstellar discs must be continuously replenished from collisional processes in rings of parent bodies (planetesimals) akin to that of the asteroid and Edgeworth-Kuiper belts in our Solar System \citep{artymowicz1997,krivov2008,wyatt2008,moromartin2013}. These underlying reservoirs of large bodies show that the presence of a stellar companion does not necessarily hinder the formation of the building blocks necessary to form planets. The presence of at least one stellar companion does not necessarily inhibit planetary formation either, since numerous planets were found in binary or multiple systems. However, the majority of these planets were found in wide binary systems \citep{mugrauer2005,daemgen2009,muterspaugh2010,bergfors2013}. In the cases where the companion was not expected to significantly affect planetary formation processes, it is not surprising that these systems were found to be hosts to planets \citep{eggenberger2007,desidera2007,duchene2010} and that the properties of circumstellar protoplanetary and debris discs in binaries are found to be similar to those hosted by single stars for separations larger than several tens of AU \citep{kraus2012,harris2012,rodriguez2015}. According to the Open Exoplanet Catalogue\footnote{\url{http://www.openexoplanetcatalogue.com/}} (during of the summer of 2016), out of the 130 multiple star systems with planets, only less than one fifth (23) are triple star systems. This is expected considering that triple star systems represent less than one fifth of the multiple star systems \citep{raghavan2010}. In most cases of triple star systems, the planet is orbiting the primary star with a secondary stellar binary surrounding the primary on a wider orbit. In addition, investigating associated debris discs and tracing the interactions between discs, planets and/or stellar companions can considerably aid our understanding of the dynamical history of a planetary system. This kind of study was carried out for the Solar System with the Nice model of \citet{gomes2005}, but also for the systems of e.g. $\beta\,$Pictoris \citep[see e.g. ][]{dent2014,nesvold2015,millarblanchaer2015,apai2015}, Fomalhaut \citep[see e.g.][]{beust2014,faramaz2015,cataldi2015,lawler2015}, HR\,8799 \citep[see e.g. ][]{moore2013,contro2015,booth2016}, $\tau\,$Ceti \citep{lawler2014}, and HD\,69830 \citep{payne2009}. Therefore, studying in detail close binary or hierarchical systems, where both planets and circumstellar emission have been found, is crucial. Five systems, out of the 23 found in triple star systems mentioned above, are associated with circumstellar dust emission. These are HD\,178911 \citep{saffe2004,kospal2009}, Fomalhaut \citep{aumann1985,kalas2008,mamajek2013,su2016}, HD\,40979 (e.g. \citealt{kospal2009,dodsonrobinson2011}), 51\,Eridani (e.g. \citealt{rivieremarichalar2014}), and 94\,Ceti (HIP\,14954, HD\,19994). Another possible candidate is L1551 IRS 5, which is a young embedded binary system and potentially a triple star system \citep{lim2006}. Due to the low number of known triple star systems with both planets and dust emission, it is important to study each of these in great detail. We will focus here on the 94\,Cet system. The third star of 94\,Cet was recently discovered by \citet{roll2011a,roll2012} with astrometry, speckle interferometry, and radial velocity measurements. The system is associated with at least one circumprimary planet \citep{eggenberger2003}, and also shows a far-infrared excess emission \citep{trilling2008,eiroa2013} that possibly originates from circumstellar dust. This is a hierarchical triple star system, at a distance of $22.6 \pm 0.1$\,pc from the Sun, where 94\,Cet\,A is an F8\,V star, and 94\,Cet\,B and C are both M dwarfs that form a binary pair that together orbits the primary on a 2029 year long orbit. The infrared excess of 94\,Cet was detected by \citet{trilling2008} with \textit{Spitzer}/MIPS at 70\,\um\ and later confirmed, at 100 and 160\,\um\, by \citet{eiroa2013} as part of the key project DUNES (DUst around NEarby Stars) of \textit{Herschel} \citep{pilbratt2010}. Infrared excesses at such long wavelengths originate from the thermal emission of dust particles surrounding the star. To improve the coverage of the spectral energy distribution (SED) we observed this system with APEX (the Atacama Pathfinder EXperiment) at 870\,\um . \citet{eiroa2013} found that the inferred dust temperature corresponds to a black-body radial distance from the primary star that corresponds to a dynamical unstable region due to the companion stars. Another hint that this emission may not be associated with 94\,Cet\,A is that there is a marginally significant offset between the expected and observed position of the source, if it is co-moving with 94\,Cet\,A. It may thus be a background source or be associated with the other stellar members of this system. The structure of this paper is as follows; we present the stellar and system properties in Section\,2, and we summarize the observations, data reduction, and the observational results in Section\,3. In Sections\,4 and 5, we discuss the nature of the excess, the extended emission, and background sources. Assuming that the excesses originate from disc(s), we apply both dynamical simulations and radiative transfer simulations and present these results in Section\,6, and summarize our conclusions in Section\,7.

Multi-component systems pose many difficult questions concerning planet formation processes. 94\,Cet provides one of few unique opportunities of a case study of a triple stellar system \citep{roll2011a,roll2012} with at least one planet \citep{eggenberger2003} and FIR excess \citep{trilling2008,eiroa2013}. For these reasons and a marginally significant offset in the PACS data we aimed to model the stellar excess and simulate the orbital dynamics of the system. \begin{itemize} \item[--] The central source fits well with that of a circumbinary disc around the companion pair (inside the dynamically stable radii) at $30.3 \pm 7.4$\,K, with a dust grain size distribution $q$ between 3 and 3.5, and fractional luminosity $f_{\rm d} = 4.6\pm 0.4\,\times 10^{-6}$. The disc extends from 3\,AU to 40\,AU from the companion stars' barycentre with a surface density distribution that was $\propto r^{2}$. Assuming $q=3.5$ and the grain size range of 8\,\um\ -- 1\,mm this corresponds to $6.0\pm 0.5\,\times 10^{-2}\,M_{\rm Moon}$. \item[--] The eastern and western extensions have the flux densities $8.7 \pm 1.3$\,mJy at 100\,\um\ and $16.1 \pm 4.1$\,mJy at 160\,\um , which corresponds well with the possibility of a circumtertiary ring of dust ($q=3.5$) with the fractional luminosity $f_{\rm d} \approx 1.4\pm 0.3\,\times 10^{-6}$ and temperature of $< 30$\,K. However the uncertainty is significant and the temperature is unconstrained. The ring was assumed to extend from an inner edge at 600 -- 650\,AU to $\sim 750$\,AU, and the initial surface density distribution of $\propto r^{-1}$. \item[--] The system and disc configurations are stable after 20 Myr. The particle discs were simulated with the symplectic integrator HJS \citep{beust2003} using $10^4$ particles for each disc. \end{itemize} Our models provide evidence for the possibility of both circumsecondary and circumteriary dust. It is possible that the lack of circumprimary (hotter) dust emission is due to additonal planets emptying the circumprimary neighbourhood and our evidence for the circumtertiary ring is only tentative. As such it would be useful to further constrain the nature of this system through additional observations. Even without \textit{Herschel} there are a few possibilities available. ALMA, for example, could probably not observe any circumtertiary dust, but may be able to confirm our findings on the circumsecondary dust, and the LMT (Large Millimeter Telescope) in Mexico can observe 94\,Cet during autumn and could reach a $3\,\sigma$ of 0.3\,mJy in 36\,hours, possibly enough to reach the dust emission.

1607.03038

1607

1607.04568_arXiv.txt

{We investigate an infrared-excess source called G2 or Dusty S-cluster Object (DSO), which moves on a highly eccentric orbit around the Galaxy's central black hole, Sgr~A*. We use, for the first time, near-infrared polarimetric imaging data to determine the nature and properties of the DSO and obtain an improved $K_\mathrm{s}$-band identification of this source in median polarimetry images of different observing years. The source started to deviate from the stellar confusion in 2008, and it does not show any flux density variability over the years we analyzed it. We measured the polarization degree and angle of the DSO between 2008 and 2012 and conclude, based on the significance analysis on polarization parameters, that it is an intrinsically polarized source ($>20\%$) with a varying polarization angle as it approaches the position of Sgr~A* . The DSO shows a near-infrared excess of $K_{\rm{s}}-L' > 3$ that remains compact close to the pericenter of its orbit. Its observed parameters and the significant polarization obtained in this work show that the DSO might be a dust-enshrouded young star, forming a bow~shock as it approaches the super massive black hole. The significantly high measured polarization degree indicates that it has a non-spherical geometry, and it can be modeled as a combination of a bow~shock with a bipolar wind of the star. We used a 3D radiative transfer model that can reproduce the observed properties of the source such as the total flux density and the polarization degree. We obtain that the change of the polarization angle can be due to an intrinsic change in the source structure. Accretion disk precession of the young star in the gravitational field of the black hole can lead to the change of the bipolar outflow and therefore the polarization angle variation. It might also be the result of the source interaction with the ambient medium. }

\label{section:Introduction} Since 2012, the focus of Galactic center (GC) observations has been set on investigating an infrared (IR) excess source detected by \cite{Gillessen2012} as a fast-moving object approaching the position of the central supermassive black hole (SMBH) of the Milky Way, Sagittarius A* (Sgr~A*). It has been interpreted as a combination of dust and core-less gas cloud called G2 \citep{Gillessen2012, Gillessen2013a, Pfuhl2015} and also DSO, standing for Dusty S-cluster Object \citep{Eckart2013}. This source moves on a highly eccentric orbit and passed its closest approach to the SMBH in May 2014 \citep{meyer2014,Valencia2015}. Given the short distance of its periapse, it has been suspected that it might produce extraordinary accretion events on to the galaxy's central black hole \citep[e.g.][]{Shcherbakov2014, Abarca2014, Scoville2013, Sadowski2013}. If the DSO is a pure gas cloud of a few Earth masses \citep{Gillessen2012}, it might have formed in the stellar cluster, possibly within the disk of young stars at a distance of few arcseconds from the GC. After forming there, it might have moved on its current remarkably eccentric orbit by gravitational interaction with massive stars \citep{Murray-Clay2012, Scoville2013}. This scenario must have happened recently (1990-2000), therefore it should have been observed during the total time of its existence. As a consequence, the pure gas scenario seems unlikely, and several authors have proposed scenarios suggesting the presence of a central star for this source \citep[e.g.,][]{Murray-Clay2012, Eckart2013, Scoville2013, Ballone2013, Phifer2013, Zajacek2014, Witzel2014, Valencia2015}. \cite{Scoville2013} proposed that the DSO is a T~Tauri star that was formed in the young stellar ring and then inserted into its current orbit. They suggested that a very dense bow~shock is produced for the T~Tauri star wind and modeled it numerically. \cite{Valencia2015} also discussed that the bright observed Br$\gamma$ emission of the DSO with a large line width might be the result of infalling material shaping a disk around the central star, which may be a T~Tauri star with an age of $\sim10^5$ yr. Considering the stellar nature, DSO would not be disrupted when reaching its closest point to the SMBH and did not need a recent formation. Using hydrodynamical simulations, \cite{Jalali2014} have shown that young stars could form very close to SMBHs within small molecular clumps on eccentric orbits around the black hole. They showed that for such orbital configurations, the gravitational potential of the SMBH and orbital (geometrical) compression increase the density of cold gas clumps to reach the threshold values suitable for star formation (see also \cite{mapelli2016} for a recent review). The $L$-band observations of DSO/G2 close to the peribothron support the idea of the compactness of this source, which means that the source cannot be a pure gas cloud \citep{Ghez2014, Witzel2014}. \cite{Eckart2013} revealed the first $K_\mathrm{s}$-band identification of \textcolor[rgb]{1,0.501961,0}{\textcolor[rgb]{1,0.501961,0}{\textcolor[rgb]{0,0,0}{the} \textcolor[rgb]{0,0,0}{DSO}} }with a magnitude of $\sim$18.9 from the ESO Very Large Telescope (VLT) continuum imaging data. Using the spectral decomposition of this source, they obtained an upper limit of $\sim 30~L_{\odot}$ for its luminosity. The $H$, $K_\mathrm{s}$, and $L$-band continuum measurements can be matched either by an unusually warm dust component at a temperature of 550-650 K or by a stellar source enclosed in the dust at a temperature of $\sim$ 450 K \citep[see Fig.~15 in][]{Eckart2013}. The $H-K_\mathrm{s} > 2.3$ color limit supports the scenario that the DSO is a dust-embedded star and not a core-less cloud of gas and dust \citep{Eckart2013}. The mass of this object is higher than what was assumed for a pure gas source, but lower than the typical mass of S-cluster stars ($20 M_{\odot}$). \cite{Valencia2015} reported NIR observations of the DSO during its approach to the SMBH at the GC, which were carried out with SINFONI at the VLT from February to September 2014. They detected spatially compact Br$\gamma$ line emission from the DSO before and after its peribothron passage and also a Br$\gamma$ line width increase, which may indicate that the DSO is a young accreting star with a dust envelope. The observational data were used to obtain the orbital parameters of this object. Comparable to the previous estimates \citep[e.g.,][]{Meyer2014a}, \cite{Valencia2015} obtained a peribothron distance of about $163\pm16$ AU with a half-axis length of about 33 mpc and an ellipticity $e=0.976$. When the DSO reaches the peribothron, tidal stretching and disruption of the envelope lead to a velocity dispersion enhancement in the accretion flow toward the central star \citep{Eckart2013, Zajacek2014}. Based on a study by \cite{Witzel2014}, the $L$-band emission of the DSO compared to the Br$\gamma$ emission measured by \cite{Pfuhl2015} is more compact. This shows that the $L$-band emission originates from an optically thick dust envelope around a central star, while the Br$\gamma$ emission is coming from the hot gas that is externally heated by ionized photons of the stars close to the DSO. However, \cite{Valencia2015} did not find evidence for significantly extended and tidally stretched Br$\gamma$ emission. The extra emission of Br$\gamma$ close to the DSO position is not connected to the DSO and is most likely emitted from other sources in the field (Peissker et al., in prep.). The DSO is not the only infrared excess source in the S-cluster of the GC, and there are more dusty sources in this region \citep{Eckart2013,Meyer2014a}. These sources might be dust-enshrouded pre-main-sequence stars that form a bow shock ahead of their path when they move through the medium with a supersonic speed. Other candidates have also been observed in the radio continuum observations of the GC \citep{yusef2015}. \\ Imaging polarimetry is a powerful technique for studying dusty environments such as core-less dusty objects and/or circumstellar dusty regions. The analysis of polarization allows us to quantitatively evaluate the object geometries and the dust properties. Intrinsic polarization can be generated only if the system is not symmetric. The asymmetry can occur when the radiation field of the star is not isotropic as a result of a geometric distortion, for instance, when the star develops a bow~shock ahead of its path, or when its photosphere surface brightness is not uniform, or in other words, is influenced by bright spots. Therefore, supplementary to considering the continuum and line emissions from the DSO, studying the light polarization can be very helpful in determining the nature and properties of this source. If the DSO is a bow~shock, the polarization is determined by the bow~shock morphology. Subsequently, the E-vectors are predicted to be perpendicular to the direction of motion if the medium is homogeneous. If the dust shell surrounding the DSO is a disk, then the resulting polarization depends on the disk inclination. In this paper we analyze the NIR polarimetric imaging data taken with NACO at the ESO VLT using its Wollaston prism to study the polarization properties of the DSO. In Sect. \ref{section:Observations} we begin with the details about the observations, then describe the data reduction and determine the position of the DSO in the images. In Sect. \ref{section:results} we present the results of the applied flux density calibration method: light curves, polarimetry measurements, and their statistical analysis. We present the performed radiative transfer model to describe the DSO polarization in Sect. \ref{dso-model}. We discuss the implications of our results in Sect. \ref{section:discussion} and finally summarize the main results of the NIR polarimetry of the DSO in Sect. \ref{section:summary}. \begin{figure*}[]% \begin{center} \subfigure{% \includegraphics[width=0.45\textwidth]{2008_referee_new.eps} } \subfigure{% \includegraphics[width=0.45\textwidth]{2009_referee_new.eps} }\\ \subfigure{% \includegraphics[width=0.45\textwidth]{2011_referee_new.eps} } \subfigure{% \includegraphics[width=0.45\textwidth]{2012_referee_new.eps} }\\ \end{center} \caption[$K_\mathrm{s}$-band deconvolved median images of the central arcsecond of the GC in polarimetry mode]{Final $K_\mathrm{s}$-band deconvolved median images of the central arcsecond at the GC in polarimetry mode ($90^{\circ}$ polarization channel) in 2008, 2009, 2011, and 2012 from top left to bottom right. The position of the DSO is shown by a circle on its orbit.} The asterisk indicates the position of Sgr~A*. In all the images north is up and east is to the left. \label{fig:allepochs-dso} \end{figure*}

\label{section:summary} We have analyzed the NIR polarized observations of a Dusty S-cluster Object (DSO/G2) on an eccentric orbit around Sgr A*. $K_\mathrm{s}$-band polarization data were available for 2008, 2009, 2011, and 2012. In all these years we clearly detected $K_\mathrm{s}$-band continuum emission in the different channels of the Wollaston prism. The data cover the polarization information of the DSO before its pericenter passage (May 2014). The source does not show significant variability in the overall $K_\mathrm{s}$-band flux density during the observed years, and its polarization degree is mostly above 20\%, which is higher than the foreground polarization measured on the surrounding stars. It appears that the polarization degree is approximately constant and the polarization angle varies as it approaches the position of Sgr A*. Based on our significance analysis, the polarization measurements of the DSO are significant in 2008, 2009, and 2012 and can be interpreted as source-intrinsic properties.\\ Since the total polarization degree is noticeably high, higher than $20\%$ for all epochs, the DSO structure is expected to deviate from the spherical symmetry. Moreover, the analysis of \citet{Valencia2015} showed that the source remains compact, meaning that it is not effected by tidal forces close to the pericenter of its orbit. \citet{Gillessen2012} and \citet{Eckart2013} discussed a NIR excess of $K_{\rm{s}}-L' > 3$, which implies the presence of a dense gaseous-dusty envelope. All of these basic observed parameters, the NIR excess, the compactness, and the significant polarization, may be reconciled within the model of a dust-enshrouded young star, to be precise, a pre-main-sequence star of class 1 \citep{Zajacek2014,2015wds..conf...27Z}, that forms a dense bow-shock layer by its supersonic motion upon approaching the supermassive black hole \citep{Zajacek2016}. The obtained polarization properties of the DSO in this work can be caused by the non-spherical geometry of a bow~shock and a bipolar wind of the star. We used the 3D radiative transfer model implementing the code Hyperion \citep{2011A&A...536A..79R} to compare the observed measurements to the model of the young stellar object forming a bow~shock. We conclude that the varying polarization angle is related to the intrinsic change of the circumstellar configuration. The change in bipolar outflow orientation may be due to the accretion disk wobbling or precession in the gravitational field of the SMBH. It can be also produced by external interaction of the DSO with the accretion flow. Although our model is simple, it can reproduce many observed properties of the DSO obtained in this work, such as the total flux density and the polarization degree. A more detailed analysis of the model will be provided in Zaja\v{c}ek et al., in prep.\\ \cite{shahzamanian2015b} showed that the Sgr~A* system exhibits a stable geometry and accretion process that is consistent with the preferred jet or wind directions. The close fly-by of the DSO, or similar dusty sources (Peissker et al., in prep.), might have an effect on the stable accretion flow onto Sgr~A* that depends on the nature of these objects. However, after the pericenter passage of the DSO, the object remained compact and its orbit Keplerian \citep{Witzel2014,Valencia2015}. Consequently, it did not lose a noticeable amount of energy and angular momentum during its closest approach to Sgr~A*, and as a result experienced weak interactions with the central black hole \citep{park2015}. However, based on hydrodynamical simulations, it may take several years for their interaction and to see a change in the activity of Sgr~A* \citep{burkert2012, schartmann2012}, either as an increase in the accretion flow rate or in the appearance of jets \citep{yuan2014}. Therefore, polarization and variability measurements of Sgr~A* are needed to be continued as they are the ideal tool to probe any change in the apparently stable system as a function of the DSO fly-by. Moreover, future polarized observations of the DSO, that is, after the pericenter passage, in the NIR can help us to better constrain the source polarization and structure.\\

1607.04568

1607

1607.03986_arXiv.txt

Proton capture cross sections in the energy range of astrophysical interest for mass region 40-54 have been calculated in the Hauser-Feshbach formalism with reaction code TALYS1.6. The density dependent M3Y effective nucleon-nucleon interaction folded with target radial matter densities from relativistic mean field approach is used to obtain the semimicroscopic optical potential. A definite normalization of potential well depths has been used over the entire mass region. The $(p,\gamma)$ rates of some reactions, important in the astrophysical scenario, are calculated using the potential in the relevant mass region.

During stellar burning, reactions involving different seed nuclei assume importance at different temperature and burning zones. The seeds in the concerned mass range A=40-54 are mainly produced during hydrostatic carbon burning and explosive oxygen burning. These then take part in synthesizing more massive elements via various reactions occurring in later phases of evolution~\cite{seed_reaction1,seed_reaction2}. The principle energy generating processes in stars i.e the pp-cycle, CNO cycle, HCNO cycle, {\em rp} process etc are reactions which require continuous addition of protons against the Coulomb barrier. Certain astrophysical sites, such as X-ray bursters, involve high flux of protons at temperature $\sim$ several GK and matter density $\sim$ $10^{6}$ g/cc that initiate a rapid proton capture process which ultimately results in a thermal burst of very short duration, peaked in X ray regime. Knowledge of the cross sections, and equivalently, the rates of the reactions occurring in these sites are required to study complete nucleosynthesis via a network calculation. However, it is difficult to measure all the essential rates in terrestrial laboratories due to unavailability and instability of required targets. Reaction rates calculated in a theoretical approach may provide necessary information to this context after proper validation of theory with the experimental data. For abundance calculation in explosive astrophysical environment, the concerned network must have to take various quantities like temperature, pressure, proton mass fraction along with forward and reverse reaction rates into account. So, we need to take care about the proper tuning of the interaction potential. Many works have been devoted so far to study the theoretical capture cross sections by constructing different nucleon-nucleus potentials. Rauscher {\em et. al.}~\cite {raus1,raus2} have calculated reaction rates in a global approach and suggested that statistical model calculation can be made improved using locally tuned nuclear properties like optical potential. Reaction rates from phenomenological approach i.e, with phenomenological global optical potentials or even with semimicroscopic optical potential with phenomenological densities give rise to uncertainty and the reaction rates have to be varied by large factors to study their effects. Prediction of rates, as has been done in Ref.~\cite{schatz1}, gave rise to uncertainties away from the stability valley. Phenomenological global optical potential should not be expected to provide an adequate description of nucleon-nucleus interaction as differences in nuclear structure among adjacent nuclei do not allow simple and smooth Z- and A- dependence of the Woods-Saxon parameters. Microscopic optical potentials obtained by folding with appropriate microscopic densities are expected to be more accurate and do not require frequent variation of the reaction rates. On the other hand, calculation of reaction rates in a microscopic or semimicroscopic with microscopic density prescription approach is far more free from these uncertainties. Recently, the semimicroscopic optical potential obtained in a folding model prescription has proved to be highly successful in explaining various nuclear phenomena. For example, Bauge {\em et al.}~\cite{bauge1,bauge2} have constructed a lane-consistent semimicroscopic optical potential to study elastic scattering and differential and total cross sections for nuclei over a broad mass range. The theoretical cross section calculation requires a complete knowledge of various ingredients such as transmission coefficients, i.e., transition probabilities (averaged over resonances of the compound nucleus formed upon radiative proton captures) between various states which in turn depend on the level schemes, lifetimes of the states, level densities, $\gamma$ ray strength functions, nuclear masses, giant dipole resonance parameters, etc. Application of statistical model requires sufficiently high nuclear level density in the compound nuclear state. The theoretical reactions are generally derived in Hauser-Feshabach (HF) formalism which assumes presence of sufficiently large number of resonances at relevant energies. The cross sections are in general sum of contribution from different reaction mechanisms depending on projectile energies. At higher energies, presence of many close and overlapping resonances allow one to calculate an average cross sections using HF approach. Sometimes there are interference effects between single resonance and direct capture. The nonresonant reaction cross sections are mainly determined by the direct capture transitions to the ground states and low excited states. In light nuclei and low energy regime, level densities are generally low, especially for targets near closed shells with widely spaced nuclear levels and close to drip lines with low particle separation energies and Q-values. Hence, application of statistical model to these nuclei at astrophysical temperatures is somewhat problematic and requires careful study. Optical potential is a very important ingredient in HF statistical model calculation. Here, we have constructed an optical potential by folding the density dependent M3Y (DDM3Y) interaction with the densities of the target. The DDM3Y interaction has proved to be successful in explaining various nuclear properties. For example, folded DDM3Y nucleon-nucleus interaction potential has been successfully used to study the incompressibility of infinite nuclear matter \cite{dnbasu1} and radioactivity lifetimes of spherical proton rich nuclei \cite{dnbasu2}. The paper is organized as follows. In the next section we have presented the framework of our calculation. In section III we discuss the results and in the last section we summarize.

To summarize, we calculated the cross sections for $(p,\gamma)$ reactions in the mass range 40-54 in the relevant Gamow energy window appropriate for low energy astrophysical environment using the well known reaction code TALYS1.6. Charge radii and binding energy values of various stable nuclei calculated using RMFT have been compared in this concerned mass region A = 40-54 with the available experimental data. The DDM3Y NN interaction is folded with target nuclear densities calculated from RMFT to construct the optical potential which is needed for HF statistical model calculation and after proper normalization it has been used to verify the theory with the observed experimental data. The $(p,\gamma)$ reaction rates are calculated and plotted along with NON-SMOKER reaction rates. The main feature of our work is to place all the nuclei considered in this mass region A = 40-54 at the same footing and to use same methodology for all of them to avoid systematic error.

1607.03986

1607

1607.02664_arXiv.txt

Due to their long mean free paths, X-rays are expected to have global impacts on the properties of the intergalactic medium (IGM) by their large scale heating and ionizing processes. At high redshifts, X-rays from Population (Pop) III binaries might have important effects on cosmic reionization and the Ly$\alpha$ forest. As a continuation of our previous work on Pop III binary X-rays \citep{Xu14}, we use the Pop III distribution and evolution from the \textit{Renaissance Simulations}, a suite of self-consistent cosmological radiation hydrodynamics simulations of the formation of the first galaxies, to calculate the X-ray luminosity density and background over the redshift range $20 \geq z \geq 7.6$. As we find that Pop III star formation continues at a low, nearly constant rate to the end of reionization, X-rays are being continuously produced at significant rates compared to other possible X-ray sources, such as AGNs and normal X-ray binaries during the same period of time. We estimate that Pop III binaries produce approximately 6 eV of energy in the X-rays per hydrogen atom. We calculate the X-ray background for different monochromatic photon energies. KeV X-rays redshift and accumulate to produce a strong X-ray background spectrum extending to roughly 500 eV. The X-ray background is strong enough to heat the IGM to $\sim 1000$ K and to ionize a few percent of the neutral hydrogen. These effects are important for an understanding of the neutral hydrogen hyperfine transition 21-cm line signatures, the Ly$\alpha$ forest, and the Thomson optical depth to the CMB.

\label{sec:introduction} One of the most important problems in astrophysics and cosmology in both theory and observation is to obtain the thermal and ionization history of the intergalactic medium (IGM) prior, during, and after the epoch of reionization. It is critical to understand cosmological observations, such as high redshift 21-cm neutral hydrogen line signatures, the Thomson scattering optical depth of the CMB, and properties of the Ly$\alpha$ forest. The appearance of the first luminous objects---likely the first stars---marks the end of the cosmic dark ages after recombination, and the beginning of the last cosmic phase transition of reionization, a prolonged process that ionizes and heats the IGM. The first generation of stars (Population III stars) are believed to have formed from metal-free gas in small dark matter halos and having a large (tens to a few hundreds M$_{\odot}$) characteristic mass \citep{Abel02, Bromm02, OShea07a, Turk09, Greif12_P3Cluster, Hirano15, Hosokawa16}. Pop III stars have short lifetimes \citep{Schaerer02} and may explode as supernovae (SNe), enriching their surrounding IGM. Once the gas metallicity passes some critical threshold, $\sim$ 10$^{-6}$ Z$_{\odot}$ if dust cooling is efficient \citep{Omukai05,Schneider06, Clark08} or $\sim$ 10$^{-3.5}$ Z$_{\odot}$ otherwise \citep{Bromm01, Smith09}, the gas can cool efficiently and lower its Jeans mass and form (Pop II) stars in clusters at much higher rates. This makes the Pop III period of a given galaxy short. Consequently it is believed that Pop III stars are not a major contributor to the UV photons responsible for reionization \citep[e.g.][]{Fan06}. (However see \citet{Wise08_Reion} and \citet{Ahn12} for their contribution to the {\em start} of reionization). X-ray radiation with its much longer mean free path than UV radiation has been considered a good candidate for the pre-ionization and pre-heating the IGM much earlier than the end of the reionization at $z \sim 6$ \citep[e.g.,][]{Oh01, Venkatesan01, Ricotti04, Ricotti05}, and for smooth heating of the IGM; e.g. \citep{Haiman11}. Its impact on the high redshift 21cm signal has been explored by many authors \citep{Pritchard07,Mirabel11,PritchardLoeb12,Visbal12,Pacucci14,Fialkov14a,Mirocha16}, and observational constraints on the cosmic X-ray background have been investigated by \citet{Fialkov16}. Recent cosmological simulations \citep{Turk09, Stacy10, Susa14} have found that metal-free star-forming clouds might fragment to form binaries in a non-negligible fraction of Pop III star forming events, making Pop III stars and remnants possible strong X-ray sources at high redshifts. Pop III in the approximate mass range $40-140 \Ms$ and $>260 \Ms$ may directly collapse to form black holes \citep[BHs;][]{Heger03}. Any strong accretion onto these massive Pop III seed BHs would lead to X-ray radiation at high redshifts \citep{Kuhlen05, Alvarez09, Tanaka12, Visbal12}. X-rays from Pop III binaries have been suggested to produce a pre-heated IGM \citep[e.g.,][] {Mirabel11, Haiman11, Visbal12, Mesinger13, Fialkov14b, Madau16}, and to partially ionize the IGM in large volumes \citep[e.g.,][]{Ostriker96, Pritchard07}. \citet{Xu14} showed that Pop III binaries might be the dominant X-ray sources at high redshifts and discussed how the IGM might be heated locally by hundred eV photons and globally by the X-ray background from keV photons. Homogeneously heating of the IGM on tens to hundreds Mpc scales by an X-ray background may be important to understand the linewidths of Ly$\alpha$ forest spectra \citep{Tytler09} and high-z 21-cm neutral hydrogen line signatures \citep{Ahn14}. Recently, \citet{Xu16_Late3} showed using numerical simulations that Pop III stars can form continuously at significant rates from redshifts higher than 20 down to the end of the reionization. This suggests that X-rays from Pop III binaries might continue to be important X-ray sources till the end of the reionization and the X-ray background may heat the IGM for a long enough period to significantly heat the IGM globally. In this {\em Letter}, we estimate the X-ray luminosities from Pop III binaries using the results of the \textit{Renaissance Simulations} \citep{Xu13, OShea15, Xu16_fesc} and calculate the X-ray background using simple models. This paper extends to lower redshifts the results first presented in \citet{Xu14} concerning the X-ray background itself; we defer to a future paper a detailed calculation of the heating and ionization of the IGM and its 21cm signature. We first describe the simulations and X-ray source and background models in Section 2. In Section 3, we present our results for the X-ray luminosity density, and the background intensity and spectra. We discuss the implications of this strong X-ray background from Pop III stars in Section 4.

\label{sec:conclusions} In this {\em Letter}, we present estimates of the X-ray emission and background from Pop III binaries derived from the Pop III star formation histories in \textit{Renaissance Simulations} beginning at $z \sim 25$ down to $z=7.6$ near the end of reionization. The key results are summarized as follows: \begin{enumerate} \item A significant amount of X-rays may be produced by Pop III binaries down to the end of reionization, as a direct result of Pop III stars forming continuously with steady rates since they begin forming at redshifts of $20-30$. \item The X-ray background builds up gradually once the first stars begin forming. The X-ray background intensity depends on the assumed source photon energies. Only keV photons contribute to the X-ray background, and become a broad band spectrum in the background. \item X-rays from Pop III binaries might continuously dominate X-rays from normal stellar X-ray binaries and AGN until the end of reionization, although this conclusion depends on the assumed fraction of Pop III stars that become X-ray binaries. \end{enumerate} The results of this {\em Letter} may have important implications on the understanding of the reionization process, the properties of the post-reionization IGM, and many cosmological observations, like the Thomson scattering optical depth of the cosmic microwave background, neutral hydrogen 21-cm transition line, and the Ly$\alpha$ forest. The Pop III binary X-ray luminosity per unit comoving volume is estimated to be $\log I_{\rm X} \simeq (\epsilon_{\rm bin}/0.5)[40.87 - 0.12(1+z)]$ in units of erg s$^{-1}$ Mpc$^{-3}$, which is likely higher than those from other X-ray sources of AGNs and normal XRBs at redshifts higher than 6 \citep{Fragos13}. This linear fit is only valid at $z \ge 10$ and has a maximum error of 0.5\% compared to the computed values. This is based on our assumption at 50\% of all Pop III stars form in binaries, and evolve through a high mass XRB phase. Since it takes a long time for the X-ray background to build up, the contribution from Pop III binaries may continue to dominate the X-ray background down to much lower redshifts. Recently \citet{Fialkov16} used the unresolved soft X-ray background to place constraints on the nature of high redshift X-ray sources. They derived constraints on the X-ray efficiency parameter $f_x$ for three source types and two reionization scenarios, where $f_x$ is defined implicitly by the relation \citep{Furlanetto06,Fragos13} \begin{equation} \frac{L_X}{SFR} = 3\times10^{40} f_X~ \rm{erg\,s}^{-1} \rm{M}_{\odot}^{-1}yr\ , \end{equation} where here $SFR$ refers to the normal (Pop II) star formation rate density. Using the Method 2 value of $L_X \approx 5 \times 10^{39} \mathrm{erg s}^{-1} \mathrm{Mpc}^{-3}$ at $z=7.6$, and the observed SFR of $10^{-2} M_{\odot} \mathrm{yr}^{-1} \mathrm{Mpc}^{-3}$ \citep{Bouwens15}, we get $f_x \approx 16$, well within observational constraints for all cases considered. It is interesting to compare the relative importance of ionizing UV and X-ray photons during the course of the simulations. To $z=7.6$, both Pop III and metal-enriched stars produce about 1.4 $\times$ 10$^{71}$ erg in ionizing UV photons (or about 6.3 $\times$ 10$^{69}$ photons) inside the refined region, of which less than 1\% (10$^{69}$ erg), is produced by Pop III stars. There are about 7.6 $\times$ 10$^{68}$ hydrogen atoms in the survey volume of the \textit{Void} run. At $z=7.6$, only $\sim$ 13.2\% hydrogen atoms have been ionized by the UV ionizing photons in the simulation \citep{Xu16_fesc}. So most of the ionizing UV photons are absorbed by hydrogen atoms that later recombine and are consumed to maintain the ionization state of high density regions. On the contrary, a large fraction of X-rays are expected to penetrate to low density regions which have recombination times longer than the Hubble time. During the same period, the Pop III binaries of the \textit{Void} simulation produced $\sim$ 6.8 $\times$ 10$^{69}$ eV of energy in X-rays inside the refined region. It is about 5\% of energy of ionizing UV photons and about 9 eV per hydrogen atom. Using Method 2, we estimate that Pop III binaries produce in total 2.1 $\times$ 10$^{72}$ eV of energy in X-rays in the entire simulation box. As there are about 3.6 $\times$ 10$^{71}$ hydrogen atoms in the simulation box, this amounts to about 6 eV of X-ray energy per hydrogen atom. This photon budget analysis shows that the X-ray component is not negligible. There is enough energy in X-rays from Pop III binaries to significantly change the thermal and ionization states of the IGM. It is very hard to know the ratio of X-ray energy used in ionizing and heating the IGM without knowing the SED of the X-ray sources and doing a detailed simulation \citep{Shull85}. \citet{Xu14} used a one-zone model to calculate heating and ionizing effects for fixed X-ray photon energies and fluxes and to estimate the temperature and hydrogen ionization fraction of the \textit{Rarepeak} simulation volume by simply assuming the X-ray background is constant since z $=$ 15. This predicted that the IGM may be heated to several hundred Kelvin and be ionized to about $0.5\%$ by $z=6$. We have shown that the X-ray luminosity and background actually increase until the end of the simulation at $z=7.6$ and are at much higher ($\sim$ 10 times) levels. The electron fraction due to X-ray ionization might therefore reach a few percent by simply extrapolating the results in \citet{Xu14}, or even more since X-rays are also efficient at helium ionization. This physics is important to the understanding the optical depth of CMB. In addition, by the end of reionization, the IGM may have been homogeneously heated by the keV photons of the X-ray background for hundreds of Myr. The additional heat input above and beyond the UV photoheating may reach thousands of Kelvin \citep{Xu14}, and that is enough to have important impacts on high redshift 21-cm signatures \citep{Ahn14} and the Ly$\alpha$ forest spectra \citep{Tytler09} at lower redshifts.

1607.02664

1607

1607.05272_arXiv.txt

A principal scientific goal of the Gemini Planet Imager (GPI) is obtaining milliarcsecond astrometry to constrain exoplanet orbits. However, astrometry of directly imaged exoplanets is subject to biases, systematic errors, and speckle noise. Here we describe an analytical procedure to forward model the signal of an exoplanet that accounts for both the observing strategy (angular and spectral differential imaging) and the data reduction method (Karhunen-Lo\`eve Image Projection algorithm). We use this forward model to measure the position of an exoplanet in a Bayesian framework employing Gaussian processes and Markov chain Monte Carlo (MCMC) to account for correlated noise. In the case of GPI data on $\beta$ Pic b, this technique, which we call Bayesian KLIP-FM Astrometry (BKA), outperforms previous techniques and yields 1$\sigma$-errors at or below the one milliarcsecond level. We validate BKA by fitting a Keplerian orbit to twelve GPI observations along with previous astrometry from other instruments. The statistical properties of the residuals confirm that BKA is accurate and correctly estimates astrometric errors. Our constraints on the orbit of $\beta$ Pic b firmly rule out the possibility of a transit of the planet at 10-$\sigma$ significance. However, we confirm that the Hill sphere of $\beta$ Pic b will transit, giving us a rare chance to probe the circumplanetary environment of a young, evolving exoplanet. We provide an ephemeris for photometric monitoring of the Hill sphere transit event, which will begin at the start of April in 2017 and finish at the end of January in 2018.

Astrometry is an essential tool for characterizing directly imaged exoplanets and their physical relationship to other elements of the planetary system in which they reside. The Gemini Planet Imager \citep[GPI;][]{macintosh14} was designed with a goal of achieving $\leq$1.8~mas astrometric accuracy \citep{graham09a}, necessary for characterizing the eccentricity distribution of exoplanet orbits from the GPI Exoplanet Survey \citep{konopacky14}. To do so, the astrometric calibration of GPI has continually been benchmarked to well-calibrated astrometric fields \citep{konopacky14}. This had led to some of the most precise astrometry on directly-imaged exoplanet systems to date \citep{mmb15, derosa15, rameau16}, allowing us to constrain or fit the first ever orbit of some of these directly-imaged exoplanets. However, limited by either the signal to noise ratio (SNR) of these exoplanets or by biases in the various data analysis algorithms, so far no astrometric study with GPI has achieved the design goal of 1.8~mas precision. The importance of understanding planetary orbits is highlighted by the $\beta$ Pictoris system, a young \citep[$\sim23$ Myr;][]{mamajek14, binks16} and nearby \citep[19.3~pc;][]{2007A&A...474..653V} system that has been extensively studied. $\beta$ Pic harbors a near edge-on debris disk that was first imaged by \citet{smith84} and which was subsequently observed to have a warp thought to be induced by a planet whose orbit is inclined relative to the debris disk \citep{burrows95a,mouillet97a,heap00a}. Additional indirect signatures of a planet were derived from variable spectral features modeled as infalling comets \citep{beust00a} and a peculiar light curve anomaly detected in 1981 \citep{lecavelier97}. \citet{lagrange09a,lagrange10} then discovered \betapicb{}, a planet at an appropriate mass ($\sim10$~$M_{\text{Jup}}$) and semi-major axis ($8-13$~AU) to be responsible for the previously observed indirect signatures of planets. A key focus of subsequent observations was determining the alignment of \betapicb{} relative to the main outer disk and the warp to determine if \betapicb{} is causing the warp. By observing the disk and planet simultaneously, \citet{lagrange12} concluded that the planet is misaligned from the main outer disk and consistent with being responsible for the warp. Additionally, \citet{dawson11} ruled out the possibility of having another giant planet in the system massive enough to cause the warp instead. Thus, \betapicb{} is responsible for the warp in the debris disk. This was confirmed in astrometric monitoring campaigns \citep{chauvin12, nielsen14, mmb15} which used homogeneous datasets to limit systematics and constrain the orbit of \betapicb{}. Refining the orbital elements of $\beta$ Pic b is not only crucial for investigating the dynamical link between the planet and the disk warp, but also because $\beta$ Pic b may transit its host star once every $\sim$20 years. To date, there are no other known systems where the physical properties of an exoplanet can be characterized by using both the direct imaging and transit techniques. Currently, the tightest constraint on the probability of transit is $\sim$0.06\%, obtained with a dedicated astrometric monitoring campaign with GPI \citep{mmb15}. However, \citet{lecavelier16} point out that GPI measurements from \citet{mmb15} have a higher position angle than astrometry from previous measurements, which could arise from a possible systematic calibration offset between GPI and previous instruments instead of actually due to \betapicb{}'s orbit being slightly inclined away from edge on. We note there currently is no evidence of a position angle offset with the GPI astrometric calibration, and the GPI astrometry of HD 95086 b are consistent with previous astrometry from other instruments \citep{rameau16}. However, it is important to more accurately compute the orbit of \betapicb{} because in late 2017 to early 2018 \citep{mmb15}, it will be at its closest projected separation from the star. The transit of the planet and/or any circumplanetary material orbiting around it could therefore be detectable. One of the obstacles in characterizing directly imaged exoplanets is that even with the newest instrumentation, the glare of the host star covers the signal of the planet. In order to subtract the point spread function (PSF) of the star and maximize the SNR of the planet, observing techniques such as Angular Differential Imaging \cite[ADI;][]{marois06} and Spectral Differential Imaging \citep[SDI;][]{marois00} and data reduction algorithms like Karhunen-Lo\`eve Image Projection \citep[KLIP;][]{soummer12,pueyo15a} are used in combination to disentangle the PSF of the star from potential astrophysical sources. However, these techniques distort the planet signal and create data reduction artifacts, which usually are nuisance parameters that need to be calibrated out to obtain unbiased astrometry. Forward modelling effects of observing techniques and data reduction algorithms on the PSF of the planet was first done in the context of ADI and LOCI, showing significant improvements in astrometry and photometry for simulated planets \citep{marois10,galicher11}. In similar contexts with ADI and LOCI, \citet{brandt13} and \citet{esposito14} used forward modelling to correct for flux loss of exoplanets and the flux and morphology of disks respectively. For classical ADI (cADI), \citet{cantalloube15} showed that forward modelling techniques can improve the sensitivity of cADI and mitigate biases. The use of Markov-Chain Monte Carlo (MCMC) in conjunction with forward modelling was presented in \citet{bottom15} for reference differential imaging. Recently, \citet{pueyo16} introduced a method, called KLIP-FM, to analytically forward model the degradation of a faint astrophysical signal that occurs when using least squares-based PSF subtraction algorithms such as KLIP that is also generally applicable to any observing strategy. Additionally, the computation of the forward model with KLIP-FM is much quicker than negative fake planet injection methods \citep{marois10, lagrange10}, as the stellar PSF subtraction algorithm, KLIP, needs only to be run once. In this paper, we demonstrate the advantages of KLIP-FM for precise astrometry and constraining planetary orbits by applying it to GPI observations of \betapicb{} reduced using KLIP and ADI+SDI. In Section \ref{sec:obs}, we describe our new astrometry technique, Bayesian KLIP-FM Astrometry (BKA), in which we forward model the PSF of the planet with KLIP-FM and then use the forward model in a Bayesian framework to measure the position of the planet while also modelling the correlated nature of the noise. In Section \ref{sec:valid}, we validate our technique by fitting an orbit to the data and showing that our astrometry and uncertainties are consistent with Keplerian motion with no obvious systematics. Finally, in Section \ref{sec:orbit}, we apply our new astrometry to constrain the orbit of \betapicb{} and place the tightest constraints yet on the transit of the planet and its Hill sphere.

In the first part of this work, we have presented a new technique for more precise and accurate astrometry of directly imaged exoplanets using a new analytical forward modelling approach in a robust statistical framework. \begin{itemize} \item Using the KLIP-FM framework presented in \citet{pueyo16}, we are able to analytically forward model the PSF of the planet through the data reduction process, giving us better information on the location of the planet. We apply KLIP-FM to GPI data on $\beta$~Pic and forward model the PSF of \betapicb{} using the open-source \texttt{pyKLIP} package. \item For a close-in planet orbiting a bright star like in the case of \betapicb{}, we are limited by correlated speckle noise in our data. We developed a Bayesian framework utilizing Gaussian processes and MCMC to account for the correlated noise and to find the position of the planet simultaneously. \item With this technique, we have achieved the most precise astrometry on \betapicb{} to date. On most of our GPI datasets, we achieve $\sim$1~mas precision on the relative separation between \betapicb{} and its host star, a $\sim$2-4 fold improvement over previous techniques using the same data \citep{mmb15}. \item In datasets where the astrometry is limited by noise and not by astrometric calibration uncertainty, Bayesian KLIP-FM Astrometry approach should improve astrometric precision. \end{itemize} In the second part of this work, we apply our Bayesian KLIP-FM Astrometry technique to the orbit of \betapicb{}. \begin{itemize} \item To validate this new astrometric technique, we used it to measure the position of \betapicb{} in twelve epochs of GPI data. We combined these twelve astrometric points with two previous astrometric monitoring campaigns and fit a Keplerian orbit using MCMC methods. We find the residuals to the fit are consistent with zero and show no apparent systematic trends, indicating that our fit is accurate and the uncertainties we estimate are reliable. \item Due to the improved PA measurements from our technique, we have the tightest constraints on the inclination of the orbit and can exclude a possible transit of \betapicb{} at 10-$\sigma$ significance. \item While the planet will not transit, we are confident the Hill sphere around \betapicb{} will transit. The Hill sphere will begin transit at the start of April in 2017 and finish transiting at the end of January in 2018 with closest approach in the end of August in 2017. The transit of \betapicb{}'s Hill sphere should be our best chance in the near future to investigate young circumplanetary material. \end{itemize} In the future, this MCMC forward modelling technique can be applied to photometry and spectral extraction alongside of astrometry of directly imaged exoplanets, allowing for improved characterization of their atmospheres. For \betapicb{}, continued monitoring of its orbit will yield more insight into the dynamics of the star system, although the planet will soon be too close to its star to be seen with current direct imaging instrumentation. However, once the planet appears on the other side, continued astrometric monitoring should be able to constrain the semi-major axis and eccentricity of the orbit much better, which will improve our understanding of how \betapicb{} perturbs the disk and if there are other planets perturbing \betapicb{}.

1607.05272

1607

1607.00382_arXiv.txt

{Galaxies located in the environment or on the line of sight towards gravitational lenses can significantly affect lensing observables, and can lead to systematic errors on the measurement of $H_0$ from the time-delay technique. We present the results of a systematic spectroscopic identification of the galaxies in the field of view of the lensed quasar \HEofor\, using the W.~M.~Keck, Gemini and ESO-Very Large telescopes. Our new catalog triples the number of known galaxy redshifts in the direct vicinity of the lens, expanding to 102 the number of measured redshifts for galaxies separated by less than 3\arcmin\,from the lens. We complement our catalog with literature data to gather redshifts up to 15\arcmin\, from the lens, and search for galaxy \cref{groups or clusters} projected towards \HEofor. We confirm that the lens is a member of a small group that includes at least 12 galaxies, and find 8 other group candidates near the line of sight of the lens. The {\it flexion shift}, namely the shift of lensed images produced by high order perturbation of the lens potential, is calculated for each galaxy/group and used to identify which objects produce the largest perturbation of the lens potential. This analysis demonstrates that i) at most three of the five brightest galaxies projected within 12\arcsec\,of the lens need to be explicitly used in the lens models, and ii) the groups can be treated in the lens model as an external tidal field (shear) contribution.}

\label{sec:intro} Ongoing and upcoming cosmological studies deeply rely on the accurate knowledge of the Hubble constant, $H_0$ \citep{Hu2005, Suyu2012, Weinberg2013}. The measurement of $H_0$ has long been controversial \citep[e.g.][]{KOC02, KochanekSchechter2004}, but in the past decade several techniques have measured $H_0$ with a relative uncertainty much smaller than 10\% \citep{Freedman2010, Humphreys2013, Suyu2013a, Riess2016}. In order to reach the goal of the next decade of cosmological experiments, and be able to e.g. unveil the nature of dark energy, it is necessary to pin down the accuracy on $H_0$ at the percent level. This is an ambitious goal and in order to identify unknown systematic errors, it is mandatory to gather several independent constraints on $H_0$ \citep{Weinberg2013, Riess2016}. The gravitational time-delay technique \citep{Refsdal1964}, applied to a large number of lensed systems, is one of the few techniques allowing one to reach percent precision on $H_0$ \citep{Suyu2012}. Among the various cosmological probes, it is also the most sensitive to $H_0$ \citep[e.g.][]{Jackson2007, Freedman2010}. By measuring the time delay $\Delta t$ between pairs of lensed images, and modeling the mass distribution of the lens galaxy, the time delay distance $D_{\Delta t}$ can be inferred. As summarized in a recent review by \cite{Treu2016}, the technique has long been plagued by poor time-delay measurements, invalid assumptions about the lens mass profile and systematic errors. However, times have changed. It has been demonstrated that an exhaustive study of a lensed quasar with high quality lightcurves \citep[B1608+656;][]{Fassnacht2002} allows the measurement of $H_0$ for a single system with a precision of 6\% \citep{Suyu2010}. In addition, it was shown that the time-delay technique leads to tight constraints on the other cosmological parameters comparable to those from contemporary Baryon Acoustic Peak studies, when each probe is combined with the Cosmic Microwave Background \citep{PLXVI, Anderson2014, Planck2015XIII, Ross2015}. The improved precision of the time delay technique stems from the combination of several ingredients. First, \cref{the COSmological MOnitoring of GRAvItational Lenses (COSMOGRAIL)} has been running for over a decade, gathering exquisite high cadence photometric data for tens of lensed quasars \citep{Eigenbrod2005, Tewes2013}. Those unprecedented high quality lightcurves combined with new curve shifting algorithms \citep{Tewes2013a} now enable time-delay measurements down to a few percent accuracy \citep{Bonvin2016, Liao2016}. Second, advanced modeling techniques that use the full surface brightness of the multiple lensed images, containing thousands of pixels as data points, are now used to constrain the lens mass distribution \citep{Suyu2009}. Third, independent constraints on the lens potential, obtained from the measurement of the lens velocity dispersion \citep{Romanowsky1999, Treu2002}, are now combined with the lens models, enabling one to reduce the impact of the mass-sheet degeneracy\footnote{The impact on cosmographic inference of other degeneracies among lens models, such as the source position transformation \citep{SS14, Unruh2016}, that does not leave the time-delay ratio invariant, still needs to be quantified.} \citep{FGS85, SS13} on the lens models. Finally, the direct lens environment and the line-of-sight galaxies are studied in detail \citep{Keeton2004,Fassnacht2006}. The observed galaxy counts in the vicinity of the lens are compared to galaxy counts from ray tracing through cosmological simulations to derive a probability distribution of the external convergence $\kappa_{\rm ext}$ produced by over- and under-densities along the line of sight \citep{Hilbert2007, Fassnacht2011}. The \HOLI project ($H_0$ Lenses in COSMOGRAIL's Wellspring) aims at achieving better than 3.5\% accuracy on $H_0$. To reach this goal, we have gathered a sample of five lenses (B\,1608$+$656, RX\,J1131$-$1231, \HEofor, \HEeleven, \WFItwenty) for which we apply our modeling technique on archival and Cycle 20 \emph{HST} data (PI Suyu). The project, together with cosmographic forecasts based on the full sample, is presented in \HOLI Paper I (Suyu et al., submitted). The first two systems have been analyzed \citep{Suyu2010, Suyu2013a}. To tackle systematic errors in the other three systems, a stellar velocity dispersion for the lenses and a study of the lens environments are needed. In this paper, we focus on the spectroscopic identification of the brightest galaxies in the field of view of \HEofor, a quadruply imaged quasar at $z_{\rm s} = 1.693 \pm 0.001$ lensed by a foreground elliptical galaxy at $z_{\rm d} = 0.4546 \pm 0.0002$ \citep{Wisotzki2002, Morgan2005, Sluse2012b}. The main objective of this work is to measure the spectroscopic redshifts of most of the bright galaxies in the central region around \HEofor\, (i.e. about 100 galaxies), a necessary observable to measure the contribution of individual galaxy halos to the surface mass density projected towards \HEofor\, \citep{Hilbert2007, Hilbert2009, Greene2013, Collet2013}. Our secondary objective is to identify major groups and/or galaxy cluster(s), as well as individual galaxies, at the redshift of the main lens but also along the line of sight, that would perturb non linearly the gravitational potential of the main lensing galaxy. For that purpose we complement our data with the spectroscopic catalog compiled by Momcheva et al. (\citeyear[][hereafter MOM15]{Momcheva2015}) that gathers redshifts of $\sim$ 400 galaxies (about 30 galaxies are duplicated with our catalog) over a 30\arcmin$\times$30\arcmin\, field centered on \HEofor. The spectroscopic redshift measurements are an important ingredient of the statistical analysis of the line of sight towards \HEofor\, carried out in the companion \HOLI Paper III (Rusu et al., submitted). This companion paper presents a weighted count analysis of the galaxies in the field of view of \HEofor\, that is compared to galaxy counts from the Canada-France-Hawaii-Telescope Legacy Survey \citep[CFHTLenS, ][]{Heymans2012} and to galaxy counts from Millennium Simulation \citep{Springel2005, Hilbert2007, Hilbert2009}. This yields a probability distribution of convergence $\kappa_{\rm ext}$ produced by the other galaxies in the field. On the other hand, the redshifts of the galaxies closest in projection to the lens are included explicitly in the multi-plane lens modeling analysis of \HEofor\, presented in \HOLI Paper IV (Wong et al., submitted). Finally, Paper V (Bonvin et al., submitted) presents the time-delay measurements of \HEofor\,and the joint cosmographic inference from the three lensed systems analyzed to-date in \HOLI. The paper is structured as follows. We present an overview of the data sets used, of the data reduction process and redshift measurements in Sect.~\ref{sec:data}. The methodology used to identify galaxy groups is explained in Sect.~\ref{sec:groups}. The galaxy groups identified with our algorithm and the spectra of the galaxies that \cref{are most likely to produce large gravitational potential perturbations} are presented in Sect.~\ref{sec:environment}. Section ~\ref{sec:model} quantifies the impact of individual galaxies and galaxy groups on the model. We use the \emph{flexion shift} to flag the systems that require explicit inclusion in the multi-plane lens models presented in \HOLI Paper IV. Finally, Sect.~\ref{sec:conclude} summarizes our main results. In this work, with the exception of the target selection that was based on $R-$band magnitude in the Vega system, photometric information comes from the deep multicolor imaging presented in \HOLI Paper III and uses the AB photometric system. For convenience, group radii and masses reported in this work assume a flat $\Lambda$CDM cosmology with cosmological parameters from \citep{Planck2015XIII}, namely $H_0=67.7 \,\kmsMpc$, $\Omega_m$ =\,0.307. We stress that this choice has no impact on the group identification as the latter does not depend on a specific choice of cosmological parameters.

\label{sec:conclude} \cref {We have used multi-objects spectrographs on ESO-VLT, Keck and Gemini telescopes to measure the redshifts of 65 galaxies (down to $i = 23$\,mag) within a field of $\sim 4\arcmin$\, radius centered on \HEofor. In addition, our spectroscopic sample contains 18 galaxies with tentative redshifts, and 46 objects which had uncertain photometric classification, but turn out to be stars in our Galaxy. We have complemented our catalog with independent spectroscopic data sets compiled by MOM15. This expands the number of confirmed (or tentative) spectroscopic redshift in the field of \HEofor\, to 425 galaxies, up to a projected distance of 15\arcmin\, from the lens. Both the spectroscopic catalog and associated spectra are made publicly available with this paper. } The analysis of this new data set, combined with deep multicolor ($ugri$) photometric data covering the same field of view and presented in the companion \HOLI Paper III (Rusu et al., submitted), yields the following important results: \begin{enumerate} \item The redshifts of the five brightest galaxies that fall within 12\arcsec\, of the lens (G1 - G5, with $i \in [19.9, 22.1]$\,mag), are measured to be $z_{\rm G1}=0.7821, z_{\rm G2}=0.7806, z_{\rm G3}=0.4191, z_{\rm G4}=0.4568, z_{\rm G5}=0.7792$, \cref{with a typical random uncertainty of $\sigma_z(\rm ran) \sim 0.0002$, and a possible systematic uncertainty $\sigma_z(\rm sys) \sim 0.0004$.} \item In order to pinpoint the galaxies that are most likely to produce high order perturbations of the gravitational potential of the main lens, we have derived the so called flexion shift $\Delta_3 x$ \citep{McCully2016} of each individual galaxy in the field. \cite{McCully2016} suggest that $\Delta_3 x \sim 10^{-4}$ arcseconds is a conservative threshold above which a perturber is susceptible producing a bias at the percent level on $H_0$ if not included explicitly in the lens model. The largest flexion shift is found for G1 for which we get $\Delta_3 x(\rm {G1})\sim 8 \times 10^{-4}$\,arcseconds. This motivates the explicit inclusion of this galaxy in all the lens models of \HEofor\, presented in the companion \HOLI Paper IV (Wong et al., submitted). \cref{The two galaxies G3 and G4 are also found to have flexion shifts close to or above $10^{-4}$\,arcseconds} such that they are also included in one of the lens models presented in \HOLI Paper IV. \item We search for galaxy groups or clusters in the field of view of \HEofor\, using an iterative algorithm similar to those developed by \cite{Wilman2005}, \cite{Calvi2011} and \cite{Ammons2014}. Our iterative method identifies group members based on the joint separation of galaxies projected on the sky and redshift space. We have searched for galaxy groups of at least 5 members in the inner 6\arcmin\,around the lens, where our spectroscopic completeness is the highest, and for groups of at least 10 members at larger distance from the lens. No evidence for a massive galaxy cluster was found, but 9 galaxy groups (0.05 $< \bar{z}_{\rm group} < 0.8$) with velocity dispersion $\sigma_{int} < 500$\,\kms\,(some groups being possibly bimodal) were identified. One of these groups includes the lensing galaxy. It has been previously reported by \cite{Wong2011} with one less member, and is independently found by \cite{Wilson2016} based on the catalog published by MOM15. \item The impact of the groups on the lens model is more difficult to determine than for individual galaxies because of the uncertainty on the position of the group centroid. Fixing the group centroid to the brightest (spectroscopically confirmed) group member yields $\Delta_3 x < 10^{-4}$\,arcseconds for every group. A similar result is found when fixing the group centroid to the luminosity/mass weighted centroid of the identified group members. The centroid uncertainty has little impact on these conclusions for most of the groups but for the group hosting the lens. In that case, a shift of the luminosity/mass weighted centroid (found $\sim 70\arcsec$ from the lens), by more than 20\arcsec\, towards the lens would yield a flexion shift a few times $10^{-4}$\,arcseconds. We think that such a shift is unlikely as we have good evidence that we identified all the members of that group down to $i\sim 22\,$ mag. \end{enumerate} Our spectroscopic study demonstrates that \HEofor\, requires an explicit inclusion of the nearest galaxy G1, while the galaxies G2-G5, produce smaller, but potentially non-negligible, perturbation of the gravitational potential of the main lens. On the other side, galaxy groups are unlikely to produce significant perturbations. This is confirmed by the weighted number counts analysis of the field of \HEofor\, presented in \HOLI Paper III, that shows that the line of sight is not particularly overdense, with an external convergence $\kappa_{\rm ext} = 0.003\pm 0.025$. The small convergence produced by the lens environment is confirmed by the weak lensing study of the field of view (Tihhanova et al., in prep) that shows that the total external convergence towards \HEofor\, is $\kappa_{\rm ext} < 0.04$ at $3\sigma$. This motivates the lens models presented in \HOLI Paper IV where only galaxy $G1$ is included explicitly in all the lens models using a mutiplane formalism, while a distribution of the convergence produced by the other galaxies (\HOLI Paper III), is used to account for the other galaxies. We are completing the analysis of the spectroscopic environment of the next two \HOLI lensed systems, \HEeleven\, and \WFItwenty. The much richer line-of-sight environment of these two systems may produce stronger systematic errors on $H_0$ if not carefully accounted for in the lens models, making spectroscopic characterization of the lens environment a key ingredient of cosmography with time-delay lenses.

1607.00382

1607

1607.07107_arXiv.txt

We extend the scalar-tensor reconstruction techniques for classical cosmology frameworks, in the context of loop quantum cosmology. After presenting in some detail how the equations are generalized in the loop quantum cosmology case, we discuss which new features and limitations does the quantum framework brings along, and we use various illustrative examples in order to demonstrate how the method works. As we show the energy density has two different classes of solutions, and one of these yields the correct classical limit while the second captures the quantum phenomena. We study in detail the scalar tensor reconstruction method for both these solutions. Also we discuss some scenarios for which the Hubble rate becomes unbounded at finite time, which corresponds for example in a case that a Big Rip occurs. As we show this issue is non-trivial and we discuss how this case should be treated in a consistent way. Finally, we investigate how the classical stability conditions for the scalar-tensor solutions are generalized in the loop quantum framework.

The strikingly unexpected observation of the late-time acceleration of the Universe in the late 90's \cite{riess} has set the stage for the construction of alternative theories of gravity to model the Universe. Up to date, it is believed that the Universe experienced two acceleration eras, the early-time acceleration and the late-time acceleration era. One characteristic feature of the early-time acceleration era is the production of a slightly red tilted scale invariant power spectrum of primordial curvature perturbations, which has recently been verified by the Planck data \cite{planck}. One of the successful theories that produce a nearly scale invariant spectrum is the inflationary scenario \cite{inflation1,inflation2,inflation2a,inflation2b,inflation2c,inflation3,inflation4,inflation5,inflation6}. However, an alternative scenario to the standard inflationary paradigm is the big bounce evolution \cite{brande1,bounce1,bounce1a,bounce1b,bounce2,bounce3,bounce4,bounce5,matterbounce1,matterbounce2,matterbounce3,matterbounce4,matterbounce5,matterbounce6,matterbounce7}, in which the initial singularity is avoided and also it is possible to produce a scale-invariant power spectrum. Particularly, it was known for quite some time that an inflationary de Sitter evolution and a contracting cosmological phase with scale factor $a(t)\sim (-t)^{2/3}$, are related by a duality \cite{duality}, and both produce scale invariant spectrum. A well known cosmological bounce which realizes a contracting phase which produces an exactly scale invariant spectrum is the matter bounce scenario \cite{matterbounce1,matterbounce2,matterbounce3,matterbounce4,matterbounce5,matterbounce6,matterbounce7}, which naturally arises in the context of Loop Quantum Cosmology (LQC) \cite{LQC1,LQC2,LQC3,LQC4,LQC5,LQC6,LQC19}, if the matter content consists of a pressureless perfect fluid. Scalar fields are frequently used in order to describe inflationary theories \cite{inflation1,inflation2,inflation2a,inflation2b,inflation2c,inflation3,inflation4,inflation5}, and in order to describe the late-time acceleration era \cite{scalarrecon0,scalarrecon1,scalarrecon2,scalarrecon3,scalarrecon4,scalarrecon5,scalarrecon6,scalarrecon7,scalarrecon8,scalarrecon9}. Particularly, several reconstruction techniques use canonical or non-canonical scalar fields \cite{scalarrecon0,scalarrecon1,scalarrecon2,scalarrecon3,scalarrecon4,scalarrecon5,scalarrecon6,scalarrecon7,scalarrecon8,scalarrecon9} in order to generate a quintessential or even a phantom late-time evolution. The purpose of this paper is to generalize the reconstruction methods of Refs. \cite{scalarrecon0,scalarrecon1,scalarrecon2,scalarrecon3} in the context of LQC. We aim to present the general method of realizing a given cosmological evolution in terms of it's Hubble rate and scale factor, and we investigate how the classical results obtained in Refs. \cite{scalarrecon0,scalarrecon1,scalarrecon2,scalarrecon3} are generalized in the case of LQC. As we will see, the LQC resulting equations are identical to the classical equations, when the classical limit is taken. Both the non-canonical and canonical scalar fields cases shall be discussed, and also we use several illustrative examples to show how the method works. We also discuss the limitations of the method and also we highlight the difference with the classical case. In addition we discuss how the method works in the case that the quantum era is considered, since there are two branches of solutions for the energy density. Also the case that the Hubble rate is unbounded is discussed in brief, in the case of a Big Rip where the unboundedness occurs at a finite time. In the context of LQC, this requires special attention, as was demonstrated in \cite{rev1}, since the Friedmann equations have a different form, however the Big Rip singularity should be avoided as it happens in all LQC frameworks, see for example \cite{nobigrip1,nobigrip1a,nobigrip1b,nobigrip1c,nobigrip2,nobigrip3,nobigrip4,haronobigrip}. We point out the problem of this issue and we discuss how the problem is actually solved by using the arguments of Ref. \cite{rev1}. This paper is organized as follows: In section II, we discuss the general reconstruction technique for non-canonical LQC scalar fields. In section III we address the stability of the general solution we obtained in section II, and we compare the results with the classical case. In section IV we investigate how phantom and oscillating cosmologies can be generated by LQC non-canonical scalar fields, and we discuss the case of unbounded Hubble rates at finite time. The canonical scalar field case is discussed in section V and the concluding remarks along with a discussion of the results follow in the end of the paper.

In this article we extended the scalar-tensor reconstruction techniques for realizing cosmological evolutions in the context of LQC. We presented the basic equations that constitute the LQC reconstruction method and we discussed the limitations of the method. Several examples were presented in order to demonstrate how the method works and also to show the new constraints that the LQC framework brings along. As we showed, it is possible to realize various cosmological scenarios and particularly certain features of a viable cosmology can be generated, for example the late and early-time acceleration era, phantom or quintessential evolution and even transitions between phantom and quintessential accelerations. As we showed the energy density has two branches of solutions with one yielding the classical limit and the other capturing the quantum phenomena. We discussed how the reconstruction method works in both these cases. We also addressed the case that the Hubble rate can be come unbounded at finite time. This issue is non-trivial and by adopting the method we used in the previous sections lead to inconsistencies. However, by using the right theoretical context presented in Ref. \cite{rev1}, the inconsistencies do not occur and formally the Rip singularities can be avoided. In the case of non-canonical scalar fields, we also addressed the stability issue of the solution we proposed, and in all cases, the stability conditions are direct generalizations of the classical stability conditions, with the two coinciding in the classical limit $\rho_c\to \infty$. We also discussed the canonical scalar field case, and we studied the case of a perfect fluid with constant equation of state parameter $w$ and also we performed a numerical analysis for an example that was difficult to study analytically. A direct promising extension of the scalar-tensor LQC reconstruction method we proposed in this paper, is to use several scalar fields. This extension will provide a framework in which several cosmological scenarios could be realized, and also there is always the possibility for some scalar fields to be phantom and with the rest being non-phantom. Also the appearance of several scalars offers more freedom in realizing various evolution scenarios, so this theoretical extension should be worked out in detail in a future work.

1607.07107

1607

1607.02514_arXiv.txt

{ We develope a computationally efficient extension of the Dyson Brownian Motion (DBM) algorithm to generate random function in $C^2$ locally. We further explain that random functions generated via DBM show an unstable growth as the traversed distance increases. This feature restricts the use of such functions considerably if they are to be used to model globally defined ones. The latter is the case if one uses random functions to model landscapes in string theory. We provide a concrete example, based on a simple axionic potential often used in cosmology, to highlight this problem and also offer an ad hoc modification of DBM that suppresses this growth to some degree. }

% \label{sec:intro}% Generating random functions locally via Dyson Brownian Motion (DBM) was first introduced in \cite{Marsh:2013qca} following prior work by Dyson \cite{Dyson:62}, see \cite{Mehta:1991} for a textbook, and applied to potentials in cosmology in \cite{Marsh:2013qca,Battefeld:2014qoa,Dias:2016slx}. The algorithm starts with a Taylor expansion of the potential to second order and, after moving a set distance from the initial expansion point, adding random perturbations to the matrix of second derivatives\footnote{Dyson Brownian motion \cite{Dyson:62} was developed primarily to introduce a ``time'' dependence into Random Matrix Theory, which enables the computation of transition probabilities. Analytic methods to compute such probabilities were derived in \cite{Pedro:2016jyd}.}. Thus, a random function is generated locally along a trajectory. The algorithm as introduced in \cite{Marsh:2013qca} generates a function in $C^1$ along a trajectory\footnote{The continuum limit of the algorithm yields a function in $C^2$. If implemented numerically and the step length of the DBM function is kept in line with the discretization used elsewhere, the functions introduced in \cite{Marsh:2013qca} are indistinguishable from functions in $C^2$. However, using functions that are in $C^2$ by construction can be computationally advantageous.} with a Hessian in the Gaussian orthogonal ensemble. In \cite{Battefeld:2014qoa}, the method was generalized to generate functions in any differentiability class by relegating the perturbations to a higher order derivative tensor, with particular emphasis on functions in $C^2$. Such functions are of particular interest to model the potential in inflationary cosmology, see \cite{Baumann:2014nda,Marsh:2015xka} for recent reviews, while enabling numerical computation of the power-spectrum of cosmological fluctuations \footnote{In \cite{Dias:2016slx}, semi-analytic methods were used to compute observables. } without generating ringing or requiring a step size for DBM in line with the one used for the discretization of the required differential equations. However, the methods introduced in \cite{Battefeld:2014qoa} are computationally more demanding than ordinary DBM, since a coordinate rotation is needed at each step. In this brief technical note we improve upon the algorithm presented in \cite{Battefeld:2014qoa} to generate a function in $C^2$ without the need to rotate the Hessian, see Sec.~\ref{Sec:extDBM}. As a consequence, field spaces with more than 100 fields, as needed to model certain landscapes in string theory/cosmology, can be handled. We also explain a severe shortcoming of functions generated via DBM or its generalization: while the elements of the Hessian have a stable distribution, the function itself performs a random walk and is thus unbounded. Simply using the appropriate distribution for a Hessian in a bounded potential, such as the axionic one investigated in \cite{Wang:2015rel}, instead of one in the Gaussian orthogonal ensemble does not alleviate the problem. This shortcoming renders DBM potentials ill suited to model many landscapes of interest in cosmology if one is interested in going beyond the first coherent patch. We explain this shortcoming in a concrete case-study in Sec.~\ref{Sec:axionic} and provide a minor modification of DBM that can suppress the unstable growth to some degree. We refer the interested reader to \cite{Marsh:2013qca,Battefeld:2014qoa,Dias:2016slx} for a more pedagogical introduction to DBM and cosmological implications, while focusing on technical aspects in this note.

We provided a computationally efficient extension of the Dyson Brownian Motion algorithm to generate random function in $C^2$ locally, as desirable for certain applications in cosmology. We further showed at the example of a simple globally defined potential that DBM potentials fail to recover basic features of the globally defined ones, due to the presence of an unstable growth of the gradient's and the potential's variance. We also showed that a minor ad hoc modification of the algorithm can weaken the unstable behaviour, such that a DBM potential can mimic a globally defined one reasonable well for some time. However, the unstable growth is still present. Thus, in their current form, DBM potentials can be used to model landscapes in String Theory or to draw conclusions in cosmological settings only for regions not much bigger than their coherence length.

1607.02514

1607

1607.01047_arXiv.txt

Based on spectroscopy and multiband wide-field observations of the gravitationally lensed quasar \hequad , we determine the probability distribution function of the external convergence $\kext$ for this system. We measure the under/overdensity of the line of sight towards the lens system and compare it to the average line of sight throughout the universe, determined by using the CFHTLenS as a control field. Aiming to constrain $\kext$ as tightly as possible, we determine under/overdensities using various combinations of relevant informative weighing schemes for the galaxy counts, such as projected distance to the lens, redshift, and stellar mass. We then convert the measured under/overdensities into a $\kext$ distribution, using ray-tracing through the Millennium Simulation. We explore several limiting magnitudes and apertures, and account for systematic and statistical uncertainties relevant to the quality of the observational data, which we further test through simulations. Our most robust estimate of $\kext$ has a median value $\kappa^\mathrm{med}_\mathrm{ext} = 0.004$ and a standard deviation of $\sigma_\kappa = 0.025$. The measured $\sigma_\kappa$ corresponds to $2.5\%$ uncertainty on the time delay distance, and hence the Hubble constant $H_0$ inference from this system. The median $\kappa^\mathrm{med}_\mathrm{ext}$ value is robust to $\sim0.005$ (i.e. $\sim0.5\%$ on $H_0$) regardless of the adopted aperture radius, limiting magnitude and weighting scheme, as long as the latter incorporates galaxy number counts, the projected distance to the main lens, and a prior on the external shear obtained from mass modelling. The availability of a well-constrained $\kext$ makes \hequad\ a valuable system for measuring cosmological parameters using strong gravitational lens time delays.

\label{section:intro} By measuring time delays between the multiple images of a source with time-varying luminosity, strong gravitational lens systems with measured time delays can be used to measure cosmological distances and the Hubble constant $H_0$ \citep{Refsdal64}. In particular, for a lens system with a strong deflector at a single redshift, one may infer the \lq{}time-delay distance\rq{} \begin{equation} \tdist = (1+\zd) \frac{\Dd \Ds}{\Dds}, \label{eq:tdist} \end{equation} where $\zd$ denotes the redshift of the foreground deflector, $\Dd$ the angular diameter distance to the deflector, $\Ds$ the angular diameter distance to the source, and $\Dds$ the angular diameter distance between the deflector and the source. The time-delay distance is primarily sensitive to the Hubble constant, i.e. $\tdist \propto H_0^{-1}$ \citep[see][for a recent review]{treu16}. Inferring cosmological distances from measured time delays also requires accurate models for the mass distribution of the main deflector and its environment, as well as for any other matter structures along the line of sight that may influence the observed images and time delays \citep{suyu10}. Galaxies very close in projection to the main deflector often cause measurable higher-order perturbations in the lensed images and time delays and require explicit models of their matter distribution. The effect of galaxies more distant in projection is primarily a small additional uniform focusing of the light from the source. Furthermore, matter underdensities along the line of sight such as voids, indicated by a low galaxy number density, cause a slight defocusing. For a strong lensing system with a main deflector at a single redshift, the net effect of the (de)focusing by these weak perturbers is equivalent (to lowest relevant order) to that of a constant external convergence\footnote{The external convergence $\kext$ may be positive or negative depending on whether focusing or defocusing outweighs the other.} term $\kext$ in the lens model for the main deflector \citep{suyu10}. This implies on the one hand that the weak perturbers' effects, i.e. the external convergence they induce, cannot be inferred from the observed strongly lensed image properties alone due to the \lq{}mass-sheet degeneracy\rq{} \citep[MSD,][]{falco85,schneider13}. On the other hand, if the external convergence is somehow determined from ancillary data, and a time-delay distance $\tdist^{(0)}$ has been inferred using a model not accounting for the effects of weak perturbers along the line of sight, the true time-delay distance $\tdist$ can simply be computed by: \begin{equation} \tdist = \frac{\tdist^{(0)}}{1 - \kext}. \label{eq:tdist_and_external_convergence} \end{equation} This relation makes clear that any statistical and systematic uncertainties in the external convergence due to structures along the line of sight directly translate into statistical and systematic errors in the inferred time delay distance and Hubble constant: \begin{equation} H_0 = (1 - \kext) H_0^{(0)}, \label{eq:H} \end{equation} where $H_0^{(0)}$ denotes the Hubble constant inferred when neglecting weak external perturbers. With reduced uncertainties on other component of the time delay distance measurement from state-of-the-art imaging, time-delay measurements, and modeling techniques of strong lens systems, the external convergence $\kext$ is now left as an important source of uncertainty on the inferred $H_0$, contributing up to $\sim 5\%$ to the error budget on $H_0$ \citep{suyu10,suyu13}. Moreover, the mean external convergence may not vanish for an ensemble of lens systems due to selection effects, causing a slight preference for lens systems with overdense lines of sights \citep{collett16}. Thus, an ensemble analysis simply assuming $\kext=0$ is expected to systematically overestimate the Hubble constant $H_0$. Accurately quantifying the distribution of mass along the line of sight requires wide-field imaging and spectroscopy \citep[e.g.,][see \citet{treu16} for a recent review]{keeton04,fassnacht06,momcheva06,fassnacht11,wong11}. \citet{suyu10} pioneered the idea of estimating a probability distribution function $P(\kext)$ by (i) measuring the galaxy number counts around a lens system, (ii) comparing the resulting counts against those of a control field to obtain relative counts, and (iii) selecting lines of sight of similar relative counts, along with their associated convergence values, from a numerical simulation of cosmic structure evolution. To this end, \citet{fassnacht11} measured the galaxy number counts in a $45\arcsec$ aperture around \hequad\ [$\alpha$(2000):~04h~38m~14.9s, $\delta$(2000):~-12$^\circ$17\arcmin14\farcs4; \citealp{wisotzki00,wisotzki02}; lens redshift $\zd = 0.455$; \citealp{morgan05}; source redshift $\zs = 1.693$; \citealp{sluse12}], and found that it is 0.89 of that on an average line of sight through their $\sim0.06 \deg^2$ control field. Both \citet[][hereafter G13]{greene13} and \citet{collett13} find that $P(\kext)$ can be most precisely constrained for lens systems along underdense lines of sight, making \hequad\ a valuable system. Recent work has focused on tightening the constraints on $P(\kext)$ with data beyond simple galaxy counts. \citet{suyu13} used the external shear inferred from lens modelling as a further constraint, which significantly affected the inferred external convergence due to the large external shear required by the lens model. G13 extended the number counts technique by considering more informative, physically relevant weights, such as galaxy redshift, stellar mass, and projected separation from the line of sight. Both of these works used ray-tracing through the Millennium Simulation \citep[][hereafter MS]{springel05,hilbert09} in order to obtain $P(\kext)$. For lines of sight which are either underdense or of common density, G13 found that the residual uncertainty $\sigma_{\kext}$ on the external convergence can be reduced to $\lesssim 0.03$, which corresponds to an uncertainty on time delay distance and hence $H_0$ comparable to that arising from the mass model of the deflector and its immediate environment. Furthermore, \citet{collett13} considered a reconstruction of the mass distribution along the line of sight using a galaxy halo model. They convert the observed environment around a lens directly into an external convergence, after calibrating for the effect of dark structures and voids by using the MS. We have collected sufficient observational data to implement these techniques for the case of \hequad . We choose to adopt the G13 approach, with several improvements. We first aim to understand and account for various sources of error in our observational data for \hequad , as well as that of CFHTLenS \citep{heymans12}, which we choose as our control field. Second, we incorporate our understanding of these uncertainties into the simulated catalogues of the MS, in order to ensure a realistic estimate of $P(\kext)$. Third, we use the MS to test the robustness of this estimate for simulated fields of similar under/overdensity. This paper is organized as follows. In Section \ref{section:data} we present the relevant observational data for \hequad\, and its reduction. In Section \ref{section:CFHT} we present an overview of our control field, CFHTLenS. In Section \ref{section:lens} we present our source detection, classification, photometric redshift and stellar mass estimation, carefully designed to match the CFHTLenS fields. In Section \ref{section:overdensity} we present our technique to measure weighted galaxy count ratios for \hequad , by accounting for relevant errors. In Section \ref{section:kappa} we use ray-tracing through the MS in order to obtain $P(\kext)$ for the measured ratios, and present our tests for robustness. We present and discuss our results in Section \ref{section:discuss}, and we conclude in Section \ref{section:concl}. We present additional details in the Appendix. The current work represents Paper III (hereafter H0LiCOW Paper III) in a series of five papers from the H0LiCOW collaboration, which together aim to obtain an accurate and precise estimate of $H_0$ from a comprehensive modelling of \hequad . An overview of this collaboration can be found in H0LiCOW Paper I (Suyu et al., submitted), and the derivation of $H_0$ is presented in H0LiCOW Paper V (Bonvin et al., submitted). Throughout this paper, we assume the MS cosmology, $\Om=0.25$, $\OL=0.75$, $h=0.73$, $\sigma_8 = 0.9$.\footnotemark \footnotetext{We estimate the impact of using a different cosmology in Section \ref{section:cosm_depend}.} We present all magnitudes in the AB system, where we use the following conversion factor between the Vega and the AB systems: $J_\mathrm{AB} = J_\mathrm{Vega} + 0.91$, $H_\mathrm{AB} = H_\mathrm{Vega} + 1.35$ and $K_{s\ \mathrm{AB}} = K_{s\ \mathrm{Vega}} + 1.83$\footnotemark \footnotetext{Results based on the MOIRCS filters, available at \url{http://www.astro.yale.edu/eazy/filters/v8/FILTER.RES.v8.R300.info.txt}}. We define all standard deviations as the semi-difference between the 84 and 16 percentiles.

\label{section:concl} In this work, we aimed to estimate a robust probability distribution function of the external convergence for \hequad , in order to enable the use of this lens system as an accurate probe of $H_0$. We used spectroscopy and multiband images of the \hequad \ field, and we used the wide component of CFHTLenS as a control field. Building on the work by G13, we refined the method in order to cope with the large fraction of masks in our control field, and we also used more robust medians rather than sums in order to compare weighted counts. We thoroughly explored sources of error in our data sets, such as mask coverage, galaxy-star classification, detection efficiency etc.; we propagated these into the computation of weighted count ratios, finding that the \hequad \ field is more overdense, in terms of number counts, than previously estimated. We used the whole extent of the MS to simulate photometric data of the same quality, and connect the MS lensing convergence catalogue to synthetic weighted count ratios estimated in a similar way. We than estimated the probability distribution function of the external convergence for fields similar in overdensity to \hequad, in a Bayesian, unbiased way. We considered multiple aperture radii and limiting magnitudes, and tested them using the MS, finding that a $45\arcsec$ aperture and a limiting magnitude of $i\leq23$ provide enough spatial coverage and depth to estimate the distribution of external convergence via the weighted counts technique. We find that our different estimates are consistent with each other at a level of $\sim 0.005$, corresponding to $\sim 0.5\%$ impact on $H_0$. Our estimate which is least affected by photometric redshifts and stellar mass uncertainties, $P(\kext|\zeta_{q_\mathrm{gal}},\zeta_{q_{1/r}},\zeta_{\gext})$, has a median of 0.004, and a standard deviation of 0.025. This uncertainty contributes $\sim2.5\%$ rms error to the value of $H_0$. We intend to employ the techniques developed in this paper for the analysis of the other H0LiCOW lens systems. In particular, \hequad\ is a rather typical line of sight, and we expect that lenses residing in comparatively overdense fields will benefit more from the use of additional constraints including photometric redshifts and stellar masses. Throughout this work, we have made extensive use of the MS. The weighted count ratios technique is designed to minimize our reliance on a particular simulation, but it will be useful to repeat this analysis by using simulations for different cosmologies and galaxy models to test any remaining dependencies. However, we expect such dependencies to be small, given that the external convergence we measure is close to zero. Assuming a simple linear deterministic galaxy bias model, the convergence inferred from a given relative galaxy number overdensity scales roughly with the mean matter density parameter $\Omega_\mathrm{m}$ and the matter density fluctuation amplitude $\sigma_8$ (see Section \ref{section:cosm_depend}). Therefore, for example, $\kappa^\mathrm{med, Planck}_\mathrm{ext} \propto \kappa^\mathrm{med,MS}_\mathrm{ext}\Omega_\mathrm{m}^\mathrm{Planck}\sigma_8^\mathrm{Planck}/(\Omega_\mathrm{m}^\mathrm{MS}\sigma_8^\mathrm{MS}) \sim 1.13\kappa^\mathrm{med,MS}_\mathrm{ext}$. For $\kappa^\mathrm{med,MS}_\mathrm{ext}=0.004$, this corresponds to $\lesssim 0.001$ impact. We leave further checks for future work, as other simulations with convergence maps become available. Recently, \citet{mccully16} presented a technique of reconstructing the external convergence without relying on a particular simulation, through a direct modelling of the field. This has the potential of further reducing the uncertainty on the external convergence. This work has produced the galaxy catalogues necessary for a future implementation of that technique. While we have accounted in this work for the presence of voids, groups and clusters statistically, through the use of the MS, our catalogue products are also used in separate works (H0LiCOW Paper II and Tihhonova et al., in prep.) to directly identify such structures.

1607.01047

1607

1607.01792_arXiv.txt

We present results from a suite of three-dimensional global hydrodynamic simulations which show that spiral density waves propagating in circumstellar disks are unstable to the growth of a parametric instability that leads to break-down of the flow into turbulence. This spiral wave instability (SWI) arises from a resonant interaction between pairs of inertial waves, or inertial-gravity waves, and the background spiral wave. The development of the instability in the linear regime involves the growth of a broad spectrum of inertial modes, with growth rates on the order of the orbital time, and results in a nonlinear saturated state in which turbulent velocity perturbations are of a similar magnitude to those induced by the spiral wave. The turbulence induces angular momentum transport, and vertical mixing, at a rate that depends locally on the amplitude of the spiral wave (we obtain a stress parameter $\alpha \sim 5 \times 10^{-4}$ in our reference model). The instability is found to operate in a wide-range of disk models, including those with isothermal or adiabatic equations of state, and in viscous disks where the dimensionless kinematic viscosity $\nu\le 10^{-5}$. This robustness suggests that the instability will have applications to a broad range of astrophysical disk-related phenomena, including those in close binary systems, planets embedded in protoplanetary disks (including Jupiter in our own Solar System) and FU Orionis outburst models. Further work is required to determine the nature of the instability, and to evaluate its observational consequences, in physically more complete disk models than we have considered in this paper.

Spiral density waves are excited in circumstellar disks by a variety of processes, and play an important role in the dynamical evolution and observational appearance of these systems. They arise as normal modes in self-gravitating disks \citep{lin64}, where they provide an efficient mechanism for the transport of mass and angular momentum through gravitational torques and an advective wave flux \citep{lyndenbell72,papaloizou91,laughlin94}. Spiral shocks can be excited by stellar companions in close binary systems, including cataclysmic variables, X-ray binaries and T Tauri binaries \citep{lin79, sawada86, artymowicz94}. The excitation of nonlinear spiral waves by giant planets also plays an important role in structuring protoplanetary disks, and in driving the migration of embedded planets \citep{bryden99,kley99,nelson00}. Models for the origin of FU Orionis outbursts have been presented where spiral waves, originating in the gravitationally unstable outer regions of protostellar disks during the early infall phase, propagate into the inner disk regions and trigger the magnetorotational instability by heating and ionizing the disk gas there \citep{gammie99,armitage01,zhu10,bae14}. Recent work has also shown that the infall of low angular momentum material onto a protostellar disk during the early infall phase can also generate global spiral waves \citep{lesur15}. In this paper, we show that spiral waves in circumstellar disks are unstable to the growth of a parametric instability that leads to the disk flow becoming turbulent. We came across this phenomenon while extending the two-dimensional hydrodynamic simulations of triggered FU Orionis outbursts in \cite{bae14} to three dimensions. The instability arises because pairs of inertial waves, or inertial-gravity waves, couple resonantly to the spiral wave, leading to the extraction of energy and angular momentum from it and the growth of a broad spectrum of inertial waves. Similar parametric instabilities, leading to the growth of inertial waves and the generation of small scale turbulence, have been reported to arise in numerous circumstances where disk fluid elements are subjected to periodic forcing. \cite{goodman93} and \cite{ryu94} showed that a parametric instability arises due to the elliptical distortion of a disk caused by an external orbiting companion. Warped disks have been shown to be subject to parametric instability \citep{gammie00, ogilvie13}, as have globally eccentric disks \citep{papaloizou05a, papaloizou05b, barker14}. In a study similar to the one that we present here, \cite{fromang07} used shearing box simulations and analytical calculations to show that axisymmetric, nonlinear sound waves traveling in disks are also subject to parametric instability. In this paper, we examine the stability of propagating $m=2$ spiral waves in circumstellar disks, by carrying out three-dimensional global hydrodynamic simulations at high resolution. We consider a broad range of simplified disk models and input physics, including both isothermal and adiabatic equations of state, viscous and inviscid disks, and vertically stratified and non-stratified density profiles. Simulations are computed with four independent numerical codes, and the outcomes are found to be in good agreement with each other. Our main result is that the spiral density waves are unstable to the growth of the aforementioned parametric instability, leading to the development of turbulence in the disk models. We refer to the instability as the spiral wave instability (SWI). The SWI is found to be robust across the full range of models that we considered, suggesting that it will be influential in the dynamical evolution of disks that contain nonlinear spiral density waves of whatever origin. This paper is organized as follows. In Section \ref{sec:background}, we discuss the theoretical background to the SWI to set the scene for the numerical simulations. We describe our numerical methods in Section \ref{sec:method}, including the basic equations solved and the disk models examined. In Section \ref{sec:cyl}, we present results from cylindrical disk models as a demonstration of the instability in the simplest global model with a non-stratified, globally isothermal initial setup. Results from vertically stratified disk models with isothermal and adiabatic equations of state are discussed in Sections \ref{sec:iso} and \ref{sec:adia}, respectively. We discuss the application of the SWI to various astrophysical disk phenomena, with a particular emphasis on protoplanetary disks, in Section \ref{sec:discussion}, and we provide concluding remarks in Section \ref{sec:conclusion}. A reader who is mainly interested in the physical manifestation of the instability, and its potential applications, can safely skip Sections \ref{sec:background} -- \ref{sec:cyl} and read from Section \ref{sec:iso} onwards.

\label{sec:conclusion} We have presented the results of high resolution, 3-D simulations of circumstellar disks that are perturbed by two-armed spiral density waves, and we have shown that a broad range of disk models are subject to a parametric instability involving the excitation of pairs of inertial waves that interact resonantly with the spiral wave. This spiral wave instability (SWI) gives rise to turbulence that transports angular momentum and causes vertical mixing. The apparent robustness of the SWI under changes to physical conditions in the disks suggests that it may arise in a broad range of astrophysical settings. Future work is required to understand and evaluate its influence on the physical evolution and observational appearance of these systems.

1607.01792

1607

1607.02875_arXiv.txt

\noindent We investigate the correlations in galaxy shapes between optical and radio wavelengths using archival observations of the COSMOS field. Cross-correlation studies between different wavebands will become increasingly important for precision cosmology as future large surveys may be dominated by systematic rather than statistical errors. In the case of weak lensing, galaxy shapes must be measured to extraordinary accuracy (shear systematics of $< 0.01\%$) in order to achieve good constraints on dark energy parameters. By using shape information from overlapping surveys in optical and radio bands, robustness to systematics may be significantly improved without loss of constraining power. Here we use HST-ACS optical data, VLA radio data, and extensive simulations to investigate both our ability to make precision measurements of source shapes from realistic radio data, and to constrain the intrinsic astrophysical scatter between the shapes of galaxies as measured in the optical and radio wavebands. By producing a new image from the VLA-COSMOS L-band radio visibility data that is well suited to galaxy shape measurements, we are able to extract precise measurements of galaxy position angles. Comparing to corresponding measurements from the HST optical image, we set a lower limit on the intrinsic astrophysical scatter in position angles, between the optical and radio bands, of $\sigma_\alpha > 0.212\pi$ radians (or $38.2^{\circ}$) at a $95\%$ confidence level.

Weak gravitational lensing analyses exploit the coherent distortion of galaxy shapes induced by the gravitational potential along the line of sight to constrain the distribution and evolution of massive structures in the Universe, providing an excellent probe of cosmology \citep[see e.g.][for a review]{2015RPPh...78h6901K}. Conducting precision cosmology tests with weak lensing (e.g.~selecting between competing models of dark energy) requires high number densities ($n_{\rm gal} \gtrsim 1 \, \mathrm{arcmin}^{-2}$) of high redshift ($z \gtrsim 1$) sources and analysis tools to measure their shapes to extraordinary accuracy. These attributes have been available to surveys at optical wavelengths for a number of years, with useful cosmological results beginning to emerge from CFHTLens \citep{2016MNRAS.456.1508K}, DES-SV \citep{2015arXiv150705552T} and DLS \citep{2015arXiv151003962J}. As source number densities increase in these and future experiments (and statistical uncertainties decrease) it is the systematic uncertainties which will come to dominate. One way of tackling the problem of such systematics is through multi-wavelength investigations, which have the potential to address both instrumental and astrophysical contamination. In particular, telescopes operating at radio wavelengths are about to undergo a significant leap in survey speed and sensitivity and, ultimately, with the Square Kilometre Array (SKA)\footnote{\url{http://www.skatelescope.org}} will be capable of world-leading weak lensing cosmology alone \citep{2015aska.confE..23B, 2016arXiv160103948B}. In addition, cross-correlation weak lensing studies between radio and optical wavebands appear highly promising, providing comparable cosmological constraints as traditional single-waveband experiments, but with the significant advantage of being much more robust to wavelength-dependent systematic effects \citep{2016MNRAS.456.3100D, 2016arXiv160103947H}. In this paper we focus on a particular aspect of these cross-correlation studies: the correlation between the optical and radio shapes of those galaxies which will be common to both catalogues. As discussed in \cite{2016arXiv160103947H} this shape co-variance constitutes a noise term on the measured weak lensing shear power spectra, and needs to be well understood to form robust cosmological parameter constraints. Furthermore, literature results on these shape correlations are divided, with \cite{2009MNRAS.399.1888B} finding a significant correlation in orientations between the shapes of objects detected in the SDSS and FIRST catalogues, and \cite{2010MNRAS.401.2572P} finding almost no correlation in the optical and radio shapes of galaxies in the Hubble Deep Field North (HDF-N). Here we measure and compare the optical and radio shapes of sources in the COSMOS field, which has deep data available in both wavebands. To facilitate an accurate comparison, we have implemented a significant re-analysis of the radio VLA-COSMOS data set. Our re-analysis significantly reduces the level of galaxy shape systematics which we found to be present in the previous analysis of \cite{2007ApJS..172...46S} (though we note that their analysis was primarily focused on measuring faint number counts and was not optimised for extracting galaxy shapes). When conducting any study which relies upon accurate galaxy shape measurements, it is important to appreciate that the data generated in radio observations is fundamentally different to that generated by imaging telescopes in optical wavebands. Radio telescopes are typically operated as interferometers, consisting of many individual antennas, from which the signals are correlated to form `visibilities'. For a pair of antennas a projected distance $d$ apart, the recorded visibility is a measurement of the flux at a specific spatial frequency on the sky. For interferometers with $N$ antennas, $N(N-1)/2$ correlations can be formed between the different baseline lengths, sampling a large number of scales and allowing us to reconstruct a partially sampled Fourier transform of the sky brightness distribution \citep[when the assumptions of the Van Cittert–-Zernike theorem,][are satisfied]{2009MNRAS.395.1558C}. This can provide distinct advantages, allowing for very good angular resolution (set by the longest baseline in the array) with a comparably large maximum field of view (set by the primary beam of the individual antennas). The incomplete sampling of spatial scales means the data must be further processed to obtain an estimate of the corresponding image plane information. The resultant Point Spread Function (PSF) from a radio interferometer (commonly referred to as a `dirty beam') is highly deterministic at $\gtrsim 700 \, \mathrm{MHz}$ frequencies (as it is set by the known sampling of the Fourier plane -- a useful property for weak lensing) but can have significant structure, with appreciable sidelobes extending across the entire sky. The convolution of this complicated beam with an unknown sky creates the `dirty image' (the simple Fourier transform to image space of the data). Deconvolving the dirty beam from the dirty image to enable source identification, flux and shape measurement is then a difficult process. A commonly used algorithm for this purpose is known as H\"ogbom-\clean. This algorithm assumes that any section of the sky is made up of a finite number of delta function point sources. Beginning with the dirty image, \clean iteratively finds the position and brightness of the brightest source in the field and subtracts flux, corresponding to the convolution of this point source with the dirty beam, directly from the visibility data. This is repeated across the image until a residual map remains with \clean components, and produces a `\clean' map \citep[see][for a detailed description and motivation]{1974A&AS...15..417H}. Additionally this procedure helps to reduce sidelobe artifacts across the radio map which can be problematic, especially when a large dynamical range in source fluxes is considered. The \clean algorithm works well at producing plausible radio images and placing sources at reliable positions. However \clean is a non-linear deconvolution technique which assumes that missing information, due to a partially sampled visibility (or Fourier) plane, can be modelled as a set of point sources. This inherent assumption may result in significant systematics when attempting to measure precision morphology as is necessary for weak lensing cosmology. \cite{2015aska.confE..30P} have tested the performance of \clean on SKA-like simulations and have found the performance to be orders of magnitude away from what is necessary for systematic uncertainties from shape measurement to be sub-dominant to statistical ones when measuring cosmic shear. In this paper, we have adopted a \clean-based imaging pipeline in our re-analysis of the VLA-COSMOS data. However, mindful of the above concerns around potential shape biases being induced by the \clean processing, we have also carried out extensive simulations to quantify the level of shape measurement bias we expect in our newly-created image. In addition to minimising systematics through cross-correlation, radio weak lensing also presents further potential advantages through additional information from polarisation \citep{2011MNRAS.410.2057B, 2015MNRAS.451..383W} and rotational velocity maps \citep{2006ApJ...650L..21M} which may be used to reduce the shot noise from intrinsic galaxy shapes and also mitigate against intrinsic alignment systematics -- a key astrophysical systematic due to correlations of galaxy shapes being imprinted during their formation process. To date two attempts have been made to measure a weak lensing signal in radio data alone. In the first, \cite{2004ApJ...617..794C} were able to detect a lensing aperture mass signal at a significance of $3.6\sigma$. In this case the Fourier-plane data from the FIRST radio survey \citep{1995ApJ...450..559B} was used to estimate the shear on radio sources directly, without imaging, modelling them using Fourier-plane shapelets \citep{2002ApJ...570..447C}. The second study, by \cite{2010MNRAS.401.2572P}, used combined radio observations from the VLA and MERLIN interferometers to extract galaxy shapes from radio images constructed by the H\"ogbom-\textsc{clean} method discussed above, where the low absolute number of detected sources precluded a significant weak lensing detection. This paper is structured as follows. In \cref{The COSMOS Field} we introduce the COSMOS data used in this study and outline the data reduction steps that we have applied. In \cref{Shape Measurements}, we describe our technique for creating simulated radio datasets and we also describe our approach to measuring the radio galaxy shapes, which includes the extraction of a model of the VLA-COSMOS PSF from the simulated datasets. An assessment of the accuracy of the shape parameter recovery from the radio simulations is provided in \cref{simulation_err}. In \cref{vla_shape_measurements} we present the shape measurements from the real VLA-COSMOS data. We then investigate the correlation between the radio-derived and optical-derived position angles of galaxies within the COSMOS field in \cref{Multi-Waveband Shape Comparison}. Finally, in \cref{concl} we summarise our conclusions.

In this paper we have performed a detailed comparison analysis of the shapes of galaxies in the COSMOS field, as measured in the optical using HST, and as measured in the radio, using the VLA. Our study has been motivated by the scientific potential of cross-correlation cosmic shear analyses of future overlapping optical and radio surveys \citep{2015aska.confE..23B, 2016arXiv160103947H, 2016arXiv160103948B, 2016arXiv160603451C}. In order to fully exploit such future cross-correlations, one needs to understand in detail the correlations in intrinsic optical and radio shapes, which is the issue that we have attempted to address in this study. In the course of our analysis, we have highlighted some of the challenges involved in extracting shapes from radio observations. For this analysis we have chosen to measure galaxy shapes from images reconstructed from the VLA radio data. In particular, we have used simulations, composed of a known distribution of galaxy shapes combined with a realistic radio interferometer observation and imaging pipeline, to show the effectiveness of position angle ($\alpha$) recovery through typical radio data reduction and image creation techniques. We have investigated the use of two variations of the widely used H\"ogbom-\clean image creation algorithm in the radio: natural and uniform visibility weighting. We find the choice in weighting scheme affects the measured shape parameters greatly. In particular the choice of a uniform weighting scheme yields largely discrepant values of multiplicative bias in ellipticity components for a given galaxy. We found this problem was diminished when natural weighting was used, which produced highly similar multiplicative bias on ellipticities, which later cancels when position angles are derived from the measured ellipticities. This allowed us to obtain unbiased estimates of the position angles, which were important for this study. Quantifying the bias in position angle ($\alpha$) achieved in this study with a linear bias model, a link to weak lensing requirements for future surveys was established via the position angle only shear estimators of \cite{2014MNRAS.445.1836W}. A key result is shown in \cref{fig:radioarecovery} which allows one to translate from the usually quoted requirements on shear bias to requirements on position-angle bias. We find that although our position angle recovery appears relatively unbiased given our measurement errors, the uncertainty on the position angle bias that we are able to achieve using current data is still larger than current weak lensing requirements. After quantifying the expected uncertainties through simulations we have applied our shape measurement pipeline to the real COSMOS radio (VLA) observations. We used the resulting radio galaxy shape measurements (and associated uncertainties) to place a lower limit on the astrophysical scatter in source position angles between the continuum radio and optical emission of the sources which are detected in both our VLA data and an optical HST study of the same field. We find this lower limit to be $\sigma_{\alpha}^{\mathrm{scatter}} \gtrsim 0.212\pi$ (or $38.2^{\circ}$) at a $95\%$ confidence level. This appears consistent with results from previous studies \citep[][]{2010MNRAS.401.2572P} considering the low absolute number of sources in our sample. % Understanding the radio-optical shape correlations is important as it will affect the noise term on the cosmological power spectra measured by radio-optical cross-correlation weak lensing studies \citep{2016arXiv160103947H, 2016arXiv160603451C}. High levels of correlation will increase the shot noise on shear measurement, but allow for cross-waveband calibration against some systematics, whilst low levels of correlation effectively increase the source number density being used to probe the cosmic shear field. In the near future deep, high-resolution optical and radio surveys such as SuperCLASS\footnote{\url{http://www.e-merlin.ac.uk/legacy/projects/superclass.html}} will further constrain this correlation, better informing forecasts and survey design considerations for SKA.

1607.02875

1607

1607.07985_arXiv.txt

Models of nova outbursts suggest that an X-ray flash should occur just after hydrogen ignition. However, this X-ray flash has never been observationally confirmed. We present four theoretical light curves of the X-ray flash for two very massive white dwarfs (WDs) of 1.380 and 1.385 $M_\sun$ and for two recurrence periods of 0.5 and 1 years. The duration of the X-ray flash is shorter for a more massive WD and for a longer recurrence period. The shortest duration of 14 hours (0.6 days) among the four cases is obtained for the $1.385~M_\sun$ WD with one year recurrence period. In general, a nova explosion is relatively weak for a very short recurrence period, which results in a rather slow evolution toward the optical peak. This slow timescale and the predictability of very short recurrence period novae give us a chance to observe X-ray flashes of recurrent novae. In this context, we report the first attempt, using the \swift observatory, to detect an X-ray flash of the recurrent nova M31N 2008-12a (0.5 or 1 year recurrence period), which resulted in the non-detection of X-ray emission during the period of 8 days before the optical detection. We discuss the impact of these observations on nova outburst theory. The X-ray flash is one of the last frontiers of nova studies and its detection is essentially important to understand the pre-optical-maximum phase. We encourage further observations.

\label{sec_introduction} A nova is a thermonuclear runaway event that occurs on an accreting white dwarf (WD) \citep[e.g.,][]{ibe82,jos93,nar80,pri95,sta74}. Figure \ref{hr} shows a schematic HR diagram for one cycle of a nova outburst on a very massive WD. The thermonuclear runaway of hydrogen sets in on an accreting WD at point A. The luminosity increases toward point B at which the nuclear luminosity ($L_{\rm nuc}$) reaches its maximum. After that, the envelope on the WD greatly expands and reaches point D (the maximum expansion of the photosphere: corresponding to the optical peak). An optically thick wind begins to blow at point C and continues until point E through D. A part of the envelope mass is lost in the wind. From point C to E, the hydrogen-rich envelope mass decreases owing to wind mass loss and nuclear burning. After point E, it decreases owing to hydrogen burning. The hydrogen burning extinguishes at point F. The decay phase of optical and near-infrared (NIR) light curves corresponds to the phase from point D to E. The supersoft X-ray phase corresponds to the phase from point E to F. These phases have been well observed in a number of novae in various wavelength bands \citep[e.g.][and references therein]{hac06kb,hac10,hac14k,hac15k,hac16k,osb15,sch11}. The evolution of novae has been modeled by the optically thick wind theory \citep{kat94h}, and their theoretical light curves for D-E-F have successfully reproduced the observed light curves including NIR, optical, UV, and supersoft X-rays. From point D to E, the optical/IR light curves are well explained in terms of free-free emission \citep{gal76}, the fluxes of which are calculated from the mass-loss rate of the optically thick winds \citep{wri75}. From point E to F, the duration of the supersoft X-ray phase is theoretically reproduced. Detailed comparison with theory and observation enables us to determine/constrain the nova parameters such as the WD mass, distance, and extinction, in many novae \citep{hac14k,hac15k,hac16k,hac16II}. Thus, the characteristic properties of a nova from D to F have been well understood in both observational and theoretical terms. The X-ray flash is the stage from point B to C, which occurs just after the hydrogen ignition \citep{kat15sh, hac16sk}, but {\it before} the optical discovery. This stage has not been theoretically studied well, partly because of numerical difficulties and partly because of insufficient observational data to guide the theoretical models. In general, we cannot know in advance when and where a nova will erupt. Thus, soft X-ray flashes have never been detected in any kind of nova with any X-ray satellite. X-ray flashes represent one of the last frontiers of nova eruption studies and their detection will open a new landscape of nova physics. The X-ray flash of novae has been predicted from theoretical models for many years \citep[e.g.,][]{sta90,kra02}, but its observation had not been attempted until recently. In an attempt to provide observational constraints on X-ray flashes \citet{mor16} analyzed MAXI/GSC (Gas Slit Camera) data obtained with 92 minute cadence for 40 novae, including recurrent novae. They deduced the upper limit of the soft X-ray fluxes spanning a period of 10 days before the optical discovery of each nova. The energy bandpass of MAXI/GSC, however, is too high (2-4 keV) to detect the supersoft X-rays \citep[blackbody temperatures up to a maximum of 120~eV, observed in nova \novak, see][]{2014A&A...563A...2H,hen15} expected during the flash. Thus their upper limits of the bolometric luminosity are much higher than the theoretically expected values \citep[$\sim 10^{38}$ erg~s$^{-1}$, see][]{kat15sh} and their approach was not effective to restrict the epoch of an X-ray flash. We carried out a coordinated, very high-cadence observing campaign with the \swift satellite \citep{2004ApJ...611.1005G} to detect the X-ray flash during the 2015 outburst of the recurrent nova \nova \citep{dar14,dar15,hen14,hen15,tan14}. This is the \textit{ideal} object to detect X-ray flashes because its recurrence period is as short as a year, possibly even half a year \citep{hen15.half.period}. Such a very short recurrence period allows us to predict the eruption date with unprecedented accuracy ($\pm1$~month) and thereby makes any observational campaigns significantly more feasible than for any other novae. We found no significant X-ray emission during the eight days before the optical discovery by \citet{2015ATel.7964....1D}. This result is not consistent with the prediction made by \citet{kat15sh}, and suggests that theoretical models are still incomplete especially in the rising phase. Because no observational detection of soft X-rays and their properties has ever been obtained in the pre-optical-maximum phase, we are unable to constrain the theoretical models. In the present paper we describe the theoretical light curves of X-ray flashes for massive WDs, and present the observational results. We also address the implication of a non-detection of a flash. This paper is organized as follows. Section \ref{section_model} describes our improved numerical calculations and presents theoretical light curves of X-ray flashes as well as the physical properties of expanding envelopes in the early phase of shell flashes. Section \ref{section_observation} describes the {\it Swift} observations of the 2015 outburst of M31N 2008-12a, which resulted in the non-detection of an X-ray flash. In Section \ref{section_implication}, we identify the reason why X-ray flash emission was not detected. Discussion and conclusion follow in Sections \ref{section_discussion} and \ref{section_conclusion}. \begin{figure} \epsscale{1.10} \plotone{f1.eps} \caption{Schematic HR diagram for one cycle of a nova outburst on a $1.38~M_\sun$ WD. A mass-accreting WD stays at point A. When unstable hydrogen burning sets in, the star becomes bright (goes up). Point B denotes the epoch of maximum nuclear luminosity. Then the envelope expands and the photospheric temperature decreases with time (goes rightward). The optically thick wind starts at point C. The photospheric radius reaches maximum at point D. A part of the envelope matter is blown away in the wind. The optically thick wind continues until point E. Hydrogen nuclear burning extinguishes at point F. Finally the star cools down to point A. Three stages, X-ray flash (from B to C), wind phase (from C to E through D), and supersoft X-ray phase (from E to F) are indicated. } \label{hr} \end{figure}

Our main results are summarized as follows. \noindent 1. In a very early phase of a recurrent nova outburst, the photospheric luminosity rises very close to the Eddington luminosity at the photosphere and the temperature reaches as high as $T_{\rm ph}\sim 10^6$ K in WDs as massive as $1.38~M_\odot$. We expect bright supersoft X-ray luminosities in this X-ray flash phase, as large as $L_{\rm X} \sim 10^{38}$ erg~s$^{-1}$. \noindent 2. We present light curves of X-ray flashes for 1.35, 1.38, and $1.385 ~M_\odot$ WDs. The duration of the X-ray flash depends on the WD mass and the recurrence period, shorter for a more massive WD, and longer for a shorter recurrence period. The duration of the X-ray flash would be a good indicator of the WD mass and mass-accretion rate because it depends sensitively on these values. \noindent 3. The optically thick wind arises at the end of X-ray flash ($\log T_{\rm ph}$ (K) $\sim 5.6$) owing to acceleration by the Fe opacity peak. As no strong wind mass loss is expected during the X-ray flash, we could observe a naked photosphere, i.e., the spectrum is close to that of blackbody with $T_{\rm ph}$. \noindent 4. We observed with a six-hour-cadence the 2015 outburst of M31N 2008-12a with \swift from eight days before the optical discovery. Although our theoretical prediction of the X-ray flash duration was long enough, as long as 0.5 -- 1.5 days, no X-ray flash was detected. \noindent 5. We examined two possible reasons for the non-detection. Absorption by the surrounding matter originated from the companion is unlikely. Instead, we suggest that the X-ray flash could have occurred before our observations started, because short recurrence period novae undergo a very slow evolution. \noindent 6. The X-ray flash is one of the last frontiers of nova studies. We encourage further attempts at observational confirmation in the near future. Any detection of X-ray flashes would be essentially important to explore the pre-optical-maximum phase and to ultimately understand the complete picture of nova eruptions.

1607.07985

1607

1607.00936_arXiv.txt

We performed a detailed spectroscopic analysis of the fullerene C$_{60}$-containing planetary nebula (PN) \mbox{Lin49} in the Small Magellanic Cloud using XSHOOTER at the ESO VLT and the \emph{Spitzer}/IRS instruments. We derived nebular abundances for nine elements. We used {\sc TLUSTY} to derive photospheric parameters for the central star. \mbox{Lin49} is C-rich and metal-deficient PN ($Z$$\sim$0.0006). The nebular abundances are in good agreement with Asymptotic Giant Branch nucleosynthesis models for stars with initial mass 1.25\,M$_{\odot}$ and metallicity $Z$ = 0.001. Using the {\sc TLUSTY} synthetic spectrum of the central star to define the heating and ionising source, we constructed the photoionisation model with {\sc CLOUDY} that matches the observed spectral energy distribution (SED) and the line fluxes in the UV to far-IR wavelength ranges simultaneously. We could not fit the $\sim$1-5\,$\mu$m SED using a model with 0.005-0.1\,$\mu$m-sized graphite grains and a constant hydrogen density shell owing to the prominent near-IR excess, while at other wavelengths the model fits the observed values reasonably well. We argue that the near-IR excess might indicate either (1) the presence of very small particles in the form of small carbon clusters, small graphite sheets, or fullerene precursors, or (2) the presence of a high-density structure surrounding the central star. We found that SMC C$_{60}$ PNe show a near-IR excess component to lesser or greater degree. This suggests that these C$_{60}$ PNe might maintain a structure nearby their central star.

The discovery of C$_{60}$ in the C-rich planetary nebula (PN) Tc1 \citep{Cami:2010aa} confirmed the presence outside the Solar System of the enigmatic molecule buckminsterfullerene C$_{60}$, first discovered by \citet{Kroto_85_C60}. Since then, C$_{60}$ has been identified towards ten other PNe in the Milky Way \citep{Cami:2010aa,Garcia-Hernandez:2010aa,Garcia-Hernandez:2011aa, Garcia-Hernandez:2012aa,Otsuka:2013aa,Otsuka:2014aa}, bringing the total to 11 detections out of a sample of 338, both C-rich and O-rich, PNe observed with the Infrared Spectrograph \citep[IRS;][]{Houck:2004aa} on the \emph{Spitzer} Space Telescope. Assuming that the evolved star content of the Milky Way is 1/3 C-rich and 2/3 O-rich \citep{Ishihara:2011aa}, it can be inferred that fullerenes occur in about 10\,\% of the Galactic C-rich PNe, although this number may be lower if a larger fraction of Galactic PNe are C-rich. For C-rich PNe, \citet{Garcia-Hernandez:2015aa} reports a detection rate of $\sim$5\,\%, $\sim$20\,\%, and $\sim$44\,\% in the Milky Way, the Large Magellanic Cloud (LMC), and Small Magellanic Cloud (SMC), respectively. This indicates that the processing of fullerenes may depend on the metallicity, with fullerenes being more often detected in low-metallicity environments. In most cases, even the two strongest C$_{60}$ resonances at 17.4\,\micron\ and 18.9\,\micron\ are rather weak with respect to the local continuum emission around these wavelengths, with the notable exception of the PN \mbox{Lin49} (Fig.~\ref{image}) in the SMC, which appears to have C$_{60}$ 17.4\,\micron\ and 18.9\,\micron\ features of very similar strength and appearance to what is seen towards Tc1. The similarities in their infrared spectra and the similar C$_{60}$ band strengths motivated us to know more about physical properties of \mbox{Lin49}. However, little is known about \mbox{Lin49}. Prior to its \emph{Spitzer}/IRS observation, \mbox{Lin49} only occurs in some catalogues as an SMC PN \citep{Lindsay_61_new,Dopita_85_kinematics,Meyssonnier_93_New,Morgan_95_UKST} until recently. The source was selected for spectroscopic follow-up with \emph{Spitzer} based on its mid-infrared IRAC photometric colours, which suggested a pre-main sequence nature (G.~Sloan, private communication). The \emph{Spitzer}/IRS spectrum revealed that \mbox{Lin49} is a C-rich dust PN, showing strong C$_{60}$ resonances at 17.4\,$\mu$m and 18.9\,$\mu$m and similar dust features such as the broad 11\,$\mu$m and 30\,$\mu$m bands seen in the other C$_{60}$-containing LMC and SMC PNe \citep{Sloan:2014aa,Ruffle_15_Spitzer}, but the physical properties of the central star and dusty nebula remain unknown. Therefore, we wanted to further characterise \mbox{Lin49} using the XSHOOTER UV-near-IR spectrograph \citep{Vernet:2011aa} on the ESO Very Large Telescope (VLT) UT2 (Kueyen), in combination with the \emph{Spitzer}/IRS spectrum. In the case of \mbox{Lin49}, the well determined distance to the SMC allows us to accurately determine the luminosity of the central star, the size of the nebula, and the total gas and dust masses in the nebula, and then clarify the current evolutionary stage of the central star and estimate the initial mass. In this study, we present a spectroscopic analysis of \mbox{Lin49} in order to study the physical conditions and chemical properties of this interesting PN. This is part of an ongoing study to understand in more depth the physical and chemical properties of fullerene-containing PNe. Although we expect that these studies give us information on why fullerenes formed and exists in these PNe, the aim of this specific paper is not to investigate the formation and processing of fullerene molecules. The remainder of this paper is organised as follows. In Section~\ref{S:S2}, we describe our XSHOOTER observation and the data reduction of the XSHOOTER spectrum and the archived \emph{Spitzer}/IRS spectrum. The results of plasma-diagnostic and ionic and elemental abundance derivations using nebular lines, derivations of photospheric properties, and spectral energy distribution (SED) fitting are described in Section~\ref{S:S3}. In Section~\ref{S:S4}, we discuss the prominent near-IR excess found in \mbox{Lin49}, and we give interpretations of this feature. We discuss the SEDs of SMC C$_{60}$ PNe and non-C$_{60}$ C-rich PNe in the SMC by comparing with the SED of \mbox{Lin49}. We compare physical properties of the C$_{60}$-containing PNe and counterparts in the SMC. Finally, we summarise the works in Section~\ref{S:S5}. \begin{figure} \centering \includegraphics[width=0.8\columnwidth]{f01.eps} \caption{Image of \mbox{Lin49} in the $z^{'}$-band and the slit positions used in the XSHOOTER observations. We observed \mbox{Lin49} on the slit positions A and B. The averaged FWHM amongst nine nearby stars is $\sim$0.69{\arcsec}. } \label{image} \end{figure}

} \subsection{Interpretations for the near infrared excess \label{S:nir}} \begin{figure} \centering \includegraphics[width=\columnwidth]{f12.eps} \caption{Near-IR excess. The grey lines and the blue dots are the residual flux densities ($\Delta$\,$F_{\lambda}$) between the observed XSHOOTER/IRS spectra and 2MASS $JHKs$, IRAC 3.6/4.5/5.8\,$\mu$m photometry bands and the corresponding values obtained from the {\sc CLOUDY} model. In the XSHOOTER and IRS spectra, we block the spectral regions except for 1.046--1.271\,$\mu$m ($J$), 1.508--1.778\,$\mu$m ($H$), 1.974-2.377\,$\mu$m ($Ks$), and 5.31--19.74\,$\mu$m. The red line is the best fit of this near-IR excess with a Plank function with a single temperature of 1250\,K. The green line is another Plank function fit with the fixed 861\,K. The $\Delta$\,$F_{\lambda}$ in the photometry bands are listed in Table~\ref{T:nir-IR}. See the text in section~\ref{S:nir}. \label{F:nir} } \end{figure} \begin{table} \centering \caption{The residual flux densities ($\Delta$\,$F_{\lambda}$) between the observed 2MASS $JHKs$ and \emph{Spitzer}/IRAC 3.6/4.5/5.8\,$\mu$m bands and the corresponding values obtained from the {\sc CLOUDY} model. \label{T:nir-IR}} \begin{tabu} to \columnwidth {@{}LCC@{}} \hline Band&$\lambda_{\rm c}$ &$\Delta$\,$F_{\lambda}$\\ &($\mu$m) &(erg s$^{-1}$ cm$^{-2}$ $\mu$m$^{-1}$)\\ \hline 2MASS $J$ &1.235 &2.83(--13) \\ 2MASS $H$ &1.662 &4.32(--13) \\ 2MASS $Ks$ &2.159 &5.87(--13) \\ IRAC-Band1 &3.600 &3.96(--13) \\ IRAC-Band2 &4.500 &3.06(--13) \\ IRAC-Band3 &5.800 &2.51(--13) \\ \hline \end{tabu} \end{table} \subsubsection{Stochastic heating of extremely small particles} For the usual dust grain sizes (e.g., $\gtrsim$ 0.01\,$\mu$m), dust temperatures are determined by solving an energy balance equation between the radiative heating owing to the central star and the cooling of grains. For such grain sizes, individual quantum events are not important. However, for very small grains, which are composed of 100 atoms or less, single photons would cause them to heat up significantly for very short time scales. This mechanism is known as stochastic heating (or quantum heating), and has been proposed to explain the spectra of the reflection nebulae NGC7023 and NGC2023 \citep{Sellgren:1984aa}, and the PNe IC418 \citep{Phillips:1984aa} and Abell 58 \citep{Koller:2001aa}. Interestingly, these reflection nebulae and IC418 show mid-IR C$_{60}$ band emission. We included the stochastic heating mechanism in {\sc CLOUDY} model, as is default for {\sc CLOUDY}. However, our model cannot fit the observed SED in the $\sim$1-5\,$\mu$m wavelength range at all. For the residual data plots in Fig.~\ref{F:nir}, it is possible to fit the excess with a Planck function with a single temperature of 1250 $\pm$ 42\, K (indicated by the red line). The luminosity and the minimum emitting radius of this component are 290 $\pm$ 80\,L$_{\odot}$ and $\sim$2\,AU. According to \citet{Sellgren:1984aa} and \citet{Whittet:2003aa}, the thermal properties of solids are described by Debye's theory and the heat capacity $C_{\rm V}$ for $T$ over Debye temperature $\Theta$ (in graphite, $\Theta$ is $\sim$500\,K) is 3$N$$k$ where $N$ is the number of the atoms in a molecule and 3$N$ is its number of degrees of freedom. For extremely small grains, an average absorbed photon energy $E_{\rm ph}$ produces the difference between the maximum and minimum temperatures $\Delta~{T}$ written by the following equation: \begin{equation} \Delta~{T} = \frac{E_{\rm ph}}{3Nk}. \end{equation} The $T_{\rm eff}$ of the central star is 30\,500 K, so the photon energy at the radiation peak is 13.05 eV (Fig.~\ref{F:CSPNsed}). $E_{\rm ph}$ could be lower than 13.05\,eV. Thus, we obtained $\Delta~{T}$ $<$ 5.05(+4)/$N$. By adopting the maximum and minimum temperatures of 1250\,K and 20\,K, we obtain a value for $N$ of $\lesssim$~39, although our estimation is very optimistic and also depends on the minimum temperature. If such a molecule formed as a honeycomb structure sheet is distributed in the nebula, the molecule's dimension is roughly $\sim$8(--4)\,$\mu$m $\times$ 8(--4)\,$\mu$m square. If the molecule is a cage not a sheet, .e.g, fullerene C$_{36}$, the size would be small; in the case of C$_{36}$, the approximate diameter is 5(--5)\,$\mu$m \citep{Piskoti:1998aa}. The SED with $\Delta~T$ of 841\,K derived by keeping $N$ of 60 (indicated by the green line) gives a better fit to the differential spectrum in $\gtrsim$~3.6\,$\mu$m. A top-down mechanism has been suggested for the formation of C$_{60}$; C$_{60}$ could be formed from the shrinkage of larger molecules, e.g., from larger clusters of PAHs \citep{Zhen_14_Laboratory,Berne_15_Top}, or HAC \citep{Duley_12_FULLERENES}. PAH clusters could form from HACs. \citet{Scott:1997ab} showed that C$_{50,60,70}$ may be produced by the decomposition of HACs. As explained in Section~\ref{S:Spitzer}, the 6-9\,$\mu$m band profile in Lin49 is very similar to the thermal emission profile of HAC as presented in Fig.~2 of \citet{Scott:1997aa}. Molecules composing of $\lesssim$~39 C-atoms might be a by-product in the decomposition process of HACs. As a pragmatic problem, with grains composing of $\lesssim$~39 C-atoms only, it could be difficult to reproduce the observed broad continuous near-IR excess feature seen in \mbox{Lin49}; to get a continuum like behaviour, enough interacting vibrational modes are necessary. Therefore, we need to examine other possible explanations for near-IR excess in \mbox{Lin49}. \subsubsection{High density structure nearby the CSPN} \label{S:highdens} \begin{figure} \includegraphics[width=\columnwidth]{f13.eps} \caption{Radial profiles of the {\te} ({\it panel a}), $n$(H) ({\it panel b}), and {\Ne} ({\it panel c}) predicted by the two density shell model. The averaged value of each physical parameter is indicated in each panel. } \label{F:physical} \end{figure} \begin{figure} \includegraphics[width=\columnwidth]{f14.eps} \caption{({\it panel a}) Comparison between the observed SED plots and the predicted SED by the two density shell model. The spectral resolution of the gas emission lines is a constant 1000 in $<$~0.3\,$\mu$m, 9200 in 0.3-1.0\,$\mu$m, 4800 in 1.0-2.5\,$\mu$m, and 100 in $>$~2.5\,$\mu$m. ({\it panel b}) Close-up plot for mid-IR wavelengths. The lines and symbols in both of panels are as defined in Fig.~\ref{F:sed}.} \label{F:sed2} \end{figure} \begin{figure} \includegraphics[width=\columnwidth]{f15.eps} \caption{Radial temperature profiles of graphite grains in the smallest and largest size bins predicted by the two density shell model. } \label{F:cloudy_td2} \end{figure} \begin{table} \centering \caption{Input parameters of the best fitting in the two density shell model and the derived properties. \label{modelT}} \begin{tabu} to \columnwidth {@{}l@{\hspace{15pt}}l@{}} \hline {Parameters of the Central star} &\multicolumn{1}{c}{Values}\\ \hline $L_{\ast}$ &6333~L$_{\odot}$\\ $R_{\ast}$ &2.84~R$_{\odot}$\\ $M_{\ast}$ &0.57~M$_{\odot}$\\ $T_{\rm eff}$ &30\,500\,K \\ $M_{V}$ &--1.70\\ $m_{V}$ &17.26\\ $\log\,g$ &3.29~cm s$^{-2}$\\ Distance &61.9~kpc\\ \hline {Parameters of the Nebula} &\multicolumn{1}{c}{Values}\\ \hline Boundary condition &Ionisation bound\\ $\epsilon$(X) &He:10.80/C:8.46/N:7.06/O:8.03,\\ &Ne:7.27/S:5.88/Cl:4.08/Ar:5.22, \\ &Fe:4.71/the others:\citet{Fishlock:2014aa}\\ Geometry &Spherical\\ Shell size &$R_{\rm in}$ = 0.00015 pc (31 AU)\\ &$R_{\rm out}$ = 0.068 pc (1.42(+4) AU)\\ $n_{\rm H}$ &See Fig.~\ref{F:physical}\\ Filling factor &0.50\\ $\log_{10} I$({\hb}) &--12.89~erg s$^{-1}$ cm$^{-2}$\\ $m_{\rm g}$ &0.11\,M$_{\odot}$ \\ \hline {Parameters of the Dust} &\multicolumn{1}{c}{Values}\\ \hline Grain &graphite only\\ Grain radius &0.005-0.10\,$\mu$m\\ $T_{d}$ &See Fig.~\ref{F:cloudy_td2}\\ $m_{d}$ &4.15(--5)\,M$_{\odot}$ \\ $m_{d}$/$m_{\rm g}$ &3.86(--4) \\ \hline \end{tabu} \end{table} In section~\ref{photomodel}, we assumed that \mbox{Lin49} does not have any sub-structures surrounding the CSPN but that this PN has a normal density nebula. However, the minimum emitting radius of the 1250\,K blackbody component suggests that the near-IR component could be emitted by a sub-structure near the central star. A similar idea was proposed for the near-IR excess at the central position of IC418; \citet{Hora:1993aa} took the near-IR $JHK$ images of this PN and found excess at the central position after subtracting the contribution from the central star. The authors argued that the excess indicates a possible compact shell interior to the main shell. \mbox{Lin49} may also have a central dense structure, which will be responsible for its near-IR excess. To test this hypothesis, we construct a two-shell model. The model is composed of an outer low density shell and an inner high density shell. For the dust distribution on the high density shell we assume the dust is composed of graphite grains with an $a^{-3.5}$ size distribution, where $a$ = 0.005-0.1\,$\mu$m. $R_{\rm out}$ for this shell corresponds to the $R_{\rm in}$ for the outer, low density shell. The photoionisation model for the low density shell has already been constructed in section~\ref{photomodel}, with the residual SED indicating the near-IR excess as presented in Fig.~\ref{F:nir}. We fit the residual SED in the wavelength range from 1 to 5\,$\mu$m using {\sc CLOUDY}. In this process, we keep the following parameters derived in the low density shell model: the filling factor, the elemental abundances and the dust composition/size distribution/abundance, and $L_{\ast}$. As the luminosity of the near-IR excess component is small compared to that of the central star ($<$~5\,$\%$ of $L_{\ast}$), we will minimise the number of free parameters by initially fixing $L_{\ast}$ (later this value will be fine-tuned). The free parameters are $R_{\rm in}$ and $n$(H), which are determined through fitting the residual 1-5\,$\mu$m SED. We combine the radial $n$(H) profiles of the low and high density shells into one (See Fig. \ref{F:physical}), and run the model with this $n$(H) profile to match the observed SED plots in UV to far-IR wavelength. For fine-tuning, we allow an increase of $L_{\ast}$ by 10\,$\%$ and a slight increase of the elemental abundances, except for C. Finally, we obtained the predicted SED as presented in Fig.~\ref{F:sed2}, which better fits the 1-5\,$\mu$m wavelength rage, compared to Fig.~\ref{F:sed}. The $\chi^{2}$ was 37 calculated from the 51 gas emission fluxes, 9 broad band fluxes, 3 far-IR flux densities at 65, 90, 120\,$\mu$m, and the $I$({\hb}). In the fourth and ninth columns of appendix Table~\ref{model2}, we compare the observed and model predicted values and list the predicted fluxes. The input parameters for the best-fitting model and the derived parameters are summarised in Table~\ref{modelT}. Since we increased $L_{\ast}$ to obtain the best fit, the core-mass of the CSPN is 0.57\,M$_{\odot}$, accordingly. As we argued before, even in this core-mass, our conclusion on the initial mass of the progenitor star does not change. The predicted elemental abundances in the two density shell model are almost consistent with the single shell model in section~\ref{photomodel}. The 1-$\sigma$ of the elemental abundances and the $m_{g}$, $m_{d}$, and $m_{d}$/$m_{g}$ is almost same value discussed in section~\ref{photomodel}. The radial profiles of {\te}, $n$(H), and {\Ne} are presented in Fig.~\ref{F:physical}. The value for {\Ne} of the high density shell is very high. However, the emitting volume of the high density shell is very small. Therefore, the volume-averaged {\Ne} (5150\,cm$^{-3}$) is not over the observed one\footnote{XSHOOTER could not resolve Lin49 and this instrument looks at the average light of this PN in each wavelength. The Paschen line analysis could suggest the presence of such high density shell indirectly.}. The radial $t_{d}$ profile is presented in Fig.~\ref{F:cloudy_td2}. The maximum temperature of the $\langle{a}\rangle$ = 0.0054\,$\mu$m grains (1450\,K) meets the requirement to fit the residual 1-5\,$\mu$m SED plots but is not over the evaporation temperature of graphite. Thus, we succeeded in explaining the near-IR excess by postulating the existence of a high density structure nearby the CSPN. If we can believe that the near-IR excess emits from this high-density structure, how did the progenitor form this during its evolution? \subsubsection{Does Lin49 have a disc?} We suggest that the near-IR emission does not originate from the nebular shell, but from a disc around the central star. According to the review of \citet{van-Winckel:2003aa}, the hot dust component ($\sim$1000\,K) in some post-AGB stars \citep[first noticed by][]{Trams:1991aa} was interpreted as evidence for significant post-AGB mass loss. However, this interpretation became untenable because dusty post-AGB mass loss would speed up the evolution such that very few objects would be observable \citep{Trams:1989aa}. At present, the most accepted formation mechanism to produce a disc around an evolved star is the binary model \citep[e.g.,][and references therein]{Kwok:2000aa}. This model states that PNe with a disc around the central star evolve from a binary system that went through a common envelope phase. During this phase, the secondary star induces the mass loss of the AGB star to occur preferentially in the orbital plane, which gives birth to a disc. Thus, if the presence of a disc is confirmed around the central star of \mbox{Lin49}, this is a strong indication that this PN evolved from a binary system. The binary disc could stably harbour the near-IR emitters near the central star for a long time. \citet{Kamath:2014aa} classified 63 SMC objects into post-AGB/red giant branch (RGB) candidates based on their SEDs in the optical to mid-IR, and they reported that 21 objects out of these are post-AGB stars and 27 show a strong near-IR excess interpreted as the presence of a circumbinary disc. \citet{Maas:2005aa} investigated elemental abundances of the Galactic 12 post-AGB stars showing near-IR excess (they interpreted as the presence of a disc); nine of them are affected by the depletion process, that is, elements with a high condensation temperature (e.g., Fe) is largely depleted and get locked in dust grains whereas elements with a low temperature remain in gas phase. Since the temperature in the disc decreases with increasing radius, at the inner radius of the disc only the elements with a high condensation temperature are caught in grains, given that the inner temperature of the disc is high. \citet{Maas:2005aa} indicated that the presence of of the depletion process and the presence of a disc are linked. From the observational results of \citet{Kamath:2014aa} \citet{Maas:2005aa}, the strong near-IR excess and the strongly depleted {\it nebular} Fe abundance in \mbox{Lin49} might be explained by the presence of a disc. Most of the Fe-atoms might be tied up in dust grains (e.g., FeO) within a disc. If \mbox{Lin49} has a disc, the large carbon molecules such as fullerene were relatively easily formed. \subsubsection{Near-IR excess in SMC C$_{60}$ PNe and counterparts} \begin{figure} \includegraphics[width=1.0\columnwidth]{f16.eps} \caption{SED plots of SMC C$_{60}$ PNe SMC13, 15, 16, 18, 24 and Lin49. The blue filled circles are the de-reddened photometric data. In each panel, we compare with the resultant SED of Lin49 synthesised in section~\ref{photomodel}. Note that in each panel Lin49's SED is scaled to the observed de-reddened flux density of each PN at the $I_{c}$ band.} \label{F:sedSMC} \end{figure} \begin{figure} \includegraphics[width=1.0\columnwidth]{f17.eps} \caption{SED plots of non-C$_{60}$ C-rich SMC PNe SMC6 and 27. The lines and symbols in both of panels are as defined in Fig.~\ref{F:sedSMC}.} \label{F:sedSMC2} \end{figure} Is the near-IR excess seen in other SMC C$_{60}$ PNe? Amongst SMC C$_{60}$ PNe, SMP24 also shows a near-IR excess. By applying an SED fit over the range from $B$-band to $\sim$1.337 GHz, \citet{Bojicic:2010aa} find that a hotter dust component ($\sim$1000 K) is necessary to fit the observed SED down to 1\,$\mu$m, apart from the hot dust component ($\sim$270\,K). The photometric data from MCPS, 2MASS, and \emph{Spitzer}/IRAC and MIPS, and the \emph{Spitzer}/IRS spectra of the SMC C$_{60}$ PNe are plotted in Fig.~\ref{F:sedSMC}. As no optical data are available for SMC1, this source was not included in the figure. According to Table~\ref{T:compSMCXS}, the effective temperatures of the SMC C$_{60}$ PNe are in the range between 30\,500\,K and 58\,000\,K and the average $T_{\rm eff}$ amongst these PNe except for SMC15 (58\,000\,K) is 34\,100\,K. The observed UV-optical wavelength SEDs in these PNe are not largely different from each other, except for SMC15. Thus, we plot the resultant synthesised SED of Lin49 (section \ref{photomodel}) as the comparison; in each panel this Lin49's SED is scaled to the observed de-redden flux density of each PN at the $I_{c}$ band. Amongst C$_{60}$ PNe, the SEDs of SMC16 and SMC24 are very similar to that of Lin49; their SEDs have a flux density peak around 2MASS $Ks$-band and the near-IR excess features are apparently as broad as that seen in \mbox{Lin49}. SMC13, 15, and 18 do not show such a broad near-IR excess as seen in \mbox{Lin49}. However, their SED slope in the range from $U$ to $Ic$ band wavelength is different from those in the range from 2MASS to IRAC band wavelength. In Fig.~\ref{F:sedSMC2}, we show the SEDs of non-C$_{60}$ C-rich PNe, SMC6 and SMC27. From appendix Table~\ref{T:compSMCXS2}, they are selected as comparison objects to the C$_{60}$ SMC PNe because the $T_{\rm eff}$s of SMC6 and SMC27 are close to those of SMC C$_{60}$ PNe. We excluded SMC11 as a comparison because \citet{Villaver:2004aa} reported $T_{\rm eff}$ of 40\,900\,K whereas we confirmed that its SED in the range from MCPS $B$ (no $m_{U}$) to IRAC bands has a peak around $Ic$-band and it can be well expressed by a single Planck function with the temperature of 3722\,K, indicating that this object could not be a PN. Although it is hard to draw a strong conclusion based on the photometric data points and mid-IR spectra only, SMC C$_{60}$ PNe seem to show a near-IR excess component to lesser or greater degree. This suggests that these C$_{60}$ PNe might maintain a structure near their central star. Meanwhile, we do not find a flux density peak around 2MASS $Ks$ in SMC non-C$_{60}$ PNe SMC6 and 27. To reach a firm conclusion, we need to obtain near-IR spectra in order to check whether or not each PN displays a near-IR excess. \subsection{Comparison of physical properties between C$_{60}$ PNe and non-C$_{60}$ C-rich PNe} \begin{table} \centering \caption{The average elemental abundances of C$_{60}$-containing PNe and non-C$_{60}$ C-rich PN in the SMC and the Milky Way (MW). The data of the SMC PNe are taken from Table~\ref{T:compSMCXS} and appendix Table~\ref{T:compSMCXS2}. The data of the MW PNe are from Tables~3 and 4 of \citet{Otsuka:2014aa}. } \begin{tabu} to \columnwidth {@{}lcccccc @{}} \hline PNe &$\epsilon$(C)&$\epsilon$(N)&$\epsilon$(O)&$\epsilon$(Ne)&$\epsilon$(S)&$\epsilon$(Ar)\\ \hline SMC C$_{60}$ PNe&8.28 &7.17&7.97&7.07&6.54&5.64\\ Lin49 &8.46 &6.93&8.11&7.18&6.02&5.48\\ SMC C-rich PNe&8.67 &7.37&8.08&7.28&6.69&5.64\\ \hline MW C$_{60}$ PNe &8.73 &7.81 &8.48 &7.85 &6.45&5.98\\ MW C-rich PNe &8.85 &8.23 &8.56 &8.06 &8.05&6.85 \\ \hline \end{tabu} \label{T:abundS} \end{table} In Table~\ref{T:abundS}, we summarise the average elemental abundances of SMC and Milky Way (MW) C-rich PNe. In section~\ref{S:compAGB}, from the view of elemental abundances, we concluded that \mbox{Lin49} and the other SMC C$_{60}$ PNe evolved from the initially 1.0-1.25\,M$_{\odot}$ stars by comparing with the AGB nucleosynthesis results of \citet{Fishlock:2014aa}. The lowest S and Ar abundances of \mbox{Lin49} amongst SMC C$_{60}$ PNe indicate that this PN is an older population in the SMC. While there are no large differences in S and Ar (because these are Type II SN products), we found that the C, N, and Ne in SMC non-C$_{60}$ PNe are greater than those in SMC C$_{60}$ PNe, indicating that non-C$_{60}$ PNe evolved from more massive stars. The C and Ne are synthesised in the He-rich intershell during the thermal pulse AGB phase and these elements together with N and $n$-capture elements are brought up to the stellar surface by the third dredge-up (TDU). The efficiency of TDU depends on the initial mass and composition and increases as larger initial mass \citep[e.g.,][]{Karakas:2010aa}. The rich Ne in non-C$_{60}$ PNe would be due to the double $\alpha$ capturing by rich $^{14}$N. Certainly, the average abundances in SMC non-C$_{60}$ PNe are close to the AGB model result for the 1.25\,M$_{\odot}$ and 1.50\,M$_{\odot}$ stars rather for the 1.0\,M$_{\odot}$ and 1.25\,M$_{\odot}$ stars. \citet{Otsuka:2014aa} argued that elemental abundances of MW C$_{60}$ PNe can be explained by AGB models for 1.5-2.5\,M$_{\odot}$ stars with the SMC metallicity (i.e., $Z$ = 0.004). Although there is a sample selection bias, the C, N, and Ne abundances in the non-C$_{60}$ MW C-rich PNe indicate that these PNe evolved from more massive stars. Thus, at this moment, we might conclude that the progenitors of C$_{60}$ PNe in both the SMC and the MW are not greater than those of non-C$_{60}$ PNe. \citet{Otsuka:2013aa,Otsuka:2014aa} reported that the MW C$_{60}$ PNe have cool central stars. The average $T_{\rm eff}$ amongst MW C$_{60}$ PNe calculated from Table~3 of \citet{Otsuka:2014aa} is 37\,780\,K, which is in excellent agreement with the average $T_{\rm eff}$ = 38\,770\,K amongst the SMC C$_{60}$ PNe (See Table~\ref{T:compSMCXS}). The $T_{\rm eff}$ of \mbox{Lin49} is the coolest amongst SMC C$_{60}$ PNe. Taking their average ionised nebula radius $r$ of 0.17$''$ (0.05\,pc in linear scale), SMC C$_{60}$ PNe are slower evolving objects than SMC non-C$_{60}$ PNe, where the average $T_{\rm eff}$ is 72\,100\,K and the average $r$ is 0.34$''$ (0.10\,pc in linear scale).

1607.00936

1607

1607.07511_arXiv.txt

Data from recent laser-shock experiments, density-functional theory (DFT) with molecular-dynamics (MD), and path-integral Monte Carlo (PIMC) simulations on carbon are compared with predictions from the neutral-pseudo-atom (NPA)+ hyper-netted-chain (HNC) approach for carbon, a complex liquid in the warm-dense matter regime. The NPA results are in good agreement, not only with high-density regimes that have been studies via PIMC, but even at low densities and low temperatures where transient covalent bonding dominates ionic correlations. Thus the `pre-peak' due to the C-C bond at $\sim$1.4-1.6 \AA$\,$ and other features found in the pair-distribution function from DFT+MD simulations at 0.86 eV and 3.7 g/cm$^3$ etc., are recovered accurately in the NPA+HNC calculations. Such C-C bonding peaks have not been captured via average-atom ion-sphere (IS) models. Evidence for an unusual liquid $\to$ vapor and metal$\to$ semi-metal transition occurring simultaneously is presented. Here a strongly correlated metallic-liquid with transient C-C bonds, i.e., carbon at density $\sim$ 1.0 g/cm$^3$ and mean ionization $Z=4$ transits abruptly to a disordered mono-atomic vapour at 7 eV, with $Z\simeq$ 3. Other cases where $Z$ drops abruptly are also noted. The nature of $Z$, its discontinuities, and the role of exchange-correlation, are reviewed. The limitations of IS models in capturing the physics of transient covalent bonding in warm dense matter are discussed.

\label{intro} A number of light elements like hydrogen, boron, carbon, nitrogen, silicon, phosphorous, etc., and their mixtures form strong covalent bonds in the solid and retain much of this bonding, even when molten~\cite{galli89,glosli99, ghiring05} and in warm-dense-matter (WDM) regimes~\cite{whitley15, kraus13, hammelCH12, DWP-Carb90}. A theory for predicting the properties of these elements and their mixtures reliably and rapidly is needed in many technological applications. Their properties under higher compressions and temperatures are also of great interest in astrophysics~\cite{hubNellis91,sherman12,driver12}. Carbon is an important component of inter-stellar matter, white dwarfs, solar and extra-solar planets. WDM carbon is basic to many technologies and inertial-confinement fusion (ICF) ablators~\cite{lindl04,benedict2014}. Many applications require the equation of state (EOS) and transport properties in experimentally inaccessible regimes. The experiments when feasible are quite demanding ~\cite{kraus13,savvati08}. The theoretical prediction requires the electronic and ionic structure factors, and their interaction potentials. At low temperatures (compared to the Fermi energy $E_F$), numerically expensive density-functional theory (DFT) calculations coupled to molecular-dynamics (MD) provide a reliable method if sufficiently large calculations can be made using suitable exchange-correlation (XC) functionals. In DFT the WDM is modeled by a sequence of quenched $N$-atom periodic crystals thermally evolved using MD. The spurious electronic band structure of each crystal is averaged over the many ionic configurations obtained from MD. Thus DFT+MD is impractical at higher-temperatures ($T$) due to the many electronic states needed. At high $T$ and high density ($\bar{\rho}$) path-integral Monte Carlo (PIMC) methods are available for small-$N$ (e.g, 24 atoms) simulations~\cite{driver12}. Since the number of ions $N$ in DFT simulations is limited, other methods are used, especially for ambient-temperature applications. Semi-empirical methods developed from the ``embedded-atom'' approach have led to ``reactive-potentials'' which include bonding, dependence on coordination, torsion and bond angle. Carbon is a complex liquid where conjugated bonds ($sp3, sp2, sp$) and weak graphite-like partial conjugation occur deploying four valance electrons, i.e. with mean ionization $Z$=4. The Brenner potentials incorporate conjugation and were used by Glosli {\it et al.,}~\cite{glosli99} in their study of the phase transition in liquid carbon. The modern bond-order potentials (LRBOP) by Los {\it et al.}~\cite{losFaso03,losGhir05} include long-range effects and are fitted to a very wide data base. However, after applying LRBOP to calculate $S(k)$ of liquid C at 0.52 eV, Kraus {\it et al.}~\cite{kraus13} stated that ``this potential appears to be too stiff. Compared to DFT+MD, we find a higher pressure and a more pronounced structure which is not found in the experiment''. Thus LRBOP is hardly promising for $T\sim E_F$ typical of WDM since $T$=0.52 eV, $T/E_F=0.017$, defines a very low-$T$ WDM where it is already inadequate. Here we use a first-principles DFT approach treating electrons with a one-body density $n(r)$, and ions with a one-body density $\rho(r)$ to present results using the neutral-pseudo-atom (NPA) model. This couples one-body Kohn-Sham calculations and its electron-ion potentials to pair-distribution functions (PDFs) via integral equations like the hyper-netted-chain (HNC) equation, or MD simulations. NPA+HNC closely recovers PIMC results, and DFT+MD results including C-C covalent bonding signatures in the ion-ion PDFs at low -$T$, but extends seamlessly to high $T$ regimes inaccessible by DFT+MD. In Sec.~\ref{monomer-dimer.sec} we present evidence for a phase transition of highly-correlated liquid carbon containing transient C-C bonds into a weakly correlated mono-atomic carbon gas. The pressure, compressibility and conductivity in the neighborhood of the transition are presented. The nature of the mean ionization $Z$, its behaviour near the transition, and comparisons of the NPA+HNC method with PIMC and DFT+MD are given. An appendix discusses details of the electronic structure near the phase transition, and the nature and constraints on the mean ionization $Z$.

1607.07511

1607

1607.05500_arXiv.txt

% {} {We perform MHD modeling of a single bright coronal loop to include the interaction with a non-uniform magnetic field. The field is stressed by random footpoint rotation in the central region and its energy is dissipated into heating by growing currents through anomalous magnetic diffusivity that switches on in the corona above a current density threshold. } {We model an entire single magnetic flux tube, in the solar atmosphere extending from the high-$\beta$ chromosphere to the low-$\beta$ corona through the steep transition region. The magnetic field expands from the chromosphere to the corona. The maximum resolution is $\sim 30$ km. } {We obtain an overall evolution typical of loop models and realistic loop emission in the EUV and X-ray bands. The plasma confined in the flux tube is heated to active region temperatures ($\sim 3$ MK) after $\sim 2/3$ hr. Upflows from the chromosphere up to $\sim 100$ km/s fill the core of the flux tube to densities above $10^9$ cm$^{-3}$. More heating is released in the low corona than the high corona and is finely structured both in space and time. }

\label{sec:intro} Coronal loops are magnetic flux tubes where million degree plasma is confined and are the building blocks of the magnetically closed part of the solar corona. Understanding them means understanding how the corona is structured and powered \citep[see][for a review]{Reale2014a}. Each coronal loop is known to evolve fairly independently of nearby ones, because the major mass and energy transport processes occur only along the magnetic field lines. This remains true when a coronal loop is modelled as a bundle of thinner fibrils. Although the fibrils show overall a collective behaviour, each of them is thermally isolated from the others. On this basis, coronal loops have been largely investigated as single isolated systems by means of one-dimensional models, where the main role of the magnetic field is to guide the mass and energy transport. This approach has been successful in describing the basic physical processes and many features observed in loops \citep[e.g.,][]{Priest1978a,Rosner1978a,Hood1979a,Nagai1980a,Peres1982a,Doschek1982a,Nagai1984a,Fisher1985c,MacNeice1986a,Hansteen1993a,Priest1998a,Antiochos1999a,Reale2000a,Muller2003a,Bradshaw2003a,Cargill2004a,Bradshaw2006a,Reale2008a,Guarrasi2010a,Reale2012b}. Triggering a loop brightening inside these one-dimensional single-loop models is by an energy input inside a tenuous and cool initial coronal atmosphere. The heating makes the temperature increase rapidly all over the loop, because of the efficient thermal conduction, and drives a strong overpressure rapidly down to the dense chromosphere. The chromosphere expands upwards and fills the coronal part of the loop with hot denser plasma, so that the loop brightens. The following evolution depends on the duration of the heat release. Continuous heating allows the loop to reach quasi-equilibrium conditions at the highest possible density. With a short heat pulse the plasma cooling becomes important: after the heating, the temperature decreases rapidly by the very efficient conduction, but the density decreases much more slowly, leading to an overdensity over most of the loop's life. The loops are observed to be bright on time scales longer than the cooling times \citep[e.g.,][]{Rosner1978a}; the question is whether the heat release is really gradual and long-lasting, or is instead made of a sequence of short and localised heat pulses distributed in the loop cross-section\citep{Klimchuk2006a,Reale2014a}. In the latter case, the real structure of a loop would be that of a bundle of thinner flux tubes whose thickness is determined by the transverse size of the heat pulse. Evidence for overdensity \citep[e.g.,][]{Lenz1999a,Winebarger2003b}, multi-thermal plasma distribution \citep[e.g.,][]{Warren2011a} and some very hot plasma \citep[e.g.,][]{Reale2009b,Testa2012c,Miceli2012a} supports an impulsive heat release, and the question is now turning to how the energy is stored and released, what is the frequency of the pulses, what is the charging mechanism, what is the local conversion mechanism, and whether by the dissipation of waves or by resistive reconnection. \aftr{Intermittent heating and/or fine structuring is also predicted and discussed by some wave dissipation models, either Alfven \citep{van-Ballegooijen2011a,Asgari-Targhi2012a,Asgari-Targhi2013a,van-Ballegooijen2014a,Cranmer2015a} or kink modes \citep{Antolin2014a,Magyar2016a}.} A second complementary two-dimensional or three-dimensional approach investigates the way the freed magnetic energy powers coronal flux tubes. The stress of a magnetic flux tube has been studied due to twisting \citep[e.g.,][]{{Rosner1978b},{Golub1980a},Klimchuk2000b,Baty2000a,Torok2003a} and braiding of the field lines \citep{Lopez-Fuentes2010a,Wilmot-Smith2011a,Bingert2011a}. Most efforts have been devoted to study the conditions and effects of the resulting kink instability \citep{Hood1979b,Zaidman1989a,Velli1990a,Baty2000a,Gerrard2001a,Torok2004a}, and to the resulting formation of current sheets \citep{Velli1997a,Kliem2004a} and relaxation due to several dissipation mechanisms \citep{Hood2009b,Bareford2013a}. MHD simulations have shown the possible importance of local instabilities in the coronal magnetic field to trigger cascades to large-scale energy release \citep{Hood2016a}. A third approach is a large-scale one that ranges from the low chromosphere to the corona. It takes the magnetograms and the observations of photospheric granules, and of their dynamics, as boundary conditions to determine the structure and evolution of the upper atmosphere. This approach is able to describe the formation and powering of coronal loops in a qualitative or semi-quantitative way \citep{Gudiksen2005a,Bingert2011a}, including the emergence of flux tubes by magnetic twisting \citep{Martinez-Sykora2008a,Martinez-Sykora2009a}, so as to reproduce several observed features, such as a constant cross-section \citep{Peter2012a}, and to help interpret and use data analysis tools \citep{Testa2012b}. Recent work has supported episodic and structured heating due to the fragmentation of current sheets and/or turbulent cascades \citep{Hansteen2015a,Dahlburg2016a}. \new{Evidence for coherent widespread twisting of magnetic flux tubes has been found on the solar disk \citep{{Wedemeyer-Bohm2012a},De-Pontieu2014a} and well-studied \citep{Levens2015a} from optical and UV observations. This represents a natural stressing mechanism of the magnetic field that eventually leads to a relaxation and a release of magnetic energy \citep[e.g.,][]{{Rosner1978b},{Golub1980a},Klimchuk2000b,Lopez-Fuentes2003a}. \new{Here we set up a numerical experiment for a coronal magnetic flux tube that is anchored in the chromosphere and progressively twisted by the rotation of the plasma at the footpoints. Our approach is a step forward from 1-D loop modeling to allow an active role for the magnetic field, including the expansion of the field lines in the transition region and the energy production from field dissipation. Our choice has been to assume a relatively simple setup but still including many ingredients of a loop heated by magnetic dissipation. A coherent rotation of the footpoints with some moderate perturbation allows us to have some control on the effects in such a complex MHD system. As new achievements with respect to other previous MHD modeling of twisted loops, our modeling includes a highly non-uniform solar atmosphere and magnetic field with the boundary in the chromosphere. A fundamental target of our work is to reproduce the full typical evolution of a coronal loop, including the chromospheric evaporation driven by magnetic heating excess. This is not an easy task because it requires high spatial resolution \citep{Bradshaw2013a} in a 3D MHD framework. Our model aims also at accurately describing the temperature stratification and evolution to synthesize observables for diagnostics and direct comparison with observations.} We consider a complete loop atmosphere with a corona connected to two thick chromospheric layers by thin transition regions, immersed in a magnetic field. The magnetic field is arranged to be mostly uniform in the corona and strongly tapering in the chromosphere, where the ratio of thermal to magnetic pressure switches from low to high values ($\beta > 1$). Therefore, our model accounts for the interaction with the magnetic field including the critical region where $\beta$ changes regime. \new{The model also includes heating mechanisms that derive from the dissipation of the magnetic field. The heating is basically determined by the anomalous diffusivity that reconnects the magnetic field above a current density threshold. The currents grow because of the progressive twisting of the magnetic field. The field is twisted by random rotational plasma motions at the loop footpoints, which drag the field. We show that this magnetic stress and dissipation drives a typical coronal loop ignition, structuring and evolution.}

\new{This work describes a possible scenario of a coronal loop heated by magnetic field stressing. This is an evolution from the standard one-dimensional single, or multi-strand, loop modeling (see Section~\ref{sec:intro}) and a step forward in self-consistent MHD loop modeling.} Single loop models describe the hydrodynamics of a coronal atmosphere linked to the chromosphere through a steep transition region confined in a curved flux tube. The curvature appears only in the formulation of the gravity. The plasma moves and transports energy only along the tube, under the effect of a prescribed heating function. Here we maintain the same plasma atmosphere but we immerse it in an ambient ``cylindrical" magnetic field, which expands from the chromosphere up into the corona, as in \cite{Guarrasi2014a}. We no longer consider a prescribed heating function, but the heating is a consequence of stressing the magnetic field. This provides a self-consistent conversion of magnetic energy into heat. Our choice has been to start from simple, but realistic, assumptions on magnetic stressing and energy conversion: the magnetic field is stressed through the twisting driven by rotational footpoint motion in layers where plasma $\beta >> 1$; twisting is believed to be quite usual in the solar atmosphere \citep{{Wedemeyer-Bohm2012a},De-Pontieu2014a}. The rotation motion is perturbed as expected in the solar surface, and this is fundamental to break the symmetries and let currents fragment into sheets \citep{Rappazzo2013a,Nickeler2013a}. The sheets are progressively intensified by the twisting and this occurs more at the loop footpoints where the magnetic field is tapered crossing the transition region to the chromosphere. In this scenario we have hypothesised a switch-on dissipation mechanism. Heating from a very high anomalous resistivity is released as soon as the current density grows above a given threshold \citep{Hood2009b}, that we set to $2.25 \times 10^{11}$ esu s$^{-1}$ cm$^{-2}$, to mimic possible turbulent cascades or MHD avalanche \citep{Rappazzo2013b,Hood2016a}. Since here we address mostly the coronal evolution, the heating is assumed (in common with some previous simulations) to be active in the corona only, just because otherwise the magnetic field in the chromosphere is rapidly dissipated and we have found no way to refurbish it. Our single loop study supports other findings from MHD modeling of solar atmosphere boxes \citep[e.g.][]{Hansteen2015a}, and provides fine details. We start from a tenuous and cool atmosphere. Where the current grows above the threshold in the corona, the plasma begins to heat above 1 MK. The heating is more steady and efficient around the loop central axis, where the temperature rises above 3 MK on average in about half an hour. At the same time, the increasing pressure gradients determine the expansion of the chromospheric layers and the tube fills with denser plasma. The density gradually rises above $10^9$ cm$^{-3}$. This evaporation is in agreement with standard single loop models. From comparison with an equivalent simulation with ever-present anomalous resistivity, we have ascertained that the presence of the switch-on heating that leads to a factor two larger evaporation speeds, even in the late steady state. This is therefore a major difference between a gradual and an impulsive heating mechanism. Another important difference is the presence of overheated plasma. At variance from the gradual-heating simulation, the switch-on heating produces some amount of plasma significantly hotter than the average, as expected from impulsive heating \citep{Klimchuk2006a} and recently detected in bright active regions \citep[e.g.][]{Reale2009b,Reale2011a,Testa2012c,{Miceli2012a}}. The heating is more intense where the magnetic field is more intense, i.e. close to the footpoints, where it expands more. This is in agreement with other MHD modeling of the solar atmosphere \citep{Gudiksen2005a,Gudiksen2005b,Bingert2011a,Bingert2013a}. The heating is more intense around the central axis where the footpoint rotates and the twisting of the magnetic field is effective. The energy release determines a progressive dissipation of the magnetic field, through the local reconnection of sheared field lines. With our choice of magnetic diffusivity, this dissipation drives the plasma to density and temperature typical of active regions already at moderate twisting angles, far from the conditions to trigger kink instabilities \citep{Hood1979b,Hood1981a,Einaudi1983a,Velli1990a,Torok2003a}. The kink instability has been suggested as a trigger mechanism for the rapid heating of coronal loops. \cite{Hood2009b} have shown that its non linear development creates current sheets, and triggers magnetic reconnection. Once reconnection starts, the current sheets fragment, resulting in the dissipation of magnetic energy across the loop cross-section, as it relaxes towards its lowest energy state. Starting from a temperature of only $10^4$~K, their results show that plasma heating up to $10^7$~K and above is possible. \cite{Botha2011a} included thermal conduction so that lower temperatures were obtained. In addition to previous twisting models, our description includes the chromosphere and the transition region at reasonable resolution. Another important ingredient is the expansion of the magnetic field from the chromosphere to the corona \citep{Rosner1978b}, which, together with the change of $\beta$ regime in the chromosphere, stresses the importance of the non-linear interaction of the plasma and the magnetic field. Our model is also able to describe a significant mass transfer from the chromosphere to the corona, an essential feature for comparison with the observed loops brightness. The plasma produces realistic emission in X-ray and EUV bands. Its evaporation and the twisting drive significant spiraling motions as recently extensively observed \citep{De-Pontieu2014a}. Our choice here is to produce loop heating with a relatively ordered magnetic stressing, i.e. the progressive random twisting due to footpoint rotation. So we are not describing an entirely chaotic magnetic stress, determined by random photospheric motions that lead to magnetic braiding \citep{Lopez-Fuentes2010a,Wilmot-Smith2011a,Bingert2011a}. Our approach allows us to keep a tighter grasp on the physical effects that lead to the loop evolution, still maintaining a reasonable description of possible coronal drivers \citep{Rosner1978b}. An entirely ordered footpoint rotation would not lead to the formation of structured currents and heating. As mentioned above, an essential ingredient to have fine structure is a random motion at the footpoints. The deriving fine structure is both in space and time. We see filamentary structures on the cross-scale of a few hundreds kilometers, but also a structured heating with spikes that reach the scale of proper flare intensities, with durations on the scale of few tens of seconds. Evidence for loop fine structure is widespread \citep[e.g.][]{Vekstein2009a,Guarrasi2010a,{Viall2011a},{Antolin2012a},{Brooks2012b},{Brooks2013a},{Cirtain2013a},{Peter2013a},{Testa2013a},Tajfirouze2016a,{Tajfirouze2016b}}. The fine temporal and spatial structure that develops within our modeling deserves further investigation and will be the subject of future research. In particular, we plan to study the effects of: radial motions; different rotation profiles; different magnetic field strengths; and different initial magnetic configurations, especially those that lead to tectonics heating \citep{Priest2002a}.

1607.05500

1607

1607.05736_arXiv.txt

We used a combination of optical and near-UV Hubble Space Telescope photometry and FLAMES/ESO-VLT high-resolution spectroscopy to characterize the stellar content of the old and massive globular cluster (GC) NGC~121 in the Small Magellanic Cloud (SMC). We report on the detection of multiple stellar populations, the first case in the SMC stellar cluster system. This result enforces the emerging scenario in which the presence of multiple stellar populations is a distinctive-feature of old and massive GCs regardless of the environment, as far as the light element distribution is concerned. We find that second population (SG) stars are more centrally concentrated than first (FG) ones. More interestingly, at odds with what typically observed in Galactic GCs, we find that NGC~121 is the only cluster so far to be dominated by FG stars that account for more than $65\%$ of the total cluster mass. In the framework where GCs were born with a $90-95\%$ of FG stars, this observational finding would suggest that either NGC~121 experienced a milder stellar mass-loss with respect to Galactic GCs or it formed a smaller fraction of SG stars.

All relatively massive ($>4-5 \times 10^4 M_{\odot}$) and old ($10-13$ Gyr) Galactic globular clusters (GCs) studied so far host multiple stellar populations showing appreciable differences in the abundance of He and several other light elements \citep[e.g., C, N, Na, O, Al, Mg, see][for a review]{gratton12}. Hereafter, we define these multiple populations as light element MPs (LE-MPs), to distinguish them from the multiple populations showing large iron abundance differences, as those observed in $\omega$Cen \citep[see e.g.][]{pancino00,ferraro04,johnson10} and in Terzan~5 \citep{ferraro09,massari12,origlia13}. These LE-MPs manifest themselves as multi-modal or broadened evolutionary sequences in color-magnitude-diagrams (CMDs), when appropriate filters (or filter combinations) are adopted \citep[see, for example,][]{piotto07,marino09,monelli13,dalex11a,dalex14a,milone13,milone15}. Different scenarios have been proposed over the years to explain the formation of LE-MPs. In the most popular formation models, LE-MPs are the result of a self-enrichment process, which likely occurred in the very early epochs of GC evolution ($\sim 100$ Myr). A second generation (SG) formed from a combination of the ejecta of stars from a first population (polluters - FG) and from a ``pristine material'' \citep{decressin07,dercole08,demink09,conroy12,denis14}. However, all of the proposed scenarios face serious problems as none can explain more than few relevant observations \citep[e.g.][]{salaris14,bastian15,renzini15}. As a matter of fact we still lack a self-consistent explanation of the physical process(es) at the basis of the multiple population formation and evolution. To this end it is necessary to use a thorough and wide approach that combines photometric, spectroscopic and kinematical information. For example, crucial insights on the very first epoch of GC evolution can be obtained by observing and comparing clusters over a wide range of properties (metallicities, mass, structural parameters etc.) and in different environments (from dwarf to giant elliptical galaxies). One of the most debated and still poorly understood topics is what determines the fraction of SG stars in GCs. By using the available collection of spectroscopic and photometric data, \citet{bastian15} showed that the fraction of SG stars is remarkably uniform ($N_{\rm SG}/N_{\rm TOT}\sim0.68\pm0.07$) regardless of the GC mass, metallicity, and present-day galactocentric distance, with the only exceptions of NGC~6362 \citep{dalex14a,mucciarelli16} and NGC5272 \citep{massari16} that show equally-populated FG and SG. To reproduce this evidence, models require that all GCs had an initial mass up to 10-100 times larger than the current one, and they eventually lost $>90-95\%$ of their original mass through tidal-stripping or gas-expulsion \citep{dercole08,conroy12}. This strong requirement, which is common to any model, is typically known as the ``mass-budget problem'' and it may have important implications on the galaxy mass-assembly scenarios. However, it should be also noted that it appears to be inconsistent with some observational results (e.g. \citealt{larsen12}, \citealt{bastian15}) and GC dynamical evolution models (e.g. \citealt{krui15}). On the same line, important efforts have been made to detect LE-MPs in other galaxies. The presence of LE-MPs has been directly observed by means of photometry and/or spectroscopy of resolved stars in the old GCs of the Large Magellanic Clouds \citep[see, for example,][]{mucciarelli09} and in the Fornax dwarf galaxy \citep{larsen14}, while it has been indirectly inferred in the massive GCs of M31 \citep{schiavon13} and M87 \citep{chung11} thanks to integrated spectro-photometric studies. Quite surprisingly, similar investigations still lack for the relatively close ($d\sim60$kpc) Small Magellanic Cloud (SMC) irregular dwarf galaxy. The SMC harbors a population of GCs with an almost continuous distribution in age up to $t\sim 8$ Gyr \citep{harris04,carrera08,dias10}. The only significant exception is the older cluster NGC~121, a system located at $\sim$ 2.3$^{\circ}$ NW from the optical center of the SMC \citep{crowl01}. It has a mass $M\simeq 3\times 10^{5} M_{\odot}$ \citep{mackey03} and age $t_{\rm AGE}\sim 11$ Gyr \citep{glatt08}. These properties make NGC~121 an extremely interesting target, as {\it it is the only old, Galactic-like GC in the SMC}. In this work we present a spectro-photometric analysis of NGC~121 with the aim of characterizing its light element chemical patterns, the presence of LE-MPs (as traced by C, N, O, N, Mg and Al abundances) and their properties. The paper is structured as follows: in Section~2 the observational data-base is presented, in Sections~3 and 4 the results obtained by means of the photometric and spectroscopic data-sets, respectively, are discussed. In Section~5 we summarize the most relevant results and draw our conclusions.

The spectro-photometric analysis of the old and massive SMC cluster NGC~121 performed in this paper, provided the following main results: \begin{itemize} \item{NGC~121 shows a clear broadening and/or splitting of the RGB sequence in CMDs where the F336W filter is used (Section~\ref{phot}; Figures~\ref{cmd2} and \ref{color}). This result represents the first evidence that NGC~121 hosts LE-MPs. This cluster is also the first system to show such a property in the SMC, thus enforcing the emerging scenario in which LE-MPs are a distinctive-feature of old and massive GCs, regardless of the environment. In fact, this feature is observed in the Milky Way, in the Fornax dwarf galaxy and, now, in both the Magellanic Clouds.} \item{SG stars are more centrally concentrated than FG stars, as typically observed in Galactic GCs (Section~\ref{phot}; Figure~\ref{rad}).} \item{NGC~121 is dominated by FG stars that account for $65\%$ of the entire population of the cluster, at variance with what found in Galactic GCs, where the fraction of FG stars is smaller than the SG one.} \item{The analysis of high-resolution spectra of five RGB stars yields an average iron abundance [Fe/H] = $-1.28$ dex ($\sigma$ = 0.06 dex), with no hints of intrinsic iron spread (Section~\ref{spec}). Homogeneous light-element abundance ratios ([O/Fe], [Na/Fe] and [Mg/Fe] -- Figure~\ref{anti}) have been also found, consistent with all the analyzed stars being O-rich/Na-poor FG stars.} \end{itemize} \citet{dercole08} and \citet{vespe13} suggested that at birth GCs should be populated by $\sim95\%$ of FG stars. Then they reach present-day number ratios because of a preferential FG star loss due to both early stellar ($t<1-2 $ Gyr) and long-term dynamical evolutions. Figure~\ref{mass} shows the location of NGC~121 in the Mass {\it vs} $N_{\rm SG}/N_{\rm TOT}$ diagram, together with the values measured in NGC~6362 \citep{dalex14a,mucciarelli16} and NGC~5272 \citep{massari16}, by using the same photometric technique to distinguish among different sub-populations. For comparison we also show the average $N_{\rm SG}/N_{\rm TOT}$ ratio computed by Bastian \& Lardo (2015) using a collection of spectroscopic and photometric measurements for 33 Galactic GCs, which stays constant at a value of $(68\pm7\%)$. It is worth noticing that the SG fraction observed in NGC~121 appears to be consistent with the value predicted by the gas-expulsion model of \cite{khalaj15} based on N-body simulations (Figure~\ref{mass}), which foresees a decreasing SG star fraction with increasing the cluster mass. Such a similarity is quite remarkable, as at the mass of NGC~121, these models underpredict by a factor of three or more the fraction of SG stars in Galactic GCs (Figure~\ref{mass}). A dominant FG sub-population in NGC~121 can be explained with a less severe mass loss (still $\sim90\%$) experienced by this cluster, when compared to Galactic GCs, due to weak tidal interactions with the intrinsically shallower potential of the SMC, and to the fact that this GC orbits at a quite large distance (about 2.5 kpc) from the SMC center. Another possible explanation is that NGC~121 lost the same fraction of FG stars as Galactic GCs, but formed a smaller fraction of SG stars, because of, e.g., a less efficient self-enrichment processes, gas retention, the impact of different environments on LE-MP formation, etc. However, on top of these speculative arguments, it should be noted that the available FG and SG population ratios for Galactic GCs can be quite uncertain and not well representative of the entire cluster population. In fact, spectroscopic-based number ratios strongly depend on the the ability to separate different sub-populations that sometimes have not discretely different abundance patterns. Moreover, they are spatially incomplete by construction, typically not sampling the cluster central regions. On the other hand, photometric derivations are generally limited to the small HST field of view and cannot account for possible large-scale radial variations. Such potential biases in the derived population ratios in most of the Galactic GCs are a serious concern. Indeed, robust estimates obtained in NGC~6362 (Figure~\ref{mass}; \citealt{dalex14a}) and NGC~5272 \citep{massari16}, for which global number ratios are available (thanks to a proper combination of HST and ground-based wide-field observations), indicate $N_{\rm SG}/N_{\rm TOT}\sim50\%$, significantly lower than the average value of $\sim 0.7$ computed by \citet{bastian15}. Concluding, NGC~121 turns out to be an interesting and intriguing case among the old, massive GCs studied so far. Understanding the origin of its peculiarities, will likely provide important insights into the formation and early evolution of GCs and multiple populations. Moreover, performing similar studies on a larger sample of clusters will allow us to shed new light on the role of the environment and host galaxy properties on the multiple population properties.

1607.05736

1607

1607.02099_arXiv.txt

\noindent \vspace*{-4mm}\newline We consider a theory with isotropic nonbirefringent Lorentz violation in the photon sector and explore the effects on the development of the electromagnetic component of extensive air showers in the Earth atmosphere. Specifically, we consider the case of a ``fast'' photon with a phase velocity larger than the maximum attainable velocity of a massive Dirac fermion (this case corresponds to a negative Lorentz-violating parameter $\kappa$ in the action). Shower photons with above-threshold energies decay promptly into electron-positron pairs, instead of decaying by the conventional production of electron-positron pairs in the background fields of atomic nuclei. This rapid production of charged leptons accelerates the shower development, decreasing the atmospheric depth of the shower maximum ($X_\text{max}$) by an amount which could be measured by cosmic-ray observatories. Precise measurements of $X_\text{max}$ could then improve existing limits on the negative Lorentz-violating parameter $\kappa$ by several orders of magnitude.

Ever since the foundation of special relativity, electrodynamics has played a fundamental role in the establishment of elementary-particle-physics theory. Experimental tests using photons have provided valuable and sensitive probes for the search of potential deviations from exact Lorentz symmetry. In fact, some of the most stringent limits on dimensionless parameters controlling Lorentz violation have been found in the photon sector. Although laboratory experiments have access to some unique signatures, the determination of most of the best limits on Lorentz-violating (LV) parameters in various sectors have taken advantage of the high energies or the large propagation distances of ``astroparticles,'' i.e., cosmic rays~\cite{KlinkhamerRisse1,KlinkhamerRisse2,KS2008,DK2015}, gamma rays \cite{KS2008}, cosmic-microwave-background photons \cite{KM2007}, and neutrinos \cite{DKM2014,Diaz:2016c}. Here, we explore the potential of ultrahigh-energy photons to test LV effects and consider photons which are produced as secondary particles in air showers. In Sec.~\ref{sec:Theory}, we give the theoretical setup for isotropic Lorentz violation in the photon sector and review existing bounds on the relevant LV parameter. In Sec.~\ref{sec:EAS}, we discuss a simple model for the electromagnetic component of extensive air showers, first for the standard Lorentz-invariant theory and, then, for the Lorentz-violating theory considered. In Sec.~\ref{sec:summary}, we present some concluding remarks. As the scope of this article is restricted to the study of isotropic Lorentz violation in electrodynamics, we only consider isotropic Lorentz-violating effects in the photon sector. Notice that the corresponding effects from the electron sector can be moved to the photon sector by suitable spacetime coordinate transformations~\cite{BaileyKostelecky2004}. Furthermore, independent studies regarding Lorentz-violating effects involving weakly interacting particles can be found for muons \cite{Noordmans-etal2014}, pions~\cite{Antonov-etal2001,Altschul2007,Boncioli-etal2015}, and neutrinos \cite{Diaz2014}.

\label{sec:summary} In this article, we considered a theory with isotropic nonbirefringent Lorentz violation in the photon sector and studied the effects on the development of extensive air showers. In particular, we focused on the consequences of a negative value for the Lorentz-violating parameter $\ka$ because positive values are already well constrained \cite{KS2008}. The relevant nonstandard process for $\ka<0$ is photon decay into an electron-positron pair \cite{KS2008}, which substantially modifies the development of the electromagnetic component of the shower. Since the photon decay length is of the order of a meter or less, photons produced by the decay of neutral pions or by subsequent Bremsstrahlung emission will promptly decay into electron-positron pairs, well before conventional pair production can take place. This nonstandard decay process accelerates the particle multiplication in the shower, which reaches its maximum earlier than for the conventional Lorentz-invariant case. Two stages are identified in this electromagnetic shower. In the first stage, the rapid Lorentz-violating decay of Bremsstrahlung-photons produced with energies above the photon-decay threshold will effectively lead to the production of three charged leptons at each generation, one ``Bremsstrahlung-lepton'' and two ``photon-decay-leptons.'' This part of the shower contains mostly electrons and positrons, and the energy per particle will rapidly fall below the threshold for photon decay. In the second stage, the particle multiplication proceeds in the conventional manner. The energy per particle starts at the Lorentz-violating photon-decay threshold energy and drops until it is below the critical value $E_c$, at which moment the particle number in the shower reaches its maximum. We have found that negative $\ka$ could drastically reduce the shower-maximum depth $X_\text{max}$ of the electromagnetic shower as well as modify its energy dependence. The absence of these features could be used to significantly improve the existing limits on negative values of $\ka$.

1607.02099

1607

1607.07991_arXiv.txt

{Gas plays a major role in the dynamical evolution of young stellar objects (YSOs). Its interaction with the dust is the key to our understanding planet formation later on in the protoplanetary disc stage. Studying the gas content is therefore a crucial step towards understanding YSO and planet formation. Such a study can be made through spectroscopic observations of emission lines in the far-infrared, where some of the most important gas coolants emit, such as the [OI] $\rm ^{3}P_{1} \rightarrow$$\rm ^{3}P_{2}$ transition at 63.18 $\rm \mu m$. } {We provide a compilation of observations of far-IR lines in 362 YSOs covering all evolutionary stages, from Class 0 to Class III with debris discs. In the present paper we focus on [OI] and o-$\rm H_{2}O$ emission at 63 $\rm \mu m$.} {We retrieved all the available \textit{Herschel}-PACS spectroscopic observations at 63 $\rm \mu m$ that used the dominant observing mode, the chop-nod technique. We provide measurements of line fluxes for the [OI] $\rm ^{3}P_{1} \rightarrow$$\rm ^{3}P_{2}$ and o-$\rm H_{2}O$ $\rm 8_{08} \rightarrow 7_{17}$ transitions at 63 $\rm \mu m$ computed using different methods. Taking advantage of the PACS IFU, we checked for spatially extended emission and also studied multiple dynamical components in line emission.} {The final compilation consists of line and continuum fluxes at 63 $\rm \mu m$ for a total of 362 young stellar objects (YSOs). We detect [OI] line emission at 63 $\rm \mu m$ in 194 sources out of 362, and line absorption in another five sources. o-$\rm H_{2}O$ was detected in 42 sources. We find evidence of extended [OI] emission in 77 sources, and detect 3$\sigma$ residual emission in 71 of them. The number of sources showing extended emission decays from Class 0 to Class II. We also searched for different components contributing to the line emission, and found evidence for multiple components in 30 sources. We explored correlations between line emission and continuum emission and found a clear correlation between WISE fluxes from 4.6 to 22 $\rm \mu m$ and [OI] line emission. We conclude that the observed emission is typically a combination of disc, envelope and jet emission.} {}

Young stellar objects (YSO) are complex sources consisting of many components, such as the central source (protostellar or stellar), an envelope made of gas and dust, a circumstellar disc, stellar and disc winds, and large-scale collimated jets. Each of the components can contribute to different observables, such as photometry and line fluxes. A detailed study is therefore needed to elucidate the contribution of each component. In the initial stages of stellar formation, Class 0 and I protostars \citep{Lada1984,Lada1987,Andre1993} are surrounded by an envelope. Discs are clearly detected around Class I sources. Class I sources later evolve to Class II sources, in which the central star is already formed and the envelope dispersed. The formation of a dust opacity hole in the inner disc leads to the formation of the so-called transitional discs \citep{Strom1989}. Many mechanisms have been used to explain the formation of the inner opacity holes, including planet formation. At 10 Myr, most primordial discs have been dispersed \citep{Strom1989}, but destructive collisions between planetesimals can repopulate the circumstellar environment with dust, resulting in the so-called debris discs. Young stellar objects can also be classified according to their masses. The so-called T Tauri stars are variable stars showing bright emission lines with stellar masses $\rm M_{*} < 2.0 M_{\odot}$, while HAeBe stars are the high-mass counterparts of T Tauri stars ($\rm 2.0 < M/M_{\odot} < 8.0$). Although gas is thought to dominate the mass budget during the primordial stages (Class 0 to II), little is known about its mass and spatial distribution, mostly because it is difficult to detect $\rm H_{2}$, which lacks a permanent dipole moment. However, to learn about the formation of planets, we need to understand the chemical evolution of gas and dust. The \textit{Herschel Space Observatory} \citep{Pilbratt2010} produced thousands of observations of YSOs during its four-year mission. The most widely used instrument was the Photodetector Array Camera and Spectrometer \citep[PACS,][]{Poglitsch2010}, which can spectroscopically observe the far-IR 50-250 $\rm \mu m$ range. Furthermore, it also performed photometric observations at 70, 100 and 160 $\rm \mu m$ with great sensitivity. One of the most interesting characteristics of the PACS spectrometer is its Integral Field Unit (IFU), divided into 25 spaxels distributed in a regular grid covering 47\arcsec $\rm \times$47\arcsec. The IFU allows us to study the spatial distribution of the continuum and line emission. Some studies have surveyed [OI], CO, OH, and $\rm H_{2}O$ emission in objects belonging to different stellar associations and moving groups using \textit{Herschel} \citep{Donaldson2012,Howard2013,Green2013,Mathews2013,Lindberg2014,Riviere2013,Riviere2014,Riviere2015}. Other studies have focused on the analysis of individual sources \citep{Meeus2010, vanKempen2010, Kempen2010, Sturm2010, Thi2010, Tilling2012, Lebreton2012, Riviere2012b, Thi2013}. However, the spatial extension of the emission was discussed in only a few cases \citep[][]{Karska2013,Karska2014B,Nisini2013,Nisini2015}. The most extensively studied wavelength range is 63.0-63.4 $\rm \mu m$, which includes two transitions, [OI] $\rm ^{3}P_{1} \rightarrow $$\rm^{3}P_{2}$ at 63.185 $\rm \mu m$ and o-$\rm H_{2}O$ $\rm 8_{08} \rightarrow 7_{17}$ at 63.325 $\rm \mu m$. [OI] emission has been detected in YSOs at all evolutionary stages, from Class 0 and I \citep{Green2013} to Class II and transitional \citep{Howard2013} and debris discs \citep{Riviere2012}. o-$\rm H_{2}O$ emission was observed around Class 0, I, II and transition discs, but not around debris discs. Understanding the spatial distribution of far-IR lines emission is crucial, since it has been shown that envelopes, protoplanetary discs, and outflows can contribute to [OI] emission \citep{vanKempen2010,Podio2012, Karska2013}. [OI] extended emission along the jet direction has been commonly observed, while molecular extended emission is observed in only a few cases \citep{vanKempen2010,Herczeg2012}. \cite{Podio2012} and \cite{Karska2013} explained the extended emission as being produced by J- and C-shocks along the jet, and noted a decay in far-IR lines intensity from Class 0/I to Class II. \cite{Howard2013} studied a sample of Class II sources in Taurus, including sources with and without a jet or an outflow. The authors found a tight correlation between continuum emission at 63 $\rm \mu m$ and [OI] emission, suggesting a disc origin for the line. However, sources with jets show a brighter [OI] emission for the same level of continuum, indicating a contribution from the jet. The authors did not find a correlation between disc mass (derived from sub-millimeter continuum emission) and [OI] line intensity, indicating that either the line is optically thick or it is a poor tracer of gas mass. \cite{Green2013} studied a sample of 30 embedded sources (Class 0 and I) from the DIGIT program \citep[see e.g.][]{Kempen2010,Sturm2010} and found a tight correlation between line intensity and $\rm L_{bol}$. In this paper, we present a compilation of 432 PACS spectroscopic observations of 362 YSOs and main-sequence stars with debris discs. We focus on the small wavelength range between 63.0 and 63.4 $\rm \mu m$, which includes the [OI] transition at 63.185 $\rm \mu m$ and the o-$\rm H_{2}O$ transition at 63.325 $\rm \mu m$. Our wavelength range selection is motivated by the fact that the [OI] transition at 63.185 $\rm \mu m$ typically is the strongest line coolant in protoplanetary discs \citep{Gorti2008}. We leave the study of other transitions observed in \textit{PACS} range mode for a future paper.

\label{ref:SumConc} We have compiled \textit{Herschel}-PACS observations of [OI] and o-$\rm H_{2}O$ at 63 $\rm \mu m$ in YSOs, including Class 0, I, II, transitional discs, and debris discs, for a total of 432 observations of 362 sources. We note that the [OI] emission line intensity, as well as detection fractions, decreases during the evolution from Class 0 to debris discs. However, we did not see a difference in [OI] emission between Class 0 and Class I, nor between Class II and transition discs. o-$\rm H_{2}O$ emission line intensity also decreases from Class 0 and I to more evolved sources (Class II and transition discs). By means of comparing the fluxes computed from the central spaxel, the central 3x3 spaxels and the integrated IFU, we detected extended emission in the [OI] line for a total of 77 sources. For those sources showing hints of extended [OI] emission, we obtained line emission maps and residual maps, and confirmed residual emission in 71 sources. The fraction of sources showing extended emission decreases dramatically from Class 0, where 63\% of the sources show extended emission, to Class II, where only 17\% of the sources show extended emission. We detected extended o-$\rm H_{2}O$ line emission in only one source. For 30 sources in the sample we were able to fit multiple components to the line emission profile, which is indicative of different contributions to the line (envelope, discs, winds, and jets). We have tested previously identified correlations in the entire sample. The [OI] line emission correlates with continuum emission at 63 $\rm \mu m$ for all classes, with the exception of of debris discs. We confirm the correlation between [OI] and o-$\rm H_{2}O$ at 63 $\rm \mu m$, and tentatively see a change in slope in the correlation between class 0 and I sources and class II sources. We have identified new correlations with continuum emission between 4.6 and 22 $\rm \mu m$, indicating an extended emitting region (from the inner disc to tens of au) as the origin of the disc contribution. \onllongtab{1}{ \begin{longtable}{llllll} \caption{\label{YSO_sample} Sample of YSOs observed with \textit{Herschel}-PACS}\\ \hline\hline Source name & RA & Dec & Sp. type$^{1}$ & Disc type & Association \\ -- & (deg) & (deg) & -- & -- & -- \\ \hline \endfirsthead \caption{continued.}\\ \hline\hline Source name & RA & Dec & Sp. type & Disc type & Association \\ -- & (deg) & (deg) & -- & -- \\ \hline \endhead \hline \endfoot \multicolumn{6}{l}{Spectral types are taken from: (1) \cite{Torres2006}, (2) \cite{Houk1975},(3) \cite{Houk1988},}\\ \multicolumn{6}{l}{(4) \cite{Gray2006}, (5) \cite{Harlan1974}, (6) \cite{Luhman2010}, (7) \cite{Mohanty2013}, (8) \cite{Mora2001},}\\ \multicolumn{6}{l}{(9) \cite{Manoj2006}, (10) \cite{Meeus2013}, (11) \cite{Houk1999}, (12) \cite{Vieira2003}, }\\ \multicolumn{6}{l}{(13) \cite{Zuckerman2004}, (14) \cite{Webb1999}, (15) \cite{Rydgren1980}, (16) \cite{Furlan2009}, }\\ \multicolumn{6}{l}{(17) \cite{Vacca2011}, (18) \cite{Luhman2007}, (19) \cite{Luhman2004}, (20) \cite{deLaReza1989}, }\\ \multicolumn{6}{l}{(21) \cite{Gregorio-Hetem1992}, (22) \cite{Houk1978}, (23) \cite{Levenhagen2006}, (24) \cite{Weinberger2004}, }\\ \multicolumn{6}{l}{(25) \cite{Houk1982}, (26) \cite{Spezzi2008}, (27) \cite{Fedele2013}, (28) \cite{Bouy2009}, }\\ \multicolumn{6}{l}{(29) \cite{Alcala2014}, (30) \cite{Lopez2011}, (31) \cite{Carpenter2006}, (32) \cite{Kohler2000}, }\\ \multicolumn{6}{l}{(33) \cite{Kraus2008}, (34) \cite{Dent2013}, (35) \cite{Hughes1994}, (36) \cite{Mortier2011},}\\ \multicolumn{6}{l}{(37) \cite{Prato2003}, (38) \cite{Preibisch1998}, (39) \cite{Preibisch2002}, (40) \cite{Wichmann1997},}\\ \multicolumn{6}{l}{(41) \cite{McClure2010}, (42) \cite{Bouvier1992}, (43) \cite{Ricci2010}, (44) \cite{Erickson2011},}\\ \multicolumn{6}{l}{(45) \cite{Martin1998}, (46) \cite{Nilsson2009}, (47) \cite{Hales2014}, (48) \cite{Zuckerman2004b}}\\ \endlastfoot \object{HD~105} & 1.468958 & -41.753067 & G0$^{1}$ & DD & TucHor \\ \object{HD~377} & 2.107291 & 6.616806 & G2$^{1}$ & DD & - \\ \object{HD~3003} & 8.182958 & -63.0315 & A0$^{2}$ & DD & TucHor \\ \object{HD~3670} & 9.736267 & -52.534284 & F5$^{2}$ & DD & - \\ \object{HD~9672} & 23.657412 & -15.676359 & A1$^{3}$ & DD & Argus \\ \object{tau~Ceti} & 26.017 & -15.937472 & G8.5$^{4}$ & DD & - \\ \object{W3~IRS5} & 36.419125 & 62.097361 & - & 0 & - \\ \object{HD~15115} & 36.567667 & 6.292556 & F4$^{5}$ & DD & - \\ \object{HD~16743} & 39.781511 & -52.934806 & F1$^{2}$ & DD & - \\ \object{SMM~J032537+30451} & 51.343 & 30.75386 & - & 0 & Per \\ \object{LDN~1448N} & 51.40204 & 30.75616 & - & 0 & Per \\ \object{L~1448-C(S)} & 51.413 & 30.73283 & - & 0/I & Per \\ \object{IRAS~03235+3004} & 51.65612 & 30.25780 & - & 0 & Per \\ \object{L~1455} & 52.0167 & 30.1733 & - & 0 & Per \\ \object{2MASS~J03283706+3113310} & 52.15454 & 31.22522 & - & I & Per \\ \object{SSTc2d~J032857.4+311416} & 52.239 & 31.23775 & - & I & Per \\ \object{SSTc2d~J032900.5+311200} & 52.25229 & 31.20022 & - & I & Per \\ \object{2MASS~J03290149+3120208} & 52.2565 & 31.33905 & - & I & Per \\ \object{2MASS~J03290773+3121575} & 52.28241 & 31.36591 & - & I & Per \\ \object{NGC~1333IRAS4A} & 52.293708 & 31.22527 & - & 0 & Per \\ \object{SSTc2dJ~032910.7+311821} & 52.2945 & 31.30572 & - & 0 & Per \\ \object{NGC~1333IRAS4B} & 52.300042 & 31.21894 & - & 0 & Per \\ \object{SSTc2d~J032912.0+311301} & 52.30025 & 31.21713 & - & 0 & Per \\ \object{SSTc2d~J032913.5+311358} & 52.30641 & 31.23283 & - & 0 & Per \\ \object{HBC~347} & 52.409875 & 24.510556 & K1$^{6}$ & III & Taurus \\ \object{LDN~1455~IRS3} & 52.001711 & 30.133690 & - & 0/I & - \\ \object{IRAS~03267+3128} & 52.46591 & 31.65166 & - & 0 & Per \\ \object{IRAS~03271+3013} & 52.56308 & 30.39705 & - & I & Per \\ \object{IRAS~03282+3035} & 52.83741 & 30.75836 & - & 0 & Per \\ \object{HD~21997} & 52.973542 & -25.614139 & A3$^{3}$ & DD & - \\ \object{IRAS~03292+3039} & 53.07483 & 30.82986 & - & 0 & Per \\ \object{epsilon~Eri} & 53.232667 & -9.45825 & K2$^{4}$ & DD & - \\ \object{IRAS~03301+3111} & 53.303333 & 31.356722 & - & I & Per \\ \object{SSTc2d~J033314.3+310710} & 53.30991 & 31.11969 & - & 0 & Per \\ \object{SSTc2d~J033316.4+310653} & 53.3185 & 31.11458 & - & 0 & Per \\ \object{B1-a} & 53.319583 & 31.132 & - & I & Per \\ \object{B1-c} & 53.32437 & 31.15886 & - & 0/I & Per \\ \object{SSTc2d~J033327.3+310710} & 53.3637 & 31.1195 & - & I & Per \\ \object{IRAS~03407+3152} & 55.9855 & 32.01466 & - & 0 & Per \\ \object{SSTc2d~J034356.8+320305} & 55.98683 & 32.05130 & - & 0 & Per \\ \object{2MASS~J03444389+3201373} & 56.18316 & 32.02672 & - & 0 & Per \\ \object{LkHa~330} & 56.45117 & 32.40331 & G2 & TD & Per \\ \object{HBC~356} & 60.808292 & 25.883306 & K2$^{6}$ & III & Taurus \\ \object{HBC~358} & 60.961833 & 26.181444 & M2$^{6}$ & III & Taurus \\ \object{IRAS~04016+2610} & 61.1794625 & 26.315663 & - & I & - \\ \object{PP~13S} & 62.671458 & 38.132028 & - & I & - \\ \object{LkCa~1} & 63.308917 & 28.319678 & M4$^{6}$ & III & Taurus \\ \object{V1096~Tau} & 63.363458 & 28.273556 & M0$^{6}$ & III & Taurus \\ \object{V773~Tau} & 63.553833 & 28.203458 & K3$^{6}$ & II & Taurus \\ \object{FM~Tau} & 63.556583 & 28.213667 & M0$^{6}$ & II & Taurus \\ \object{CW~Tau} & 63.570833 & 28.182722 & K3$^{6}$ & II & Taurus \\ \object{CIDA~1} & 63.573375 & 28.102694 & M5.5$^{6}$ & II & Taurus \\ \object{CX~Tau} & 63.699417 & 26.803058 & M3$^{6}$ & TD & Taurus \\ \object{LkCa~3} & 63.699875 & 27.876292 & M1$^{6}$ & III & Taurus \\ \object{FO~Tau} & 63.705375 & 28.2085 & M2$^{6}$ & T & Taurus \\ \object{CIDA-2} & 63.7715 & 28.146167 & M5.5$^{6}$ & III & Taurus \\ \object{LkCa~4} & 64.117125 & 28.126614 & K7$^{6}$ & III & Taurus \\ \object{CY~Tau} & 64.390542 & 28.346361 & M1.5$^{6}$ & II & Taurus \\ \object{LkCa~5} & 64.41225 & 28.550142 & M2$^{6}$ & III & Taurus \\ \object{IRAS~04158+2805} & 64.74225 & 28.206528 & - & I & Taurus \\ \object{FQ~Tau} & 64.803375 & 28.492528 & M3$^{6}$ & II & Taurus \\ \object{BP~Tau} & 64.815958 & 29.107472 & K7$^{6}$ & II & Taurus \\ \object{V819~Tau} & 64.859417 & 28.437306 & K7$^{6}$ & II/III & Taurus \\ \object{FR~Tau} & 64.897708 & 28.456056 & - & TD & Taurus \\ \object{LkCa~7} & 64.921958 & 27.830136 & K7$^{6}$ & III & Taurus \\ \object{IRAS~04169+2702} & 64.99333 & 27.16583 & - & I & Taurus \\ \object{IRAS~04181+2654} & 65.2975 & 27.01916 & - & I & Taurus \\ \object{DE~Tau} & 65.481833 & 27.918361 & M1$^{6}$ & II & Taurus \\ \object{IRAM~04191} & 65.487083 & 15.496083 & - & 0 & - \\ \object{RY~Tau} & 65.489167 & 28.443206 & K1$^{6}$ & TD & Taurus \\ \object{HD~283572} & 65.495208 & 28.301844 & G5$^{6}$ & III & Taurus \\ \object{T~Tau} & 65.497625 & 19.535103 & M3$^{6}$ & I/II & Taurus \\ \object{2MASS~J04220069+2657324} & 65.5025 & 26.95888 & - & I & Taurus \\ \object{FS~Tau} & 65.509083 & 26.958472 & M0$^{6}$ & II/Flat & Taurus \\ \object{FS~TauB} & 65.509071 & 26.958469 & - & I & Taurus \\ \object{FU~Tau} & 65.8975 & 25.05075 & M7.25$^{6}$ & II & Taurus \\ \object{FT~Tau} & 65.913292 & 24.93725 & M3$^{6}$ & II & Taurus \\ \object{IP~Tau} & 66.237833 & 27.199028 & M0$^{6}$ & TD & Taurus \\ \object{J1-4872} & 66.323667 & 26.297336 & K7/M1$^{6}$ & III & Taurus \\ \object{F~VTau} & 66.723042 & 26.115111 & K5$^{6}$& II & Taurus \\ \object{DG~TauB} & 66.760667 & 26.091861 & - & I & Taurus \\ \object{DF~Tau} & 66.762833 & 25.706472 & M2$^{6}$ & II & Taurus \\ \object{DG~Tau} & 66.769583 & 26.104528 & K6$^{6}$ & II & Taurus \\ \object{IRAS~04248+2612} & 66.98875 & 26.32166 & - & I & Taurus \\ \object{IRAS~04264+2433} & 67.375 & 24.66527 & - & I & Taurus \\ \object{FW~Tau} & 67.373792 & 26.281444 & M4$^{6}$ & III & Taurus \\ \object{DH~Tau} & 67.425083 & 26.548111 & M2$^{6}$ & II & Taurus \\ \object{IQ~Tau} & 67.464833 & 26.112472 & M0.5$^{6}$ & II & Taurus \\ \object{CFHT~20} & 67.497958 & 24.552167 & M5$^{6}$ & II & Taurus \\ \object{UX~Tau} & 67.516667 & 18.230389 & K5$^{6}$ & TD & Taurus \\ \object{FXTau} & 67.623375 & 24.445833 & M1$^{6}$ & II & Taurus \\ \object{DK~Tau} & 67.684333 & 26.023556 & K6$^{6}$ & II & Taurus \\ \object{ZZ~Tau} & 67.714083 & 24.706194 & M3$^{6}$ & II & Taurus \\ \object{ZZ~TauIRS} & 67.715458 & 24.696528 & M5$^{6}$ & II & Taurus \\ \object{IRAS~04287+1801} & 67.891987 & 18.134694 & - & I & Taurus \\ \object{V927~Tau} & 67.84925 & 24.181369 & M4.7$^{6}$ & III & Taurus \\ \object{HL~Tau} & 67.9101 & 18.232681 & K7$^{6}$ & II & Taurus \\ \object{XZ~Tau} & 67.9145 & 18.232139 & M2$^{6}$ & II & Taurus \\ \object{HK~Tau} & 67.960708 & 24.405019 & M0.5$^{6}$ & II & Taurus \\ \object{V710~Tau} & 67.990833 & 18.35975 & M0.5$^{6}$ & II & Taurus \\ \object{Haro~6-13} & 68.064208 & 24.48325 & M0$^{6}$ & II & Taurus \\ \object{GG~Tau} & 68.126458 & 17.527944 & M5.5$^{6}$ & II & Taurus \\ \object{UZ~Tau} & 68.178458 & 25.875389 & M1$^{6}$ & II & Taurus \\ \object{GH~Tau} & 68.276792 & 24.162361 & M2$^{6}$ & II & Taurus \\ \object{V807~Tau} & 68.277621 & 24.165278 & K5$^{6}$ & II & Taurus \\ \object{GI~Tau} & 68.392958 & 24.353194 & K7$^{6}$ & II & Taurus \\ \object{GK~Tau} & 68.394 & 24.351611 & K7$^{6}$ & II & Taurus \\ \object{DL~Tau} & 68.41275 & 25.343944 & K7$^{6}$ & II & Taurus \\ \object{HN~Tau} & 68.41395 & 17.86454 & K5$^{6}$ & II & Taurus \\ \object{DM~Tau} & 68.453 & 18.169444 & M1$^{6}$ & T & Taurus \\ \object{CI~Tau} & 68.466667 & 22.841722 & K7$^{6}$ & II & Taurus \\ \object{AA~Tau} & 68.730917 & 24.481444 & K7$^{6}$ & II & Taurus \\ \object{HO~Tau} & 68.834167 & 22.537389 & M0.5$^{6}$ & II & Taurus \\ \object{FF~Tau} & 68.837083 & 22.906722 & K7$^{6}$ & III & Taurus \\ \object{DN~Tau} & 68.864042 & 24.249694 & M0$^{6}$ & TD & Taurus \\ \object{IRAS~04325+2402} & 68.89708 & 24.13861 & - & I & Taurus \\ \object{HP~Tau} & 68.969917 & 22.906417 & K3$^{6}$ & II & Taurus \\ \object{J04381486} & 69.561917 & 26.194417 & M7.25$^{6}$ & II & Taurus \\ \object{GM~Tau} & 69.588917 & 26.153806 & M6.5$^{6}$ & II & Taurus \\ \object{DO~Tau} & 69.619083 & 26.180389 & M0$^{6}$ & II & Taurus \\ \object{HV~Tau} & 69.647 & 26.177397 & - & I & Taurus \\ \object{TMR~1B} & 69.807875 & 25.889 & - & I & Taurus \\ \object{TMR~1} & 69.807917 & 25.889056 & - & I & Taurus \\ \object{VY~Tau} & 69.822542 & 22.798167 & M0$^{6}$ & II & Taurus \\ \object{LkCa~15} & 69.824167 & 22.350972 & K5$^{6}$ & TD & Taurus \\ \object{2MASS~J04393364+2359212} & 69.890208 & 23.989222 & M5$^{6}$ & II & Taurus \\ \object{IRAS~04365+2535} & 69.89421 & 25.69685 & - & I & Taurus \\ \object{BD~Tau4} & 69.947833 & 26.028 & M7$^{7}$ & II & Taurus \\ \object{L~1527} & 69.97458 & 26.0527 & - & 0 & Taurus \\ \object{IRAS~04381+2540} & 70.302708 & 25.77663 & - & I & Taurus \\ \object{CoKu~Tau4} & 70.32 & 28.666833 & M1.5$^{6}$ & II & Taurus \\ \object{IRAS~04385+2550} & 70.41175 & 25.940778 & M0$^{6}$ & II & Taurus \\ \object{DP~Tau} & 70.657083 & 25.260417 & M0.5$^{6}$ & II & Taurus \\ \object{GO~Tau} & 70.762875 & 25.338556 & M0$^{6}$ & II & Taurus \\ \object{DQTau} & 71.721042 & 17.000056 & M0$^{6}$ & II & Taurus \\ \object{Haro6-37} & 71.74575 & 17.043944 & K7$^{6}$ & II & Taurus \\ \object{DS~Tau} & 71.950458 & 29.420681 & K5$^{6}$ & II & Taurus \\ \object{UY~Aur} & 72.947417 & 30.787083 & M0$^{6}$ & II & Taurus \\ \object{GM~Aur} & 73.79575 & 30.366528 & K7$^{6}$ & T & Taurus \\ \object{AB~Aur} & 73.940958 & 30.551214 & B9$^{6}$ & II & Taurus \\ \object{SU~Aur} & 73.997417 & 30.5671 & G2$^{6}$ & II & Taurus \\ \object{HD~31648} & 74.692777 & 29.843611 & A5$^{8}$ & HAeBe & - \\ \object{HD~32297} & 75.614319 & 7.461022 & A0$^{1}$& DD & - \\ \object{V836Tau} & 75.7775 & 25.388808 & K7$^{6}$ & TD & Taurus \\ \object{RW~Aur} & 76.956417 & 30.401408 & K3$^{6}$ & II & Taurus \\ \object{HD~35187} & 81.004875 & 24.960444 & A2+A7$^{9}$ & HAeBe & - \\ \object{HD~35841} & 81.652442 & -22.489923 & F3$^{3}$ & DD & GAYA2 \\ \object{HD~36112} & 82.614707 & 25.332523 & A5$^{10}$ & HAeBe & Taurus \\ \object{HD~36910} & 83.993612 & 24.748359 & F3$^{10}$ & HAeBe & Taurus \\ \object{V833~Ori} & 84.575417 & -7.040611 & - & I & - \\ \object{HD~245906} & 84.877 & 26.331972 & A6$^{9}$ & II & - \\ \object{[SMZ2000]~L1643-S3MMS1} & 84.982917 & -7.507778 & - & 0 & - \\ \object{Re50~NNIRS} & 85.114167 & -7.458667 & - & I & - \\ \object{HD~38120} & 85.799542 & -4.997194 & B9$^{11}$ & HAeBe & - \\ \object{HD~38207} & 85.837331 & -20.189297 & F2$^{3}$ & DD & - \\ \object{HD~38206} & 85.840293 & -18.557475 & A0$^{3}$ & DD & - \\ \object{V1647~Ori} & 86.5547 & -0.101333 & - & I & - \\ \object{HR~1998} & 86.73892 & -14.82195 & A2$^{4}$ & DD & - \\ \object{NGC~2071IR} & 86.768333 & 0.36361 & - & 0 & - \\ \object{$\beta$~Pic} & 86.821208 & -51.066528 & A6$^{4}$ & DD & BPMG \\ \object{R~Mon} & 99.791458 & 8.736028 & B8$^{4}$ & HAeBe & - \\ \object{HD~50138} & 102.889167 & -6.9665 & A1$^{11}$ & HAeBe & - \\ \object{PDS~27} & 109.89975 & -17.655 & B2$^{12}$ & HAeBe & - \\ \object{HD~61005} & 113.94775 & -32.203889 & G8$^{2}$ & DD & - \\ \object{Bran~76} & 117.648333 & -33.106639 & - & I & - \\ \object{RECX~1} & 129.234292 & -78.945972 & K4$^{13}$ & III & EtaCha \\ \object{RECX~14} & 130.37625 & -78.885139 & M4$^{13}$ & TD & EtaCha \\ \object{RECX~3} & 130.404292 & -79.058444 & M3$^{13}$ & TD & EtaCha \\ \object{RECX~4} & 130.598875 & -79.0675 & K7$^{13}$ & TD & EtaCha \\ \object{RECX~5} & 130.612958 & -78.963306 & M5$^{13}$ & TD & EtaCha \\ \object{RECX~6} & 130.661667 & -78.911889 & M2$^{13}$ & III & EtaCha \\ \object{RECX~8} & 130.800958 & -79.070083 & A7$^{2}$ & III & EtaCha \\ \object{RECX~15} & 130.827417 & -79.088389 & M2$^{13}$ & II & EtaCha \\ \object{J0844.2-7833} & 131.038125 & -78.562694 & M5.5$^{13}$ & II & EtaCha \\ \object{RECX~9} & 131.06825 & -78.985583 & M4$^{13}$ & TD & EtaCha \\ \object{RECX~10} & 131.132833 & -78.775333 & K7$^{13}$ & III & EtaCha \\ \object{[MGL99]~IRS1757} & 131.644833 & 43.908472 & - & I & - \\ \object{RECX~11} & 131.756917 & -78.992917 & K4$^{13}$ & II & EtaCha \\ \object{RECX~12} & 131.986542 & -78.914778 & M2$^{13}$ & III & EtaCha \\ \object{TWA~07} & 160.625458 & -33.671167 & M1$^{14}$ & DD & TWA \\ \object{SZ~Cha} & 164.569875 & -77.288083 & K0$^{15}$ & TD & Cha \\ \object{CR~Cha} & 164.779125 & -77.027889 & K2$^{16}$ & TD & Cha \\ \object{TWA~01} & 165.466333 & -34.704722 & M2.5$^{17}$ & TD & TWA \\ \object{CS~Cha} & 165.603792 & -77.559917 & K6$^{16}$ & TD & Cha \\ \object{CHX~7} & 166.564208 & -77.365806 & G5$^{1}$ & TD & Cha \\ \object{Eso-HA~559} & 166.606458 & -76.561639 & M5$^{18}$ & TD & ChaI \\ \object{2MASS~J11064658-7722325} & 166.694084 & -77.375710 & - & I & - \\ \object{[NC98]~Cha~HA~1} & 166.8195 & -77.598139 & M8$^{19}$ & II & ChaI \\ \object{Sz~18} & 166.829792 & -76.051333 & M2.5$^{19}$ & TD & Cha \\ \object{HD~97048} & 167.013833 & -77.654861 & A0$^{10}$ & HAeBe & ChaI \\ \object{HP~Cha} & 167.062917 & -77.564778 & G7$^{18}$ & Flat & ChaI \\ \object{Sz~27} & 167.162708 & -77.267833 & K8$^{18}$ & TD & ChaI \\ \object{TWA~02AB} & 167.307542 & -30.027722 & M2$^{1}$ & III & TWA \\ \object{GM~Cha} & 167.3687 & -76.5578 & - & I & Cha I \\ \object{T42} & 167.472542 & -76.57375 & K4.7$^{19}$ & II & ChaI \\ \object{WW~Cha} & 167.500458 & -76.58275 & K5$^{18}$ & TD & ChaI \\ \object{TWA~03A} & 167.616167 & -37.531111 & M3$^{20}$ & TD & TWA \\ \object{Hn~13} & 167.733208 & -76.759056 & M6$^{19}$ & II & ChaI \\ \object{CV~Cha} & 168.1155 & -76.739528 & G9$^{18}$ & II & ChaI \\ \object{CHX~22} & 168.177875 & -77.373056 & G8$^{18}$ & TD & ChaI \\ \object{Sz~45} & 169.404208 & -77.07725 & M0.5$^{19}$ & TD & ChaI \\ \object{TWA13AB} & 170.321833 & -34.779306 & M1$^{1}$ & III & TWA \\ \object{HD~98800B} & 170.522083 & -24.777583 & M5$^{21}$ & DD/TD & TWA \\ \object{HD~98922} & 170.631958 & -53.369861 & B9$^{22}$ & HAeBe & - \\ \object{HD~100453} & 173.27325 & -54.324583 & A9$^{12}$ & HAeBe & - \\ \object{HD~100546} & 173.356 & -70.194778 & B9$^{23}$ & HAeBe & Ass Sco OB 2-4 \\ \object{T~Cha} & 179.306375 & -79.35875 & K0$^{1}$ & TD & Cha \\ \object{HD~104237} & 180.021167 & -78.192944 & A4$^{10}$ & HAeBe & - \\ \object{IRAS~11590-6452} & 180.404583 & -65.14833 & - & 0 & - \\ \object{TWA~23} & 181.864083 & -32.783417 & M1$^{24}$ & III & TWA \\ \object{MML~17} & 185.638458 & -53.563611 & G2 & DD$^{1}$ & Ass Sco OB 2-4 \\ \object{TWA~10} & 188.767708 & -41.610722 & M2.5$^{14}$ & III & TWA \\ \object{HR~4796A} & 189.004298 & -39.869505 & A0$^{25}$ & DD & TWA \\ \object{DK~Cha} & 193.321792 & -77.119639 & - & I & ChaII \\ \object{ChaII~J125342.86-771511.5} & 193.428583 & -77.253194 & - & I & ChaII \\ \object{ChaII~J125633.66-764545.3} & 194.14025 & -76.762583 & M1$^{26}$ & II & ChaII \\ \object{ChaII~J125711.77-764011.3} & 194.299042 & -76.669806 & M0$^{26}$ & II & ChaII \\ \object{ChaII~J125806.78-770909.4} & 194.52825 & -77.152611 & M7$^{26}$ & II & ChaII \\ \object{2MASSJ12590656-7707401} & 194.7774 & -77.1277 & - & I & ChaII \\ \object{ChaII~J130055.36-771022.1} & 195.230667 & -77.172806 & M3$^{26}$ & II & ChaII \\ \object{ChaII~J130158.94-775121.7} & 195.495583 & -77.856028 & K8.5$^{26}$ & II & ChaII \\ \object{ChaII~J130222.85-773449.3} & 195.595208 & -77.580361 & M5$^{26}$ & II & ChaII \\ \object{ChaII~J130424.92-775230.1} & 196.103833 & -77.875028 & M2.5$^{26}$ & II & ChaII \\ \object{Hn~24} & 196.232292 & -77.66375 & M0$^{26}$ & TD & ChaII \\ \object{ChaII~J130508.53-773342.4} & 196.285542 & -77.561778 & M2.5$^{26}$ & II & ChaII \\ \object{ChaII~J130512.69-773052.3} & 196.302875 & -77.514528 & M1$^{26}$ & II & ChaII \\ \object{ChaII~J130520.68-773901.4} & 196.336167 & -77.650389 & K5$^{26}$ & II & ChaII \\ \object{ChaII~J130521.66-773810.0} & 196.34025 & -77.636111 & -- & II & ChaII \\ \object{ChaII~J130529.04-774140.1} & 196.371 & -77.694472 & -- & II & ChaII \\ \object{ChaII~J130718.05-774052.9} & 196.825208 & -77.681361 & M4.5$^{26}$ & II & ChaII \\ \object{ChaII~J130748.51-774121.4} & 196.952125 & -77.689278 & M2$^{26}$ & II & ChaII \\ \object{ChaII~J130806.28-775505.2} & 197.026167 & -77.918111 & K5$^{26}$ & II & ChaII \\ \object{ChaII~J130827.17-774323.2} & 197.113208 & -77.723111 & M4.5$^{26}$ & II & ChaII \\ \object{HD~114082} & 197.317467 & -60.308347 & F3$^{2}$ & DD & LCC \\ \object{ChaII~J130950.38-775723.9} & 197.459917 & -77.956639 & M2.5$^{26}$ & II & ChaII \\ \object{HD~131835} & 224.226958 & -35.695472 & A2$^{25}$ & DD & - \\ \object{SAO~206462} & 228.951831 & -37.154452 & F4$^{10}$ & HAeBe & - \\ \object{HIP~76310} & 233.817125 & -25.734167 & A0$^{3}$ & DD & UpSco \\ \object{HD~139614} & 235.19325 & -42.498194 & A7$^{10}$ & HAeBe & Ass Sco OB 2-3 \\ \object{HT~Lup} & 236.303625 & -34.291833 & K3$^{1}$ & II & LupusI \\ \object{HD~141569} & 237.490625 & -3.921222 & B9.5$^{11}$ & DD & MBM37 \\ \object{G327-0.6} & 238.286005 & -54.616659 & - & - & - \\ \object{HIP~77911} & 238.673333 & -22.76625 & B9$^{3}$ & DD & UpSco \\ \object{HD~142666} & 239.166763 & -22.027781 & A8$^{3}$ & HAeBe & AssScoOB2-2 \\ \object{HD~142527} & 239.174542 & -42.323139 & F6$^{22}$ & HAeBe & - \\ \object{RU~Lup} & 239.176292 & -37.820972 & G5$^{27}$ & II & LupusI \\ \object{USco~J155729.9-225843} & 239.374458 & -22.978806 & M4$^{28}$ & II & UpSco \\ \object{Sz~84} & 239.510542 & -37.60075 & M5$^{29}$ & TD & Cha \\ \object{USco~J155829.8-231007} & 239.624292 & -23.168722 & M3$^{29}$ & II & UpSco \\ \object{RY~Lup} & 239.868292 & -40.364222 & K4$^{30}$ & II & Ass Sco OB 2-3 \\ \object{1RXS~J160044.7-234330} & 240.18625 & -23.725 & M2$^{31}$ & II & UpSco \\ \object{EX~Lup} & 240.772875 & -40.307056 & - & I & Lupus III \\ \object{USco~J160357.6-203105} & 240.990292 & -20.518222 & K5$^{32}$ & II & UpSco \\ \object{USco~J160357.9-194210} & 240.991417 & -19.703028 & M2$^{28}$ & II & UpSco \\ \object{USco~J160421.7-213028} & 241.09025 & -21.507889 & K2$^{33}$ & II & UpSco \\ \object{J160532.1-193315} & 241.383958 & -19.554444 & M5$^{28}$ & III & UpSco \\ \object{USco~J160545.4-202308} & 241.439458 & -20.385583 & M2$^{28}$ & II & UpSco \\ \object{USco~J160600.6-195711} & 241.502625 & -19.953056 & M5$^{34}$ & II & UpSco \\ \object{Sco~PMS31} & 241.5915 & -19.479056 & M0.5$^{34}$ & II & UpSco \\ \object{HD~144432} & 241.7415 & -27.719389 & A9$^{25}$ & HAeBe & Ass Sco OB 2-2 \\ \object{Sz~91} & 241.798375 & -39.063083 & M0.5$^{35}$ & TD & Lupus \\ \object{UScoJ160823.2-193001} & 242.096667 & -19.500278 & K9 & II & UpSco \\ \object{HD~144668} & 242.142875 & -39.105083 & A1-A2$^{27}$ & HAeBe & LupusIII \\ \object{Sz~111} & 242.227875 & -39.628639 & M1.5$^{35}$ & TD & Lupus \\ \object{UScoJ160959.4-180009} & 242.49725 & -18.002528 & M4$^{28}$ & II & UpSco \\ \object{SST~Lup} & 242.623333 & -39.370833 & M5$^{36}$ & TD & Lupus \\ \object{AS~205} & 242.880625 & -18.640583 & K5$^{37}$ & II & UpSco \\ \object{HIP~79439} & 243.183792 & -19.502833 & B9V$^{3}$ & DD & UpSco \\ \object{UScoJ~161411.0-230536} & 243.546167 & -23.093389 & K0$^{38}$ & II & UpSco \\ \object{USco~J161420.2-190648} & 243.584583 & -19.113361 & K5$^{39}$ & II & UpSco \\ \object{RXJ~1615.3-3255} & 243.834292 & -32.918083 & K5$^{40}$ & TD & Lupus \\ \object{HIP~79878} & 244.567375 & -28.041694 & A0$^{25}$ & DD & UpSco \\ \object{HIP~80088} & 245.209292 & -22.594083 & A9$^{3}$ & DD & UpSco \\ \object{2MASS~J16230923-2417047} & 245.788458 & -24.284639 & G0$^{41}$ & TD & Oph \\ \object{Doar~21} & 246.512596 & -24.393344 & K0$^{42}$ & II & Oph \\ \object{GSS~30-IRS1} & 246.589167 & -24.384528 & - & I & Oph \\ \object{GSS~31} & 246.597417 & -24.349917 & K0$^{42}$ & II & Oph \\ \object{DoAr25} & 246.59867 & -24.72053 & K5$^{43}$ & II & Oph \\ \object{VLA~1623-243} & 246.61 & -24.408333 & - & 0 & Oph \\ \object{WL~12} & 246.684167 & -24.580111 & - & I & Oph \\ \object{DoAr~28} & 246.697583 & -23.247833 & K5$^{41}$ & TD & Oph \\ \object{WL~2} & 246.702083 & -24.477417 & - & 0/I & Oph \\ \object{Oph~01} & 246.7462 & -24.5842 & - & I & Oph \\ \object{Elias~29} & 246.789167 & -24.621833 & - & I & Oph \\ \object{SR~21} & 246.792833 & -24.320139 & G3$^{43}$ & TD & Oph \\ \object{IRS~44} & 246.8725 & -24.654472 & - & 0/I & Oph \\ \object{IRS~46} & 246.8725 & -24.654472 & - & 0/I & Oph \\ \object{IRS~48} & 246.904958 & -24.509722 & A0$^{44}$ & HAeBe & Oph \\ \object{GY~314} & 246.914292 & -24.654306 & M0$^{45}$ & - & Oph \\ \object{WSB~60} & 247.068792 & -24.616111 & M4$^{43}$ & TD & Oph \\ \object{DoAr~44} & 247.889417 & -24.460361 & K3$^{42}$ & TD & Oph \\ \object{IRS~63} & 247.898333 & -24.024806 & - & 0/I & Oph \\ \object{L1689~SNO2} & 247.967 & -24.937667 & - & I & Oph \\ \object{16289-2457} & 247.978083 & -25.056611 & - & 0/I & Oph \\ \object{2MASS~J16320099-2456419} & 248.0041 & -24.9451 & - & I & Oph \\ \object{16293-2424} & 248.08775 & -24.509944 & - & 0/I & Oph \\ \object{V346~Nor} & 248.134125 & -44.925194 & - & I & - \\ \object{RNO~90} & 248.538208 & -15.804667 & G5$^{43}$ & II & Oph \\ \object{HBC~650} & 248.62216 & -15.78372 & - & I & - \\ \object{HD~150193} & 250.074667 & -23.895889 & A2$^{10}$ & HAeBe & Ass Sco OB 2-2 \\ \object{Sco~01} & 251.7427 & -9.58883 & - & I & \\ \object{KK~Oph} & 257.533583 & -27.255056 & A6$^{10}$ & HAeBe & - \\ \object{NGC~6334-I} & 260.222037 & -35.78329 & - & - & - \\ \object{HD~158352} & 262.206895 & 0.330625 & A7$^{10}$ & DD & - \\ \object{HD~158643} & 262.853975 & -23.962643 & A0$^{3}$ & HAeBe & - \\ \object{HD~163296} & 269.0887 & -21.956078 & A1$^{8}$ & HAeBe & - \\ \object{HD~164249} & 270.764208 & -51.649011 & F5$^{46}$ & DD & BPMG \\ \object{W33A} & 273.662980 & -17.868719 & - & - & - \\ \object{[SER2000]~L483} & 274.37475 & 4.660917 & - & 0/I & - \\ \object{HD~169142} & 276.124078 & -29.780381 & A7$^{10}$ & HAeBe & - \\ \object{[MAM2011]~Aqu-MM2} & 277.2659 & -1.65041 & - & 0 & Aqu \\ \object{[MAM2011]~Aqu-MM4} & 277.2858 & -1.51188 & - & 0 & Aqu \\ \object{[MAM2011]~SerpS-MM1} & 277.407 & -1.84938 & - & 0 & Ser \\ \object{Serpens~SMM1a} & 277.4575 & 1.25572 & - & 0 & - \\ \object{EC~82} & 277.487042 & 1.24625 & - & 0/I & Serpens \\ \object{Serpens-SMM4} & 277.4882 & 1.220302 & - & 0 & Serpens \\ \object{Serpens-SMM3} & 277.497083 & 1.233806 & - & 0 & Serpens \\ \object{Serpens-SMM18} & 277.51721 & -2.05058 & - & 0 & Serpens \\ \object{[MAM2011]~Aqu-MM6} & 277.6045 & -1.90372 & - & 0 & Aqu \\ \object{[MAM2011]~Aqu-MM7} & 277.6192 & -1.94658 & - & 0 & Aqu \\ \object{[MAM2011]~Aqu-MM8} & 277.6209 & -1.93483 & - & 0 & Aqu \\ \object{[MAM2011]~Aqu-MM14} & 277.708 & -1.93502 & - & 0 & Aqu \\ \object{[MAM2011]~W40-MM3} & 277.7892 & -2.10680 & - & 0 & W40 \\ \object{[MAM2011]~W40-MM5} & 277.7931 & -2.06400 & - & 0 & W40 \\ \object{[MAM2011]~W40-MM26} & 277.9439 & -2.07291 & - & 0 & W40 \\ \object{[MAM2011]~W40-MM27} & 277.9449 & -2.03886 & - & 0 & W40 \\ \object{[MAM2011]~W40-MM28} & 277.9495 & -2.027 & - & 0 & W40 \\ \object{[MAM2011]~W40-MM34} & 277.9885 & -2.00769 & - & 0/I & W40 \\ \object{[MAM2011]~W40-MM36} & 278.0556 & -1.95822 & - & 0 & W40 \\ \object{Vega} & 279.23475 & 38.783694 & A0$^{4}$ & DD & - \\ \object{HD~172555} & 281.362083 & -64.871258 & A5$^{4}$ & DD & BPMG \\ \object{RXJ~18523-3700} & 283.072083 & -37.003306 & K7$^{47}$ & DD & - \\ \object{G34.26+0.15} & 283.328128 & 1.249232 & - & - & - \\ \object{SCra} & 285.285833 & -36.955556 & - & I & Cra \\ \object{R~CrA-IRS5A} & 285.450417 & -36.956306 & - & I & Cra \\ \object{R~CrA-IRS5N} & 285.451921 & -36.954133 & - & 0 & Cra \\ \object{R~CrA-IRS~7A} & 285.480417 & -36.954722 & - & - & Cra \\ \object{SMM~1C} & 285.48042 & -36.95464 & - & I & Cra \\ \object{R~CrA} & 285.473568 & -36.952187 & B5$^{4}$ & HAeBe & Cra \\ \object{R~CrA-IRS~7B} & 285.485 & -36.957861 & - & I & Cra \\ \object{CrA~01} & 285.7444 & -37.1266 & - & I & Cra \\ \object{HD~179218} & 287.796875 & 15.787667 & A0$^{12}$ & HAeBe & - \\ \object{LDN~723-mm} & 289.47375 & 19.2055 & - & 0 & - \\ \object{HD~181296} & 290.713375 & -54.423931 & A0$^{46}$ & DD & BPMG \\ \object{HD~181327} & 290.745583 & -54.538056 & F6$^{46}$ & DD & BPMG \\ \object{Parsamian~21} & 292.254 & 9.645000 & - & I & - \\ \object{[SER2000]~B335} & 294.2529 & 7.5689 & - & 0 & - \\ \object{HD~191089} & 302.271727 & -26.224036 & F5$^{25}$ & DD & - \\ \object{HD~192758} & 304.565792 & -42.860082 & F0$^{22}$ & DD & Argus \\ \object{AFGL~2591} & 307.353139 & 40.189365 & - & - & - \\ \object{DR~21~(OH)} & 309.753208 & 42.380500 & - & - & - \\ \object{LDN~1157-mm} & 309.7758 & 68.0375 & - & 0 & - \\ \object{AU~Mic} & 311.289708 & -31.340889 & M1$^{1}$ & DD & BPMG \\ \object{HH~381IRS} & 314.589208 & 52.490806 & - & I & - \\ \object{HD~203024} & 319.012583 & 68.914472 & A5$^{8}$ & DD & - \\ \object{L~1014} & 321.0312 & 49.98583 & - & 0 & - \\ \object{HH~354IRS} & 331.710417 & 59.046389 & - & I & - \\ \object{V733~Cep} & 343.388583 & 62.539889 & - & I & Cep OB \\ \object{Fomalhaut} & 344.412708 & -29.622222 & A4$^{4}$ & DD & - \\ \object{HR~8799} & 346.869625 & 21.13425 & F0$^{4}$ & DD & Col \\ \object{NGC~7538~IRS1} & 348.4387 & 61.46944 & - & - & - \\ \object{HD~221853} & 353.90064 & 8.382618 & F0$^{48}$ & DD & LA \\ \hline \end{longtable} } \onllongtab{2}{ \begin{longtable}{llllccc} \caption{\label{YSO_fluxes} Fluxes at 63 $\rm \mu m$ for YSOs in the sample}\\ \hline\hline Source name & RA & Dec & obs ID & $\rm F_ {[OI]}$ & $\rm F_{H {2}O}$ & $\rm F_{63\mu m}$ \\ -- & (deg) & (deg) & -- & ($\rm 10^{-18}W/m^{2}$) & ($\rm 10^{-18}W/m^{2}$) & (Jy) \\ \hline \endfirsthead \caption{continued.}\\ \hline\hline Source name & RA & Dec & obs ID & $\rm F _{[OI]}$ & $\rm F_{H {2}O}$ & $\rm F_{63\mu m}$ \\ -- & (deg) & (deg) & -- & ($\rm 10^{ -18}W/m^{2}$) & ($\rm 10^{ -18}W/m^{2}$) & (Jy) \\ \hline \endhead \hline \endfoot \hline \multicolumn{6}{l}{$\rm ^{*}$: observation was mis-pointed.}\\ \multicolumn{6}{l}{$\rm ^{**}$: source is located outside the central spaxel.}\\ \multicolumn{6}{l}{$\rm ^{***}$: the source shows absorption in the central spaxel.} \endlastfoot HD 105 & 1.468958 & -41.753 & 1342199239 & $\rm < 6$ & $\rm < 6$ & $\rm < 0.5$ \\ HD 377 & 2.107292 & 6.616806 & 1342212530 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.5$ \\ HD 3003 & 8.182958 & -63.0315 & 1342199240 & $\rm < 7$ & $\rm < 7$ & $\rm < 1.0$ \\ HD 3670 & 9.736267 & -52.5342 & 1342269870 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.5$ \\ HD 9672 & 23.65741 & 5.6763 & 1342188424 & $\rm < 11$ & $\rm < 11$ & $\rm 1.9 \pm 0.4$ \\ tau Ceti & 26.017 & 5.9374 & 1342189611 & $\rm < 8$ & $\rm < 8$ & $\rm < 0.624$ \\ W3 IRS5 & 36.41916 & 62.0975 & 1342229091 & $\rm 22026 \pm 520$ & $\rm < 632$ & $\rm 6958 \pm 20$ \\ W3 IRS5 & 36.41916 & 62.0975 & 1342229093 & $\rm 25543 \pm 3329$ & $\rm < 1000$ & $\rm 8461 \pm 61$ \\ HD 15115 & 36.567667 & 6.292556 & 1342225584 & $\rm < 7.6$ & $\rm < 7.6$ & $\rm < 0.50$ \\ HD 16743 & 39.78151 & -52.9348 & 1342264255 & $\rm < 10$ & $\rm < 10$ & $\rm < 0.57$ \\ SMM J032537+30451 & 51.343 & 30.75386 & 1342263508 & $\rm 1668 \pm 8$ & $\rm 37 \pm 6$ & $\rm 17.1 \pm 1.1$ \\ LDN 1448N & 51.40204 & 30.75616 & 1342263506 & $\rm 1609 \pm 10$ & $\rm 38 \pm 5$ & $\rm 22 \pm 1$ \\ L1448-C(S) & 51.413 & 30.73283 & 1342263510$^{*,**}$ & $\rm 545 \pm 13$ & $\rm 168 \pm 7$ & $\rm 24.4 \pm 1.2$ \\ LDN 1455 & 52.0027 & 30.133 & 1342204122 & $\rm 66 \pm 14$ & $\rm < 16$ & $\rm 1.8 \pm 0.4$ \\ IRAS 03235+3004 & 51.65612 & 30.2578 & 1342264250 & $\rm 238 \pm 11$ & $\rm < 13$ & $\rm 5.9 \pm 0.8$ \\ LDN 1455 IRS3 & 51.913644 & 30.217410 & 1342214677 & $\rm 117 \pm 16$ & $\rm < 21$ & $ 39.0 \pm 0.5 $\\ 2MASS J03283706+3113310 & 52.15454 & 31.22522 & 1342264248 & $\rm 188 \pm 9$ & $\rm < 15$ & $\rm 66 \pm 1$ \\ SSTc2d J032857.4+311416 & 52.239 & 31.23775 & 1342264247 & $\rm < 30$ & $\rm < 39$ & $\rm 60.5 \pm 1.7$ \\ SSTc2d J032900.5+311200 & 52.25229 & 31.20022 & 1342264245 & $\rm 114 \pm 9$ & $\rm < 17$ & $\rm < 2.7$ \\ 2MASS J03290149+3120208 & 52.2565 & 31.33905 & 1342264243 & $\rm 1791 \pm 11$ & $\rm 70 \pm 7$ & $\rm 20.8 \pm 0.7$ \\ 2MASS J03290773+3121575 & 52.28241 & 31.36591 & 1342267612 & $\rm 1465 \pm 8$ & $\rm 136 \pm 6$ & $\rm 59.9 \pm 0.9$ \\ NGC 1333 IRAS 4A & 52.29375 & 31.22525 & 1342216083 & $\rm 119 \pm 4$ & $\rm 32 \pm 5$ & $\rm 15.1 \pm 0.5$ \\ & 52.29375 & 31.22525 & 1342216084 & $\rm 141 \pm 20$ & $\rm < 12$ & $\rm 15.1 \pm 0.8$ \\ SSTc2d J032910.7+311821 & 52.2945 & 31.30572 & 1342267616 & $\rm < 28$ & $\rm < 15$ & $\rm 15.9 \pm 0.8$ \\ NGC 1333 IRAS 4B & 52.300042 & 31.21894 & 1342216178 & $\rm 111 \pm 23$ & $\rm < 29$ & $\rm 5.7 \pm 0.7$ \\ SSTc2d J032912.0+311301 & 52.30025 & 31.21713 & 1342267608$^{*,**}$ & $\rm 260 \pm 7$ & $\rm 258 \pm 4$ & $\rm < 2.4$ \\ SSTc2d J032913.5+311358 & 52.30641 & 31.23283 & 1342267610 & $\rm 54 \pm 11$ & $\rm < 14$ & $\rm < 3$ \\ HBC 347 & 52.40987 & 24.51055 & 1342192136 & $\rm < 14$ & $\rm < 14$ & $\rm < 1.2$ \\ IRAS 03267+3128 & 52.46591 & 31.65166 & 1342267614 & $\rm < 20$ & $\rm < 12$ & $\rm < 3.7$ \\ IRAS 03271+3013 & 52.56308 & 30.39705 & 1342263513 & $\rm 277 \pm 14$ & $\rm 48 \pm 5$ & $\rm 8.6 \pm 1.1$ \\ IRAS 03282+3035 & 52.83741 & 30.75836 & 1342263515 & $\rm 146 \pm 13$ & $\rm < 16$ & $\rm 5.1 \pm 0.8$ \\ HD 21997 & 52.97354 & -25.6141 & 1342247735 & $\rm < 11$ & $\rm < 11$ & $\rm < 0.9$ \\ IRAS 03292+3039 & 53.07483 & 30.82986 & 1342265448 & $\rm 163 \pm 9$ & $\rm < 15$ & $\rm < 2.8$ \\ epsilon Eri & 53.23266 & -9.45825 & 1342191348 & $\rm < 9$ & $\rm < 9$ & $\rm < 0.6$ \\ IRAS03301+3111 & 53.30333 & 31.35672 & 1342215668 & $\rm 387 \pm 18$ & $\rm 39 \pm 13$ & $\rm 5.1 \pm 0.6$ \\ B1-c & 53.32437 & 31.15886 & 1342216213 & $\rm < 21$ & $\rm < 21$ & $\rm 13.0 \pm 0.5$ \\ SSTc2d J033314.3+310710 & 53.30991 & 31.11969 & 1342263487 & $\rm 93 \pm 8$ & $\rm < 17$ & $\rm < 2.6$ \\ SSTc2d J033316.4+310653 & 53.3185 & 31.11458 & 1342265450 & $\rm 397 \pm 13$ & $\rm < 17$ & $\rm < 2.5$ \\ B1-a & 53.31958 & 31.132 & 1342216182 & $\rm 358 \pm 15$ & $\rm 59 \pm 19$ & $\rm 8.4 \pm 0.5$ \\ SSTc2d J033327.3+310710 & 53.3637 & 31.1195 & 1342265452 & $\rm 1065 \pm 5$ & $\rm < 15$ & $\rm < 3.6$ \\ IRAS 03407+3152 & 55.9855 & 32.01466 & 1342265454 & $\rm 231 \pm 9$ & $\rm < 14$ & $\rm 2.3 \pm 0.68$ \\ SSTc2d J034356.8+320305 & 55.98683 & 32.0513 & 1342265456 & $\rm 172 \pm 8$ & $\rm < 16$ & $\rm 4.7 \pm 1.2$ \\ 2MASS J03444389+3201373 & 56.18316 & 32.02672 & 1342265702 & $\rm 601 \pm 7$ & $\rm < 10$ & $\rm 8.4 \pm 0.7$ \\ LkHa 330 & 56.45116 & 32.4033 & 1342239377 & $\rm < 10$ & $\rm < 10$ & $\rm 12.3 \pm 0.2$ \\ & 56.45117 & 32.40331 & 1342238941 & $\rm < 57$ & $\rm < 57$ & $\rm 12.6 \pm 1.2$ \\ HBC 356 & 60.80829 & 25.8833 & 1342214359 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.0$ \\ HBC 358 & 60.96183 & 26.18144 & 1342204347 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.0$ \\ & 60.96183 & 26.18144 & 1342214680 & $\rm < 13$ & $\rm < 13$ & $\rm < 1.4$ \\ IRAS 04016+2610 & 61.179055 & 26.315706 & 1342216216 & $\rm 430 \pm 22$ & $\rm < 28$ & $\rm 42.2 \pm 0.5$ \\ & 61.17916 & 26.31583 & 1342204348 & $\rm 396 \pm 6$ & $\rm 31 \pm 5$ & $\rm 43.1 \pm 0.5$ \\ PP 13S & 62.671458 & 38.132028 & 1342263505 & $\rm 1108 \pm 41$ & $\rm < 64$ & $\rm 54 \pm 1$ \\ LkCa 1 & 63.30891 & 28.31967 & 1342214679 & $\rm < 10$ & $\rm < 10$ & $\rm < 1.4$ \\ V1096 Tau & 63.36345 & 28.27355 & 1342214678 & $\rm < 13$ & $\rm < 13$ & $\rm < 1.2$ \\ V773 Tau & 63.55383 & 28.20345 & 1342216217 & $\rm 80 \pm 4$ & $\rm < 7$ & $\rm 0.8 \pm 0.2$ \\ FM Tau & 63.55658 & 28.21366 & 1342216218 & $\rm < 12$ & $\rm < 12$ & $\rm < 1.20$ \\ CW Tau & 63.57083 & 28.18272 & 1342216221 & $\rm 84 \pm 6$ & $\rm < 12$ & $\rm 1.5 \pm 0.4$ \\ CIDA1 & 63.57337 & 28.10269 & 1342268646 & $\rm < 4$ & $\rm < 4$ & $\rm < 0.4$ \\ CX Tau & 63.69941 & 26.80305 & 1342225729 & $\rm < 9$ & $\rm < 9$ & $\rm < 1.1$ \\ LkCa 3 & 63.69987 & 27.87629 & 1342216220 & $\rm < 10$ & $\rm < 10$ & $\rm < 1.2$ \\ FO Tau & 63.70537 & 28.2085 & 1342216645 & $\rm < 14$ & $\rm < 14$ & $\rm < 1.8$ \\ CIDA-2 & 63.7715 & 28.14616 & 1342216643 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.2$ \\ LkCa 4 & 64.11712 & 28.12661 & 1342216642 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.2$ \\ CY Tau & 64.39054 & 28.34636 & 1342192794 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.0$ \\ LkCa 5 & 64.41225 & 28.55014 & 1342216641 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.3$ \\ IRAS 04158+2805 & 64.74225 & 28.20652 & 1342192793 & $\rm 55 \pm 4$ & $\rm < 11$ & $\rm < 1.3$ \\ FQ Tau & 64.80337 & 28.49252 & 1342192795 & $\rm < 10$ & $\rm < 10$ & $\rm < 1.2$ \\ BP Tau & 64.81595 & 29.10747 & 1342192796$^{*}$ & $\rm 8 \pm 2$ & $\rm < 9$ & $\rm < 1.0$ \\ & 64.81595 & 29.10747 & 1342225728 & $\rm 9 \pm 3$ & $\rm 11 \pm 3$ & $\rm < 0.8$ \\ V819 Tau & 64.85941 & 28.4373 & 1342216651 & $\rm < 12$ & $\rm < 12$ & $\rm < 1.2$ \\ FR Tau & 64.8977 & 28.45605 & 1342263516 & $\rm < 3$ & $\rm < 3$ & $\rm < 0.2$ \\ LkCa 7 & 64.92195 & 27.83013 & 1342216649 & $\rm < 12$ & $\rm < 12$ & $\rm < 1.3$ \\ IRAS 04169+2702 & 64.99333 & 27.16583 & 1342265457 & $\rm 608 \pm 7$ & $\rm < 23$ & $\rm 16.9 \pm 0.7$ \\ IRAS 04181+2654 & 65.2975 & 27.01916 & 1342265459 & $\rm 206 \pm 6$ & $\rm 24 \pm 3$ & $\rm 2.7 \pm 0.8$ \\ DE Tau & 65.48183 & 27.91836 & 1342192797 & $\rm < 13$ & $\rm < 13$ & $\rm 1.4 \pm 0.5$ \\ IRAM04191 & 65.48708 & 15.49608 & 1342216654 & $\rm 160 \pm 23$ & $\rm < 24$ & $\rm < 1.7$ \\ RY Tau & 65.48916 & 28.4432 & 1342190361 & $\rm 100 \pm 6$ & $\rm 22 \pm 7$ & $\rm 10.3 \pm 0.6$ \\ HD 283572 & 65.4952 & 28.30184 & 1342216646 & $\rm < 11$ & $\rm < 11$ & $\rm < 0.9$ \\ T Tau & 65.49762 & 19.5351 & 1342190353 & $\rm 8339 \pm 15$ & $\rm 288 \pm 14$ & $\rm 98.1 \pm 1.1$ \\ 2MASS J04220069+2657324 & 65.5025 & 26.95888 & 1342265461 & $\rm 91 \pm 8$ & $\rm 22 \pm 5$ & $\rm 3.7 \pm 0.9$ \\ FS Tau & 65.50908 & 26.95847 & 1342192791 & $\rm 378 \pm 4$ & $\rm 20 \pm 4$ & $\rm 2.4 \pm 0.5$ \\ FS Tau B & 65.509071 & 26.958469 & 1342192791 & $\rm 43 \pm 6$ & $\rm < 16$ & $\rm 2.1 \pm 0.5$ \\ FU Tau & 65.8975 & 25.05075 & 1342264241 & $\rm < 3$ & $\rm < 3$ & $\rm < 21 $ \\ FT Tau & 65.91329 & 24.93725 & 1342192790 & $\rm 20 \pm 6$ & $\rm < 12$ & $\rm < 1.2$ \\ IP Tau & 66.23783 & 27.19902 & 1342225756 & $\rm < 8$ & $\rm < 8$ & $\rm < 0.7$ \\ J1-4872 & 66.32366 & 26.29733 & 1342216653 & $\rm < 12$ & $\rm < 12$ & $\rm < 1.2$ \\ FV Tau & 66.72304 & 26.11511 & 1342239720 & $\rm < 72$ & $\rm < 72$ & $\rm 2.35 \pm 1.6$ \\ DG Tau B & 66.76066 & 26.09186 & 1342192798$^{*}$ & $\rm 130 \pm 5$ & $\rm < 10$ & $\rm 5.7 \pm 0.5$ \\ & 66.76083 & 26.09166 & 1342265463 & $\rm 515 \pm 11$ & $\rm 24 \pm 6$ & $\rm 16.0 \pm 1.0$ \\ DF Tau & 66.76283 & 25.70647 & 1342190359 & $\rm 52 \pm 7$ & $\rm < 13$ & $\rm 0 \pm 0$ \\ DG Tau & 66.76958 & 26.10452 & 1342190382$^{*,**}$ & $\rm 606 \pm 7$ & $\rm < 16$ & $\rm 5.6 \pm 0.4$ \\ & 66.76958 & 26.10452 & 1342225730 & $\rm 1534 \pm 15$ & $\rm < 16$ & $\rm 25.5 \pm 0.4$ \\ IRAS 04248+2612 & 66.98875 & 26.32166 & 1342265466 & $\rm 173 \pm 10$ & $\rm < 17$ & $\rm 5.2 \pm 1.1$ \\ IRAS 04264+2433 & 67.375 & 24.66527 & 1342265468 & $\rm 569 \pm 9$ & $\rm 35 \pm 9$ & $\rm 5.4 \pm 1.0$ \\ FW Tau & 67.37379 & 26.28144 & 1342225735 & $\rm < 10$ & $\rm < 10$ & $\rm < 1.4$ \\ DH Tau & 67.42508 & 26.54811 & 1342225734 & $\rm < 14$ & $\rm < 14$ & $\rm < 1.5$ \\ IQ Tau & 67.46483 & 26.11247 & 1342239721 & $\rm < 38$ & $\rm < 38$ & $\rm < 2.4$ \\ & 67.46483 & 26.11247 & 1342192135$^{*}$ & $\rm < 12$ & $\rm < 12$ & $\rm < 1.2$ \\ & 67.46483 & 26.11247 & 1342225733 & $\rm 16 \pm 3$ & $\rm 11 \pm 3$ & $\rm 0.7 \pm 0.3$ \\ CFHT20 & 67.49795 & 24.55216 & 1342265469 & $\rm < 4$ & $\rm < 4$ & $\rm < 2.5$ \\ UX Tau & 67.51662 & 18.23038 & 1342204350 & $\rm 38 \pm 4$ & $\rm < 10$ & $\rm 3.3 \pm 0.4$ \\ & 67.51662 & 18.23038 & 1342214357 & $\rm 35 \pm 3$ & $\rm < 10$ & $\rm 3.6 \pm 0.3$ \\ & 67.51666 & 18.23038 & 1342239724 & $\rm < 63$ & $\rm < 63$ & $\rm 4.1 \pm 1.4$ \\ FX Tau & 67.62337 & 24.44583 & 1342192800 & $\rm < 13$ & $\rm < 13$ & $\rm < 1.2$ \\ DK Tau & 67.68433 & 26.02355 & 1342192132$^{*}$ & $\rm < 12$ & $\rm < 12$ & $\rm < 1.4$ \\ & 67.68433 & 26.02355 & 1342225732 & $\rm 21 \pm 4$ & $\rm < 6$ & $\rm 0.9 \pm 0.2$ \\ ZZ Tau & 67.71408 & 24.70619 & 1342192799 & $\rm < 13$ & $\rm < 13$ & $\rm < 1.0$ \\ ZZ TauIRS & 67.71545 & 24.69652 & 1342240153 & $\rm < 61$ & $\rm < 61$ & $\rm 2.4 \pm 1.3$ \\ V927 Tau & 67.84925 & 24.18136 & 1342225763 & $\rm < 9$ & $\rm < 9$ & $\rm < 0.8$ \\ IRAS 04287+1801 & 67.892047 & 18.134481 & 1342192805 & $\rm 4619 \pm 77$ & $\rm < 89$ & $\rm 358 \pm 3$ \\ HL Tau & 67.9101 & 18.232681 & 1342190351$^{**}$ & $\rm 510 \pm 54$ & $\rm 62 \pm 10$ & $\rm 65.5 \pm 3.2$ \\ XZ Tau & 67.9145 & 18.23213 & 1342190351$^{**}$ & $\rm 242 \pm 6$ & $\rm 15 \pm 4$ & $\rm 3.1 \pm 0.3$ \\ HK Tau & 67.9607 & 24.40501 & 1342225736 & $\rm 37 \pm 3$ & $\rm < 9$ & $\rm 2.3 \pm 0.3$ \\ V710 Tau & 67.99083 & 18.35975 & 1342192804 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.0$ \\ Haro 6-13 & 68.0642 & 24.48325 & 1342239763 & $\rm < 78$ & $\rm < 78$ & $\rm 8.4 \pm 1.8$ \\ & 68.0642 & 24.48325 & 1342192128 & $\rm 64 \pm 5$ & $\rm < 11$ & $\rm 4.6 \pm 0.4$ \\ GG Tau & 68.12645 & 17.52794 & 1342192121 & $\rm 53 \pm 5$ & $\rm < 13$ & $\rm 3.2 \pm 0.5$ \\ UZ Tau & 68.17845 & 25.87538 & 1342240155 & $\rm < 74$ & $\rm < 74$ & $\rm < 4.7$ \\ & 68.1787 & 25.87572 & 1342192131$^{*,**}$ & $\rm 13 \pm 4$ & $\rm < 13$ & $\rm < 1.4$ \\ GH Tau & 68.27679 & 24.16236 & 1342192801$^{*,**}$ & $\rm < 9$ & $\rm < 9$ & $\rm < 1.0$ \\ V807 Tau & 68.27762 & 24.16528 & 1342192801 & $\rm 31 \pm 6$ & $\rm < 10$ & $\rm < 0.7$ \\ GI Tau & 68.39295 & 24.35319 & 1342225760$^{*,**}$ & $\rm < 7$ & $\rm < 7$ & $\rm < 0.6$ \\ GK Tau & 68.394 & 24.35161 & 1342239764$^{*}$ & $\rm < 50$ & $\rm < 50$ & $\rm < 3.3$ \\ DL Tau & 68.41275 & 25.34394 & 1342190355$^{*,**}$ & $\rm < 11$ & $\rm < 11$ & $\rm < 1.5$ \\ & 68.41275 & 25.34394 & 1342240154 & $\rm < 56$ & $\rm < 56$ & $\rm < 3.9$ \\ & 68.41275 & 25.34394 & 1342225800 & $\rm 27 \pm 3$ & $\rm < 6$ & $\rm 1.1 \pm 0.2$ \\ HN Tau & 68.41395 & 17.86454 & 1342225796 & $\rm 50 \pm 3$ & $\rm < 7$ & $\rm 0.9 \pm 0.2$ \\ DM Tau & 68.453 & 18.16944 & 1342192123$^{*}$ & $\rm < 10$ & $\rm < 10$ & $\rm < 1.2$ \\ & 68.453 & 18.16944 & 1342239748 & $\rm < 44$ & $\rm < 44$ & $\rm < 2.6$ \\ & 68.453 & 18.16944 & 1342225825 & $\rm 11 \pm 3$ & $\rm < 6$ & $\rm 0.8 \pm 0.2$ \\ CI Tau & 68.46666 & 22.84172 & 1342225799 & $\rm < 66$ & $\rm < 66$ & $\rm < 4.6$ \\ & 68.46666 & 22.84172 & 1342192125 & $\rm 27 \pm 4$ & $\rm < 10$ & $\rm 1.5 \pm 0.4$ \\ AA Tau & 68.73091 & 24.48144 & 1342190357$\rm ^{*}$ & $\rm < 11$ & $\rm 5 \pm 2$ & $\rm < 1.2$ \\ & 68.73091 & 24.48144 & 1342240152 & $\rm < 51$ & $\rm < 51$ & $\rm < 3.1$ \\ & 68.73091 & 24.48144 & 1342225758 & $\rm 25 \pm 2$ & $\rm 9 \pm 2$ & $\rm 1.0 \pm 0.2$ \\ HO Tau & 68.83416 & 22.53738 & 1342192803 & $\rm < 13$ & $\rm < 13$ & $\rm < 1.4$ \\ FF Tau & 68.83708 & 22.90672 & 1342192802 & $\rm < 10$ & $\rm < 10$ & $\rm < 1.6$ \\ DN Tau & 68.86404 & 24.24969 & 1342192127$^{*}$ & $\rm < 9$ & $\rm < 9$ & $\rm < 1.7$ \\ & 68.86404 & 24.24969 & 1342240151 & $\rm < 38$ & $\rm < 38$ & $\rm < 2.5$ \\ & 68.86404 & 24.24969 & 1342225757 & $\rm 4 \pm 1$ & $\rm < 6$ & $\rm 0.75 \pm 0.25$ \\ IRAS 04325+2402 & 68.89708 & 24.13861 & 1342267857 & $\rm 479 \pm 11$ & $\rm < 14$ & $\rm 11.0 \pm 0.8$ \\ HP Tau & 68.96991 & 22.90641 & 1342240150 & $\rm < 62$ & $\rm < 62$ & $\rm 4.3 \pm 1.3$ \\ J04381486 & 69.56191 & 26.19441 & 1342265470 & $\rm < 3$ & $\rm < 3$ & $\rm < 0.2$ \\ GMTau & 69.58891 & 26.1538 & 1342264239 & $\rm < 3$ & $\rm < 3$ & $\rm < 0.2$ \\ DO Tau & 69.61908 & 26.18038 & 1342240156$^{*,**}$ & $\rm < 68$ & $\rm < 68$ & $\rm 6.6 \pm 1.4$ \\ & 69.61908 & 26.1804 & 1342190385 & $\rm 63 \pm 8$ & $\rm < 13$ & $\rm < 1.54$ \\ HV Tau & 69.647 & 26.17739 & 1342225801$^{*}$ & $\rm 40 \pm 4$ & $\rm < 7$ & $\rm < 0.8$ \\ TMR 1$^{*}$ & 69.80791 & 25.88905 & 1342192985 & $\rm 401 \pm 19$ & $\rm 178 \pm 40$ & $\rm 9.1 \pm 0.5$ \\ TMR 1B & 69.80708 & 25.88916 & 1342225834$^{*,**}$ & $\rm 360 \pm 7$ & $\rm 45 \pm 6$ & $\rm 20.4 \pm 0.6$ \\ VY Tau & 69.82254 & 22.79816 & 1342192989 & $\rm < 10$ & $\rm < 10$ & $\rm < 1.3$ \\ LkCa 15 & 69.82416 & 22.35097 & 1342240149$^{*}$ & $\rm < 73$ & $\rm < 73$ & $\rm < 4.8$ \\ & 69.82416 & 22.35096 & 1342190387$^{*}$ & $\rm < 11$ & $\rm < 11$ & $\rm < 1.3$ \\ & 69.82416 & 22.35097 & 1342225798 & $\rm 12 \pm 2$ & $\rm < 6$ & $\rm 1.2 \pm 0.2$ \\ 2MASS J04393364+2359212 & 69.8902 & 23.98922 & 1342263934 & $\rm < 3$ & $\rm < 3$ & $\rm 0.12 \pm 0.07$ \\ IRAS 04365+2535 & 69.89541 & 25.69583 & 1342225832$^{**}$ & $\rm 446 \pm 13$ & $\rm < 31$ & $\rm 36.6 \pm 0.7$ \\ & 69.89583 & 25.69597 & 1342192987$^{*}$ & $\rm 531 \pm41$ & $\rm < 50$ & $\rm 32.7 \pm 0.7$ \\ BD Tau4 & 69.94783 & 26.028 & 1342264240 & $\rm < 3$ & $\rm < 3$ & $\rm < 0.2$ \\ L1527 & 69.97458 & 26.05272 & 1342192983 & $\rm 261 \pm 6$ & $\rm < 18$ & $\rm 7.0 \pm 0.4$ \\ & 69.973865 & 26.052661 & 1342192981 & $\rm 243 \pm 19$ & $\rm < 22$ & $\rm 6.8 \pm 0.5$ \\ IRAS 04381+2540 & 70.30166 & 25.77666 & 1342225830$^{*,**}$ & $\rm 111 \pm 9$ & $\rm < 21$ & $\rm 1.1 \pm 0.5$ \\ & 70.30166 & 25.77666 & 1342225803 & $\rm 923 \pm 18$ & $\rm < 18$ & $\rm 9.3 \pm 0.7$ \\ CoKuTau 4 & 70.32004 & 28.66669 & 1342191360 & $\rm 20 \pm 5$ & $\rm < 12$ & $\rm < 1.4$ \\ & 70.32004 & 28.66669 & 1342225837 & $\rm 23 \pm 1$ & $\rm < 4$ & $\rm 1.0 \pm 0.2$ \\ IRAS 04385+2550 & 70.41175 & 25.94077 & 1342240157 & $\rm < 71$ & $\rm < 71$ & $\rm < 4.6$ \\ & 70.41175 & 25.94076 & 1342225828 & $\rm 66 \pm 3$ & $\rm < 8$ & $\rm 2.6 \pm 0.3$ \\ DP Tau & 70.65708 & 25.26041 & 1342191362$^{*,**}$ & $\rm < 11$ & $\rm < 11$ & $\rm < 1.1$ \\ GO Tau & 70.76287 & 25.33855 & 1342191361 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.2$ \\ & 70.76287 & 25.33855 & 1342225826 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.9$ \\ DQ Tau & 71.72104 & 17.00005 & 1342225806 & $\rm 25 \pm 5$ & $\rm < 13$ & $\rm 1.4 \pm 0.5$ \\ Haro 6-37 & 71.74575 & 17.04394 & 1342225805 & $\rm < 15$ & $\rm < 15$ & $\rm < 1.6$ \\ DS Tau & 71.95045 & 29.42068 & 1342225851$^{*,**}$ & $\rm < 6$ & $\rm < 6$ & $\rm < 0.7$ \\ UY Aur & 72.94741 & 30.78708 & 1342193206$^{*,**}$ & $\rm < 11$ & $\rm < 11$ & $\rm < 1.1$ \\ & 72.94741 & 30.78708 & 1342215699 & $\rm 338 \pm 6$ & $\rm 13 \pm 3$ & $\rm 5.6 \pm 0.4$ \\ GM Aur & 73.79575 & 30.36652 & 1342243657 & $\rm < 80$ & $\rm < 80$ & $\rm < 5.7$ \\ & 73.79579 & 30.36645 & 1342191357$^{*}$ & $\rm 23 \pm 4$ & $\rm < 12$ & $\rm 1.7 \pm 0.5$ \\ AB Aur & 73.94095 & 30.55121 & 1342191355 & $\rm 667 \pm 20$ & $\rm < 55$ & $\rm 94.3 \pm 0.6$ \\ & 73.941 & 30.55119 & 1342217842$^{*}$ & $\rm 359 \pm 27$ & $\rm < 34$ & $\rm 50.3 \pm 0.8$ \\ SU Aur & 73.99741 & 30.5671 & 1342217844 & $\rm 86 \pm 3$ & $\rm < 8$ & $\rm 6.0 \pm 0.4$ \\ HD 31648 & 74.69277 & 29.84361 & 1342226002 & $\rm 101 \pm 3$ & $\rm 10 \pm 2$ & $\rm 12.4 \pm 0.2$ \\ HD 32297 & 75.61431 & 7.461022 & 1342217849 & $\rm < 9$ & $\rm < 9$ & $\rm 1.2 \pm 0.3$ \\ V836 Tau & 75.7775 & 25.3888 & 1342227634 & $\rm < 7$ & $\rm < 7$ & $\rm < 1.1$ \\ RW Aur & 76.95641 & 30.4014 & 1342191359 & $\rm 158 \pm 9$ & $\rm < 14$ & $\rm 1.9 \pm 0.4$ \\ HD 35187 & 81.00487 & 24.96044 & 1342217846 & $\rm < 26$ & $\rm < 26$ & $\rm 4.4 \pm 0.6$ \\ & 81.00488 & 24.96043 & 1342226900 & $\rm 27 \pm 3$ & $\rm < 8$ & $\rm 4.5 \pm 0.3$ \\ HD 35841 & 81.65244 & -22.4899 & 1342270644 & $\rm < 12$ & $\rm < 12$ & $\rm < 0.8$ \\ HD 36112 & 82.6147 & 25.33252 & 1342227635 & $\rm 39 \pm 3$ & $\rm < 9$ & $\rm 18.5 \pm 0.4$ \\ & 82.6147 & 25.33252 & 1342228247 & $\rm < 19$ & $\rm < 19$ & $\rm 16.9 \pm 0.4$ \\ HD 36910 & 83.99361 & 24.74835 & 1342227638 & $\rm 49 \pm 3$ & $\rm < 7$ & $\rm 19.8 \pm 0.2$ \\ V833 Ori & 84.57541 & -7.04061 & 1342265949 & $\rm < 120$ & $\rm < 120$ & $\rm 155 \pm 2.8$ \\ HD 245906 & 84.877 & 26.33197 & 1342228528 & $\rm 88 \pm 24$ & $\rm < 29$ & $\rm < 1.0$ \\ $\rm [SMZ2000]$ L1643-S3 MMS 1 & 84.98291 & -7.50777 & 1342226195 & $\rm 1554 \pm 8$ & $\rm 82 \pm 17$ & $\rm 364 \pm 1$ \\ Re50 NN IRS & 85.114167 & -7.458667 & 1342265946 & $\rm 1054 \pm 32$ & $\rm < 65$ & $\rm 104 \pm 2$ \\ V1647 Ori & 86.55475 & -0.101333 & 1342267864 & $\rm 139 \pm 32$ & $\rm < 570$ & $\rm 29 \pm 1$ \\ HD 38120 & 85.79954 & -4.99719 & 1342226212 & $\rm 66 \pm 16$ & $\rm < 20$ & $\rm 9.2 \pm 0.5$ \\ HD 38207 & 85.83733 & -20.1892 & 1342270645 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.5$ \\ HD 38206 & 85.84029 & 8.5574 & 1342270646 & $\rm < 11$ & $\rm < 11$ & $\rm < 0.8$ \\ HR 1998 & 86.73892 & 4.8219 & 1342226192 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.9$ \\ NGC 2071 IR & 86.76833 & 0.363611 & 1342218761$^{*}$ & $\rm 7996 \pm 39$ & $\rm 517 \pm 23$ & $\rm 402 \pm 2$ \\ $\beta$ Pic & 86.8212 & -51.0665 & 1342188425 & $\rm 11 \pm 3$ & $\rm < 10$ & $\rm 9.9 \pm 0.2$ \\ R Mon & 99.79145 & 8.736028 & 1342250903 & $\rm 2370 \pm 82$ & $\rm < 97$ & $\rm 96.6 \pm 1.9$ \\ HD 50138 & 102.8891 & -6.9665 & 1342206991 & $\rm 2103 \pm 16$ & $\rm < 23$ & $\rm 8.5 \pm 0.7$ \\ PDS 27 & 109.8997 & 7.655 & 1342251034 & $\rm < 74$ & $\rm < 74$ & $\rm 55 \pm 1.0$ \\ HD 61005 & 113.94775 & -32.203889 & 1342230911 & $\rm < 8.2$ & $\rm < 8.2$ & $\rm 0.4 \pm 0.1$ \\ Bran 76 & 117.648333 & -33.106639 & 1342254930 & $\rm < 61$ & $\rm < 61$ & $\rm < 4.2$ \\ RECX 1 & 129.2342 & -78.9459 & 1342210391 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.5$ \\ RECX 14 & 130.3762 & -78.8851 & 1342210390 & $\rm < 6$ & $\rm < 6$ & $\rm < 0.8$ \\ RECX 3 & 130.4042 & -79.0584 & 1342210389 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.1$ \\ RECX 4 & 130.5988 & -79.0675 & 1342199241 & $\rm < 5$ & $\rm < 5$ & $\rm < 0.7$ \\ RECX 5 & 130.6129 & -78.9633 & 1342210392 & $\rm < 9$ & $\rm < 9$ & $\rm < 1.0$ \\ RECX 6 & 130.6616 & -78.9118 & 1342223114 & $\rm < 9$ & $\rm < 9$ & $\rm < 1.1$ \\ RECX 8 & 130.8009 & -79.07 & 1342223113 & $\rm < 10$ & $\rm < 10$ & $\rm < 1.1$ \\ RECX 15 & 130.8274 & -79.0883 & 1342186314 & $\rm 30 \pm 6$ & $\rm < 10$ & $\rm < 0.6$ \\ RECX 15 & 130.8274 & -79.0883 & 1342210388 & $\rm 24 \pm 2$ & $\rm < 6$ & $\rm < 1.2$ \\ J0844.2-7833 & 131.0381 & -78.5626 & 1342210387 & $\rm < 7$ & $\rm < 7$ & $\rm < 1.0$ \\ RECX 9 & 131.0682 & -78.9855 & 1342223112 & $\rm < 8$ & $\rm < 8$ & $\rm < 0.6$ \\ RECX 10 & 131.1328 & -78.7753 & 1342223111 & $\rm < 13$ & $\rm < 13$ & $\rm < 1.3$ \\ $\rm [MGL99]$ IRS 17 57 & 131.6445 & -43.9084 & 1342211844 & $\rm 8156 \pm 12$ & $\rm 84 \pm 18$ & $\rm 206 \pm 1$ \\ RECX 11 & 131.7569 & -78.9929 & 1342223115 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.7$ \\ RECX 12 & 131.9865 & -78.9147 & 1342223110 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.2$ \\ TWA 07 & 160.6254 & -33.6711 & 1342199411 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.8$ \\ Sz Cha & 164.5698 & -77.288 & 1342225560 & $\rm < 76$ & $\rm < 76$ & $\rm < 5.0$ \\ & 164.5698 & -77.288 & 1342233478 & $\rm 15 \pm 2$ & $\rm < 6$ & $\rm 3.8 \pm 0.1$ \\ CR Cha & 164.7791 & -77.0278 & 1342232614 & $\rm 16 \pm 3$ & $\rm < 6$ & $\rm 1.6 \pm 0.1$ \\ TWA 01 & 165.4662 & -34.7047 & 1342248544 & $\rm 32 \pm 5$ & $\rm < 15$ & $\rm 3.4 \pm 0.3$ \\ & 165.4663 & -34.7047 & 1342187127 & $\rm 35 \pm 4$ & $\rm < 10$ & $\rm 2.9 \pm 0.3$ \\ CS Cha & 165.6037 & -77.5599 & 1342233479 & $\rm < 65$ & $\rm < 65$ & $\rm < 4.2$ \\ & 165.6037 & -77.5599 & 1342233480 & $\rm 16 \pm 2$ & $\rm < 5$ & $\rm 3.5 \pm 0.1$ \\ CHX 7 & 166.5642 & -77.3657 & 1342232584 & $\rm < 61$ & $\rm < 61$ & $\rm < 4.0$ \\ & 166.5642 & -77.3658 & 1342233477$^{**}$ & $\rm 15 \pm 2$ & $\rm < 7$ & $\rm 0.31 \pm 0.09$ \\ Eso H alpha 559 & 166.6064 & -76.5616 & 1342263489 & $\rm < 4$ & $\rm < 4$ & $\rm < 0.3$ \\ 2MASS J11064658-7722325 & 166.694084 & -77.37571 & 1342210187 & $\rm 0 \pm 0$ & $\rm 0 \pm 0$ & $\rm 0 \pm 0$ \\ \object{[NC98]~Cha~HA~1} & 166.8195 & -77.598139 & 1342263459 & $\rm < 3$ & $\rm < 3$ & $\rm < 0.3$ \\ Sz 18 & 166.8297 & -76.0513 & 1342232290 & $\rm < 32$ & $\rm < 32$ & $\rm < 2.1$ \\ & 166.8297 & -76.0513 & 1342232585 & $\rm < 6$ & $\rm < 6$ & $\rm 0.55 \pm 0.1$ \\ HD 97048 & 167.0138 & -77.6548 & 1342199412$^{*}$ & $\rm 853 \pm 18$ & $\rm < 26$ & $\rm 42.9 \pm 0.7$ \\ & 167.0139 & -77.6548 & 1342188436 & $\rm 1481 \pm 4$ & $\rm < 11$ & $\rm 64.1 \pm 0.4$ \\ HP Cha & 167.0629 & -77.5647 & 1342233473 & $\rm < 61$ & $\rm < 61$ & $\rm 8.9 \pm 1.3$ \\ Sz27 & 167.1627 & -77.2678 & 1342233476 & $\rm 10 \pm 1$ & $\rm < 4$ & $\rm 0.4 \pm 0.1$ \\ TWA 02AB & 167.3075 & -30.0277 & 1342199410 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.9$ \\ GM Cha & 167.3687 & -76.5578 & 1342267619 & $\rm 221 \pm 13$ & $\rm < 17$ & $\rm 11 \pm 1$ \\ T 42 & 167.4725 & -76.5737 & 1342232291 & $\rm 493 \pm 59$ & $\rm < 74$ & $\rm 21 \pm 2$ \\ WW Cha & 167.5004 & -76.5827 & 1342232292 & $\rm 422 \pm 55$ & $\rm < 70$ & $\rm 35 \pm 2$ \\ TWA 03A & 167.6161 & -37.5311 & 1342209871 & $\rm < 6$ & $\rm < 6$ & $\rm 0.8 \pm 0.2$ \\ Hn13 & 167.7332 & -76.759 & 1342263492 & $\rm 4 \pm 1$ & $\rm < 3$ & $\rm < 0.20$ \\ CV Cha & 168.1155 & -76.7395 & 1342232293 & $\rm < 78$ & $\rm < 78$ & $\rm < 5.4$ \\ CHX22 & 168.1778 & -77.373 & 1342233474 & $\rm 24 \pm 1$ & $\rm < 4$ & $\rm 0.33 \pm 0.09$ \\ Sz 45 & 169.4042 & -77.0772 & 1342231724 & $\rm < 31$ & $\rm < 31$ & $\rm < 2.1$ \\ & 169.4042 & -77.0772 & 1342233475 & $\rm < 5$ & $\rm < 5$ & $\rm 0.9 \pm 0.1$ \\ TWA 13AB & 170.3218 & -34.7793 & 1342210382 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.6$ \\ HD 98800B & 170.522 & -24.7777 & 1342223821 & $\rm < 7$ & $\rm < 7$ & $\rm 7.6 \pm 0.3$ \\ & 170.522 & -24.7775 & 1342199409 & $\rm 8 \pm 2$ & $\rm < 5$ & $\rm 6.2 \pm 0.2$ \\ HD 98922 & 170.6319 & -53.3698 & 1342210385 & $\rm 197 \pm 16$ & $\rm < 20$ & $\rm 4.1 \pm 0.5$ \\ HD 100453 & 173.2732 & -54.3245 & 1342203059 & $\rm 58 \pm 5$ & $\rm < 13$ & $\rm 37.9 \pm 0.9$ \\ & 173.2732 & -54.3245 & 1342212228 & $\rm 34 \pm 3$ & $\rm < 9$ & $\rm 28.6 \pm 0.3$ \\ & 173.2732 & -54.3245 & 1342211695 & $\rm 99 \pm 32$ & $\rm < 21$ & $\rm 32.8 \pm 0.50$ \\ HD 100546 & 173.356 & -70.1947 & 1342188038 & $\rm 5462 \pm 72$ & $\rm < 88$ & $\rm 180.2 \pm 0.7$ \\ & 173.356 & -70.1947 & 1342188438 & $\rm 5645 \pm 8$ & $\rm < 22$ & $\rm 180 \pm 2$ \\ T Cha & 179.3063 & -79.3587 & 1342232294 & $\rm 48 \pm 4$ & $\rm < 12$ & $\rm 6.7 \pm 0.2$ \\ HD 104237 & 180.0211 & -78.1929 & 1342207819 & $\rm 85 \pm 17$ & $\rm < 23$ & $\rm 10.6 \pm 0.5$ \\ & 180.0211 & -78.1929 & 1342212234 & $\rm 80 \pm 4$ & $\rm < 11$ & $\rm 10.5 \pm 0.3$ \\ IRAS 11590-6452 & 180.400965 & -65.147898 & 1342212230 & $\rm 1354 \pm 41$ & $\rm < 46$ & $\rm 53 \pm 1$ \\ & 180.4012 & -65.148 & 1342212232$^{*}$ & $\rm 1340 \pm 34$ & $\rm 80 \pm 21$ & $\rm 52 \pm 1$ \\ TWA 23 & 181.864 & -32.7834 & 1342213143 & $\rm < 9$ & $\rm < 9$ & $\rm < 1.0$ \\ MML 17 & 185.638458 & -53.5636 & 1342226186 & $\rm < 8.4$ & $\rm < 8.4$ & $\rm < 0.5$ \\ TWA 10 & 188.7677 & -41.6107 & 1342203443 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.9$ \\ HR 4796A & 189.0042 & -39.8695 & 1342199242 & $\rm < 7$ & $\rm < 7$ & $\rm 6.7 \pm 0.2$ \\ DK Cha & 193.3217 & -77.1196 & 1342226006 & $\rm 2716 \pm 78$ & $\rm < 94$ & $\rm 125 \pm 2$ \\ & 193.3216 & -77.1196 & 1342188039$^{*}$ & $\rm 1330 \pm 119$ & $\rm < 98$ & $\rm 37 \pm 3$ \\ ChaII-J125342.86-771511.5 & 193.4285 & -77.2531 & 1342226005 & $\rm 73 \pm 3$ & $\rm < 8$ & $\rm 2.5 \pm 0.3$ \\ ChaII-J125633.66-764545.3 & 194.1402 & -76.7625 & 1342226009 & $\rm < 8$ & $\rm < 8$ & $\rm < 1.0$ \\ ChaII-J125711.77-764011.3 & 194.299 & -76.6698 & 1342226010 & $\rm < 9$ & $\rm < 9$ & $\rm < 1.0$ \\ ChaII-J125806.78-770909.4 & 194.5282 & -77.1526 & 1342226007 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.8$ \\ 2MASS J12590656-7707401 & 194.7774 & -77.1277 & 1342265694 & $\rm 139 \pm 9$ & $\rm 10 \pm 3$ & $\rm 9.8 \pm 0.7$ \\ ChaII-J130055.36-771022.1 & 195.2306 & -77.1728 & 1342226008 & $\rm < 6$ & $\rm < 6$ & $\rm 0.7 \pm 0.2$ \\ ChaII-J130158.94-775121.7 & 195.4955 & -77.856 & 1342229799 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.8$ \\ ChaII-J130222.85-773449.3 & 195.5952 & -77.5803 & 1342227072 & $\rm < 6$ & $\rm < 6$ & $\rm < 0.7$ \\ ChaII-J130424.92-775230.1 & 196.1038 & -77.875 & 1342229800 & $\rm 10 \pm 3$ & $\rm < 8$ & $\rm < 0.7$ \\ Hn24 & 196.2322 & -77.6637 & 1342235656$^{*}$ & $\rm 9 \pm 2$ & $\rm < 5$ & $\rm < 0.4$ \\ ChaII-J130508.53-773342.4 & 196.2855 & -77.5617 & 1342227071 & $\rm < 8$ & $\rm < 8$ & $\rm < 0.8$ \\ ChaII-J130512.69-773052.3 & 196.3028 & -77.5145 & 1342229829 & $\rm < 7$ & $\rm < 7$ & $\rm < 1.0$ \\ ChaII-J130520.68-773901.4 & 196.3361 & -77.6503 & 1342229831 & $\rm 15 \pm 5$ & $\rm < 8$ & $\rm < 1.3$ \\ ChaII-J130521.66-773810.0 & 196.3402 & -77.6361 & 1342229830 & $\rm 6 \pm 2$ & $\rm < 7$ & $\rm < 0.8$ \\ ChaII-J130529.04-774140.1 & 196.371 & -77.6944 & 1342229832 & $\rm < 8$ & $\rm < 8$ & $\rm < 0.7$ \\ ChaII-J130718.05-774052.9 & 196.8252 & -77.6813 & 1342229833 & $\rm < 8$ & $\rm < 8$ & $\rm < 0.7$ \\ ChaII-J130748.51-774121.4 & 196.9521 & -77.6892 & 1342228418 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.8$ \\ ChaII-J130806.28-775505.2 & 197.0261 & -77.9181 & 1342229801 & $\rm 15 \pm 4$ & $\rm 6 \pm 2$ & $\rm < 0.8$ \\ ChaII-J130827.17-774323.2 & 197.1132 & -77.7231 & 1342228419 & $\rm < 8$ & $\rm < 8$ & $\rm < 0.9$ \\ HD 114082 & 197.3174 & -60.3083 & 1342265693 & $\rm < 8$ & $\rm < 8$ & $\rm 0.5 \pm 0.2$ \\ ChaII-J130950.38-775723.9 & 197.4599 & -77.9566 & 1342228420 & $\rm < 6$ & $\rm < 6$ & $\rm < 0.7$ \\ HD 131835 & 224.2269 & -35.6954 & 1342248686 & $\rm < 12$ & $\rm < 12$ & $\rm < 1.2$ \\ SAO 206462 & 228.9518 & -37.1544 & 1342190370 & $\rm 43 \pm 5$ & $\rm < 13$ & $\rm 33.0 \pm 0.6$ \\ & 228.9518 & -37.1544 & 1342213921 & $\rm < 26$ & $\rm < 26$ & $\rm 28.3 \pm 0.4$ \\ HIP 76310 & 233.8171 & -25.7341 & 1342191303 & $\rm < 13$ & $\rm < 13$ & $\rm < 1.2$ \\ HD 139614 & 235.1932 & -42.4981 & 1342191300 & $\rm 41 \pm 7$ & $\rm < 17$ & $\rm 19.8 \pm 0.5$ \\ & 235.1932 & -42.4981 & 1342215683 & $\rm 18 \pm 6$ & $\rm < 20$ & $\rm 20.0 \pm 0.4$ \\ HT Lup & 236.3036 & -34.2918 & 1342213920 & $\rm < 19$ & $\rm < 19$ & $\rm 2.2 \pm 0.5$ \\ HD 141569 & 237.4906 & -3.92121 & 1342190376 & $\rm 191 \pm 4$ & $\rm < 11$ & $\rm 3.4 \pm 0.4$ \\ & 237.4906 & -3.92122 & 1342213913 & $\rm 203 \pm 19$ & $\rm < 20$ & $\rm 4.1 \pm 0.5$ \\ G327-0.6$^{***}$ & 238.286005 & -54.616659 & 1342216202 & $\rm -3105 \pm 136$ & $\rm < 335$ & $\rm 882 \pm 8$ \\ HIP 77911 & 238.6733 & -22.7662 & 1342214223 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.6$ \\ HD 142666 & 239.1667 & -22.0277 & 1342213916$^{*}$ & $\rm < 16$ & $\rm < 16$ & $\rm 0 \pm 0$ \\ & 239.1667 & -22.0277 & 1342214224$^{**}$ & $\rm 13 \pm 2$ & $\rm < 8$ & $\rm 0.9 \pm 0.3$ \\ HD 142527 & 239.1745 & -42.3231 & 1342216173 & $\rm 49 \pm 3$ & $\rm < 8$ & $\rm 108.4 \pm 0.3$ \\ & 239.1745 & -42.3231 & 1342216174 & $\rm < 36$ & $\rm < 36$ & $\rm 97.6 \pm 0.8$ \\ RU Lup & 239.1762 & -37.8209 & 1342215682 & $\rm 185 \pm 20$ & $\rm < 19$ & $\rm 5.7 \pm 0.5$ \\ USco J155729.9-225843 & 239.3744 & -22.9788 & 1342214222 & $\rm < 6$ & $\rm < 6$ & $\rm < 0.7$ \\ Sz84 & 239.5105 & -37.6007 & 1342229826 & $\rm < 5$ & $\rm < 5$ & $\rm 0.4 \pm 0.1$ \\ USco J155829.8-231007 & 239.6242 & -23.1687 & 1342203460 & $\rm < 10$ & $\rm < 10$ & $\rm < 1.5$ \\ RY Lup & 239.8682 & -40.3642 & 1342216171 & $\rm < 20$ & $\rm < 20$ & $\rm 6.18 \pm 0.5$ \\ 1RXSJ160044.7-234330 & 240.1862 & -23.725 & 1342213760 & $\rm < 10$ & $\rm < 10$ & $\rm < 1.1$ \\ EX Lup & 240.772875 & -40.30706 & 1342266967 & $\rm < 53$ & $\rm < 53$ & $\rm < 4.5$ \\ USco J160357.6-203105 & 240.9902 & -20.5182 & 1342214227 & $\rm < 9$ & $\rm < 9$ & $\rm < 0.7$ \\ USco J160357.9-194210 & 240.9914 & 9.703 & 1342214226 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.7$ \\ USco J160421.7-213028 & 241.0902 & -21.5078 & 1342215681 & $\rm 25 \pm 2$ & $\rm < 7$ & $\rm 1.7 \pm 0.2$ \\ J160532.1-193315 & 241.384 & -19.5544 & 1342215637 & $\rm < 11$ & $\rm < 11$ & $\rm < 1.0$ \\ USco J160545.4-202308 & 241.4394 & -20.3855 & 1342216197 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.6$ \\ USco J160600.6-195711 & 241.5026 & 9.953 & 1342215736 & $\rm < 7$ & $\rm < 7$ & $\rm < 1.0$ \\ ScoPMS 31 & 241.5915 & 9.479 & 1342216196 & $\rm < 9$ & $\rm < 9$ & $\rm < 0.8$ \\ HD 144432 & 241.7415 & -27.7193 & 1342213919 & $\rm < 18$ & $\rm < 18$ & $\rm 5.59 \pm 0.5$ \\ Sz 91 & 241.7983 & -39.063 & 1342229827 & $\rm 8 \pm 2$ & $\rm < 5$ & $\rm 0.7 \pm 0.1$ \\ USco J160823.2-193001 & 242.0966 & 9.5002 & 1342216193 & $\rm < 8$ & $\rm < 8$ & $\rm < 1.1$ \\ USco J160823.2-193001 & 242.0966 & 9.5002 & 1342216194 & $\rm < 9$ & $\rm < 9$ & $\rm < 1.1$ \\ HD 144668 & 242.1428 & -39.105 & 1342192146 & $\rm 128 \pm 6$ & $\rm < 15$ & $\rm 6.2 \pm 0.5$ \\ & 242.1428 & -39.105 & 1342215641 & $\rm 132 \pm 14$ & $\rm < 19$ & $\rm 6.1 \pm 0.4$ \\ Sz111 & 242.2278 & -39.6286 & 1342229828 & $\rm 10 \pm 1$ & $\rm < 4$ & $\rm 1.39 \pm 0.09$ \\ USco J160959.4-180009 & 242.4972 & 8.0025 & 1342216188 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.9$ \\ SSTLup & 242.6233 & -39.3708 & 1342241709 & $\rm < 5$ & $\rm < 5$ & $\rm < 0.4$ \\ AS205 & 242.8806 & 8.6405 & 1342215737 & $\rm 222 \pm 21$ & $\rm < 24$ & $\rm 19.8 \pm 0.6$ \\ HIP 79439 & 243.1837 & 9.5028 & 1342216190 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.6$ \\ USco J161411.0-230536 & 243.5461 & -23.0933 & 1342216169 & $\rm < 8$ & $\rm < 8$ & $\rm < 0.7$ \\ USco J161420.2-190648 & 243.5845 & 9.1133 & 1342216191 & $\rm 43 \pm 4$ & $\rm < 8$ & $\rm 1.1 \pm 0.3$ \\ RXJ1615.3-3255 & 243.8342 & -32.918 & 1342229825 & $\rm 17 \pm 2$ & $\rm < 5$ & $\rm 1.3 \pm 0.1$ \\ HIP 79878 & 244.5673 & -28.0416 & 1342216170 & $\rm < 6$ & $\rm < 6$ & $\rm < 0.7$ \\ HIP 80088 & 245.2092 & -22.594 & 1342216168 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.7$ \\ 2MASS J16230923-2417047 & 245.7884 & -24.2846 & 1342250127 & $\rm 19 \pm 2$ & $\rm < 4$ & $\rm 2.6 \pm 0.1$ \\ Doar 21 & 246.5125 & -24.3933 & 1342240163 & $\rm 385 \pm 56$ & $\rm < 60$ & $\rm 8.6 \pm 1.3$ \\ GSS30-IRS1 & 246.5891 & -24.3845 & 1342215678 & $\rm 1836 \pm 51$ & $\rm 334 \pm 53$ & $\rm 148 \pm 1$ \\ GSS 31 & 246.5974 & -24.3499 & 1342240164 & $\rm < 94$ & $\rm < 94$ & $\rm 5.2 \pm 2.1$ \\ DoAr 25 & 246.5986 & -24.7205 & 1342241708 & $\rm < 68$ & $\rm < 68$ & $\rm < 4.6$ \\ VLA1623-243 & 246.61 & -24.4083 & 1342213918 & $\rm 399 \pm 19$ & $\rm < 20$ & $\rm 13.1 \pm 0.5$ \\ WL12 & 246.6841 & -24.5801 & 1342228187 & $\rm 494 \pm 18$ & $\rm 50 \pm 15$ & $\rm 13.6 \pm 0.5$ \\ DoAr 28 & 246.6975 & -23.2478 & 1342241707 & $\rm 10 \pm 1$ & $\rm < 5$ & $\rm 0.9 \pm 0.1$ \\ & 246.6978 & -23.2485 & 1342229823 & $\rm < 30$ & $\rm < 30$ & $\rm < 2.0$ \\ WL 2 & 246.702 & -24.4774 & 1342242631 & $\rm < 60$ & $\rm < 60$ & $\rm < 4.1$ \\ Oph 01 & 246.7462 & -24.5842 & 1342266925 & $\rm 3109 \pm 19$ & $\rm < 12$ & $\rm 28.2 \pm 0.6$ \\ Elias 29 & 246.7891 & -24.6218 & 1342228519 & $\rm 1840 \pm 69$ & $\rm 511 \pm 71$ & $\rm 141 \pm 1$ \\ SR 21 & 246.7928 & -24.3201 & 1342229824 & $\rm 13 \pm 4$ & $\rm < 13$ & $\rm 36.8 \pm 0.3$ \\ & 246.7928 & -24.3201 & 1342227209 & $\rm < 21$ & $\rm < 21$ & $\rm 33.9 \pm 0.5$ \\ IRS 44 & 246.8725 & -24.6544 & 1342228474 & $\rm 436 \pm 37$ & $\rm < 37$ & $\rm 17.9 \pm 0.9$ \\ IRS 46 & 246.8725 & -24.6544 & 1342228474 & $\rm 76 \pm 23$ & $\rm < 22$ & $\rm 4.5 \pm 0.4$ \\ IRS 48 & 246.9049 & -24.5097 & 1342227069 & $\rm 322 \pm 19$ & $\rm < 25$ & $\rm 44.8 \pm 0.5$ \\ GY 314 & 246.9142 & -24.6543 & 1342242632 & $\rm 113 \pm 35$ & $\rm < 61$ & $\rm 3.1 \pm 1.3$ \\ WSB 60 & 247.0687 & -24.6161 & 1342242633 & $\rm < 36$ & $\rm < 36$ & $\rm < 2.4$ \\ & 247.0687 & -24.6161 & 1342250128 & $\rm < 6$ & $\rm < 6$ & $\rm 0.8 \pm 0.1$ \\ DoAr 44 & 247.8893 & -24.4603 & 1342241269 & $\rm < 68$ & $\rm < 68$ & $\rm 5.2 \pm 2.0$ \\ & 247.8894 & -24.4603 & 1342250578 & $\rm 25 \pm 2$ & $\rm 8 \pm 1$ & $\rm 4.86 \pm 0.08$ \\ IRS 63 & 247.8983 & -24.0248 & 1342228473 & $\rm 205 \pm 18$ & $\rm < 20$ & $\rm 16.8 \pm 0.57$ \\ L1689S NO2 & 247.967 & -24.9376 & 1342241270$^{*}$ & $\rm 235 \pm 70$ & $\rm < 73$ & $\rm 4.3 \pm 1.6$ \\ 16289-2457 & 247.978 & -25.0566 & 1342242634 & $\rm < 78$ & $\rm < 78$ & $\rm < 5.1$ \\ 2MASS J16320099-2456419 & 248.0041 & -24.9451 & 1342263469 & $\rm 1092 \pm 12$ & $\rm 48 \pm 7$ & $\rm 87.4 \pm 0.9$ \\ 16293-2424 & 248.0877 & -24.5099 & 1342241268 & $\rm < 72$ & $\rm < 72$ & $\rm < 4.5$ \\ V346 Nor & 248.134125 & -44.925194 & 1342267622 & $\rm 382 \pm 36$ & $\rm < 69$ & $\rm 96.4 \pm 1.6$ \\ RNO 90 & 248.5382 & 5.8046 & 1342228206 & $\rm 128 \pm 20$ & $\rm < 18$ & $\rm 3.5 \pm 0.5$ \\ HBC 650 & 248.6221 & -15.7837 & 1342215639 & $\rm 811 \pm 6$ & $\rm < 17$ & $\rm 27.2 \pm 0.6$ \\ HD 150193 & 250.0746 & -23.8958 & 1342227068 & $\rm 18 \pm 5 $ & $\rm < 19$ & $\rm 7.6 \pm 0.5$ \\ & 250.0746 & -23.8958 & 1342216625 & $\rm 23 \pm 3$ & $\rm < 9$ & $\rm 6.8 \pm 0.2$ \\ Sco01 & 251.7427 & -9.58883 & 1342267176 & $\rm 223 \pm 9$ & $\rm < 17$ & $\rm 8.5 \pm 0.6$ \\ KK Oph & 257.5335 & -27.255 & 1342192148 & $\rm 168 \pm 6$ & $\rm < 15$ & $\rm 4.6 \pm 0.4$ \\ NGC 6334-I$^{***}$ & 260.222037 & -35.78329 & 1342239385 & $\rm 5313 \pm 151$ & $\rm < 742$ & $\rm 6407 \pm 17$ \\ HD 158352 & 262.2068 & 0.330625 & 1342190377 & $\rm < 11$ & $\rm < 11$ & $\rm < 0.9$ \\ HD 158643 & 262.8539 & -23.9626 & 1342217821 & $\rm 49 \pm 3$ & $\rm < 7$ & $\rm 1.1 \pm 0.3$ \\ HD 163296 & 269.0887 & -21.956 & 1342192161 & $\rm 206 \pm 5$ & $\rm < 14$ & $\rm 19.2 \pm 0.6$ \\ & 269.0887 & -21.956 & 1342243512 & $\rm 197 \pm 6$ & $\rm 14 \pm 4$ & $\rm 18.2 \pm 0.3$ \\ & 269.0887 & -21.956 & 1342217819 & $\rm 211 \pm 18$ & $\rm < 21$ & $\rm 18.1 \pm 0.5$ \\ HD 164249 & 270.7642 & -51.649 & 1342215648 & $\rm < 8$ & $\rm < 8$ & $\rm < 1.0$ \\ W33A$^{**,***}$ & 273.662980 & -17.868719 & 1342239713 & $\rm -3302 \pm 145$ & $\rm < 429$ & $\rm 1155 \pm 15$ \\ $\rm [SER2000]$ L483 & 274.3745 & -4.66097 & 1342192156 & $\rm 385 \pm 6$ & $\rm 58 \pm 10$ & $\rm 67.6 \pm 0.6$ \\ HD 169142 & 276.124 & -29.7803 & 1342186310 & $\rm 81 \pm 5$ & $\rm < 13$ & $\rm 20.2 \pm 0.4$ \\ & 276.124 & -29.7803 & 1342206987$^{*}$ & $\rm < 23$ & $\rm < 23$ & $\rm 11.5 \pm 0.5$ \\ $\rm [MAM2011]$ Aqu-MM2 & 277.2659 & -1.65041 & 1342254233 & $\rm 48 \pm 15$ & $\rm 9.9 \pm 2.4$ & $\rm < 2.8$ \\ $\rm [MAM2011]$ Aqu-MM4 & 277.2858 & -1.51188 & 1342254271$^{*,**}$ & $\rm 48 \pm 13$ & $\rm < 14$ & $\rm 3.2 \pm 0.8$ \\ $\rm [MAM2011]$ SerpS-MM1 & 277.407 & -1.84938 & 1342254230 & $\rm 105 \pm 6$ & $\rm < 14$ & $\rm 14.8 \pm 0.7$ \\ Serpens-SMM1a & 277.4575 & 1.255694 & 1342207781 & $\rm 3329 \pm 25$ & $\rm < 61$ & $\rm 147 \pm 1$ \\ EC82 & 277.487 & 1.24625 & 1342192975 & $\rm < 23$ & $\rm < 23$ & $\rm 8.4 \pm 0.5$ \\ Serpens-SMM4 & 277.4882 & 1.220302 & 1342193217 & $\rm 148 \pm 22$ & $\rm < 26$ & $\rm < 1.5$ \\ Serpens-SMM3 & 277.497 & 1.233806 & 1342193216 & $\rm 275 \pm 29$ & $\rm < 21$ & $\rm 3.3 \pm 0.5$ \\ & 277.4966 & 1.233416 & 1342207779 & $\rm 472 \pm 12$ & $\rm < 31$ & $\rm 31 \pm 2$ \\ Serpens-SMM18 & 277.51736 & -2.050743 & 1342254222 & $\rm 113 \pm 8$ & $\rm < 18$ & $\rm 13.9 \pm 0.8$ \\ $\rm [MAM2011]$ Aqu-MM6 & 277.6045 & -1.90372 & 1342254227$^{**}$ & $\rm < 21$ & $\rm < 14$ & $\rm < 2.4$ \\ $\rm [MAM2011]$ Aqu-MM7 & 277.6192 & -1.94658 & 1342254224$^{**}$ & $\rm 49 \pm 14$ & $\rm < 12$ & $\rm < 4.2$ \\ $\rm [MAM2011]$ Aqu-MM8 & 277.6209 & -1.93483 & 1342254228 & $\rm < 25$ & $\rm < 15$ & $\rm < 2.9$ \\ $\rm [MAM2011]$ Aqu-MM14 & 277.708 & -1.93502 & 1342254273 & $\rm < 20$ & $\rm < 19$ & $\rm < 2.2$ \\ $\rm [MAM2011]$ W40-MM3 & 277.7892 & -2.1068 & 1342254220 & $\rm 59 \pm 7$ & $\rm < 14$ & $\rm 10.3 \pm 0.6$ \\ $\rm [MAM2011]$ W40-MM5 & 277.7931 & -2.064 & 1342254268 & $\rm 802 \pm 12$ & $\rm < 17$ & $\rm 52 \pm 2$ \\ $\rm [MAM2011]$ W40-MM26 & 277.9439 & -2.07291 & 1342254266 & $\rm 1582 \pm 10$ & $\rm < 19$ & $\rm 7.2 \pm 1.1$ \\ $\rm [MAM2011]$ W40-MM27 & 277.9449 & -2.03886 & 1342254260 & $\rm 736 \pm 10$ & $\rm < 22$ & $\rm 10.8 \pm 0.7$ \\ $\rm [MAM2011]$ W40-MM28 & 277.9495 & -2.027 & 1342254264 & $\rm 656 \pm 7$ & $\rm < 14$ & $\rm 4.7 \pm 0.8$ \\ $\rm [MAM2011]$ W40-MM34 & 277.9885 & -2.00769 & 1342254259 & $\rm < 25$ & $\rm < 12$ & $\rm < 2.1$ \\ $\rm [MAM2011]$ W40-MM36 & 278.0556 & -1.95822 & 1342254263 & $\rm 76 \pm 10$ & $\rm < 17$ & $\rm 2.3 \pm 0.8$ \\ Vega & 279.2347 & 38.78369 & 1342188033 & $\rm < 10$ & $\rm < 10$ & $\rm 1.1 \pm 0.2$ \\ HD 172555 & 281.362 & -64.8712 & 1342215649 & $\rm 9 \pm 3$ & $\rm < 6$ & $\rm < 0.8$ \\ & 281.362 & -64.8712 & 1342228417 & $\rm 10 \pm 3$ & $\rm < 9$ & $\rm < 0.9$ \\ RXJ18523-3700 & 283.07208 & -37.0033 & 1342216163 & $\rm 18 \pm 2$ & $\rm < 6$ & $\rm 1.8 \pm 0.1$ \\ G34.26+0.15$^{***}$ & 283.328128 & 1.249232 & 1342209733 & $\rm -13757 \pm 286$ & $\rm < 511$ & $\rm 3049 \pm 15$ \\ S Cra & 285.2858 & -36.9555 & 1342207809 & $\rm 429 \pm 22$ & $\rm < 24$ & $\rm 21.1 \pm 0.6$ \\ RCrA-IRS 5A & 285.4504 & -36.9563 & 1342207806 & $\rm 1940 \pm 18$ & $\rm 170 \pm 48$ & $\rm 23 \pm 2$ \\ RCrA-IRS 5N & 285.451921 & -36.954133 & 1342207806 & $\rm 880 \pm 20$ & $\rm < 21$ & $\rm 12 \pm 1$ \\ RCrA-IRS 7A & 285.480417 & -36.954722 & 1342206990 & $\rm 4790 \pm 70$ & $\rm < 68$ & $\rm 122 \pm 6$ \\ SMM 1C & 285.48042 & -36.95464 & 1342206990 & $\rm 8730 \pm 60$ & $\rm 0 \pm 0$ & $\rm 95 \pm 6$ \\ R CrA & 285.473568 & -36.95218 & 1342206990 & $\rm 2340 \pm 50$ & $\rm < 58$ & $\rm 104 \pm 2$ \\ R CrA-IRS 7B & 285.485 & -36.9578 & 1342207807 & $\rm 4768 \pm 71$ & $\rm 686 \pm 210$ & $\rm 80 \pm 3$ \\ CrA 01 & 285.7444 & -37.1266 & 1342254253 & $\rm 1416 \pm 7$ & $\rm < 16$ & $\rm 47 \pm 0.9$ \\ HD 179218 & 287.7968 & 15.78766 & 1342208884 & $\rm 213 \pm 17$ & $\rm < 22$ & $\rm 27.0 \pm 0.4$ \\ LDN 723-mm & 289.47375 & 19.2055 & 1342208918 & $\rm 228 \pm 7$ & $\rm < 24$ & $\rm 7.7 \pm 0.6$ \\ HD 181296 & 290.7133 & -54.4239 & 1342209730 & $\rm < 6$ & $\rm < 6$ & $\rm < 0.6$ \\ HD 181327 & 290.7455 & -54.538 & 1342186311 & $\rm < 8$ & $\rm < 8$ & $\rm 1.8 \pm 0.3$ \\ Parsamian 21 & 292.254 & 9.645 & 1342254615 & $\rm < 53$ & $\rm < 53$ & $\rm 12 \pm 1$ \\ $\rm [SER2000]$ B335 & 294.2529 & 7.5689 & 1342208889 & $\rm 428 \pm 25$ & $\rm < 20$ & $\rm 10.8 \pm 0.5$ \\ HD 191089 & 302.2717 & -26.224 & 1342268180 & $\rm < 14$ & $\rm < 14$ & $\rm < 0.9$ \\ HD 192758 & 304.5657 & -42.86 & 1342216659 & $\rm < 9$ & $\rm < 9$ & $\rm < 0.7$ \\ AFGL 2591$^{***}$ & 307.353139 & 40.189365 & 1342208938 & $\rm -3202 \pm 89$ & $\rm < 480$ & $\rm 3397 \pm 12$ \\ DR 21 (OH)& 309.753208 & 42.380500 & 1342209400 & $\rm 6269 \pm 113$ & $\rm < 237$ & $\rm 665 \pm 7$ \\ LDN 1157-mm & 309.7758 & 68.0375 & 1342208909 & $\rm 596 \pm 21$ & $\rm < 25$ & $\rm 12.2 \pm 0.6$ \\ AU Mic & 311.2897 & -31.3408 & 1342193195 & $\rm < 8$ & $\rm < 8$ & $\rm < 0.5$ \\ HH 381 IRS & 314.589208 & 52.490806 & 1342258845$^{*}$ & $\rm 210 \pm 41$ & $\rm < 72$ & $\rm 47 \pm 2$ \\ HD 203024 & 319.0125 & 68.91447 & 1342206975 & $\rm < 25$ & $\rm < 25$ & $\rm 2.7 \pm 0.5$ \\ L1014 & 321.0312 & 49.98583 & 1342208911 & $\rm 76 \pm 16$ & $\rm < 22$ & $\rm < 1.5$ \\ HH 354 IRS & 331.710417 & 59.046389 & 1342262014 & $\rm 602.0 \pm 21$ & $\rm < 69$ & $\rm 56 \pm 1$ \\ V733 Cep & 343.388583 & 62.539889 & 1342262017 & $\rm < 72$ & $\rm < 72$ & $\rm < 4.8$ \\ Fomalhaut & 344.4127 & -29.6222 & 1342210402 & $\rm < 8$ & $\rm < 8$ & $\rm 1.5 \pm 0.2$ \\ HR 8799 & 346.869625 & 21.13425 & 1342212242 & $\rm < 7$ & $\rm < 7$ & $\rm < 0.4$ \\ NGC 7538 IRS1 & 348.4387 & 61.46944 & 1342211545$^{*}$ & $\rm 45248 \pm 147$ & $\rm < 440$ & $\rm 3300 \pm 11$ \\ HD 221853 & 353.9006 & 8.382618 & 1342212528 & $\rm < 10$ & $\rm < 10$ & $\rm < 0.7$ \\ \end{longtable} } \onllongtab{3}{ \begin{longtable}{llcccc} \caption{\label{Tab:YSO_fluxes_ext} [OI] fluxes at 63 $\rm \mu m$ from central spaxel, central 3x3 spaxels and IFU integrated }\\ \hline\hline Source name & obs ID & $\rm F _{[OI]}$ & $\rm F _{[OI]3x3}$ & $\rm F _{[OI]5x5}$ & Ext?$\rm ^{*}$ \\ -- & -- & ($\rm 10^{ -18}W/m^{2}$) & ($\rm 10^{ -18}W/m^{2}$) & ($\rm 10^{ -18}W/m^{2}$) & -- \\ \hline \endfirsthead \caption{continued.}\\ \hline\hline Source name & obs ID & $\rm F _{[OI]}$ & $\rm F _{[OI]3x3}$ & $\rm F _{[OI]5x5}$ & Ext?$\rm ^{*}$ \\ -- & -- & ($\rm 10^{ -18}W/m^{2}$) & ($\rm 10^{ -18}W/m^{2}$) & ($\rm 10^{ -18}W/m^{2}$) & -- \\ \hline \endhead \hline \endfoot W3IRS5 & 1342229093 & 25543.0$\rm \pm$3328.9 & 373064.6$\rm \pm$28956.2 & 1027951.6$\rm \pm$34137.0 & YYY \\ & 1342229091 & 22026.4$\rm \pm$520.2 & 304283.8$\rm \pm$1250.4 & 728668.6$\rm \pm$1775.3 & YYY \\ SMM J032537+30451 & 1342263508 & 1668.3$\rm \pm$8.2 & 2260.4$\rm \pm$28.9 & 2368.6$\rm \pm$57.8 & YYN \\ LDN1448N & 1342263506 & 1609.5$\rm \pm$10.0 & 4001.2$\rm \pm$30.0 & 4927.1$\rm \pm$66.8 & YYY \\ L1448-C(S) & 1342263510 & 545.1$\rm \pm$12.6 & 1054.2$\rm \pm$31.3 & 1418.6$\rm \pm$95.2 & YYY \\ IRAS 03235+3004 & 1342264250 & 238.0$\rm \pm$11.0 & 361.0$\rm \pm$25.8 & 492.8$\rm \pm$72.9 & YYN \\ LDN 1455 & 1342204122 & 65.0$\rm \pm$14.0 & $\rm <$37.9 & $\rm <$102.2 & NNN \\ 2MASSJ03283706+3113310 & 1342264248 & 187.8$\rm \pm$8.6 & 134.6$\rm \pm$17.8 & $\rm <$99.0 & NNN \\ SSTc2dJ032857.4+311416 & 1342264247 & $\rm <$30.0 & $\rm <$60.2 & 1239.7$\rm \pm$55.7 & NYY \\ SSTc2dJ032900.5+311200 & 1342264245 & 114.5$\rm \pm$9.1 & 343.8$\rm \pm$29.6 & 690.2$\rm \pm$76.1 & YYY \\ 2MASSJ03290149+3120208 & 1342264243 & 1791.2$\rm \pm$10.7 & 4074.3$\rm \pm$34.3 & 6597.4$\rm \pm$51.9 & YYY \\ 2MASSJ03290773+3121575 & 1342267612 & 1465.3$\rm \pm$8.5 & 7009.4$\rm \pm$47.5 & 21728.9$\rm \pm$98.5 & YYY \\ NGC1333IRAS4A & 1342216084 & 140.6$\rm \pm$20.1 & 419.7$\rm \pm$54.7 & 485.9$\rm \pm$126.9 & YNN \\ & 1342216083 & 119.3$\rm \pm$5.2 & 362.7$\rm \pm$10.6 & 183.8$\rm \pm$41.1 & YNN \\ SSTc2dJ032910.7+311821 & 1342267616 & $\rm <$28.1 & 999.2$\rm \pm$23.6 & 2501.7$\rm \pm$62.1 & YYY \\ SSTc2dJ032912.0+311301 & 1342267608 & 260.3$\rm \pm$7.4 & 110.8$\rm \pm$12.4 & $\rm <$167.5 & NNN \\ NGC1333IRAS4B & 1342216178 & 111.0$\rm \pm$23.0 & 235.0$\rm \pm$42.6 & $\rm <$88.3 & NNN \\ SSTc2dJ032913.5+311358 & 1342267610 & 54.0$\rm \pm$10.6 & $\rm <$67.6 & $\rm <$107.4 & NNN \\ IRAS03271+3013 & 1342263513 & 277.3$\rm \pm$14.2 & 299.7$\rm \pm$36.8 & 215.8$\rm \pm$68.5 & NNN \\ IRAS03282+3035 & 1342263515 & 146.4$\rm \pm$13.4 & 258.9$\rm \pm$33.3 & $\rm <$129.9 & YNN \\ IRAS03292+3039 & 1342265448 & 163.1$\rm \pm$9.0 & 211.8$\rm \pm$16.9 & 255.5$\rm \pm$52.2 & NNN \\ IRAS03301+3111 & 1342215668 & 386.6$\rm \pm$18.1 & 379.2$\rm \pm$43.7 & 445.4$\rm \pm$76.8 & NNN \\ SSTc2dJ033314.3+310710 & 1342263487 & 92.7$\rm \pm$7.7 & 186.5$\rm \pm$24.3 & 493.9$\rm \pm$77.0 & YYY \\ SSTc2dJ033316.4+310653 & 1342265450 & 397.0$\rm \pm$12.6 & 1076.1$\rm \pm$29.9 & 1586.3$\rm \pm$84.5 & YYY \\ B1-a & 1342216182 & 358.0$\rm \pm$14.8 & 745.4$\rm \pm$47.1 & 1008.5$\rm \pm$119.3 & YYN \\ SSTc2dJ033327.3+310710 & 1342265452 & 1065.3$\rm \pm$5.2 & 1229.3$\rm \pm$13.4 & 1259.6$\rm \pm$49.4 & YYN \\ IRAS03407+3152 & 1342265454 & 230.7$\rm \pm$9.4 & 596.3$\rm \pm$33.5 & 1267.0$\rm \pm$57.1 & YYY \\ SSTc2dJ034356.8+320305 & 1342265456 & 172.0$\rm \pm$8.5 & 378.9$\rm \pm$24.1 & 554.4$\rm \pm$48.4 & YYY \\ 2MASSJ03444389+3201373 & 1342265702 & 600.5$\rm \pm$6.9 & 868.1$\rm \pm$13.5 & 758.6$\rm \pm$72.2 & YNN \\ IRAS04016+2610 & 1342216216 & 429.9$\rm \pm$21.8 & 508.2$\rm \pm$43.1 & 598.5$\rm \pm$97.8 & NNN \\ & 1342204348 & 395.7$\rm \pm$5.6 & 406.4$\rm \pm$15.4 & 382.6$\rm \pm$22.9 & NNN \\ PP13S & 1342263505 & 1102.0$\rm \pm$41.0 & 1277.9$\rm \pm$90.0 & 1460.1$\rm \pm$175.7 & NNN \\ V773Tau & 1342216217 & 79.4$\rm \pm$3.5 & 81.6$\rm \pm$6.8 & 84.6$\rm \pm$13.0 & NNN \\ FMTau & 1342216218 & $\rm <$11.6 & 19.8$\rm \pm$6.5 & $\rm <$51.9 & YNN \\ CWTau & 1342216221 & 83.4$\rm \pm$5.8 & 101.7$\rm \pm$8.2 & 62.0$\rm \pm$13.3 & NNN \\ IRAS04158+2805 & 1342192793 & 55.1$\rm \pm$4.2 & 53.6$\rm \pm$8.4 & 63.1$\rm \pm$17.8 & NNN \\ BPTau & 1342192796 & 6.8$\rm \pm$2.4 & $\rm <$11.7 & $\rm <$49.3 & NNN \\ & 1342225728 & 9.3$\rm \pm$2.9 & 10.8$\rm \pm$3.1 & $\rm <$28.5 & NNN \\ IRAS04169+2702 & 1342265457 & 608.3$\rm \pm$7.0 & 951.3$\rm \pm$25.2 & 1362.3$\rm \pm$61.0 & YYY \\ IRAS04181+2654 & 1342265459 & 206.0$\rm \pm$6.4 & 235.4$\rm \pm$15.4 & 213.1$\rm \pm$28.8 & NNN \\ IRAM04191 & 1342216654 & 161.2$\rm \pm$22.7 & 176.7$\rm \pm$35.8 & 269.4$\rm \pm$93.3 & NNN \\ RYTau & 1342190361 & 99.6$\rm \pm$5.8 & 85.7$\rm \pm$10.0 & 85.5$\rm \pm$23.2 & NNN \\ TTau & 1342190353 & 8337.1$\rm \pm$15.2 & 16591.8$\rm \pm$58.0 & 20023.0$\rm \pm$69.0 & YYY \\ 2MASSJ04220069+2657324 & 1342265461 & 91.0$\rm \pm$7.9 & 144.4$\rm \pm$20.3 & 662.5$\rm \pm$58.4 & NYY \\ FSTau & 1342192791 & 378.3$\rm \pm$3.9 & 482.6$\rm \pm$8.0 & 601.3$\rm \pm$24.0 & YYY \\ FTTau & 1342192790 & 19.6$\rm \pm$6.2 & $\rm <$23.3 & $\rm <$55.1 & NNN \\ DGTauB & 1342192798 & 130.0$\rm \pm$4.8 & 582.7$\rm \pm$14.9 & 739.2$\rm \pm$29.3 & YYY \\ & 1342265463 & 514.7$\rm \pm$10.5 & 722.6$\rm \pm$32.4 & 810.6$\rm \pm$58.2 & YYN \\ DFTau & 1342190359 & 51.9$\rm \pm$7.0 & 35.9$\rm \pm$9.0 & $\rm <$70.3 & NNN \\ DGTau & 1342190382 & 605.7$\rm \pm$7.3 & 1492.7$\rm \pm$20.8 & 1709.3$\rm \pm$42.1 & YYY \\ & 1342225730 & 1533.0$\rm \pm$14.7 & 1682.4$\rm \pm$45.0 & 1921.8$\rm \pm$89.6 & YYN \\ IRAS04248+2612 & 1342265466 & 172.8$\rm \pm$9.5 & 207.0$\rm \pm$20.5 & 293.9$\rm \pm$44.9 & NNN \\ FWTau & 1342225735 & $\rm <$10.0 & 17.0$\rm \pm$5.6 & $\rm <$44.7 & YNN \\ IRAS04264+2433 & 1342265468 & 569.3$\rm \pm$9.3 & 628.0$\rm \pm$22.4 & 709.8$\rm \pm$38.6 & NYN \\ IQTau & 1342192135 & $\rm <$11.8 & 15.6$\rm \pm$4.6 & $\rm <$61.7 & YNN \\ & 1342225733 & 16.0$\rm \pm$2.6 & $\rm <$19.0 & $\rm <$35.6 & NNN \\ UXTau & 1342204350 & 37.7$\rm \pm$3.6 & 32.6$\rm \pm$6.0 & $\rm <$45.7 & NNN \\ & 1342214357 & 34.7$\rm \pm$3.1 & 29.6$\rm \pm$4.7 & $\rm <$55.1 & NNN \\ DKTau & 1342225732 & 21.2$\rm \pm$3.6 & $\rm <$16.3 & $\rm <$30.2 & NNN \\ IRAS04287+1801 & 1342192805 & 4611.6$\rm \pm$77.2 & 6547.5$\rm \pm$203.3 & 8233.6$\rm \pm$171.5 & YYY \\ HLTau & 1342190351 & 510.5$\rm \pm$10.5 & 859.8$\rm \pm$32.7 & 1413.5$\rm \pm$38.5 & YYY \\ HKTau & 1342225736 & 36.6$\rm \pm$3.3 & 35.2$\rm \pm$6.6 & 44.5$\rm \pm$11.1 & NNN \\ Haro6-13 & 1342192128 & 63.4$\rm \pm$5.3 & 47.5$\rm \pm$6.1 & $\rm <$48.4 & NNN \\ GGTau & 1342192121 & 51.9$\rm \pm$4.9 & 52.2$\rm \pm$8.0 & 84.9$\rm \pm$23.8 & NNN \\ UZTau & 1342192131 & $\rm <$13.1 & 48.2$\rm \pm$13.9 & $\rm <$62.8 & YNN \\ GITau & 1342225760 & $\rm <$7.1 & 39.7$\rm \pm$8.4 & $\rm <$34.6 & YNN \\ DLTau & 1342225800 & 26.9$\rm \pm$2.7 & 23.3$\rm \pm$3.4 & 18.1$\rm \pm$7.9 & NNN \\ HNTau & 1342225796 & 49.4$\rm \pm$2.8 & 65.8$\rm \pm$6.9 & 58.3$\rm \pm$15.1 & NNN \\ DMTau & 1342225825 & 10.8$\rm \pm$3.1 & $\rm <$12.7 & $\rm <$29.4 & NNN \\ CITau & 1342192125 & 27.0$\rm \pm$3.9 & 21.4$\rm \pm$6.6 & $\rm <$62.0 & NNN \\ AATau & 1342190357 & $\rm <$10.9 & 32.6$\rm \pm$9.9 & $\rm <$48.6 & YNN \\ & 1342225758 & 24.3$\rm \pm$2.1 & 17.8$\rm \pm$2.2 & $\rm <$29.8 & NNN \\ DNTau & 1342225757 & 4.0$\rm \pm$1.2 & $\rm <$14.8 & $\rm <$30.5 & NNN \\ IRAS04325+2402 & 1342267857 & 478.9$\rm \pm$10.7 & 634.0$\rm \pm$25.6 & 670.1$\rm \pm$43.7 & YYN \\ DOTau & 1342190385 & 62.8$\rm \pm$8.5 & 158.0$\rm \pm$14.4 & 267.5$\rm \pm$61.4 & YYN \\ HVTau & 1342225801 & 39.3$\rm \pm$4.3 & 83.9$\rm \pm$7.4 & 112.3$\rm \pm$30.1 & YNN \\ TMR1B & 1342225834 & 360.2$\rm \pm$6.6 & 704.8$\rm \pm$27.5 & 738.8$\rm \pm$76.3 & YYN \\ TMR1 & 1342192985 & 401.3$\rm \pm$19.4 & 711.4$\rm \pm$58.2 & 743.1$\rm \pm$138.3 & YNN \\ LkCa15 & 1342225798 & 11.3$\rm \pm$2.2 & $\rm <$10.6 & $\rm <$22.4 & NNN \\ IRAS04365+2535 & 1342225832 & 446.0$\rm \pm$13.1 & 816.2$\rm \pm$26.8 & 892.1$\rm \pm$56.3 & YYN \\ & 1342192987 & 529.8$\rm \pm$31.1 & 875.0$\rm \pm$59.7 & 858.9$\rm \pm$106.9 & YNN \\ L1527 & 1342192981 & 243.3$\rm \pm$19.4 & 778.3$\rm \pm$67.1 & 1259.1$\rm \pm$108.3 & YYY \\ & 1342192983 & 260.9$\rm \pm$5.8 & 728.7$\rm \pm$17.0 & 1227.8$\rm \pm$55.3 & YYY \\ IRAS04381+2540 & 1342225803 & 922.9$\rm \pm$17.1 & 1268.3$\rm \pm$50.4 & 1514.5$\rm \pm$114.9 & YYN \\ & 1342225830 & 111.4$\rm \pm$8.9 & 1321.5$\rm \pm$16.4 & 1504.7$\rm \pm$71.1 & YYN \\ CoKuTau4 & 1342191360 & 19.6$\rm \pm$5.3 & $\rm <$14.9 & $\rm <$56.0 & NNN \\ & 1342225837 & 23.1$\rm \pm$1.4 & 16.9$\rm \pm$1.5 & $\rm <$24.1 & NNN \\ IRAS04385+2550 & 1342225828 & 66.1$\rm \pm$2.8 & 75.9$\rm \pm$6.3 & 86.6$\rm \pm$14.7 & NNN \\ DPTau & 1342191362 & $\rm <$10.7 & 123.9$\rm \pm$11.9 & 143.4$\rm \pm$21.9 & YYN \\ DQTau & 1342225806 & 24.8$\rm \pm$4.7 & 33.6$\rm \pm$8.1 & $\rm <$57.1 & NNN \\ Haro6-37 & 1342225805 & $\rm <$15.1 & 11.7$\rm \pm$3.0 & $\rm <$58.7 & YNN \\ UYAur & 1342193206 & $\rm <$10.6 & 343.4$\rm \pm$9.4 & 354.2$\rm \pm$20.0 & YYN \\ & 1342215699 & 338.3$\rm \pm$5.8 & 370.1$\rm \pm$10.2 & 386.1$\rm \pm$24.5 & NNN \\ GMAur & 1342191357 & 23.4$\rm \pm$4.4 & 56.5$\rm \pm$15.6 & $\rm <$60.9 & NNN \\ ABAur & 1342191355 & 666.5$\rm \pm$19.6 & 849.5$\rm \pm$10.6 & 919.4$\rm \pm$19.9 & YYY \\ & 1342217842 & 358.6$\rm \pm$27.0 & 888.5$\rm \pm$64.4 & 1018.1$\rm \pm$132.7 & YYN \\ SUAur & 1342217844 & 85.7$\rm \pm$3.1 & 114.4$\rm \pm$6.0 & 127.0$\rm \pm$11.8 & YYN \\ HD31648 & 1342226002 & 101.3$\rm \pm$2.5 & 92.9$\rm \pm$5.1 & 105.3$\rm \pm$14.3 & NNN \\ RWAur & 1342191359 & 157.5$\rm \pm$9.3 & 214.0$\rm \pm$20.6 & 249.2$\rm \pm$38.9 & NNN \\ HD35187 & 1342226900 & 27.2$\rm \pm$2.9 & 34.1$\rm \pm$3.7 & 30.9$\rm \pm$9.1 & NNN \\ HD36112 & 1342227635 & 39.3$\rm \pm$3.5 & 40.6$\rm \pm$4.3 & 47.4$\rm \pm$13.2 & NNN \\ HD36910 & 1342227638 & 49.0$\rm \pm$3.1 & 51.0$\rm \pm$6.3 & $\rm <$35.3 & NNN \\ HD245906 & 1342228528 & 87.6$\rm \pm$24.1 & 197.2$\rm \pm$59.8 & $\rm <$157.4 & NNN \\ $\rm [SMZ2000]$L1643-S3MMS1 & 1342226195 & 1554.1$\rm \pm$7.9 & 2116.5$\rm \pm$16.2 & 2462.8$\rm \pm$61.2 & YYY \\ Re50NNIRS & 1342265946 & 1054.4$\rm \pm$31.8 & 1176.1$\rm \pm$68.0 & 1297.0$\rm \pm$178.6 & NNN \\ HD38120 & 1342226212 & 66.0$\rm \pm$15.9 & $\rm <$45.0 & $\rm <$99.8 & NNN \\ V1647Ori & 1342267864 & 139.2$\rm \pm$32.0 & 95.1$\rm \pm$26.0 & $\rm <$332.2 & NNN \\ NGC2071IR & 1342218761 & 7995.8$\rm \pm$39.3 & 40864.6$\rm \pm$154.3 & 60785.1$\rm \pm$199.4 & YYY \\ HD50138 & 1342206991 & 2103.4$\rm \pm$16.1 & 2125.3$\rm \pm$41.9 & 2464.9$\rm \pm$87.5 & NYY \\ RECX4 & 1342199241 & 1.0$\rm \pm$0.3 & $\rm <$13.1 & $\rm <$33.7 & NNN \\ RECX15 & 1342186314 & 30.1$\rm \pm$6.0 & 30.7$\rm \pm$7.7 & $\rm <$41.0 & NNN \\ RECX15 & 1342210388 & 23.8$\rm \pm$2.5 & 19.4$\rm \pm$3.6 & $\rm <$20.4 & NNN \\ $\rm MGL99]$IRS1757 & 1342211844 & 8155.9$\rm \pm$11.5 & 18459.8$\rm \pm$47.9 & 25508.3$\rm \pm$79.9 & YYY \\ SzCha & 1342233478 & 14.5$\rm \pm$1.8 & $\rm <$18.4 & $\rm <$27.0 & NNN \\ CRCha & 1342232614 & 16.4$\rm \pm$2.6 & $\rm <$17.7 & $\rm <$25.8 & NNN \\ TWA01 & 1342248544 & 32.4$\rm \pm$4.7 & 33.1$\rm \pm$4.3 & 33.8$\rm \pm$10.9 & NNN \\ & 1342187127 & 34.8$\rm \pm$3.9 & 55.1$\rm \pm$11.1 & $\rm <$33.3 & NNN \\ CSCha & 1342233480 & 16.5$\rm \pm$1.8 & $\rm <$19.8 & $\rm <$22.6 & NNN \\ CHX7 & 1342233477 & 15.4$\rm \pm$1.7 & $\rm <$29.8 & 64.7$\rm \pm$10.6 & NYY \\ HD97048 & 1342199412 & 852.7$\rm \pm$18.3 & 991.3$\rm \pm$47.5 & 1230.0$\rm \pm$94.4 & NYN \\ & 1342188436 & 1481.4$\rm \pm$3.7 & 1415.3$\rm \pm$6.1 & 1485.9$\rm \pm$18.9 & NNY \\ Sz27 & 1342233476 & 9.8$\rm \pm$1.3 & $\rm <$18.3 & $\rm <$21.3 & NNN \\ GMCha & 1342267619 & 221.4$\rm \pm$13.3 & 283.2$\rm \pm$26.4 & 555.4$\rm \pm$65.2 & NYY \\ T42 & 1342232291 & 492.8$\rm \pm$58.6 & 985.6$\rm \pm$287.2 & 1447.4$\rm \pm$357.1 & NNN \\ WWCha & 1342232292 & 422.1$\rm \pm$55.3 & $\rm <$239.3 & $\rm <$389.6 & NNN \\ Hn13 & 1342263492 & 4.1$\rm \pm$1.1 & $\rm <$13.4 & $\rm <$18.4 & NNN \\ CHX22 & 1342233474 & 24.4$\rm \pm$1.4 & 21.4$\rm \pm$6.2 & $\rm <$28.3 & NNN \\ HD98800B & 1342223821 & $\rm <$7.1 & 15.3$\rm \pm$3.5 & $\rm <$39.5 & YNN \\ & 1342199409 & 7.9$\rm \pm$2.0 & $\rm <$11.4 & $\rm <$24.3 & NNN \\ HD98922 & 1342210385 & 196.8$\rm \pm$16.5 & 260.6$\rm \pm$55.8 & $\rm <$68.8 & NNN \\ HD100453 & 1342203059 & 58.0$\rm \pm$5.2 & 51.0$\rm \pm$6.5 & $\rm <$51.4 & NNN \\ & 1342212228 & 34.0$\rm \pm$3.0 & 36.7$\rm \pm$5.4 & $\rm <$42.3 & NNN \\ & 1342211695 & 113.4$\rm \pm$39.2 & $\rm <$44.5 & $\rm <$80.8 & NNN \\ HD100546 & 1342188038 & 5461.5$\rm \pm$71.5 & 5339.0$\rm \pm$92.2 & 5536.9$\rm \pm$124.6 & NNN \\ & 1342188438 & 5645.1$\rm \pm$8.3 & 5605.1$\rm \pm$11.8 & 5923.6$\rm \pm$20.1 & NYY \\ TCha & 1342232294 & 48.3$\rm \pm$3.9 & 64.9$\rm \pm$18.4 & 56.8$\rm \pm$15.8 & NNN \\ HD104237 & 1342207819 & 85.1$\rm \pm$16.9 & $\rm <$47.7 & $\rm <$107.3 & NNN \\ & 1342212234 & 79.5$\rm \pm$4.0 & 71.3$\rm \pm$5.9 & 96.0$\rm \pm$18.4 & NNN \\ IRAS11590-6452 & 1342212230 & 1352.1$\rm \pm$41.3 & 1817.4$\rm \pm$73.5 & 2376.8$\rm \pm$117.4 & YYY \\ & 1342212232 & 1340.4$\rm \pm$33.9 & 1935.1$\rm \pm$81.1 & 2229.8$\rm \pm$101.6 & YYN \\ DK Cha & 1342188039 & 1329.5$\rm \pm$119.1 & 3095.9$\rm \pm$302.7 & 3696.5$\rm \pm$219.7 & YYN \\ & 1342226006 & 2715.9$\rm \pm$77.9 & 3197.7$\rm \pm$239.2 & 3773.0$\rm \pm$169.4 & NYN \\ ChaII-J125342.86-771511.5 & 1342226005 & 72.9$\rm \pm$3.2 & 72.4$\rm \pm$8.8 & 70.7$\rm \pm$11.6 & NNN \\ 2MASSJ12590656-7707401 & 1342265694 & 139.1$\rm \pm$8.9 & 144.4$\rm \pm$27.8 & 150.2$\rm \pm$41.4 & NNN \\ ChaII-J130055.36-771022.1 & 1342226008 & $\rm <$6.5 & $\rm <$13.0 & 60.9$\rm \pm$16.7 & NYY \\ ChaII-J130424.92-775230.1 & 1342229800 & 10.0$\rm \pm$3.1 & $\rm <$14.1 & $\rm <$40.5 & NNN \\ Hn24 & 1342235656 & 8.7$\rm \pm$1.7 & $\rm <$18.2 & $\rm <$30.0 & NNN \\ ChaII-J130520.68-773901.4 & 1342229831 & 14.9$\rm \pm$4.6 & $\rm <$11.3 & $\rm <$30.5 & NNN \\ ChaII-J130521.66-773810.0 & 1342229830 & 6.4$\rm \pm$2.0 & $\rm <$9.4 & $\rm <$30.0 & NNN \\ ChaII-J130806.28-775505.2 & 1342229801 & 15.0$\rm \pm$4.3 & $\rm <$16.3 & $\rm <$30.8 & NNN \\ ChaII-J130827.17-774323.2 & 1342228419 & $\rm <$7.8 & 4.6$\rm \pm$1.0 & $\rm <$35.2 & YNN \\ SAO206462 & 1342190370 & 42.7$\rm \pm$5.0 & 45.9$\rm \pm$8.6 & $\rm <$60.1 & NNN \\ HD139614 & 1342191300 & 41.4$\rm \pm$6.7 & 55.9$\rm \pm$13.5 & $\rm <$46.4 & NNN \\ & 1342215683 & 19.0$\rm \pm$5.9 & $\rm <$53.6 & $\rm <$106.0 & NNN \\ HD141569 & 1342190376 & 190.6$\rm \pm$4.3 & 218.3$\rm \pm$8.6 & 199.2$\rm \pm$28.5 & NNN \\ & 1342213913 & 202.9$\rm \pm$19.1 & 220.4$\rm \pm$33.6 & $\rm <$80.1 & NNN \\ HD142666 & 1342214224 & 13.5$\rm \pm$2.4 & 22.2$\rm \pm$4.5 & $\rm <$29.9 & NNN \\ HD142527 & 1342216173 & 49.5$\rm \pm$3.0 & 44.6$\rm \pm$5.8 & 39.9$\rm \pm$10.9 & NNN \\ RULup & 1342215682 & 184.5$\rm \pm$20.3 & 133.0$\rm \pm$42.5 & $\rm <$105.4 & NNN \\ RYLup & 1342216171 & $\rm <$20.3 & 60.9$\rm \pm$14.4 & $\rm <$71.6 & YNN \\ UScoJ160421.7-213028 & 1342215681 & 24.9$\rm \pm$2.2 & 31.2$\rm \pm$4.8 & 35.7$\rm \pm$11.9 & NNN \\ Sz91 & 1342229827 & 8.2$\rm \pm$1.5 & $\rm <$19.5 & 39.5$\rm \pm$10.4 & NNY \\ HD144668 & 1342192146 & 128.1$\rm \pm$5.6 & 110.7$\rm \pm$6.8 & 94.6$\rm \pm$19.4 & NNN \\ & 1342215641 & 131.6$\rm \pm$14.3 & 147.3$\rm \pm$44.6 & $\rm <$83.2 & NNN \\ Sz111 & 1342229828 & 9.7$\rm \pm$1.4 & $\rm <$17.4 & $\rm <$23.8 & NNN \\ AS205 & 1342215737 & 222.0$\rm \pm$21.3 & 248.3$\rm \pm$55.5 & $\rm <$107.2 & NNN \\ UScoJ161420.2-190648 & 1342216191 & 43.1$\rm \pm$3.7 & 43.5$\rm \pm$5.8 & 62.0$\rm \pm$20.1 & NNN \\ RXJ1615.3-3255 & 1342229825 & 16.8$\rm \pm$1.7 & $\rm <$20.1 & $\rm <$23.0 & NNN \\ 2MASSJ16230923-2417047 & 1342250127 & 19.4$\rm \pm$1.6 & 32.4$\rm \pm$9.1 & 40.5$\rm \pm$9.2 & NNN \\ Doar21 & 1342240163 & 385.1$\rm \pm$56.2 & 1515.5$\rm \pm$294.7 & 1575.7$\rm \pm$354.2 & YYN \\ GSS30-IRS1 & 1342215678 & 1836.3$\rm \pm$51.4 & 4000.3$\rm \pm$100.1 & 10189.4$\rm \pm$165.8 & YYY \\ VLA1623-243 & 1342213918 & 398.6$\rm \pm$19.1 & 1653.6$\rm \pm$64.2 & 3606.3$\rm \pm$95.1 & YYY \\ WL12 & 1342228187 & 494.4$\rm \pm$17.6 & 408.9$\rm \pm$31.9 & 309.8$\rm \pm$74.0 & NNN \\ DoAr28 & 1342241707 & 10.1$\rm \pm$1.5 & $\rm <$18.8 & $\rm <$24.0 & NNN \\ Oph01 & 1342266925 & 3109.3$\rm \pm$9.8 & 7549.8$\rm \pm$32.1 & 7938.9$\rm \pm$60.2 & YYY \\ Elias29 & 1342228519 & 1840.0$\rm \pm$69.3 & 2521.7$\rm \pm$92.8 & 3562.4$\rm \pm$129.5 & YYY \\ SR21 & 1342229824 & 12.6$\rm \pm$4.1 & $\rm <$38.3 & $\rm <$54.3 & NNN \\ IRS46 & 1342228474 & 75.3$\rm \pm$22.6 & 718.4$\rm \pm$77.4 & 1530.7$\rm \pm$119.8 & YYY \\ IRS44 & 1342228474 & 436.3$\rm \pm$36.7 & 807.5$\rm \pm$56.1 & 1530.7$\rm \pm$119.8 & YYY \\ IRS48 & 1342227069 & 322.3$\rm \pm$18.7 & 362.4$\rm \pm$42.6 & 457.1$\rm \pm$80.8 & NNN \\ GY314 & 1342242632 & 112.7$\rm \pm$35.0 & $\rm <$168.3 & $\rm <$281.0 & NNN \\ DoAr44 & 1342250578 & 25.3$\rm \pm$1.5 & 32.3$\rm \pm$9.7 & $\rm <$25.9 & NNN \\ IRS63 & 1342228473 & 204.6$\rm \pm$18.4 & 316.9$\rm \pm$72.0 & 669.8$\rm \pm$173.3 & NNN \\ L1689SNO2 & 1342241270 & 235.5$\rm \pm$70.4 & $\rm <$199.6 & $\rm <$288.4 & NNN \\ 2MASSJ16320099-2456419 & 1342263469 & 1091.7$\rm \pm$11.8 & 1329.5$\rm \pm$14.8 & 1354.9$\rm \pm$55.2 & YYN \\ V346Nor & 1342267622 & 381.9$\rm \pm$35.6 & 966.5$\rm \pm$57.3 & 1482.9$\rm \pm$136.1 & YYY \\ RNO90 & 1342228206 & 128.1$\rm \pm$19.8 & $\rm <$47.8 & $\rm <$72.2 & NNN \\ HBC650 & 1342215639 & 811.2$\rm \pm$5.7 & 1040.9$\rm \pm$24.2 & 1055.5$\rm \pm$76.7 & YYN \\ HD150193 & 1342216625 & 22.8$\rm \pm$3.4 & 52.0$\rm \pm$10.2 & $\rm <$40.6 & NNN \\ Sco01 & 1342267176 & 223.3$\rm \pm$8.9 & 272.7$\rm \pm$33.7 & 274.2$\rm \pm$46.5 & NNN \\ KKOph & 1342192148 & 168.2$\rm \pm$5.9 & 186.1$\rm \pm$8.6 & 218.0$\rm \pm$23.4 & NNN \\ HD158643 & 1342217821 & 48.8$\rm \pm$2.8 & 65.1$\rm \pm$6.2 & $\rm <$38.4 & NNN \\ HD163296 & 1342192161 & 206.4$\rm \pm$5.3 & 209.9$\rm \pm$12.2 & 197.4$\rm \pm$22.4 & NNN \\ & 1342243512 & 197.1$\rm \pm$5.5 & 228.8$\rm \pm$10.4 & 237.5$\rm \pm$17.3 & NNN \\ & 1342217819 & 210.5$\rm \pm$18.3 & 220.7$\rm \pm$56.6 & $\rm <$104.9 & NNN \\ $\rm [SER2000]$L483 & 1342192156 & 384.9$\rm \pm$5.9 & 595.4$\rm \pm$23.9 & 774.6$\rm \pm$69.2 & YYN \\ HD169142 & 1342186310 & 80.6$\rm \pm$4.6 & 123.0$\rm \pm$10.0 & 157.6$\rm \pm$26.3 & YNN \\ $\rm [MAM2011]$Aqu-MM2 & 1342254233 & 47.7$\rm \pm$14.5 & 172.0$\rm \pm$40.0 & 212.1$\rm \pm$68.3 & NNN \\ $\rm [MAM2011]$Aqu-MM4 & 1342254271 & 48.4$\rm \pm$12.5 & 157.2$\rm \pm$22.6 & $\rm <$150.0 & YNN \\ $\rm [MAM2011]$SerpS-MM1 & 1342254230 & 105.5$\rm \pm$6.0 & 311.4$\rm \pm$39.9 & 504.7$\rm \pm$60.8 & YYN \\ Serpens-SMM1a & 1342207781 & 3329.5$\rm \pm$24.9 & 6982.9$\rm \pm$78.3 & 8213.7$\rm \pm$139.3 & YYY \\ Serpens-SMM4 & 1342193217 & 146.3$\rm \pm$22.5 & 349.6$\rm \pm$40.3 & 2858.9$\rm \pm$112.5 & YYY \\ Serpens-SMM3 & 1342207779 & 472.1$\rm \pm$12.2 & 1557.0$\rm \pm$18.9 & 3455.2$\rm \pm$79.5 & YYY \\ Serpens-SMM3 & 1342193216 & 121.3$\rm \pm$24.9 & 931.5$\rm \pm$46.1 & 2247.4$\rm \pm$93.0 & YYY \\ Serpens-SMM18 & 1342254222 & 113.2$\rm \pm$7.6 & 2483.5$\rm \pm$93.5 & 5440.7$\rm \pm$103.7 & YYY \\ $\rm [MAM2011]$Aqu-MM7 & 1342254224 & 49.1$\rm \pm$13.8 & $\rm <$63.5 & $\rm <$139.2 & NNN \\ $\rm MAM2011]$Aqu-MM8 & 1342254228 & $\rm <$25.2 & 26.3$\rm \pm$7.5 & $\rm <$116.4 & YNN \\ $\rm [MAM2011$]W40-MM3 & 1342254220 & 59.5$\rm \pm$6.6 & 706.3$\rm \pm$18.5 & 2777.0$\rm \pm$49.7 & YYY \\ $\rm [MAM2011]$W40-MM5 & 1342254268 & 801.8$\rm \pm$11.7 & 4489.4$\rm \pm$24.1 & 10708.7$\rm \pm$46.1 & YYY \\ $\rm [MAM2011]$W40-MM26 & 1342254266 & 1581.5$\rm \pm$10.0 & 10123.5$\rm \pm$26.4 & 27140.0$\rm \pm$97.4 & YYY \\ $\rm [MAM2011]$W40-MM27 & 1342254260 & 736.2$\rm \pm$10.1 & 6070.4$\rm \pm$27.7 & 20842.8$\rm \pm$70.2 & YYY \\ $\rm [MAM2011]$W40-MM28 & 1342254264 & 656.3$\rm \pm$7.1 & 5504.3$\rm \pm$19.7 & 13798.5$\rm \pm$65.8 & YYY \\ $\rm [MAM2011]$W40-MM36 & 1342254263 & 75.7$\rm \pm$10.2 & 160.0$\rm \pm$23.5 & 255.9$\rm \pm$38.0 & YYN \\ HD172555 & 1342215649 & 8.6$\rm \pm$2.6 & 25.2$\rm \pm$8.3 & $\rm <$32.0 & NNN \\ & 1342228417 & 10.5$\rm \pm$3.2 & $\rm <$13.1 & $\rm <$32.8 & NNN \\ SCra & 1342207809 & 429.5$\rm \pm$22.4 & 514.0$\rm \pm$58.6 & 827.8$\rm \pm$141.9 & NNN \\ IRS5A & 1342207806 & 1943.7$\rm \pm$18.6 & 3467.4$\rm \pm$183.1 & 5015.8$\rm \pm$187.0 & YYY \\ & 1342207806 & 1939.8$\rm \pm$18.4 & 3638.1$\rm \pm$60.6 & 5488.0$\rm \pm$91.2 & YYY \\ IRS7A & 1342206990 & 8706.1$\rm \pm$61.7 & 26366.2$\rm \pm$345.7 & 47729.5$\rm \pm$320.4 & YYY \\ IRS7B & 1342207807 & 4768.1$\rm \pm$71.2 & 26576.1$\rm \pm$360.1 & 41090.1$\rm \pm$283.0 & YYY \\ CrA01 & 1342254253 & 1416.3$\rm \pm$7.3 & 1946.1$\rm \pm$36.7 & 2314.6$\rm \pm$42.3 & YYY \\ HD179218 & 1342208884 & 213.1$\rm \pm$16.5 & 229.9$\rm \pm$42.1 & $\rm <$96.4 & NNN \\ LDN723-mm & 1342208918 & 227.9$\rm \pm$7.2 & 250.3$\rm \pm$19.1 & 306.2$\rm \pm$47.6 & NNN \\ $\rm [SER2000]$B335 & 1342208889 & 421.5$\rm \pm$25.9 & 553.9$\rm \pm$56.2 & 606.7$\rm \pm$125.9 & NNN \\ AFGL 2591 & 1342208938 & $\rm <$428.9 & 27712.4$\rm \pm$292.1 & 46314.9$\rm \pm$377.2 & YYY \\ LDN1157-mm & 1342208909 & 595.0$\rm \pm$21.2 & 771.5$\rm \pm$40.7 & 885.3$\rm \pm$88.2 & YYN \\ HH381IRS & 1342258845 & 209.6$\rm \pm$41.0 & $\rm <$136.5 & 480.5$\rm \pm$117.8 & NNY \\ L1014 & 1342208911 & 75.8$\rm \pm$15.5 & $\rm <$46.3 & $\rm <$109.4 & NNN \\ HH354IRS & 1342262014 & 602.1$\rm \pm$21.4 & 639.8$\rm \pm$47.3 & $\rm <$317.4 & NNN \\ NGC 7538 IRS1 & 1342211545 & 45247.8$\rm \pm$146.9 & 148957.8$\rm \pm$610.7 & 257177.9$\rm \pm$689.3 & YYY \\ \end{longtable} \tablefoot{(*): the extended emission column shows the results for the 3x3 spaxels versus 1 spaxel/ 5x5 versus 1 spaxel and / 5x5 versus 3x3 spaxels. \\ } } \begin{table*}[!t] \scriptsize \caption{Fitted parameters for [OI] lines with multiple components} \label{Tab:multiG} \centering \begin{tabular}{lllllllllllll} % \hline \hline Source& peak$\rm_{0}$ & $\rm \lambda_{0,0}$ & $\rm v_{0}$ & $\rm \sigma_{0}$ & peak$\rm_{1}$ & $\rm \lambda_{0,1}$ & $\rm v_{1}$ & $\rm \sigma_{1}$ & peak$\rm_{2}$ & $\rm \lambda_{0,2}$ & $\rm v_{2}$ & $\rm \sigma_{2}$ \\ & (Jy) & ($\rm \mu m$) & (km/s) & ($\rm \mu m$) & (Jy) & ($\rm \mu m$) & (km/s) & ($\rm \mu m$) & (Jy) & ($\rm \mu m$) & (km/s) & ($\rm \mu m$) \\ \hline T Tau & 15.52 & 63.159 & -123 & 0.0074 & 309.76 & 63.182 & -14 & 0.0075 & 209.03 & 63.19 & 24 & 0.0101 \\ FS Tau$\rm ^{*}$ & 16.15 & 63.183 & -10 & 0.007 & 7.78 & 63.189 & 19 & 0.0122 & 0.0 & 0.0 & 0.0 & 0.0 \\ DG Tau$\rm ^{*}$& 17.62 & 63.181 & -20 & 0.0072 & 14.3 & 63.182 & -13 & 0.0149 & 0.0 & 0.0 & 0.0 & 0.0 \\ HL Tau$\rm ^{*}$ & 25.11 & 63.197 & 55 & 0.0083 & 3.92 & 63.185 & 0.0 & 0.0279 & 0.0 & 0.0 & 0.0 & 0.0 \\ TMR 1 & 2.58 & 63.156 & -189 & 0.012 & 10.36 & 63.203 & 85 & 0.0185 & 0.0 & 0.0 & 0.0 & 0.0\\ UY Aur$\rm ^{*}$ & 12.77 & 63.183 & -10 & 0.0065 & 8.03 & 63.188 & 16 & 0.0133 & 0.0 & 0.0 & 0.0 & 0.0 \\ AB Aur$\rm ^{*}$ & 40.46 & 63.182 & -15 & 0.0071 & 9.42 & 63.194 & 41 & 0.0074 & 0.0 & 0.0 & 0.0 & 0.0 \\ HD 50138 & 3.33 & 63.144 & -197 & 0.0083 & 68.75 & 63.179 & -30 & 0.015 & 10.94 & 63.214 & 137 & 0.0077 \\ R Mon & 4.03 & 63.139 & -219 & 0.004 & 61.89 & 63.180 & -22 & 0.0154 & 14.22 & 63.212 & 126 & 0.0099 \\ HD 97048$\rm ^{*}$ & 77.27 & 63.192 & 33 & 0.0074 & 19.49 & 63.19 & 25 & 0.0116 & 0.0 & 0.0 & 0.0 & 0.0\\ T 42 & 9.73 & 63.169 & -77 & 0.0051 & 12.6 & 63.191 & 28 & 0.0164 & 0.0 & 0.0 & 0.0 & 0.0 \\ WW Cha & 12.38 & 63.178 & -35 & 0.013 & 5.16 & 63.207 & 105 & 0.0121 & 0.0 & 0.0 & 0.0 & 0.0 \\ HD 100546 & 149.42 & 63.184 & -7 & 0.0155 & 66.07 & 63.210 & 121 & 0.008 & 0.0 & 0.0 & 0.0 & 0.0 \\ ChaII-J12531723-7707107 & 53.95 & 63.163 & -105 & 0.0103 & 74.84 & 63.187 & 8 & 0.0095 & 14.92 & 63.213 & 131 & 0.013 \\ GSS30 IRS1 & 3.91 & 63.142 & -206 & 0.0152 & 28.76 & 63.184 & -4 & 0.0137 & 31.69 & 63.203 & 88 & 0.0178 \\ Elias 29 & 7.58 & 63.166 & -89 & 0.0044 & 48.41 & 63.186 & 3 & 0.0198 & 0.0 & 0.0 & 0.0 & 0.0 \\ RCrA IRS5A & 14.33 & 63.167 & -86 & 0.0117 & 51.29 & 63.198 & 61 & 0.0157 & 0.0 & 0.0 & 0.0 & 0.0 \\ RCrA IRS7C & 23.05 & 63.166 & -90 & 0.004 & 197.92 & 63.188 & 12 & 0.0211 & 36.5 & 63.206 & 99 & 0.0086 \\ RCrA IRS7B & 4.96 & 63.233 & 227 & 0.0961 & 79.13 & 63.187 & 10 & 0.0203 & 55.44 & 63.210 & 119 & 0.0128 \\ Re 50 NN IRS$\rm ^{*}$ & 45.99 & 63.184 & -5 & 0.010 & 9.20 & 63.200 & 71 & 0.014 & 0.0 & 0.0 & 0.0 & 0.0 \\ Cra01 & 3.64 & 63.164 & -100 & 0.005 & 68.61 & 63.190 & 24 & 0.011 & 0.0 & 0.0 & 0.0 & 0.0 \\ 2MASS J16320099-2456419$\rm ^{*}$ & 57.10 & 63.189 & 19 & 0.00829 & 8.63 & 63.192 & 33 & 0.014 & 0.0 & 0.0 & 0.0 & 0.0 \\ $\rm [MAM2011]$ W40-MM5$\rm ^{*}$ & 27.13 & 63.183 & -9 & 0.0075 & 25.13 & 63.194 & 43 & 0.00827 & 0.0 & 0.0 & 0.0 & 0.0 \\ IRAS 04325+2402$\rm ^{*}$ & 32.15 & 63.185 & 0 & 0.0074 & 2.98 & 63.20 & 71 & 0.003 & 0.0 & 0.0 & 0.0 & 0.0 \\ L1448-C(S)$\rm ^{*}$ & 28.29 & 63.186 & 5 & 0.0088 & 5.60 & 63.204 & 90 & 0.0069 & 0.0 & 0.0 & 0.0 & 0.0 \\ LDN 1448N$\rm ^{*}$ & 82.94 & 63.189 & 19 & 0.0090 & 8.08 & 63.184 & -5 & 0.015 & 0.0 & 0.0 & 0.0 & 0.0 \\ SMM J032537+30451$\rm ^{*}$ & 41.18 & 63.188 & 14 & 0.008 & 50.41 & 63.185 & 0 & 0.012 & 0.0 & 0.0 & 0.0 & 0.0 \\ VELA IRS 17 & 375.22 & 63.190 & 24 & 0.099 & 108.50 & 63.197 & 57 & 0.005 & 3.84 & 63.218 & 157 & 0.0078 \\ Serpens-SMM3 & 19.00 & 63.155 & -142 & 0.0093 & 24.18 & 63.189 & 19 & 0.011 & 11.52 & 63.218 & 157 & 0.0077 \\ Serpens-SMM1a$\rm ^{*}$ & 29.68 & 63.182 & -14 & 0.016 & 144.41 & 63.190 & 24 & 0.0091 & 7.60 & 63.210 & 119 & 0.0076 \\ 2MASS J23134531+6128116 & 225.25 & 63.168 & -81 & 0.010 & 2205 & 63.175 & -47 & 0.0099 & 1.38 & 63.197 & 57 & 0.024 \\ IRAS 11590-6452 & 15.54 & 63.160 & -119 & 0.0083 & 5051 & 63.182 & -14 & 0.010 & 3.99 & 63.187 & 9 & 0.029 \\ \hline \end{tabular} \tablefoot{(*): the velocity separation between the components is smaller than the spectral resolution ($\rm \sim 88~km/s$).} \end{table*}

1607.07991

1607

1607.07443_arXiv.txt

As several large single-dish radio surveys begin operation within the coming decade, a wealth of radio data will become available and provide a new window to the Universe. In order to fully exploit the potential of these data sets, it is important to understand the systematic effects associated with the instrument and the analysis pipeline. A common approach to tackle this is to forward-model the entire system -- from the hardware to the analysis of the data products. For this purpose, we introduce two newly developed, open-source Python packages: the HI Data Emulator (\hide) and the Signal Extraction and Emission Kartographer (\seek) for simulating and processing single-dish radio survey data. \hide forward-models the process of collecting astronomical radio signals in a single-dish radio telescope instrument and outputs pixel-level time-ordered-data. \seek processes the time-ordered-data, removes artifacts from Radio Frequency Interference (RFI), automatically applies flux calibration, and aims to recover the astronomical radio signal. The two packages can be used separately or together depending on the application. Their modular and flexible nature allows easy adaptation to other instruments and data sets. We describe the basic architecture of the two packages and examine in detail the noise and RFI modeling in \hide, as well as the implementation of gain calibration and RFI mitigation in \seek. We then apply \hide \& \seek to forward-model a Galactic survey in the frequency range 990 -- 1260 MHz based on data taken at the Bleien Observatory. For this survey, we expect to cover 70\% of the full sky and achieve a median signal-to-noise ratio of approximately 5 -- 6 in the cleanest channels including systematic uncertainties. However, we also point out the potential challenges of high RFI contamination and baseline removal when examining the early data from the Bleien Observatory. The fully documented \hide \& \seek packages are available at \url{http://hideseek.phys.ethz.ch/} and are published under the GPLv3 license on GitHub.

\label{sec:introduction} Forward-modeling has become a common approach in various fields of astronomy where mock data sets are simulated and analyzed in parallel with the science data. This has become especially prevalent in cosmology where large data sets are used and high precision is required. Prominent examples are analyses of the cosmic microwave background \citep{reinecke2006simulation}, spectroscopy \citep{nord2016spokes} and weak gravitational lensing \cite[e.g.][]{bridle2009handbook, refregier2014way, bruderer2015calibrated, peterson2015simulation}. These forward-modeling pipelines simulate the astrophysical signals, the instrument response and the data reduction process in order to understand any systematic biases from hardware or software and to estimate statistical errors in the measurement chain. In this paper, we implement this forward-modeling approach for single-dish radio surveys. Several single-dish radio surveys are being planned for the next decades with the goal of mapping the H$_{\rm I}$ neutral hydrogen in the Universe \citep{Battye2013, Santos2015, Bigot-Sazy2016}. We develop two software packages: the {\texttt{HI Data Emulator}\xspace} (\hide) and the {\texttt{Signal Extraction and Emission Kartographer}\xspace} (\seek). \hide forward models the entire radio survey system chain, while \seek processes both the simulated data and the observed survey data in a reproducible and consistent way. Various sophisticated simulation and data reduction pipeline packages for radio astronomy exist \citep[e.g.,][]{swinbank2015lofar, mcmullin2007casa, dodson2016imaging}. However, many of them are either non-open source or project-specific. \hide \& \seek are developed in a different angle -- the initial functionalities are rather simple, but can be expanded easily as the codes are designed with a high level of modularity, flexibility and transparency, in a pure Python implementation with rigorous testing. Developing the two packages simultaneously has the advantage that the individual components of one pipeline can be cross validated against its counter part in the other pipeline. \hide \& \seek are developed based on the hardware system and data products from the 7m telescope at the Bleien Observatory as described in \cite[][hereafter C16]{chang2016an}. This framework is then used to forward-model a Galactic survey in the frequency range 990 -- 1260 MHz conducted at the Bleien Observatory for testing and science verification purposes. Such an analysis allows us to forecast the expected power of this survey with the existing hardware system at Bleien. Comparing the results of the forward model and data also helps to identify areas that require improvements in \hide \& \seek as well as the hardware system. This paper is organized as follows. In Section \ref{sec:pipelines} we first describe the basic architecture and design of \hide \& \seek. Detailed implementations for specific functionalities are described in \ref{sec:beam_convolution}, \ref{sec:rfi} and \ref{sec:flux_calibration}. In Section \ref{sec:application2bgs}, we apply \hide \& \seek to forward-model a survey based on early data taken at the Bleien Observatory. This includes customizing the various functionalities to this specific survey and providing a forecast for the expected outcome of the survey. Finally, we conclude in Sections \ref{sec:conclusion}. In \ref{sec:comparison} we show an example of how we applied \seek to process part of the early data from the Bleien Observatory and what we learn comparing these results to the \hide simulations. Information for downloading and installing \hide \& \seek, as well as the default file format is described in \ref{sec:distribution} and \ref{sec:fileformat}, respectively.

\label{sec:conclusion} In this paper, we present two software packages, \hide \& \seek, for simulating and processing data from single-dish radio surveys. \hide simulates the entire system chain of a radio telescope from the astronomical signal to the time-ordered-data (TOD). \seek on the other hand processes simulated and real observed TOD to \healpix maps while calibrating the signal and automatically masking contamination from radio frequency interference (RFI). The two packages together provide an end-to-end forward-modeling framework that can be used to systematically understand and test the various steps in the data processing procedure. They can also be used independently if only simulation or data processing is needed. The packages were developed based on the data taken at the Bleien Observatory, but can be easily adjusted to model and process similar projects. The other strength of using both \hide \& \seek is the possibility to perform cross-validation between the simulation and the analysis algorithm. We demonstrate how the two packages are used together to study systematic effects such as imperfect baseline removal and RFI leakage. We present the main architecture of the codes and how typical data is simulated and processed. We then describe more specifically how we apply the two packages to a forward-modeling exercise of a Galactic survey at 990 -- 1260 MHz. We use \hide to simulate the entire survey with the main survey characteristics matched to the Science Verification data from the Bleien Observatory. We then process this simulated data with \seek, using settings close to that used for the observed data. The result of this forward model is a forecast of the expected output from an idealized Galactic survey, with simplifications that are well-understood. We predict a median sky coverage of 50\% and a median signal-to-noise ratio of 2.3. In the channels from 990 MHz to 1010 MHz we expect a total sky coverage up to 70\% with median signal-to-noise ratio of approximately 9 without systematic errors and 5 -- 6 with systematic errors. \hide \& \seek follow common software engineering best practices. Being compliant with well-established coding standards, they offer great flexibility for defining data processing pipelines. Although some of the current implementations in \hide \& \seek are relatively simple and more general compared to existing software such as \cite{offringa2010lofar, peck2013serpent, mcmullin2007casa}, they have the advantage of being open-source code with a rigorous structure, and thus provide an easy-to-use foundation to build upon for more complicated functionalities. We have developed both packages in pure Python and increase the performance of computationally intensive parts by just-in-time compilation. By doing so, we are able to perform the RFI mitigation of the TOD in \seek at an rate of 190-200 GB/h/CPU. Using \ivy's parallelization scheme, we can furthermore distribute the workload to multiple cores and make efficient use of the available hardware. This allows us to easily process the expected data volume from a five-month survey on a modern laptop. In \hide we implemented an efficient beam convolution on the sphere by using \textit{Quaternions} combined with a KD-Tree. Creating the simulations for the Galactic survey with \hide took around 3 hours on a single-core of an average notebook. Further information such as documentation of the software can be found at \url{http://hideseek.phys.ethz.ch/}. \appendix

1607.07443

1607

1607.02266_arXiv.txt

We present Submillimeter Array (SMA) 1.35~mm subarcsecond angular resolution observations toward the \LK\ intermediate-mass star-forming region. The dust emission arises from a filamentary structure of $\sim$5~arcsec ($\sim$4500~au) enclosing VLA~1-3 and MM~1, perpendicular to the different outflows detected in the region. The most evolved objects are located at the southeastern edge of the dust filamentary structure and the youngest ones at the northeastern edge. The circumstellar structures around VLA~1, VLA~3, and MM~1 have radii between $\sim$200 and $\sim$375~au and masses in the $\sim$0.08--0.3\msun\ range. The 1.35~mm emission of VLA~2 arises from an unresolved (r$\la 135$~au) circumstellar disk with a mass of $\sim$0.02\msun. This source is powering a compact ($\sim$4000~au), low radial velocity ($\sim$7\kms) SiO bipolar outflow, close to the plane of the sky. We conclude that this outflow is the ``large-scale" counterpart of the short-lived, episodic, bipolar outflow observed through H$_2$O masers at much smaller scales ($\sim $180~au), and that has been created by the accumulation of the ejection of several episodic collimated events of material. The circumstellar gas around VLA~2 and VLA~3 is hot ($\sim$130~K) and exhibits velocity gradients that could trace rotation. There is a bridge of warm and dense molecular gas connecting VLA~2 and VLA~3. We discuss the possibility that this bridge could trace a stream of gas between VLA~3 and VLA~2, increasing the accretion rate onto VLA~2 to explain why this source has an important outflow activity.

Accretion disks and mass-loss processes with the presence of magnetic fields govern the formation of low-mass stars \citep[e.g.][]{mac08, fra14, rao14}. These mass-loss processes are non-steady, presenting pulsed events that are probably related to episodic increases in the accretion rates \citep[e.g.][]{re01, pec10, aud14}. In the case of high-mass protostars, there is increasing evidence that the associated outflows are also non-steady. This is seen, in particular, with the detection through Very Long Baseline Interferometry (VLBI) H$_2$O maser observations of short-lived (tens of years), episodic outflow events, which are interpreted as due to variability in the accretion processes as in the case of low-mass protostars \citep[e.g.][]{tor01, san12, tri13, car15}. \begin{figure*} \includegraphics[width=17cm]{LkHa234_1mm_2.eps} \caption{{\it Panel (a):} Contour map of the 1.35~mm continuum emission toward \LK. {\it Panel (b):} same as (a) but using only the visibilities for baselines longer than 100~k$\lambda$. Contours are $-4$, 4, 8, 14, and from 20 to 170 in steps of 10, times 0.6~\mJy. The crosses show the position of the YSOs previously reported: \LK, VLA~1, 2, 3A and 3B are taken from \citet[][VLA 3A is the source closer to the dust peak]{tri04} and MM~1 from \citet{fue01}. The black ellipses show the synthesised beams for each map. } \label{Fig-cont} \end{figure*} \begin{table*} \caption{1.35 mm continuum sources\label{T1}} \begin{tabular}{cccrlrcrr} \hline & \multicolumn{1}{c}{$\alpha$(J2000)} & \multicolumn{1}{c}{$\delta$(J2000)} && \multicolumn{2}{c}{Deconvolved Size$^a$} & \\ \cline{2-3} & \cline{4-5} & \multicolumn{1}{c}{$21^{\rm h} 43^{\rm m}$} & \multicolumn{1}{c}{$66\degr 6'$} & \multicolumn{1}{c}{$S_{\nu}$} & \multicolumn{1}{c}{Major axis, Minor axis} & \multicolumn{1}{c}{P.A.} & \multicolumn{1}{c}{$M${\em$^b$}} & \multicolumn{1}{c}{$n({\rm H_2})$} & \multicolumn{1}{c}{$N({\rm H_2})$} \\ Source & \multicolumn{1}{c}{($^{\rm s}$)} & \multicolumn{1}{c}{($''$)} & \multicolumn{1}{c}{(mJy)} & \multicolumn{1}{c}{(arcsec$\times$arcsec)} & \multicolumn{1}{c}{(\DG)} & \multicolumn{1}{c}{(M$_{\odot}$)} & \multicolumn{1}{c}{(\cmt)} & \multicolumn{1}{c}{(\cmd)} \\ \hline VLA 3 & 6.461 & 55.00 &139\Mm3 & 0.44\Mm0.02$\times$0.26\Mm0.02 & $-22$\Mm12 & 0.25 & $1.8 \, 10^9$ & $5.4\,10^{24}$ \\ VLA 2 & 6.321 & 55.95 & 12\Mm1 & $\la$0.3$^c$ & & 0.023 & $\ga2 \, 10^8$ & $\ga6\,10^{23}$\\ MM 1 & 6.292& 57.26 & 24\Mm2 & 0.83\Mm0.15$\times$0.24\Mm0.15 & $-4$\Mm10 & 0.09--0.29 & (3--9)$ \, 10^8$ & (1--4)$\,10^{24}$ \\ VLA 1 & 6.084& 58.06 & 22\Mm2 & 0.74\Mm0.15$\times$0.23\Mm0.11 & 35\Mm12 & 0.08--0.26 & (3--10)$ \, 10^8$ & (1--4)$\,10^{24}$\\ \hline \end{tabular} \\ {\em $^a$} Obtained by fitting a Gaussian with the CASA imfit task to each source in the image from Fig.~\ref{Fig-cont}b. The major and minor axis are the full width at half maximum values of the deconvolved size. \\ {\em $^b$} Mass estimated from the dust emission assuming: optically thin emission, a gas to dust ratio of 100, a dust opacity per unit of dust mass of 1.1~cm$^2$/g for VLA~2--3 and 0.9~cm$^2/$g for MM~1 and VLA~1, which are the computed values by \citet{Ossenkopf94} at the observed wavelength and for dust particles with thin and thick ices, respectively. See Section 3.1 for the value of temperature used for each source. \\ {\em $^c$} The source appears to be unresolved. The Gaussian fit yields an uncertainty for the major axis of $0.09$~arcsec. We adopt an upper limit which is three times this uncertainty (i.e., at 3-$\sigma$ level). \\ \end{table*} Very recently, we extended these VLBI water maser studies towards the cluster of intermediate-mass young stellar objects (YSOs) close to the optically visible Herbig Be star LkH$\alpha$234 in the NGC~7129 star-forming region \citep{tor14}. The YSOs in this cluster (including the star LkH$\alpha$234) are grouped within a radius of $\sim$5~arcsec ($\sim$4500~au at a distance of 0.9~kpc; see Kato et al. 2011 and references therein). Five of these YSOs show radio continuum emission: LkH$\alpha$~234, VLA~1, VLA~2, VLA~3A, and VLA~3B \citep{tri04}, all of them, excepting VLA~1, with mid-infrared counterparts \citep{kat11}. From the absence of near-infrared emission in the J, H, and K-bands of VLA~2 and VLA~3A+3B (hereafter VLA~3), but being bright in the $L'$, $M'$, and mid-infrared bands, \citet{kat11} concluded that these two objects are highly embedded protostars, having also associated H$_2$O maser emission \citep[][see also Fig. 1 in Torrelles et al. 2014 showing the spatial distribution of the YSOs in the cluster]{tof95, ume02, tri04, mar05}. Our multi-epoch VLBI H$_2$O maser observations \citep{tor14} revealed a very compact ($\sim$0.2~arcsec, $\sim$180~au), short-lived (kinematic age of $\sim$40~yr), episodic, and highly collimated water maser bipolar outflow emerging from VLA~2. These results predicted the presence of an accretion disk associated with VLA~2 as well as a relatively compact bipolar molecular outflow when observed through thermal molecular lines. In this paper, we present Submillimeter Array (SMA) dust continuum, and thermal molecular line observations at $\sim$1~mm wavelengths carried out towards VLA~2 to detect and characterise the expected disk-YSO-outflow system at scales of a few hundred au (Section~2). We show the observational results in Section~3, discussing them in Section~4. The main conclusions of this research are presented in Section~5.

We report SMA 1.35~mm continuum and spectral line observations (angular resolution $\sim$0.5-0.8~arcsec) towards the intermediate-mass star-forming region around the optically visible star LkH$\alpha$~234. We have detected a dust ridge of $\sim$5~arcsec ($\sim$4500~au) size containing the cluster of YSOs VLA~1, VLA~2, VLA~3A, VLA~3B, and MM1. The mass of the dust ridge is $\sim$7~M$_\odot$ while the circumstellar masses around these YSOs are in the range of $\sim$0.02--0.3~M$_\odot$. We find an evolutionary star formation sequence along the dust ridge, with the most evolved objects located at the southeastern edge of the ridge and the youngest ones located at the northeastern edge. The outflow activity of the different YSOs in the region occurs approximately perpendicular to the dust ridge. This overall trend of orientation of the outflows could be explained by the presence of a ``large-scale" magnetic field perpendicular to the dust ridge. Our SMA observations reveal towards VLA~2, a compact ($\sim$4~arcsec, $\sim$4000~au), SiO bipolar outflow, moving close to the plane of the sky in the northeast-southwest direction. The kinematic age of the SiO outflow is $\sim$350~yr. We conclude that this outflow is the ``large-scale" counterpart of the much more compact ($\sim$0.2~arcsec, $\sim$180~au), short-lived ($\sim$40~yr), episodic, bipolar H$_2$O maser outflow previously reported with VLBI observations towards VLA~2. We propose that the bipolar outflow seen in SiO has been created by the accumulation of material through repetitive ejections of collimated gas as the one observed through the H$_2$O masers. There is a bridge of molecular gas connecting VLA~3 and VLA~2 having SiO emission at ambient velocities. We discuss the possibility that this SiO ambient emission in the VLA~3--2 molecular bridge is due to an interaction of the gas transfer between VLA~3 and VLA~2. If the gas transfer occurs from VLA~3 to VLA~2 it could increase the accretion rate onto VLA~2, explaining the observed high outflow activity in this source.

1607.02266

1607

1607.07675_arXiv.txt

{We use a sample built on the SDSS DR7 catalogue and the bulge-disc decomposition of Simard et al. (2011) to study how the bulge and disc components contribute to the parent galaxy's star formation activity, by determining its position in the star formation rate (SFR) - stellar mass (M$_{\star}$) plane at 0.02$<z<$0.1. For this purpose, we use the bulge and disc colours as proxy for their SFRs. We study the mean galaxy bulge-total mass ratio (B/T) as a function of the residual from the MS ($\Delta_{MS}$) and find that the B/T-$\Delta_{MS}$ relation exhibits a parabola-like shape with the peak of the MS corresponding to the lowest B/Ts at any stellar mass. The lower and upper envelop of the MS are populated by galaxies with similar B/T, velocity dispersion and concentration ($R_{90}/R_{50}$) values. The mean values of such distributions indicate that the majority of the galaxies are characterised by classical bulges and not pseudo-bulges. Bulges above the MS are characterised by blue colours or, when red, by a high level of dust obscuration, thus indicating that in both cases they are actively star forming. When on the MS or below it, bulges are mostly red and dead. At stellar masses above $10^{10.5} $M$_{\odot}$, bulges on the MS or in the green valley tend to be significantly redder than their counterparts in the quiescence region, despite similar levels of dust obscuration. This could be explained with different age or metallicity content, suggesting different evolutionary paths for bulges on the MS and green valley with respect to those in the quiescence region. The disc color anti-correlates at any mass with the distance from the MS, getting redder when approaching the MS lower envelope and the quiescence region. The anti-correlation flattens as a function of the stellar mass, likely due to a higher level of dust obscuration in massive SF galaxies. We conclude that the position of a galaxy in the LogSFR-LogM$_{\star}$ plane depends on the star formation activity of its components: above the MS both bulge and disk are actively star forming. The nuclear activity is the first to be suppressed, moving the galaxies on the MS. Once the disk stops forming stars as well, the galaxy moves below the MS and eventually to the quiescence region. This is confirmed by a large fraction ($\sim45\%$) of passive galaxies with a secure two component morphology, coexisting with a population of pure spheroidals. Our findings are qualitatively in agreement with the compaction-depletion scenario, in which subsequent phases of gas inflow in the centre of a galaxy and depletion due to high star formation activity move the galaxy across the MS before the final quenching episode takes place. }

The Main Sequence (MS) of star-forming galaxies (SFGs) is a linear relation between the star formation rate (SFR) of a galaxy and its stellar mass (M$_{\star}$). This relation is close to linear, with a small scatter (0.2-0.3 dex) over a wide stellar mass range (e.g., Brinchmann et al. 2004; Noeske et al. 2007a; Daddi et al. 2007; Noeske et al. 2007b; Elbaz et al. 2007; Salim et al. 2007; Whitaker et al. 2012; Speagle et al. 2014; Pannella et al. 2015). It has also been shown that, while the slope of the relation remains $\sim1$ up to $z\sim2$, the normalisation changes drastically with redshift, where the SFR of galaxies at $z\sim 2$ is almost 20 times larger than that of galaxies at $z\sim0$. (e.g., Schreiber et al. 2015). The physical processes that cause this decrease in SF activity are generally grouped under the term {\textit{quenching}}, a topic at the core of galaxy evolution studies. Quenching works by prohibiting SF in a galaxy, by acting on the cold gas reservoir. Processes that expel cold gas from galaxies such as galactic winds driven by supernovae and massive stars are efficient in dark matter (DM) haloes with $M < 10^{12}$M$_{\odot}$ \citep[e.g.][]{1986ApJ...303...39D,2000MNRAS.317..697E}. Above this mass threshold, more powerful outflows are required to counter the deeper DM halo potential well. In such haloes, accreting central BHs can quench star formation by either heating the gas in the disk \citep[quasar mode;][]{2006ApJS..163....1H,2006ApJ...652..864H,2010MNRAS.404..180D,2010A&A...518L.155F,2012ARA&A..50..455F,2014A&A...562A..21C}, or mechanically removing the gas from the inner regions through powerful radio jets \citep[radio mode;][]{2006MNRAS.365...11C,2006MNRAS.370..645B,2006MNRAS.366..499D,2007ARA&A..45..117M,2008MNRAS.390.1399B,2013MNRAS.436.3031V,2005Natur.433..604D,2007ApJ...658...65C,2009Natur.460..213C}. Besides ejecting the gas, the effects of quenching can also be caused by a lack of cold gas supply. In DM halos with $M > 10^{12}$ M$_{\odot}$, the transition from cold to hot accretion \citep{2006MNRAS.368....2D} would prohibit further SF \citep{2014ApJ...789L..21C}. On the other hand, \cite{2009ApJ...707..250M}, propose morphological quenching where the growth of central mass concentration, i.e., a massive bulge, could simply stabilise a gas disc against fragmentation, thus preventing SF without affecting the cold gas reservoir in the galaxy. It has also been proposed that quenching could be a combination of all these processes. For example, using their cosmological zoom-in simulations, \cite{2015MNRAS.450.2327Z} show that quenching of high redshift SFGs is preceded by a compaction phase that creates the so called blue nuggets: SFGs morphologically similar to quiescent ones. This compaction phase is caused by strong inflows to the centre due to minor mergers, and/or counter--rotating gas and violent disc instabilities. The availability of dense cold gas in the centre results in a high SF and consequent stellar/supernova and/or AGN feedback causing quenching. \cite{2015Sci...348..314T} propose that episodes of compaction, quenching and replenishment are responsible for the position of a galaxy in a $\pm 0.3$ dex region around the MS. The final quenching episode results in a passive galaxy when the replenishment time is longer than the depletion time, typical of massive haloes, or low redshifts. There is increasing evidence in support of this picture from an observations. At $z\gtrsim2$, SFGs with the highest central gas densities are remarkably compact and have high S\'ersic indices ($n$) and spheroidal morphologies \citep{2011ApJ...742...96W,2014ApJ...791...52B,2013ApJ...766...15P,2013ApJ...768...92S,2013AAS...22131307W,2016ApJ...828...27N}. These galaxies resemble the quiescent population at the same redshift but are radically different from other SFGs that have irregular and clumpy appearances (\citealt{2004ApJ...604L..21E}; \citealt{2008ApJ...687...59G}; \citealt{2015ApJ...800...39G}). At low redshift, less massive SFGs are mainly pure disks, while more massive SFGs are characterised by a bulge+disc structure (\citealt{2012ApJ...753..114W}; \citealt{2012MNRAS.427.1666B}; \citealt{2014MNRAS.444.1660B}; \citealt{2014ApJ...788...11L}). \citet{2011ApJ...742...96W} analyse the dependence of galaxy structure (size and S\'{e}rsic index) on the position of the galaxies with respect to the MS ridge at $z=0-2$. They find that the S\'{e}rsic index tends to be roughly the same, $n\sim1$, within $3\sigma$ from the MS, i.e., there is no significant gradient of $n$ across the MS. They observe an increase of $n$ only in the starburst region. While Cheung et al. (2012) confirm that the S\'ersic index most sharply discriminates between the red sequence and the blue cloud, they also find a large number of blue outliers ($\approx40\%$) in their $n>2.3$ galaxy sample at $z\sim 0.65$. Concurrently, using a sample of local galaxies \cite{Guo:2015kp} report that that bulges/bars are responsible for the large dispersion of the sSFR, but only for massive SFGs. These recent results imply that a significant fraction of SFGs must have a bulge+disk morphology. Under the assumption that the star formation activity is confined in the disc, \cite{2014ApJ...785L..36A} show that defining the sSFR (specific SFR) with respect to the disc mass rather than the total galaxy mass decreases (and sometimes erases) the dependence of sSFR on stellar mass. This suggests that the slope of the sSFR - M$_{\star}$ relation reflects the increase of the bulge prominence with stellar mass, and also points towards the importance of treating galaxies as multicomponent systems. The aim of this paper is to study the relation between galaxies structural parameters and location with respect to the MS in the local Universe. The choice of performing this study on a local galaxy sample is dictated by the fact that the Sloan Digital Sky Survey \citep[SDSS;][]{2002AJ....124.1810S} provides the required statistics to dissect the MS brick by brick to study with high accuracy the nature of its scatter as a function of the galaxy morphology. In addition, the availability of accurate morphological classification and bulge/disc decomposition (Simard et al. 2011) of the SDSS galaxies allows to study the role and the interconnection of the individual galaxy components. This is the first paper of a series, which follow the evolution of the relation of the galaxy structural parameter and the scatter across the MS in the $0-1.2$ redshift window. The paper is organised as follows. In Section 1 we describe our dataset. In Section 2 we analyse the reliability of the morphological classification and of the colours of the bulge and disc in individual galaxies. In Section 3 we present our results and in Section 4 we present our discussion and conclusions. Throughout the study, the following cosmology is assumed: $H_0=70.0\ km\ s^{-1}Mpc^{-1}$, $\Omega_m$=0.3 and $\Omega_{\Lambda}$=0.7.

\label{discussion} We used the bulge-disc decomposition of Simard et al. (2011) to study the link between a galaxies' structural parameters and their location in the LogSFR-LogM$_{\star}$ plane. We use the colour of bulges and disks to understand how these individual components influence the evolution of the galaxy as a whole. Our sample is drawn from the spectroscopic sample of SDSS DR7, in the redshift range $0.02<z<0.1$ and M$_{\star}>10^9$M$_{\odot}$, and the main findings are as follows: \begin{itemize} \item The MS of star forming galaxies is populated by galaxies that, at every stellar mass, have the lowest B/Ts. At low stellar masses, LogM$_{\star}<$10.2M$_{\odot}$, MS galaxies are pure discs, while for more massive galaxies the prominence of the bulge component increases with increasing stellar mass. \item The upper and lower envelopes of the MS are populated by galaxies characterised by intermediate B/Ts. This is robust against different decompositions, and independent of the modelling of the bar component; \item Bulges in the upper envelop of the MS are characterised by blue colours at low stellar masses, or red colours and large dust obscuration at high stellar masses. This is consistent with high SF activity in the central region of the galaxy. \item The study of the mean bulge velocity dispersion and galaxy concentration parameter indicate that galaxies populating the upper and lower envelope of the MS are structurally similar. The values of the concentration parameter, in particular, suggest that these galaxies are characterised by classical bulges rather than pseudo-bulges. \item In the low mass regime, M$_{\star}< 10^{10}$M$_{\odot}$, disc and bulge colours show a similar behaviour at fixed stellar mass, becoming progressively redder from the upper envelop of the MS to the passive region. Nevertheless, the reddening of the bulge component is steeper than for discs. For M$_{\star}> 10^{10}$M$_{\odot}$ disks and bulges become bluer when going form the MS towards its upper envelop, despite a less pronounced total variation of the colours. The trend of the bulge colour reverses in the lower envelop of the MS, where bulges are redder than in the passive region. \item The population of passive galaxies is largely made of genuine bulge + disc systems (at least 45$\%$). \end{itemize} Our results point to a tight link between the distribution of galaxies around the MS and their structural parameters. Contrary to Wuyts et al. (2011) and Cheung et al. (2012), we find that blue bulgy star forming galaxies are not outliers in the distribution of galaxies on the MS. They occupy the upper tail of the MS distribution leading to a progressive increase in B/T above the MS. In particular, Wuyts et al. (2011) find that the S\'ersic index remains $\sim 1$ in the region of the MS. We find, instead, that the MS location corresponds to the lowest value of the B/T at any mass. This discrepancy could be due to the inability of a single S\'ersic model to capture small bulges in disc dominated galaxies. In agreement with previous results, we find a large percentage of bulge dominated systems in the high mass SF region, where the MS scatter tends to increase \citep[i.e.][]{2014ApJ...788...11L,2014ApJ...795..104W,2016MNRAS.455.2839E}. However, the B/T-$\Delta_{MS}$ relation at fixed stellar mass holds in the stellar mass range $10^{9-11} $M$_{\odot}$, and seems to flatten where the MS itself tends to disappear. We also investigate how the presence of a bar influences the scatter around the MS. Overall we conclude that bars do not affect the scatter of the MS or the B/T -$\Delta_{MS}$ relation. In fact, this relation holds for both the unbarred and the barred G09 subsamples. This suggests, partly in disagreement with the findings of \cite{Guo:2015kp}, that the predominance of the bulge in a galaxy is intrinsically related to the location of a galaxy around the MS, and this is true at any stellar mass independent of the presence of a bar. Nevertheless, we do not exclude the possibility that a large fraction of bars among very massive SFGs can contribute to an increase in the scatter of the MS, as suggested by \cite{Guo:2015kp}. \begin{figure} \includegraphics[width=0.95\columnwidth,trim={17cm 0cm 0cm 0cm},clip ]{abramson.eps} \caption{B/T ratio as a function of the distance from the MS, represented by the vertical dotted line. Here $\Delta_{MS}$ has been computed from the sSFR normalised to M$_{disc}$. Different bins of stellar mass are indicated with different colours. } \label{ssfr} \end{figure} Our study shows that the location of galaxies in the LogSFR-LogM$_{\star}$ plane and, in particular, with respect to the MS of SFGs, is determined by the combined activity (or inactivity) of the individual galaxy components. This was previously suggested by \cite{2014ApJ...785L..36A}, who pointed out that the MS becomes roughly linear if one assumes that the SF activity takes place only in the disc rather than in the entire galaxy. However, our results suggest otherwise. The effect of normalising the SF activity to the disc mass leads to a completely flat and tight relation only for pure disc galaxies, that, as shown here, dominate the core of the MS distribution. Extending this normalisation to galaxies of intermediate morphology has the additional effect of increasing the scatter by $\sim0.1$ dex in any mass bin. We suggest the following explanation: \begin{equation} SFR_{gal}=SFR_{disc}+SFR_{bulge} \end{equation} thus: \begin{equation} \frac{SFR_{gal}}{M_{disc}}=sSFR_{disc}+\frac{SFR_{bulge}}{M_{disc}} \end{equation} \noindent where M$_{disc}$ is the disc mass and sSFR$_{disc}$ is the specific SFR of the disc. The last equation leads to the following effects. \begin{itemize} \item[-]For pure disc galaxies, which dominate the core of the MS, the mass of the disc equals the mass of the galaxy and SFR$_{bulge}\sim 0$, hence SFR$_{gal}/$M$_{disc}$ = sSFR$_{disc}$. Since the MS is nearly linear, sSFR$_{disc}$-M$_{\star}$ relation is flat with a very small scatter. \item[-]For intermediate B/T galaxies in the upper envelope of the MS, thus with a blue star forming bulge, SFR$_{bulge}> 0$ and M$_{disc}< $ M$_{star}$. This implies that SFR$_{gal}/$M$_{disc}$ = sSFR$_{disc}$ + SFR$_{bulge}/$M$_{disc}$. Therefore such galaxies are displaced well above the sSFR$_{disc}$-M$_{\star}$ relation by the contribution of SFR$_{bulge}/$M$_{disc}$. The larger the SFR of the bulge, the larger the displacement. \item[-]For intermediate B/T galaxies in the lower envelope of the MS, thus with a red quiescent bulge, SFR$_{bulge}\sim 0$, hence SFR$_{gal}/$M$_{disc}$ = sSFR$_{disc}$. They scatter around the sSFR$_{disc}-$M$_{\star}$ relation in the same way as around the MS. \end{itemize} In Fig. \ref{ssfr}, we show the relation between B/T and distance from the MS computed as the difference between the sSFR$_{disk}$ of galaxies and the sSFR$_{disk}$ of galaxies in the MS. This is done for galaxies that are pure disks, or have a secure bulge+disc structure. The B/T trends in Fig. \ref{ssfr} are similar to the ones in Fig. \ref{fig9}: a steep increase of the mean B/T in the upper envelop of the MS due to the displacement of bulgy SFGs above the sSFR$_{disc}-$M$_{\star}$ relation. We underline that the B/T in the passive region is lower that the one observed in Fig. 5, as here pure spheroidal galaxies are excluded from the sample. We find that the overall effect of neglecting the SF activity of the bulge component is to increase the scatter of the sSFR-M$_{\star}$ relation. It has been proposed that minor mergers or violent disc instabilities could favour the flow of cold gas from the disc towards the galaxy centre and, thus, cause an overall compaction of the system. The high SF activity in the centre would lead to a compact, bulgy, star forming object \citep{2014MNRAS.438.1870D,2015MNRAS.450.2327Z}. \cite{2015Sci...348..314T} study the evolution of galaxies in this scenario using zoom-in simulations and they find that galaxies at high redshift undergo subsequent phases of compaction and depletion of the gas reservoir which ultimately leads to quenching. The compaction phase causes high SF in the central region of galaxies, and hence could favour bulge growth. This phase is then followed by depletion due to gas exhaustion. Such subsequent phases can move a galaxy across the MS: towards the upper envelope during compaction, and towards the lower envelope during depletion. Complete quiescence can be reached once the bulge reaches a given mass threshold corresponding to no more inflows in massive halos, or AGN feedback. Our findings on the B/T in the LogSFR-LogM$\star$ plane and bulge/disk colours can be related to such a scenario, and are also consistent with the observed gradient of molecular and atomic gas fraction across the MS, as seen by \cite{2016MNRAS.462.1749S}. Integral Field Spectroscopy surveys like MaNGA \citep{2015ApJ...798....7B} and CALIFA \citep{2012A&A...538A...8S} will greatly advance our understanding of the evolution of individual galaxy components, and how they impact the galaxy as a whole. This will help us in drawing a better picture of the quenching mechanism, and in understanding the CSFH at any epoch.

1607.07675

1607

1607.06390_arXiv.txt

We estimate the conventional astrophysical emission from dwarf spheroidal satellite galaxies (dSphs) of the Milky Way, focusing on millisecond pulsars (MSPs), and evaluate the potential for confusion with dark matter (DM) annihilation signatures at GeV energies. In low-density stellar environments, such as dSphs, the abundance of MSPs is expected to be proportional to stellar mass. Accordingly, we construct the \gammaRayHyph luminosity function of MSPs in the Milky Way disk, where $>90$ individual MSPs have been detected with the \textit{Fermi} Large Area Telescope (LAT), and scale this luminosity function to the stellar masses of 30 dSphs to estimate the cumulative emission from their MSP populations. We predict that MSPs within the highest stellar mass dSphs, Fornax and Sculptor, produce a \gammaRayHyph flux $>500$~MeV of $\sim10^{-11}$~ph~cm$^{-2}$~s$^{-1}$, which is a factor $\sim10$ below the current LAT sensitivity at high Galactic latitudes. The MSP emission in ultra-faint dSphs, including targets with the largest J-factors, is typically several orders of magnitude lower, suggesting that these targets will remain clean targets for indirect DM searches in the foreseeable future. For a DM particle of mass 25~GeV annihilating to $b$ quarks at the thermal relic cross section (consistent with DM interpretations of the Galactic Center excess), we find that the expected \gammaRayHyph emission due to DM exceeds that of MSPs in all of the target dSphs. Using the same Milky Way MSP population model, we also estimate the Galactic foreground MSP coincidence probability along the same sightlines to the dSphs.

\label{sec:intro} Searches for the annihilation products of dark matter (DM) are now testing significant portions of the theoretically motivated parameter space for weakly interacting massive particles. The rapid progress of indirect DM searches can be attributed to a large number of astrophysical probes that have become available over the last decade \citep[reviewed by][]{Gaskins:2016cha}. Among these, the Large Area Telescope (LAT) on board \textit{Fermi} has played a vital role due to its full-sky coverage, sensitivity, and energy range relevant for DM searches at the electroweak scale. LAT data have been used for numerous DM searches involving a variety of astrophysical objects, including dwarf spheroidal satellite galaxies (dSphs) of the Milky Way (MW). dSphs are especially promising targets for indirect DM searches due to their (1) substantial DM content \citep[e.g.,][]{mat98, Simon:2007dq} and proximity, (2) distribution over a range of Galactic latitudes, including regions with low diffuse foreground emission, and (3) dearth of non-thermal production mechanisms. No \gammaRayHyph signal has been conclusively associated with dSphs, either individually or as a population, and the corresponding upper limits have been used to set competitive constraints on DM annihilation \citep[summarized by][]{Charles:2016pgz}. For example, a joint analysis of 15 dSphs with 6 years of LAT data excluded DM particles annihilating at the canonical thermal relic cross section in some annihilation channels for DM masses up to 100 GeV \citep{ack15}. Although the non-DM \gammaRayHyph emission from dSphs is expected to be low, no empirical measurement and few quantitative estimates for this contribution have been previously available. dSphs have old stellar populations \citep[e.g.,][]{2012ApJ...753L..21B,2014ApJ...789..147W} and low gas content \citep{2009ApJ...696..385G,2014ApJ...795L...5S}, and therefore contain few sites for non-thermal radiation from cosmic-ray (CR) interactions. However, their ancient stellar populations might include small populations of \gammaRayHyph-emitting millisecond pulsars (MSPs), which have characteristic ages of several Gyr based upon their measured spin periods and period derivatives ($\tau \equiv P / 2 \dot{P}$). MSPs are luminous sources that account for nearly half of LAT-detected pulsars. In addition, 25 MW globular clusters, which have similar-age stellar populations to dSphs, have been detected by the collective emission of their MSP populations \citep{abd10,Hooper:2016rap}. MSPs exhibit hard spectral indices $\sim1.5$ and spectral cut-offs around 3~GeV \citep{abd13,cho14}. As a consequence, their intensity peaks in the GeV range, where the LAT sensitivity is highest. The characteristic spectral shape of MSPs is also similar to that of the Galactic Center excess, and many authors have investigated the contribution of MSPs to that signal \citep[e.g.,][]{Abazajian:2012pn,Brandt:2015,bar15,Lee:2015fea,hoo15}. In this Letter, we estimate the conventional astrophysical emission intrinsic to dSphs, focusing on MSPs, and evaluate the potential for confusion with DM annihilation signatures at GeV energies.

\label{sec:summary} dSphs have commonly been regarded as astrophysically ``clean'' targets for which the detection of excess \gammaRays would constitute compelling evidence for particle DM. However, conventional astrophysical emission must be present at some level, and in this Letter we predict the contribution from MSPs. Under the assumption that MSPs in both the MW disk and in dSphs originate mainly from primordial binary systems (in contrast with globular clusters), we scale the LF of Galactic MSPs to the stellar masses of dSphs to quantify their MSP populations. We estimate that MSP emission within the highest stellar mass dSphs, Fornax and Sculptor, is a factor $\sim10$ below the current LAT sensitivity threshold (Figure~\ref{fig:sum_fig}). The MSP emission within ultra-faint dSphs (including targets with the largest J-factors) is several orders of magnitude lower. Moreover, for a DM particle of mass 25~GeV annihilating to $b$ quarks at the thermal relic cross section (consistent with DM interpretations of the Galactic Center excess), the expected \gammaRayHyph emission due to DM exceeds that of the MSP population in all of the dSphs considered here. At the current LAT sensitivity, the more likely source of confusion is a Galactic foreground MSP along the same line of sight, although the probability of this alignment is typically $\lesssim10\%$ per target dSph (Figure~\ref{fig:survival}). The LAT sensitivity to DM annihilation in dSphs is anticipated to improve by nearly an order of magnitude over the coming decade due to increasing LAT exposure, more precise J-factor measurements from deep spectroscopy, and additional dSph targets discovered in optical surveys such as LSST \citep{Charles:2016pgz}. These forecasts are based on a combined likelihood analysis weighting dSphs by their J-factors, as appropriate for DM searches. Since we do not expect that many nearby and high stellar mass dSphs remain to be discovered, the sensitivity to the MSP contribution may not improve as quickly as for DM signals.

1607.06390

1607

1607.03867_arXiv.txt

We analyze extragalactic extinction profiles derived through gamma-ray burst afterglows, using a dust model specifically constructed on the assumption that dust grains are not immutable but respond time-dependently to the local physics. Such a model includes core-mantle spherical particles of mixed chemical composition (silicate core, sp$^2$ and sp$^3$ carbonaceous layers), and an additional molecular component, in the form of free-flying polycyclic aromatic hydrocarbons. We fit most of the observed extinction profiles. Failures occur for lines of sight presenting remarkable rises blueward the bump. We find a tendency in the carbon chemical structure to become more aliphatic with the galactic activity, and to some extent with increasing redshifts. Moreover, the contribution of the molecular component to the total extinction is more important in younger objects. The results of the fitting procedure (either successes and failures) may be naturally interpreted through an evolutionary prescription based on the carbon cycle in the interstellar medium of galaxies.

A powerful way of studying dust properties, and in principle to discriminate among different dust production sources, is through the extinction of stellar light as a function of wavelength. Accurate measurements of InterStellar Extinction Curves (ISECs) are almost exclusively limited to our Galaxy (e.g., \citealt{FM07}, and galaxies in the Local Group (e.g., \citealt{C15}) because at greater distances it becomes impossible to obtain the photometry or spectroscopy of individual stars needed for extinction determinations. Derivations of interstellar dust properties of the interstellar medium in high redshift galaxies focus on the study of either the whole galaxy, or within absorption systems along a single line of sight to distant quasars. In the first case dust extinction is inferred from composite emission spectra, and it is the result of blended events of gas and dust absorption, emission and scattering. Thus, the resultant observed stellar dimming is not produced by a true extinction law, but rather by a dust attenuation profile that is highly dependent on the geometric distribution of the dust, gas and stars (e.g., \citealt{C00}). Direct measurements of the interstellar dust properties in high redshift galaxies performed along a single line of sight to quasars typically suffer from the relatively complex spectral shape of these objects, which makes distinguishing dust from intrinsic color variation very challenging (e.g., \citealt{H13}). Such difficulties are in part relieved observing the deep and point-like line of sight to a Gamma-Ray Burst (GRB) afterglow. GRBs are intense bursts of extremely high energy observed in distant galaxies, the brightest explosive phenomena in the Universe, followed by a longer-lived afterglow emitted at longer wavelengths from X-rays to the radio domain. Bursts that last more than a couple of seconds are known as long-duration GRBs, and are associated with core-collapse supernova explosions. Long duration GRBs have spectra with simple shapes, basically a featureless broken power-law from the X-ray to the near infrared, with a rising and decaying behavior (e.g., \citealt{GAO15}). They have high intrinsic brightness, and generally occur in dense, star forming environments, making these events ideal in the study of extinction properties of interstellar dust on a cosmological scale. To interpret the extinction curves of GRB hosts many studies assume empirical templates, such as the average galactic ISEC (e.g., \citealt{S11}). An alternative approach that does not require such a strong assumption has been put forward by \citet{Li08} and \citet{LL10}, who exploited a parametrized functional form for the normalized (to the visible) extinction, the so-called Drude approach. The description consists of three separate contributions to far-ultraviolet extinction rise, 217.5~nm bump, and near-infrared/visible extinction, respectively, shaped by four dimensionless free parameters. This prescription is similar to the \citet{FM07} parametrization from which differs essentially for the lower number of free parameters, and the more extended (at lower frequencies) domain of applicability. One of the major problems in the study of GRB dust is the paucity of photometric or spectroscopic data, that in some cases are outnumbered by model parameters. In those situations \citet{LL10} measure the goodness of their model fit reducing the $\chi^2$ by the number of data (instead of the free parameters). This is conducive to possibly incomplete pictures of the observables. However, such an unsatisfactory situation is alleviated by the \emph{a priori} choice of the extinction profiles, either templates (less number of parameters, but morphologically stiff) and empirical parametrizations (more flexibility). It is unfortunate that the single parameter \citet{CCM} recipe is only valid for the Milky Way Galaxy (MWG), failing when applied even to our closest neighbourhoods \citep{G03}. The sparsity of data implies that, given the complexity of the underlying physics and chemistry, any physical dust forward model totally overfits the data, and therefore it cannot be applied directly to the observations. Thus, we assume the extinction profile derived by \citet{LL10} using the empirical description inferred via the Drude approach, and (try to) unfold such synthetic description of dust into physically well-grounded properties. The validity of the physical properties we obtain therefore obviously relies on the validity of the assumed Drude model. The sample of GRB host galaxies considered in the \citet{LL10} study are located at redshifts $z \la 2$. Then, the use of both templates and empirical representations relies on the reasonable assumption that the nature of dust has not changed that much in the last 10~Gyr. In fact, gathering together data collected in the last decades, it is evident that, whatever the method exploited to derive dust attenuation, the amount of dust in the Universe increases from the earliest epochs to reach a maximum at $z \sim 1 - 1.5$ \citep{B12}. At about $z \sim 3.5$ the dust content reaches about the same level as measured in the local Universe. Beyond $z = 4$ the dust presence fades with increasing redshifts. The dust attenuation peak appears to be delayed by approximately $2-3$ Gyr with respect to the cosmic Star Formation Rate (SFR) density (e.g., \citealt{B10}). A natural way to reconcile the epoch of the maximum SFR with the dust attenuation peak is to suppose that dust is produced by intermediate-mass long-lived stars in the redshift range $z \sim 1.5 - 3$, linked with the fast drop of dust at $z < 1$ because of the decline of new star formation from $z\sim 2$. Even allowing for a net positive contribution from supernovae to the dust budget (see e.g., \citealt{M15}), we still expect interstellar dust to be mostly composed of the familiar heavy elements, i.e. "silicates" and "carbons" \citep{D03}. How these materials are assembled in a grain is still a matter of debate (see \citealt{CP12} for a brief review). Nevertheless, when re-read in physical terms, a threefold representation such as the one put forward by \citet{LL10} ~\textendash~ or \citet{FM07} ~\textendash~ may be interpreted as the superposition of a population of "classical" grains and a (macro-)molecular component of free-flying PAHs (see \citealt{M13}), contributing to both bump ($\pi^\star \leftarrow \pi$ transitions) and far-ultraviolet rise (low energy side of $\sigma^\star \leftarrow \sigma$ transitions). Our purpose in this work is to explore the sensitivity of normalized ISEC shapes observed during GRB events to the relative abundances of dust grain components, using a dust model in which the various component contributing to the interstellar extinction are related and respond to the local physical and chemical conditions in the interstellar medium of galaxies. In Section \ref{EDM} we present data and modelling procedure, in Section \ref{Res} we show and comment our results, and finally the conclusions are reported in the last Section.

We analyzed a relatively small sample of lines of sight towards GRBs exploiting the [CM]$^2$ dust model, and we interpret the results of the fitting procedure through an evolutionary prescription based on the carbon cycle in the interstellar medium of galaxies. The novel aspect of this work is that in the dust model used, the PAHs and other atoms and fragments of erosion are part of the natural circulation of carbon in the interstellar medium between gas and solid phases. In this latter phase the mantle thickness and its chemical composition are determined by the local physical conditions. It appears that all known types of observed ISECs can be accounted for on the basis of the [CM]$^2$ model (see e.g., \citealt{Z11}). \begin{figure} \includegraphics[width=\hsize]{fig4} \caption{Relation between the mantle aromatic fraction $f_{sp^{2}}$ and the carbon column density locked in PAHs normalized to the visual extinction. The blue triangle indicates the line of sight GRB061121.} \label{four} \end{figure} The results of the fitting procedure shown in the preceding Section may be naturally interpreted in the framework of the carbon cycle outlined above. The results shown in Figure \ref{one} seem to suggest that carbonaceous mantles are increasingly aliphatic with increasing SFR considered as a reliable measure of the galactic activity. Thus, the mantle annealing timescale must be almost always much longer than the average time between mantle shattering events. The annealing time is defined by the competition between energetic processing by ultraviolet radiation and re-hydrogenation by hot H atomic gas, turning aromatic carbonaceous material back into aliphatic. In the MWG, re-hydrogenation was deemed to be by far negligible, but this appears not to be the case in the majority of lines of sight towards the GRB events considered here. Since the SFR increases with increasing redshifts (in the range of redshifts considered in this study), "far" dust loses the aromatic character pervading quiescent galaxies such as the MWG. It is important to realize, however, that dust properties determined using GRBs are not totally representative of the conditions in the host galaxies. Moreover, sometimes host associations are uncertain (e.g., \citealt{H10}). Although the considered lines of sight appears to be sufficiently typical to provide correlation with global galactic properties, this is not possibly true for the line of sight towards the event GRB061121, in which the most MWG-like ISEC within the sample occurs. The prevalent aromatization of dust mantles and the abundance in free-flying PAHs (with respect to the visual extinction) are not consistent with the high SFR~$= 27~M_\odot$~yr$^{-1}$ present in the host galaxy. This discrepancy may be alleviated decreasing the sp$^3 \to$~sp$^2$ conversion time, in competition with re-hydrogenation induced by the large activity of the galaxy as summarized by its SFR. The shape of the ISEC is unaffected as long as the ratio between collisional and photo-darkening rates has kept unchanged. Collisional rates are proportional to the product $n_{\rm H} T_k^{1/2}$, while the annealing rate to the intensity of the local radiation field $\cal{X}$ (e.g., in units of the Habing's field). Thus, the increase of $\cal{X}$ by a factor e.g., hundred is balanced by a corresponding increase in the gas pressure (incidentally consistent with the required larger radiation energy deposition). The ISEC towards GRB061121 appears to be then consistent with the presence of a photon-dominated region, not at all incongruous since GRB events are often associated with star-forming regions. It is worthwhile to note that GRB061121 presents the lowest charge dispersion and one of the highest positive charge of the PAH mixture. More interesting is the well-defined increase of $N_{\rm C}^{\rm {PAH}}/A_V$ with the redshift $z$. Such increase should be read in relative terms, as a major contribution of PAHs with respect to the classical dust grains. If we compare the aromatization fraction of mantles with the fractional concentration of PAHs (Figure \ref{four}) we note the existence of two separate regimes, one characterized by $f_{sp^{2}}$ close to zero, the other one associated with positive $f_{sp^{2}}$ values. In the first case, it is not apparent any clear correlation, while in the second case $f_{sp^{2}}$ grows with $N_{\rm C}^{\rm {PAH}}/A_V$. This latter trends is characteristic of low activity galaxies, as supported by the anticorrelation of $f_{sp^{2}}$ and SFRs. In the first case, the situation is reminiscent of the environments of the LMC and SMC, i.e. the carbon recycling time must be much shorter than the local photo-darkening time. From a physical point of view the uncorrelation of $f_{sp^{2}}$ with the PAH concentration evidences a potential problem in our description of the carbon cycle in galaxies posed by the sharp decline of the source of aromatic material in the diffuse interstellar medium, since most dust mantles are destroyed before they experience substantial annealing. Since PAH formation rates in the cold, carbon-rich winds of evolved stars are too slow, the injection time being approximately 2 Gyr, such component must grow in the diffuse medium or be formed by another yet unknown mechanism. Such a mechanism may be connected to the fate of the plethora of aliphatic fragments released by shattering events. They could evaporate into polyyines, and thereafter quickly photodissociated as proposed by \citet{DW84}, or survive as a population of nanoparticles for a significant time contributing trough $\sigma^\star \leftarrow \sigma$ transitions to the far-ultraviolet rise (but not to the bump). This additional component, missing in the current version of the [CM]$^2$ model may recover the failures in those lines of sight exhibiting extraordinary far-ultraviolet rises. Over time, depending on the competition of collisional (destruction) and radiative (structural transformation) events this population might be (partially) converted into PAHs. The lack of correlation of the 217.5~nm bump with redshifts evidenced by \citet{LL10} may be easily explained by the combined action of the classic and molecular components. In other words within the [CM]$^2$ model an increase in the PAH column density does not translate straightforwardly in an increase of the bump intensity. In conclusion we model the ISECs inferred along a sample of lines of sight to GRB afterglows with a synthetic population of dust grains consisting of core-mantle particles and a collection of free-flying PAHs, providing excellent fits. While this result is not particularly interesting, the retrieval of dust physical properties through the application of the [CM]$^2$ forward model produces a solid base on which to discuss the nature of dust in the local and distant Universe. The major results of these work are: (1) there is tendency in the chemical structure of carbon dust to become more aliphatic with the galactic activity, and to some extent with increasing redshifts; (2) the contribution of the molecular component (PAHs) to the total extinction is more important at early times. Along some lines of sight the lack of any relation between mantle aromatic fractions and relative abundances of PAHs suggest the existence of a transient aliphatic carbon component.

1607.03867

1607

1607.08172_arXiv.txt

{X-ray observations of protostellar jets show evidence of strong shocks heating the plasma up to temperatures of a few million degrees. In some cases, the shocked features appear to be stationary. They are interpreted as shock diamonds.} {We aim at investigating the physics that guides the formation of X-ray emitting stationary shocks in protostellar jets, the role of the magnetic field in determining the location, stability, and detectability in X-rays of these shocks, and the physical properties of the shocked plasma.} {We performed a set of 2.5-dimensional magnetohydrodynamic numerical simulations modelling supersonic jets ramming into a magnetized medium and explored different configurations of the magnetic field. The model takes into account the most relevant physical effects, namely thermal conduction and radiative losses. We compared the model results with observations, via the emission measure and the X-ray luminosity synthesized from the simulations.} {Our model explains the formation of X-ray emitting stationary shocks in a natural way. The magnetic field collimates the plasma at the base of the jet and forms there a magnetic nozzle. After an initial transient, the nozzle leads to the formation of a shock diamond at its exit which is stationary over the time covered by the simulations ($\sim 40-60$~yr; comparable with time scales of the observations). The shock generates a point-like X-ray source located close to the base of the jet with luminosity comparable with that inferred from X-ray observations of protostellar jets. For the range of parameters explored, the evolution of the post-shock plasma is dominated by the radiative cooling, whereas the thermal conduction slightly affects the structure of the shock.} {}

The early stages of a star birth are characterized by a variety of mass ejection phenomena, including outflows and collimated jets that are strongly related with the accretion process developed in the context of the star-disc interaction. In fact, magnetohydrodynamic (MHD) centrifugal models for jet launching \citep{bla82, pud83} indicated that protostellar jets could provide a valid solution to the angular momentum problem (see \citealt{bac02} for a wider description) via vertical transport along the ordered component of the strong magnetic field threading the disk. According to the widely accepted magneto-centrifugal launching scenario \citep{fer06,gom93}, outflows in young stars are driven from the inner portion of accretion discs and dense plasma from the disc is collimated into jets. This scenario was challenged by several models based on different lines of evidence from the observations suggesting the presence of dust in jets due to material from the outer part of the disc \citep{pod09}. Thus several MHD ejection sites probably coexist in young stars, and the difficulty is to determine the relative contribution of each to the observed jet (for a complete review confronting observations and theory see \citealt{fra14}). The general consensus is that magnetic fields play a fundamental role in launching, collimating and stabilizing the plasma of jets. This idea was recently corroborated by scaled laboratory experiments that are representative of young stellar object outflows \citep{alb14}. These experiments revealed that stable and narrow collimation of the entire flow can result from the presence of a poloidal magnetic field. A key observational diagnostic to discriminate between different theories could be the detection of possible signatures of rotation in protostellar jets. Pushing the limits of observational resolution, several authors described the detection of asymmetric Doppler shifts in emission lines from opposite borders of the flow in different objects \citep{bac02,woi05,cof04,cof07,cof15}. However, this possibility is still undergoing active debate. For example, some inconsistencies in rotation signatures were found at different positions along the jet in a UV study of RW Aur \citep{cof12} and in a near-infrared study of DG Tau \citep{whi14} where systematic transverse velocity gradients could not be identified. Alternative interpretations include asymmetric shocking and/or jet precession \citep[e.g.,][]{sok05,cer06,cor09}. Usually jets from young stars are revealed by the presence of a chain of knots, forming the so-called Herbig-Haro (HH) objects. The observations of multiple HH objects showed a knotty structure along the jet axis interpreted as the consequence of the pulsing nature of the ejection of material by the star (e.g. \citealt{rag90,rag07}; \citealt{bon10b,bon10a}; and references therein). After been ejected, the trains of blobs forming the jet move through the ambient medium and they may interact and produce shocks and complex structures that are observed at different wavelengths. In particular, observations showed evidence of faint X-ray emitting sources forming within the jet (e.g. \citealt{pra01,fav02,bal03,pra04,tsu04,gud05,ste09}). The origin of this X-ray emission was investigated through hydrodynamic models which have shown that the observations are consistent with the production of strong shocks that heat the plasma up to temperatures of a few million degrees \citep{bon07,bon10b,bon10a}. In some cases, the shocked features appeared to be stationary and located close to the base of the jet (e.g. HH154, \citealt{fav06}; DG Tau, \citealt{gud05}). In one of the best-studied X-ray jet, HH154, the X-ray emission consisted of a bright stationary component and a fainter and slightly variable elongated component \citep[see][]{fav06}. On the base of hydrodynamic modelling, this source was interpreted as a shock diamond formed just after a nozzle from which the jet originated \citep{bon11}. This scenario was recently supported by scaled laboratory experiments showing the formation of a shock at the base of a laboratory jet that may explain the X-ray emission features observed at the base of some protostellar jets \citep{alb14}. Previous MHD models of protostellar jets were aimed at studying the effect of the magnetic field on the dynamical evolution of HH objects (e.g. \citealt{cer97,sul00,sto00}). These studies therefore have focused mainly on the dynamical aspects and the evolution of the jet rather than on obtaining predictions of the emitted X-ray spectrum. Some authors also studied the dynamics considering the pulsing nature of the jets (e.g. \citealt{rag98,col06}). These models usually include the radiative cooling but not the effects of thermal conduction, being these studies mostly focussing on the optical emission of jets where the thermal conduction efficiency is lower. Here we aim at studying the formation of quasi-stationary X-ray emitting sources close to the base of protostellar jets through detailed 2.5-dimensional (2.5D) MHD simulations. We propose a new MHD model which describes the propagation of a jet through a magnetic nozzle which ram with supersonic speed into an initially isothermal and homogeneous magnetized less dense medium. The MHD model takes into account, for the first time, the relevant physical effects, including the radiative losses from optically thin plasma and the magnetic field oriented thermal conduction. The latter is expected to have some effect on the jet dynamics in the presence of plasma at a few MK as that in the presence of strong shocks. In particular the thermal conduction can play an important role in determining the structure of shock diamonds. We compare the results with observations via the emission measure and the X-ray luminosity, synthesized from the simulations. These studies are important to better understand the structure of HH objects and, more specifically, to determine the jet and interstellar magnetic field structure, and may give some insight on the still debated jet ejection and collimation mechanisms. The organization of this paper is as follows. In Section 2, we describe the MHD model and the numerical setup. The results of our numerical simulations are described in Section 3. Finally, discussion and conclusions are presented in Section 4.

The analysis of the observations of HH154 in three different epochs with Chandra \citep{bal03,fav06,sch11,bon11} revealed a faint and elongated X-ray source at some 150 AU from the protostar that appeared to be quasi-stationary over a time base of $\sim 8$ yr without appreciable proper motion and variability of X-ray luminosity and temperature. Also, the best-studied bright, central X-ray jet of DG Tau seemed to be stationary on timescales of several years on spatial scales of about 30 AU from the central star \citep{sch08}. One of the viable models proposed by \citet{bon11} on the base of hydrodynamic modelling, explained the observations of HH154 describing a shock diamond formed at the opening of a nozzle and producing a X-ray stationary source. They propose a magnetic nozzle as the origin of the shock and they derive a magnetic field strength $B \approx 5$ mG in the magnetic nozzle at the base of the jet. Following the above line of research, we proposed here a new MHD model which describes the propagation of a jet through a magnetic nozzle which rams with a supersonic speed into an initially isothermal and homogeneous magnetized medium. Our MHD model takes into account, for the first time, the relevant physical effects, including the radiative losses from optically thin plasma and the magnetic field oriented thermal conduction. We investigated how the magnetic nozzle contributes to the jet collimation and, possibly, to the formation of a shock diamond at the exit of the nozzle. To this end, we performed an extensive exploration of the parameter space that describes the model. These results allowed us to study and diagnose the properties of protostellar jets over a broad range of physical conditions and to determine the physical properties of the shocked plasma. The different parameters considered are shown in Table~\ref{parameters}. We found that a minimum magnetic field ($\sim1-2$ mG) is necessary to collimate the plasma and form a train of shock diamonds. We selected a magnetic field strength of 5 mG as reference case. For lower values of magnetic field strength, we found that the shock forms at larger distances from the driving star and the beam radius is wider than those usually observed. We found that the stronger is the magnetic field and the lower is the flow velocity, the closer forms the shock to the base of the jet. The summary of results shown in Table~\ref{results} could be very useful in some cases to constrain part of the main physical parameters of protostellar jets from those already known. We derived the physical parameters of a protostellar jet that can give rise to stationary X-ray sources at the base of the jet consistent with observations of HH objects. We found that, in most of the cases explored, quasi-stationary X-ray emission originates from the first shock diamond close to the base of the jet. We obtained shock temperatures of $\sim 1-5 \cdot 10^6$~K, in excellent agreement with the X-ray results of \citet{fav02} and \citet{bal03}. X-ray emission from several HH objects was detected with both the XMM-Newton and Chandra satellites: HH2 in Orion \citep{pra01}, HH154 in Taurus \citep{fav02,bal03}, HH168 in Cepheus A \citep{pra05}, and HH80 in Sagittarius \citep{lop13}. They showed luminosities of $L_\mathrm{X} \approx 10^{29}-10^{30}$ erg s$^{-1}$. Other cases well studied, as Taurus jets of L1551 IRS-5, DG Tau, and RY Tau, showed luminous ($L_\mathrm{X} \approx 10^{28}-10^{29}$ erg s$^{-1}$) X-ray sources at distances corresponding to 30-140 AU from the driving star \citep{fav02,bal03,gud08,sch11,ski11}. The parameters used in our model and the luminosity values synthesized from the model results are in good agreement with those observed: our reference case predicts a luminosity $L_\mathrm{X} \approx 10^{29}$ erg s$^{-1}$ which is in good agreement with observed values. The cases with higher observed luminosities, as HH154, are in good agreement with models M9-M11 (with lower velocities), and M6-M8 (with magnetic field twisting). We obtain the highest luminosity in the model M8, the model with the highest jet rotational velocity. In additon, we investigated the effect of jet rotation on the structure of shock diamonds. Several detections of gradients in the radial velocity profile across jets from T Tauri stars were reported \citep{bac02,cof04,cof07,woi05}. These velocity shifts might be interpreted as signatures of jet rotation about its symmetry axis. For example, \citet{bac02} derived toroidal velocities of the emitting regions between 6 and 15 km s$^{-1}$, depending on position. \citet{woi05} gave higher toroidal velocities in the range 5-30 km s$^{-1}$. They interpreted these velocity asymmetries as rotation signatures in the region where the jet has been collimated but has not yet manifestly interacted with the environment. Considering our model, this region corresponds to the bottom part of the velocity maps in Fig.~\ref{twist_comp} with values $5-30$ km s$^{-1}$ for model M6, $6-40$ km s$^{-1}$ for model M7 and $7-50$ km s$^{-1}$ for model M8. Alternative interpretations include asymmetric shocking and/or jet precession \citep[e.g.,][]{sok05,cer06,cor09}. We derived the angular momentum loss rate at the base of the jet for the three models, M6, M7 and M8 as $\dot J_{\mathrm{j},\omega} = \int \rho_{\mathrm{j}} v_{\mathrm{j}} v_\varphi r \, \mathrm{d}A$, where $\rho_{\mathrm{j}}$ and $v_{\mathrm{j}}$ are the mass density and jet velocity, respectively, $v_\varphi$ is the rotational velocity, $r$ is the radius and $\mathrm{d}A$ is the cross sectional area of the incoming jet plasma. The values obtained range between $1.63 \cdot 10^{-5}$ $M_{\odot}$ yr$^{-1}$ AU km s$^{-1}$ for model M6 and $3.25 \cdot 10^{-5}$ $M_{\odot}$ yr$^{-1}$ AU km s$^{-1}$ for model M8. These values refer to the flux carried away by the jet and they are in good agreement with that estimated by \citet{bac02}, namely, $3.8\cdot 10^{-5}$ $M_{\odot}$ yr$^{-1}$ AU km s$^{-1}$. Thus, our model could be a useful tool for the investigation of the still debated rotation of the jets. We also explored the role of thermal conduction and radiative losses in determining the structure of the shock diamonds by performing some extra simulations calculated with these physical processes turned ``on'' or ``off''. We found that the radiative losses dominate the evolution of the shocked plasma in the diamonds. The main effect is to decrease the $EM$ of plasma with temperature larger than $10^6$~K, thus reducing its X-ray luminosity. The thermal conduction plays a minor role slightly contrasting the cooling of hot plasma due to radiative losses. The comparison between our model results and the observational findings showed that the model reproduces most of the physical properties observed in the X-ray emission of protostellar jets (temperature, emission measure, X-ray luminosity, etc.). Thus we showed the feasibility of the physical principle on which our model is based: a supersonic protostellar jet leads to X-ray emission from a stationary shock diamond, formed after the jet is collimated by the magnetic field, consistent with the observations of several HH objects. We conclude that our model provides a simple and natural explanation for the origin of stationary X-ray sources at the base of protostellar jets. Therefore, the comparison of our MHD model results with the X-ray observations could provide a fundamental tool to investigate the role of the magnetic field on the protostellar jet dynamics and X-ray emission.

1607.08172

1607

1607.06735_arXiv.txt

Combining information from weak sources, such as known pulsars, for gravitational wave detection, is an attractive approach to improve detection efficiency. We propose an optimal statistic for a general ensemble of signals and apply it to an ensemble of known pulsars. Our method combines $\mathcal F$-statistic values from individual pulsars using weights proportional to each pulsar's expected optimal signal-to-noise ratio to improve the detection efficiency. We also point out that to detect at least one pulsar within an ensemble, different thresholds should be designed for each source based on the expected signal strength. The performance of our proposed detection statistic is demonstrated using simulated sources, with the assumption that all pulsar ellipticities belong to a common (yet unknown) distribution. Comparing with an equal-weight strategy and with individual source approaches, we show that the weighted combination of all known pulsars, where weights are assigned based on the pulsars' known information, such as sky location, frequency and distance, as well as the detector sensitivity, always provides a more sensitive detection statistic.

Pulsars are believed to be rapidly rotating neutron stars (NSs) that can emit continuous gravitational wave (GW) radiation if their mass distributions are asymmetric~\cite{1979PhRvD..20..351Z}. Observations from first-generation GW detectors have placed upper limits on the amplitude of these GWs from the known galactic millisecond pulsars. This in turn allows constraints to be placed on the ellipticities of these NSs~\citep{2010ApJ...713..671A}. With the advanced detector era having recently begun with Advanced LIGO~\cite{2015CQGra..32g4001L} in operation and Advanced Virgo~\cite{2015CQGra..32b4001A}, and KAGRA~\cite{Aso:2013} close behind, we will soon be able to make observations of these sources with significantly increased sensitivity. For each pulsar with known sky location and assumed GW phasing (as inferred from arrival times of its radio pulses), time and frequency-domain matched-filtering approaches~\cite{2005PhRvD..72j2002D,Pitkin:2011cl,1998PhRvD..58f3001J,2014CQGra..31p5014A,2014ApJ...785..119A} are commonly applied. The former has been used within the LIGO-Virgo Collaboration for the known pulsar searches and applies a Bayesian marginalization strategy to the unknown system parameters~\cite{2005PhRvD..72j2002D}. The latter, frequency-domain approach, known as the $\mathcal{F}$-statistic ~\cite{1998PhRvD..58f3001J} performs an analytical maximization of the likelihood over the unknown parameters of each pulsar and it is this method that we make use of for the remainder of this paper. Combining sources to improve detection probability is an attractive approach to weak signal detection (e.g.~detecting NS ellipticity from analysis of the GW stochastic background~\cite{2014PhRvD..89l3008T} and detecting gravitational wave memory using binary black hole mergers~\cite{2016arXiv160501415L}). Since GW detectors currently study ${\sim}200$ known pulsars, the existing detection strategy for this relatively large ensemble can be viewed as trying to detect each one separately, and then waiting for the first detection to appear. This is certainly the most obvious strategy to take, but not obviously the most optimal. Cutler and Schutz (CS)~\cite{2005PhRvD..72f3006C} proposed an alternative: first sum the $\mathcal{F}$ statistic from each pulsar, and then use that sum as a new detection statistic. In this initial study, CS used an equal weight for all the pulsars to be combined. One issue with this approach is that including pulsars which are likely to emit relatively weak GWs decreases the signal-to-noise ratio (SNR) of the combined statistic. As indicated in their paper, the SNR of the combined statistic decreases if the detection ensemble includes weak sources where the squared SNR is less than half of the average squared SNR for all observed pulsars. Therefore, to more efficiently detect GWs from an ensemble of all known pulsars, it seems sensible to investigate the effects of giving nonequal weights to the pulsars within the ensemble. In this paper, we generalize the idea proposed by CS, by considering the prior distribution of GW strengths from the pulsars within the ensemble. After a brief introduction to pulsar GW emission and the $\mathcal{F}$-statistic, we apply the general theory of hypothesis testing, and obtain a Neyman-Pearson criterion for detecting GWs from an ensemble of pulsars. This leads to an optimal detection statistic, which in idealized situations (i.e., when our prior knowledge of the signal and our model for the noise are an accurate representation of reality) provides the highest detection probability with a given false-alarm probability. As we show, this statistic can in some cases be approximated by linearly combining $\mathcal{F}$-statistic values from the ensemble of pulsars with appropriate weights. We assume that the ellipticities of pulsars follow a common (yet unknown) intrinsic distribution and that the orientation of their rotation axes is isotropically distributed. We then draw on our knowledge of their sky location, distance from the Earth, and their rotation frequency to construct prior distributions on the expected GW amplitudes from our known pulsars. Since the intrinsic ellipticity distribution remains unknown, we model it as a simple exponential distribution, but perform tests using both exponential and Gaussian distributions. This paper is organized as follows. In Sec.~\ref{review}, we briefly review the form of GW emission from individual pulsars and the $\mathcal{F}$ statistic; in Sec.~\ref{f_statistic}, we introduce the optimal statistic for a general ensemble of pulsars and discuss how it may apply to a set of pulsars in idealized situations; in Sec.~\ref{sim}, we test our statistic on two possible intrinsic distributions of pulsar ellipticity. We summarize our main conclusions in Sec.~\ref{res}.

\label{res} We have proposed a novel weighted-combination detection statistic for GWs from an ensemble of known pulsars. The aim of this approach is to improve the detection efficiency of GWs over that of individual pulsar detection based on the ${\mathcal F}$-statistic applied to single pulsars. The general argument behind the combination detection strategy is that a group of sources should be more detectable than an individual one if they share certain characteristics. We have shown that our general optimal statistic for the weighted combination of GW signals outperforms all other approaches. We have shown that to more efficiently detect GW signals emitted from a ensemble of pulsars, each source within the ensemble could be assigned a different detection statistic threshold based on the expected signal strength. Furthermore, by assuming that the SNRs of all sources are constant or follow exponential distributions, we have shown that the linearly weighted-combination statistic is very close to being optimal and is robust to the choice of prior SNR distributions. These analytic and simple Monte Carlo test predictions are consistent with results obtained from simulations of known pulsars. We have also used the ROC function to determine the sensitivity of a range of possible search strategies where the detection probability between approaches is compared as a function of false-alarm probability. To demonstrate the performance of the new weighted-combination detection method for the Advanced detectors era, we have compared the detection efficiency of the linearly weighted-combination method versus the equal-combination and individual detection method. We have done this by simulating GW signals emitted from the 195 known pulsars within the sensitive frequency band of Advanced LIGO and Virgo. We assume that the intrinsic pulsar parameter ellipticity $\epsilon^2$ follows a common distribution in these simulations. The true form of the ellipticity distribution and its associated parameters are unknown. We have chosen to use both exponential and Gaussian distributions with mean values corresponding to ellipticities $\epsilon \sim 10^{-8}$, a value consistent with the initial GW era nondetection of pulsar signals and a possible advanced era detection. In general, the combination methods return better detection efficiency than a method that simply considers the closest or brightest pulsar. Being consistent with results of simple Monte Carlo tests, the most efficient method in simulations for known pulsars involves combining all known pulsars with weights $\propto \overline{\rho}$, the expected value of the optimal SNR of each pulsar. For the specific case where $\overline{\epsilon} \sim 1.5\times 10^{-8}$, for one year observation of the Advanced detector network, we find that $P_{\rm DE}{\sim}0.95$ given $P_{\rm FA}=0.0001$. In this case, the improvement by our proposed combined method could be up to a factor of $\sim 4$ compared with other methods. These results are consistent with the case of taking into account the measurement errors of pulsar distances. An important feature of the proposed combination method is that it is very flexible. Using the new method it is simple to include more observed pulsars or updated source information (e.g. distance or orientation parameters) , without recalculating any individual detection ${\mathcal F}$-statistic values. However, we would expect that a fully Bayesian approach for combining all known pulsars may be more sensitive albeit at an increased computational cost. The flagship known pulsar analysis within the GW community is a Bayesian approach~\cite{2005PhRvD..72j2002D,Pitkin:2011cl,2009CQGra..26t4013P,2014CQGra..31f5002W}. We note that it is likely that a comprehensive Bayesian approach to combining all known pulsars into a single analysis may produce a truly optimal result. Besides all of the information discussed above, one could also consider the uncertainty of the major assumption (model) of this work: that all pulsars' ellipticity values follow a common but unknown distribution. A hierarchical Bayesian approach would allow us to naturally investigate the true priors governing the distribution. In this case the form of the prior would be represented as a possible model and the parameters governing that distribution would be the ``hype'' parameters of that model. We could also apply Bayesian model selection to distinguish between different prior distributions e.g. exponential vs Gaussian or power law, etc... However, it is unclear how constraining such an analysis would be and we hope to tackle this problem in future studies. Beyond the detection of GWs emitted by a ensemble of pulsars, the posterior probability of all parameters could be output from a Bayesian approach. In future studies we hope to investigate such a Bayesian application to the detection of GWs from the ensemble of known pulsars.

1607.06735

1607

1607.03326_arXiv.txt

The accretion mechanism producing the short flares observed from the Supergiant Fast X-ray Transients (SFXT) is still highly debated and forms a major part in our attempts to place these X-ray binaries in the wider context of the High Mass X-ray Binaries. We report on a 216 ks \inte\/ observation of the SFXT \src\ (August 24-27, 2014) simultaneous with two fixed-time observations with \xmm\ (33ks and 20ks) performed around the putative periastron passage, in order to investigate the accretion regime and the wind properties during this orbital phase.\\ During these observations, the source has shown luminosity variations, from $\rm \sim4\times10^{34} erg~s^{-1}$ to $\rm \sim10^{36} erg~s^{-1}$, linked to spectral properties changes. The soft X-ray continuum is well modeled by a power law with a photon index varying from $\sim$1.2 up to $\sim$1.7 and with high values of the column density in the range $\rm \sim 2-4\times10^{23} cm^{-2}$. We report on the presence of iron lines at $\sim$ 6.8-7.1 keV suggesting that the X-ray flux is produced by accretion of matter from the companion wind characterized by density and temperature inhomogeneities.

\label{sec:intro} \src\/ was first reported as an unidentified transient source in the third \inte\/ IBIS/ISGRI survey with a flux of 4 mCrab and 3.2 mCrab in the energy range 20-40 keV and 40-100 keV, respectively (Bird et al. 2007, Bodaghee et al. 2007). The \sw\/ XRT follow-up observations performed during a flare (on 2009 June 10) showed that the X-ray spectrum of the source was, at that time, well described by an absorbed power law model with a $N_{\rm H}$$\simeq$8$\times$10$^{22}$~cm$^{-2}$ in excess of the expected Galactic value and a photon index $\Gamma$$\simeq$0.56 (Grupe et al. 2009). On the basis of its transient and recurrent nature, its short and intense flares and a dynamic range of $\sim10^2$, this source has been classified as a candidate SFXT (Fiocchi et al. 2010). An orbital period corresponding to $\sim$ 10 days has been derived by Corbet et al. (2010) making use of \sw\//BAT data. \src\/ was observed by \xmm\/ on 2011 February 20 (corresponding to an orbital phase of $\sim$0.1) for a total exposure time of $\sim$22 ks (Bozzo et al. 2012). The analysis of these data showed a flux variation of a factor $\sim$10 without significant variation of the spectral parameters (N$_H$ and $\Gamma$). The average spectrum was well fitted with an absorbed power law model with a column density of $\rm \sim 17.5\times10^{22}cm^{-2}$, a photon index of $\sim 1.5$ and unabsorbed 2-10 keV flux of $\rm 1.7\times10^{-11}erg~cm^{-2}s^{-1}$. The source was also within the field of view of \sax\/ in 1998: the MECS X-ray data showed a frequent microactivity typical of the intermediate state of SFXT and a weak flare with a duration of $\sim$4.6 ks (Fiocchi et al. 2013). During these observations the photon index of the power law model remained constant while the absorption column density was highly variable, spanning from $\sim$3 to $\sim$20 $\rm \times10^{22}cm^{-2}$ across the transition from the low emission level ($\rm F_{2-10keV}\sim3\times10^{-12}erg~cm^{-2}~s^{-1}$) to the peak of the flare ($\rm F_{2-10keV}\sim10^{-10}erg~cm^{-2}~s^{-1}$). Romano et al. 2013 reported on the spectral analysis of a flare that occurred on 2009 June 10, and has observed with the \sw\//XRT instrument. During the brightest X-ray emission (unabsorbed flux of $\rm 4.2\times10^{-10}erg~cm^{-2}~s^{-1}$) the photon index was $\sim$0.65 and the column density was $\rm \sim$9$\times$10$^{22}$~cm$^{-2}$ in excess of the Galactic one. IR observations allowed confirmation of the nature of the companion as a O8I spectral type star (Coleiro et al. 2013) and determined the source distance of 7.2$\pm$0.3 kpc (Persi et al. 2015). At soft X-ray energies a long term monitoring (2011-2013) with \sw\//XRT allowed a detailed study of the emission outside the bright outbursts, identifying two low emission levels, both well described with a power law model: the first has a photon index of $\sim1.35$ and a column density of $\rm \sim13.6\times$10$^{22}$ cm$^{-2}$ at an observed flux of $\rm F_{2-10keV}\sim16\times10^{-12}erg~cm^{-2}~s^{-1}$, while the second one is fitted with a photon index of $\sim0.3$ and a column density of $\rm \sim1.5\times$10$^{22}$ cm$^{-2}$ at an observed flux of $\rm F_{2-10keV}\sim1.1\times10^{-12} erg~cm^{-2}~s^{-1}$. These observations allowed an estimate of the lower limit of the dynamic range in this source of $\sim$750 (Romano et al. 2014b).\\ This source has been reported as a SFXT characterized by an intermediate orbital period and a low flux variability in the review of Walter et al. (2015). Outbursts from \src\/ usually occur near the periastron passage, at a restricted phase range of its orbital period (10.068$\pm$0.002\ days, Fiocchi et al. 2013), allowing the use of fixed-time observations. In this paper we report on the spectral results for two \xmm\/ (Jansen et al. 2001) observations, performed quasi-simultaneously with a long \inte\/ (Winkler et al. 2003) observation performed at periastron (phase = 0.5).

The \xmm\/ observations have allowed us to perform an in-depth investigation of the transient source \src\/ at two different orbital phases. We also followed the source variability in detail, revealing changes in its spectral shape. The photon index shows significant variations, with values ranging from $\sim$1.2 during high flux intervals (states EGI and A$_2$) to $\sim$1.7 during a low state % (see Table 1 and Fig.~\ref{fig:contour}, top panel) in the first \xmm\/ observation (at an orbital phase of 0.4). This spectral softening at low luminosity is in agreement with the standard behavior observed in SFXTs. Indeed X-ray spectra during very strong flares are usually well described by a flat power law ($\Gamma\sim0-1$) while the photon index increases to values of $\Gamma\sim1-2$ at lower luminosities of $\rm \sim10^{33-34} erg~s^{-1}$ (see Romano et al. 2011, 2014b, Sidoli et al 2011). Changes to photon index corresponding to changing luminosity are not observed during the second \xmm\/ observation at orbital phase $\sim\/$0.6. In Fig.~\ref{fig:plotPAR} we show the spectral index (top panel) and the column densities (bottom panel) against unabsorbed fluxes in the 2-10 keV energy range, using spectral analysis of time-selected states reported in Table 1 and Table 2 and archival results. % In this way we can track the unabsorbed flux variations by two order of magnitude. We show parameter values for the first \xmm\/ observation in black points, for the second \xmm\/ observation in red, for the \sax\/ data in green (from Fiocchi et al. 2013), for the \sw\/ XRT data in blue (from Romano et al. 2013) and \xmm\/ data from Bozzo et al. (2012) in magenta. Since these data all cover similar energy ranges, the derived N$\rm _H$ values should be comparable. Squares with circles indicate spectral parameters measured during average low emission levels, squares alone indicate parameters measured during an active period. The analysis of different flux states confirms changes in the column density, previously observed using \sax\/ data (Fiocchi et al. 2013) and highlights the variation of spectral index: from the top panel of Fig.~\ref{fig:plotPAR} it is clear that the photon index shows significant variations without any clear correlation with unabsorbed flux. The bottom panel of Fig.~\ref{fig:plotPAR} shows that the N$_H$ values at flux lower than $\rm \sim 10^{-11} erg~s^{-1} cm ^{-2}$ and greater than $\rm \sim4\times 10^{-10} erg~s^{-1} cm ^{-2}$ do not confirm the linear correlation between flux and column density observed in this object in the past (Fiocchi et al. 2013), when we consider two orders of magnitude in flux. In a restricted range of fluxes, from $\rm \sim 10^{-11} erg~s^{-1} cm ^{-2}$ to $\rm \sim2\times 10^{-10} erg~s^{-1} cm ^{-2}$ this correlation still persists. To investigate the possible column absorption and the orbital phase correlation, we report in Fig.~\ref{fig:phase} the column density against the orbital phase. We display column density values according with colours in Fig.~\ref{fig:plotPAR}. Squares with circles indicate spectral parameters measured during average low emission levels, squares alone indicate parameters measured during an active period. Spectral parameters values obtained from spectra with long exposure time (covering $\sim$ one phase) are not included in this plot. Fig.~\ref{fig:phase} shows that there are significantly higher values of the column density during the active time interval corresponding to the orbital phase of 0.4. We note that the average low emission level (squares with circles in Fig.~\ref{fig:phase}) show a maximum at phase $\sim\/$0.4 and a minimum at phase $\sim$ 0.95. These data show there could be two levels of the density variations: the first corrisponding to the average low emission states and the second considering the active periods. During the low emission levels (circles of Fig.~\ref{fig:phase}), the $\rm N_H$ follows the same behavior that the IBIS intensity (23-50 keV) has versus the orbital phase, with a maximum value at phase $\sim$0.5 and lower values at phases $\sim$0.1 and $\sim$1.0. During the active periods, the $\rm N_H$ variations are not correlated with the orbital phase and could indicate changes in the accreting material on the neutron star. Obtained $\rm N_H$ values rule out that the observed low emission level (flux lower than $\sim 10^{-11} erg s^{-1} cm ^{-2}$) can be due to obscuration of the emitting region by circumstellar material, as infact there are values of the column density during low emission level consistent with the ones during the active time intervals (see column density of state with rate lower than 0.4 c/s and A$_2$ state). This behavior suggests that the $\rm N_H$ values in the low emission levels could be an indication of the matter distribution along the orbit, while additional mechanisms come into play during flaring activity. \\ The iron fluorescence lines show an interesting evolution: the centroid is at $\sim$6.8 keV when the source is in the JKL state while shifts up to $\sim$7.1 keV at higher fluxes. As the iron line centroid is correlated with 2-10 keV unabsorbed flux (see Table 1 and Table 2), line emissions at $\sim$ 6.8 keV and $\sim$ 7.1 keV could come from highly ionised iron ions: the ionization level is higher than Fe$_{XXV}$ and Fe$_{XX}$ for the state A$_2$ and JKL, respectively (Kallman et al. 2010). Since the theory predicts that the iron line intensity ratio $\rm I_{K_{\beta}}$/$I_{K_{\alpha}}$ is $\sim0.13~ph~cm^{-2}s^{-1}$ (Kallman et al. 2010), the lack of a strong iron line at 6.4 keV during time intervals A$_2$ and JKL exclude that the observed iron line at $\sim$ 6.8-7.1 keV can be fluorescence iron line K$_{\beta}$, not respecting this iron line intensity ratio. This behavior suggests that the X-ray flux produced by accretion onto the neutron star partly ionized the clump matter. \\ The limited statistics of our \src\/ \xmm\/ data prevent us from studying the expected linear correlations between the continuum flux and the iron line flux or between the Fe equivalent width and the continuum parameters (N$_H$ and luminosity), as reported by Gimenez-Garcia et al. (2015) and Torrejon et al. (2010). This work has shown a complex picture that is compatible with accretion from an inhomogeneus wind (in’t Zand 2005, Walter \& Zurita Heras 2007). The transient emission produced by accretion of matter from the companion wind indicates change in the wind density and temperature, not clearly correlated with the orbital phase. The inhomogeneities in the accreting material are able to give a physical interpretation of the short flares observed in both the \xmm\/ data and the previous \sax\/ ones (Fiocchi et al. 2013).\\ Conversely, the clumpy wind model alone is not able to explain the few days long flare observed with \inte\/ from this source (Fiocchi et al. 2010). This evidence confirms that additional mechanisms are nedeed to explain the extreme variability seen in SFXT (Bozzo et al. 2014, Lutovinov et al. 2013). \\ The two proposed additional mechanisms to inhomogeneus wind are the quasi spherical accretion model (Shakura et al. 2012, 2014) or the centrifugal and/or magnetic gating accretion (Bozzo et al. 2008, 2016). At this stage, for \src\/ both mechanisms cannot be ruled out, indeed: \\ 1) the theory of wind accretion in HMXB hosting a magnetic neutron star with transitions driven by centrifugal and magnetic barrier (Bozzo et al. 2008, 2016) requires an high magnetic field to explain the observed dynamic range (greater than $\sim10^{14} G$): unfortunately, the magnetic field in \src\/ is unknown. \\ 2)the quasi spherical accretion model (Shakura et al. 2012, 2014) concerns the accretion onto slowly rotating X-ray pulsars: the spin period of \src\/ is unknown. Furthermore this theory predicts two regimes of accretion at the critical X-ray luminosity value of $\rm \sim4\times10^{36} erg~s^{-1}$. The present \xmm\/ data and the previous \sw/XRT results (Romano et al. 2013) allowed to extend the studied luminosity range, spanning from $\rm \sim6\times10^{33} erg~s^{-1}$ to $\rm \sim3\times10^{36} erg~s^{-1}$. Unfortunately, the investigated luminosity values are always lower than critical value % preventing to study X-ray behaviour at very high luminosity. Finally, we note that the accretion radius and the magnetospheric radius are highly sensitive to variations in the wind velocity and this wind velocity can significantly drop or be completely halted close to the neutron star when the matter is ionizated (Krticka et al. 2015, Ducci et al. 2010), making the comparison data-model complicated.

1607.03326

1607

1607.06818_arXiv.txt

Modeling the stochastic gravitational wave background from various astrophysical sources is a key objective in view of upcoming observations with ground- and space-based gravitational wave observatories such as Advanced LIGO, VIRGO, eLISA and PTA. We develop a synthetic model framework that follows the evolution of single and binary compact objects in an astrophysical context. We describe the formation and merger rates of binaries, the evolution of their orbital parameters with time and the spectrum of emitted gravitational waves at different stages of binary evolution. Our approach is modular and allows us to test and constrain different ingredients of the model, including stellar evolution, black hole formation scenarios and the properties of binary systems. We use this framework in the context of a particularly well-motivated astrophysical setup to calculate the gravitational wave background from several types of sources, including inspiraling stellar-mass binary black holes that have not merged during a Hubble time. We find that this signal, albeit weak, has a characteristic shape that can help constrain the properties of binary black holes in a way complementary to observations of the background from merger events. We discuss possible applications of our framework in the context of other gravitational wave sources, such as supermassive black holes.

The recent detection of the gravitational wave (GW) sources GW150914 \citep{2016PhRvL.116f1102A} and GW151226 \citep{2016ApJ...818L..22A} opened the era of gravitational wave astronomy, and has provided the first direct confirmation of the existence of black holes (BHs), and in particular binary BHs (BBH) that merge within the age of the Universe. Based on the rate of BBH mergers inferred from these detections \citep{2016arXiv160203842A,2016arXiv160604856T} many more sources are expected to be discovered in the second and third Advanced LIGO observing runs \citep{2016arXiv160604856T}. Ground-based interferometers such as Advanced LIGO, which is already gathering data, as well as VIRGO and KAGRA which are expected to become operational in the near future, are sensitive to gravitational waves in the frequency range $\sim 30 - 1000$ Hz, and are designed to detect mergers of BBH and binary neutrons stars (NSs), as well as the gravitational wave background from unresolved mergers of these binary compact objects \citep{2016arXiv160203847T}. Pulsar timing array (PTA) networks \citep{2013CQGra..30v4009K,2015MNRAS.453.2576L} may detect the GW background produced by merging super-massive BHs (SMBH), topological defects such as cosmic strings, and individually resolvable SMBHs in the frequency range $\sim 10^{-9}-10^{-8}$ Hz. The frequency ranges $\sim 10^{-4} - 10^{-1}$ Hz and $\sim 10^{-1} - 10$ Hz will be explored by the space-based eLISA \citep{PhysRevLett.116.231101} and DECIGO \citep{2011CQGra..28i4011K} observatories, respectively, planned to be launched in the next decade. The frequency coverage that will be attained when all of these observatories are operational suggests the possibility of multi-wavelength GW astronomy \citep{2016PhRvL.116w1102S,2016arXiv160501341N}, where the same source can be observed by different observatories as the merger proceeds. Detections of individual sources, such as GW150914 and GW151226 are invaluable in studying the properties of compact objects and constraining gravity under extreme conditions. The masses and spins of observed BBH already begin to inform astrophysical models of BH formation \citep{2016ApJ...818L..22A,2016arXiv160204531B,2016arXiv160604889A} and future detections may provide information on the equation of state of NSs \citep{2015PhRvD..92b3012A}. Moreover, the waveforms of individual merger events allow to place stringent constraints on extensions to General Relativity \citep{2016arXiv160604856T}. Another component that can be detected with GW observatories is the gravitational wave background from unresolved, merging and inspiraling sources. This component will allow to study the compact object population from a different viewpoint, in particular by constraining the distribution of the binary parameters and their formation mechanisms \citep{2016arXiv160203847T}. The background from unresolved binary compact objects has in general three components: (a) the signal emitted by core-collapse supernovae (SNe) \citep{1985PhRvL..55..891S,1999MNRAS.303..247F}; (b) the contribution from objects that are about to merge (usually referring to inspiral, merger and ringdown phases) \citep{2008PhRvD..77j4017A,2011PhRvD..84h4037A} and (c) the contribution from inspiraling binaries which do not merge during a Hubble time but which still emit gravitational radiation, resulting in a circularisation and shrinking of their orbit. Component (b) is perhaps the most extensively studied, both in the domain of stellar-mass BBH and NSs as well as SMBH in view of its importance for predicting the signal of merger events such as GW150914 and GW151226. While the waveform of a single isolated merge is well understood, many uncertainties remain, in particular regarding the merger rates (which are related to the properties of the progenitors). GW background from stellar-mass BBH is expected to be detected by Advanced LIGO \citep{2016arXiv160203847T}, while the signal from SMBH is beginning to be constrained by PTA experiments and will be further probed by the eLISA satellite \citep{2015MNRAS.453.2576L,PhysRevLett.116.231101}. Finally, the GW signal from SN collapse is difficult to estimate due to uncertainties in the collapse mechanism \citep{2013ApJ...766...43M,2013ApJ...768..115O,2015PhRvD..92f3005C}. While the contribution from merging compact binary systems is dominant, most binaries are not expected to merge within a Hubble time. They will, however, emit gravitational radiation while slowly approaching each other and, depending on the merger rate, the source mass and redshift distribution and the initial orbital parameters, might be detectable with future generations of GW observatories. It is important to stress that the evolution of massive stars and compact objects is affected by their environment. Interactions with other stars in a dense star cluster might be an inportant channel for creating heavy stellar-mass BHs \citep{2016arXiv160604889A} and the growth rate of SMBH is clearly related to the properties of its host galaxy (e.g. \citep{2012Sci...337..544V}). The complexity of the different astrophysical processes involved in producing the GW background and their vastly different length and time scales lead to great difficulties in constructing a model that can be easily tested against upcoming data. Moreover, it is often challenging to estimate the relative importance of the various uncertainties involved. In this paper we develop a general framework for calculating the GW background from binary compact objects in an astrophysical context. As will be discussed below, many of the ingredients of this calculation are highly uncertain, therefore we tried to construct a modular approach to the problem, allowing to narrow down on one kind of uncertainty at a time. We then apply this approach to inspiraling stellar-mass BBH and binary NSs that have not merged during a Hubble time. The core of our method is in describing the number density of binary systems in terms of the continuity equation in the space of orbital parametres of the binary. A similar approach was used by Refs. \citep{1994MNRAS.268..841B,1995MNRAS.274..115M,2001MNRAS.327..531I} to study the GW background from high-mass binary pulsars in our Galaxy. In this work we go beyond the steady-state solution assumed in these studies and treat multiple source classes. This paper is structured as follows: section \ref{sec:Model} describes our synthetic approach: we start with some basic definitions in section \ref{sec:general}. We then define the number densities and formation rates of single and binary compact objects and derive the equations for the evolution of binary orbital parameters in section \ref{sec:evo_nums}. We discuss our complete synthetic model in section \ref{sec:complete_set}. Section \ref{sec:GWback} is an application of our approach to the calculation of several GW backgrounds in the context of a particular astrophysical model. In section \ref{sec:spectrum} we review the GW energy spectrum from inspiraling and merging sources, in section \ref{sec:astro} we outline our astrophysical model and in sections \ref{sec:ins_BH} and \ref{sec:ins_NS} we calculate the GW background from inspiraling and merging stellar-mass BBH and inspiraling binary NSs. We conclude in section \ref{sec:discussion}.

\label{sec:discussion} In this paper we developed a synthetic approach that allows to model the evolution of binary compact objects and the GW background they produce. We described the evolution of the number density of binaries and their interactions in the space of orbital parameters with a set of continuity equations providing also for a source and sink terms due to formation and mergers of binary systems. While we used a specific astrophysical model to estimate the GW background from inspiraling and merging BBH, our approach is modular and any ingredient can be updated or tested against competing models. In particular, we can study the rate of formation of compact objects resulting from different astrophysical models, the rate of binary formation, the distribution of initial orbital parameters and the time evolution. We use our approach to calculate for the first time the GW background from inspiraling BBH. The signal we predict is very weak and, moreover, is not in the frequency range of any of the current or planned GW observatories. Nevertheless, the characteristic shape of the transition from inspiral to merger dominated signal might provide very interesting constraints on the entire population of BBH. Since the majority of BBH are not expected to merge within a Hubble time and are thus beyond the observational capabilites of ground-based interferometers such as Advanced LIGO and VIRGO, the signal we predict offers a unique handle on the properties of this population. We stress that this work does not include an exhaustive treatment of the various astrophysical effects, such as BH production mechanisms, the effects of rotation, binary co-evolution and possible influence of dense environments. Our treatment of the various GW backgrounds is also far from complete, as we did not include the contribution from SN collapse nor merging NSs. Moreover, we assumed that all the binaries consist of equal-mass objects and did not calculate the background due to inspiraling and merging BH-NS binaries. An extensive study of these topics and the estimate of the associated uncertainties are left for future work. Finally, we expect this framework to be useful for different classes of GW sources not discussed here, such as SMBH binaries during the later stages of the merger. The signal from inspiraling SMBH which take longer than the age of the Universe to merge falls in the frequency range accesible with PTA ($10^{-9}-10^{-8}$ Hz), while merging SMBH will be observable with eLISA. The evolution of the eccentricity of these systems may be affected by their environment and some of the SMBH binaries may enter the observable frequency band while still retaining a non-negligible eccentricity which will have an imprint on the GW background \citep{2007PThPh.117..241E,2008ApJ...686..432S,2016ApJ...817...70T}. These questions can be treated within the formalism described in this paper and we plan to study them in future work.

1607.06818

1607

1607.03110_arXiv.txt

We propose a new mechanism for thermal dark matter freezeout, termed \emph{Co-Decaying Dark Matter}. Multi-component dark sectors with degenerate particles and out-of-equilibrium decays can co-decay to obtain the observed relic density. The dark matter density is exponentially depleted through the decay of nearly degenerate particles, rather than from Boltzmann suppression. The relic abundance is set by the dark matter annihilation cross-section, which is predicted to be boosted, and the decay rate of the dark sector particles. The mechanism is viable in a broad range of dark matter parameter space, with a robust prediction of an enhanced indirect detection signal. Finally, we present a simple model that realizes co-decaying dark matter.

The nature of dark matter (DM) is one of the most important open questions in physics. The possibility that dark matter is a thermal relic with mass around the weak scale is intriguing, but has been under significant experimental pressure from direct detection~\cite{Cushman:2013zza,Akerib:2015rjg,Aprile:2015uzo} and at the LHC~\cite{Askew:2014kqa}. This motivates the study of models which are not constrained by these searches, but can still be discovered by indirect detection, where limits are weaker and have made rapid progress in recent years~\cite{Buckley:2013bha}. Mechanisms for thermal dark matter freezeout usually rely on the DM remaining in chemical and thermal equilibrium with the Standard Model (SM) bath while non-relativistic, which leads to depletion of DM through Boltzmann suppression. In this work we consider the possibility that part of the dark sector decays out of equilibrium with the SM. This delays the exponential suppression of the DM density well beyond the point where the DM candidate becomes non-relativistic. The mechanism, which we refer to as \emph{Co-Decaying Dark Matter}, has the following properties: \begin{enumerate} \item The dark sector has decoupled from the SM before it becomes non-relativistic. \item The lightest dark sector particle decays into the SM out of equilibrium. \item The dark sector contains additional particles that are (approximately) degenerate with the decaying particle, and remain in chemical and thermal equilibrium with it until freezeout. One or more of these particles are DM candidates. \end{enumerate} Co-decaying DM will be a generic feature of large dark sectors in which the lightest state decays. To illustrate the idea, we will focus on the simplified case of two degenerate dark sector particles: $A$ will be the DM candidate, and $B$ will be the decaying state, with sizable annihilations $AA \to BB$. After the dark sector decouples from the SM bath, the $A$ and $B$ comoving entropy density is conserved, and their number density does not exponentially deplete when they become non-relativistic (in contrast to the Weakly Interacting Massive Particle (WIMP)). Instead, the exponential suppression is delayed until the $B$'s begin decaying: \begin{equation} n _A \sim n _B \propto e ^{ - \Gamma_B t } \simeq e ^{ - \frac{1}{2} \Gamma_B/H } \,, \end{equation} where $ n _{A,B} $ is the number density, $ \Gamma_B $ is the decay rate of the $ B $ particle, and $H$ is the Hubble parameter. The $ A $ population tracks the $ B $ population until the $ AA \rightarrow BB $ process cannot keep up with the expansion of the universe. At this point the $ A $ population freezes out and the $ B $'s continue to decay. The relic density of $A$ is then set by both the annihilation rate, $\langle \sigma v \rangle$, as well as the $B$ decay rate, $\Gamma_B$. A schematic illustration of the timeline for co-decaying DM is shown in Fig.~\ref{fig:scematic}. \begin{figure}[t!] \begin{center} \includegraphics[width=.4 \textwidth]{Timeline.pdf} \end{center} \caption{Co-decay dark matter timeline. At $ T _d $ the SM and dark sector decouple; at $ T _\Gamma $ the decay of $ B$'s begin to deplete the dark sector density; and at $T _f $ the $ AA \leftrightarrow BB $ process freezes out, resulting in a relic abundance for the $ A $ particles. } \label{fig:scematic} \end{figure} The delay in the starting point of exponential suppression from the temperature in which DM becomes non-relativistic to the temperature at which $B$-decay begins, causes freezeout to occur at later times than the WIMP. The DM relic density has less time to redshift to today, and therefore, must have a smaller density at freezeout. In order to match the observed DM relic abundance a larger annihilation cross-section is required. This leads to a boosted indirect detection signal relative to WIMP models. Previous work on multi-component dark sectors where interactions within the dark sector are necessary to get the correct dark matter relic abundance is extensive. Some examples including co-annihilating~\cite{Griest:1990kh,Baker:2015qna}, Secluded~\cite{Pospelov:2007mp}, SIMP~\cite{Hochberg:2014dra,Hochberg:2014kqa}, Cannibalizing~\cite{Pappadopulo:2016pkp,Kuflik:2015isi,Bernal:2015xba,Bernal:2015ova,Carlson:1992fn,Farina:2016llk} and Forbidden~\cite{Griest:1990kh,D'Agnolo:2015koa} DM. Additionally, models of particle decays affecting the relic abundance have been considered in~\cite{Feng:2003xh,Kaplinghat:2005sy,Farina:2015uea,Moroi:1999zb,Acharya:2009zt,Hall:2009bx,Berlin:2016vnh,Morrissey:2009ur,Cohen:2010kn,Bandyopadhyay:2011qm,Farina:2016llk}. The freezeout mechanism of co-decaying DM is unique, with differing phenomenology. Furthermore, we emphasize that while we are mainly interested in the implications on dark matter, the dynamics studied here have a broad impact and can take place for any thermal relic. In this Letter we study the co-decaying DM mechanism. We present an intuitive estimate of the relic density and check the results numerically using the Boltzmann equations. The constraints and signals of co-decaying DM are described, with a significant enhancement in the indirect detection signature. We conclude by presenting an explicit model realizing the phenomena.

1607.03110

1607

1607.07922_arXiv.txt

The most accepted interpretation of Brown Dwarfs (BDs) is that they are failed stars \citep{Cushing2014}, because, although it is assumed they formed like stars, their masses are too small to permit the fusion of hydrogen in their nucleus. This characteristic allows to separate BDs from main sequence stars based on their masses: because a star must reach a critical mass to be able to burn its hydrogen, which varies from $0.07\ M_\odot$ for solar metallicity to $0.09\ M_\odot$ for lower metallicities \citep{Burrows2001}, any star with a mass $< 70\ M_J$ (where $M_J$ is the mass of Jupiter) is a BD \citep{Bate2006}. However, determining a lower mass limit for a BD is more difficult. In practice, the consensus to adopt the critical mass for the fusion of deuterium, which is around $13\ M_J$ \citep{Bate2006}, is arbitrary, because theoretically the lowest mass a BD could have may be just a few $M_J$ \citep{Larson1969,Rees1976,Silk1977a,Silk1977,Boss1988}. Interestingly, this mass is also typical of massive exoplanets, and, since there is no obvious upper-mass limit for an exoplanet, hence, persists the problem of distinguishing between to two objects. In this poster, using a large sample of ``well-studied'' exoplanets, and comparing with a large sample of ``confirmed'' BDs available in the literature, we probe a mass range common to both classes of objects, looking for evidence of a difference between their respective physical structures, as reflected by their mass-radius relations (hereafter MRRs). Our study concentrates on two questions: 1) At what mass boundary should we expect to see a variation in the MRR that would be consistent with a difference of structure between exoplanets and BDs? 2) Is there a special intermediate mass range where these two classes of objects are likely to overlap in mass? In particular, we propose a lower-mass limit for BDs based on the Self-Gravitating (SG) limit, which marks the moment the self-gravity of matter begins to affect significantly the structure of a body \citep{Padmanabhan1993}. In addition to the MRR, the distance of a planet from its host star could also reveal something about its formation process \citep{Lissauer1993}. For the exoplanets, this last parameter is fundamental to identify Hot Jupiters \citep{Johnson2009}, while for the BDs, this parameter can be used to test the ``BD's desert'' hypothesis, which according to some authors \citep[e.g.,][]{Grether2006} might be related to different formation processes for exoplanets and BDs.

\begin{itemize} \item We conclude that the unsignificant change of radius of exoplanets above the SG limit is the characteristic signature of objects formed by LMH. \item As for the nature of these objects we propose that they could be either giant gas planets with a dominant layer of LMH, more massive than what is assumed to exist in Jupiter and Saturn, or genuine very low-mass BDs. \end{itemize}

1607.07922

1607

1607.02393_arXiv.txt

MATISSE represents a great opportunity to image the environment around massive and evolved stars. This will allow one to put constraints on the circumstellar structure, on the mass ejection of dust and its reorganization, and on the dust-nature and formation processes. MATISSE measurements will often be pivotal for the understanding of large multiwavelength datasets on the same targets collected through many high-angular resolution facilities at ESO like sub-millimeter interferometry (ALMA), near-infrared adaptive optics (NACO, SPHERE), interferometry (PIONIER, GRAVITY), spectroscopy (CRIRES), and mid-infrared imaging (VISIR). Among main sequence and evolved stars, several cases of interest have been identified that we describe in this paper.

First of all, we need to emphasize that MATISSE will mainly focus on dusty stars, as the mid-infrared is the perfect match to collect data on the dust quantity and composition around a given object. To illustrate that, we selected a few topics of interest that will be developed in this paper, and counted the number of stars that will be observable with VLTI/MATISSE. We included a quantity of non-dusty targets (regular WR stars and Be stars) in the sample to compare the performances of MATISSE with its prime targets. To do so, one need to take into account the sensitivity of the instrument itself in the L- and N-bands\cite{Matter2016}, of course, but also the VLTI infrastructure subsystems sensitivity, which can be sometimes the limiting factor, especially for red targets. We need therefore to take into account: the telescope guiding limit in the V-band (V=13.5 and V=17 for ATs STRAP tip-tilt and UTs MACAO AO, respectively), and the fringe tracker limit in the K-band for GRA4MAT and other devices like IRIS (K=7.5 for ATs and K=10 for UTs). We considered star lists on Wolf-Rayet stars (WR), Red supergiant stars (RSG), Asymptotic Giant Branch stars (AGB), Be stars and B[e] supergiant stars visible from Paranal ; we extracted them from the following catalogs: \begin{itemize} \item Rosslowe \& Crowther (2014) for WR stars\cite{2015MNRAS.447.2322R}, \item Hoffleit (1991) for RSG stars\cite{1991bsc..book.....H}, \item we scanned through ADS publications on the topic for AGB stars, \item Frémat et al. (2005)\cite{2005A&A...440..305F} and Yudin (2001)\cite{2001A&A...368..912Y} for Be stars, \item For B[e] supergiant stars we used our own catalog of stars. \end{itemize} When available, we retrieved the SED of each object (V, K, L and N magnitudes) and ISO/IRAS spectra (if it exist). Magnitudes of these targets were compared with the theoretical limits of MATISSE, considering spectrally-usable data, i.e. data with a SNR$\geq3$ per spectral channel during the expected optimal DIT and exposure time combination, with and without fringe tracker\cite{Matter2016}. The result of this study is shown in Fig.~\ref{fig:targets}. One can see that, out of our selected catalogs, roughly one third of the considered AGB targets, two third of the supergiant B[e] stars, and half of the considered red supergiant stars can be observed with MATISSE in its low spectral resolution mode (R$\approx$35). We are likely limited by our sample size in these cases. On the other hand, very red targets like dusty Wolf-Rayet stars, or very blue targets like naked Wolf-Rayet stars or Be stars, will be more difficult to observe with MATISSE due to either the sensitivity limits of VLTI in the V band (very red targets), or to the MATISSE sensitivity limit (very blue targets). In both cases, the adjunction of a K-band fringe tracker to MATISSE, like the foreseen GRA4MAT project of ESO, will improve the situation by a large factor (especially when using the ATs). That matter of fact is even more striking when using the high-spectral resolution, where we will be able to resolve the kinematic motion of gaz in the circumstellar envelopes\cite{}. From a handful of targets that will be observable with MATISSE alone, the adjunction of an external fringe tracker will allow the instrument to go at full throttle for the delivery of scientific results. \begin{figure}[htbp] \begin{center} \begin{tabular}{ccc} \includegraphics[width=.47\textwidth]{Histo_LR.png}& \includegraphics[width=.47\textwidth]{Histo_HR.png} \end{tabular} \end{center} \caption[Number of targets observable] { \label{fig:targets} {\bf Left:} Number of targets observable with MATISSE in low spectral resolution (R$\approx$35) for different types of stars, both in L band (\texttt{L\_noft}, dark blue bars) and N band (\texttt{N\_noft}, light blue bars). The transparent bars show the improvement of observability when using an external fringe tracker (\texttt{L\_ft} and \texttt{N\_ft}), like e.g. the foreseen project GRA4MAT. {\bf Right:} The same plot for the high resolution of MATISSE (R$\approx$4500 around the hydrogen Brackett $\alpha$ line).} \end{figure}

1607.02393

1607

1607.06265_arXiv.txt

{The CO$^+$ reactive ion is thought to be a tracer of the boundary between a \hii\ region and the hot molecular gas. In this study, we present the spatial distribution of the CO$^+$ rotational emission toward the \MonR\ star-forming region. The CO$^+$ emission presents a clumpy ring-like morphology, arising from a narrow dense layer around the \hii\ region. We compare the CO$^+$ distribution with other species present in photon-dominated regions (PDR), such as [\ciii] 158~$\mu$m, H$_2$ S(3) rotational line at 9.3~$\mu$m, polycyclic aromatic hydrocarbons (PAHs) and HCO$^+$. We find that the CO$^+$ emission is spatially coincident with the PAHs and [\ciii] emission. This confirms that the CO$^+$ emission arises from a narrow dense layer of the \hi/H$_2$ interface. We have determined the CO$^+$ fractional abundance, relative to \cii\, toward three positions. The abundances range from 0.1 to 1.9 $\times 10^{-10}$ and are in good agreement with previous chemical model, which predicts that the production of CO$^+$ in PDRs only occurs in dense regions with high UV fields. The CO$^+$ linewidth is larger than those found in molecular gas tracers, and their central velocity are blue-shifted with respect to the molecular gas velocity. We interpret this as a hint that the CO$^+$ is probing photo-evaporating clump surfaces.}

} Reactive ions are destroyed in almost every collision with H and H$_2$ and recombine rapidly with e$^-$. These compounds present enhanced abundances toward the hot layers of the photon-dominated regions (PDRs), where the far ultraviolet (FUV) field is only partially attenuated and maintains high abundances of the parent species \cii\ and S$^+$ \citep{sternberg-Dalgarno1995}. In particular, \citet{sternberg-Dalgarno1995} predict a high CO$^+$ abundance at the \hi/H$_2$ interface ($A_\mathrm{V}\approx1$~mag) of dense PDRs, where it is mainly produced by the \cii\ + OH $\rightarrow$ CO$^+$ + H reaction. So far, the CO$^+$ ion has been detected in several PDRs; such as the M17SW, Orion Bar, NGC7027, NGC7023 (\citealt{latter1993}; \citealt{stoerzer1995}; \citealt{fuente1997}; \citealt{fuente2003}), G29.96, MonR2 \citep{rizzo2003} and S140 \citep{savage2004}. However, all the detections have been obtained after long integrations toward a single position and they lack the information on the spatial distribution. The \MonR\ star-forming region, located at 830~pc (\citealt{herbst1976}), contains an ultracompact (UC)~\hii\ region surrounded by a series of PDRs with different physical conditions \citep{pilleri2013, trevino2014}. The main PDR, corresponding to IRS~1 (hereafter IF), is irradiated by a high UV field of $G_0>10^{5}$ (in units of the Habing flux; \citealt{habing1968}), and presents high densities ($>10^{5}$~cm$^{-3}$) and kinetic temperatures ($T_{\mathrm{k}}\approx 600$~K; \citealt{berne2009}). A second PDR, associated with the molecular peak MP2, is detected $40\arcsec$ north from IF, and shows chemical properties similar to those found in low- to mid-UV irradiated PDRs (\citealt{ginard2012}). Due to its proximity and physical conditions, \MonR\ turns to be an excellent candidate to study the \hi/H$_2$ interface. CO$^+$ is thought to be a good PDR tracer, and its distribution is potentially an excellent diagnostic tool to learn about the physical structure of these regions. In this paper, we present a study of the CO$^+$ ($J=2$--1) transition line toward \MonR\ and compare its spatial distribution with \emph{Spitzer} data reported by \citet{berne2009}, \emph{Herschel} data from \citet{pilleri2012} and \citet{ossenkopf2013} and the HCO$^+$ and H$^{13}$CO$^+$ molecules from \citet{trevino2014}. \begin{figure*}[t] \centering \includegraphics[width=4.4cm]{fig_paper_overview_1.pdf} \hspace{-0.3cm}\includegraphics[width=4.4cm]{fig_paper_overview_2.pdf} \hspace{-0.3cm}\includegraphics[width=4.4cm]{fig_paper_overview_3.pdf} \hspace{-0.3cm}\includegraphics[width=4.4cm]{fig_paper_overview_4.pdf} \caption{Color image shows the integrated emission (in K~\kms) of the CO$^+$ line, with the original (11\arcsec, \emph{Panel A}) and smoothed (16\arcsec, \emph{Panel B} to \emph{D}) angular resolution. In \emph{Panels A} to \emph{D}, the gray contour levels range from 40\% to 100\%, in steps of 10\% of the intensity peak, where the lower contour level corresponds to a S/N$=5\sigma$. The yellow contour indicates the 3$\sigma$ emission. The blue squares indicate the IF and MP2 positions, where IF corresponds to $\alpha$(J2000)$=06\mathrm{h}07\mathrm{m}46.2\mathrm{s}$, $\delta$(J2000)$=-06^{\circ}23'08.3''$. \emph{Panel~A} shows the H$^{13}$CO$^+$\ (3--2) emission (black contours) tracing the molecular gas \citep{trevino2014}. \emph{Panel~B} shows the [Ne{\small II}] emission (red contours) tracing the \hii\ region and the emission of the H$_2$ S(3) rotational line at 9.7~$\mu$m (black contours). Black contours in \emph{Panel~C} show the PAHs (11.3~$\mu$m) emission. \emph{Panel~D} shows the [\ciii] emission at 158~$\mu m$ (black contours, \citealt{pilleri2014}). The \emph{Spitzer} data are explained in \citet{berne2009}.} \label{figureA1} \end{figure*}

} We present a CO$^+$ map toward \MonR\ star-forming region. This is the first map ever reported of this reactive ion. The spatial distribution of CO$^+$ consists of a ring-like structure (similar to PAHs), tracing the layer between the \hii\ region and the molecular gas. The maps reveal a clumpy structure in the hot layer of the mainly atomic gas. Previous works (\citealt{young2000}; \citealt{goicoechea2016}) suggest that there exist fragmentation in the photodissociation front. This is, there are not uniform layers between the \hii\ region and the molecular cloud, but they present clumps that allow the radiation to penetrate deeper into the cloud. In this scenario where the PDR is conformed by a series of clumps, the emission of PDR tracers would be related to the external layers of dense clumps being photo-evaporated by the UV radiation. Despite the moderate angular resolution of our observations, we find hints that favor this scenario: (a) the spatial distribution of the CO$^+$ as observed in the higher-angular resolution (11$''$) map (\emph{Panel~A} of Fig.~\ref{figureA1}) suggests that the CO$^+$ emission is coming from the illuminated surface of the H$^{13}$CO$^{+}$ clumps, and b) the excess of blue-shifted emission seen for the PDR tracers in comparison with the molecular tracers. Considering that the molecular gas is located behind the UC~\hii\ region \citep{pilleri2014} and the chemical segregation, the difference in velocity between tracers can be explained if the PDR is formed by dense condensations that are being photo-evaporated. In this case, the photo-evaporated gas (PDR tracers) would be ejected toward us, and therefore, blue-shifted with respect to the molecular gas. Future higher angular resolution observations will help to confirm or discard this scenario. Finally, we have determined $X$[CO$^+$] in three positions. Toward IF we derive an abundance of a few $10^{-11}$, in agreement with chemical model predictions \citep{sternberg-Dalgarno1995} for $n_{\mathrm{H}}\sim 10^6$~cm$^{-3}$ and G$_{0}\sim5\times 10^5$. Toward MP2 we do not detect CO$^+$ emission with an upper limit to the CO$^+$ abundance of $<4\times 10^{-11}$. Abundances of $10^{-11}$--$10^{-10}$ had been previously observed in PDRs with $G_0>10^3$ Habing field (M17SW: \citealt{latter1993}, \citealt{stoerzer1995}; Orion Bar: \citealt{fuente1997}; NGC\,7023: \citealt{fuente2003}; G29.96$-$0.02: \citealt{rizzo2003}). The non-detection of CO$^+$ in MP2, together with the abundances found in the other PDRs, suggest that the production of CO$^+$ only occurs in dense regions with high radiation fields. High UV fields ($G_0>10^3$) and $n_{\mathrm{H}}$ ($>2\times 10^{4}$~cm$^{-3}$) are required to achieve gas temperatures $\geq300$~K that are necessary to produce high abundances of OH in the external layer of the PDR ($A_\mathrm{{V}}\sim1$~mag; \citealt{stauber-Bruderer2009}). This is also consistent with the non-detection of CO$^+$ in the Horsehead PDR \citep{goicoechea2009}, where the UV field is $G_0\sim100$ and presents chemical properties similar to MP2 \citep{ginard2012}. A counter-example that challenge this interpretation could be the detection of CO$^+$ toward S140 where the incident UV field is estimated to be $G_0\sim100$--300 \citep{savage2004}. However, the number of CO$^+$ detections is scarce and the statistics is not enough to draw firm conclusions between the relation of the CO$^+$ and the physical properties ($n_{\mathrm{H}}$ and $G_0$) of PDRs. A larger sample of objects need to be studied, including maps to characterize and understand the spatial distribution of the CO$^+$ in different environments.

1607.06265

1607

1607.04056_arXiv.txt

We present the results of 325 MHz GMRT observations of a super-cluster field, known to contain five Abell clusters at redshift $z \sim 0.2$. We achieve a nominal sensitivity of $34\,\umu$Jy beam$^{-1}$ toward the phase centre. We compile a catalogue of 3257 sources with flux densities in the range $183\,\umu\rm{Jy}\,-\,1.5\,\rm{Jy}$ within the entire $\sim 6.5$ square degree field of view. Subsequently, we use available survey data at other frequencies to derive the spectral index distribution for a sub-sample of these sources, recovering two distinct populations -- a dominant population which exhibit spectral index trends typical of steep-spectrum synchrotron emission, and a smaller population of sources with typically flat or rising spectra. We identify a number of sources with ultra-steep spectra or rising spectra for further analysis, finding two candidate high-redshift radio galaxies and three gigahertz-peaked-spectrum radio sources. Finally, we derive the Euclidean-normalised differential source counts using the catalogue compiled in this work, for sources with flux densities in excess of $223 \, \umu$Jy. Our differential source counts are consistent with both previous observations at this frequency and models of the low-frequency source population. These represent the deepest source counts yet derived at 325 MHz. Our source counts exhibit the well-known flattening at mJy flux densities, consistent with an emerging population of star-forming galaxies; we also find marginal evidence of a downturn at flux densities below $308 \, \umu$Jy, a feature so far only seen at 1.4 GHz.

Deep radio surveys of the extragalactic source population are powerful tools with which to probe a wide range of source populations across a variety of environments and redshifts. In previous decades, optical surveys have been preferred for the study of the formation, interactions and evolution of galaxies. However, radio emission is important for galaxy population studies, as the synchrotron emission is a clear indicator of magnetic fields from star formation and active galactic nuclei (AGN). Additionally, the radio emission is essentially unaffected by dust obscuration and as such provides a powerful tracer of the evolution of star-forming galaxies and AGN with redshift. All-sky surveys -- for example the Very Large Array (VLA) Faint Images of the Radio Sky at Twenty-Centimetres (FIRST; \citealt{1994ASPC...61..165B}) survey and the NRAO VLA Sky Survey (NVSS; \citealt{1998AJ....115.1693C}) at 1.4 GHz, and the 325 MHz Westerbork Northern Sky Survey (WENSS; \citealt{1997A&AS..124..259R}) -- have been effective in identifying large numbers of bright radio sources and have led to studies of the populations they represent. At higher frequencies ($\nu \gtrsim 1.4$ GHz) and/or higher flux densities ($S \gtrsim 1-10$ mJy at 1.4 GHz) the dominant population of sources are radio-loud AGN (for example \citealt{1984ApJ...287..461C}, \citealt{1994ASPC...61..165B}, \citealt{1998AJ....115.1693C}, \citealt{1999MNRAS.305..297G}, \citealt{2005ApJ...624..135A}, \citealt{2008ApJ...681.1129B}). Moving to fainter flux densities, the contribution from star-forming galaxies (SFG) and radio-quiet (RQ) AGN become increasingly important, and these sources are believed to dominate at $S \lesssim 0.1$ mJy (for example \citealt{1999ApJ...526L..73R}, \citealt{2005MNRAS.358.1159M}, \citealt{2007ASPC..380..205P}, \citealt{2008ApJS..179...95M}, \citealt{2009MNRAS.397..281I}, \citealt{2009ApJ...694..235P}, \citealt{2013MNRAS.436.3759B}, \citealt{2015MNRAS.452.1263P}). The physics of low-luminosity SFG is still poorly understood; all-sky surveys are too shallow to recover sufficient numbers of these faint sources to infer much detail. Increasingly deep surveys such as the VLA-COSMOS survey (for example \citealt{2004AJ....128.1974S}) as well as smaller, deep fields (e.g. \citealt{2004A&A...424..371M}, \citealt{2013ApJS..205...13M}) have recovered sources down to $\sim \umu$Jy flux densities at 1.4 GHz. Radio emission from SFG is comprised of two components: synchrotron emission dominates at low frequencies, whereas thermal bremsstrahlung (free-free emission) from the ionized interstellar medium dominates at higher frequencies (for example \citealt{1992ARA&A..30..575C}, \citealt{2002A&A...392..377B}, \citealt{2008A&A...477...95C}). Whilst synchrotron emission presents itself with typical spectral index\footnote{Adopting the convention $S \propto \nu^{\alpha}$.} of $\alpha \simeq -0.8$, free-free emission has a flatter spectrum (typically $\alpha \simeq -0.1$). In addition to these features at high frequency, SFG spectra exhibit a number of other features below 1 GHz, with bends and inversions detected in some spectra \citep{2010MNRAS.405..887C}. Surveys at low frequencies (for example \citealt{1991PhDT.......241W}, \citealt{2008MNRAS.383...75G}, \citealt{2008MNRAS.387.1037G}, \citealt{2009MNRAS.395..269S}, \citealt{2009AJ....137.4846O}, \citealt{2010iska.meetE..51S}, \citealt{2013MNRAS.435..650M}, \citealt{2014MNRAS.443.2590S}) open a new window for study, offering a number of advantages over their higher-frequency counterparts. Low-frequency observations are powerful at detecting (ultra-)steep spectrum sources, which are often galaxies at high redshift (\citealt{2008A&ARv..15...67M} and references therein). Low-frequency observations also enable detailed studies of the radio synchrotron spectral index, which allows for more precise characterisation of the source properties. In this work, we present the results of a deep 325 MHz continuum survey of a super-cluster field, performed using the Giant Metrewave Radio Telescope (GMRT) and carried out as part of the Super-Cluster Assisted Shear Survey (Super-CLASS). The remainder of this paper is divided as follows: we firstly introduce the Super-CLASS project in \S\ref{sec:proj}; subsequently we detail the observations and data reduction methodology in \S\ref{sec:obs}. We present our results in \S\ref{sec:RES}, including a sample from our source catalogue; we verify the catalogue and analyse the statistical properties in \S\ref{sec:AN}. In \S\ref{sec:DISC} we derive the spectral index distribution and identify sources with steep and/or inverted spectra for further study, which may constitute ultra-steep spectrum radio sources or gigahertz-peaked spectrum radio sources. We derive the source counts distribution from our catalogue, as well as evaluate the various bias and incompleteness corrections that must be applied, in \S\ref{sec:src}. Finally, we draw our conclusions in \S\ref{sec:CONC}. All errors are quoted to $1\sigma$. We adopt the spectral index convention that $S \propto \nu^{\alpha}$ where the radio spectral index $\alpha < 0$. We assume a concordance cosmology of H$_0 = 73 \rm{km} \rm{s}^{-1}$ Mpc$^{-1}$, $\Omega_{\rm{m}} = 0.27$, $\Omega_{\rm{\Lambda}} = 0.73$. At a redshift of $z = 0.2$, representative of the constituent clusters of the Super-CLASS super-cluster, an angular size of 1 arcsecond corresponds to a physical size of 3.2 kpc. \subsection{The Super-Cluster Assisted Shear Survey}\label{sec:proj} Weak lensing in the radio regime is emerging as promising cosmological probe, as many of the issues which strongly affect optical lensing studies (such as atmospheric effects and anisotropic PSFs) are negated by shifting to radio wavelengths. While measurements of cosmic shear on scales of $1-4$ degrees have been made at radio wavelengths (for example \citealt{2004ApJ...617..794C}, and see also \citealt{2010MNRAS.401.2572P}) previous shear surveys have been severely limited by resolution, field of view, and low source counts. A large, high-resolution catalogue is required to disentangle the shear signal (a factor $\sim0.01$) from intrinsic source ellipticity (typically $\sim0.3$). Recent work by \cite{2016MNRAS.456.3100D} has demonstrated that biases in shear measurements can be mitigated by cross-correlating both radio and optical survey data Of the current generation of instruments, perhaps the best suited to studies of cosmic shear is the expanded Multi-Element Remote-Linked Interferometer Network (e-MERLIN). A number of legacy surveys are currently underway with e-MERLIN, including the Super-CLASS project\footnote{For more details, see \url{http://www.e-merlin.ac.uk/legacy/projects/superclass.html}}. Super-CLASS is a wide-area, deep e-MERLIN legacy survey at L-band (reference frequency 1.4 GHz) targeting a region of sky known to contain five moderate-redshift $(z\sim0.2)$ Abell clusters -- A968, A981, A998, A1005 and A1006 \citep{1989ApJS...70....1A}. Some observational properties of these clusters are listed in Table \ref{tab:clusters}. Hereafter this region is referred to as the Super-CLASS field. The principal goal of the project is to detect the effects of cosmic shear in a super-cluster environment, where the increased level of structure should allow for a statistically significant detection of shear over a wide range of scales. However, a number of ancillary science goals exist, such as investigation of polarization properties of AGN and SFG, studies of cosmic magnetism in cluster- and super-cluster environments, classifying the galaxy population in the super-cluster, and detailed studies of the radio source population at $\umu$Jy flux densities. A wide range of other instruments are involved, including the Karl G. Jansky Very Large Array (JVLA) and LOw-Frequency Array (LOFAR; \citealt{2013A&A...556A...2V}) in the radio band, and a number of optical and mm-/sub-mm wavelength telescopes. The weak lensing component of the survey is similar in manner (although with lower source densities and on a smaller field) to those that may ultimately be conducted by the Square Kilometre Array (SKA; \url{http://skatelescope.org}). For more details on weak lensing with the SKA see for example \citet{2015aska.confE..23B}. See \citet{harrison2016arxiv} for cosmology forecasts from weak lensing with the SKA; for simulated catalogues see \citet{bonaldi2016arxiv}. \begin{table} \begin{center} \caption{Properties of galaxy clusters constituting the Super-CLASS super-cluster.} \label{tab:clusters} \begin{threeparttable} \begin{tabular}{cccccc} \hline & & & \\ Name & RA & Dec & $z$ & $L_x$ (0.1-2.4 keV) \\ & (J2000) & (J2000) & & $[\times10^{44}$ erg s$^{-1}]$\\ \hline Abell 968 & 10$^{\rm{h}}$21$^{\rm{m}}$09.5$^{\rm{s}}$ & +68$\degree$15$^{\prime}$53$^{\prime\prime}$ & 0.195 & 0.401 \\ Abell 981 & 10$^{\rm{h}}$24$^{\rm{m}}$24.8$^{\rm{s}}$ & +68$\degree$06$^{\prime}$47$^{\prime\prime}$ & 0.202 & 1.670 \\ Abell 998 & 10$^{\rm{h}}$26$^{\rm{m}}$17.0$^{\rm{s}}$ & +67$\degree$57$^{\prime}$44$^{\prime\prime}$ & 0.203 & 0.411 \\ Abell 1005 & 10$^{\rm{h}}$27$^{\rm{m}}$29.1$^{\rm{s}}$ & +68$\degree$13$^{\prime}$42$^{\prime\prime}$ & 0.200 & 0.268 \\ Abell 1006 & 10$^{\rm{h}}$27$^{\rm{m}}$37.2$^{\rm{s}}$ & +67$\degree$02$^{\prime}$41$^{\prime\prime}$ & 0.204 & 1.320 \\ \hline \end{tabular} References: \begin{tablenotes} \item{Redshift, $z$: \citet{1990ApJ...365...66H}} \item{X-ray luminosity, $L_x$: BAX database; \citet{2004A&A...424.1097S}} \end{tablenotes} \end{threeparttable} \end{center} \end{table} \begin{table} \begin{center} \caption{Summary of GMRT pointings on the Super-CLASS field. For each pointing, we list the integration time $(\tau_{\rm{int}})$ as well as the percentage of data that was used (i.e. unflagged) in the final image, as well as the predicted thermal noise.} \label{tab:pointing_summary} \scalebox{0.85}{ \begin{tabular}{ccccccc} \hline & & & & \multicolumn{2}{c}{$\sigma_{\rm{rms}} [\umu$Jy beam$^{-1}]$ } \\ RA & Dec & $\tau_{\rm{int}}$ & Unflagged & Thermal & Measured \\ (J2000) & (J2000) & $[$s$]$ & $[$\%$]$ & & \\ \hline 10$^{\rm{h}}$21$^{\rm{m}}$56.13$^{\rm{s}}$ & 68$\degree$09$^{\prime}$44.8$^{\prime\prime}$ & $14.8\times10^3$ & 66.0 & 33.9 & 46.4 \\ 10$^{\rm{h}}$24$^{\rm{m}}$17.12$^{\rm{s}}$ & 67$\degree$35$^{\prime}$26.7$^{\prime\prime}$ & $15.3\times10^3$ & 62.9 & 34.2 & 42.9 \\ 10$^{\rm{h}}$26$^{\rm{m}}$31.43$^{\rm{s}}$ & 67$\degree$01$^{\prime}$01.4$^{\prime\prime}$ & $15.0\times10^3$ & 63.8 & 34.3 & 42.8 \\ 10$^{\rm{h}}$28$^{\rm{m}}$50.28$^{\rm{s}}$ & 67$\degree$35$^{\prime}$25.0$^{\prime\prime}$ & $14.5\times10^3$ & 64.7 & 34.6 & 44.0 \\ 10$^{\rm{h}}$26$^{\rm{m}}$32.15$^{\rm{s}}$ & 68$\degree$09$^{\prime}$57.5$^{\prime\prime}$ & $14.5\times10^3$ & 68.0 & 33.8 & 41.5 \\ 10$^{\rm{h}}$24$^{\rm{m}}$15.01$^{\rm{s}}$ & 68$\degree$44$^{\prime}$23.0$^{\prime\prime}$ & $14.6\times10^3$ & 66.8 & 34.0 & 44.8 \\ \hline \end{tabular} } \end{center} \end{table} \begin{figure} \centering \includegraphics[width=0.4\textwidth]{Figures/superclass_uv_double.pdf} \label{fig:uvcovg_full} \caption{Typical \emph{uv}-coverage of a single pointing on the Super-CLASS field. \emph{Top panel:} Plot showing coverage of the full \emph{uv-}plane. \emph{Bottom panel:} Zoom on the central portion of the \emph{uv-}plane, illustrating the coverage on short baselines.} \label{fig:uvcovg} \end{figure}

\label{sec:CONC} We report deep 325 MHz GMRT observations of the Super-CLASS field, a region of sky known to contain 5 Abell clusters. We achieve a nominal sensitivity of $34\,\umu$Jy beam$^{-1}$ toward the centre of the field, the deepest study conducted at this frequency to-date. From our mosaicked image, which covers approximately 6.5 square degrees, we recover a catalogue of 3257 sources down to a flux density limit of $183 \, \umu$Jy. We have compared with available catalogues from other surveys, and identified 335 sources in common with the NVSS, after accounting for sources which become resolved in the higher resolution GMRT data. The spectral index distribution indicates two populations within the data: a more numerous population of sources centred around a spectral index of $\alpha = -0.81$, which dominate at higher flux densities, and a less numerous, fainter population of sources with flat- and rising spectra, which appear to dominate the fainter flux densities. However, this spectral index sample is currently severely limited by the sensitivity of the high-frequency reference. Adopting the definition of an ultra-steep spectrum object as a radio source with $\alpha < -1.3$ (as is typically the case in the literature) we find a total of four ultra-steep spectrum radio sources which have counterparts in the NVSS. Of these, we identify two AGN-type sources whose steep spectra may be the result of our superior sensitivity to faint diffuse emission, and two candidate HzRGs. Additionally, we have identified three GPS sources, and a further candidate GPS source which is not identified elsewhere in the literature aside from the NVSS. One of these GPS sources is the blazar candidate CGRaBS J1015+6728; we recover an integrated flux density of $3.09\pm0.31$ mJy, to our knowledge the first flux density measurement for this object below 1 GHz. We model the spectra of these GPS sources using a broken power-law; our fits suggest break frequencies in the $2-3$\,GHz range. Finally, we derive the Euclidean-normalised differential source counts for sources with flux densities in excess of $223\,\umu$Jy. We have performed a rigorous mathematical treatment of the various biases in the differential source counts, including Eddington bias and resolution bias. The differential source counts appear to be consistent with predictions from numerical models which account for steep-spectrum sources (AGN, FSRQ and BL Lac objects) and star-forming galaxies (spiral galaxies and starburst galaxies). We have also provided a new empirical model for describing the Euclidean-normalised differential source counts that is valid between $242 \, \umu$Jy and 0.74 Jy. The source counts distribution derived in this work exhibits marginal evidence of a secondary downturn at flux densities below $308\,\umu$Jy. This corresponds well with the flux density regime where this feature has been detected at 1.4 GHz. To our knowledge, this is the first detection of this feature in the differential source counts at 325 MHz. \subsection{Acknowledgements} We thank our anonymous referee for their comments, which have helped improve the quality of the scientific output of this paper. We also thank the operators and engineers of the GMRT who made these observations possible. The GMRT is operated by the National Centre for Radio Astrophysics (NCRA) of the Tata Institute of Fundamental Research (TIFR), India. CJR wishes to thank H.~Intema for many helpful conversations during the data reduction process. This work has made use of the NASA/IPAC Extragalactic Database (NED), operated by JPL under contract with NASA, as well as NASA's Astrophysics Data System (ADS) and the Cosmological Calculator developed by \citet{2006PASP..118.1711W}. This work has also used data from the Owens Valley Radio Observatory (OVRO) blazar monitoring database, available at \url{http://www.astro.caltech.edu/ovroblazars/}. CJR gratefully acknowledges funding support from the United Kingdom Science \& Technology Facilities Council (STFC). AMS gratefully acknowledges support from the European Research Council under grant ERC-2012-StG-307215 LODESTONE. IH, MLB and CD are grateful to the European Research Council for support through EC FP7 grant number 280127. MLB also thanks the STFC for the award of Advanced and Halliday fellowships (grant number ST/I005129/1).

1607.04056

1607

1607.03440_arXiv.txt

\noindent We study a potential genetic relationship of comets C/1846 O1~and~C/1973\,D1,~whose~\mbox{apparent}~\mbox{orbital} similarity~was~tested~by~Kres\'ak~(1982)~only~\mbox{statistically},\,\mbox{using}~the~\mbox{Southworth-Hawkins}\,(1963)\,\mbox{criterion} $D$. Our orbit determination for C/1846\,O1 shows its period was\,$\sim$500\,yr,\,$\sim$30 times shorter than that of C/1973~D1. Formerly unrecognized, this incongruity makes the objects'\,common origin~less~likely. Long-term orbit integration suggests that, if related, the two comets would~have~to~have~sepa\-rated far away from the Sun (probably $\sim$700~AU) 21~millennia ago and, unlike C/1973~D1, C/1846~O1 would~have~to~have been subjected to a complex orbital evolution. Given the chance of encountering Jupiter to $\sim$0.6~AU some 400~days after perihelion, C/1846~O1 and C/1973~D1~may~have~been~perturbed, during their return in the 15th millennium BCE, into {\vspace{-0.01cm}}orbits that were, respectively, smaller and larger than was the parent's, with a net difference of more than 0.002~(AU)$^{-1}$ in $1/a$. Whereas C/1973~D1 was on the way to its 1973 perihelion, C/1846~O1 should have been subjected to recurring encounters with Jupiter, during which the orbital period continued to shorten by integral multiples of the Jovian orbital period, a process called {\it high-order orbital-cascade resonance\/}. While~the~inte\-grated perturbation effect of C/1846~O1 by Jupiter does not appear to reduce the comet's orbital period~to below $\sim$1200~yr by the mid-19th~century, we find that orbital-cascade resonance offers an attractive mechanism for rapid inward drifting of aphelion especially among dynamically new comets.{\vspace{0.05cm}}

}} We recently expounded the relationship between fragmentation and the orbital properties for two well-known groups of genetically related long-period comets:\ one was the pair of C/1988 F1 (Levy) and C/1988 J1 (Shoemaker-Holt) and the other was the trio of C/1988 A1 (Liller), C/1996 Q1 (Tabur), and C/2015 F3 (SWAN) (\mbox{Sekanina} \& Kracht 2016, referred to hereafter as Paper~1). Prior to the discovery of the two groups, a genetic relationship as the provenance for close orbital similarity among comets was a subject of controversy, especially in the 1970s and early 1980s. In a contribution to the debate, Kres\'ak (1982) corroborated{\vspace{-0.075cm}} and extended Whipple's (1977) criticism of \"{O}pik's (1971) conclusion on the omnipresence of groups of related comets. Kres\'ak argued that there was no compelling evidence for the existence of a single pair or group of long-period comets that derive from a common parent, other than the Kreutz system of sungrazers.{\vspace{0.05cm}} The approach employed in the pre-1988 debate was always statistical in nature. In particular, Kres\'ak (1982) used a $D$-criterion, introduced by Southworth \& Hawkins (1963) in their investigation of meteor streams. As a function of differences in the five orbital elements --- the argument of perihelion $\omega$, the longitude of the ascending node $\Omega$, the inclination $i$, the perihelion distance $q$, and the eccentricity $e$ --- the $D$-criterion allows one to express the degree of similarity between two orbits in one-dimensional phase space. Objects in orbits of the same spatial orientation that are identical in size and shape have \mbox{$D = 0$}. The $D$-values of the genetically related 1988 pair and the 1988--2015 trio, referred to above, are listed in Table 1. They never exceed $\sim$0.008, just as the $D$-values for the Kreutz system's most tightly associated members (such as C/1882 R1 and C/1965 S1). The prime subject of Kres\'ak's study was a distribution of $D$-values among 546 comets [a majority extracted from Marsden (1979) and several added] that arrived at perihelion before the end of 1980 and whose orbital periods exceeded 200 years; the Kreutz sungrazers were excluded. After finding 38 pairs (and several chains) of comets with \mbox{$D < 0.3$} and comparing them with three independent distributions of 546 randomly generated orbits (accounting in part for observational selection effects), Kres\'ak concluded that the set of long-period comets exhibited no signs of nonrandom distribution. In particular, he judged orbital similarity of the comet pair with the least $D$-value of 0.084 --- C/1846~O1 (de Vico-Hind) and C/1973~D1 (Kohoutek) --- not to be statistically significant on the grounds that comparison with the least $D$-values in the random samples, 0.101--0.120, suggested, on the average, a $\sim$20\% expectation that this was a random pair as well.{\hspace{0.21cm}} Unfortunately, Kres\'ak's failure to remove from his statistics the grossly inferior orbits of comets observed in early times marred his main results, and he addressed a much more meaningful subset of 438 long-period comets from the period of 1800--1980 rather inadequately. Our closer examination of this subset shows that the minimum $D$-value in the three random samples then moves up to a range of 0.113--0.134 and there is only an 11\% expectation for C/1846~O1 and C/1973~D1 being a random pair. Moreover, if this pair is removed from the set, the next pair's $D$-value of 0.129 is consistent with an average random sample with an expectation of 60\%. \begin{table}[t] % \vspace{-3.9cm} \hspace{4.23cm} \centerline{ \scalebox{1}{ \includegraphics{t1_deVico.ps}}} % \vspace{-22.1cm} \end{table} Interestingly, in an investigation that extended that~of Kres\'ak (1982), Lindblad (1985) found that an updated set of long-period comets displayed, at best, only margin\-ally greater orbital similarity than did random samples.{\hspace{0.3cm}} From the statistical standpoint, the pair of C/1846~O1 and C/1973~D1 appears to be something of an oddball; at first sight, the orbital differences are not so plainly minute as those of obvious fragments of a common parent (as, e.g., C/1988~F1 and C/1988~J1; Table~2), yet they are not conforming to a random distribution so readily as the other fortuitous pairs on Kres\'ak's (1982) list. Comparison of the 1846--1973 pair's $D$-value of 0.084 with those in Table~1 suggests that this pair is orbitally bound together one order of magnitude less tightly than the 1988--2015 trio and two orders of magnitude less tightly than the 1988 pair. On the other hand, it should be emphasized that some members of the Kreutz system, although genetically related beyond any doubt, have orbits far less similar and their $D$-values much larger than the 1846--1973 pair. For example, \mbox{$D = 0.223$} for the pair of C/1963~R1 (Pereyra) and C/1965~S1 (Ikeya-Seki), both Kreutz sungrazers with very accurately determined orbits, although classified by Marsden (1967) as members of different Kreutz subgroups. \begin{table}[b] % \vspace{-3.8cm} \hspace{4.22cm} \centerline{ \scalebox{1}{ \includegraphics{t2_deVico.ps}}} % \vspace{-19.18cm} \end{table} In Paper 1 we noted that the 1988 pair (C/1988 F1 and C/1988 J1) was with high probability a single comet less than one half the orbital period before discovery, while the parent of the 1988--2015 trio was likely to have split near the previous perihelion passage. The process of fragmentation of the Kreutz system began at least two (Sekanina \& Chodas 2004){\vspace{-0.07cm}} and possibly many more (Marsden 1989; \"{O}pik 1966) revolutions about the Sun before the 19th and 20th century clusters were observed. It appears that the $D$-criterion, as a measure of orbital similarity, increases with the time elapsed since the fragmentation event, a trend that is by no means surprising. By the same token, however, the $D$-criterion proves an unreliable tool in an effort to investigate a genetic relationship between two particular comets of unknown histories, and for other than statistical purposes appears to be useless. Indeed, in meteor astronomy --- for which the $D$-criterion was developed --- its application is limited to statistical studies only and is therefore fully justified. This experience suggests that a much more rigorous approach --- pursued below --- is required in order to gain insight into the fundamental issue of our interest:\ {\it Are comets C/1846~O1 and C/1973~D1 genetically related?}

1607.03440

1607

1607.03489_arXiv.txt

As far as we are aware, this represents the first case of high-resolution spectra obtained from precisely picked-out small areas across stellar surfaces. Although the method is observationally quite demanding, requiring to combine both high-precision photometry, accurate radial velocities, and high-fidelity spectra, it is thus feasible already with existing facilities (at least for brighter stars). The method will become much more practical with the impending advent of new spectrometers at the largest telescopes \citep{pepeetal14, strassmeieretal15, zerbietal14}. And, in particular, the numerous ongoing and planned photometric surveys for transiting exoplanets are likely to discover additional bright host stars. Such might be of also special spectral types: metal-poor ones, rapidly rotating, with strong stellar winds, or other. If the planet would happen to cross a starspot, even spatially resolved spectra (with their magnetic signatures) of such surface features would become attainable, given sufficient spectral quality. Observing time for such studies is virtually guaranteed since spectroscopy of exoplanets transiting bright stars is a high-priority task in studies of exoplanet atmospheres, and data required for stellar analyses will be obtained concurrently. Finally -- and perhaps not less important -- is to point out that 3-dimensional stellar simulations should include also predictions of spatially resolved spectral line profiles. Such have not commonly been calculated, presumably because their practical observability has not been realized. Once spectral data of sufficiently high fidelity are obtained, those might well become the most sensitive diagnostics for such modeling.

1607.03489

1607

1607.03506_arXiv.txt

We analyze the coupled quintessence in the light of the linear dynamical systems theory, with two different interactions: i) proportional to the energy density of the dark energy and ii) proportional to the sum of the energy densities of the dark matter and dark energy. The results presented here enlarge the previous analyses in the literature, wherein the interaction has been only proportional to the energy density of the dark matter. In the first case it is possible to get the well-known sequence of cosmological eras. For the second interaction only the radiation and the dark energy era can be described by the fixed points. Therefore, from the point-of-view of the dynamical system theory, the interaction proportional to the sum of the energy densities of the dark matter and dark energy does not describe the universe we live in.

Sixty eight percent of our universe \cite{Ade:2015xua} consists of a still mysterious component called ``dark energy'' (DE), which is believed to be responsible for the present acceleration of the universe \cite{reiss1998, perlmutter1999}. In addition to ordinary matter, the remaining $27\%$ of the energy content of the universe is a form of matter that interacts in principle only gravitationally, known as dark matter (DM). Among a wide range of alternatives for the dark energy, which includes the cosmological constant, scalar or vector fields \cite{ArmendarizPicon:2000dh,Padmanabhan:2002cp,Bagla:2002yn,Brax1999,Copeland2000,Landim:2015upa,ArmendarizPicon:2004pm,Koivisto:2008xf,Bamba:2008ja,Emelyanov:2011ze,Emelyanov:2011wn,Emelyanov:2011kn,Kouwn:2015cdw}, holographic dark energy \cite{Hsu:2004ri,Li:2004rb,Nojiri:2005pu,Pavon:2005yx,Wang:2005jx,Wang:2005pk,Wang:2005ph,Wang:2007ak,Landim:2015hqa,Li:2009bn,Li:2009zs,Li:2011sd,Wang:2016och}, metastable dark energy \cite{Stojkovic:2007dw,Landim:2016isc,Greenwood:2008qp,Abdalla:2012ug,Shafieloo:2016bpk}, modifications of gravity and different kinds of cosmological fluids \cite{copeland2006dynamics, Nojiri:2010wj, Bamba:2012cp, dvali2000, yin2005,Jamali:2016zww,Capozziello:2013bma}, the usage of a canonical scalar field, called ``quintessence'', is a viable candidate \cite{peebles1988,ratra1988,Frieman1992,Frieman1995,Caldwell:1997ii}. In addition, the two components of the dark sector may interact with each other \cite{Wetterich:1994bg,Amendola:1999er,Farrar:2003uw,Guo:2004vg,Cai:2004dk,Guo:2004xx,Bi:2004ns,Gumjudpai:2005ry,yin2005,Wang:2005jx,Wang:2005pk,Wang:2005ph,Wang:2007ak,micheletti2009,Costa:2014pba,Shahalam:2015sja,Nunes:2016dlj,Sola:2016ecz} (see \cite{Wang:2016lxa} for a review) and the interaction can eventually alleviate the coincidence problem \cite{Zimdahl:2001ar,Chimento:2003iea}. When a scalar field is in the presence of a barotropic fluid (with equation of state $w_m=p_m/\rho_m$, where $p_m$ is the pressure and $\rho_m$ is the energy density of the fluid) the relevant evolution equations can be converted into an autonomous system. Such approach is a good tool to analyze asymptotic states of cosmological models and it has been done for uncoupled dark energy (quintessence, tachyon field, phantom field and vector dark energy, for instance \cite{copeland1998,ng2001,Copeland:2004hq,Zhai2005,DeSantiago:2012nk,Azreg-Ainou:2013jxa,Landim:2016dxh,Alho:2015ila}) and coupled dark energy \cite{Amendola:1999er,Gumjudpai:2005ry,TsujikawaGeneral,amendola2006challenges,ChenPhantom,Mahata:2015lja,Landim:2015poa,Landim:2015uda}. The coupling assumed for the quintessence field has been proportional to the energy density of the dark matter $\rho_m$. However, there are other possibilities as for instance the coupling proportional to the energy density of the dark energy $\rho_\phi$ or the sum of the two energy densities $\rho_m+\rho_\phi$. Similar kernels have been widely studied in the literature \cite{Abdalla:2007rd,He:2008tn,He:2008si,Valiviita:2008iv,Abdalla:2009mt,Gavela:2009cy,He:2010im,Marcondes:2016reb}. In particular, the dark energy evolution at high redshifts measured by the BOSS-SDSS Collaboration \cite{Delubac:2014aqe} shows a deviation from the cosmological constant which can be explained assuming interacting dark energy models \cite{Abdalla:2014cla}. A dynamical analysis remained to be done for these two kernels. In this paper we use the linear dynamical systems theory to investigate the critical points that come from the evolution equations for the quintessence, assuming the interaction between DE and DM proportional to i) $\rho_\phi $ and ii) $\rho_\phi+\rho_m$. We found that in the case i) there are fixed points that can describe the sequence of three cosmological eras. In the second case either radiation era or dark energy era can be described by fixed points, but the matter-dominated universe is absent. The remainder of this paper is structured as follows. In Sect. \ref{de} we present the basics of the interacting dark energy and the dynamical analysis theory. The quintessence dynamics is presented in Sect. \ref{quint} and the dynamical system theory is used to study the coupled quintessence in Sect. \ref{AS}, wherein the critical points are shown. We summarize our results in Sect. \ref{conclu}. We use Planck units ($\hbar=c =M_{pl}=1$) throughout the text.

In the light of the linear dynamical systems theory we have studied coupled quintessence with dark matter with two different interactions: i) proportional to the energy density of the dark energy $\rho_\phi$ and ii) proportional to the sum of the two energy densities $\rho_m+\rho_\phi$. The results presented here enlarge the previous analysis in the literature, wherein the interaction has been only proportional to the energy density of the dark matter. In the case i) the transition of cosmological eras is fully achieved with a suitable sequence of fixed points. In the second case either radiation era or dark energy era can be described by the fixed points, but not the matter-dominated universe. Therefore, the second interaction does not provide the cosmological sequence: radiation $\rightarrow$ matter $\rightarrow$ dark energy. This is not the first time that an interaction proportional to the sum of the energy densities leads to cosmological disasters. A phenomenological model with that coupling suffers early-time instability for $w_d>-1$, as shown in \cite{Valiviita:2008iv,He:2008si}. Further analysis for high redshifts and different coupling are summarized in \cite{Wang:2016lxa}.

1607.03506

1607

1607.06005_arXiv.txt

We determine the asymptotic conditions under which the Boussinesq approximation is valid for oscillatory convection in a rapidly rotating fluid. In the astrophysically relevant parameter regime of small Prandtl number, we show that the Boussinesq prediction for the onset of convection is valid only under much more restrictive conditions than those that are usually assumed. In the case of an ideal gas, we recover the Boussinesq results only if the ratio of the domain height to a typical scale height is much smaller than the Prandtl number. This requires an extremely shallow domain in the astrophysical parameter regime. Other commonly-used ``sound-proof'' approximations generally perform no better than the Boussinesq approximation. The exception is a particular implementation of the \PI\ approximation, which predicts the correct instability threshold beyond the range of validity of the Boussinesq approximation.

Most astrophysical objects contain regions in which heat is transported by convection. The numerical modelling of these convective flows (which are usually turbulent) is difficult because of the stiffness of the governing equations caused by the presence of acoustic waves. Almost all convection models therefore filter out these waves by using a ``sound-proof'' set of equations, such as the Boussinesq, anelastic or \PI\ equations. Each of these sound-proof approximations is founded on certain physical assumptions, which may not be valid in all cases of interest. Specifically, the Boussinesq equations \citep[\eg][]{SpiegelVeronis60} are valid only for small perturbations to the thermodynamic variables in systems with small vertical lengthscales (in particular, the domain height must be much smaller than all of the thermodynamic scale heights). The anelastic equations require less restrictive assumptions, but do require the fluid to be nearly adiabatically stratified \citep[\eg][]{OguraPhillips62,LippsHemler82}. The \PI\ equations were introduced by \citet{Durran89} as an improvement to the anelastic equations. Although they are formally valid only under the same conditions as the anelastic equations, albeit for stratifications stronger than anticipated by standard asymptotics \citep{Klein-etal10}, in some cases they are found to better approximate the true dynamics \citep[\eg][]{Achatz-etal10}. This situation is further complicated by the fact that there are several different versions of both the anelastic and \PI\ equations currently in use, with no general consensus on which is ``best'' \citep[\eg][]{Brown-etal12,Vasil-etal13}. The interiors of stars and gaseous planets are characterised by rapid rotation and low viscosity. In this parameter regime, convection is often oscillatory close to onset \citep[\eg][]{Jones-etal09}. The first studies of oscillatory convection were performed under the Boussinesq approximation \citep{Chandrasekhar53}, but it has never been determined in precisely what asymptotic limit the Boussinesq and fully compressible results agree. Perhaps surprisingly, some implementations of the anelastic approximation exhibit unphysical ``negative Rayleigh number'' convection in this parameter regime \citep{Drew-etal95,Calkins-etal15}. As other implementations do not exhibit this unphysical behaviour \citep{Jones-etal09}, it seems that oscillatory convection is an important test case for comparing different sound-proof models. In this paper, we perform a careful analysis of the onset of oscillatory convection in the fully compressible system, in order to determine the precise conditions under which the Boussinesq results are valid. We find that these conditions are much more stringent than anticipated from the usual heuristic arguments. In particular, it is \emph{not} sufficient that the vertical scale of the domain is much smaller than the thermodynamic scale heights. This analysis is then extended to a very general set of sound-proof equations, which includes the anelastic and \PI\ approximations as special cases. In doing so, we introduce a simple but powerful technique that can be used to standardise the anelastic and \PI\ equations, building on an observation of \citet{OguraPhillips62} about energy conservation, thus removing any ambiguity in their formulation. Our standardised anelastic and \PI\ equations are both free from the unphysical behaviour noted by \citet{Calkins-etal15}. However, the standardised \PI\ equations are the only ``sound-proof'' system that correctly predicts the threshold for oscillatory convection on larger vertical scales than the Boussinesq approximation.

\label{sec:discussion} Although we have referred to \eqs(\ref{eq:PI_mom})--(\ref{eq:PI_T}) as ``sound-proof'', we have not explicitly shown that acoustic oscillations are absent from these equations. That this is the case can be demonstrated, for example, by deriving their exact dispersion relation for perturbations about some idealised background state. However, it can also be deduced simply from \eq(\ref{eq:PI_rho}), which describes how the volume of a displaced parcel of fluid instantaneously adjusts to a value determined by the entropy inside the parcel and the local properties of the background state. By neglecting the effect of pressure perturbations on this adjustment, we remove the dynamical degree of freedom that permits the parcel to oscillate acoustically. A mathematical proof of this statement, which implicitly assumes the conservation of mass and entropy, has been given by \citet{Durran08}. The direct relation between $\rho_1$ and $s_1$ imposes a constraint on the velocity field, which can be deduced from \eqs(\ref{eq:PI_mass})--(\ref{eq:PI_rho}): \begin{equation} \frac{1}{\alpha H_{ps}\rho_0^2}\bnabla\bcdot(\rho_0\boldsymbol{u}) + \boldsymbol{u}\bcdot\bnabla s_0 = \frac{q_1}{T_0}. \label{eq:constraint} \end{equation} Although $\alpha$ can be taken as any function of height, there are two significant special cases, which correspond to the particular choices $\alpha=0$ and $\alpha=1$. When $\alpha=0$, \eq(\ref{eq:constraint}) reduces to the anelastic velocity constraint (\ref{eq:anelastic2}) and, in fact, our equations become almost identical to the particular version of the anelastic equations derived by \citet{Lantz92} and by \citet{BraginskyRoberts95}, which for brevity we will call the LBR equations. However, an important difference is that in our sound-proof system, with $\alpha=0$, \eq(\ref{eq:PI_rho}) states that $\rho_1=0$, whereas in the LBR equations $\rho_1$ is given as a function of the other thermodynamic perturbations. Formally this implies that mass is not conserved in the LBR equations, whereas mass conservation is built in to our linearised system. Nevertheless, because $\rho_1$ does not appear explicitly in \eqs(\ref{eq:PI_mom}) and (\ref{eq:PI_T}), our equations are in fact mathematically equivalent to the LBR equations, and the only difference is conceptual. \citet{Lantz92} argued that these equations are the most natural generalisation of the Boussinesq equations to a density-stratified background and, indeed, we see that \eq(\ref{eq:N2_alpha}) becomes identical to its Boussinesq counterpart (\ref{eq:Bous}) in the case where $\alpha=0$ (this means that the curves in \fig\ref{fig:Bushby} labelled as ``Boussinesq'' also correspond to the $\alpha=0$ case of our sound-proof equations). A significant practical advantage of the LBR equations over other sound-proof approximations is that they can be solved without ever having to explicitly calculate the pressure perturbation $p_1$, provided that perturbations to the heat flux are calculated from the gradient of $s_1$ rather than $T_1$ \citep{Lantz92,BraginskyRoberts95}. In previous derivations of the LBR equations, this approximation has been justified by appealing to ``turbulent diffusion'' of entropy. However, setting $\alpha=0$ implies a direct relation between $T_1$ and $s_1$ in \eq(\ref{eq:PI_T}). So in our version of the anelastic system there is no need to approximate the heat flux; instead we have an approximate equation of state (\ref{eq:PI_T}) whose form is dictated by the energy principle (\ref{eq:energy}). Whilst the choice $\alpha=0$ reduces \eqs(\ref{eq:PI_mom})--(\ref{eq:PI_T}) to the LBR anelastic equations, \eq(\ref{eq:N2_alpha}) suggests that the ``best'' sound-proof model is actually given by $\alpha = 1$, because in that case we exactly recover the corresponding compressible result (\eq\ref{eq:N2}) from \eq(\ref{eq:N2_alpha}). (The choice $\alpha = -1$ can be discounted on physical grounds, because it would imply a positive correlation between $\rho_1$ and $s_1$.) If we set $\alpha=1$ then \eqs(\ref{eq:PI_mom})--(\ref{eq:PI_T}) are equivalent to the linearisation of the ``thermodynamically consistent'' version of the \PI\ equations obtained by \citet{KleinPauluis12} (see also the ``generalized \PI'' equations of \citet{Vasil-etal13}). In fact, the results of \citet{KleinPauluis12} suggest how the method that we have proposed for standardising sound-proof models can be extended into the nonlinear regime. This is an essential step towards modelling systems that are well above the onset of convection, such as the interiors of stars, but is beyond the scope of the current paper. A very similar set of \PI\ equations was recently studied numerically by \citet{Lecoanet-etal14}. An important difference, however, is in the definition of the temperature perturbation $T_1$. We believe that many of the discrepancies \citet{Lecoanet-etal14} found between their solutions of the \PI\ equations and the fully compressible equations result from the incorrect definition of $T_1$ that they used, but further work will be required to confirm this. Having said that, not all of the issues with the \PI\ equations identified by \citet{Lecoanet-etal14} can be solved simply by redefining $T_1$. In particular, the velocity constraint (\ref{eq:constraint}) becomes ill-posed for horizontally-invariant perturbations in a fluid with impenetrable, thermally conducting boundaries. This difficulty does not arise in the stability analysis presented here, because convective motions necessarily have a finite horizontal lengthscale. In general, the \PI\ equations seem best suited to describing fluid motions on small horizontal scales, and are less accurate for motions on large horizontal scales \citep[see \eg][]{Durran08}. One possible resolution is to allow the background state to be time dependent \citep[\eg][]{ONeillKlein14}, but this is beyond the scope of the analysis presented here. We would like to thank Professor Rupert Klein, as well as the anonymous referees, for their helpful comments and suggestions.

1607.06005

1607

1607.03416_arXiv.txt

X-ray bursts are thermonuclear flashes on the surface of accreting neutron stars and reliable burst models are needed to interpret observations in terms of properties of the neutron star and the binary system. We investigate the dependence of X-ray burst models on uncertainties in (p,$\gamma$), ($\alpha$,$\gamma$), and ($\alpha$,p) nuclear reaction rates using fully self-consistent burst models that account for the feedbacks between changes in nuclear energy generation and changes in astrophysical conditions. A two-step approach first identified sensitive nuclear reaction rates in a single-zone model with ignition conditions chosen to match calculations with a state-of-the-art 1D multi-zone model based on the {\Kepler} stellar evolution code. All relevant reaction rates on neutron deficient isotopes up to mass 106 were individually varied by a factor of 100 up and down. Calculations of the 84 highest impact reaction rate changes were then repeated in the 1D multi-zone model. We find a number of uncertain reaction rates that affect predictions of light curves and burst ashes significantly. The results provide insights into the nuclear processes that shape X-ray burst observables and guidance for future nuclear physics work to reduce nuclear uncertainties in X-ray burst models.

\label{sect:introduction} Type I X-ray bursts are the most frequently observed thermonuclear explosions in nature \citep{Schatz2006a,Strohmayer2006,Lewin93,Parikh2013}. They take place on the surface of accreting neutron stars in low-mass X-ray binary systems for a certain range of mass transfer rates, generally within two orders of magnitude below the Eddington mass accretion rate ($\dot{M}_\mathrm{Edd}\approx2\times10^{-8}\,\mathrm{M}_{\odot}/\textrm{yr}$; \citealt{Woosley1976,Joss1977,Fujimoto1981,Strohmayer2006}). Since these events are not cataclysmic, bursts will repeat with recurrence times ranging from hours to days. The importance of understanding X-ray bursts as a probe of neutron star properties and the underlying physics has been discussed extensively~\citep{Lewin93,Steiner2010,Zamfir2012,Ozel2013,Guver2013}. The bursts are powered by the triple-$\alpha$ reaction, the $\alpha$\textsl{p}-process and the rapid proton capture process (\textsl{rp}-process; \citealt{Wallace1981, VanWormer1994,Schatz1998,Schatz2001,Fisker2008,Woosley2004a,Jose2010}). These nuclear processes involve hundreds of nuclear species from stable isotopes to the proton drip line. Models predicting burst light curves and the composition of the burst ashes are therefore sensitive to a broad range of uncertain nuclear structure and reaction properties near and beyond the current frontier of experimental knowledge. This limits interpretation of the vast body of observational data that has been accumulated \citep{Galloway2008}, for example in the MINBAR data base which will eventually contain X-ray light curve data of over 5000 X-ray bursts \citep{Galloway2010}. Models with reliable nuclear physics are needed to validate the astrophysical model assumptions through comparison with observations, to guide future model developments towards an understanding of the full variety of observed bursting behavior, and to constrain parameters such as distance, accretion rate, accreted composition, and neutron star properties \citep{Heger2007,Galloway2004}. An example for the latter are recent attempts to constrain the neutron star surface gravity by matching a set of model bursts to observed light curves \citep{Zamfir2012}. Whereas the X-ray light curve is the main direct observable of X-ray bursts, accurate nuclear physics is also needed to predict the composition of the burst ashes. Reliable calculations of this composition are required to predict potential spectroscopic signatures in the X-ray burst light curve from small amounts of ejected material \citep{Weinberg2006} and to predict the composition of the neutron star crust, which in mass accreting systems is made in part or entirely out of X-ray burst ashes. Of particular importance is the amount of $^{12}$C, which may reignite at greater depth \citep{Cumming2001,Strohmayer2002,Cumming2006} and explain the origin of occasionally observed superbursts \citep{Keek2008,Keek2011}. The thermal transport properties of the neutron star crust, as well as the amount of heating and cooling through weak interaction processes \citep{Haensel2008,Gupta2007, Schatz2014} also depend sensitively on composition. The goal of this paper is to identify important reaction rates that affect observables and nucleosynthesis. Besides nuclear reaction rates, burst models also depend on $\beta$-decay rates and masses of neutron deficient nuclei. All relevant $\beta$-decay rates have been determined experimentally \citep{Schatz2006a}, and corrections due to the high densities and temperatures reached in X-ray bursts are predicted to be small in most cases \citep{FFN1982,Pruet2003}. However, the validity of these corrections remains to be evaluated in detail. The majority of the relevant nuclear masses has been measured as well. The impact of remaining mass uncertainties has been discussed elsewhere \citep{Schatz2006,Parikh2009,Kankainen2012}. However, the rates for the vast majority of nuclear reactions occurring in X-ray bursts have not been determined experimentally and have large uncertainties. We therefore focus here on the sensitivity of X-ray burst models to nuclear reaction rates to provide guidance for future experimental efforts aimed at reducing the remaining nuclear physics uncertainties in X-ray burst models. Comprehensive studies of the sensitivity of X-ray burst models to nuclear reaction rates have so far been limited to single-zone post-processing studies. \citet{Parikh2008} varied 3500 nuclear processes individually, and concurrently in a Monte Carlo approach, using fixed temperature and density profiles from various previously published X-ray burst models. An important result was the lack of additional sensitivities due to correlations among reaction rates in the Monte Carlo approach, justifying the single reaction rate variation approach to nuclear sensitivity studies for X-ray bursts that we also adopt here. Post-processing studies with fixed temperature and density profiles, however, are not adequate to study nuclear sensitivities in X-ray bursts \citep{Thielemann2001} because the entire reaction sequence contributes to the energy release driving the burst. Changes in reaction rates inevitably lead to changes in energy production, thereby to changes in the temperature and density evolution of the burst sequence. It is therefore essential to ensure consistency between the nuclear physics input and the temperature and density evolution. Owing to computational limitations, however, sensitivity studies in fully self-consistent dynamic X-ray burst models have so far been limited to variations of a few individual reactions or a few groups of reactions. \citet{Thielemann2001} varied the proton capture rates on \iso{Si}{27}, \iso{S}{31}, \iso{Ar}{35}, and \iso{Ca}{38} together, rather than individually, in a 1D multi-zone X-ray burst model and demonstrated their strong influence on the light curve, in contrast to a post processing study that found these reactions to be unimportant \citep{Iliadis1999}. \citet{Fisker2004a} varied the \iso{S}{30}($\alpha$,p) and \iso{Ar}{34}($\alpha$,p) reaction rates and demonstrated their impact on doubly peaked burst profiles. \citet{Woosley2004a} varied groups of $\beta$-decay rates to simulate the potential impact of proton capture rate uncertainties on the burst light curve and found a very strong sensitivity. However, this only provides a crude estimate of how the overall processing speed of the rp-process may affect observables. It does not provide any insights into which proton capture rates may be important, nor whether proton capture rates are important at all. \citet{Fisker2006} determined a new lower limit of the \iso{O}{15}($\alpha$,$\gamma$) reaction rate and showed that bursts disappear altogether when that reaction rate is at that lower limit. \citet{Davids11}, however, did not find such an effect using a different X-ray burst model. More recently \citet{Keek2014} explored the impact of variations of the 3$\alpha$, $^{15}$O($\alpha$,$\gamma$) and $^{18}$Ne($\alpha$,$\gamma$) reaction rates on the transition from unstable to stable burning in a multizone X-ray burst model. Here we present a comprehensive study of the sensitivity of X-ray burst models to 1931 nuclear reaction rates using for the first time self-consistent X-ray burst models that account for the coupling between nuclear energy generation and the astrophysical conditions that determine the reaction sequences. Our approach is enabled by the use of two models. A calibrated, self-consistent one-zone model \citep{Schatz2001} is used to explore variations of all reaction rates. A subset of relevant rates is then investigated further using a state-of-the-art 1D multi-zone burst model based on the \textsc{Kepler} code \citep{Woosley2004a}.

We present results from the first large-scale investigation of the influence of nuclear reaction rate uncertainties on X-ray burst light curves and ashes that uses a self-consistent multi-zone 1D X-ray burst model. This approach fully accounts for the influence of reaction rate changes on temperature and density conditions, radiation transport, and compositional inertia along a sequence of bursts. The results serve as a road map for nuclear experimental and theoretical work towards reducing nuclear uncertainties in X-ray burst models that describe bursts in the common mixed hydrogen and helium burning regime. The most important reaction rate uncertainties identified here are the ones that affect the burst light curve (listed in Tab.~\ref{tbl:MZ_LCrank}), the major components of the ashes, $^{12}$C production, and the production of odd mass isotopes (Tab.~\ref{tbl:MZ_ABrank}). We note that the light curve variations we find from varying a number of single reaction rates with reasonable multiplication factors are comparable or larger than the effects of different surface gravities \cite{Zamfir2012}. Clearly reaction rate uncertainties have to be addressed before burst light curve tails can be used to constrain surface gravity and therefore neutron star compactness. We also identified a number of reactions that are not important (Fig.~\ref{FigMZ_path}). We also used a one-zone X-ray burst model as closely matched to {\Kepler} ZM as possible to maximize the number of identified critical reactions with a limited number of multi-zone calculations. For this one-zone model we identified the complete set of important reaction rate uncertainties (Tab.~\ref{tbl:SZ_LCrank}, Tab.~\ref{tbl:SZ_ABrank} and Fig.~\ref{FigSZ_path}). This is the first complete identification of important reaction rates for mixed hydrogen and helium burning bursts in a fully self-consistent one-zone burst model. We emphasize that our goal was not to calculate realistic uncertainties in a statistical sense, but to flag important reaction rates, and provide one data point on the dependence of observables on this reaction rate. Modifications to the light curve and composition for different variation factors may then be roughly estimated using our results. For a precise analysis calculations would have to be repeated for a particular reaction rate uncertainty, and uncertainty correlations with temperature also had to be included. More work is needed to quantify the uncertainties of the important reaction rates identified in this work, and to develop targeted approaches to reduce these uncertainties using experiments or nuclear theory. In addition, more sensitivity studies along the lines of this work are needed to arrive at a complete picture of the nuclear physics needs for X-ray burst models. This includes variations of additional reaction rates in the multi-zone burst model investigated here, as the selection based on the single-zone model sensitivity may be incomplete. In addition, sensitivity studies of burst in other burning regimes should be performed.

1607.03416

1607

1607.00848_arXiv.txt

We present new maps of emission-line flux distributions and kinematics in both ionized (traced by H\,{\sc i} and \feii\ lines) and molecular (\hm) gas of the inner 0.7$\times$0.7~kpc$^2$ of the galaxy NGC\,4303, with a spatial resolution 40--80\,pc and velocity resolution 90--150\,\kms obtained from near-IR integral field specroscopy using the VLT instrument SINFONI. The most promiment feature is a 200--250\,pc ring of circum-nuclear star-forming regions. The emission from ionized and molecular gas shows distinct flux distributions: while the strongest H\,{\sc i} and \feii\ emission comes from regions in the west side of the ring (ages $\sim$4\,Myr), the \hm\ emission is strongest at the nucleus and in the east side of the ring (ages $>$10\,Myr). % We find that regions of enhanced hot H$_2$ emission are anti-correlated with those of enhanced \feii\ and H\,{\sc i} emission, which can be attributed to post starburst regions that do not have ionizing photons anymore but still are hot enough ($\approx$\,2000\,K) to excite the H$_2$ molecule. The line ratios are consistent with the presence of an AGN at the nucleus. The youngest regions have stellar masses in the range 0.3-1.5$\times10^5$\,M$_\odot$ and ionized and hot molecular gas masses of $\sim0.25-1.2\times 10^{4}$~M$_\odot$ and $\sim2.5-5~M_\odot$, respectively. The stellar and gas velocity fields show a rotation pattern, with the gas presenting larger velocity amplitudes than the stars, with a deviation observed for the \hm\ along the nuclear bar, where increased velocity dispersion is also observed, possibly associated with non circular motions along the bar. The stars in the ring show smaller velocity dispersion than the surroundings, that can be attributed to a cooler dynamics due to their recent formation from cool gas.

\label{intro} Star formation in the circumnuclear regions of galaxies and its connection with the existence of nuclear young stellar clusters and/or Active Galactic Nuclei (AGN) has been the subject of many studies over the past several decades since early models invoking dynamical resonances in a rotating bar potential \citep{combes85}, bars within bars \citep{sho89}, and the direct feeding of AGN due to stellar winds or cloud-cloud collisions in the vicinity of the nucleus \citep{norman88,sho90}. Some models \citep[e.g.][]{heller94,knapen95} suggest that gas could flow inwards from the ring, creating a disk of gas that, if massive enough, would become unstable. Under this scenario, a massive black hole in the nucleus could be fed triggering an AGN \citep{sho89,fukuda98}. Regarding star formation in nuclear rings, two scenarios have been proposed: the ``pop-corn'' \citep{elmegreen94} and the ``pearls on a string'' \citep{boker08}. The ``pop-corn'' scenario assumes that the cold molecular gas is accumulated in a circumnuclear resonance ring. If massive enough, the ring becomes gravitational unstable, fragmenting in clumps and forming stellar clusters at random positions. On the other hand, in the ``pearls-on-a-string" scenario, new stars are exclusively formed in the regions where the gas enters the ring (i.e. the regions of maximum gas density). These young clusters evolve passively as they orbit along the ring, producing a string of aging clusters. Independent on the formation mechanism of the stellar clusters, the role of these clusters as they move along the ring can be important. Stellar winds and subsequent supernovae explosions can affect their surrounding interstellar medium, removing gas and halting subsequent star formation, as well as generating shocks that will reduce the angular momentum of the gas, that could eventually fall towards the center, feeding the AGN and/or forming a nuclear star cluster. Detailed multi-wavelength two-dimensional spectroscopic studies in nearby galaxies with the adequate spatial resolution ($\sim$10 pc, or less) are needed to further investigate the evolutionary scenarios mentioned above. CO interferometric maps (PdBI) have shown the presence of a wide range of molecular gas structures in the central kpc region of nearby galaxies with an AGN \citep{garcia-burillo05}. According to these studies, most of the molecular gas is concentrated in the form of a circumnuclear ring of several hundreds pc to kpc size, while $\approx$33\% of the galaxies show the evidence for a direct gas fueling into the AGN down to scales of $\sim$50 pc \citep{garcia-burillo12}. Of particular importance, the near-infrared (near-IR) bands allow the study of the multi-phase gas, from molecular (H$_2$) to shocked partially-ionized (Fe\,{\sc ii}) to highly ionized gas (e.g. Ca\,{\sc viii}), that can trace a number of different physical structures and mechanisms, from molecular gas reservoirs to AGN outflows. In the near-IR, recent studies using integral field spectroscopy \citep[IFS,][]{genzel95}, have been particularly useful to map the stellar and gas kinematics, as well the gas excitation and distribution of the different gas phases \citep[e.g.][]{n1068-exc,barbosa14}. In nearby Seyfert galaxies, compact (scales of few tens of parsecs) molecular gas disks \citep{mazzalay14,mrk1066-kin,hicks09}, streaming motions towards the nucleus \citep[e.g.][]{davies14,diniz15}, ionized gas outflows \citep[e.g.][]{iserlohe13,sb10} and young stellar populations \citep[e.g.][]{davies07} have been mapped in the nuclear regions. In addition, studies of circumnuclear star-forming rings at hundred of pcs from the nucleus in nearby spirals \citep{boker08,falcon-barroso14,laan13a,laan15} have focused on establishing the reality of the proposed evolutionary scenarios for the star-forming rings, such as the ``pop-corn'' scenario \citep{elmegreen94} and the ``pearls-on-a-string" scenario \citep{boker08}. All these previous studies have been focused on either nearby luminous Seyfert galaxies where the output energy is dominated by the AGN, or galaxies with luminous circumnuclear star-forming rings. Here we perform a similar study for NGC\,4303, a nearby galaxy with both a (low-luminosity) AGN and a star-forming ring, and that has also a young massive cluster at the nucleus. In addition, multi-wavelength observations \citep{colina97,colina99,colina02} reveal a spiral of circumnuclear star-forming regions (CNSFRs) that can be traced all the way into the inner few parsecs \citep[e.g.][]{colina00,jimenez03}, suggesting it could be the feeding channel to the AGN and nuclear star cluster. Finally, a high-velocity nuclear outflow extending up to $\approx$120\,pc to the north-east of the nucleus has also been observed in optical line emission of [O\,{\sc iii}] \citep{colina99}. NGC\,4303 is at a distance of 16.1~Mpc \citep{colina99}, and is classified as a SB(rs)bc \citep{vaucouleurs91}. It is abundant in cold molecular gas \citep{schinnerer02} that shows the global distribution expected for the gas flow in a strong, large-scale bar, and the two-arm spiral structure in the inner kiloparsec can be explained by a density wave activated by the potential of that bar. In this paper, we present near-infrared IFS of the inner 350 pc radius of NGC\,4303, obtained with the integral field spectrograph SINFONI at the VLT in order to map the kinematics and excitation properties of the different phases of the interstellar medium in the circumnuclear region. It is the first time that such mapping is provided in the near-IR. This work is organized as follows: Section~2 presents the observations and data reduction procedure, while in Sec.~3, we present maps for emission-line flux distributions and ratios, as well as for the gas and stellar kinematics. The results are discussed in Sec.~4 and Sec.~5 presents the conclusions of this work.

We have presented new emission-line flux and velocity maps, as well stellar velocity maps of the inner 0.7 kpc$\times$0.7 kpc of the nearby spiral galaxy NGC\,4303, at spatial resolutions of 78 pc, 47 pc, and 36 pc, for the J, H and K-bands, respectively, using near-IR integral field spectroscopy with the VLT instrument SINFONI. The observations cover the nucleus and circumnuclear star-forming ring (CNSFR) with radius $\approx$\,200--250\,pc. The main conclusions of this work are: \begin{itemize} \item The near-IR emission-line flux distributions delineate the CNSFR, with emission-line knots in the flux distributions of the different gas phases observed at different locations along the ring. The \hm\ and \feii\ emission lines show emission peaks at the nucleus, while the H\,{\sc i} recombination lines are dominated by emission from the CNSFR; \item Along the ring, the strongest H\,{\sc i} and \feii\ emission are observed mostly in the west side of the ring and at a region to the south-east, while the strongest emission in \hm\ seems to be anti-correlated with them, being observed mostly to the east; \item The properties: \feii/\br\ and \hm/\br line ratios, EqW and $\sigma$ of the emission lines along the star-forming ring support an age difference between the west and east sides of the CNSFR, with the former being younger (2.5--7.5\,Myr) than the latter (10--25\,Myr); \item The distribution of the star-forming regions and their age differences do not support fully the ``pop-corn'' and the ``pearls-in-a-string" scenarios for star formation in CNSFRs. The star formation in the CNSFR of NGC\,4303 appears to be instead episodic with stars forming quasi-simultaneously over a large sector of the ring (covering $\sim$\,300\,pc along the ring), aging as they rotate with an orbital time of several Myr; \item Assuming the star-forming regions in the CNSFR are closer to instantaneous bursts with ages of about 4~Myr, we derive masses for the clusters in the range 0.3-1.5$ \times10^5$\,M$_\odot$. The corresponding masses for the associated ionized and hot molecular gas are about $\sim0.25-1.2 \times 10^{4}$~M$_\odot$ and $\sim0.25-0.5 \times 10^{1}$~M$_\odot$, respectively. For stellar populations with older ages of up to 10 Myr, the corresponding stellar masses will increase by up to factor of hundred; \item The \feii\ emission shows an elongation to $\approx$120\,pc north-east of the nucleus that could be associated with the previously known optical ([OIII]) outflow; \item The \hm\ emission shows an elongation to $\approx$120\,pc west of the nucleus that could be the hot counterpart of the already known cold molecular gas ``bridge" that connects the nucleus with the large circumnuclear molecular gas reservoir; \item The near-IR emission-line ratios (\feii/\br\ and \hm/\br) of the nucleus are consistent with the presence of an AGN and/or a SNe-dominated star-forming region. Since there is no evidence for an aged stellar cluster in the nucleus, the line ratios are interpreted as due to the combined effect of X-ray radiation and shocks at the base of the ionization cone of the AGN. Higher angular resolution spectroscopy is required to further explore this scenario; \item The stellar velocity field is well reproduced by a model of a rotating disk with an inclination $i=33^\circ$ relative to the plane of the sky, major axis oriented along PA$\sim135^\circ$, and with a velocity amplitude of about 160\,\kms. The stars associated to the CNSFR show smaller velocity dispersion than the surroundings, revealing a cooler dynamical stellar population in the ring, consistent with their recent formation from cold gas; \item The gas velocity fields are also dominated by rotation, similar to that observed for the stars but with a larger amplitude. A significant deviation from rotation was observed for the \hm\ emission in a region extending $\approx$\,120\,pc to the south-west, along the orientation of the nuclear bar. A higher \hm\ velocity dispersion is also observed at this location and is attributed to the presence of more than one kinematic component associated with non-circular motions along the nuclear bar. \end{itemize} All fits files for the emission-line flux distributions, velocity fields and velocity dispersion maps are available online as supplementary material.

1607.00848

1607

1607.02884_arXiv.txt

{The optical morphology of galaxies holds the cumulative record of their assembly history, and techniques for its quantitative characterization offer a promising avenue toward understanding galaxy formation and evolution. However, the morphology of star-forming galaxies is generally dictated by the youngest stellar component, which can readily overshine faint structural/morphological features in the older underlying stellar background (e.g., relics from recent minor mergers) that could hold important insights into the galaxy build-up process. Stripping off galaxy images from the emission from stellar populations younger than an adjustable age cutoff \tcut\ can therefore provide a valuable tool in extragalactic research. RemoveYoung (\ry), a publicly available tool that is presented here, exploits the combined power of integral field spectroscopy (IFS) and spectral population synthesis (\sps) toward this goal. Two-dimensional (2D) post-processing of \sps\ models to IFS data cubes with \ry\ permits computation of the spectral energy, surface brightness, and stellar surface density distribution of stellar populations older than a user-defined \tcut. This suggests a variety of applications of star-forming galaxies, such as interacting or merging galaxy pairs and lower mass starburst galaxies near and far; these include blue compact and tidal dwarf galaxies. }

} How the assembly history of galaxies is imprinted on their present-day optical morphology is one of the most tantalizing enigmas in extragalactic research. Morphology holds the cumulative record of complex and highly interlinked processes operating on different temporal and spatial scales across cosmic time, such as quasi-monolithic gas collapse into classical bulges and galaxy spheroids, gentle gas dissipation into galactic disks, hierarchical growth via minor and major mergers, and environmentally modulated star formation (SF) in galaxy pairs/groups and clusters \citep[see][for a review]{KormendyKennicutt04}. Various quantitative morphology indicators \citep[\rem{QMIs}; see][for a review]{Conselice14} have been proposed in recent decades and extensively employed for the characterization of large extragalactic probes in the quest of elucidating the link between morphology, structure, intrinsic physical properties (e.g., stellar mass \mstar\ and surface mass density \sstar, metallicity) and the evolutionary and dynamical status of galaxies. A first-order approximation in these studies is that the optical surface brightness $\mu$ traces the \sstar, which essentially presumes that the stellar mass-to-light ratio (\ml) spans a rather narrow range of values across the galaxy extent. Whereas this assumption is certainly justified for quiescent galaxies or systems with a smooth star formation history (SFH), it cannot be maintained for systems exhibiting a high specific star formation rate (sSFR), such as isolated and interacting starburst galaxies. The optical appearance of such galaxies primarily reflects the 2D distribution of the young ($\la$100 Myr) stellar component, which owing to its very low \ml\ throughout overshines the older underlying stellar background that is dynamically dominant. As an example, in a typical blue compact dwarf (BCD) galaxy the centrally confined starburst component dictates the observed line-of-sight intensity and contributes 50\% to 80\% of the total optical emission \citep[][]{P96a,C01,GdPM2005}. In BCDs and their higher-$z$ analogs \citep[e.g., {\sl green peas};][]{Cardamone2009,Izotov2011,Amorin2012,Amorin2015} SF typically dominates down to $\mu \simeq$ 24.5-25.0 $B$ \sbb, i.e., almost out to the Holmberg radius. This strong disparity between $\mu$ and \sstar\ in star-forming galaxies has a twofold effect: First, starburst emission can strongly impact light concentration indices \citep[e.g.,][]{Morgan1958,Abraham1996} commonly used in \rem{QMI} studies. In a typical BCD, for instance, the ignition of a central starburst shrinks the optical effective radius by $\sim$70\% \citep{P06} and can mimick a S\'ersic profile with a high ($\eta$=2-4) exponent \citep[][see also Bergvall \& \"Ostlin 2002]{P96a}, eventually leading to its erroneous classification as an early-type galaxy. Secondly, a tiny substrate of luminous young stars can readily masquerade fainter morphological features that potentially hold key insights into the recent assembly history and structural properties of galaxies (e.g., shells and ripples as relics from minor mergers; cf. Schweizer \& Seitzer (1988) or intrinsically faint bars). \begin{figure*} \begin{picture}(18.4,6.7) \put(2.0,0.0){\includegraphics[width=12cm, angle=0, viewport=80 470 460 800]{Figures/PlotSpectra_reduced.pdf}} \PutLabel{10.6}{6.2}{\mcap a} \PutLabel{13.0}{6.2}{\mcap b} \PutLabel{13.80}{3.0}{\mcap c} \end{picture} \caption[]{ \rem{a)} Example of the application of \ry\ to the best-fitting synthetic SED (light blue) obtained with \SL\ for a SDSS-DR7 \citep{yor00} spectrum (orange) for three \tcut\ values: 30 Myr (blue), 0.5 Gyr (green), and 5 Gyr (red). The transmission curves of the five SDSS filters $u$, $g$, $r$, $i,$ and $z$ are depicted with shaded areas. Panels \rem{b} and \rem{c} show, respectively, the luminosity ($x_j$) and mass ($\mu_j$) contribution (\%) of individual SSPs (1 \dots $j$) to the best-fitting population vector. The light-gray vertical lines in the upper-right panel show the ages available in the SSP library for four metallicities (between 0.05 and 0.95 \zsun) and vertical arrows depict the applied \tcut\ values. From panels \rem{b} and \rem{c} it is apparent that, according to the best-fitting population vector, SSPs younger than 30 Myr contribute less than 2\% of \mstar\ but nearly 35\% of the observed intensity at the normalization wavelength of 4020 \AA. As a result, suppression of these young SSPs results in a significant dimming by 0.56, 0.27, 0.22, and 0.19 mag in the SDSS $u$, $g$, $r,$ and $i$ bands, respectively. } \label{fig:RY-spectrum} \end{figure*} In light of such considerations, a technique permitting suppression from galaxy images of the luminosity contribution from stars younger than an adjustable age cutoff \tcut\ appears especially useful. Such a tool would not only be valuable to studies of \rem{QMI}s, but also to those exploring SF patterns back to a well-defined time interval (e.g., since the infall of a galaxy onto a cluster). This task is obviously out of reach with standard techniques, such as near-infrared (NIR) imaging \citep[e.g.,][]{Noeske03} or state-of-the-art modeling of the observed spectral energy distribution \citep[SED; e.g.,][]{Wuyts2012}. In this article, we present a publicly available\footnote{A thoroughly documented version of the code is available at {\tt http://www.spectralsynthesis.org}.} tool, RemoveYoung (\ry), which exploits the combined power of integral field spectroscopy (IFS) and spectral population synthesis (\sps) with the goal of stripping off galaxy images from the luminosity contribution from stellar populations younger than a user-defined age cutoff. The concept is outlined in Sect.~\ref{RY-meth} and illustrated through its application to IFS data in Sect.~\ref{RY-IFS}, with a discussion of its potential merits to various subjects of extragalactic research following in Sect.~\ref{disc}.

} Before providing an outline of potential applications of \ry, we offer some cautionary notes: Evidently, the output from \ry\ sensitively relies on the quality and soundness of solutions from \sps\ models. These are known to suffer from a number of deficiencies, such as the notorious age-metallicity degeneracy, which may propagate into hardly predictable and as yet poorly explored biases in best-fitting SFHs. Additionally, the rather restricted number of library SSPs allowed by state-of-the-art \sps\ codes (at maximum 300 for \SL) permits a rather coarse coverage of the age and metallicity parameter space, which results in a strongly discretized approximation to the true SFH of a galaxy. Quite obviously, \ry\ cannot offer a better time resolution than that of the SSP library used for the input \sps\ models. Consequently, \ry\ can fully unfold its potential only in conjunction with next-generation \sps\ codes that are capable of significantly alleviating the above shortcomings. Another aspect to bear in mind when interpreting the output from \ry\ is that the detectability of morphological relics (e.g., past SF episodes) does not only depend on their spectrophotometric fading, but also on dynamical processes. For example, differential disk rotation acts to erase signatures of an accreted satellite within a few rotational periods ($\sim$250 Myr for a MW-like system). Likewise, \sstar\ enhancements due to minor mergers could gradually disperse in the presence of processes that eventually contribute to inside-out galaxy growth, such as bar-driven stellar migration \citep{SellwoodBinney2002,Berentzen07,Roskar08,SB14} in normal galaxies, diffusion of newly formed stars \citep{P02,PO12} in dwarf galaxies, or outwardly propagating SF \citep{G16c}. On the other hand, the briefly outlined applications of \ry\ on IFS data (Sect.~\ref{RY-IFS}) illustrate the potential of the code in various fields of extragalactic research. For example, a natural by-product of \ry\ is the removal of nebular emission and the quantification of its effect on broadband photometry. Already, with a conservative \tcut\ on the order of the main-sequence lifetime of ionizing stars (10-30 Myr), \ry\ permits partial suppression of bright SF regions from synthetic images; this opens new avenues to structural and \rem{QMI} studies of starburst galaxies where the young stellar component and nebular emission frequently dominate within the optical extent. Combined with deep IFS, \ry\ may thus be regarded as an analogous yet more powerful approach to the structural properties of high-sSFR systems than NIR photometry. This is not only because the latter is expensive in terms of observational time and data reduction effort, but also because it {\sl per se} does not permit complete suppression of stellar populations younger than an adjustable age cutoff. Additionally, the $K$-band \ml\ of an instantaneously formed stellar population varies by $\sim$2 dex within 100 Myr, depending on metallicity and initial stellar mass function, which implies that an accurate determination of \mstar\ requires even in NIR wavelengths prior knowledge of the SFH. The adjustable age cutoff allowed by \ry\ is a key advantage here, which together with an extensive set of output quantities (e.g., \mstar\ and synthetic images for an unlimited number of broad- and narrowbands) suggests a broad range of applications. Examples follow. \rem{i)} Subtraction of the SF component from BCDs and other starburst galaxies, allowing for improved determinations of the central intensity distribution (and, henceforth, the gravitational potential) of the underlying host galaxy. An unresolved question in this context is whether the latter shows an extended flat core, which under certain assumptions would imply a central minimum in the radial stellar density distribution \citep[][see also Noeske et al. 2003]{P96a}. \rem{ii)} Spatial progression of SF activity in galaxies: This subject encompasses several particular aspects, ranging from the hypothesis of unidirectional SF propagation in cometary galaxies \citep{P98,P08}, to the formation history of multiple generations of young stellar clusters (SCs) in isolated and interacting starburst galaxies \citep[e.g.,][]{Ostlin03,Adamo11,Whitmore07}, to the outward propagation of a SF front in collisional ring galaxies \citep{ASM96,Romano2008}. Adaptive removal of SCs and/or a more diffusely distributed young stellar substrate for a set of increasing \tcut\ can, in principle, provide a powerful technique for the reconstruction of SF propagation patterns and their role in galaxy build-up. \rem{iii)} Tidal dwarf galaxy (TDG) formation: Are these entities forming through gas collapse within a pre-existing gravitational potential from tidally ejected stars or out of a purely gaseous self-gravitating component? \citep[e.g.,][see also Duc 2012 for a review]{Weilbacher02}. Subtraction of the SF component with \ry\ could add decisive constraints for discriminating between both scenarios. \rem{iv)} Galaxy evolution in clusters: Galaxies falling onto clusters may experience a complex SFH involving, for example, an initial starburst episode followed by ram-pressure induced cessation of SF \citep[e.g.,][]{Poggianti99} in some cases accompanied by kpc-long, UV-emitting SF tails \citep[e.g.,][]{Hester2010,Kenney14}. \ry\ offers a handy tool to explore galaxy evolution in clusters back to several $10^8$ yr, i.e., over the critical phase of ram-pressure stripping. \rem{v)} Last but not least, a natural application of \ry\ concerns robust determinations of \rem{QMI} sets, such as CAS \citep[concentration-asymmetry-smoothness;][]{Conselice2003} and the Gini coefficient \citep{Lotz2004}, after decontamination of IFS data cubes from SF or directly from \sstar\ maps. The above examples outline the potential and wide range of applications of \ry\ toward deciphering the galaxy assembly history in the modern era of integral field spectroscopy.

1607.02884

1607

1607.05811.txt

Remote investigations of the ancient solar system matter has been traditionally carried out through the observations of long-period (LP) comets that are less affected by solar irradiation than the short-period counterparts orbiting much closer to the Sun. Here we summarize the results of our decade-long survey of the distant activity of LP comets. We found that the most important separation in the dataset is based on the dynamical nature of the objects. Dynamically new comets are characterized by a higher level of activity on average: the most active new comets in our sample can be characterized by \afrho{} values $>$3--4 higher than that of our most active returning comets. New comets develop more symmetric comae, suggesting a generally isotropic outflow. Contrary to this, the coma of recurrent comets can be less symmetrical, ocassionally exhibiting negative slope parameters, suggesting sudden variations in matter production. The morphological appearance of the observed comets is rather diverse. A surprisingly large fraction of the comets have long, teniouos tails, but the presence of impressive tails does not show a clear correlation with the brightness of the comets.

The origin and behaviour of comets are related to the entire Solar System, its general history and environment via several aspects. It is widely accepted that in the early Solar System a large number of comet-like bodies orbited in the Trans-Neptunian region and beyond, through the Oort-cloud. The group of Trans-Neptunian Objects (TNOs) was recognized as sources for short-period comets and probably various groups of asteroids (e.g. Duncan et al. 2004, Eicher 2013a). The recent exploration of the TNOs by the Herschel space observatory revealed the size and albedos of a handful objects (Lacerda et al. 2014), supporting that comet-like bodies are still present among the TNOs including other objects of different nature (e.g. Fornasier et al. 2013, Duffard et al. 2014). In recent years, the presence of similar comet clouds was suggested in several extrasolar systems, characterized by a prominent infrared excess due to cold debris (e.g. Beichman et al. 2005, Greaves and Wyatt, 2010). These observations show that comets are a common by-products of solar system formation, and they preserve the matter from the outskirts of young solar systems for a long time. Even, the relation of comet dust to ISM relics was suggested in the case of Hale--Bopp (Wooden et al. 2000). %Mg-Rich Silicate Crystals in Comet Haleâ??Bopp: ISM Relics or Solar Nebula Condensates? % The observations of cometary activity have been the traditional way of remote investigations. Since the long-period (LP) comets suffered the least solar irradiation, they can be the ideal targets for studying the ancient matter in a fairly intact state. Thank to the automated telescopes (e.g. Spacewatch, LINEAR, LONEOS, CSS/MLS/SSS, Pan-STARRS), relatively bright, large perihelion-distance comets are now regularly discovered often several years before the perihelion passage. Hence the study of distant cometary activity has become possible more regularly than before. Observations at heliocentric distances revealed that comets can be significantly active well beyond the snow line. The sublimation of water ice is excluded in this region because it can only be efficient at a few AU from the Sun, inside indicatively Jupiterâ??s orbit at 5.2 AU (Meech \& Svoren 2004). The main mechanisms that have been proposed to explain cometary activity of comets at large heliocentric distances are: the transition phase between amorphous and crystalline water ice (Prialnik 1992, Capria et al. 2002), the annealing of amorphous water ice (Meech et al. 2009), and the sublimation of more volatile admixtures like CO and/or CO2. In the past two decades, observational campaigns have been made to reveal the distant activity of comets, most by with observations of the Jupiter family (Meech 1991, % O'Ceallaigh et al. 1995; Meech \& Hainaut 1997; % Lowry et al. 1999; % Szab\'o et al. 2001; % Szab\'o et al. 2002; % Korsun \& Ch{\"o}rny 2003; % Tozzi et al. 2003; % Szab\'o et al. 2008; % Mazzotta Epifani et al. 2009; % Meech et al. 2009; % Mazzotta Epifani et al. 2010; % Korsun et al. 2010; % Szab\'o et al. 2012; % Mazzotta Epifani et al. 2014; % Shi et al. 2014; % Ivanova et al. 2015). % However, there is still a lack in the observations of LP comets beyond 5.2 AU. Most importantly, there is a natural deficit of LP comets observed on the inward orbit, because distant comets used to have been discovered near the perihelion (roughly before 2000). For example, as of writing this paper, there are only four comets with well documented observations covering at least 1 AU on the inward orbit beyond 5.2 AU (C/1995 O1 (Fulle at al. 1998), C/2003 WT42 (Korsun et al. 2010), C/2006 S3 (Rousselot et al. 2014), and C/2012 S1 (Kri\v{s}andov\'a et al. 2014). The physical mechanism that drives cometary activity is quite different at these large heliocentric distances. The sublimation of water ice is excluded in this region because this process can only be efficient at a few AU from the Sun, inside Jupiterâ??s orbit (Meech \& Svoren 2004). The main mechanisms that have been proposed to explain cometary activity of comets at large heliocentric distances are: the transition phase between amorphous and crystalline water ice (Prialnik 1992, Capria et al. 2002), the annealing of amorphous water ice (Meech et al. 2009), and the sublimation of more volatile admixtures like CO and/or CO2. According to the recognition by Oort, a dynamically new comet is usually defined as a comet on $a>10,000$ AU orbit, or $P>1$~million yr (Eicher, 2013b, Mazzotta Epifani 2014). These comets represent the early stage of the inward migration (this is why they are called as ``new'' comets). They reside very far from the Sun, and spend the waste majority ($>$99.9\%{}) of their lifetime even outside the heliosphere. Therefore, these comets are exposed to marginal solar irradiation, and only very little modifications by the solar wind. One can expect that the behavior of the dynamically new comets differs from that of the other comets (called as ``returning comets''), which investigation is in the focus of the present paper. As a result of our decade-long survey about the distant activity of LP comets, we gained 150 images of 50 comets showing activity beyond the snow line. The observations are still on-going with the same instruments, but at the present stage, we can already answer some important questions related to the activity of these comets. The observations are presented in the context of the following questions: \begin{enumerate} \item{} What is the behaviour of the long-period comets at large heliocentric distances? How does the activity evolve and cease? \item{} Are the activity profiles similar to each other during the inward and outward orbit? \item{} What kind of specific correlations can be recognized between the activity parameters ($Af\rho$, slope, tail characteristics), and between these parameters and the ephemerides? \item{} Can distinct groups be recognized by the activity parameters? \item{} What can we deduce from short time-scale variations such as outbursts or rapid evolution of matter production? \end{enumerate} The paper is structured as follows. The observations and reduction steps are described in Sect.\ 2, while Sect.\ 3 deals with the detailed observational results. The discussion of the results is given in Sect.\ 4.

\begin{figure*} \includegraphics[angle=270,width=8.0cm]{a6.eps} %\caption{} %\end{figure} %\begin{figure} \hfill\includegraphics[angle=270,width=8.0cm]{a11.eps} \caption{Left panel: Dependence of the slope on the heliocentric radius. Right panel: the same as a function of \afrho. The color coding is the same as in Fig. 3.} \end{figure*} The observation program we report here looks back to more than 15 years. There are a few comets which have been observed both before and after the perihelion with long coverage (most importantly C/2008 S3, C/2002 VQ94, C/2006 S3). The systematic and comparative study of the activity of LP comets beyond 5 AU requires observations covering at least 15-25 years, most importantly because requires systematic observations of comets which are discovered well before the perihelion. The work we present here can be considered as a basis for such a survey, since it already contains C/2008 S3 with a long inward$+$outward observation history (and a few comets with few points inward and outwards, e.g. C/2002 VQ94, C/2006 S3). %Although the observation program we report here looks back to more than 15 years, there are still few comets which have been %observed both before and after the perihelion with long coverage. The main conclusion is that a systematic and comparative study %of the activity of LP comets beyond 5 AU requires observations covering at least 20-25 years, most importantly because requires %systematic observations of comets which are discovered well before the perihelion. About ten comets in the presented sample are worth further follow-up lasting for many years, or even decades. These comets are typically around the perihelion, or somewhat after it now, have a well covered inward history in our data, and following the outward activity is highly desirable (C/2006~S3, C/2008~S3, C/2010~S1, C/2012~K8, C/2012~LP26, C/2012~U1, C/2013~G9, C/2013~P3). The most promising objects for further observations are those ones which are still before perihelion. C/2010~U3, being a unique comet for its early discovery, will reach the perihelion only in 2019. C/2013~X1 will have a very close perihelion in 2016, approaching the Sun to 1.3 AU. C/2011~KP36 and C/2010~U3 will reach their perihelion also in 2016. The brightest targets (C/2005~L3, C/2006~S3, C/2010~S1, C/2010~U3), where the size of the nucleus can be expected between those of comets Halley (Hainaut et al. 2004) and Hale--Bopp (Szab\'o et al. 2012), offer a good detectability of the naked nucleus after the cessation of the activity. A visual inspection suggests that the appearance of the observed comets is rather diverse. With \afrho~values spanning over 2 orders of magnitude (indicatively 100--10,000 cm), there are comets which are compact or diffuse -- apparently mostly regardless to the brightness or the \afrho. A surprisingly large fraction of the comets have tails, which are quite impressive in some cases (see Fig. 1 for the best examples), and while these tails are rather tenious, they were observed around fainter comets as well, and they were not necessarily associated with the brightest comets. These visual findings can be quantitatively analysed by means of \afrho and slope parameters. We will evaluate the results from the three most intriguing aspects in the followings. \subsection{The \afrho{}--$R$ activity profiles} The evolution of \afrho with $R$ is plotted in Fig. 2, left and right panels. The left panel shows the pre-perihelion activity, while the right panel shows the same, but for the post-perihelion orbit. The heliocentric distance in the left panel increases leftwards, therefore we can imagine that the Sun is at the center, between the two panels. The symbol and color coding discriminate the recurrent comets (green boxes) and Oort-cloud comets (open red circles). The general $\log$\afrho--$\log R$ profile is roughly linear for the most comets (Fig. 2), implying an underlying power law, which we can confirm for a large set of LP comets. Interestingly, the rate of increasing of the activity seems to depend on the onset level of the baseline activity level. Comets which exhibit high \afrho values far from the Sun tend to follow a shallower activity history, while other comets which are relatively fainter/less active at large solar distances can evolve more spectacularly, following a steeper power index. This observation is also compatible with the remarkable comet disappointments from the previous decades, when impressive comets, discovered at large solar distances with extraordinary activity level tend to slowly evolve and significantly underachieved the peak brightness predicted from early data (C/1973~E1 (Kohoutek), C/1989~X1 (Austin), C/2001~Q4 (NEAT)). Our dataset suggests that this is a usual pattern of LP comets which exhibit distant pre-perihelion activity of an extraordinary level. A similar behaviour is suggested for the post-perihelion phase, but may be with a less expresed difference between more and less active comets, while the more active ones seem to exhaust slowlier. The slow fading is a well known observation also for periodic comets, and is usually explained by the heat inertia of the nucleus. \subsection{Comparison of dynamically new and recurrent comets} Based on their database of comet observations ($\sim$50 comets over a range of 1 to $\sim$30 AU), Meech \& Hainaut (1997) have shown that dynamically young comets are intrinsically brighter, exhibiting dust comae and activity at large distances in the region where water ice sublimation is not possible. The observations by date support this tendency in general, and our results also support this conclusion. In the two panels of Figures 3 and 4, we plot the dependence of \afrho on various orbital and morphological parameters (slope, heliocentric distance, semi-major axis). Note that there are more comets in these figures as in Fig. 2, since here we plot those comets which are represented by only one or two observations in the current dataset. The cloud of points representing dynamically new comets (Fig 3, right panel, red points) lies convincingly above that of recurrent comets (green dots in the same figure panel), regardless to the solar radius. The difference seems to be most prominent between 5--7 AU, where the dynamically new comets tend to advance up to a factor of $\approx 4$ higher limit in \afrho, in comparison to recurrent comets. This observation is compatible to the presence of a huge amount of volatiles on the surface of new comets, in respect to returning comets. The slope parameter also shows an interesting difference between returning and new comets (Fig. 4). Both groups tends to exhibit slightly negative slope, typically between 0 and $-1$, and also, both groups have comets which were observed at a non-steady state, with a slightly positive slope value. The difference between the two groups is that the observations of negative slope usually represents the returning comets in the sample. Several features can lead to positive slopes, such as fan-like or spinning features in the coma, and the sudden and usualy temporal decrease of the activity when the outer coma is relatively richer in dust, in respect to the inner coma. The increased occurrence of positive slope in the returning group reflects that such uncommon morphological features and/or puffing activity is more common in the returning group, while activity of the new comets is usualy more regular. It is hard to say wether there were sudden changes in the activity level or morphological parameters during the observed period of the target comets, simply because there is not enough dense coverage for continuous monitoring of such effects. However, several comets have observations separated at 1-2-3 days apart (see Table 2), and during these observations, the observed parameters did not change exceeding the observational errors. Therefore, we have no evidence of sudden events in the observed data series. In essence, the most important groupping of the dataset is related to the recurrent/new nature of the comet. Dynamically new comets are characterized by a higher level of activity in average, with more regular and smoothly evolving matter production. The more symmetrical appearance suggests also a more isotropic outflow, too. %%%%%%%%%%%%%%%%%5 %The negative slope reflects an interaction between overstreaming matter and the radiation pressure modifies this value down to about -2, according to the models by Jewitt \& Meech (1987). %%%%%%%%%%%%%%%%%%%%%%%%%% Our results can be summarized as follow. \begin{enumerate} \item We present 152 observations of 50 comets from 103 nights, showing activity at solar distances between 5--15 AU. \afrho{}, \afrhomax{} and slope parameters were determined for all observations. \item{} All comets showed a coma and in many cases a tail, which sometimes exceeded 10$^\prime$. \item{} The \afrho{} value of the most active recurrent comet was surpassed by that of 5 new comets. The average value of \afrho{} of recurrent comets significantly exceeded that of the recurrent comets. We explained this observation with an increased amounts of volatiles on the surface of the new comets, demonstrating that dynamically new comets are also new in composition. \item{} All new comets showed a negative slope parameter and usualy a quite symmetric coma. The observed positive slope parameters, that we interpreted as signs of activity variations, or jets, arcs and other internal structures, were observed only in recurrent comets. \item{} The evolution of \afrho{} roughly followed a power law of $R$ for the most comets. Those comets which were characterized by larger \afrho{} values at larger solar distances tended to follow a flatter power. This suggested that comets which had very significant activity around 10--15 AU tended to less increase in activity with decreasing solar distance, thus representing the ''comet disappointment'' scenario. \item{} The analogous ``inverse'' comet disappointment behavior was observed post-perihelion, but with a less expressed difference between the most and least active comets. \item{} We examined the dependence of the studied activity parameters on orbital elements, the actual ephemerides at the observations and the visibility of a tail. We did not find significant subgroups, the only significant grupping in the sample was the new/recurrent type of the comet. \item{} Since the fully covered observation of a comet, extending beyond 10--15 AU solar distances on both half orbits lasts decades, we proposed comet observations with very long time coverage of 20--25 years, to be able to compare the inward and outward behavior of the same comets in an extensive sample. \end{enumerate}

1607.05811

1607

1607.00309_arXiv.txt

{ The Moscow State University Extensive Air Shower (EAS-MSU) array studied high-energy cosmic rays with primary energies $\sim (1 - 500)$~PeV in the Northern hemisphere. The EAS-MSU data are being revisited following recently found indications to an excess of muonless showers, which may be interpreted as the first observation of cosmic gamma rays at $\sim 100$~PeV. In this paper, we present a complete Monte-Carlo model of the surface detector which results in a good agreement between data and simulations. The model allows us to study the performance of the detector and will be used \red to obtain physical results in further studies. \black}

\label{sec:intro} The EAS-MSU array \cite{EAS-MSU} was established in late 1950s and has been upgraded in early 1980s. The array aimed at investigations of extensive air showers (EAS) produced by primary particles in the energy range $(10^{15}-5\times10^{17})$~eV. The array operated until 1990 and its main results have been published, notably the discovery of the knee in the cosmic ray spectrum \cite{knee} by the early version of the installation, results on the primary \red spectrum \cite{EAS-MSU-spec} and \black chemical composition in the knee energy region \cite{EAS-MSU-results, EAS-MSU-results1}. A unique feature of the array was the presence of large-area underground muon detectors sensitive to muons with energies $\gtrsim 10$~GeV. Quite recently, an analysis of the data of these detectors, \red whose area and energy threshold have \black no analogs in modern installations, revealed an excess of muonless events which may be interpreted as an evidence for showers initiated by primary photons \cite{EAS-MSU-gamma1, EAS-MSU-gamma2, EAS-MSU-gamma3}. If confirmed, this observation would mean the discovery of first cosmic gamma rays above 100~TeV. Therefore, it calls for a careful reanalysis. Indeed, muonless or muon-poor events may appear as rare fluctuations of usual, hadron-induced air showers. The estimate of this background represents a crucial ingredient in a reliable study of photon candidate events. Previous studies \cite{EAS-MSU-gamma1, EAS-MSU-gamma2, EAS-MSU-gamma3} used a simplified model of the detector which, in principle, might underestimate rare fluctuations in the muon content of hadronic showers. Here, we develop a modern Monte-Carlo description of the installation. It is based on the air-shower simulation with CORSIKA \cite{CORSIKA} supplemented by a model of the detector. As is customary in modern EAS experiments, see e.g.\ Ref.~\cite{Ben}, the result of a simulation run is recorded in precisely the same format as the real data. This record is further processed by the usual reconstruction routine, so that real and simulated events are processed with the same analysis code. A simulation of this kind was not technically possible at the time the bulk of the installation's results were obtained. The aim of this paper is to present the simulation, to estimate the performance of the installation and to demonstrate a good agreement between real and simulated data in terms of basic reconstruction parameters, which opens the way to use the Monte-Carlo description in further studies of the EAS-MSU data. Results of these physics studies will be reported in further publications. The rest of the paper is organized as follows. In Sec.~\ref{sec:array}, the EAS-MSU array is described in detail. Section~\ref{sec:reconstruction} discusses the event reconstruction procedure, including quality cuts used in the analysis of both real data and simulated events. In Sec.~\ref{sec:MC}, we turn to the Monte-Carlo procedure and describe both the air-shower simulation and modelling of the detector. Section~\ref{sec:comparison} presents a comparison between simulated and real events, as well as between thrown and reconstructed parameters. We briefly conclude in Sec.~\ref{sec:concl}.

\label{sec:concl} To summarize, we presented a full chain of Monte-Carlo simulations of air showers developing in the atmosphere, being detected by the EAS-MSU array and reconstructed by the procedure equivalent to that used for the real data. We have verified that the simulation describes well the distribution of the data in basic parameters. Assuming a modern hadronic-interaction model, we obtained a relation between the reconstructed shower size $N_{e}$ and the primary energy $E$. We used our Monte-Carlo event sets to estimate the accuracy of the reconstruction of the shower geometry, $N_{e}$ and energy. Forthcoming studies will use these simulations to reanalyze data on $E_{\mu}>10$~GeV muons in the air showers recorded by the installation, which will open the possibility to test, within modern frameworks, the origin of the apparent excess of muonless events seen in the EAS-MSU data.

1607.00309

1607

1607.03000_arXiv.txt

Stellar evolution theory has been extraordinarily successful at explaining the different phases under which stars form, evolve and die. While the strongest constraints have traditionally come from binary stars, the advent of asteroseismology is bringing unique measures in well-characterised stars. For stellar populations in general, however, only photometric measures are usually available, and the comparison with the predictions of stellar evolution theory have mostly been qualitative. For instance, the geometrical shapes of isochrones have been used to infer ages of coeval populations, but without any proper statistical basis. In this chapter we provide a pedagogical review on a Bayesian formalism to make quantitative inferences on the properties of single, binary and small ensembles of stars, including unresolved populations. As an example, we show how stellar evolution theory can be used in a rigorous way as a prior information to measure the ages of stars between the ZAMS and the Helium flash, and their uncertainties, using photometric data only.

\label{dvg:sec1} In this chapter a brief summary is presented of the uses of stellar evolution theory to infer properties of single stars (Section~\ref{dvg:single}), of detached binary stars whose components are assumed to have evolved independently of each other (Section~\ref{dvg:double}), and coeval stellar populations such as (presumably) those in clusters (Section~\ref{dvg:coeval}). When the fundamental properties of the star(s) in question are known (mass, absolute luminosity, effective temperature, etc), stellar tracks computed for this particular (set of) star(s) can be used to infer further properties, such as ages. In general, however, one wishes to use the predictions of stellar evolution to infer these properties. The data at hand are usually magnitudes and colours, hence the interpretation of the features in the colour-magnitude diagrammes (CMDs) is carried out with isochrones rather than stellar tracks. The thorny issue of transforming isochrones from/to the theoretical diagram to/from observed CMDs will not be dealt with here (see the contributions by Cassisi, and by Lebreton, Goupil and Montalb\'an in this volume), and constitute one of the sources of systematic uncertainties. One also has to bear in mind that while many efforts have been placed to find the best transformations, the differences observed cannot (yet?) be fully ascribed to either systematics in the observations or in missing/wrong physics in the stellar evolutionary calculations. For instance the recent analysis by \cite{vandenberg2010} for some globular clusters, and by \cite{an2007} for open clusters show that while some CMDs can be well fitted, other colour-magnitude combinations of the \textsl{same} clusters show anomalies which go well beyond the corrections for systematics, rotation, activity, transformations, metallicity, etc. Empirical bolometric corrections are another source of uncertainty \citep{torres2010} as are the systematics in the determination of effective temperatures \citep[\eg,][]{ramirezmelendez2005}. Likewise, it does make sense to adopt a standard set of values \citep{harmanec2011} with nominal values to avoid some of the systematics arising with the adoption of different key values (solar radius, mass, etc). Section~\ref{dvg:composite} deals with the general problem of inverting the CMDs of a mixture of resolved stellar populations to infer their distribution of ages and hence their chemical and star formation rate histories, while Section~\ref{dvg:pixels} is a very brief discussion on CMDs of pixels in unresolved stellar populations. Section~\ref{dvg:stats} closes this chapter with a discussion on some statistical issues in the interpretation of CMDs. We will limit the scope of this chapter to stars from the main sequence to the Helium flash, as the predictions in this range appear to be the most robust ones. The white dwarf phase can also be used in a rather robust way (provided the cooling and the equation of state are properly characterised), as described by T. von Hippel in this volume.

1607.03000

1607

1607.01005_arXiv.txt

These notes show and comment the examples that have been used to validate the CosmicFish code. We compare the results obtained with the code to several other results available in literature finding an overall good level of agreement. We will update this set of notes when relevant modifications to the CosmicFish code will be released or other validation examples are worked out. \\ The CosmicFish code and the package to produce all the validation results presented here are publicly available at~\url{http://cosmicfish.github.io}. \\ The present version is based on CosmicFish Jun16.

In~\cite{Raveri:2016xof, Raveri:2016leq} we introduced the CosmicFish code as a powerful tool to perform forecast on many different models with future cosmological experiments. \\ In this set of notes we show the validation pipeline that was used for the code. We compared the results obtained with the CosmicFish code to other results in literature. We find an overall good level of agreement. \\ Together with these notes we release a CosmicFish package that contains the relevant code to produce all the results presented here. This package is going to be updated as new validation results become available. The CosmicFish code and its validation package are publicly available at~\url{http://cosmicfish.github.io}.

1607.01005

1607

1607.04285_arXiv.txt

We attempt to interpret existing data on the evolution of the UV luminosity function and UV colours, $\beta$, of galaxies at $5 \leq z \leq 8$, to improve our understanding of their dust content and ISM properties. To this aim, we post-process the results of a cosmological hydrodynamical simulation with a chemical evolution model, which includes dust formation by supernovae and intermediate mass stars, dust destruction in supernova shocks, and grain growth by accretion of gas-phase elements in dense gas. We find that observations require a steep, Small Magellanic Cloud-like extinction curve and a clumpy dust distribution, where stellar populations younger than 15 Myr are still embedded in their dusty natal clouds. Investigating the scatter in the colour distribution and stellar mass, we find that the observed trends can be explained by the presence of two populations: younger, less massive galaxies where dust enrichment is mainly due to stellar sources, and massive, more chemically evolved ones, where efficient grain growth provides the dominant contribution to the total dust mass. Computing the IR-excess - UV color relation we find that all but the dustiest model galaxies follow a relation shallower than the Meurer et al. (1999) one, usually adopted to correct the observed UV luminosities of high-$z$ galaxies for the effects of dust extinction. As a result, their total star formation rates might have been over-estimated. Our study illustrates the importance to incorporate a proper treatment of dust in simulations of high-$z$ galaxies, and that massive, dusty, UV-faint galaxies might have already appeared at $z \lesssim 7$.

In the last decade, data from the {\it Hubble Space Telescope}\footnote{http://www.stsci.edu} (HST), especially after the advent of the Wide Field Camera (WFC3), allowed us to collect large samples of galaxies at $z \sim 7 - 8$, with smaller samples extending up to $z \sim 9 -11$ \citep{McLure2013, Bouwens2014, Oesch2014, Bouwens2015a, McLeod2015, Finkelstein2015}, among which the two most distant spectroscopically confirmed galaxies at $z = 8.68$ \citep{Zitrin2015} and $z = 11.1$ \citep{Oesch2016}. Since spectroscopic observations of galaxies at $z>6$ with ground-based telescopes are very challenging, observers have developed alternative, photometry-based techniques to both select high-$z$ candidates and estimate their physical properties. For example, the total stellar mass, the stellar age and the ongoing star formation rate (SFR) can be estimated from spectral energy distribution (SED) fitting and colour index analyses. Two key quantities are generally used to characterize the properties of the first galaxies and of their dominant stellar populations: the UV luminosity function (LF) and the observed UV spectral slope, $\beta$ ($f_\lambda \propto \lambda^\beta$, \citealt{meurer1999}). The LF, defined as the number density of galaxies per unit magnitude, provides important constraints on star formation efficiencies at different redshifts and on their evolutionary status, especially at early times \citep{Bouwens2011}. As the first structures collapse and assemble their stellar content, the inter-stellar medium (ISM) is progressively enriched with metals and dust. Dust extinction affects the UV luminosity and should leave a signature in the galaxy LF. While for low-redshift galaxies dust extinction can be corrected by measuring the far infrared emission (FIR), observations of high-$z$ galaxies with millimeter (mm) telescopes, such as the {\it Atacama Large Millimeter Array} (ALMA) and the {\it Plateau de Bure Interferometer} (PdBI), have mostly provided upper limits on the rest-frame FIR emission of $z > 6$ UV-selected galaxies \citep{Kanekar2013, Ouchi2013, Ota2014, Schaerer2015, Maiolino2015, Zavala2015}, with one notable exception \citep{Watson2015, Knudsen2016}. For this reason, it has become a common practice to estimate the effects of dust extinction using the observed $\beta$ slope, or UV colour. Despite a vigorous debate in the past ten years, recent observational results appear to converge on a common trend for the shape and the evolution of the UV LF in the redshift range $4 < z < 8$ \citep{Bouwens2015a}, down to an AB magnitude of $-16$ at $z = 4, 5$ and of $\sim -17$ at $z = 6, 7$ and $8$. At even higher redshifts, even the deepest observations in blank fields can only probe the bright-end of the LF, providing important constraints on the volume density of the most luminous galaxies with $M_{\rm UV}< -20$ at $z = 9$ and $10$ \citep{Bouwens2015a}. An efficient way to push the observations to fainter luminosities is to exploit the gravitational lensing magnification of massive galaxy clusters. Results from the HST programs CLASH and Hubble Frontier Fields (HFF) have increased the statistics of candidate galaxies at the highest redshifts, providing better constraints on the evolution of the faint-end slope of the LF and placing the first limits on the LF at $z \sim 10$ \citep{Atek2015, McLeod2015, McLeod2016, Livermore2016} . It is custumary to fit the LF with a Schechter function, which has a power-law behaviour with slope $\alpha$ at the faint-end, an exponential cut-off brighter than a characteristic luminosity (magnitude) $L_\ast$ ($M_\ast$) and a volume density of $\phi_\ast$ at this characteristic luminosity, \begin{equation} \frac{dn}{dL} = \phi(L) = \left(\frac{\phi_\ast}{L_\ast}\right)\left(\frac{L}{L_\ast}\right)^\alpha e^{-L/L_\ast}. \end{equation} In general, the evolution of the LF with redshift is characterized by means of variations of these Schechter parameters and is consistent with a steady growth in the volume density and luminosity of galaxies with time. In particular, there is a significant evidence for a steepening of the faint-end slope with $z$, in agreement with the predicted steepening of the halo mass function, a modest evolution of $M_\ast$ and a decrease of $\phi_\ast$ from $z \sim 4$ to $z \sim 7$ \citep{Bouwens2015a}. Some observations at $z = 9$ and $10$ suggest a faster evolution, and that the luminosity densities inferred from current samples are $\sim 2$ times lower than the values extrapolated from the trends at $4 < z < 8$ \citep{Bouwens2015a}. Other studies support a smoother evolution from $z = 8$ to 9 \citep{McLeod2015, McLeod2016, Finkelstein2015}. Indeed, the recent discovery of GN-z11, a luminous galaxy with $\rm M_{UV} = - 22.1$ at $z = 11$ \citep{Oesch2016} may indicate that the LF at the very bright end does not follow a Schechter functional form, possibly due to less efficient feedback at very high redshifts \citep{Bowler2014, Dayal2014,Finkelstein2015b,Waters2016}. To convert the observed UV luminosity to a SFR and compare the above findings to theoretical predictions, dust correction is usually estimated using the observed $\beta$ slopes and the so-called IRX-$\beta$ relationship by \citet{meurer1999}, who proved that, at $z <3$, the amount of SED reddening directly correlates with the $\beta$ value, as also confirmed by independent theoretical predictions (e.g. \citealt{Wilkins2012b}). Although this relation has been calibrated on starburst galaxies at low redshifts, and assumes a constant mean intrinsic slope of $\beta = -2.23$, this procedure has been widely adopted in high-$z$ galaxy surveys \citep[see][]{Bouwens2012}. However, the value of $\beta$ is also a function of important properties of the stellar populations, such as their ages, metallicity and initial mass function (IMF). Although with large uncertainties, observational trends have been reported which quantify the dependence of $\beta$ on the UV luminosity and redshift \citep{Stanway2005, Wilkins2011, Finkelstein2012, Bouwens2012, Castellano2012, Dunlop2013}. In general, the observations are consistent with a decreasing reddening towards lower luminosities and higher redshift. A coherent analysis of the observed $\beta$ for galaxies in a wide redshift range, from $z \sim 4$ to $z \sim 7$, has been recently made by \citet{Bouwens2014}, who confirm a strong evidence for a dependence of the average $\beta$ on the UV luminosity, the so-called Colour-Magnitude-Relation (CMR, \citealt{Rogers2014}), with brighter galaxies being redder and fainter galaxies being bluer, and a flattening of the relation at luminosities faintward of $M_{\rm UV} \sim -19$. They also report a small but clear evolution with time, with galaxies at fixed luminosity becoming bluer with $z$. For the faint galaxies with $-19 < M_{\rm UV} < -17$, the mean $\beta$ at $z \sim 4, 5$ and $6$ is $-2.03$, $-2.14$ and $-2.24$ respectively. Extrapolation of this trend to $z \sim 7$ and $8$ suggests mean values of $-2.35$ and $-2.45$, consistent - within the errors - with the observed ones. Theoretical studies have attempted to interpret the data by means of numerical simulations or semi-analytical models. \citet{Wilkins2012b} explored the sensitivity of the intrinsic $\beta$ slopes to the IMF and to the recent star formation and metal enrichment histories of low-$z$ galaxies. They find a distribution of $\beta$ values with a scatter of $0.3$, which introduces an uncertainty in the inferred dust attenuation when a constant intrinsic slope is assumed. This scatter is significantly reduced for galaxies at $z \sim 6$, but the mean intrinsic $\beta$ decreases with $z$. If this is not properly taken into account and the locally calibrated relation is applied, dust attenuation is systematically underestimated \citep{Wilkins2013}. \citet{GonzalezPerez2013} have demonstrated the dependence of the galaxy UV colours on the adopted dust properties and, in particular, on the dust extinction curve. With the aim of intepreting high-$z$ Lyman-$\alpha$ emitters (LAEs) and Lyman Break Galaxies (LBGs) observations, \citet{Dayal2010a} and \citet{Dayal2012} used a numerical simulation to derive intrinsic galaxy properties and a semi-analytical model to estimate dust attenuation. They explored the resulting UV LF and the dependence of $\beta$ on the galaxy UV luminosity with and without dust attenuation. They found that dust attenuation improves the agreement with the observations, but the observed CMR was not reproduced by the model results. More recently, \citet{Khakhaleva2016} post-processed the results of cosmological simulations with a simple dust model that assumes a constant dust-to-metal mass ratio in the neutral gas and that dust is instantaneously sublimated in hot ionized regions. They used a Monte Carlo radiative transfer code to predict UV attenuation and IR re-emission of their model galaxies. By means of a detailed comparison with observations at $5 \le z \le 10$, they concluded that, in order to assess the effects of dust in the ISM of high-$z$ galaxies, the complex interplay of dust creation and destruction processes should be fully incorporated into numerical simulations. With this aim, in \citet{mancini2015} we have presented a semi-numerical model which includes a physically motivated description of dust evolution, accounting for dust enrichment by Supernovae (SNe) and Asymptotic Giant Branch (AGB) stars, the effects of dust destruction by SN shocks and grain growth in the dense cold phase of the ISM (see also \citealt{Valiante2009, Valiante2011, deBennassuti2014}). We then compared the model predictions with the limits on the dust mass inferred from mm-observations of $z > 6$ galaxies, deriving interesting constraints on the properties of their ISM and on the nature of dust at high-$z$. Here we extend this previous investigation with the goal of intepreting the observed UV luminosities and colours of galaxies at $5 < z < 8$. The paper is organized as follows. In Section \ref{sec:method} we describe our method and the assumptions made to compute the dust content and luminous properties of the simulated galaxies. In Section \ref{sec:results} we first discuss the predicted physical properties of the galaxies at $5 \le z \le 8$. Then we derive the UV LFs and $\beta$ slopes assuming no dust extinction, and discussing the dependence of the results on the extinction model. In Section \ref{sec:comparison} we introduce the model that better reproduce the observed UV LFs and CMR. We analyze the origin of the scatter around the mean values at different $z$, both in the CMR and in the stellar mass - UV luminosity relation. We compute the IR excess and dust attenuation factors, comparing with observationally inferred correlations. Finally, in Section~\ref{sec:conclusions} we summarize the results and draw the main conclusions.

\label{sec:conclusions} The main goal of the present study is to provide a consistent framework to interpret the observed evolution of the UV LFs and galaxy colours over the redshift range $5 \leq z \leq 8$. To this aim, we have used a semi-numerical approach to post-process the output of a cosmological simulation with a chemical evolution model with dust. Our approach allows us to follow dust enrichment by stellar sources (SNe and AGB stars), dust destruction in the diffuse gas by SN shocks, and grain growth in dense molecular clouds. The model has already been applied by \citet{mancini2015} to interpret current observational constraints on the dust mass inferred from ALMA and Plateau de Bure observations of normal star forming galaxies at $z > 6$. Here we extend the analysis to investigate how dust properties affect the UV LFs and galaxy colours at high-$z$. Our main findings can be summarized as follows: \begin{itemize} \item The comparison between model predictions and observations at $5 \lesssim z \lesssim 8$ shows that, while the ISM dust has a negligible effect on the galaxy UV LFs at $z \sim 7$ and $8$, it reduces the number of galaxies with $M_{\rm UV} \ge -18$ and $\ge -19$ at $z \sim 5$ and 6 to values in very good agreement with observations. The observed CMR and its dependence on $z$ suggest a steep extinction curve in the wavelength range $1500 \, \AA \le \lambda \le 3000 \, \AA$, and that stars with age $\leq 15$~Myr are embedded in their dense molecular natal clouds and their UV luminosity suffers a larger dust extinction. \item The scatter in the colour distribution around the mean CMR increases with luminosity and cosmic time, consistent with observations. At $z \lesssim 6$, galaxies with $-19 \leq M_{\rm UV} \leq -18$ (where we have adequate statistics, given our simulation volume and resolution) are a mix of intrinsically faint blue galaxies and of red objects which have suffered strong dust extinction. The latter population grows with time, as a result of more efficient grain-growth in their ISM, and lie off the mean $M_{\rm star} - M_{\rm UV}$ relation inferred from observations at $z \sim 5,$ and 6. This is supported by the recent evidence for a population of massive UV-faint objects that makes a non negligible contribution to the stellar mass function at $z \lesssim 5.5$ \citep{Grazian2015}. \item By analyizing the properties of the simulated galaxies in the IRX - $\beta$ plane, we find that young, less massive galaxies, where the ISM dust is mostly contributed by stellar sources, follow a relation which is much flatter than the commonly adopted \citet{meurer1999} relation, consistent with their steep extinction curve. Massive dusty galaxies, which have experienced efficient grain growth in their ISM, introduce a considerable scatter in the IRX at a given UV continuum slope, slowly shifting towards the \citet{meurer1999} relation at $z \lesssim 6$. \item At $z \sim 7$ and 8, dust attenuation factors are better estimated assuming a flatter IRX - $\beta$ relation, such as the one recently inferred by \citet{Talia2015} or predicted for the SMC curve \citep{Pettini1998}. At $z \lesssim 6$, it is very hard to infer the proper dust attenuation from the UV slope alone, as galaxies with $-2 < \beta < -1.5$ can have vastly different IRX. \end{itemize} Our analysis suggests that {\it the total star formation rate density at high-$z$ may be overestimated if dust attenuation factors are derived using the \citet{meurer1999} relation, and that more realistic dust correction for young galaxies, which have not yet experienced major dust enrichment, can be derived from their UV colours using a flatter IRX - $\beta$ relation, such as the one inferred by \citet{Talia2015} or implied by the SMC curve. However, once grain growth starts to dominate dust enrichment, a population of massive, dusty, and UV faint galaxies appears at $z \lesssim 6$. These galaxies increase the scatter in the $\beta - M_{\rm UV}$, $M_{\rm star} - \beta$ and IRX - $\beta$ planes and slowly shift towards the \citet{meurer1999} relation. } We do not find dusty, massive, UV-faint galaxies at $z \sim 7$ and 8, but we can not exclude this to be an effect of the limited volume of our simulation. In fact, at the same redshifts we also underpredict the bright-end of the observed UV LFs, even assuming no dust extinction. Despite these limitations, our study shows that current high-$z$ observations on the evolution of galaxy colours already provide important constraints on the nature of dust and on its complex evolution and spatial distribution in the interstellar medium. The next mandatory step is to incorporate these processes directly into numerical simulations.

1607.04285

1607

1607.03146_arXiv.txt

The cosmological information contained in anisotropic galaxy clustering measurements can often be compressed into a small number of parameters whose posterior distribution is well described by a Gaussian. We present a general methodology to combine these estimates into a single set of consensus constraints that encode the total information of the individual measurements, taking into account the full covariance between the different methods. We illustrate this technique by applying it to combine the results obtained from different clustering analyses, including measurements of the signature of baryon acoustic oscillations (BAO) and redshift-space distortions (RSD), based on a set of mock catalogues of the final SDSS-III Baryon Oscillation Spectroscopic Survey (BOSS). Our results show that the region of the parameter space allowed by the consensus constraints is smaller than that of the individual methods, highlighting the importance of performing multiple analyses on galaxy surveys even when the measurements are highly correlated. This paper is part of a set that analyses the final galaxy clustering dataset from BOSS. The methodology presented here is used in \citet{Acacia2016} to produce the final cosmological constraints from BOSS.

\label{sec:intro} Over the past decades the size and quality of galaxy redshift surveys has increased dramatically. Thanks to these data sets, the information from the large-scale structure (LSS) of the Universe has played a central role in establishing the current cosmological paradigm, the $\Lambda$CDM model \citep[e.g.][]{Tegmark2004, Eisenstein2005, Cole2005, Anderson2012, Anderson2013, Anderson2014}. Several methods can be used to extract the information encoded in the large-scale distribution of galaxies. The power spectrum, $P(k)$, and its Fourier transform, the two-point correlation function $\xi(s)$, have been the preferred tools for LSS analyses. The anisotropies in these measurements caused by redshift-space distortions (RSD) and the Alcock--Paczynski effect \citep{Alcock1979} can be studied by means of their Legendre multipoles \citep[e.g.][]{Padmanabhan2008} or using the clustering wedges statistic \citep{Kazin2012}. Thanks to the combined information of baryon acoustic oscillations (BAO) and RSD, anisotropic clustering measurements can simultaneously constrain the expansion history of the Universe and the growth of density fluctuations, thus offering one of the most powerful cosmological probes. The potential of LSS observations as cosmological probes has led to the construction of increasingly larger galaxy catalogues. Examples of these new surveys include the completed Baryon Oscillation Spectroscopic Survey \citep[BOSS;][]{Dawson2013}, which is part of the Sloan Digital Sky Survey III \citep[SDSS-III;][]{Eisenstein2011}, the on-going SDSS-IV extended Baryon Oscillation Spectroscopic Survey \citep[eBOSS;][]{Dawson2016} and future surveys such as the Hobby Eberly Telescope Dark Energy Experiment \citep[HETDEX;][]{Hill:2008mv}, the Dark Energy Spectroscopic Instrument \citep[DESI;][]{Levi:2013gra} and the ESA space mission \emph{Euclid} \citep{Laureijs:2011gra}. As the construction of galaxy surveys requires a considerable amount of resources from the community, substantial efforts are put into maximizing the information extracted from the obtained data sets. This problem has often been posed as that of determining which statistic is the best to extract cosmological information (e.g. power spectrum vs. correlation function), often based on a simple metric or figure of merit. However, although the results obtained by applying different statistics to a given data set are highly correlated, as they are based on estimators and each measurement is analysed over a limited range of scales, they do not contain exactly the same information or are affected by noise in the same way. This means that, if the covariance between the different measurements is correctly taken into account, additional information could be obtained by combining the results inferred from different methods. In most cases, the cosmological information contained in the clustering measurements can be condensed into a small number of parameters whose posterior distribution is well described by a multivariate Gaussian. In this case, the obtained constraints can be represented by the mean values of these parameters and their respective covariance matrices. The analyses of the final BOSS galaxy samples of our companion papers are examples of this situation \citep{Beutler2016a, Beutler2016b, Grieb2016, Ross2016, Sanchez2016, Satpathy2016}. The BAO and RSD information obtained in these analyses can be expressed as constraints on the ratio of the comoving angular diameter distance to the sound horizon at the drag redshift, $D_{\rm M}(z)/r_{\rm d}$, the product of the Hubble parameter and the sound horizon, $H(z)\times r_{\rm d}$, and the growth-rate of cosmic structures, characterized by the combination $f\sigma_8(z)$, where $f(z)$ is the logarithmic growth rate and $\sigma_8(z)$ represents the linear rms mass fluctuation in spheres of radius $8\,h^{-1}{\rm Mpc}$. Here we present a general methodology to combine several Gaussian posterior distributions into a single set of consensus constraints representing their joint information, taking into account the full covariance between the different estimates. We illustrate this technique by applying it to the results inferred from the application of the same clustering analyses performed on the final BOSS galaxy samples to 996 {\sc Multidark-Patchy} ({\sc MD-Patchy}) mock galaxy catalogues reproducing the properties of the survey \citep{Kitaura2016}. The obtained consensus distributions represent a gain in constraining power with respect to the results of the individual methods, highlighting the importance of performing multiple analyses on galaxy surveys. The methodology presented here is used in our companion paper \citet{Acacia2016} to combine the cosmological information from the different analyses methods applied to the final BOSS galaxy samples \citep{Beutler2016a,Beutler2016b,Grieb2016,Ross2016,Sanchez2016,Satpathy2016} into a final set of consensus constraints. The structure of the paper is as follows, in Section \ref{sec:combination} we present the general scheme for the combination of different Gaussian posterior distributions into a set of consensus constraints that encode the full information provided by these estimates. We consider the cases in which the posterior distributions cover the same parameter spaces and when they differ. In Section~\ref{sec:app} we illustrate this procedure by applying it to the results obtained from different BAO and RSD measurements from a set of BOSS mock catalogues. Finally, Section~\ref{sec:conclusions} contains our main conclusions.

\label{sec:conclusions} We presented a general framework to combine the information of multiple Gaussian posterior distributions into a set of consensus constraints representing their joint information. This methodology can be applied to combine the cosmological information obtained from different clustering measurements based on the same galaxy sample, which can often be expressed as Gaussian constraints on a small number of parameters. The application of this technique requires the knowledge of the full cross-covariances of the different methods. For clustering measurements, this information can be obtained using a brute-force approach, applying the same methods being combined to a set of mock galaxy catalogues and measuring the correlations between the obtained results. \begin{figure} \includegraphics[width=0.45\textwidth]{figs/fig8.pdf} \caption{ Correlation matrix corresponding to the covariance of the full consensus constraints in our three redshift bins recovered from our BOSS mock catalogues. } \label{fig:cov_zbins} \end{figure} We illustrate our technique by applying it to combine the results obtained from different BAO-only and BAO+RSD measurements from an ensemble of mock catalogues of the final BOSS. The obtained consensus constraints represent a reduction in the allowed region of the parameter space with respect to the results of the individual methods. This shows the value of using the combination of the results of multiple clustering analyses as a strategy to maximise the constraining power of galaxy surveys. In our companion paper \citet{Acacia2016}, the methodology described here is used to obtain a set of consensus constraints that encode the results obtained by applying the same methods studied here to the final BOSS galaxy samples. These results are then used to explore the cosmological implications of the data set in combination with the information from cosmic microwave background and Type Ia supernovae data. We anticipate that the procedure detailed here can help to optimize the use of the cosmological information encoded in future clustering and lensing analyses.

1607.03146

1607

1607.01972_arXiv.txt

We analyse the heating of stellar discs by non axisymmetric structures and giant molecular clouds (GMCs) in $N$-body simulations of growing disc galaxies. The analysis resolves long-standing discrepancies between models and data by demonstrating the importance of distinguishing between measured age-velocity dispersion relations (AVRs) and the heating histories of the stars that make up the AVR. We fit both AVRs and heating histories with formulae $\propto t^\beta$ and determine the exponents $\beta_R$ and $\beta_z$ derived from in-plane and vertical AVRs and $\tilde{\beta}_R$ and $\tilde{\beta}_z$ from heating histories. Values of $\beta_z$ are in almost all simulations larger than values of $\tilde{\beta}_z$, wheras values of $\beta_R$ are similar to or mildly larger than values of $\tilde{\beta}_R$. Moreover, values of $\beta_z$ ($\tilde{\beta}_z$) are generally larger than values of $\beta_R$ ($\tilde{\beta}_R$). The dominant cause of these relations is the decline over the life of the disc in importance of GMCs as heating agents relative to spiral structure and the bar. We examine how age errors and biases in solar neighbourhood surveys influence the measured AVR: they tend to decrease $\beta$ values by smearing out ages and thus measured dispersions. We compare AVRs and velocity ellipsoid shapes $\sigma_z/\sigma_R$ from simulations to Solar-neighbourhood data. We conclude that for the expected disc mass and dark halo structure, combined GMC and spiral/bar heating can explain the AVR of the Galactic thin disc. Strong departures of the disc mass or the dark halo structure from expectation spoil fits to the data.

When the stars in the Solar neighbourhood (Snhd) are binned by age, the velocity dispersion of each bin increases with its age. This age-velocity dispersion relation (AVR) has been known and studied for decades (e.g. \citealp{stromberg,parenago, wielen}) and similar relations have now also been inferred for external galactic discs (\citealp{beasley} for M33, \citealp{dorman} for M31). It is generally agreed that understanding the physics that establishes the AVR would be a significant step towards understanding how galaxies form and evolve. Despite many efforts, the shape of the AVR is still not adequately constrained. The major constraints on the AVR come from observations of stars in the Snhd. The measured ages of stars suffer from substantial errors, and samples of stars with measured ages typically have a selection function that favours young stars over old \citep{nordstroem}. Alternatively modelling the velocity dispersions as functions of stellar colour has been used to determine the shape of the AVR \citep[hereafter AB09]{ab09}. Whereas the vertical density profile of the Milky Way (MW) is well fitted by a double-exponential in $|z|$, the AVR is typically described as a simple power-law in age $\sigma(\tau)\propto \tau^{\beta}$ with exponent $\beta$. \cite{quillen} claimed to detect a jump in the vertical AVR for ages $\tau> 10\gyr$ and connected this jump to the double-exponential nature of the density profile. However, more recent studies find that the AVR in the Snhd can be reasonably described by a single power-law (\citealp{holmberg}, AB09). Moreover, \citet{sb09} demonstrated that the observed vertical density profile is fitted well by a model in which the histories of star formation and disc heating are continuous, and \cite{bovy12} showed that subsets of stars selected at different points in the $([\alpha/\hbox{Fe}],[\hbox{Fe/H}])$ abundance plane have scale heights that vary continuously, with no evidence for a dichotomy. To constrain the heating processes responsible for the Snhd AVR, three diagnostics have been used in addition to the value of $\beta$: (i) the axis ratios $\sigma_z/\sigma_R$ of the velocity ellipsoids of old and young populations; (ii) the magnitude $\sigma_{{\rm old}}$ of $\sigma_z$ for the oldest populations; (iii) the vertical density profile of the MW's disc. Several physical processes work together to establish the AVR. The goals of this paper are to deepen our understanding of the collaboration of secular heating processes and to show how the AVR responds to one or the other process playing a more prominent role. External heating due to interaction with dark substructure is not considered in this paper. \citet{spitzer} showed that scattering of stars off small-scale irregularities in the potential of the Galactic disc would transform the nearly circular orbits of young stars, which reflect their origin from the gas disc, into orbits with higher radial and vertical energy. The discovery of Giant molecular clouds (GMCs) with masses $M_{\rm GMC}\sim 10^{5-7}\msun$ provided a suitable candidate for the heating agent. Using analytic arguments, \citet{lacey} showed that GMC heating could have contributed significantly to the observed AVR, but that the dispersion $\sigma_{{\rm old}}$ of the oldest stars could be reproduced only if the masses and / or number density of GMCs was significantly higher in the past than they are now. Lacey further derived values $\sigma_z/\sigma_R\sim 0.8$ and $\beta\sim 0.25$ that are, respectively, larger and smaller than the data indicate. \citet{ida} used analytical calculations more sophisticated than that of Lacey to show that scattering by GMCs in discs actually yields values $\sigma_z/\sigma_R\sim 0.5-0.6$ that are consistent with observations. This was confirmed with particle simulations by \citet{shiidsuka} and \citet{haenninen}: to obtain correct values it is essential to take into account that in a thin disc impact parameters are concentrated towards the galactic plane, whereas Lacey had assumed an isotropic distribution of impact parameters \citep{sellwood}. However, the conflict between observations and Lacey's values for $\beta$ and $\sigma_{{\rm old}}$ remained \citep{haenninen}. GMCs are not the only sources of a non-axisymmetric component to the gravitational field experienced by disc stars, and any such component will heat the disc. \citet{barbanis} showed that spiral structure could significantly heat the disc, when the spiral pattern has either a very high density contrast or is of a transient and recurring nature. Two-dimensional simulations of discs continuously fed with young, cold particles showed that transient and recurrent spiral structure is always present in star forming discs and provides sufficient heating to explain the in-plane AVR \citep{selcar, carsel}. From such two-dimensional simulations and analytical arguments for vertical cloud heating, \citet{carlberg} concluded that a combination of GMC and spiral heating could explain the observations. However, spirals do not directly increase $\sigma_z$ significantly (e.g.\ \citealp{sellwood13, martinez}), and the question remained open whether deflections by GMCs can convert in-plane heat to vertical heat in an appropriate manner. \citet{jenkins} used analytic arguments to examine this question for growing discs. They concluded that the observed value of $\sigma_z/\sigma_R$ could be explained, but the observed value of $\sigma_{{\rm old}}$ was problematic. Another continuous secular heating process that has been discussed in the literature is heating by the bar \citep{saha, grand}. The interaction of galactic discs with satellite galaxies and the corresponding dark matter substructure can also cause disc heating \citep{velazquez}. However, in the case of the MW this process has likely only made a minor contribution to the observed AVRs \citep{just}. Another contributor to the AVR could be a decline over cosmic time in the velocity dispersion of stars at their time of birth as discs have become less gas-rich and and less turbulent \citep{bournaud, forbes}. These models are motivated by observations of gas kinematics in disc galaxies at various redshifts, mostly based on H$\alpha$ emission. These observations have revealed significantly ($\sim3-10$ times) higher velocity dispersions $\sigma$ at redshifts $z\sim 2-4$ than in corresponding observations of the local universe (e.g. \citealp{sins}), and a decline of $\sigma$ with decreasing redshift \citep{wisn}. It is, however, unclear how the kinematics of young stars which form from cold gas, relate to these observations. Fully cosmological hydrodynamical simulations of galaxy formation have recently reached reasonable levels of success in reproducing MW like disc galaxies and the AVRs in some of these simulations have been studied \citep{house, bird, martig, grand}. At $z=0$ the stellar populations in the majority of these simulations are significantly hotter than the stars in the MW at all ages (but see model \emph{g92} of Martig et al.). Especially young stars have overly high $\sigma_z$ which has been linked to numerical effects and insufficient resolution and shown to depend on the specifics of the numerical sub-grid star-formation models \citep{house, martig}. Better agreement with the Snhd AVR is generally found for galaxies unaffected by mergers. Martig et al. find that the thin disc stars in their more successful models are born cold and heated with time, which they attribute to heating by spirals and overdensities and possibly the coupling between transient spirals and weak bending waves \citep{masset}. Note that GMCs are not resolved in these simulations. For many years it was assumed that the chemodynamical evolution of any annulus of the Galactic disc could be modelled in isolation of other annuli. Now there is clear evidence that radial migration of stars within discs is an important process \citep{sellwoodb,roskar, sb09, kordopatis}, with the consequence that the production of a hot population in one annulus, for example through the action of a bar, can subsequently endow a distant annulus with a hot population that could not have been locally heated. On account of radial migration, it is essential to understand the origin of the AVR \emph{globally}, that is by tracking the evolution of the disc at all radii. In general we expect the mean birth radius of a coeval cohort of Snhd stars will decrease with increasing age, and on account of the radial gradient in velocity dispersion the decrease in birth radius will be reflected in the AVR \citep{sb09}. In this paper we use the simulations presented in \citet[hereafter ABS16]{abs16} to study the formation of the AVR. Unlike the previously cited studies, these simulations include simultaneously all the following important aspects: growing discs with multiple coeval populations, GMCs, recurring spiral structure with evolving properties, a bar, an evolving GMC-to-stellar mass fraction, radial migration and sufficiently cold young stars. ABS16 showed that although the vertical profiles of their models do not show a thick disc like that of the MW, some models do provide quite good fits to the AVR of the Snhd. Hence they concluded that the thick disc requires additional sources of heat, but the thin disc can be explained by combined GMC and spiral heating. They showed that the efficiency of GMC heating declines over time because the fraction of the disc's mass contained in GMCs falls steadily as a consequence of a declining star-formation rate (SFR) and a growing disc mass. Their simulations are thus a promising tool to study what shapes the AVR in thin galactic discs. Two major conclusions will emerge from our study: (i) biased ages and age uncertainties cause measured AVRs to deviate significantly from the true AVRs; (ii) it is vital to distinguish between an AVR $\sigma(\tau)$, which gives velocity dispersion as a function of age for stars that are now co-located, and a \emph{heating history} $\sigma(t-t_{\rm b})$, which gives the velocity dispersion as a function of time for a cohort of currently co-located stars that were born at a given time $t_{\rm b}$. Whereas the AVR $\sigma(\tau)$ for $\tau\simeq4.5\gyr$ quantifies the current kinematics of stars born contemporaneously with the Sun, the heating history $\sigma(t-t_{\rm b})$ for $t-t_{\rm b}\simeq4.5\gyr$ quantifies the kinematics of stars $4.5\gyr$ after they were born, which would be $5.5\gyr$ ago in the case of $10\gyr$ old stars in the disc. If stars were born into a statistically stationary environment providing heating processes which are constant in time, the cohort born $10\gyr$ ago would $5.5\gyr$ ago have been in the same dynamical state that the Sun's cohort is in now. That is, given a stationary environment the AVR would be the same function of $\tau$ that the heating history is of $t-t_{\rm b}$. If a galaxy undergoes a major merger, stars born before and after the merger will undergo different heating histories $\sigma(t-t_{\rm b})$. Here we argue, that even in the absence of mergers or declining birth dispersions, the thin discs of galaxies change beyond recognition over cosmological timescales, so the environment is very far from stationary, and the heating experienced by stars born $10\gyr$ ago during the first Gyr of their lives was very different from the environment experienced by recently born stars during the first Gyr of their lives. Consequently heating histories are described by entirely different functions from the AVR. Nevertheless we will find that both AVRs and heating histories can be well approximated by the modified power law \begin{equation} \sigma(x)=\sigma_{10}\left({{x + x_1} \over {10\gyr + x_1}}\right)^{\beta} \label{eq:heatlaw} \end{equation} used by AB09, with $x=\tau$ or $t$. To differentiate between parameters derived from AVRs and heating histories, we will mark the latter with a tilde, i.e. $\tilde{\beta}$, $\tilde{\sigma}_{10}$ etc. We will find that the indices $\beta$ and $\tilde{\beta}$ of these power laws are often dissimilar. Moreover, we find that in the case of a heating history the value of $\tilde{\sigma}_{10}$ can evolve strongly with the time $t_{\rm b}$ of the cohort's birth, whereas in most models $\tilde{\beta}$ evolves only mildly. Our paper is organised as follows: In Section \ref{sec:simulations} we briefly describe the simulations. In Section \ref{sec:biases} we examine the effects of observational age errors and biases on the AVR. In Section \ref{sec:AVR} we describe the model AVRs and compare them to local data. Topics discussed include the uncertainties of the comparisons arising from azimuthal variations in the model AVRs (Sections \ref{sec:azimuth} and \ref{sec:specifics}), and the diagnostic content of the axis ratios of velocity ellipsoids (Section \ref{sec:arat}), and power-law fits to AVRs (Section \ref{sec:heatIndex}). In Section \ref{sec:heathistory} we consider the heating histories for different populations of coeval model stars and show how these relate to AVRs. In Section~\ref{sec:discuss} we relate our findings to the physics of star scattering, and we conclude in Section \ref{sec:conclude}. \tabcolsep=4.5pt \begin{table*} \vspace{-0cm} \caption{List of models analysed in this paper. {\it 1st Column}: Model Name; {\it 2nd Column}: Initial Conditions; {\it 3rd Column}: Total baryonic IC mass $M_{\rm{b,i}}$; {\it 4th Column}: IC DM halo scalelength {$a_{\rm halo}$}; {\it 5th Column}: IC radial disc scalelength $h_{R,{\rm disc}}$; {\it 6th Column}: IC vertical disc scaleheight $z_{0,{\rm disc}}$; {\it 7th Column}: GMCs Yes/No; {\it 8th Column}: Cutoff: Adaptive (no new particles in bar region) or fixed (pre-defined evolving inner cutoff for new particles); {\it 9th Column}: Total inserted baryonic model mass $M_{\rm f}$ (including initial baryonic mass); {\it 10th Column}: Initial disc scalelength $h_{R, {\rm i}}$; {\it 11th Column}: Final disc scalelength $h_{R, {\rm f}}$; {\it 12th Column}: Scalelength growth parameter $\xi$; {\it 13th Column}: Exponential decay timescale $t_{\rm SFR}$ for the star formation rate; {\it 14th Column}: Initial velocity dispersion for inserted stellar particles, {$\sigma_0$}; {\it 15th Column}: GMC star formation efficiency $\zeta$; } \begin{tabular}{@{}ccccccccccccccc@{}}\hline 1st & 2nd & 3rd & 4th & 5th & 6th &7th & 8th & 9th & 10th & 11th & 12th & 13th & 14th & 15th \\ {Name}&{ICs} &{$M_{\rm{b,i}}$} &{$a_{\rm halo}$}&{$h_{R,{\rm disc}}$}&{$z_{0,{\rm disc}}$}&{GMCs}&{Cutoff}&{$M_{\rm f}/\msun$}&{$h_{R, {\rm i}}$}& {$h_{R, {\rm f}}$}& {$\xi$} & {$t_{\rm SFR}$}&{$\sigma_0$}& {$\zeta$}\\ & &{$[10^{9}\msun]$}&{$\kpc$} &{$\kpc$} &{$\kpc$} & & &{$[10^{10}]$} &{$\kpc$} & {$\kpc$} & & {$[\rm Gyr]$}&{$[\kms]$} &\\ \hline Y1 & Y & 5 & 30.2 & 1.5 & 0.1 & Yes & Adap & 5 & 1.5 & 4.3 & 0.5 & 8.0 & 6 & 0.08 \\ Y1s2 & Y & 5 & 30.2 & 1.5 & 0.1 & Yes & Adap & 5 & 1.5 & 4.3 & 0.5 & 16.0 & 6 & 0.08 \\ Y1$\zeta $- & Y & 5 & 30.2 & 1.5 & 0.1 & Yes & Adap & 5 & 1.5 & 4.3 & 0.5 & 8.0 & 6 & 0.04 \\ Y1f$\sigma$ & Y & 5 & 30.2 & 1.5 & 0.1 & Yes & Fix & 5 & 1.5 & 4.3 & 0.5 & 8.0 & 10 & 0.08 \\ Y2 & Y & 5 & 30.2 & 1.5 & 0.1 & Yes & Adap & 5 & 2.5 & 2.5 & 0.0 & 8.0 & 6 & 0.08 \\ Y2Mb- & Y & 5 & 30.2 & 1.5 & 0.1 & Yes & Adap & 3 & 2.5 & 2.5 & 0.0 & 8.0 & 6 & 0.08 \\ Y2Mb+ & Y & 5 & 30.2 & 1.5 & 0.1 & Yes & Adap & 7.5 & 2.5 & 2.5 & 0.0 & 8.0 & 6 & 0.08 \\ Y4f$\zeta$- & Y & 5 & 30.2 & 1.5 & 0.1 & Yes & Fix & 5 & 1.5 & 2.2 & 0.5 & 8.0 & 6 & 0.04 \\ YG1 & YG & 5 & 30.2 & 1.5 & 0.1 & Yes & Adap & 5 & 1.5 & 4.3 & 0.5 & 8.0 & 6 & 0.08 \\ YN1 & Y & 5 & 30.2 & 1.5 & 0.1 & No & Adap & 5 & 1.5 & 4.3 & 0.5 & 8.0 & 6 & -- \\ A2$\tau$ & A & 10 & 30.2 & 1.5 & 0.8 & Yes & Adap & 5 & 2.5 & 2.5 & 0.0 & 8.0&$6 + 30\e^{-t/1.5\gyr}$ & 0.08 \\ E2 & E & 15 & 30.2 & 2.5 & 1.2 & Yes & Adap & 5 & 2.5 & 2.5 & 0.0 & 8.0 & 6 & 0.08 \\ F2 & F & 5 & 51.7 & 1.5 & 0.1 & Yes & Adap & 5 & 2.5 & 2.5 & 0.0 & 8.0 & 6 & 0.08 \\ \hline \end{tabular} \label{modeltable} \end{table*}

\label{sec:conclude} In this paper, we have used a series of $N$-body simulations of growing disc galaxies (ABS16) to study (i) age-velocity-dispersion relations (AVRs) and (ii) the heating histories of the coeval cohorts of stars which make up the AVRs. As these models feature heavy GMC particles, secular heating is dominated by a combination of scattering of stars off GMCs and non-axisymmetric disc structures. To be able to compare these simulations to observational data from the Snhd, we analysed the impact on the AVR of biases and errors in measured stellar ages. Stars with ages $\tau\sim2\gyr$ are very much over-represented in the GCS data (Fig.~\ref{gcs}). Scattering of such stars to young ages artificially boosts $\sigma(\tau)$ at the youngest ages, and depresses $\sigma(\tau)$ at the oldest ages. When a power law in $\tau$ is fitted to the measured AVR, lower values of the exponent $\beta$ are recovered than would be in the absence of errors (see also \citealp{martig}). The reduction in $\beta$ is particularly marked in the case of $\sigma_z$ (Fig.~\ref{heatfits}). On account of spiral structure and bars, AVRs vary with azimuth as well as radius. Fig.~\ref{azi} quantifies the extent of this azimuthal variation, which must be borne in mind when considering whether a given model is consistent with data for the Snhd, which are measured at one particular azimuth. After taking azimuthal variation into account, we concluded that the GCS data are consistent with some models in ABS16 that have the expected disc mass ($5\times10^{10}\msun$) and the cosmologically motivated dark halo ($M=10^{12}\msun$, $a=30.2\kpc$). Models with a significantly different disc mass or a less concentrated dark halo are inconsistent with data for the Snhd. The data also favour the model that starts with a massive, extended thick disc over models in which (a rather inadequate) thick disc forms as a consequence of powerful non-axisymmetries developing in the thin disc. As we do not self-consistently form appropriate thick discs and as we lack heating by dark matter substructure, which may contribute a minor part of the observed disc heating, we are not able to put tight constraints on our model parameters. AVRs vary with radius. At locations currently inside the bar, the AVR's index $\beta_R$ is generally very small, $\beta_R<0.1$, as there are no circular orbits and young stars thus acquire high $\sigma_R$ rapidly. At the end of the bar, $\beta_R$ rises abruptly and is thereafter constant or slowly rising with $R$ (Fig.~\ref{heatrad}). By contrast, $\beta_z$ sometimes increases and sometimes decreases at the end of the bar. A buckling bar can lead to exceptionally high $\sigma_z$ for young stars in the bar regions. The heating history, $\sigma(t-t_{\rm b})$, of stars now in the Snhd that were born at time $t_{\rm b}$ can also be fitted by the power-law \eqref{eq:heatlaw}. We mark the corresponding parameters with a tilde. For standard models, the heating history depends on $t_{\rm b}$ more strongly in the case of $\sigma_z$ than $\sigma_R$. Smaller values of the exponent $\tilde{\beta}$ are required to fit heating histories than AVRs. In fact, values of $\tilde{\beta}_R$ are consistent with the predictions of dynamics in the case that the fluctuating gravitational potential is a stationary random process. The values of $\tilde{\beta}_z$ are generally somewhat larger than is consistent with a stationary random process, but in agreement with numerical simulations of stationary GMC heating \citep{haenninen}. The AVR reflects the history of star and bar formation. The past SFR strongly affects the AVR for two reasons: the time integral of the SFR determines the mass of the disc, and thus the fraction of the gravitational force on a star that derives from the disc rather than the dark halo. At early times this fraction is small, so spiral arm formation is already suppressed by a low value of $\sigma_R$. As the mass of the disc increases, spiral structure increases $\sigma_R$ to keep Toomre's $Q$ nearly constant. If the SFR is rapidly declining, the rate at which $\sigma_R$ increases will decline rapidly, and a relatively small value of $\tilde{\beta}_R$ will be required to fit the heating histories of the oldest stellar groups. In contrast, the vertical heating is dominated by GMCs. By analysing histories of stars that end up at $R=8\kpc$, we showed that the heating histories of older stars reach higher $\sigma_z$ after $10\gyr$ ($\tilde{\sigma}_{10}$) than those of younger stars. The corresponding $\tilde{\beta}_z$ values only vary mildly with $t_{\rm b}$. This decline in heating efficiency is connected to the declining influence of GMCs, the total mass of which declines due to a declining SFR and the mass fraction of which declines due to a growing stellar disc. When coeval cohorts, whose $\tilde{\sigma}_{10}$ values decline with $t_{\rm b}$, are superposed to form an AVR, a value of $\beta_z$ in excess of $0.5$ is needed to fit the curve. By combining all these results we have been able to clarify long standing discrepancies between the observed AVR and theoretical predictions for combined spiral and GMC heating. Some of our models correctly reproduce the general shape of both $\sigma_R(\tau)$ \emph{and} $\sigma_z(\tau)$ as observed in the Snhd and thus also the ratio of the two components. The key ingredient is that each coeval cohort of stars that contributes to the AVR has undergone a different heating history and the AVR is not produced by a single stationary heating law. We conclude that combined GMC and spiral/bar heating has likely shaped the MW thin disc AVR.

1607.01972

1607

1607.05283_arXiv.txt

\noindent The high energy events observed at the IceCube Neutrino Observatory have triggered many investigations interpreting the highly energetic neutrinos detected as decay products of heavy unstable Dark Matter particles. However, while very detailed treatments of the IceCube phenomenology exist, only a few references focus on the (non-trivial) Dark Matter production part -- and all of those rely on relatively complicated new models which are not always testable directly. We instead investigate two of the most minimal scenarios possible, where the operator responsible for the IceCube events is directly involved in Dark Matter production. We show that the simplest (four-dimensional) operator is not powerful enough to accommodate all constraints. A more non-minimal setting (at mass dimension six), however, can do both fitting all the data and also allowing for a comparatively small parameter space only, parts of which can be in reach of future observations. We conclude that minimalistic approaches can be enough to explain all data required, while complicated new physics seems not to be required by IceCube.

Introduction} The IceCube Neutrino Observatory, a neutrino telescope located at the Amundsen-Scott South Pole Station, is a unique window to observe highly energetic neutrinos reaching the Earth's surface, originating from sources as close as the upper regions of the atmosphere up to extra-galactic objects~\cite{Waxman:2014vja}. Its applications to closer sources range from a more precise determination of the atmospheric neutrino flux~\cite{Aartsen:2015rwa} over measuring the properties of active neutrinos~\cite{Aartsen:2014yll} and constraining those of sterile neutrinos~\cite{TheIceCube:2016oqi} to astrophysical findings such as the shadowing effect of the moon on cosmic rays~\cite{MoonShadow}. As for the wider sources, IceCube's goal is to investigate several types of astrophysical neutrino emitters, its possible applications ranging from astrophysical point sources~\cite{Aartsen:2016tpb} over Dark Matter annihilation~\cite{Aartsen:2016pfc} to supernovae~\cite{Aartsen:2015trq}. Finally, also certain exotic particles may leave visible signatures in the detector, such as magnetic monopoles~\cite{Aartsen:2015exf}. A big surprise in the data taken between 2010 and 2013 was the detections of three very high energy events, reported in Refs.~\cite{Aartsen:2013bka,Aartsen:2013jdh,Aartsen:2014gkd}. These events have been under such scrutiny and have generated such an amount of interest, that they have even been given names after characters of the Sesame Street~\cite{SesameStreet} for better recognition: Ernie (1.14~PeV), Bert (1.04~PeV), and Big Bird (2.2~PeV). The origin of these very high energy events is still unclear, though. The initial discussion was immediately targeting various astrophysical sources, see Refs.~\cite{Cholis:2012kq,Anchordoqui:2013dnh,Murase:2014tsa} for comprehensive treatments and extensive lists of references. However, in the particle physics community, great interest arose instead in relating the detections to the physics of Dark Matter (DM), in order to address one of the most fascinating topics in all of science. It had been argued that such high energy events probably cannot originate from DM \emph{annihilation}~\cite{Bai:2013nga} because of the unitarity bound~\cite{Griest:1989wd,Hui:2001wy}. Thus, although this bound may be circumvented~\cite{Zavala:2014dla}, most works have focused on DM \emph{decay} instead. Looking at the literature, most authors consider the decay of superheavy DM-type particles~\cite{Feldstein:2013kka,Bai:2013nga,Esmaili:2013gha,Ema:2013nda,Bhattacharya:2014vwa,Fong:2014bsa,Rott:2014kfa,Dudas:2014bca,Murase:2015gea,Anchordoqui:2015lqa,Ko:2015nma,Aisati:2015vma,Higaki:2014dwa,Roland:2015yoa,Esmaili:2014rma,Boucenna:2015tra}, although some work has also been presented on lower-mass candidates boosted to high energies~\cite{Bhattacharya:2014yha,Kopp:2015bfa}. In general, depending on the interaction between DM and Standard Model particles, the decays of superheavy DM particles may be able to account for the whole TeV--PeV IceCube diffuse neutrino flux (see for instance Ref.~\cite{Esmaili:2014rma}) -- or at least for part of it, as shown in Ref.~\cite{Boucenna:2015tra}, where the TeV neutrinos events are explained in terms of an astrophysical power-law flux (two-component flux). While all kinds of phenomenological aspects of the signal are considered, like e.g.\ its variation with the DM profile~\cite{Chianese:2016opp}, most settings are not specified very accurately from the particle physics side, making it tempting to unify the treatments based on a set of effective operators mediating DM decay~\cite{Esmaili:2014rma,Boucenna:2015tra}. Although the IceCube part has been treated in great detail, the literature on how to produce such a type of DM in the first place appears a bit scarce in comparison. Nevertheless, there are some notable exceptions which treat the full course of DM production down to an analysis of the IceCube signal: in the examples found, the DM particles are e.g.\ produced in a secluded sector~\cite{Dev:2016qbd}, by freeze-out with resonantly enhanced annihilations~\cite{DiBari:2016guw}, or via \emph{freeze-in}~\cite{Fiorentin:2016avj,Roland:2015yoa,Fong:2014bsa}. It is this latter mechanism we would also like to focus on in our current work. While Ref.~\cite{Fiorentin:2016avj} investigated a full model based on left-right symmetry, we go the opposite way and try to be very minimalistic by asking the question which of the possible operators mediating DM decay~\cite{Esmaili:2014rma,Boucenna:2015tra} could \emph{at the same time} be responsible for DM production in the early Universe. We will in particular focus on the 4-dimensional operator discussed in Ref.~\cite{Esmaili:2014rma}, which allows the DM particle to directly decay into a neutrino and a SM Higgs, as well as on an alternative leptophilic 6-dimensional operator which has a somewhat richer phenomenology and features the same predictions as the one discussed in Ref.~\cite{Boucenna:2015tra}. As we will see, while the minimal ($d=4$)-operator is in fact not sufficient to bring DM production in accordance with the IceCube signal (unless both parts are completely disentangled, as in Ref.~\cite{Fiorentin:2016avj}), the ($d=6$)-operator turns out to be powerful enough: not only can it accommodate for all data and bounds, but it actually leaves us with a potentially testable allowed window. We therefore show that, beyond the ingredients needed for DM production and (of course) a candidate DM particle, no complicated new physics is needed to ensure both consistency and testability. This paper is structured as follows. We start by introducing the basic underlying setup in Sec.~\ref{sec:basic}, before giving a general discussion on the necessary characteristics of decaying DM in Sec.~\ref{sec:general}. DM production with the different operators is discussed in detail in Sec.~\ref{sec:DM-production}, before our numerical results are presented and discussed in Sec.~\ref{sec:Results}. We finally conclude in Sec.~\ref{sec:conc}. Technical details are given in App.~\ref{app:A}, which lists the explicit expressions for all matrix elements used in the computation of DM production.

1607.05283

1607

1607.07637_arXiv.txt

The photometric sky quality of Mt.~Shatdzhatmaz, the site of Sternberg Astronomical Institute Caucasian Observatory 2.5~m telescope, is characterized here by the statistics of the night-time sky brightness and extinction. The data were obtained as a by-product of atmospheric optical turbulence measurements with the MASS (Multi-Aperture Scintillation Sensor) device conducted in 2007--2013. The factors biasing night-sky brightness measurements are considered and a technique to reduce their impact on the statistics is proposed. The single-band photometric estimations provided by MASS are easy to transform to the standard photometric bands. The median moonless night-sky brightness is $22.1$, $21.1$, $20.3$, and $19.0$ mag per square arcsec for the $B$, $V$, $R$, and $I$ spectral bands, respectively. The median extinction coefficients for the same photometric bands are $0.28$, $0.17$, $0.13$, and $0.09$ mag. The best atmospheric transparency is observed in winter.

It is well known that the efficiency of classical ground-based astronomical observations (resolution and limiting magnitude) depends strongly on the atmospheric seeing~\citep{Bowen1964}, but other astroclimatic parameters also play an important role. An overall telescope productivity is proportional to the clear sky fraction. Observations of faint object depend strongly on the night-sky brightness, which includes a contribution from the light pollution. The accuracy of photometric observations is determined by temporal and spatial stability of the atmospheric extinction. Since the Mt.~Shatdzhatmaz on North Caucasus had been chosen as a place for the new 2.5~m telescope, we initiated in 2007 a long-term monitoring of the atmospheric seeing and other astroclimatic parameters at this site, pursuing two main goals: first, support the operation of the 2.5~m telescope, and, second, gain a better understanding of the astroclimate of North Caucasus. The first results of these measurements are presented in \citep{2010MNRAS} where the used hardware and technique are also described. The summary results of the optical turbulence studies of the 2007--2013 campaign are given in the paper \citep{2014PASP}. In this paper we present the results of measurements of the photometric parameters, the night-sky brightness and the atmospheric extinction. The night-sky background was measured during optical turbulence monitoring with the MASS device \citep{2007bMNRAS} because it is required to properly calculate scintillation indices \citep{2003MNRAS, 2007aMNRAS}. More than $30\,000$ such sky background estimates were obtained during the campaign. Since MASS is essentially a fast high-precision photoelectric photometer, it also allows us to estimate atmospheric extinction. To this end, the optical turbulence program was complemented by special extinction measurements in 2009. The first results of these measurements are given by \citet{Voziakova2012AstL}. More comprehensive atmospheric extinction data are presented in Section~\ref{sec:extinction}. In the same Section, it is shown how to transform the extinction measured in the MASS spectral band to the standard $U$, $B$, $V$, $R$, and $I$ photometric system. The statistics of sky background estimates are given in Section~\ref{sec:background}. They are transformed to the conventional units of stellar magnitude per square arcsec using the known magnitudes of the MASS program stars and the atmospheric extinction. The results presented in this paper allow us to comprehensively characterise the Sternberg Astronomical Institute (SAI) observatory, facilitating the optimum scheduling of the 2.5~m telescope. Monitoring of the astroclimatic parameters will help to operate the telescope in the most efficient way.

The above results illustrate that a slight modification of the optical turbulence measurement algorithm allows us to obtain additional information from MASS/DIMM measurements. A possibility to study photometric properties of the atmosphere with MASS/DIMM was pointed out by us a long time ago \citep{2007bMNRAS}. In order to determine the extinction in the MASS photometric band, the widely used method of pairs was adopted. Its ability to yield nearly instantaneous extinction estimates with minimum resources was fully employed. This method requires knowledge of precise above-atmosphere instrumental magnitudes of the reference stars. At least a year-long cycle of measurements is needed to obtain a consistent solution of the respective system of photometric equations. Meanwhile, the weather at Mt.~Shatdzhatmaz is such that a significant part of observing time occurs on partially clear nights, when the use of other methods of extinction estimates is not feasible. The transformation of the measured MASS band extinction coefficients into the standard $B$, $V$, $R$, and $I$ photometric system leads to the following median extinction values: $0.28$, $0.17$, $0.13$, and $0.09$~mag, respectively. It is worth to recall that a specific task to monitor the night sky brightness has not been initially planned and, as result, the data presented in Section~\ref{sec:background} suffer from uncertainty and incompleteness. In particular, the unaccounted for light scattering in the telescope and instrument led to significant errors in the individual estimates of the sky brightness during moonless periods. Still, it was possible to derive reasonable statistical estimates of the night sky background. Being transformed into the standard photometric system, they are $22.1$, $21.1$, $20.3$, and $19.0$~mag per square arcsec in the $B$, $V$, $R$, and $I$ bands, respectively. The scattered light and star contamination are not important while measuring the sky background with moon above horizon. However, building a consistent model of the moon background at our site, similar to \citep{Jones2013AA}, is not straightforward because it contains many input parameters and requires additional measurements. Monitoring of the photometric atmospheric properties will definitely be continued. The algorithm of the background measurement has already been improved. Measurements are now performed at 27 arcmin offset from the program star and in two positions. The extinction cycle now includes two low altitude stars in different parts of the sky. At the same time, the emphasis of photometric measurements has now shifted to the real time use of the data for flexible planning of observations, while achieving statistical completeness has a lower priority. Special programs may acquire more time during optical turbulence monitoring and optimally will include a feedback from observations schedulers working for the telescopes at this site.

1607.07637

1607

1607.02965_arXiv.txt

BlackGEM is an array of telescopes, currently under development at the Radboud University Nijmegen and at NOVA (Netherlands Research School for Astronomy). It targets the detection of the optical counterparts of gravitational waves. The first three BlackGEM telescopes are planned to be installed in 2018 at the La Silla observatory (Chile). A single prototype telescope, named MeerLICHT, will already be commissioned early 2017 in Sutherland (South Africa) to provide an optical complement for the MeerKAT radio array. The BlackGEM array consists of, initially, a set of three robotic 65-cm wide-field telescopes. Each telescope is equipped with a single STA1600 CCD detector with 10.5k x 10.5k 9-micron pixels that covers a 2.7 square degrees field of view. The cryostats for housing these detectors are developed and built at the KU Leuven University (Belgium). The operational model of BlackGEM requires long periods of reliable hands-off operation. Therefore, we designed the cryostats for long vacuum hold time and we make use of a closed-cycle cooling system, based on Polycold PCC Joule-Thomson coolers. A single programmable logic controller (PLC) controls the cryogenic systems of several BlackGEM telescopes simultaneously, resulting in a highly reliable, cost-efficient and maintenance-friendly system. PLC-based cryostat control offers some distinct advantages, especially for a robotic facility. Apart of temperature monitoring and control, the PLC also monitors the vacuum quality, the power supply and the status of the PCC coolers (compressor power consumption and temperature, pressure in the gas lines, etc.). Furthermore, it provides an alarming system and safe and reproducible procedures for automatic cool down and warm up. The communication between PLC and higher-level software takes place via the OPC-UA protocol, offering a simple to implement, yet very powerful interface. Finally, a touch-panel display on the PLC provides the operator with a user-friendly and robust technical interface. In this contribution, we present the design of the BlackGEM cryostats and of the PLC-based control system.

\label{sec:intro} The announcement earlier this year of the first direct detection of gravitational waves by the LIGO detectors brought gravitational wave astrophysics to the centre of science and research. In order to better understand these gravitational wave events and to maximize the science return from their detection, it will be essential to observe them in the electromagnetic domain as well. For that reason, the BlackGEM project \cite{bloemen2016} was started by the Department of Astrophysics of the Radboud University Nijmegen (NL), partnering with NOVA (NL) and the Institute of Astronomy of the KU Leuven University (BE). BlackGEM is an array of telescopes dedicated to and optimized for the search of the optical counterparts of gravitational wave events, like the merging of two neutron stars or the collision of a neutron star with a black hole. \begin{wrapfigure}{R}{9cm} \resizebox{9cm}{!}{\includegraphics{blackgem_telescope.pdf}} \caption{ \label{fig:telescope} BlackGEM telescope design.} \end{wrapfigure} However, the observations of these electromagnetic counterparts are not straightforward. These sources will probably be faint ($M_{V}\sim22$) and transienting quickly ($\sim$1~day). The large error window of GW detections by facilities like LIGO and VIRGO (typically $\sim$100 square degrees) requires BlackGEM to search with high sensitivity over a very large field of view. Therefore, each telescope will provide seeing-limited imaging over a wide field of 2.7 square degrees. This gives a total sky coverage of 8 square degrees for an array of 3 telescopes. It is envisioned that in a later stage and with the entrance of new partners in the project, the array will be extended with more telescopes, leading to an even larger total coverage. The array consists of identical Dall-Kirkham Cassegrain telescopes with a 65-cm primary mirror ($f/5.5$) in a carbon fibre structure. The telescopes will be fully robotic with a goal of one-year periods of maintenance-free and unattended operation. Fig.~\ref{fig:telescope} shows a design model of the BlackGEM telescope. We plan to install and commission BlackGEM at the ESO La Silla Observatory in Chile early 2018. Currently, a prototype for the BlackGEM telescopes is already under construction. After assembly, integration and testing in the Netherlands, this telescope named MeerLICHT, will be installed in Sutherland (South Africa) at the MeerKAT radio array site. There it will complement the radio observations from MeerKAT, the precursor to the Square Kilometre Array (SKA). MeerLICHT will co-point the same field as the radio dishes and provide simultaneous optical observations. This contribution focuses on the design of the cryostat that houses the CCD detector and the associated cryostat monitoring and control system. The outline of the paper is as follows: in section 2 we introduce the BlackGEM/MeerLICHT (hereafter solely referred to as BlackGEM) cryostats as well as their detector system, section 3 discusses the thermo-mechanical design of the cryostat, in section 4 we present the PLC-based cryostat control system, followed by a performance discussion in section 5 and some conclusions in section 6.

\label{sec:conclusions} We presented the design of the cryostat and cryostat control for the large BlackGEM detectors. A prototype has been built and successfully characterised, and will now be installed on BlackGEM's precursor, the MeerLICHT telescope. We argue that PLC technology is very well suited for controlling detector cryostats, a task that requires highly dependable hard- and software. This includes accurate temperature stabilisation as well as robotic control of closed-cycle coolers, and a convenient OPC-UA interface to high-level software.

1607.02965

1607

1607.03478_arXiv.txt

In this work the homogeneous and isotropic Universe of Friedmann-Robertson-Walker is studied in the presence of two fluids: stiff matter and radiation described by the Schutz's formalism. We obtain to the classic case the behaviour of the scale factor of the universe. For the quantum case the wave packets are constructed and the wave function of the universe is found.

% Quantum cosmology \cite{Israel, Wiltshire} is one of the aspects explored in the quantum gravity program. Its application to the Universe as a whole takes place at very early times the universe was small indeed. For exemple, in the Planck time, about $10^{-44}$ s after the Big Bang, its size was the order of $10^{-33}$ cm suggesting that quantum effects dominated \cite{Atkatz}. In general the quantum cosmological models are built considering a finite number of degrees of freedom in the presence of geometrical symmetry inherent to the universe. The quantization of a system subjected to these conditions is called quantization in a minisuperspace. Furthermore, with the quantum cosmology it is possible to study a real specific system in which quantum effects of gravity are fundamental. As the early Universe is the best of laboratories where different quantum theories of gravitation can be tested, quantum cosmology has as one of its motivations to serve as a powerful auxiliary tool in the construction of a final and fundamental quantum theory of gravity. \par Considering only the cosmological point of view, quantum cosmology plays an important role in solving a serious problem of the standard cosmological model: The existence of an initial singularity. Mechanisms such as quantum tunneling used in more simple cosmological models or in sophisticated approaches such as volume quantization provided by the loop formalism are examples of consistent solutions of this problem \cite{Bojowald}. In addition, the theory allows us to establish the initial conditions for inflation, for primordial perturbations and for spontaneous symmetry breaking \cite{Konstantin,Pad}. That is, the quantum cosmology is a theory of initial conditions. \par Minisuperspace quantization is a reasonable approximation to describes certain quantum gravitational effects. It is expected however that the inclusion of other degrees of freedom in the problem must refine the results obtained and lead to a more accurate description of the phenomenon studied, even lead to additional technical complications. This addition of new degrees of freedom can be done in several different ways. In the hydrodynamic description of the material content, consistent and obvious way to refine the description of the Universe is to use not just one but two fluids \cite {Fracalossi, Tovar}. \par In this paper we choose stiff matter ($\alpha = 1$) and radiation ($\alpha = 1/3$) as matter content, which play an important role in the early Universe. The energy density of this fluid in the cosmological gauge ($N(t) = 1$) is proportional to $1/a(t)^4$. As the energy density of the stiff matter \cite{Zeldovich}, in the same gauge, is proportional to $1/a(t)^6$, there must have been a time before the radiation in which the stiff matter dominated. In fact, the abundance of particles produced after the Big Bang due to expansion and cooling of the Universe is an implication of the presence of this fluid with $\alpha = 1$ at this stage \cite{Kamionkowski, Joyce, Salati, Pallis, Gomez, Pallis2, Germano}. \par The material content is treated here according to the Schutz formalism \cite{Schutz1, Schutz2}. With this description we can make the degrees of freedom of the fluid plays the role of time in the theory, transforming the Wheeler-DeWitt equation, which governs the dynamic behaviour of the system, in a Schr\"{o}dinger-like equation whose solution is a time-dependent wave function, among other variables. So the complex problem of the absence of a variable linked to the time evolution in quantum gravity is solved giving to time a purely phenomenological character leading to a well-defined Hilbert space structure. This paper is organized as follows: Section (\ref{class}) presents the Schutz's formalism in a classical FRW cosmological model with stiff matter and radiation. The analytical evolution of the scale factor of the universe, $a(t)$, is obtained. We emphasize that, since this point, time is introduced following the Schutz's formalism. We present next, Section (\ref{quant}), the quantum model and we obtain the equation that governs the dynamics of the model, the Wheeler-DeWitt equation, which is solved. A detailed discussion and summary of these results is provided in Section (\ref{concl}).

\label{concl} In this paper, we study the quantum FRW model with stiff matter and radiation and we solve the Wheeler-DeWitt equation in minisuperspace to obtain the wave function of the corresponding universe. The perfect fluid is described by the Schutz's canonical formalism, which allows to attribute dynamical degrees of freedom to matter. In the classical model, the universe of null curvature expands forever from an initial singularity, as can be seen in equations (\ref{a_conforme}), (\ref{t_cosmico}) and in Figure \ref{classico}. Indeed, this singular point can be avoided in quantum cosmology. In quantum treatment the model of fluids was investigated on the basis of the associated Wheeler-DeWitt equation. The introduction of a time variable phenomenologically using the degrees of freedom of radiation fluid allows to obtain a Hamiltonian constraint linear in one of the momenta and the Wheeler-DeWitt equation can be reduced to a Schr\"odinger like equation, whose solution is the wave function of the universe. Applying the boundary conditions, this function was explicitly written and the wave packets was constructed analytically. The behavior of the wave packet (\ref{16}) is showed in Figures \ref{psi_a_sigma}, \ref{psi_a_t} and \ref{psi_sigma_t} for each two pairs of variables $a$, $\sigma$ and $t$. The packet goes to $0$ as the variables grow, but the decaying rate depends on the values of the third fixed variable and the free parameters $\gamma$ and $\xi$. It is worth to note that in equation (\ref{16}) the variable $t$ appears many times accompanied by the parameter $\gamma$ and $\sigma$ by $\xi$. This fact can be used as a manner to regulate how fast the packet approaches zero for the directions $\sigma$ and $t$. In Figure \ref{psi_1dim} we show the cross section of the wave packet and see how different combinations of values for the variables and the parameters can change the intensity and position of the peak, as the decaying rate. The dependence in relation to the free parameter is showed in Figure \ref{psi_1dim_gamma_xi}. The behavior of the wave packet, shown in Figures (\ref{psi_a_sigma}) and (\ref{psi_a_t}), makes it clear that $\Psi(a,\sigma,t)$ tends to zero when $a$ tends to zero. Because of this, the probability density $P(a)$ tends to zero when $a$ tends to zero, so that the classical singularity is avoided, unlike what happens in the classical case. As perspective to future work, the investigation of new scenarios containing other types of fluid is the natural way. It is hoped that this will highlight the role played by new degrees of freedom added to cosmological models, as well as contribute to the discussion about the quantization procedures applied to cosmology.

1607.03478

1607

1607.03152_arXiv.txt

We develop a new methodology called double-probe analysis with the aim of minimizing informative priors in the estimation of cosmological parameters. Using our new methodology, we extract the dark-energy-model-independent cosmological constraints from the joint data sets of Baryon Oscillation Spectroscopic Survey (BOSS) galaxy sample and Planck cosmic microwave background (CMB) measurement. We measure the mean values and covariance matrix of $\{R$, $l_a$, $\Omega_b h^2$, $n_s$, $log(A_s)$, $\Omega_k$, $H(z)$, $D_A(z)$, $f(z)\sigma_8(z)\}$, which give an efficient summary of Planck data and 2-point statistics from BOSS galaxy sample. The CMB shift parameters are $R=\sqrt{\Omega_m H_0^2}\,r(z_*)$, and $l_a=\pi r(z_*)/r_s(z_*)$, where $z_*$ is the redshift at the last scattering surface, and $r(z_*)$ and $r_s(z_*)$ denote our comoving distance to $z_*$ and sound horizon at $z_*$ respectively; $\Omega_b$ is the baryon fraction at $z=0$. The advantage of this method is that we do not need to put informative priors on the cosmological parameters that galaxy clustering is not able to constrain well, i.e. $\Omega_b h^2$ and $n_s$. Using our double-probe results, we obtain $\Omega_m=0.304\pm0.009$, $H_0=68.2\pm0.7$, and $\sigma_8=0.806\pm0.014$ assuming $\Lambda$CDM; $\Omega_k=0.002\pm0.003$ assuming oCDM; $w=-1.04\pm0.06$ assuming $w$CDM; $\Omega_k=0.002\pm0.003$ and $w=-1.00\pm0.07$ assuming o$w$CDM; and $w_0=-0.84\pm0.22$ and $w_a=-0.66\pm0.68$ assuming $w_0w_a$CDM. The results show no tension with the flat $\Lambda$CDM cosmological paradigm. By comparing with the full-likelihood analyses with fixed dark energy models, we demonstrate that the double-probe method provides robust cosmological parameter constraints which can be conveniently used to study dark energy models. We extend our study to measure the sum of neutrino mass using different methodologies including double probe analysis (introduced in this study), the full-likelihood analysis, and single probe analysis. From the double probe analysis, we obtain $\Sigma m_\nu<0.10/0.22$ (68\%/95\%) assuming $\Lambda$CDM and $\Sigma m_\nu<0.26/0.52$ (68\%/95\%) assuming $w$CDM. This paper is part of a set that analyses the final galaxy clustering dataset from BOSS.

\label{sec:intro} We have entered the era of precision cosmology along with the dramatically increasing amount of sky surveys, including the cosmic microwave background (CMB; e.g., \citealt{Bennett:2012zja,Ade:2013sjv}), supernovae (SNe; \citealt{Riess:1998cb,Perlmutter:1998np}), weak lensing (e.g., see \citealt{VanWaerbeke:2003uq} for a review), and large-scale structure from galaxy redshift surveys, e.g. 2dF Galaxy Redshift Survey (2dFGRS; \citealt{Colless:2001gk,Colless:2003wz}, Sloan Digital Sky Survey (SDSS, \citealt{York:2000gk,Abazajian:2008wr}, WiggleZ \citep{Drinkwater:2009sd, Parkinson:2012vd}, and the Baryon Oscillation Spectroscopic Survey (BOSS; \citealt{Dawson:2012va,Alam:2015mbd}) of the SDSS-III \citep{Eisenstein:2011sa}. The future galaxy redshift surveys, e.g. Euclid\footnote{http://sci.esa.int/euclid} \citep{Laureijs:2011gra}, Dark Energy Spectroscopic Instrument \footnote{http://desi.lbl.gov/} (DESI; \citealt{Schlegel:2011zz}), and WFIRST\footnote{http://wfirst.gsfc.nasa.gov/} \citep{Green:2012mj}, will collect data at least an order of magnitude more. It is critical to develop the methodologies which could reliably extract the cosmological information from such large amount of data. The galaxy redshifts samples have been analysed studied in a cosmological context (see, e.g., \citealt{Tegmark:2003uf,Hutsi:2005qv,Padmanabhan:2006cia,Blake:2006kv,Percival:2007yw,Percival:2009xn,Reid:2009xm, Montesano:2011bp, Eisenstein:2005su,Okumura:2007br,Cabre:2008sz,Martinez:2008iu,Sanchez:2009jq,Kazin:2009cj,Chuang:2010dv,Samushia:2011cs,Padmanabhan:2012hf,Xu:2012fw, 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}). \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 Hubble expansion rate $H(z)$ and angular-diameter distance $D_A(z)$. \cite{Chuang:2011fy} demonstrated the feasibility of measuring $H(z)$ 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 and it has been improved later on in \cite{Chuang:2012ad,Chuang:2012qt} upgrading the methodology and modelling to measure $H(z)$, $D_A(z)$, the normalized growth rate $f(z)\sigma_8(z)$, and the physical matter density $\Omega_m h^2$ from the same data. Analyses have been perform to measure $H(z)$, $D_A(z)$, and $f(z)\sigma_8(z)$ from earlier data release of SDSS BOSS galaxy sample \cite{Reid:2012sw,Chuang:2013hya,Wang:2014qoa,Anderson:2013zyy,Beutler:2013yhm,Chuang:2013wga,Samushia:2013yga}. There are some cosmological parameters, e.g. $\Omega_bh^2$ (the physical baryon fraction) and $n_s$ (the scalar index of the power law primordial fluctuation), not well constrained by galaxy clustering analysis. We usually use priors adopted from CMB measurements or fix those to the best fit values obtained from CMB while doing Markov Chain Monte Carlo (MCMC) analysis. There would be some concern of missing weak degeneracies between these parameters and those measured. These could lead to incorrect constraints if models with very different predictions are tested, or double-counting when combining with CMB measurements. One might add some systematics error budget to be safe from the potential bias (e.g., see \cite{Anderson:2013zyy}). An alternative approach is to use a very wide priors, e.g. 5 or 10 $\sigma$ flat priors from CMB, to minimize the potential systematics bias from priors (e.g., see \cite{Chuang:2010dv,Chuang:2011fy}). However, the approach would obtain weaker constraints due to the wide priors. In this study, we test the ways in which LSS constraints are combined with CMB data, focussing on the information content, and the priors used when analysing LSS data. Since CMB data can be summarized with few parameters (e.g., see \cite{Wang:2007mza}), we use the joint data set from Planck and BOSS to extract the cosmological constraints without fixing dark energy models. By combining the CMB data and the BOSS data in the upstream of the data analysis to constrain the cosmological constraints, we call our method "double-probe analysis". Our companion paper, \cite{Chuang16}, constrains geometric and growth information from the BOSS data alone independent of the CMB data, thereby dubbed "single-probe", and combines with the CMB data in the downstream of the analysis. Note that we assume there is no early time dark energy or dark energy clustering in this study. $\Omega_bh^2$ and $n_s$ will be well constrained by CMB so that we will obtain the cosmological constraints without concerning the problem of priors. The only input parameter which is not well constrained by our analysis is the galaxy bias on which is applied a wide flat prior. In principle, our methodology extract the cosmological constraints from the joint data set with the optimal way since we do not need to include the uncertainty introduced by the priors. In addition to constraining dark energy model parameters, we extend our study to constrain neutrino masses. High energy physics experiments provides with the squared of mass differences between neutrino species from oscillation neutrino experiments. Latest results are $\Delta m^2_{21}=7.53\pm 0.18 \times 10^{-5} eV^2$ and $\Delta m^2_{32}=2.44\pm 0.06 \times 10^{-3} eV^2$ for the normal hierarchy ($m_3\gg m_2 \simeq m_1$) and $\Delta m^2_{32}=2.52\pm 0.07 \times 10^{-3} eV^2$ for the inverted mass hierarchy ($m_3\ll m_2 \simeq m_1$) (\citealt{1674-1137-38-9-090001}). Cosmology shows as a unique tool for the measurement of the sum of neutrino masses $\Sigma m_{\nu}$, since this quantity affects the expansion rate and the way structures form and evolve. $\Sigma m_\nu$ estimations from galaxy clustering has been widely studied theoretically (see \citealt{Hu:1997mj,Lesgourgues:2006nd} for a review) and with different samples such as WiggleZ (see \citealt{Riemer-Sorensen:2013jsa,Cuesta:2015iho}) or SDSS data (see \citealt{Aubourg:2014yra,Beutler:2014yhv,Reid:2009nq,Thomas:2009ae,Zhao:2012xw}). At late times, massive neutrinos can damp the formation of cosmic structure on small scales due to the free-streaming effect \citep{Dolgov:2002wy}. Existing in the form of radiation in the early Universe, neutrinos shift the epoch of the matter-radiation equality thus changing the shape of the cosmic microwave background (CMB) angular power spectrum. They affect CMB via the so called Early Integrated Sachs Wolfe Effect and they influence gravitational lensing measurements (e.g., see \citealt{Lesgourgues:2005yv}). Recent publications have attempted to constrain $\Sigma m_\nu$ , imposing upper limits \citep{Seljak:2006bg,Hinshaw:2008kr,Dunkley:2008ie,Reid:2009nq,Komatsu:2010fb,Saito3,Tereno:2008mm,Gong:2008pg,Ichiki:2008ye,Li:2008vf,dePutter:2012sh,Xia:2012na,Sanchez:2012sg,Giusarma:2013pmn} and some hints of lower limits using cluster abundance results \citep{Ade:2013lmv,Battye:2013xqa,Wyman:2013lza,Burenin:2013wg,Rozo:2013hha}. We measure the sum of neutrino mass using different methodologies including double probe analysis (introduced in this study), the full-likelihood analysis, and single probe analysis (\citealt{Chuang16}; companion paper). This paper is organized as follows. In Section \ref{sec:data}, we introduce the Planck data, 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 joint CMB and galaxy clustering analysis. In Section \ref{sec:results}, we present our double-probe cosmological measurements. In Section \ref{sec:use}, we demonstrate how to derive cosmological constraints from our measurements with some given dark energy model. In Section \ref{sec:full_run}, opposite to the manner of dark energy model independent method, we present the results from the full-likelihood analysis with fixing dark energy models. In Section \ref{sec:mnu}, we measure the sum of neutrino mass with different methodologies. We summarize and conclude in Section \ref{sec:conclusion}.

1607.03152

1607

1607.03963_arXiv.txt

The six-degree obliquity of the sun suggests that either an asymmetry was present in the solar system's formation environment, or an external torque has misaligned the angular momentum vectors of the sun and the planets. However, the exact origin of this obliquity remains an open question. \citet{BB16a} have recently shown that the physical alignment of distant Kuiper Belt orbits can be explained by a $5-20\,m_{\oplus}$ planet on a distant, eccentric, and inclined orbit, with an approximate perihelion distance of $\sim250\,$AU. Using an analytic model for secular interactions between Planet Nine and the remaining giant planets, here we show that a planet with similar parameters can naturally generate the observed obliquity as well as the specific pole position of the sun's spin axis, from a nearly aligned initial state. Thus, Planet Nine offers a testable explanation for the otherwise mysterious spin-orbit misalignment of the solar system.

\label{sect1} The axis of rotation of the sun is offset by six degrees from the invariable plane of the solar system \citep{SouSou12}. In contrast, planetary orbits have an RMS inclination slightly smaller than one degree\footnote{An exception to the observed orbital coplanarity of the planets is Mercury, whose inclination is subject to chaotic evolution \citep{Laskar94, BatyginHolmanMorby15}}, rendering the solar obliquity a considerable outlier. The origin of this misalignment between the sun's rotation axis and the angular momentum vector of the solar system has been recognized as a a longstanding question \citep{Kuip56, Trem91, heller93}, and remains elusive to this day. With the advent of extensive exoplanetary observations, it has become apparent that significant spin-orbit misalignments are common, at least among transiting systems for which the stellar obliquity can be determined using the Rossiter-McLaughlin effect \citep{Rossiter, McLaugh}. Numerous such observations of planetary systems hosting hot Jupiters have revealed spin-orbit misalignments spanning tens of degrees \citep{Hebrard, Winn, Albrecht}, even including observations of retrograde planets \citep{Narita, Winn09, Bayliss, Winn11}. Thus, when viewed in the extrasolar context, the solar system seems hardly misaligned. However, within the framework of the nebular hypothesis, the expectation for the offset between the angular momentum vectors of the planets and sun is to be negligible, unless a specific physical mechanism induces a misalignment. Furthermore, the significance of the solar obliquity is supported by the contrasting relative coplanarity of the planets. Because there is no directly observed stellar companion to the sun (or any other known gravitational influence capable of providing an external torque on the solar system sufficient to produce a six-degree misalignment over its multi-billion-year lifetime \citealt{heller93}), virtually all explanations for the solar obliquity thus far have invoked mechanisms inherent to the nebular stage of evolution. In particular, interactions between the magnetosphere of a young star and its protostellar disk can potentially lead to a wide range of stellar obliquities while leaving the coplanarity of the tilted disk intact \citep{Lai11}. Yet another possible mechanism by which the solar obliquity could be attained in the absence of external torque is an initial asymmetry in the mass distribution of the protostellar core. Accordingly, asymmetric infall of turbulent protosolar material has been proposed as a mechanism for the sun to have acquired an axial tilt upon formation \citep{Bate, Fielding15}. However, the capacity of these mechanisms to overcome the re-aligning effects of accretion, as well as gravitational and magnetic coupling, remains an open question \citep{Lai11, Spalding14, Spalding15}. In principle, solar obliquity could have been excited through a temporary, extrinsic gravitational torque early in the solar system's lifetime. That is, an encounter with a passing star or molecular cloud could have tilted the disc or planets with respect to the sun \citep{heller93,Adams10}. Alternatively, the sun may have had a primordial stellar companion, capable of early star-disc misalignment \citep{Bat12,Spalding14,Lai14}. To this end, ALMA observations of misaligned disks in stellar binaries \citep{Jensen, Williams} have provided evidence for the feasibility of this effect. Although individually sensible, a general qualitative drawback of all of the above mechanisms is that they are only testable when applied to the extrasolar population of planets, and it is difficult to discern which (if any) of the aforementioned processes operated in our solar system. Recently, \citet{BB16a} determined that the spatial clustering of the orbits of Kuiper Belt objects with semi-major axis $a\gtrsim250\,$AU can be understood if the solar system hosts an additional $m_{9}=5-20\,m_{\oplus}$ planet on a distant, eccentric orbit. Here, we refer to this object as Planet Nine. The orbital parameters of this planet reside somewhere along a swath of parameter space spanning hundreds of AU in semi-major axis, significant eccentricity, and tens of degrees of inclination, with a perihelion distance of roughly $q_{9}\sim250\,$AU \citep{BB16b}. In this work, we explore the possibility that this distant, planetary-mass body is fully or partially responsible for the peculiar spin axis of the sun. Induction of solar obliquity of some magnitude is an inescapable consequence of the existence of Planet Nine. That is, the effect of a distant perturber residing on an inclined orbit is to exert a mean-field torque on the remaining planets of the solar system, over a timespan of $\sim4.5$ Gyr. In this manner, the gravitational influence of Planet Nine induces precession of the angular momentum vectors of the sun and planets about the total angular momentum vector of the solar system. Provided that angular momentum exchange between the solar spin axis and the planetary orbits occurs on a much longer timescale, this process leads to a differential misalignment of the sun and planets. Below, we quantify this mechanism with an eye towards explaining the tilt of the solar spin axis with respect to the orbital angular momentum vector of the planets. The paper is organized as follows. Section (\ref{sect2}) describes the dynamical model. We report our findings in section (\ref{sect3}). We conclude and discuss our results in section (\ref{sect4}). Throughout the manuscript, we adopt the following notation. Similarly named quantities (e.g. $a$, $e$, $i$) related to Planet Nine are denoted with a subscript ``9", whereas those corresponding to the Sun's angular momentum vector in the inertial frame are denoted with a tilde. Solar quantities measured with respect to the solar system's invariable plane are given the subscript $\odot$. \begin{figure} \includegraphics[width=0.48\textwidth]{PhatGeo.pdf} \caption{Geometric setup of the dynamical model. The orbits of the planets are treated as gravitationally interacting rings. All planets except Planet Nine are assumed to have circular, mutually coplanar orbits, and are represented as a single inner massive wire. The sun is shown as a yellow sphere, and elements are not to scale. Black, grey, and dotted lines are respectively above, on, and below the inertial reference plane. The pink arrows demonstrate the precession direction of the angular momentum vector of the inner orbit, $L_{\text{in}}$, around the total angular momentum vector of the solar system $L_{\text{total}}$. Red and blue arrows represent the differential change in longitudes of ascending node of the orbits and inclination, respectively. Although not shown in the figure, the tilting of the oblate sun is modeled as the tilting of an inner test ring. Over the course of 4.5 billion years, differential precession of the orbits induces a several-degree solar obliquity with respect to the final plane of the planets.} \label{setup} \end{figure}

\label{sect4} Applying the well-established analytic methods of secular theory, we have demonstrated that a solar obliquity of order several degrees is an expected observable effect of Planet Nine. Moreover, for a range of masses and orbits of Planet Nine that are broadly consistent with those predicted by \cite{BB16a, BB16b}, Planet Nine is capable of reproducing the observed solar obliquity of $6$ degrees, from a nearly coplanar configuration. The existence of Planet Nine therefore provides a tangible explanation for the spin-orbit misalignment of the solar system. Within the context of the Planet Nine hypothesis, a strictly null tilt of the solar spin-axis is disallowed. However, as already mentioned above, in addition to the long-term gravitational torques exerted by Planet Nine, numerous other physical processes are thought to generate stellar obliquities (see e.g. \citealt{CridaBat} and the references therein). A related question then, concerns the role of Planet Nine with respect to every other plausible misalignment mechanism. Within the context of our model, this question is informed by the present-day offset between the longitudes of ascending node of Planet Nine and the Sun, $\Delta\,\Omega$. Particularly, if we assume that the solar system formed in a configuration that was strictly co-planar with the sun's equator, the observable combination of the parameters $m_9,a_9,e_9,i_9$ maps onto a unique value of the observable parameter $\Delta\,\Omega$. Importantly, our calculations suggest that if the orbit of Planet Nine resides in approximately the same plane as the orbits of the $a\gtrsim250\,$AU Kuiper belt objects (which inform the existence of Planet Nine in the first place), then the inferred range of $\Delta\,\Omega$ and Planet Nine's expected orbital elements are incompatible with an exactly co-linear initial state of the solar spin axis. Instead, backwards integrations of the equations of motion suggest that a primordial spin-orbit misalignment of the same order as the RMS spread of the planetary inclination ($i\sim1\,\deg$) is consistent with the likely orbital configuration of Planet Nine. In either case, our results contextualize the primordial solar obliquity within the emerging extrasolar trend of small spin-orbit misalignments in flat planetary systems \citep{MorWin}, and bring the computed value closer to the expectations of the nebular hypothesis. However, we note that at present, the range of unconstrained parameters also allows for evolutionary sequences in which Planet Nine's contribution does not play a dominant role in exciting the solar obliquity. The integrable nature of the calculations performed in this work imply that observational characterization of Planet Nine's orbit will not only verify the expansion of the solar system's planetary album, but will yield remarkable new insights into the state of the solar system, at the time of its formation. That is, if Planet Nine is discovered in a configuration that contradicts a strictly aligned initial condition of the solar spin axis and planetary angular momentum, calculations of the type performed herein can be used to deduce the true primordial obliquity of the sun. In turn, this information can potentially constrain the mode of magnetospheric interactions between the young sun and the solar nebula \citep{Konigl91, Lai11, Spalding15}, as well as place meaningful limits on the existence of a putative primordial stellar companion of the sun \citep{Bat12,XiangGruessPapaloizou}. Finally, this work provides not only a crude test of the likely parameters of Planet Nine, but also a test of the viability of the Planet Nine hypothesis. By definition, Planet Nine is hypothesized to be a planet having parameters sufficient to induce the observed orbital clustering of Kuiper belt objects with semi-major axis $a>250$ AU \citep{BB16a}. According to this definition, Planet Nine must occupy a narrow swath in $a-e$ space such that $q_{9}\sim250\,$AU, and its mass must reside in the approximate range $m_9=5-20\,m_{\oplus}$. If Planet Nine were found to induce a solar obliquity significantly higher than the observed value, the Planet Nine hypothesis could be readily rejected. Instead, here we have demonstrated that, over the lifetime of the solar system, Planet Nine typically excites a solar obliquity that is similar to what is observed, giving additional credence to the Planet Nine hypothesis.

1607.03963

1607

1607.08240_arXiv.txt

The discovery of extraterrestrial neutrinos in the $\sim$ 30 TeV -- PeV energy range by IceCube provides new constraints on high energy astrophysics. An important background to the signal are the prompt neutrinos which originate from the decay of charm hadrons produced by high energy cosmic-ray particles interacting in the Earth's atmosphere. It is conventional to use pQCD calculations of charm hadroproduction based on gluon splitting $g \to c \bar c$ alone. However, QCD predicts an additional ``intrinsic" component of the heavy quark distribution which arises from diagrams where heavy quarks are multiply connected to the proton's valence quarks. We estimate the prompt neutrino spectrum due to intrinsic charm. We find that the atmospheric prompt neutrino flux from intrinsic charm is comparable to the pQCD contribution once we normalize the intrinsic charm differential cross sections to the ISR and the LEBC-MPS collaboration data. In future, IceCube will constrain the intrinsic charm content of the proton and will contribute to one of the major uncertainties in high energy physics phenomenology.

\label{sec:Introduction} Astrophysical neutrinos ($\equiv$ $\nu + \bar{\nu}$) discovered by IceCube provide new insights on profound astrophysics and particle physics questions\,\cite{Aartsen:2013bka,Aartsen:2013jdh,Aartsen:2014gkd,Aartsen:2015knd,Aartsen:2015rwa,Aartsen:2015zva,Aartsen:2016xlq}. Many astrophysical models have been proposed to explain these events\,\cite{Laha:2013eev,Murase:2013ffa,Anchordoqui:2013dnh,Kistler:2013my,Ahlers:2013xia,Anchordoqui:2014pca,Fang:2014qva,Kashiyama:2014rza,Baerwald:2014zga,Ando:2015bva,Emig:2015dma,Tamborra:2015fzv,Kistler:2015ywn,Kistler:2015oae,Kistler:2016ask,Moharana:2016mkl,Halzen:2016uaj,Pagliaroli:2016lgg,Anchordoqui:2016dcp,Dey:2016psn,Fang:2016hyv,Palladino:2016zoe,Senno:2015tsn} and to constrain various processes\,\cite{Bustamante:2015waa,Arguelles:2015dca,Vincent:2016nut,Kopp:2015bfa,Dasgupta:2012bd,Murase:2015gea,Rott:2014kfa,Esmaili:2014rma,Bhattacharya:2014vwa,Bhattacharya:2014yha,Ng:2014pca,Dutta:2015dka,Ko:2015nma,Aisati:2015ova,Boucenna:2015tra,Roland:2015yoa,Cherry:2016jol,Dev:2016uxj,Dev:2016qbd,Ema:2016zzu,DiBari:2016guw,Dey:2016sht}. These have spurred development of new signatures such as the through-going tracks caused by $\tau$ leptons\,\cite{Kistler:2016ask} and the echo technique\,\cite{Li:2016kra}. IceCube has detected an excess of neutrinos over the atmospheric neutrino background; however: how well do we know the background? The contribution of conventional atmospheric neutrinos, produced from the decays of $\pi$'s and $K$'s, is known to $\lesssim$ 10\% precision\,\cite{Honda:2006qj,Barr:2004br}. The major background uncertainty comes from $pp \rightarrow c X$, which results in prompt neutrinos produced from the decay of charm hadrons\,\cite{Bugaev:1989we,Gondolo:1995fq,Pasquali:1998ji,Gelmini:1999ve,Gelmini:1999xq,Candia:2003ay,Martin:2003us,Berghaus:2007hp,Enberg:2008te,Gaisser:2013ira,AhnICRC2013,Lipari:2013taa,Engel:2015dxa,Bhattacharya:2015jpa,Gauld:2015kvh,Gauld:2015yia,Garzelli:2015psa,Fedynitch:2015zma,Fedynitch:2015zbe,bausimpact,Bhattacharya:2016jce,Gaisser:2016obt,Halzen:2016pwl,Halzen:2016thi}. The flavor ratio of prompt neutrinos is $\nu_e : \nu_\mu : \nu_\tau \approx 1:1:0.1$, and $\nu : \bar{\nu} = 1 : 1$. Most calculations of the prompt neutrino spectrum from charm hadroproduction are based on perturbative QCD (pQCD) gluon splitting $g \to c\bar c$ alone\,\cite{Pasquali:1998ji,Gelmini:1999ve,Gelmini:1999xq,Candia:2003ay,Martin:2003us,Berghaus:2007hp,Enberg:2008te,Gaisser:2013ira,AhnICRC2013,Engel:2015dxa,Bhattacharya:2015jpa,Gauld:2015yia,Gauld:2015kvh,Garzelli:2015psa,Fedynitch:2015zma,Fedynitch:2015zbe,bausimpact,Bhattacharya:2016jce,Fedynitch:2016nup}. Inclusion of nonperturbative effects, for e.g., intrinsic charm, have received much less consideration\,\cite{Bugaev:1989we,Gondolo:1995fq}. Recently attention has been drawn to this issue by Refs.\,\cite{Halzen:2016pwl,Halzen:2016thi}. We calculate this component by using improved theoretical and experimental input. The important distinction between intrinsic charm and gluon splitting is that intrinsic charm uses the incoming proton energy much more efficiently due to its harder $d\sigma/ dx_F$ distribution. Inclusion of the nonperturbative effects are important since the amount of intrinsic charm is an important uncertainty in QCD simulations. Due to its inherent non-perturbative nature, it has not yet been calculated from first principles, and thus its normalization must be inferred from experiment. Experiments have not yet decisively measured the normalization of intrinsic charm in the proton, which typically dominates the differential cross section at $x_F \gtrsim$ 0.4. Various experimental techniques have been suggested for measuring atmospheric prompt neutrinos\,\cite{Gelmini:2002sw,Beacom:2004jb,Gandhi:2005at,Desiati:2010wt}. These studies illustrate how measurements can constrain the underlying QCD mechanism in regions of the parameter space where it is difficult to obtain constraints from colliders\,\cite{bausimpact}. IceCube compares the prompt neutrino spectrum derived by Enberg, Reno and Sarcevic (with modifications by Gaisser) (ERS w/G)\,\cite{Enberg:2008te,Gaisser:2013ira,Gaisser:2016obt} with their data. The present upper limits on the prompt neutrinos are near the nominal predictions\,\cite{Aartsen:2014muf,Aartsen:2015zva}. An additional contribution to the prompt neutrino spectrum can change the interpretation of the astrophysical neutrinos. In this paper, we calculate the prompt neutrino contribution from intrinsic charm after normalizing the differential cross section to the ISR and the LEBC-MPS collaboration data\,\cite{Chauvat:1987kb,Ammar:1988ta}. This contribution must be added to the $g\rightarrow c\bar{c}$ contribution to obtain the total atmospheric prompt neutrino spectrum. We show that the prompt neutrino flux from intrinsic charm is comparable to the pQCD contribution. The inclusion of this component as a background in the atmospheric neutrino flux can have important implications on the flux and spectral shape of the ``IceCube excess neutrinos". IceCube sensitivity is in the ball-park of the neutrino flux due to intrinsic charm\,\cite{Brodsky:1980pb,Brodsky:2015fna}. We emphasize that IceCube can test these differential cross sections which have proven to be difficult to measure in colliders. This synergy between IceCube and the collider searches\,\cite{Rostami:2015iva,Bailas:2015jlc,Boettcher:2015sqn,Lipatov:2016feu} can constrain the normalization of the intrinsic charm contribution and solve a $\sim$36 year old puzzle in QCD.

\label{sec:conclusions} The landmark discovery of astrophysical neutrinos by IceCube opens up a new era. Due to the atmospheric veto employed by IceCube, any atmospheric neutrino spectrum shows an up v/s down asymmetry. The excess of neutrinos unveiled by IceCube is isotropic implying the astrophysical origin of these events. Careful consideration of the atmospheric neutrino background will impact the astrophysical neutrino flux interpretation. The neutrino backgrounds considered so far by IceCube are the conventional atmospheric and prompt neutrinos predicted by $g\rightarrow c \bar{c}$. Intrinsic charm, rigorously predicted by QCD, has strong theoretical justification and experimental indications. We find that this often neglected component can be as large as the component due to pQCD $g \rightarrow c \bar{c}$ without violating any direct experimental constraints. This has important implications in interpreting the astrophysical neutrino flux, and inferring the atmospheric prompt neutrino component. We present our calculation of the neutrino flux due to intrinsic charm in Fig.\,\ref{fig:Comparison prompt spectrum} after normalizing to the ISR and the LEBC-MPS collaboration data. We show the atmospheric prompt neutrino flux due to three different scenarios. The atmospheric prompt neutrino flux due to intrinsic charm is comparable to that due to pQCD. Our calculation is lower than Refs.\,\cite{Halzen:2016pwl,Halzen:2016thi} as we use improved theoretical and experimental input. The measurement of atmospheric $\nu_e + \bar{\nu}_e$ at higher energies is the most promising channel to discover prompt neutrinos and constrain the intrinsic charm of the proton (Fig.\,\ref{fig:Atmospheric neutrino measurement} left). The comparison of the total atmospheric flux with the $\nu_\mu + \bar{\nu}_\mu$ data, including the intrinsic charm contribution, is shown in Fig.\,\ref{fig:Atmospheric neutrino measurement} (right). The total atmospheric neutrino flux including intrinsic charm can dominate the pQCD contribution at energies $\gtrsim$ 200 TeV and $\gtrsim$ 2 PeV for $\nu_e + \bar{\nu}_e$ and $\nu_\mu + \bar{\nu}_\mu$ respectively. The conventional atmospheric $\nu_e + \bar{\nu}_e$ flux is lower, implying that the prompt component is more visible in this channel. We estimate that a measurement of the atmospheric $\nu_e + \bar{\nu}_e$ flux at $\sim$ 200 TeV at $\sim$ 50\% accuracy will cleanly distinguish between the pQCD and intrinsic charm component. The current upper limit on prompt neutrinos is 1.06 times the ERS w/G flux. The neutrino flux due to intrinsic charm is at the same level as the ERS w/G flux implying that IceCube can constrain intrinsic charm of the proton. This shows that IceCube can constrain QCD predictions in regions of parameter space which have been difficult to constrain in colliders for decades. The multi-pronged approach consisting of IceCube data, collider physics, and global analysis will help us constrain the intrinsic charm of the proton, and solve a $\sim$ 36 year old problem in QCD. The use of neutrinos, a weakly interacting particle, to constrain the strong interactions will also highlight the importance of cross disciplinary searches in physics. \vspace{-0.5 cm}

1607.08240

1607

1607.01682_arXiv.txt

We derive an exact formulation of the multivariate integral representing the single-visit obscurational and photometric completeness joint probability density function for arbitrary distributions for planetary parameters. We present a derivation of the region of nonzero values of this function which extends previous work, and discuss time and computational complexity costs and benefits of the method. We present a working implementation, and demonstrate excellent agreement between this approach and Monte Carlo simulation results.

Obscurational completeness was introduced by \citet{BrownOC} as a necessary, but not sufficient, condition for detection of an exoplanet. Assuming distributions for semimajor axis and eccentricity of planetary orbits, Brown defined obscurational completeness as the probability of a planet falling outside a telescope's central obscuration, thus becoming potentially observable. This concept was expanded \citep{BrownSVPOC} to include selection effects due to photometric restrictions on exoplanet observability introduced by telescope optics. Completeness can be extended to indirect exoplanet detection methods like reflex astrometry \citep{BrownRA} and account for successive observations \citep{BrownNew}. Benefits of completeness studies include realistic expectations of mission outcomes based on objective terms for search power and a scientific metric to inform and optimize mission designs \citep{BrownPhO}. These studies have been used in mission analysis and design for a variety of proposed exoplanet observatories \citep{BrownRA, LindlerTPF, SavranskyDRM, SavranskyPFM, SavranskySMD, Stark2014, Brown2015, Stark2015}. Single-visit completeness is determined by the assumption that an exoplanet is observable if its angular separation from its star is greater than the observatory's inner working angle (IWA) and less than the observatory's outer working angle (OWA) while also being illuminated such that the difference in brightness between the star and planet ($\Delta$mag) is below a limiting value ($\Delta\mathrm{mag}_{0}$). The IWA and OWA represent the minimum and maximum angular separation of the field of view. $\Delta\mathrm{mag}_{0}$ is where unresolvable confusion between the planet signal and noise occur. For simple cases an analytic functional representation of completeness may be possible. With the exception of \citet{Agol}, previous approaches to calculating completeness have relied on Monte Carlo trials because of the complexity of the assumed distributions. Probability distributions are assumed for the orbital elements and physical properties necessary for determining separation and $\Delta\mathrm{mag}$. A large, equal number of samples (\citet{BrownSVPOC} used 100 million) is generated from each of the distributions. Successive function evaluations are made, including solving Kepler's equation iteratively, leading to the calculation of star-planet separation and $\Delta\mathrm{mag}$ for each set of samples. A two-dimensional histogram of these values is constructed which gives the expectation, or relative frequency of occurrence, of separation and $\Delta\mathrm{mag}$ for each bin of the two-dimensional histogram. Dividing the expectation values by the area of each bin gives the joint probability density function. Integrating with respect to separation and $\Delta\mathrm{mag}$ over this joint probability density function yields a cumulative density function which gives the completeness, or probability that an observatory with given $\Delta\mathrm{mag}_{0}$, IWA, and OWA, observing a specific star for the first time, will detect a planet belonging to the assumed population. The Monte Carlo trial approach of finding the expectation is analogous to numerical integration. Increasing the number of samples, $n$, in the Monte Carlo trial approach results in reduced error which goes as $O\left(n^{-1/2}\right)$ \citep{Davis}. To reduce the error by one decimal place, the number of samples must be increased by a factor of 100. For one-dimensional numerical integration, the simple Riemann sum, $O\left(n^{-1}\right)$; Newton-Cotes quadrature, better than $O\left(n^{-2}\right)$; or Gaussian quadrature all have better error performance for $n$ sample points than Monte Carlo integration. Multidimensional integrals may be numerically integrated using a composition of one-dimensional integrals or product rules which will also have better error performance than Monte Carlo integration. In terms of time complexity, the Monte Carlo trial approach requires sampling of quantities, solving Kepler's equation iteratively, additional function evaluations to get separation and $\Delta\mathrm{mag}$, and sorting these values into a two-dimensional histogram. All of these operations can be performed in polynomial time or better. Numerical integration algorithms require functional evaluations at sample points, determination of weights, multiplying the weights with functional evaluations, and summing these values. All of these operations can be performed in polynomial time. Increasing the dimension of an integral exponentially increases the number of samples required which increases the computational time. Multidimensional integrals higher than about dimension three are better computed using Monte Carlo integration due to the computational time. The bivariate distribution sampled by Monte Carlo trials is a function of non-independent variables. The completeness distribution must be sampled fully to find any one point of the joint probability distribution accurately. Because of the number of parameters involved and the potential wide range of values these parameters may take, full sampling requires a large number of Monte Carlo trials. Increasing numbers of Monte Carlo trials increases the computational time of accurately determining completeness. For exoplanet mission simulators, it is desirable to produce completeness values quickly without sacrificing accuracy. Numerical integration of lower dimensional integrals would require fewer sample points and give better error and computational time performance for a single point of the joint probability density function. We present a functional approach to determining single-visit completeness which avoids the undersampling problem inherent in the Monte Carlo simulation. This approach also allows for calculation of a single point of the completeness joint probability density function without simulation of the entire phase space. We begin by presenting the necessary assumptions which allow the description of completeness in functional terms. These functional expressions are made to be as general as possible. We then discuss computational considerations of this approach and provide comparisons to the Monte Carlo trial approach.

We have derived an analytical approach for finding the completeness joint probability density function which avoids the undersampling inherent in Monte Carlo approaches and allows computation of a single point of this joint probability density function without simulation of the entire phase space. This allows a quicker, more accurate computation of the double integral giving completeness because the function is evaluated only at points necessary for computing the integral compared to sampling the entire parameter space fully with Monte Carlo trials. This approach is dependent on the assumptions of closed Keplerian orbits; orbital poles distributed uniformly over a spherical volume with respect to the observer; planet-planet interactions neglected; very large distance to target star; and independence of the distributions of geometric albedo, planetary radius, phase angle, and orbital radius. We have shown good agreement between this approach and the Monte Carlo approach. This approach will allow researchers more accurate computation of single-visit photometric and obscurational completeness, which will improve the estimate of the number of extrasolar planets discovered by direct-imaging planet-finding mission simulations and better inform mission design.

1607.01682

1607

1607.01011_arXiv.txt

{ We present spatially-resolved Atacama Large Millimeter/sub-millimeter Array (ALMA) 870~$\mu$m dust continuum maps of six massive, compact, dusty star-forming galaxies (SFGs) at $z\sim2.5$. These galaxies are selected for their small rest-frame optical sizes ($r_{\rm e, F160W}\sim1.6$~kpc) and high stellar-mass densities that suggest that they are direct progenitors of compact quiescent galaxies at $z\sim2$. The deep observations yield high far-infrared (FIR) luminosities of L$_{\rm IR}=10^{12.3-12.8}$~L$_{\odot}$ and star formation rates (SFRs) of SFR~$=200-700$~\suny, consistent with those of typical star-forming ``main sequence'' galaxies. The high-spatial resolution (FWHM$\sim0\farcs12-0\farcs18$) ALMA and HST photometry are combined to construct deconvolved, mean radial profiles of their stellar mass and (UV+IR) SFR. We find that the dusty, nuclear IR-SFR overwhelmingly dominates the bolometric SFR up to $r\sim5$~kpc, by a factor of over 100$\times$ from the unobscured UV-SFR. Furthermore, the effective radius of the mean SFR profile ($r_{\rm e, SFR}\sim1$~kpc) is $\sim$30\% smaller than that of the stellar mass profile. The implied structural evolution, if such nuclear starburst last for the estimated gas depletion time of $\Delta t=\pm100$~Myr, is a 4$\times$ increase of the stellar mass density within the central 1~kpc and a 1.6$\times$ decrease of the half-mass radius. This structural evolution fully supports dissipation-driven, formation scenarios in which strong nuclear starbursts transform larger, star-forming progenitors into compact quiescent galaxies.}

\begin{figure*}[t] \centering \includegraphics[width=8.1cm,angle=0.]{fig1a.eps} \hspace{1cm} \includegraphics[width=8cm,angle=0.]{fig1b.eps} \caption{\label{selectiondiag} {\it Left:} SFR--mass diagram for galaxies in CANDELS GOODS-S at $2<z<3$. The grey-scale density bins map the location of the SFR-MS. The solid black and dashed blue lines depict the best fit and 2.5$\times$ and 5$\times$ limits above and below the SFR-MS. The blue circles depict the compact SFGs observed with ALMA. The subpanels in the bottom-left corner show the $5''\times5''$ ACS/WFC3 $zJH$ images of the ALMA galaxies. The red dashed line marks the threshold in sSFR (\lssfr$<-1$) used to identify quiescent galaxies (red circles). {\it Right:} mass--size distribution for the same galaxies as in the left panel. The dashed line marks the {\it compactness} threshold, log$(\Sigma_{1.5})=10.4~M_{\odot}$kpc$^{-1.5}$.} \end{figure*} The majority of SFGs follow a relatively tight, almost linear relation between SFR and stellar mass, usually referred to as the star-formation ``main sequence'' that seems to be in place since $z\sim5-6$ (SF-MS; e.g., \citealt{mainseq}; \citealt{whitaker12b}). The ubiquitous and tight SF-MS suggests that the majority of the stars are formed in a predominantly smooth, {\it secular} mode. Furthermore, there is also evidence that, despite their wide range of sizes and morphologies, most of the stars in SFGs are formed in disks which are growing from the inside out, thus increasing their sizes with cosmic time (\citealt{wuyts13}; \citealt{nelson13, nelson15}). The progressive structural growth in the SF-MS is consistent with the classic notion of galaxy formation in a $\Lambda$CDM Universe in which gas accreted from dark matter halos cools and forms new stars in disks with increasingly larger scale lengths with cosmic time (e.g., \citealt{fall80}; \citealt{mo98}). A challenge to this simplified picture are the small sizes ($r_{e}\sim1$~kpc) of the first massive quiescent galaxies at $z\gtrsim1.5-3$ (e.g., \citealt{vdw14} and references therein). On one hand, their small sizes might be the consequence of having smaller star-forming progenitors formed at earlier times when the Universe was more dense (i.e., more concentrated haloes and higher gas fractions). On the other hand, compact quiescent galaxies could form in strongly dissipative processes, triggered by mergers or interaction-driven disk instabilities that cause a substantial growth of the nuclear stellar density as a result of gas-rich starbursts (\citealt{hopkins08a}; \citealt{dekel09b}). Both scenarios imply the formation of compact SFGs as the last stage before quenching star formation, but the predictions differ on whether these compact SFGs would exhibit extended SFR profiles, driving the inside-out size growth, or compact star-forming regions (starbursts) triggered by the dissipative phase. Such compact SFGs have been identified in sizable numbers and their small stellar sizes, steep mass profiles and obscured SFR properties have been confirmed by multiple studies (e.g., \citealt{barro13, barro14a}; \citealt{dokkum15}). However, direct measurements of their spatial distribution of the star formation relative to the mass profile, needed to discriminate between the two formation scenarios discussed above, are still inconclusive. These measurements have proven very difficult because even spatially resolved UV and optical SFR indicators based on HST observations are significantly affected by the high dust obscuration, particularly in galaxy centers (\citealt{wuyts12}; \citealt{tacchella15}), and FIR observations, sensitive to ionizing radiation re-emitted by the dust, usually have very poor spatial resolution. Modern (sub-)millimeter/radio interferometers such as ALMA and JVLA have opened a new window into this regime and enable us to measure the dust emission with high sensitivity and similar spatial resolution as those from HST observations. Here, we exploit a joint analysis of the high spatial resolution HST/ACS and WFC3 and ALMA continuum imaging to simultaneously characterize the UV- and IR- SFR profiles and the stellar mass profiles of 6 compact SFGs at $z\sim2.5$. Throughout this paper, we quote magnitudes in the AB system, assume a \cite{chabrier} initial mass function (IMF), and adopt the following cosmological parameters: ($\Omega_{M}$,$\Omega_{\Lambda}$,$h$) = (0.3, 0.7, 0.7).

\begin{figure*} \centering \includegraphics[width=6.15cm,angle=0.]{fig5a.eps} \hspace{0.3cm} \includegraphics[width=5.5cm,angle=0.]{fig5b.eps} \includegraphics[width=5.5cm,angle=0.]{fig5c.eps} \caption{\label{evoltracks} {\it Left:} Evolution of the $\Sigma_{\rm M}$ profile (grey contour) of compact SFGs assuming that their $\Sigma_{\rm SFR}$ (blue contour) remains constant during $\Delta t=\pm100$~Myr (light-to-dark green lines). The magenta contour indicates the region below the ALMA detection limit. The magenta line shows a possible extended star-forming component (see \S~4.2). The red dashed line shows the mean stellar mass profile of compact quiescent galaxies at $z\sim2$. The cyan dashed line indicates the $\Sigma_{\rm SFR}$ of typical SFGs at $z\sim1$ from \citet{nelson15} (see also \citet{tacchella15} at $z\sim2$). The arrows indicate the $r_{\rm e}$. {\it Middle:} The blue contour and the red points indicate the loci of star-forming and quiescent galaxies at $z\sim2$, and the blue and red lines depict their best-fit scaling relations from \citet{vdw14}. The light-to-dark green circles indicate the size-mass evolution of the stellar mass profile in panel a. The black arrows show the direction of the structural evolution in the Vela simulations during the ``wet compaction'' phase. {\it Right:} Same as middle panel but showing the \sigone-mass evolution. The average relations for SFGs and quiescent galaxies at $z\sim2$ are from Barro et al. (2016b).} \end{figure*} In a simplified picture of galaxy growth, the average structural evolution of SFGs proceeds roughly along their well-defined scaling relations (blue arrows in Figure~\ref{evoltracks}; e.g., \citealt{dokkum15}; Barro et al. 2016b). In this picture, massive compact quiescent galaxies at $z\sim2$ would be descendants of smaller SFGs at higher-z that achieve such high stellar densities by continuously growing in stellar mass and size fueled by extended SFR profiles. Alternatively, these SFGs could deviate from the smooth track due to dissipative processes that would rapidly increase their concentrations and potentially decrease their half-mass radii in strong nuclear starbursts (\citealt{dekel13b}; \citealt{wellons15}). The {\it secular} and dissipation-driven scenarios are not mutually exclusive. However, we aim to understand whether the massive dense cores of compact quiescent galaxies are primarily formed in dissipative processes. The strong nuclear starbursts embedded in larger stellar mass profiles found in compact SFGs are indeed an excellent match to the dissipation-driven scenario. The light-to-dark green lines and circles in Figure~\ref{evoltracks} show the predicted change in the stellar mass profile and the evolutionary tracks in $r_{\rm e}$ and central mass density for compact SFGs due to star formation, assuming that their SFR profiles remain constant during $\Delta t=200$~Myr (approximately $t_{\rm dpl}$). The significant stellar mass growth within the inner $r\lesssim2$~kpc decreases the half-mass radius by $1.6\times$ from $r_{\rm e,mass}=1.9$ to $1.2$~kpc, while the central density within $r\leq1$~kpc increases by $\sim4\times$ from log(\sigone)~$=$~9.7 to 10.3~M$_{\odot}$~kpc$^{-2}$. If compact SFGs had more extended star formation at $r\gtrsim3$~kpc the evolution of \sigone~would be the same, while $r_{\rm e, mass}$ would decrease less (e.g., $\sim$7\% for the magenta line). These evolutionary tracks are very similar to the predictions of the Vela simulations during the ``wet compaction'' phase (black arrows; e.g., \citealt{zolotov15}; \citealt{tacchella16}) and contrast with the expected evolution for typical SF-MS galaxies, which have extended SFR profiles with $\sim100\times$ lower central $\Sigma_{\rm SFR}$ (cyan line in Figure~\ref{evoltracks}a) and thus favor a more gradual increase of \sigone~and a positive size evolution. The short depletion times of compact SFGs and the similarity with the mean stellar mass profile of quiescent galaxies at $z\sim2$ (red dashed line in Figure~5a) suggest that the nuclear starburst is unlikely to continue for more than a few hundred Myrs, either because no further gas is accreted into the galaxy center or because the dense stellar component stabilizes the gas to prevent further star formation and eventually leads to galaxy quenching. This scenario is consistent with previous results indicating that the formation of a dense core precedes the shut down of star formation (e.g., \citealt{cheung12}; \citealt{dokkum14}), and suggests that, at high redshift, both quenching and the dense cores are simultaneous consequences of enhanced periods of nuclear star formation that cause a rapid depletion of the gas reservoirs.

1607.01011

1607

1607.04919_arXiv.txt

We report the discovery of a high mass-ratio planet $q=0.012$, i.e., 13 times higher than the Jupiter/Sun ratio. The host mass is not presently measured but can be determined or strongly constrained from adaptive optics imaging. The planet was discovered in a small archival study of high-magnification events in pure-survey microlensing data, which was unbiased by the presence of anomalies. The fact that it was previously unnoticed may indicate that more such planets lie in archival data and could be discovered by similar systematic study. In order to understand the transition from predominantly survey+followup to predominately survey-only planet detections, we conduct the first analysis of these detections in the observational $(s,q)$ plane. Here $s$ is projected separation in units of the Einstein radius. We find some evidence that survey+followup is relatively more sensitive to planets near the Einstein ring, but that there is no statistical difference in sensitivity by mass ratio.

For the first decade of microlens planet detections, beginning with OGLE-2003-BLG-235Lb \citep{ob03235}, the great majority of detections required a combination of survey and followup data. This is a consequence of two effects. First, the survey coverage was generally too sparse to characterize the planetary anomalies in the detected events \citep{gouldloeb}. Second, thanks to aggressive alert capability, pioneered by the Optical Gravitational Lensing Experiment (OGLE) Early Warning System (EWS, \citealt{ews1,ews2}), it became possible to organize intensive followup of planet-sensitive events -- or even ongoing planetary anomalies -- and so obtain sufficient time resolution to detect and characterize planets. However, as surveys have become more powerful over the past decade, they have become increasingly capable of detecting planets without followup observations. That is, making use of larger cameras, the surveys are able to monitor fairly wide areas at cadences of up to several times per hour. While still substantially lower than followup observations of the handful of events that were monitored by followup groups, this is still adequate to detect most planets (provided that the anomalies occur when the survey is observing). Very simple reasoning given below, which is supported by detailed simulations \citep{Zhu:2014}, leads one to expect that the transition from survey+followup to survey-only mode implies a corresponding transition from planets detected primarily in high-magnification events via central and resonant caustics to planets primarily detected in lower magnification events via planetary caustics. High-magnification events are intrinsically sensitive to planets because they probe the so-called ``central caustic'' that lies close to (or overlays) the position of the host \citep{griest98}. Planets that are separated from the hosts by substantially more (less) than the Einstein radius generate one (two) other caustics that are typically much larger than the central caustic and thus have a higher cross section for anomalous deviations from a point-lens light curve due to a random source trajectory. However, for high-magnification events, the source is by definition passing close to the host and hence close to or over the central caustic. For planet-host separations that are comparable to the Einstein radius, the two sets of caustics merge into a single (and larger) ``resonant caustic'', which is even more likely to generate anomalous deviations of a high-magnification event. For many years, the Microlensing Follow Up Network ($\mu$FUN) employed a strategy based on this high planet sensitivity of high-magnification events. They made detailed analyses of alerts of ongoing events from the OGLE and the Microlensing Observations in Astrophysics (MOA) teams to predict high-magnification events and then mobilized followup observations over the predicted peak. \citet{gould10} showed that $\mu$FUN was able to get substantial data over peak for about 50\% of all identified events with maximum magnification $A_\max>200$, but that its success rate dropped off dramatically at lower magnification, i.e., even for $100<A_\max<200$. The reason for this drop off was fundamentally limited observing resources: there are twice as many events $A_\max>100$ compared to $A_\max>200$, and monitoring the full-width half-maximum requires twice as much observing time. Hence, observations grow quadratically with effective magnification cutoff. By contrast, because planetary caustics are typically much larger than central caustics, most planets detected in survey-only mode are expected to be from anomalies generated by the former, which occur primarily in garden-variety (rather than high-mag) events \citep{Zhu:2014}. For example, \citet{ob120406} detected a large planetary caustic in OGLE-2012-BLG-0406 based purely upon OGLE data, while \citet{moabin1} detected one in MOA-bin-1 based mostly on MOA data. In the latter case it would have been completely impossible to discover the planet by survey+followup mode because the ``primary event'' (due to the host) was so weak that it was never detected in the data. Nevertheless, there has been a steady stream of survey-only detections of planets in high-magnification events as well. The first of these was MOA-2007-BLG-192Lb, a magnification $A_\max> 200$ event, which required a combination of MOA and OGLE data \citep{mb07192}. The first planet detected by combining three surveys (MOA, OGLE, Wise), MOA-2011-BLG-322Lb, was also via a central caustic, although in this case the caustic was very large so that the magnification did not have to be extremely large $(A_\max\sim 20)$ \citep{mb11322}. Similarly, \citet{ob150954} detected a large central caustic due to the large planet OGLE-2015-BLG-0954Lb despite modest peak magnification of the underlying event $A_\max\sim 20$. This case was notable because high-cadence data from the Korea Microlensing Telescope Network (KMTNet) captured the caustic entrance despite the extremely short source self-crossing time, $t_*=16\,$min. There also exist 2 planets MOA-2008-BLG-379Lb \citep{Suzuki2014} and OGLE-2012-BLG-0724Lb \citep{Hirao2016} that were detected by the OGLE+MOA surveys through the high-magnification channel. KMTNet is still in the process of testing its reduction pipeline. Motivated by the above experience, the KMTNet team focused its tests on high-magnification events identified as such on the OGLE web page. In addition to exposing the reduction algorithms to a wide range of brightnesses, this testing has the added advantage that there is a high probability to find planets. Here we report on the first planet found by these tests from among the first seven high-mag events that were examined: OGLE-2016-BLG-(0261,0353,0471,0528,0572,0596,0612). These events were chosen to have model point-lens magnifications $A>20$ and modeled peak times $2457439<t_0<2457492$. The lower limit was set by the beginning of the KMTNet observing season and the upper limit was the time of the last OGLE update when the seven events were selected.

OGLE-2016-BLG-0596Lb is a very high mass-ratio $(q=0.0117)$ planet that lies projected very close to the Einstein ring $(s=1.075)$, which consequently generated a huge resonant caustic that required 16 days for the source to traverse. The underlying event was of quite high magnification $(A_{\rm point-lens}\sim 100)$, which led to pronounced features at peak. It therefore would seem to be extremely easy to discover. While the data set posted on the OGLE web site are adversely affected by the nearby variable, it is still the case that a free fit to these data leads to a solution qualitatively similar to the one presented here (except that it lacks a measurement of $\rho$). It is therefore striking that none of the automated programs nor active individual investigators that query this site noticed this event (or at least they did not alert the community to what they found as they do for a wide range of other events, many less interesting). This indicate the possibility that there may be many other planets ``hidden in plain sight'' in existing data. This is also supported by the planet discoveries MOA-2008-BLG-397Lb \citep{Suzuki2014}, OGLE-2008-BLG-355Lb \citep{Koshimoto2014}, and MOA-2010-BLG-353Lb \citep{Rattenbury2015}, for which the planetary signals were not noticed during the progress of events. These three characteristics, high-magnification (which is usually associated with survey+followup rather than survey-only mode), very high mass ratio, and apparent failure of both machine and by-eye recognition of the planetary perturbation, prompt us to address two questions. First, how do the real (as opposed to theoretical) planet sensitivities differ between survey-only and survey+followup modes. Second, why was this planet discovered only based on systematic analysis and what does this imply about the need for such systematic analysis of all events? \subsection{Summary of Microlens Planet Detections in the Observational $(s,q)$ Plane} Many papers contain figures that summarize microlensing planet detections in the physical plane of planet-mass versus projected separation (with the latter sometimes normalized by the snow line), e.g., Figure 1 of \citet{mb13605}. And there are many studies that show plots of {\it planet sensitivity} in the observational $(s,q)$ plane, (e.g., \citealt{gaudi02,gould10}). But to our knowledge, there are no published figures (or even figures shown at conferences) showing the census of microlensing planet discoveries on this plane. \begin{figure} \includegraphics[width=\columnwidth]{fig8.eps} \caption{ Cumulative microlensing planet detections by log mass ratio $\log(q)$, with top normalized and bottom unnormalized. Green shows the 18 planets that were detected and characterized by surveys, while magenta show the 27 planets that required significant followup observations for detection and/or characterization. Black is total. The green and magenta curves are not statistically distinguishable. } \label{fig:eight} \end{figure} Figure~\ref{fig:seven} illustrates the position of OGLE-2016-BLG-0596 (green pentagon) among the 44 previously published planets (or, to the extent we have such knowledge, submitted for publication). Discovered bodies are defined to be ``planets'' if their measured or best-estimated mass $m_p<13\,M_{\rm jup}$ and if they are known to orbit a more massive body\footnote{To facilitate comparison with future compilations, we list here the 45 planets used to construct this figure and those that follow. We compress, e.g., OGLE-2003-BLG-235Lb to OB03235 for compactness and only use ``b,c'' for multiple planets: OB03235, OB05071, OB05169, OB05390, MB06bin1, OB06109b, OB06109c, MB07192, MB07400, OB07349, OB07368, MB08310, MB08379, OB08092, OB08355, MB09266, MB09319, MB09387, MB10073, MB10328, MB10353, MB10477, MB11028, MB11262, MB11293, MB11322, OB110251, OB110265, OB120026b, OB120026c, OB120358, OB120406, OB120455, OB120563, OB120724, MB13220, MB13605, OB130102, OB130341, OB140124, OB141760, OB150051, OB150954, OB150966, OB160596.}. Planets are color-coded by discovery method: discovered by followup observations (blue), discovered (or discoverable) in survey-only observations but requiring followup for full-characterization (magenta), fully (or essentially fully) characterized by survey observations (green). The shapes of the symbols indicate the type of caustic that gave rise to the planetary perturbation: circles, squares, and triangles for planetary, central, and resonant caustics, respectively. In many cases, solutions with $(q,s)$ and $(q,1/s)$ yield almost equally good fits to the data \citep{griest98}. In these cases, the two solutions are shown as open symbols in order to diminish their individual visual ``weight'' relative to the filled symbols used when this degeneracy is broken. Hence, OGLE-2016-BLG-0596Lb would be a green filled triangle if it were not being singled out by making it larger pentagon. \begin{figure} \includegraphics[width=\columnwidth]{fig9.eps} \caption{ Cumulative microlensing planet detections by absolute value of the log projected separation (normalized to $\theta_\e$) $|\log(s)|$, with top normalized and bottom unnormalized. Colors are the same as in Figure~\ref{fig:eight}. The gap between the green (survey) and magenta (followup) curves has a 8.5\% probability of being random. If real, this indicates that followup observation have been relatively more sensitive to planets near the Einstein ring while surveys are more sensitive to those further from the Einstein ring. } \label{fig:nine} \end{figure} The most striking feature of this figure is that, in sharp contrast to the triangular appearance of high-magnification-event planet-sensitivity plots (e.g., \citealt{gould10}) and to ``double pronged'' low-magnification sensitivity plots (e.g., \citealt{gaudi02}), this detection plot looks basically like a cross, with a vertical band of detections near $\log s\sim 0$ and a horizontal band near $\log(q)\sim -2.5$. The part of this structure at high mass ratio $\log(q)>-2$ is easily explained: companions with high mass ratio are, a priori, most likely stars or brown dwarfs (BDs) and can only be claimed as ``planets'' if the host mass is known to be low. This in turn usually requires a measurement of the microlens parallax, which for ground based observations is much more likely if there is a large caustic and so $s\sim 1$. We note that there are 4 planet detections in the region $(\log(s)>+0.15, \log(q)<-3)$, while there is no detection in the opposite quadrant $(\log(s)<-0.15, \log(q)<-3)$. All the 4 planets derive from planetary caustics and 3 of them are pure survey detections: MOA-2011-BLG-028Lb\footnote{We note that this event's light curve does contain some followup data, but it is not essential for characterizing the planet.}, OGLE-2008-BLG-092Lb, and MOA-2013-BLG-605Lb \citep{mb11028,ob08092,mb13605}. The remaining planet, OGLE-2005-BLG-390Lb \citep{ob05390}, dates from an era when followup groups intensively monitored the wings of events, primarily due to the paucity of better targets. Thus we may expect that surveys will gradually fill in this quadrant. The difference in the detection rates between the quadrants with $\log(s)>+0.15$ and $\log(s)<-0.15$ can be explained by the difference in the size of the planetary caustics with $s<1$ and $s>1$. In the case of $s>1$, there exist a single planetary caustic. In the case of $s<1$, on the other hand, there exist two sets of planetary caustics and each one is smaller than the planetary caustic with $s>1$. As a result, the planetary caustic with $s>1$ has a larger cross section and thus higher sensitivity. Furthermore, smaller caustic size of planets with $s<1$ makes planetary signals tend to be heavily affected by finite-source effects, which diminish planetary signals, while signals of planets with $s>1$ can survive and show up in the wings of light curves. Actually, all 4 events with planet detections via planetary-caustic perturbations are involved with large source stars, i.e.\ giant and subgiant stars for which finite-source effects are important. \begin{figure} \includegraphics[width=\columnwidth]{fig10.eps} \caption{ Cumulative microlensing planet detections by year of discovery, with top normalized and bottom unnormalized. Colors are the same as in Figure~\ref{fig:eight}. Followup discoveries (magenta) have dropped off dramatically since 2013 } \label{fig:ten} \end{figure} Apart from this quadrant, it is not obvious that surveys are probing a different part of parameter space from the previously dominant survey+followup mode. To further investigate this, we show in Figures~\ref{fig:eight} and \ref{fig:nine} the cumulative distributions of planets by log mass ratio $\log(q)$ and (absolute value of) log separation $|\log(s)|$. In this case we distinguish only between events that could be fully characterized by survey observations (green) and those that required significant followup (including auto-followup by surveys). These distributions generally appear quite similar. For the mass ratio distribution, the greatest difference (0.259) is at $\log(q)=-2.319$, which is very typical (Kolmogorov-Smirnoff (KS) probability 40\%). The greatest difference for the separation distribution (0.334 at $|\log(s)|=0.124$) has a KS probability of 8.5\%. This may be indicative of a real difference. If so, the difference would be that pure-survey is relatively more efficient at finding widely separated lenses, which was already hinted at by inspection of the $(q,s)$ scatter plot. Finally, in Figure~\ref{fig:ten}, we show cumulative distributions by year of discovery. One might expect that with the massive ramp-up of surveys, survey-only discoveries would move strongly ahead of survey+followup. This expectation is confirmed in its sign but not its magnitude by Figure~\ref{fig:ten}. It shows that in (2014, 2015, 2016) there have been (2,2,1) and (0,1,0) discoveries by survey-only and survey+followup, respectively. This is certainly not a complete accounting, in part because 2016 has just begun and in part because historically there has been a considerable delay in microlensing planet publications for a variety of reasons. For example, of the 28 planets discovered prior to 2012, the number with delays (publication year minus discovery year) of $(0,1,\ldots,9)$ years was $N=(1,5,9,5,1,2,4,0,0,1)$. In the history of microlensing, there has been only one planet published during the discovery year, OGLE-2005-BLG-071Lb \citep{ob05071}. Hence, we will only get a full picture of this transition after a few years. \subsection{Challenges to the By-Eye and By-Machine discovery of OGLE-2016-BLG-0596} There are three interrelated reasons why OGLE-2016-BLG-0596 may have escaped notice as a potentially planetary event until the KMTNet data for this event were examined (for reasons unrelated to any apparent anomaly). First, it is relatively faint at peak. Second, it has a variable baseline. Third, it was not announced as a microlensing event until one day after the peak. As a general rule, high-magnification events are singled out for intensive followup observations only if they are still rising. When such intensive observations would have been % conducted, they would have immediately revealed the anomalous nature of the event, probably triggering additional observations. This is how many of the planets discovered by $\mu$FUN were found. While $\mu$FUN itself is now semi-dormant, its protocols are directly relevant here because what is of interest is whether there is prima facie evidence for a population of missed planets during past years, during most of which $\mu$FUN was active. Now, in fact, OGLE-2016-BLG-0596 met the criteria for an OGLE alert 24 hours previously, but no alert was issued because of caution due to the variable baseline. Nevertheless, even if such an alert had been issued, it would not have triggered any followup observations because (due to the anomaly) the event would have appeared to have already peaked at that time. Finally, the variability of the baseline may have influenced modelers and followup groups to discount the evident irregularities in the light curve near peak as being due to data artifacts. This could have been exacerbated by the faintness of the event, which increases both the formal error bars and the probability of centroiding errors (hence irregular photometry) due to bright blends. Both of these effects reduce the confidence of modelers that apparent anomalies in online ``quick look'' photometry are due to physical effects. It is nevertheless a fact that when the original OGLE data are modeled, they show a clear signal for a massive planet or low-mass BD, which would trigger a re-reduction of the data, such as the one we report here. We therefore conclude that while OGLE-2016-BLG-0596 has some near-unique features that increased the difficulty of recognizing it as a planetary event, such recognition was clearly feasible. Hence, we do indeed regard this event as prima facie evidence for more such events in archival data, particularly OGLE-IV data 2010-2015.

1607.04919

1607

1607.03825_arXiv.txt

The Rosetta probe around comet 67P/Churyumov-Gerasimenko (67P) reveals an anisotropic dust distribution of the inner coma with jet-like structures. The physical processes leading to jet formation are under debate, with most models for cometary activity focusing on localised emission sources, such as cliffs or terraced regions. Here we suggest, by correlating high-resolution simulations of the dust environment around 67P with observations, that the anisotropy and the background dust density of 67P originate from dust released across the entire sunlit surface of the nucleus rather than from few isolated sources. We trace back trajectories from coma regions with high local dust density in space to the non-spherical nucleus and identify two mechanisms of jet formation: areas with local concavity in either two dimensions or only one. Pits and craters are examples of the first case, the neck region of the bilobed nucleus of 67P for the latter one. The conjunction of multiple sources in addition to dust released from all other sunlit areas results in a high correlation coefficient ($\sim$0.8) of the predictions with observations during a complete diurnal rotation period of 67P.

Solar illumination drives the gas and dust emission from comets by sublimating the ice of the frozen dust-gas conglomerate. Embedded dust particles are accelerated by the expanding gas in the outer porous mantle of the comet and form the innermost coma around the nucleus (\cite{Huebner2006}). Starting with the 1P/Halley flyby, regions with increased dust intensity within several cometary radii have been imaged and connected to active and inactive surface areas (\cite{Whipple1982,Keller1994,Combi2012,BruckSyal2013,Belton2013}), or alternatively to the general shape of the nucleus (\cite{Crifo2002,Zakharov2009}). Specific surface features linked to jet formation include active pits (\cite{Vincent2015}), terraced regions (\cite{Farnham2013}), and cliffs (\cite{Vincent2015a}). In addition cometary {dust and gas emission} highly varies according to the rapidly changing temperature conditions driven by solar illumination (\cite{Ali-Lagoa2015a,Lara2015}). So far, the complex {dust} activity with temporally and spatially varying source patterns has precluded a detailed prediction of the dust distribution observed in the innermost coma of 67P besides for two specific cases (\cite{Kramer2015a,Marschall2015}). The conclusions drawn from these studies are limited by the lack of testing the theoretical prediction across a complete diurnal rotation of the nucleus. One obstacle for a predictive model has been the computational complexity to describe the cometary gas and dust in three dimensions emitted from a rotating cometary nucleus. Here, we report results for a high-resolution three-dimensional model of the dust environment around 67P predicted across the entire rotation period of the nucleus, which besides the time-dependent gas-dust interaction, also takes into account the rotation and detailed shape of the nucleus (\cite{Kramer2015,Kramer2015a}).

The residual differences between a strictly homogeneous emission model and observations could be explained by a varying surface composition, local variations in outgassing rates, the accuracy of the shape model, or {light scattering across different particle sizes (\cite{Fink2015}).} {Obtaining the absolute dust density requires to further constrain the dust particle acceleration and the gas pressure within meters from the surface by observations.} The homogeneous activity model already reproduces the {relative dust densities}, including jet-like features, at locations observed by Rosetta in the innermost coma around 67P. This indicates that the homogeneous model is a suitable assumption which highly correlates with the imaged coma light-intensities. The three-dimensional dust tracing analysis shows that photographed jets are often a conjunction of multiple aligned jet structures on top of a non-collimated, homogeneously released dust background. Adding a locally varying activity profile on top of the homogeneous model is in principle possible, but restricting the dust emission to only surface concavities reduces the agreement with observations. {A rapid acceleration of dust in the vicinity of the surface results in a coma mirroring the surface orography.} This supports the hypothesis that dust emanates from fractures and subsurface pores in the mantle (\cite{Vincent2015a}). { For a slower lift-off velocity, the rotation of the nucleus leads to a dispersion and drift of the dust, \cite{Kramer2015a}.} A possible implication of the underlying homogeneous emission model is that the overall surface ablation of the illuminated nucleus proceeds in a more uniform way (\cite{Cheng2013,Schulz2015}) compared to models of isolated dust emission sources.

1607.03825

1607

1607.01820_arXiv.txt

The unexpected energy spectrum of the positron/electron ratio is interpreted astrophysically, with a possible exception of the 100-300 GeV range. The data indicate that this ratio, after a decline between $0.5-8$ GeV, rises steadily with a trend towards saturation at 200-400GeV. These observations (except for the trend) appear to be in conflict with the diffusive shock acceleration (DSA) mechanism, operating in a \emph{single} supernova remnant (SNR) shock. We argue that $e^{+}/e^{-}$ ratio can still be explained by the DSA if positrons are accelerated in a \emph{subset} of SNR shocks which: (i) propagate in clumpy gas media, and (ii) are modified by accelerated CR \emph{protons}. The protons penetrate into the dense gas clumps upstream to produce positrons and, \emph{charge the clumps positively}. The induced electric field expels positrons into the upstream plasma where they are shock-accelerated. Since the shock is modified, these positrons develop a harder spectrum than that of the CR electrons accelerated in other SNRs. Mixing these populations explains the increase in the $e^{+}/e^{-}$ ratio at $E>8$ GeV. It decreases at $E<8$ GeV because of a subshock weakening which also results from the shock modification. Contrary to the expelled positrons, most of the antiprotons, electrons, and heavier nuclei, are left unaccelerated inside the clumps. Scenarios for the 100-300 GeV AMS-02 fraction exceeding the model prediction, including, but not limited to, possible dark matter contribution, are also discussed.

} Recent measurements of a positron/electron, $e^{+}/\left(e^{-}+e^{+}\right)$, excess in the $8-300$ GeV range by Pamela, Fermi-LAT and AMS-02 \citep{Pamela_Pos_09,Fermi_Pos_12,AMS_SepElPos14,AMS02_2014} has added fuel to the hotly contested race for elusive dark matter (DM) signatures in rapidly improving cosmic ray (CR) data \citep{Hooper09,HooperDM_AMS13,Berezinsky2015JPhCS}. Indeed, conventional acceleration schemes, even the most promising of them all, the diffusive shock acceleration (DSA), has not yet suggested any viable mechanism for the $e^{+}/\left(e^{-}+e^{+}\right)$ anomaly, free of tension with the antiproton spectra and other secondaries \citep{Mertsch2009PhRvL,Kachelriess11,CholisHooper2014PhRvD}. In addition to the surprising excess at high energies, the $e^{+}/\left(e^{-}+e^{+}\right)$ ratio has a distinct minimum at $\approx8$GeV which is not easier to explain making a minimum of assumptions. Both features appear at odds with the \emph{single source} DSA operation, which predicts similar rigidity ($R=$momentum/charge) spectra for all primary species. Moreover, there are also other well documented exceptions, namely the He$^{++}/p$ and $C/p$ ratios that both show a $\sim R^{0.1}$ growth, also seemingly inconsistent with the DSA \citep{Adriani11,AMS02He2015PhRvL,AMS02andBessPtoHe2015}. Less pronounced than $e^{+}/\left(e^{-}+e^{+}\right)$, but not less astonishing at first glance, these anomalies can be explained by the difference in charge to mass ratio \citep{MDSPamela12}. Other scenarios are possible but require additional assumptions, such as inhomogeneity of the SNR environment \citep{Ohira11,Ohira15,Drury12} or multiple sources with adjusted spectral indices (see \citep{Serpico15,Tomassetti2015ApJ,Ohira15} for a recent discussion). In fact, the mass to charge based explanation of the $\approx0.1$ difference in rigidity indices has been given only for the He/$p$ spectrum, while the C/$p$ was measured with sufficient accuracy only recently \citep{AMS02He2015PhRvL} and turned out to be identical to the He/$p$ rigidity spectrum. Thus, the mechanism suggested by \citep{MDSPamela12} predicted the $C/p$ spectrum since He and C have the same mass to charge ratio. If this injection mechanism is correct, the latest AMS-02 data speak against a direct carbon acceleration from grains \citep{Meyer97}. The mass to charge selectivity of the DSA which works for He/$p$ and C/$p$ does not apply to the $e^{+}/e^{-}$ fraction. It, therefore, seems logical to look for a possible \emph{charge-sign} dependence of the SNR-DSA production of CRs, including the $e^{+}/\left(e^{-}+e^{+}\right)$ anomaly. We call it 'anomaly' rather than 'excess' (also encountered in the literature) since the ratio rises with the particle energy only at $E>8$GeV. Below this energy, it \emph{declines, }thus creating a \emph{deficit}. The decline, the rise and the clear minimum between them (at $8$ GeV) are all pivotal to the mechanism proposed here. These aspects are intrinsic to a \emph{single-source} mechanism proposed, revealing unique characteristics of the accelerator. By contrast, assuming two or more independent positron contributions to the spectrum (as, e.g., in refs.\citep{ErlykWolf2013APh,Mertsch2014PhRvD,Cowsik2014ApJ}), one fits the nonmonotonic positron fraction, but with no constraints on the underlying acceleration mechanisms. The position of the minimum in the positron fraction is then coincidental, and the fit does not add credibility to the model predictions for the higher energy data points yet to come. We will return to this point in the Discussion section. A vast majority of conventional scenarios for the $e^{+}/\left(e^{-}+e^{+}\right)$ excess (including the present one) invoke secondary positrons. They are produced by galactic CR protons colliding with an ambient gas near an SNR accelerator, e.g. \citep{Fujita2009PhRvD}, elsewhere in the galaxy, e.g.,\citep{Waxman2013PhRvL,Cowsik2014ApJ}, or are immediately involved in the SNR shock acceleration, \citep{Blasi2009PhRvL,Mertsch2014PhRvD,CholisHooper2014PhRvD}. Some of these scenarios face the unmatched antiprotons and other secondaries in the data, as discussed, e.g., in \citep{Kachelriess11,VladimirMoskPamela11,CholisHooper2014PhRvD}. Improvements along these lines have recently been achieved by using Monte Carlo $pp$ collision event generators, e.g. \citep{Kohri2016PTEP}. However, improved cross sections of $pp$ collisions do not shed light on the \emph{physics} of $e^{+}/\left(e^{-}+e^{+}\right)$ anomaly, particularly the minimum at 8 GeV. This spectrum complexity hints at richer physics than a mere production of secondary $e^{+}$ and $\bar{p}$ power-law spectra from the primary CR power-law. We propose and investigate the idea that the physics of the $e^{+}/\left(e^{-}+e^{+}\right)$ fraction and, by implication, that of the $\bar{p}/p$ unmatched fraction, is in the charge-sign asymmetry of particle acceleration. The subsequent particle propagation through the galaxy or multiple accelerators plays no significant role in the phenomenon, as they act equally on all species. This proposition is particularly consistent with a scenario wherein almost all the positrons contributing to the observed $e^{+}/\left(e^{-}+e^{+}\right)$ ratio are produced in a single SNR of a particular kind described further in the paper. Electrons in the $e^{+}/\left(e^{-}+e^{+}\right)$ fraction may in part originate from an ensemble of other remnants. However, the single-source explanation for the $e^{+}/\left(e^{-}+e^{+}\right)$ anomaly is \emph{generic} to all SNRs of the kind and thus is equally consistent with its multi-source origin. From the Occam's razor perspective, this mechanism is preferred over those requiring \emph{different }types of sources, to state the obvious. A striking exception to the proposed scenario is the 100-300 GeV range where the current AMS-02 points significantly exceed our model predictions. This region then requires an independent source atop of the SNR contribution, that can be of a dark matter annihilation/decay or pulsar origin. Further model improvements are planned to see if simplifications made in its current version are responsible for the difference, but it appears to be unlikely. The proposed mechanism relies on the following two aspects of the DSA. The first one is the injection process whereby particles become supra-thermal and may then cross and re-cross the shock front, thus gaining more energy. The proposed injection mechanism is charge-sign asymmetric. It differs from the conventional DSA in which the injection efficiency primarily depends on the mass-to-charge ratio but not so much on the sign of the charge (see, however, the Discussion section). We will argue that the charge-sign dependence of injection arises when the shock propagates into an interstellar medium (ISM) containing clumps of dense molecular gas (MC, for short). The second aspect of the proposed mechanism concerns the phenomenon of nonlinear shock modification which is known to make the spectrum of low-energy particles steeper and that of the high-energy particles flatter than the canonical $p^{-4}$ spectrum produced by strong but unmodified shocks. Consequently, in the modified shocks, a point $p=p_{4}$ exists in the particle energy spectrum where the index is equal to four. Assuming that the bulk of galactic CR electrons are accelerated in conventional shocks, thus having $p^{-4}$ source spectra, the ratio of the modified positron spectrum to unmodified electron spectrum will show the required nonmonotonic behavior with a minimum at $p=p_{4}$. In a customary $p^{4}f$$\left(p\right)$ normalization, the individual positron spectrum is, therefore, the same as that of the $e^{+}/e^{-}$ ratio. Therefore, it conicides with the \emph{proton }spectrum, provided all the species are \emph{relativistic}. An analytic solution places the \emph{proton} $p^{4}f$$\left(p\right)$ minimum at $\lesssim10$ GeV/c (see Fig.5 in Ref. \citep{MDru01}), depending weakly on the shock Mach number, $M$, proton maximum energy, $E_{{\rm max}}$, and their injection rate. However, $M\gtrsim10$ and $E_{\max}\gtrsim1$ TeV conditions are required, along with some minimum proton injection, for the solution to transition into a strongly nonlinear regime (often called efficient acceleration). Although the minimum in the spectrum looks encouraging for explaining the nonmonotonic $e^{+}/e^{-}$ ratio, it was obtained for protons and needs to be reconsidered for positrons in the 1-10 GeV/c momentum range. The reason for that is a different momentum dependence of positron and proton diffusivity (positrons enter a relativistic transport regime at much lower energy than protons). Once an SNR shock is strongly modified, MCs in its precursor will survive the sub-shock UV and X- radiation, severely diminished in such shocks. At the same time, shock-accelerated CR protons illuminate the MC well before the subshock encounter. These CRs generate positrons (along with other secondaries) in the MC interior by colliding with the dense gas material. The CR protons also charge the MC \emph{positively;} as a result, many positively charged particles abandon the MC, while negatively charged particles remain inside. Being charged by the shock-accelerated protons, the MC thus acquires a positive potential which creates a charge-sign asymmetry for the subsequent particle injection into the DSA. Plasmas are intolerant to external charges and immediately restore charge neutrality. Nevertheless, a large and dense MC needs to build up a strong electric field to restore charge neutrality. Due to the high rigidity of CRs, their density in the MC interior increases almost simultaneously with that in the exterior, as a CR-loaded shock approaches the MC. However, by contrast with a strongly ionized exterior, where the plasma resistivity is negligible, the electron-ion, ion-neutral (and, in the case of very dense clouds, also electron-neutral) collisions inside the MC, provide significanty resistivity to the neutralizing electric current. Therefore, a strong macroscopic electric field is generated in response to the CR penetration. This field expels the secondary positrons most efficiently as the lightest positively charged species \textendash{} although it also shields the MC from low-energy CR protons. The mechanism outlined above implies that negatively charged primaries and secondaries have much better chances to stay in an MC than positively charged particles. When the subshock eventually reaches the MC, the subshock engulfs it, e.g., \citep{Inoue12,Draine2010BookISM}. What was inside of the MC, is transferred downstream unprocessed by the subshock. Therefore, the negatively charged particles in its interior largely evade acceleration. This charge-sign asymmetry of particle injection into the DSA explains why there is no $\bar{p}/p$ excess, similar to that of $e^{+}/e^{-}$ . It follows that the positron spectrum results from several interwoven processes. We will consider them separately and study their linkage. The remainder of the paper is organized as follows. Sec.\ref{sec:InteractionOfCRwithMC} deals with a spatial distribution of CR in a shock precursor, their propagation inside of an MC and electrodynamic processes that the CR induce there. In Sec.\ref{sec:Spectrum-of-shock-accelerated} we discuss and estimate the distribution of secondary positrons as they come out of the MC and become subject to the DSA. The spectrum of accelerated positrons is calculated from low to high energies, and the nature of the 8 GeV minimum is elucidated. We briefly discuss some of the alternative explanations of the $e^{+}/\left(e^{-}+e^{+}\right)$ excess in Sec.\ref{sec:Discussion}, followed by the Conclusion section, Sec.\ref{sec:Conclusions}.

} The objectives of this paper have been a detailed explanation of the $e^{+}/e^{-}$ energy spectrum and understanding of the charge-sign dependent particle injection and shock acceleration. The principal results of our study are: \begin{enumerate} \item assuming that an SNR shock environment contains clumps of weakly ionized dense molecular gas (MC), we investigated the effects of their illumination by shock accelerated protons before the shock traverses the MC. The main effects are the following: \begin{enumerate} \item an MC of size $L_{{\rm MC}}$ is charged (positively) by penetrating protons to$\sim\left(L_{{\rm MC}}/pc\right)\left(V_{sh}/c\right)\left(1eV/T_{e}\right)^{3/2}\left(n_{CR}/cm^{-3}\right)$GV, eq.(\ref{eq:fiMAX}) \item secondary positrons produced in $pp$ collisions inside the MC are pre-accelerated by the MC electric potential and expelled from the MC to become a seed population for the DSA \item most of the negatively charged secondaries, such as $\bar{p}$, along with electrons and heavier nuclei, remain locked inside the MC \end{enumerate} \item assuming that the shock Mach number, the proton injection rate, and their cut-off momentum exceeds the threshold of efficient acceleration regime \citep{MDru01}, we calculated the spectrum of injected positrons and, concomitantly, electrons \begin{enumerate} \item the momentum spectra of accelerated leptons have a concave form, characteristic for nonlinear shock acceleration, which physically corresponds to the steepening at low momenta, due to the subshock reduction, and hardening at high momenta, due to acceleration in the smooth part of the precursor flow \item the crossover region between the trends in (a) is also directly related to the change in the proton transport (from $\kappa\propto p^{2}$ to $\kappa\propto p$) and respective contribution to the CR partial pressure in a mildly-relativistic regime. The crossover pinpoints the 8 GeV minimum in the $e^{+}/\left(e^{+}+e^{-}\right)$ fraction measured by AMS-02 \item due to the nonlinear subshock reduction, the MC crosses it virtually unshocked so that secondary $\bar{p}$ and, in part, heavier nuclei accumulated in its interior largely evade shock acceleration \end{enumerate} \end{enumerate} Some important physical aspects of the proposed mechanism have not been elaborated. These include, but are not limited to, the following \begin{enumerate} \item calculation of energy distribution of runaway positrons preaccelerated in MC before their injection into the DSA \item calculation of electron injection for this kind of shock environment \item evaluation of conditions for the runaway gas breakdown in MC with associated pair production and calculation of the yield of this process \item escape of secondary antiprotons, generated in outer regions of MC or with sufficient energy, to the ambient plasma and their subsequent diffusive acceleration \item integration of the present calculations of positron spectra into the available fully nonlinear DSA solutions \item study of the MC interaction with a supersonic flow in modified shock precursor, bow shock formation and implications for additional particle injection \end{enumerate} Implementation of items (1), (2) and (5) will be particularly useful when AMS-02 gathers more statistics in the $>10^{2}$GeV range, so that the positron fraction saturation level can be more accurately compared with the prediction of the improved model. \appendix

1607.01820

1607

1607.04402_arXiv.txt

The EAGLE cosmological simulations reproduce the observed galaxy stellar mass function and many galaxy properties. In this work, we study the dust-related properties of present-day EAGLE galaxies through mock observations in the far-infrared and submm wavelength ranges obtained with the 3D dust radiative transfer code \SKIRT. To prepare an EAGLE galaxy for radiative transfer processing, we derive a diffuse dust distribution from the gas particles and we re-sample the star-forming gas particles and the youngest star particles into star-forming regions that are assigned dedicated emission templates. We select a set of redshift-zero EAGLE galaxies that matches the $K$-band luminosity distribution of the galaxies in the \emph{Herschel} Reference Survey (HRS), a volume-limited sample of about 300 normal galaxies in the Local Universe. We find overall agreement of the EAGLE dust scaling relations with those observed in the HRS, such as the dust-to-stellar mass ratio versus stellar mass and versus $\mathrm{NUV}-r$ colour relations. A discrepancy in the $f_{250}/f_{350}$ versus $f_{350}/f_{500}$ submm colour-colour relation implies that part of the simulated dust is insufficiently heated, likely because of limitations in our sub-grid model for star-forming regions. We also investigate the effect of adjusting the metal-to-dust ratio and the covering factor of the photodissociation regions surrounding the star-forming cores. We are able to constrain the important dust-related parameters in our method, informing the calculation of dust attenuation for EAGLE galaxies in the UV and optical domain.

Cosmological simulations are a valuable tool in the study of how galaxies form and evolve. Recently, hydrodynamical simulations of the formation of galaxies in cosmologically representative volumes have succeeded in reproducing many -- but not all -- observed properties of galaxies and of the intergalactic medium to unprecedented levels of agreement \citep[e.g.,][]{LeBrun2014, Vogelsberger2014, Schaye2015}. The mass resolution for baryonic matter in these simulations is on the order of $10^6$ solar masses. Physical processes on unresolved scales (including star formation and stellar feedback) are handled through sub-grid prescriptions. Zoom-in simulations \citep[e.g.,][]{Hopkins2014, Wang2015, McKinnon2016, Sawala2016} offer a better resolution, however, they still use similar sub-grid prescriptions. Inevitably these limitations lead to uncertainties in some of the simulation predictions. By comparing simulation results and observations we hope to examine the empirical scaling laws, deduce improved sub-grid prescriptions, and eventually, to further our understanding of the underlying physical processes. Because properties of real galaxies are derived from observed quantities (i.e.\ fluxes), they may be subject to unknown systematic biases. Making mock observations of simulated galaxies enables direct comparison to observational data, and helps to characterise the systematics involved in the transformation between intrinsic and observed quantities \citep[see, e.g.,][]{Hayward2015, Guidi2015}. Extinction by dust grains residing in the interstellar medium (ISM) can substantially influence the flux detected from a galaxy in the UV and optical wavelength ranges. It is very hard to estimate the dust mass in a galaxy based solely on the information at these wavelengths, and thus it is difficult to account accurately for the dust obscuration effect \citep[e.g.,][]{Disney1989, Byun1994}. To alleviate this limitation, one can turn to the far-infrared (FIR) to submm wavelength range. In this window, the continuum spectra of star-forming galaxies are dominated by thermal emission from dust grains that reprocess the UV/optical radiation, providing an independent and more direct measurement of the amount of dust in a galaxy. This additional information is especially useful for constraining the dust modelling of numerically simulated galaxies that have no explicit dust component. On the other hand, accurately predicting dust emission from a simulated galaxy requires solving a nontrivial 3D radiative transfer problem \citep[see, e.g.,][]{Whitney2011, Steinacker2013}. \ch{Several authors have performed UV to submm radiative transfer modelling for up to a few dozen simulations of isolated galaxies or galaxy mergers \citep[e.g.,][]{Narayanan2010a, Narayanan2010b, Jonsson2010, Scannapieco2010, Hayward2011, Hayward2012, Hayward2013, Hayward2014, Hayward2015, Robitaille2012, Dominguez-Tenreiro2014, Saftly2015}. While \citet{Torrey2015} do produce mock images and SEDs for thousands of present-day galaxies in the cosmological simulation Illustris \citep{Vogelsberger2014}, they do not include dust emission, limiting the observables to UV, optical and near-infrared (NIR) wavelengths. In this present work, we study the effects of dust in the full UV to submm wavelength range. We use simulated galaxies that were evolved as part of a cosmologically representative volume, and that are available in sufficiently large numbers to allow a statistically relevant confrontation with observations.} \ch{Specifically}, we concentrate on the FIR and dust-related properties of the present-day galaxies produced by the EAGLE simulations \citep{Schaye2015, Crain2015}. EAGLE is a suite of hydrodynamical simulations of the formation of galaxies in cosmologically representative volumes, with sub-grid models for radiative cooling, star formation, stellar mass loss, and feedback from stars and accreting black holes. The sub-grid physics recipes are calibrated to reproduce the present-day galaxy stellar mass function and galaxy sizes, and show good agreement with many observables not considered in the calibration, including present-day specific star-formation rates, passive fractions, the Tully-Fisher relation \citep{Schaye2015}, and the neutral gas content \citep{Bahe2016}. The simulations also track the observed evolution of the galaxy stellar mass function out to redshift $z = 7$ \citep{Furlong2015} and reproduce the observed optical colours for galaxies in the Local Universe \citep{Trayford2015, Trayford2016}. We use the \emph{Herschel} Reference Survey \citep{Boselli2010} (HRS), a volume-limited sample of about 300 `normal' galaxies in the Local Universe, as a reference for observed dust properties. We select a set of redshift-zero EAGLE galaxies that matches the $K$-band luminosity distribution of the HRS galaxies, and we use the 3D dust radiative transfer code \SKIRT\ \citep{Baes2011, Camps2015a} to calculate observable properties for these galaxies from UV to submm wavelengths. We compare the stellar mass, dust mass, and star-formation rate derived from our mock observations through standard tracers with the intrinsic EAGLE values, and we compare the EAGLE dust scaling relations with those observed for HRS galaxies presented by \citet{Boselli2012} and \citet{Cortese2012}. Finally, we investigate the effect of varying dust-related parameters in our post-processing procedure. This allows us to constrain these parameters, thus informing the calculation of dust attenuation for EAGLE galaxies in the UV and optical domain by \citet{Trayford2016}. In Sect.~\ref{Methods.sec} we provide some background on the EAGLE simulations and the \SKIRT\ radiative transfer code, and we describe how the EAGLE results were exported to and post-processed by \SKIRT, with some details relayed to the appendices. In Sect.~\ref{Results.sec} we present and discuss the results of our analysis, and in Sect.~\ref{Conclusions.sec} we summarise and conclude.

\label{Conclusions.sec} We calculated mock observations in the wavelength range from UV to submm for simulated galaxies extracted from the EAGLE suite of cosmological simulations using the radiative transfer code SKIRT. To help overcome some of the resolution limitations of the simulations, we employed sub-grid models for the star-forming regions and for the diffuse dust distribution. We also took special care to mimic the effects of instrumental properties and observational limitations when calculating band-integrated fluxes, which are important especially in submm bands. To validate our method, and at the same time confront the properties of the simulated galaxies with observed galaxies, we selected a set of present-day EAGLE galaxies that matches the $K$-band luminosity distribution and overall morphological type classification (using the sSFR as a proxy) of the galaxies in the \emph{Herschel} Reference Survey (HRS), a volume-limited sample of about 300 normal galaxies in the Local Universe. We evaluated some intrinsic properties of our selected galaxies (Fig.~\ref{IntrinsicProperties.fig}), calculated by summing over the particles, and confirmed that the stellar masses, star-formation rates and dust masses fall in the expected range, while the average gas metallicities are above the metallicities observed in comparable galaxies. We evaluated some relevant tracers by comparing the values derived from our mock observations to the corresponding intrinsic values (Figs.~\ref{StellarAndDustMass.fig} and \ref{StarFormationRate.fig}). We found that the \citet{Zibetti2009} calibration of stellar mass versus $g$ and $i$ band luminosity used by \citet{Cortese2012} underestimates stellar mass by about 0.25 dex, in line with the systematic uncertainty on the stellar mass-to-light ratio relation reported by other authors. Furthermore, the \citet{Cortese2012} recipe for deriving dust mass from the three SPIRE fluxes produces an offset that seems to be mostly caused by differences in the assumed dust absorption coefficient at 350~\micron. The star-formation indicators based on respectively NUV, 24~\micron, and integrated infrared fluxes perform fairly well, although the 24~\micron\ estimates are consistently low, most likely because some of the dust in our model is too cold. We then studied dust scaling relations, including the $f_{250}/f_{350}$ versus $f_{350}/f_{500}$ submm colour-colour relation (Fig.~\ref{Boselli.fig}) and the dust-to-stellar mass ratio versus stellar mass and versus $\mathrm{NUV}-r$ colour relations (Fig.~\ref{Cortese.fig}), comparing the properties derived from mock observations of the EAGLE galaxies with those observed for HRS galaxies. Using our `standard' set of post-processing parameter values, we found good correspondence between the EAGLE and HRS scaling relations, with one important caveat. The submm colours indicate that the average temperature of the dust in our EAGLE galaxy models is lower than observed. We concluded that, even with the sub-grid techniques in our procedures, our model does not fully capture the clumpiness of the dust distribution in galaxies, so that an insufficient amount of dust is irradiated by the strong radiation fields present within and near star-forming regions. We investigated the effects of varying the assumed dust-to-metal fraction, $f_\mathrm{dust}$ (Fig.~\ref{ParameterStudy_fdust.fig}), and PDR covering fraction, $f_\mathrm{PDR}$ (Fig.~\ref{ParameterStudy_fpdr.fig}), in our post-processing procedure. The first parameter controls the diffuse dust, while the latter controls the dust near young stellar populations. We found that the effects on the scaling relations are consistent with expectations, although it is hard to determine unambiguously optimal values for both parameters because of the partial degeneracy of the effects. We settled on $f_\mathrm{dust}=0.3$ and $f_\mathrm{PDR}=0.1$ as our standard values, noting that these values also depend on the properties of the dust in our model. These parameter values are compatible with observed values of $f_\mathrm{dust}$ \citep{Dwek1998, Brinchmann2013, Zafar2013} and values of $f_\mathrm{PDR}$ suggested by other authors \citep{Jonsson2010}. In conclusion, our analysis has shown that, in spite of some limitations, the EAGLE simulations can reproduce infrared and submm observations through a physically motivated post-processing procedure. Furthermore, we have used the mock submm observations to constrain the important dust-related parameters in our method, leading to a consistent calculation of dust attenuation in the UV and optical domain by \citet{Trayford2016}. While we studied present-day galaxies in this work, our post-processing method is equally applicable to galaxies at higher redshifts, and could be readily adapted to other hydrodynamical simulations. The method presented here opens many opportunities for future work. We plan to add infrared and submm fluxes for a large subset of the EAGLE galaxies at various redshifts to the public EAGLE database \citep{McAlpine2016} as a service to other researchers. We could study the morphology and structure of mock EAGLE galaxy images in various wavelength bands from UV to submm, comparing those properties with resolved observations in the same bands. It would also be instructive to post-process some of the more resolved galaxies in the zoom-in simulations \citep{Sawala2016, Oppenheimer2016} based on the EAGLE code. In a more distant future, the astrophysical community will undoubtedly develop more advanced simulation techniques. Future simulations of galaxy evolution and assembly may use sub-grid recipes on a smaller scale to model a cold phase in the ISM, which will allow more accurate modeling of the clumpy dust structure during post-processing. Radiative transfer codes may include sub-grid models of star-forming regions that are connected to the overall RT model in a self-consistent manner, as opposed to employing `disconnected' SED templates. We hope that our current work will help inform such future efforts.

1607.04402

1607

1607.06631_arXiv.txt

The motivation of the present work is to reconstruct a dark energy model through the {\it dimensionless dark energy function} $X(z)$, which is the dark energy density in units of its present value. In this paper, we have shown that a scalar field $\phi$ having a phenomenologically chosen $X(z)$ can give rise to a transition from a decelerated to an accelerated phase of expansion for the universe. We have examined the possibility of constraining various cosmological parameters (such as the deceleration parameter and the effective equation of state parameter) by comparing our theoretical model with the latest Type Ia Supernova (SN Ia), Baryon Acoustic Oscillations (BAO) and Cosmic Microwave Background (CMB) radiation observations. Using the joint analysis of the SN Ia+BAO/CMB dataset, we have also reconstructed the scalar potential from the parametrized $X(z)$. The relevant potential is found, which comes to be a polynomial in $\phi$. From our analysis, it has been found that the present model favors the standard $\Lambda$CDM model within $1\sigma$ confidence level.

The various cosmological observations such as Type Ia Supernovae~\cite{SN,Riess:1998cb}, cosmic microwave background (CMB) radiation~\cite{Planck:2015xua, Ade:2015lrj, Ade:2014xna, Ade:2015tva, Array:2015xqh, Komatsu:2010fb, Hinshaw:2012aka}, large scale structure~\cite{LSS,Seljak:2004xh}, baryon acoustic oscillations (BAO)~\cite{Eisenstein:2005su}, and weak lensing~\cite{Jain:2003tba} have supported that the expansion of the current universe is accelerating. All of these observations also strongly indicate that the alleged acceleration is rather a recent phenomenon and the universe was decelerating in the past. Two representative approaches have been proposed to account for the late-time cosmic acceleration. The first approach is to assume the existence of ``dark energy'' (DE) in the framework of general relativity. The second approach is to consider the modification of gravity on the large scale (for reviews on the issues of DE and the modified theories of gravitation, see, for example, \cite{Nojiri:2010wj,Nojiri:2006ri,Book-Capozziello-Faraoni,Capozziello:2011et,delaCruzDombriz:2012xy,Bamba:2012cp,Joyce:2014kja, Koyama:2015vza,Bamba:2015uma}). In this work, we will concentrate only on the first approach and consider DE to be responsible for this accelerated phenomenon. There are some excellent review articles where various DE models have been comprehensively discussed \cite{3,4,4a,5}. The simplest candidate of DE a is cosmological constant $\Lambda$ whose energy density remains constant with time and its {\it equation of state} (EoS) parameter is, $\omega_{\Lambda}=-1$. However, the models based upon cosmological constant suffer from the {\it fine tuning} and the cosmological {\it coincidence} problems \cite{sw, Steinhardt}. Scalar field models with generic features can alleviate these problems and provide the late-time evolution of the universe (see Ref. \cite{4} for a review). Scalar field models are very popular as the simplest generalization of cosmological constant is provided by a scalar field, dubbed as quintessence field, which can drive the acceleration with some suitably chosen potentials. In this case, one needs some degree of fine tuning of the initial conditions to account for the accelerated expansion of the universe and none of the potentials really have proper theoretical support from field theory explaining their origin (for review, see \cite{6}). In the last decade, an enormous number of DE models were explored to explain the origin of this late time acceleration of the universe and none of these models have very strong observational evidence \cite{4}. Therefore, the search is on for a suitable DE model and the present study is one of them.\\ \par In Ref. \cite{ellis}, Ellis and Madsen had discussed about {\it reconstruction} method to find the scalar field potential. Recently, this method finds a very wide application in current research in cosmology. However, there are two types of reconstruction, namely, parametric and non-parametric. The parametric reconstruction method is an attempt to build up a model by assuming a specific evolution scenario for a model parameter and then estimate the values of the parameters from different observational datasets. On the other hand, the non-parametric reconstruction method does not require any specific assumption for the parameters and finds the nature of cosmic evolution directly from observational dataset.\\ \par In the context of DE, the reconstruction method was first discussed in \cite{starpr}, where Starobinsky determined the scalar field potential from the observational dataset from the behavior of density perturbations in dust-like matter. Some other earlier works on reconstruction have been discussed in \cite{tdsprl,relz} where the dataset of cosmological distance measurement has been used. In practice, a large number of dynamical models have been proposed for DE in which the properties of DE component are generally summarized as a perfect fluid with a time-dependent EoS parameter $\omega_{\phi}(z)$. In building up the DE model by the parametric reconstruction method, efforts are normally made through the DE EoS parameter. In literature, there are many examples where the authors had proposed different redshift parametrizations of $\omega_{\phi}$ to fit with observational data \cite{lz1,lz2,cpl1,cpl2} (for review, see also Refs. \cite{jmang,kshi,aam,aam2}). However, it has been found that the parametrization of the energy density $\rho_{\phi}(z)$ (which depends on its EoS parameter through an integral) provides tighter constraints than $\omega_{\phi}(z)$ from the same observational dataset (for details, see Refs.\cite{mt0,mt1,mt2}). Recently, many investigations have been performed to find the actual functional form of $\omega_{\phi}$ directly from the available datasets \cite{np1,np2,np3,np4}. However, the problem with this method is that the parameters of interest usually contain noisy data. The present work uses the idea of parametrizing the DE density, where we have presented a parametric reconstruction of the DE function $X(z)$ (which is basically the DE density in units of its present value) to study the essential properties of DE. The basic properties of this chosen $X(z)$ has been discussed in detail in the next section. The functional form of $X(z)$ depends on the model parameters which have been constrained from the observational datasets. The constraints on the model parameters are obtained by using various observational datasets (namely, SN Ia, BAO and CMB) and $\chi^{2}$ minimization technique. With the estimated values of model parameters, we have then reconstructed the deceleration parameter and the EoS parameter at the $1\sigma$ and $2\sigma$ confidence levels. Furthermore, we have also tried to reconstruct the scalar potential $V(\phi)$ directly from the dark energy function $X(z)$. Clearly, the present study enables us to construct the scalar field potential without assuming its functional form. This is one of the main objectives of the present work. We have found that the results obtained in this work are consistent with the recent observations and the model do not deviate very far from the $\Lambda$CDM model at the present epoch. \\ \par The outline of the paper is as follows. In the next section, we have presented the basic formalism of a flat FRW cosmology along with the definitions of different cosmological parameters. We have then solved the field equations for this toy model using a specific choice of the dark energy function $X(z)$. The observational datasets and methodology are discussed in section \ref{data}. The main results of this analysis are summarized in section \ref{result}. Finally, in the last section, we have presented our main conclusions.

In this paper, we have focused on a quintessence model in which the scalar field is considered as a candidate of dark energy. It has been shown that for a spatially flat FRW universe, we can construct a presently accelerating model of the universe with the history of a deceleration in the past by considering a specific choice of the dimensionless dark energy function $X(z)$. The motivation behind this particular choice of $X(z)$ has been discussed in details in section \ref{sec2} and for this specific ansatz, we have solved the field equations and have obtained the expressions for different cosmological parameters, such as $H(z)$, $q(z)$ and $\omega_{eff}(z)$. As mentioned earlier that the model parameters ($\alpha$ and $\beta$) are a good indicator of deviation of the present model from cosmological constant as for $\alpha=0$ and $\beta=0$ the model mimics the $\Lambda$CDM model. We have also constrained the model parameters using the SN Ia+BAO/CMB(WMAP7) and SN Ia+BAO+CMB(Planck) datasets to study the different properties of this model extensively. It is evident from table \ref{tab:fntable1} that the best-fit values of $\alpha$ and $\beta$ are very close to zero. So, our analysis indicates that the reconstructed $\omega_{\phi}(z)$ is very close to the $\Lambda$CDM value at the present epoch. In summary, using SN Ia+BAO/CMB dataset jointly, we have then reconstructed various parameters (e.g., $q(z)$, $\omega_{eff}(z)$ and $\omega_{\phi}(z)$) as well as the quintessence potential $V(\phi)$ directly from the chosen $X(z)$, which describes the properties of the dark energy. The resulting cosmological scenarios are found to be very interesting. It has been found that the evolution of $q(z)$ in this model shows a smooth transition from a decelerated to an accelerated phase of expansion of the universe at late times. As discussed in section \ref{result}, it has been found that our reconstructed results of $q(z)$ and $\omega_{\phi}(z)$ are in good agreement with the previous works \cite{aam,jmang,kshi,aam2}. For completeness of the work, we have also derived the form of the effective scalar field potential $V(\phi)$, in terms of $\phi$, for this model and the resulting potential is found to be a polynomial in $\phi$. \\ \par From the present investigation, it can be concluded that the SN Ia+BAO/CMB dataset although supports the concordance $\Lambda$CDM model at the $1\sigma$ confidence level, but it favors the scalar field dark energy model as well. In other words, it is well worth emphasizing that the observational datasets are not yet good enough to strongly distinguish present dark energy model from the $\Lambda$CDM model at present. With the progress of the observational techniques as well as the data analysis methods in the future, we hope that the parameters in $X(z)$ can be constrained more precisely, which will improve our understanding about the nature of dark energy. The present analysis is one preliminary step towards that direction. In future, we plan to test this parametric form of $X(z)$ in scalar-tensor theories of gravity.

1607.06631

1607

1607.05823_arXiv.txt

\noindent\noindent With respect to earlier investigations, the theory of multi-component, concentric, copolar, axisymmetric, rigidly rotating polytropes is improved and extended, including subsystems with nonzero density on the boundary and subsystems with intersecting boundaries. The formulation is restricted to two subsystems for simplicity but, in principle, can be extended to $N$ subsystems. Equilibrium configurations are independent of the nature of the fluid i.e. collisional or collisionless, provided the polytropic index lies within the range, $1/2\le n\le5$, as in one-component systems. The solution of the equilibrium equations is expanded in power series, which can be continued up to the boundary and outside via starting points placed at increasingly larger distance from the centre of mass. A detailed analysis is devoted to special cases where the solution of the equilibrium equations can be expressed analytically. Finally a guidance example is shown, involving homogeneous subsystems with intersecting boundaries, where a substantially flattened component extends outside a slightly flattened one.

\label{intro} Special attention has been devoted to polytropes since more than a century, and is continuing to be devoted at present. In fact, polytropic models are useful not only for the rough estimates of some processes in real stars, but also in the precise investigation of some matters of principle, such as the effect of increasing central condensation on nonradial pulsation models [26], the structure of supermassive and superdense stars [42], the collapse of stellar cores [17], the fission and the equatorial shedding of matter due to rotation [24], [23], [22], [30], [33], [28], [19], the oscillation and stability of rapidly rotating stellar models [38], [2], [3], [32], the structure of neutron stars and quark stars [16], [25], the expression of the total mass of a rotating configuration as a function of the central density [35], equilibrium configurations and triaxiality in stars and stellar systems [39], [40]. For exhaustive review and complete references, an interested reader is addressed to classical textbooks concerning stellar structure [14] and, in addition, applications in astrophysics and related fields [21]. Among others, polytropes can provide a continuous transition from null to infinite matter concentration i.e. homogeneous and Roche (mass points surrounded by vanishing atmospheres) models, distorted by rotation and/or tidal interaction [24], Chap.\,IX, \S235. On the other hand, large-scale stellar systems (galaxies and galaxy clusters) appear to be made of at least two subsystems interacting only via gravitation e.g., bulge+disk, core+halo, visible baryonic (including leptons)+dark nonbaryonic matter. In most applications, each component is treated separately and the model is a simple superposition of the two matter distributions e.g., [12]. To this respect, models involving two-component polytropes make a further step in that each subsystem readjusts itself in presence of tidal interaction [10], [5], [7] and, in addition, both collisional and collisionless fluids can be described within the range of polytropic index, $1/2\le n\le5$, [39] where $n=0,5,$ relates to homogeneous and Roche or Plummer models, respectively. The current paper aims to improve and extend earlier investigation on two-component polytropes [10], [5], [7], mainly focusing on the following points: (i) formulation of the theory where due effort is devoted to the connection with one-component polytropes; (ii) extension of the theory to subsystems with nonzero density on the boundary; (iii) extension of the theory to subsystems with intersecting boundaries; (iv) extension of the theory to collisionless fluids; (v) series expansion of the solution of equilibrium equations, which can be continued up to the boundary and outside, via starting points placed at increasingly larger distance from the centre of mass; (vi) detailed analysis of special cases where the solution of equilibrium equations can be expressed analytically; (vii) a guidance example involving homogeneous subsystems with intersecting boundaries, where a substantially flattened component extends outside a slightly flattened one. The paper is organized as follows. Points (i)-(vi) mentioned above make the subject of Section \ref{bath}, where each argument is discussed in one or more Subsections. The guidance example stated on point (vii) is shown in Section \ref{guex}. The discussion is performed in Section \ref{disc} and the conclusion is drawn in Section \ref{conc}. Further details on the formulation of the theory can be seen in the Appendix.

\label{conc} The theory of EC2 polytropes has been reformulated. The method used in earlier investigations [36], [29], [6] for series expansion of the solution of EC1 equation and related associated equations has been extended to the solution of EC2 equation and related associated equations. In addition, special cases where the solution can be expressed analytically have been considered in detail: $(n_i,n_j)=(0,0), (0,1), (1,0),(1,1), (5,0)$. Subsystems with nonvanishing density on the boundary and subsystems with intersecting boundaries have also been included in the general theory, with regard to both collisional and collisionless fluids. A selected class of configurations has been defined in terms of central densities, $\lambda_w$, scaling radii, $\alpha_w$, rotation parameters, $\upsilon_w$, where $w=i,j$. Accordingly, it has been realized subsystems belong to one among the following states: (1) rotating to a different extent and showing different scaling radius; (2) rotating to a different extent but showing equal scaling radius; (3) rotating to the same extent but showing different scaling radius; (4) rotating to the same extent and showing equal scaling radius. With regard to the noncommon region, it has been realized the solution of the EC2 equation and associated EC2 equations can be expressed similarly to its counterpart related to EC1 polytropes. It has been recognized the contrary holds for the common region, where the solution of EC2 equation and associated EC2 equations involves both subsystems instead of only one, leaving aside a few special cases. The main results of the current paper may be summarized in the following points. \begin{description} \item[(i)] Subsystems with nonvanishing density on the boundary are included in the description, which could be useful in e.g., modelling neutron stars and strange quark stars [16], [25] and computing the total mass of a rotating configuration as a function of the central density [35]. \item[(ii)] Subsystems with intersecting boundaries are included in the description, which could be useful for representing special configurations e.g., a substantially flattened subsystem extending outside a slightly flattened one. \item[(iii)] Special cases where the results can be expressed analytically are analysed in detail with the addition of a guidance example involving homogeneous configurations for simplicity, but with intersecting boundaries. \end{description}

1607.05823

1607

1607.07367_arXiv.txt

{The discontinuous Galerkin finite element method (DG-FEM) is successfully applied to treat a broad variety of transport problems numerically. In this work, we use the full capacity of the DG-FEM to solve the radiative transfer equation in spherical symmetry. We present a discontinuous Galerkin method to directly solve the spherically-symmetric radiative transfer equation as a two-dimensional problem. The transport equation in spherical atmospheres is more complicated than in the plane-parallel case due to the appearance of an additional derivative with respect to the polar angle. The DG-FEM formalism allows for the exact integration of arbitrarily complex scattering phase functions, independent of the angular mesh resolution. We show that the discontinuous Galerkin method is able to describe accurately the radiative transfer in extended atmospheres and to capture discontinuities or complex scattering behaviour which might be present in the solution of certain radiative transfer tasks and can, therefore, cause severe numerical problems for other radiative transfer solution methods. }

Many stars, such as AGB stars, have extended atmospheres which has important physical and observational implications. Especially the radiation field at large distances from the inner stellar disk becomes very dilute and is confined primarily to a narrow solid angle around the radial direction \citep{Mihalas1978}. Additionally, planets can also feature extended atmospheres in which case the use of the usual assumption of a plane-parallel atmosphere is no longer valid. This includes, for example, the extended and spherically stratified atmosphere of Pluto \citep{Gladstone2016Sci...351.8866G}. Compared to the radiative transfer equation in the plane-parallel approximation, the transfer equation in spherical symmetry is more complicated and more difficult to handle. This is especially caused by the appearance of a derivative with respect to the polar angle, which is not present in the plane-parallel geometry. Other cases, in which the plane-parallel radiative transfer equation becomes much more complicated, include systems with a moving medium, where photons are subject to the Doppler effect and aberration \citep{Mihalas1978}. Depending on the transport coefficients, the general radiative transport equation can have a very different mathematical character. As described in \citet{Kanschat2009nmmr.book.....K}, it is a hyperbolic wave equation in regions without matter, an elliptic diffusion equation in case of an optically thick medium, or a parabolic equation when the light is strongly peaked in the forward direction. It is, therefore, very difficult to find numerical methods, which are capable of dealing with these different (or even mixed) types of behaviour. The problem of a grey spherical atmosphere in radiative equilibrium has been studied first by \citet{Kosirev1934MNRAS..94..430K} and \citet{Chandrasekhar1934MNRAS..94..444C}. An extension of the iterative moment method for plane-parallel atmospheres using variable Eddington factors \citep{Auer1970MNRAS.149...65A} to the problem of a spherically-symmetric atmosphere restricted to isotropic scattering has been introduced by \citet{Hummer1971MNRAS}. Additionally, approximative solutions based on a generalised Eddington approximation have been developed by \citet{Lucy1971ApJ...163...95L} and \citet{Unno1976PASJ...28..347U}, respectively. Alternatively, the long \citep{Cannon1970ApJ...161..255C} or short characteristic methods \citep{Kunasz1988JQSRT..39...67K} can be used. Combined with e.g. an operator perturbation method to account for the unknown source function -- such as the accelerated lamdba iteration \citep{Cannon1973JQSRT..13..627C} -- they can be used to solve the radiative transfer equation in multiple dimensions \citep{Olson1987JQSRT..38..325O, Kunasz1988JQSRT..39....1K}. In combination with the work of \citet{Hummer1971MNRAS}, operator splitting methods have also been directly utilised for spherically-symmetric atmospheres by \citet{Kubat1994A&A...287..179K}. Nowadays, the computationally demanding Monte Carlo techniques can solve very efficiently the radiative transfer problem for a broad variety of applications ranging from stellar atmospheres \citep[e.g.][]{1985ApJ...288..679A,1993ApJ...405..738L,1999A&A...344..282L}, circumstellar disks \citep[e.g.][and the references therein]{2004A&A...417..793P}, supernovae \citep[e.g.][]{2005A&A...429...19L} to dusty objects \citep[e.g.][and the references therein]{2013ARA&A..51...63S}. Their advantages are, for example, that they are intrinsically three-dimensional, can handle easily complex geometries and density distributions, and have a low algorithmic complexity. However, in some special cases with very optically thick regions the Monte Carlo approach becomes very slow because the number of photon interactions, i.e. scattering, increases exponentially. In these situations grid-based solution techniques, e.g. a finite difference, finite element, or finite volume method, for the differential equation of radiative transfer are more favourable even if there are algorithmically quite complex. Moreover, grid-based techniques allowing for higher order approximations as well as error control and they can be formulated to be intrinsically flux conserving \citep{Hesthaven2008nodal, Henning2001IAUS..200..567H}. In this study, we present a numerical method for a direct solution of the radiative transfer equation in spherical symmetry by means of a finite element method. Finite element methods have already been previously used to solve the radiative transfer equation, most notably in the three-dimensional case by \citet{Kanschat1996PhDT........73K} or \citet{Richling2001A&A...380..776R}. In these cases, a continuous Galerkin finite element was used to discretise the transport operator. Such an approach, however, causes problems in cases where the transport equation has a dominant hyperbolic character, which may lead to discontinuities within the solution. It is possible to stabilise the numerical scheme by introducing a streamline diffusion method, for example. The streamline diffusion method essentially transforms the hyperbolic equation into a parabolic one by adding a small diffusion in the direction of photon transport. While this stabilises the numerical discretisation, it also adds a free parameter to the method, which has to be chosen carefully in such a way, that it stabilises the numerical scheme, but does not result in a decreased accuracy of the solution (see e.g. \citet{Eriksson1996cde..book.....E,Kanschat1996PhDT........73K}). An alternative way to deal with such problems is to soften the requirements for the solution of the finite element method, in particular the requirement of continuity of the solution across element boundaries. One such numerical approach in the context of finite element methods is the discontinuous Galerkin method. The discontinuous Galerkin finite element method (DG-FEM) was first introduced by \citet{Reed_Hill_1973} to study neutron transport problems. Since then, the method has been successfully applied to a broad variety of parabolic, hyperbolic, and elliptic problems (e.g. \citet{Hesthaven2008nodal, Cockburn2003}). For three-dimensional internal heat transfer problems, the discretisation with a discontinuous Galerkin method has been studied by \citet{Cui2004NHTB...46..399C}. For astrophysical applications, discontinuous Galerkin methods have, for example, been used by \citet{Dykema1996ApJ...457..892D} and \citet{Castor1992ApJ...387..561C}. In this work we use a DG-FEM to directly solve the radiative transfer equation with arbitrarily complex scattering phase functions in spherical symmetry as a two-dimensional problem. In contrast to other solution methods, the scattering integral can thereby be evaluated exactly, independent of the angular mesh resolution. The development of the numerical scheme is outlined for the formulation of the DG approach in Sects. \ref{sec:rt_eq} \& \ref{sec:dg_fem}. Its applicability to several test problems is presented in Sect. \ref{sec:test_problems}. Remarks on the numerical efficiency are given in Sect. \ref{sec:efficiency}, followed by an outlook (Sect. \ref{sec:outlook}) and a summary (Sect. \ref{sec:summary}).

\label{sec:summary} In this work we presented a numerical method to directly solve the radiative transfer equation in spherical symmetry. We used a discontinuous Galerkin finite element method to discretise the transport equation as a two-dimensional boundary value problem. Arbitrary complex scattering phase function can, in particular, be integrated exactly using this approach. The discretisation yields a sparse system of equations which can be solved by standard iterative methods with suitable preconditioners. The applicability of the DG-FEM is verified on a set of different test problems. The discontinuous Galerkin method is able to directly describe the spherical dilution of the radiation field in an extended atmosphere. We show that the DG-FEM is able to accurately capture discontinuities which might appear in the solution of the transfer equation for special test problems. We also compared our method with previously published radiative transfer calculations for a spherical, isotropically scattering atmosphere. The agreement between the different approaches was found to be excellent. Consequently, the discontinuous Galerkin method is perfectly suitable to treat accurately the radiative transfer in a spherical atmosphere even with e.g. discontinuities in the radiation field or complex scattering behaviour, which might cause severe numerical problems for other radiative transfer solution methods.

1607.07367

1607

1607.05681_arXiv.txt

\noindent Local helioseismology has opened new frontiers in our quest for understanding of the internal dynamics and dynamo on the Sun. Local helioseismology reconstructs subsurface structures and flows by extracting coherent signals of acoustic waves traveling through the interior and carrying information about subsurface perturbations and flows, from stochastic oscillations observed on the surface. The initial analysis of the subsurface flow maps reconstructed from the five years of SDO/HMI data by time-distance helioseismology reveals the great potential for studying and understanding of the dynamics of the quiet Sun and active regions, and the evolution with the solar cycle. In particular, our results show that the emergence and evolution of active regions are accompanied by multi-scale flow patterns, and that the meridional flows display the North-South asymmetry closely correlating with the magnetic activity. The latitudinal variations of the meridional circulation speed, which are probably related to the large-scale converging flows, are mostly confined in shallow subsurface layers. Therefore, these variations do not necessarily affect the magnetic flux transport. The North-South asymmetry is also pronounced in the variations of the differential rotation ('torsional oscillations'). The calculations of a proxy of the subsurface kinetic helicity density show that the helicity does not vary during the solar cycle, and that supergranulation is a likely source of the near-surface helicity.

Observations of solar oscillations provide a unique opportunity to obtain information about the structure and dynamics of the solar interior beneath the visible surface. The oscillations with a characteristic period of five minutes represent acoustic waves stochastically excited by the turbulent convection in a shallow subsurface layer. The excitation mechanism has not been completely understood. However, recent numerical simulations have shown that the waves can be excited due to the interaction of turbulent vortex tubes ubiquitously generated in the intergranular lanes \citep{Kitiashvili2011}. These stochastic waves produce chaotic oscillation patterns on the solar surface. However, a spectral analysis of the time series of these patterns reveals that most of the oscillation power is concentrated in a set of normal modes (Fig.~\ref{fig1}a). These modes represent standing acoustic waves trapped in the subsurface layers by their reflection between the sharp density gradient near the surface, and the increasing sound speed in the interior. The depth of the inner reflection depends on the horizontal wavelength of the oscillations. The horizontal wavelength, $\lambda_h$, is usually represented in terms of the spherical harmonic degree, $\ell = 2\pi R/\lambda_h$. The oscillation frequency is expressed in terms of cyclic frequency $\nu=\omega/2\pi$. In the $\ell-\nu$ diagram shown in Figure~\ref{fig1}a, the lowest ridge represents the surface gravity mode ($f$-mode). The other ridges are acoustic modes of various radial order $n$, which is equal to the number of nodes along the radius. This number is higher for higher frequency ridges. The time-series of solar oscillations have been obtained almost uninterruptedly since 1995 from the ground-based network GONG and space mission SOHO (Solar and Heliospheric Observatory) and SDO (Solar Dynamics Observatory). The oscillation frequencies are routinely measured from 72- and 108-day time series by fitting the modal lines which are used for inferring variations of the sound speed, asphericity, and differential rotation rate. This approach called `global helioseismology' has provided important information about the structure, composition and dynamics of the solar interior. In particular, it was led to the discovery of a sharp radial gradient of the differential rotation at the base of the convection zone \citep{Kosovichev1996a}, the so-called tachocline, the near-surface rotational shear layer\citep{Schou1998}, subsurface zonal flows migrating with the solar activity cycle\citep{Kosovichev1997}. Recent analysis of the high-degree oscillation modes revealed a sharp gradient of the sound speed in a narrow 30-Mm deep layer just beneath the solar surface \citep{Reiter2015}. This layer (called `leptocline', \citep{Godier2001}) presumably plays an important role in the solar dynamo \citep{Pipin2011}. \begin{figure}[t] \begin{center} \includegraphics[width=\linewidth]{kosovichev_figure_01.eps} \caption{ a) The power spectrum of solar oscillations as a function of the angular degree $\ell$, and cyclic frequency, $\nu$. The enhanced power corresponds to the normal oscillation modes of the Sun. b) The cross-covariance function ('time-distance diagram') of solar oscillations as a function of the distance between the correlation points on the solar surface and the time lag of the cross-covariance. The lowest ridge is formed by acoustic wave packets traveling between two surface points ('source' and 'receiver')through the solar interior (so-called, the first skip); the higher ridges are formed by the wave packets arriving to the receiver after additional reflections from the surface (the 'second' skip, and so on). }\label{fig1} \end{center} \end{figure} It is important to note that while the oscillation power spectrum extends into the high-frequency region (10 mHz and higher), only the ridge parts with the frequency below the acoustic cut-off frequency (which is approximately at 5.2 mHz) represent the normal modes. The higher frequency parts correspond to so-called `pseudo-modes' . The pseudo-modes are formed by interference between the waves traveling from the excitation sources directly to the surface and the waves which come to the same surface location after reflection in the interior. The pseudo-mode ridges are close to the mode ridges (so that the ridges look continuous) because the excitation sources are located very close to the surface where the oscillations are observed. The pseudo-mode frequencies depend on details of the excitation mechanism and on the wave interaction with the solar atmosphere. Therefore, so-far, only the normal modes have been used for the reconstruction of solar subsurfaces. The primary restriction of global helioseismology is that it can only reconstruct the azimuthally averaged properties of the interior. This is not sufficient for the understanding of the solar dynamics and magnetism. The three-dimensional structure of the solar subsurfaces can be reconstructed by techniques of local helioseismology. One of these techniques, called `ring-diagram analysis' \citep{Gough1983} is based on measuring frequency shifts in local (typically $15\times 15$ degrees) areas, and uses the global helioseismology description of the mode frequency sensitivity to local sound-speed variations and flows. This techniques allows us to reconstruct the solar subsurfaces with relatively low spatial resolution in shallow regions. The reconstruction with higher spatial resolution and much deeper in the interior can be achieved by methods based on extracting coherent wave signals and measuring variations of the wave travel times or phase shift. These techniques called time-distance helioseismology \citep{Duvall1993} and acoustic holography \citep{Lindsey2000} employ cross-covariance functions of solar oscillations instead of the power spectral analysis. The discovery that coherent signals, such as wave packets, can be extracted from the cross-covariance functions of the stochastic solar oscillations (Fig.~\ref{fig1}b) was made by Duvall \citep{Duvall1993}. This approach was then developed in helioseismology, terrestrial seismology, and other disciplines, and in broader applications is called `ambient noise imaging'. The foundation of this approach is based on the property of cross-covariance functions to represent wave signals corresponding to point sources. Roughly speaking the cross-covariance function can be considered as the Green's function of the solar wave equation. In real solar conditions this is only an approximation because of the limited frequency bandwidth of solar oscillations and inhomogeneities of the solar structures and distribution of the stochastic sources. A complete theory of this approach of helioseismology has not been developed. It requires extensive studies of wave interaction with turbulence, flows and magnetic field. Nevertheless, the initial results based on relatively simple descriptions of wave propagation have provided important insights into the three-dimensional structures and flow patterns of the solar subsurfaces. The primary focus of these studies is mapping the flow patterns associated with the solar cycle, and formation and evolution of active regions. In the current state of local helioseismology the systematic errors as well as effects of the stochastic realization noise have not been fully investigated. These studies require substantial effort for modeling the wave dynamics in realistic solar conditions, and require 3D MHD simulations on large supercomputer systems. The validation and testing of the time-distance technique have been performed by comparing the helioseismic inversions in the shallowest layer with the surface flows obtained by a local correlation tracking technique \citep{Liu2013}, and through the analysis and inversion of numerical simulation data for subsurface flows and sound-speed variations \citep{Birch2011,Hartlep2013,Parchevsky2009,Parchevsky2014}. The testing for regions with strong magnetic field has not yet been completed. However, the simulations of the wave propagation in sunspot models showed that one of the primary effects in sunspot regions is the wave reflection from deeper layers, compared to the quiet-Sun regions, where the plasma parameter, $\beta=8\pi P/B^2$, the ratio of the gas pressure to magnetic pressure, is equal to unity (this layer also corresponds to the deeper photospheric surface of sunspots, known as the Wilson depression). Below the Wilson depression level the gas pressure dominates, and the helioseismic acoustic waves behave like fast MHD waves: the wave speed becomes anisotropic, and also depends on the temperature stratification beneath sunspots. The magnetic and temperature effects have not been separated in the wave-speed inversion results \citep{Kosovichev2000}. This is an important task of local helioseismology. One of the difficulties is that the limited computer power has not allowed to simulated the sunspot models sufficiently large and deep for the helioseismology testing, so that the wave properties are not affected by the boundary conditions of the simulations. A comparison of the wave-speed inversions obtained for a sunspot regions by the time-distance and ring-diagram techniques (see \citep{Kosovichev2012,Kosovichev2011} and references therein) shows a good qualitative agreement: both inversions show a two-layer structure with a layer of reduced wave speed beneath the surface followed by a layer of an increased wave speed. However, the depth of these layers is different, perhaps, because of the drastically different spatial resolution, and different contributions of magnetic field. A comparison of the time-distance and acoustic holography inversion results has been performed by using artificial simulation data \citep{Birch2011,Parchevsky2014}. \begin{figure}[t] \begin{center} \includegraphics[width=\linewidth]{kosovichev_figure_02.eps} \caption{ a) A scheme of the Time-Distance Helioseismology Pipeline implemented at the Joint Science Operations Center (JSOC) for Solar Dynamics Observatory at Stanford University \citep{Zhao2012}; b) Illustration of the surface locations of the individual patches used for inferences of the subsurface structure and flows; the total 25 patches are used to cover $120\times 120$ degrees of the disk area. }\label{fig2} \end{center} \end{figure}

The initial analysis of the subsurface flow maps reconstructed from the five years of SDO/HMI data by time-distance helioseismology reveals the great potential for studying and understanding the dynamics of the quiet-Sun and active regions, and the evolution with the solar cycle. In particular, our results show that the emergence and evolution of active regions are accompanied by multi-scale flow patterns. Beneath the sunspot, during their formation, we observe appearance of flows converging towards the sunspot center. and also the 'moat'-like flows diverging from the active region in the surrounding regions. On the larger scale, revealed by averaging the high-resolution flow maps, we find a pattern of flows converging towards the active region. This pattern is formed when the active region is fully developed. On the global-Sun scale, the flow maps allow us to investigate the structure and evolution of the meridional flows. In particular, we find that the meridional flows display the North-South asymmetry closely correlating with the magnetic activity. The latitudinal variations of the meridional circulation speed, which are probably related to the large-scale converging flows, are mostly confined in a shallow subsurface layers. Therefore, these variations do not necessarily affect the magnetic flux transport. The North-South asymmetry is also pronounced in the variations of the differential rotation ('torsional oscillation'). The calculations of a proxy of the subsurface kinetic helicity density show that the helicity does not vary during the solar cycle, and that the supergranulation is a likely source of the near-surface helicity. These initial results are obtained from the analysis of a small sample of flow maps produced by the SDO/HMI time-distance helioseismology pipeline. Further detailed investigations are required for understanding the complicated subsurface dynamics of the Sun.

1607.05681

1607

1607.04963_arXiv.txt

Huge astrospheres or stellar wind bubbles influence the propagation of cosmic rays at energies up to the TeV range and can act as small-scale sinks decreasing the cosmic ray flux. We model such a sink (in 2D) by a sphere of radius 10\,pc embedded within a sphere of a radius of 1\,kpc. The cosmic ray flux is calculated by means of backward stochastic differential equations from an observer, which is located at $r_{0}$, to the outer boundary. It turns out that such small-scale sinks can influence the cosmic ray flux at the observer's location by a few permille (i.e\ a few 0.1\%), which is in the range of the observations by IceCube, Milagro and other large area telescopes.

Large area cosmic ray detectors like the Tibet Air shower experiment, IceCube/IceTop, Milagro and HAWC, among others observe a multipole like anisotropy of the high energy cosmic ray flux (CRF) over the entire sky \citep{Iuppa-etal-2013,Abeysekara-etal-2014,BenZvi-2014,Desiati-2014,DiSciascio-Iuppa-2014,diSciascio-2015,Lopez-Barquero-etal-2015,Icecube-etal-2016}. The energies of interest are in the TeV range which vary on small-scales by a few permille (\permille) for details see \citet{Toscano-etal-2012} and \citet{Iuppa-DiSciascio-2012}. At higher energies (PeV) anisotropies were also found \citep{Giacinit-Sigl-2012,Zotov-Kulikov-2012,Aartsen-etal-2013,Glushkov-Pravdin-2013} which may be still of Galactic origin. Even at energies around EeV, anisotropies in the arrival directions are obsereved \citep[e.g.\ ][]{Abreu-etal-2011,Abreu-etal-2013}, which may be at the transition to an extragalactic origin of the cosmic rays. These high energies are not taken under consideration here. A few explanations have been proposed, either related to interstellar magnetic field variations \citep{Amenomori-etal-2011}, intermediate turbulence \citep{Biermann-etal-2015} due to the heliotail \citep{Desiati-Lazarian-2014,Zhang-etal-2014,Pogorelov-etal-2015,Schawdron-etal-2015}. A detailed analysis of the power spectrum is discussed in \citet{Ahlers-Mertsch-2015} where thet authors showed that the strength of the power spectrum is related to the diffusion tensor. \citet{Harari-etal-2016} discussed turbulent magnetic fields as the cause of small angular scale variations, while \citet{Battaner-etal-2015} correlated the anisotropy to the global cosmic ray flux. There are small-scale variations (tens of degrees) and even tiny-scale variations about a degree or less in the interstellar medium caused by astrospheres, planetary nebulae and similar inhomogeneities \citep{Stanimirocic-etal-2010,Haverkorn-Spangler-2013,Linsky-Redfield-2014}. In the following we explain how such variations -- due to the presence of astrospheres -- act as small-scale sinks (\s) of CRF in the interstellar medium, can lead to such anisotropies. Huge astrospheres or stellar wind bubbles have been discussed for example in \citet{Mackey-etal-2015} and \citet{Scherer-etal-2016}, especially their influence to the CRF was studied in \citet{Scherer-etal-2015a} for the case of $\lambda$ Cephei. The latter authors found that the CRF at energies up to 100\,TeV is affected on scales below 1\,pc along the stagnation line of the astrosphere. Because the discussed astrosphere of $\lambda$ Cephei is special, in the sense that the relative motion between the star and the ISM is high (about 80\,km/s), and the bow shock distance is about 1\,pc. Most of the astrospheres of hot stars do not show any relative motion and build stellar wind bubbles of the order of 10-100\,pc. These bubbles have very high compression ratios \citep{Toala-Arthur-2011} and thus can effectively act as sinks for the CRF. The CRF is already affected directly beyond the bowshock, see \citet{Scherer-etal-2011}, \citet{Strauss-etal-2013}, and \citet{Luo-etal-2015}, and thus the effective modulation in astrospheres starts directly behind the bow shock, and not as in the helioshere at the heliopause \citep{Kota-Jokipii-2014,Potgieter-2014}. Here we setup a model where we study the transport of cosmic rays (CRs), when there is a \s between the outer boundary and the observer. In section~\ref{sec:1} we present the model in detail, while in section~\ref{sec:2} the numerical scheme is discussed. In section~\ref{sec:3} we study the results for a range of appropriate parameters (given in Table~\ref{tab:1} below) and finally we give a r\'esum\'e in section~\ref{sec:4}

1607.04963

1607

1607.06682_arXiv.txt

Spectra of comet C/2014 Q2 (Lovejoy) were taken with a low resolution spectrograph mounted on the 0.5 m telescope at the Mount Abu Infrared Observatory (MIRO), India during January to May 2015 covering the perihelion and post-perihelion periods. The spectra showed strong molecular emission bands (C$_{2}$, C$_{3}$ and CN) in January, close to perihelion. We have obtained the scale lengths for these molecules by fitting the Haser model to the observed column densities. The variation of gas production rates and production rate ratios with heliocentric distance were studied. The extent of the dust continuum using the $Af\rho$ parameter and its variation with the heliocentric distance were also investigated. The comet is seen to become more active in the post-perihelion phase, thereby showing an asymmetric behaviour about the perihelion.

Comets are cold icy bodies in the solar system that were formed in the solar nebula and are considered to be the signature bodies to understand the formation of the solar system. As the comet nucleus makes its journey towards its parent star, the ices start sublimating giving rise to a mixture of gas and dust which form the coma. For comets at heliocentric distances less than 3 AU, the visible band spectrum shows strong molecular emission bands riding on the continuum radiation scattered by the cometary dust. Studying these molecular emission bands has been an important part of the cometary study. \\ Comet C/2014 Q2 (Lovejoy), an Oort cloud comet, was discovered by Terry Lovejoy in August 2014 using an 8 inch telescope when the comet had a visual magnitude of 15. The comet brightened to a magnitude of 4 in the visual band in January 2015 being closest to Earth at 0.469 AU on January 7. The comet's perihelion was on the 30th of January 2015 at a heliocentric distance of 1.29 AU. According to the data from JPL Horizons database, the comet orbit has an eccentricity of 0.9976, orbital inclination of 80.3$^{\degree}$, semi-major axis of 576.34 AU and orbital period of 13700 years. \\ \citet{b1} have carried out a detailed survey of 85 comets during the years 1976-1992. They have studied the variation of gas production rates and dust to gas ratio with heliocentric distances. They have defined certain limits in the gas ratios in order to classify the comets as to whether they are depleted in carbon-chain molecules or not. According to their study, most of the carbon depleted comets are from the Jupiter family. They also say that CN is produced from the nucleus as well as from the dust in the comet's coma. Several spectroscopic surveys have been carried out, e.g. \citet{Newburn}, \citet{Fink_2009}, \citet{b2}. \citet{b2} have spectroscopically studied five comets. They found a linear correlation between production rates of C$_{2}(\Delta \nu=0)$, C$_{2}(\Delta \nu=1)$ and C$_{3}$ with respect to CN. No correlation was found between the production rate ratios and heliocentric distance. \citet{Langland} have given new statistical methods in cometary spectroscopy for the extraction and subtraction of the sky using the comet frames itself. They have given a dynamical classification of comets by studying 26 different comets. More recently, \citet{Cochran2012} have compiled and reported the spectroscopic results for 130 comets observed from McDonald Observatory. They have found remarkable similarity in the composition of most of the comets. They quote that the carbon chain depleted comets can be from any dynamical class. However they have not found any of them from the Halley type comets (Tisserand parameter $<$ 2 and period $<$ 200 years) due to significantly low number of these types in their sample. \\ In our study of comet C/2014 Q2, we have determined the production rates for various gas species, $Af\rho$ and dust to gas ratios at different heliocentric distances.

We have carried out spectral study of the comet C/2014 Q2 (Lovejoy) which shows strong emission lines of C$_{2}$, C$_{3}$ and CN. The molecular production rates and the quantity Af$\rho$ are estimated. We arrive at the following conclusions: \begin{itemize} \item The comet shows a large positive asymmetry, which indicates increase in the post-perihelion activity. \item Considering the limits set by \citet{b1} for the carbon abundance, C/2014 Q2 is found to be in the typical class of comets. \item The dust to gas ratio remains fairly constant within the observed heliocentric range which is in agreement with \citet{b1} \item The $Af\rho$ values indicate that the dust distribution favours smaller dust particles near perihelion which depletes as the comet moves away. \item New Haser model scale lengths were obtained which best fit the observed data. These differ significantly from the previously quoted values. This might be due to the comet's intrinsic nature itself or affected by the solar activity at the time of observation. \end{itemize} A more exhaustive study is required to confirm some of the conclusions drawn here.

1607.06682

1607

1607.04687_arXiv.txt

We present results of an investigation into the formation of nitrogen-bearing molecules in the atmosphere of Titan. We extend a previous model \citep{li14,li15} to cover the region below the tropopause, so the new model treats the atmosphere from Titan's surface to an altitude of 1500 km. We consider the effects of condensation and sublimation using a continuous, numerically stable method. This is coupled with parameterized treatments of the sedimentation of the aerosols and their condensates, and the formation of haze particles. These processes affect the abundances of heavier species such as the nitrogen-bearing molecules, but have less effect on the abundances of lighter molecules. Removal of molecules to form aerosols also plays a role in determining the mixing ratios, in particular of HNC, HC$_3$N and HCN. We find good agreement with the recently detected mixing ratios of C$_2$H$_5$CN, with condensation playing an important role in determining the abundance of this molecule below 500 km. Of particular interest is the chemistry of acrylonitrile (C$_2$H$_3$CN) which has been suggested by \cite{slc15} as a molecule that could form biological membranes in an oxygen-deficient environment. With the inclusion of haze formation we find good agreement of our model predictions of acrylonitrile with the available observations.

A major goal of planetary exploration is to obtain a fundamental understanding of planetary environments, both as they are currently and as they were in the past. This knowledge can be used to explore the questions of (a) how conditions for planetary habitability arose and (b) the origins of life. Titan is a unique object of study in this quest. Other than Earth itself, and Pluto \citep[which has also been observed to have photochemically produced haze;][]{stern15,gladstone16}, Titan is the only solar system body demonstrated to have complex organic chemistry occurring today. Its atmospheric properties--- (1) a thick N$_2$ atmosphere, (2) a reducing atmospheric composition, (3) energy sources for driving disequilibrium chemistry and (4) an aerosol layer for shielding the surface from solar UV radiation---suggest it is a counterpart of the early Earth, before the latter's reducing atmosphere was eradicated by the emergence and evolution of life \citep{ct99,lunine05,lm08}. A significant number of photochemical models have been developed to investigate the distribution of hydrocarbons in Titan's atmosphere \citep{strobel74,yap84,lara96,wa04,delahaye08,lavvas08a,lavvas08b,krasnopolsky09}. Recently, more constraints have been placed on the abundance of hydrocarbons and nitriles in the mesosphere of Titan (500 -- 1000 km) from Cassini/UVIS stellar occultations \citep{koskinen11,kammer13}. In combination with the updated version of Cassini/ CIRS limb view \citep{vinatier10}, the complete profiles of C$_2$H$_2$, C$_2$H$_4$, C$_6$H$_6$, HCN, HC$_3$N are revealed for the first time. C$_3$ compounds, including C$_3$H$_6$, were modeled by \cite{li15}, and the agreement with observations \citep{nixon13} is satisfactory. The chemistry of many of these nitrogen molecules has recently been modeled by \cite{loison15}. In this paper we introduce our updated Titan chemical model that includes the formation of such potentially astrobiologically important molecules as acrylonitrile. In addition to the usual gas phase chemistry, it also includes a numerically stable treatment of the condensation and sublimation, allowing the formation and destruction of ices in the lower atmosphere to be tracked. Haze formation is also included in a parameterized fashion, allowing for the permanent removal of molecules from the atmosphere. We present here the effects of condensation on the nitrogen chemistry. The interaction of hydrocarbons and nitrile species in the condensed phase is complex and is beyond the scope of this paper \citep[see, for example, Figures 1 and 2 of][]{anderson16}. We begin with describing our updated model and in particular our treatment of condensation and sublimation (Section~\ref{sec:model}). We use this updated model to consider the chemistry in Titan's atmosphere from the surface of the moon to an altitude of 1500 km. We explore how condensation processes and haze formation affect the predicted gas phase abundances of observable molecules (Section~\ref{sec:results}). We also consider where the condensates form within the atmosphere (Section~\ref{sec:cond}). Section~\ref{sec:conc} presents our conclusions.

\label{sec:conc} The removal of molecules by condensation plays an important role in determining the gas phase composition of Titan's atmosphere, as well as creating new aerosols. Condensates are found throughout the atmosphere. For the majority of molecules, condensation is most efficient below the tropopause. Larger molecules, and in particular nitrogen-bearing molecules have another condensation peak between 200 and 600 km. Relatively high abundances of condensates can also be present above 500 km if the gas phase abundance of a given molecule is high, e.g.\ HC$_3$N, HC$_5$N, CH$_3$CN and C$_2$H$_5$CN. These molecules condense in the region where Titan's haze forms. The effect is enhanced if it is assumed that some molecules can be permanently removed from the gas by being incorporated into aerosol particles. This mechanism was able to bring the abundances of HC$_3$N, HCN, HNC, CH$_3$CN and C$_2$H$_5$CN into good agreement with the observations below 600 km. Although Titan possesses a rich organic chemistry it is unclear whether this could lead to life. Photochemically produced compounds on Titan, principally acetylene, ethane and organic solids, would release energy when consumed with atmospheric hydrogen, which is also a photochemical product. \cite{mckaysmith05} speculate on the possibility of widespread methanogenic life in liquid methane on Titan. On Earth fixed nitrogen is often a limiting nutrient. Our work shows that an abundant supply of fixed nitrogen, including species of considerable complexity, is available from atmospheric photochemistry. Creating the kinds of lipid membranes that form the basis of lie on Earth depends on the presence of liquid water. Titan's atmosphere contains little oxygen and the surface temperature is well below that at which liquid water can survive. Instead surface liquids are hydrocarbons \citep[e.g.][]{hayes16}. Therefore any astrobiological processes, if present, are likely to be quite different to those on Earth. A recent paper by \cite{slc15} suggests that as alternative to lipids, membranes could be formed from small nitrogen-bearing organic molecules such as acrylonitrile (C$_2$H$_3$CN). Stevenson et al.\ calculate that a membrane composed of acrylonitrile molecules would be thermodynamically stable at cryogenic temperatures and would have a high energy barrier to decomposition. All of our models predict abundances of C$_2$H$_3$CN that are in agreement with observations above 500 km. Below this condensation and incorporation into haze are required to bring the predicted mixing ratios down to the values inferred from observations \cite{cordiner15}. If acrylonitrile were to be involved in life formation it needs to reach the surface of Titan. Our predicted flux of this molecule onto Titan's surface is 1.5 \x 10$^7$ molecules cm$^{-2}$ s$^{-1}$, or $\sim$ 41.5 gcm$^{-2}$/Gyr, a quantity that is potentially of biological importance.

1607.04687

1607

1607.03919_arXiv.txt

We show that the angular momentum exchange mechanism governing the evolution of mass-transferring binary stars does not apply to Roche lobe filling planets, because most of the angular momentum of the mass-transferring stream is absorbed by the host star. Apart from a correction for the difference in specific angular momentum of the stream and the centre of mass of the planet, the orbit does not expand much on Roche lobe overflow. We explore the conditions for dynamically unstable Roche lobe overflow as a function of planetary mass and mass and radius (age) of host star and equation of state of planet. For a Sun-like host, gas giant planets in a range of mass and entropy can undergo dynamical mass transfer. Examples of the evolution of the mass transfer process are given. Dynamic mass transfer of rocky planets depends somewhat sensitively on equation of state used. Silicate planets in the range $1 \ M_{\bigoplus} <M_{\mathrm{p}} < 10 \ M_{\bigoplus} $ typically go through a phase of dynamical mass transfer before settling to slow overflow when their mass drops to less than $1 \ M_{\bigoplus}$.

Though constituting only a small fraction of the total exoplanet population, exoplanets orbiting close to their host stars pose interesting challenges to theoretical models for the formation and evolution of planetary systems. Since the hosts generally rotate more slowly than the planet orbit, tidal interaction causes the planets to lose angular momentum. Depending on the somewhat uncertain strength of tidal friction (for a review see Ogilvie 2014) the planets on the closest observed orbits, on the order of a few days, would spiral into their host within a few billion years (e.g. Raymond, Barnes \& Mandell 2008; Levrard, Winisdoerffer \& Chabrier 2009; Jackson, Barnes \& Greenberg 2009). The loss of planets by spiral-in has been invoked to explain the orbital distribution of close-in exoplanets (Jackson et al. 2009) and the dearth of close-in exoplanets around fast rotating stars (Teitler \& K\"onigl 2014) from the observations (McQuillan, Mazeh \& Aigrain 2013). The final fate of an angular momentum losing planet depends on its mass, mean density, and composition. It also depends sensitively on the details of the angular momentum balance during the Roche overflow process, for which different assumptions have been made in previous work. In several studies it is assumed that the planet is rapidly disrupted once it fills its Roche lobe, and the material of planet is then accreted on to host star (e.g. Jackson et al. 2009; Rappaport et al. 2013; Teitler \& K\"onigl 2014). Metzger, Giannios \& Spiegel (2012), on the other hand, found that the mass transfer of the planet-star (hot Jupiter) system will be stable, occurring on the slow tidal evolution time-scale, if $ \bar{\rho}_{\mathrm{p}} / \bar{\rho}_\star \lesssim 1 $, where $\bar{\rho}_{\mathrm{p}}$ and $\bar{\rho}_\star$ are mean density of the planet and that of the host star, respectively. The possible outcomes of the spiral-in process of hot Jupiter are conveniently classified into three cases (Metzger et al. 2012). With decreasing orbital separation the planet can either reach its Roche limit and disrupt before physically entering the star, or it can spiral in `whole'. In the former case the planet loses mass either as a slow process governed by the orbital evolution under tidal interaction, or dynamically, evolving on an orbital time-scale. Whether Roche lobe overflow takes place before reaching the stellar surface depends on ratio of the mean density of the planet $\bar{\rho}_{\mathrm{p}}$ to that of the star $\bar{\rho}_\star$. If $\bar{\rho}_{\mathrm{p}} / \bar{\rho}_\star \lesssim 5 $, the planet reaches its Roche limit outside the host star. If the planet has a higher mean density, $\bar{\rho}_{\mathrm{p}} / \bar{\rho}_\star \gtrsim 5 $, it would fill its Roche lobe only below the stellar surface. In this case a direct merger occurs between the planet and the host star (Metzger et al. 2012). As in the case of mass transfer in binary stars (Paczy{\'n}ski 1971; Frank, King \& Raine 1992), the time-scale on which Roche lobe overflow takes place depends critically on the adiabatic mass-radius relation of the planet. If loss of mass causes the radius of the planet to decrease in size more slowly than the Roche radius, mass transfer is unstable on a dynamical time-scale (`dynamical mass transfer', Paczy{\'n}ski 1965, Paczy{\'n}ski, Zi{\'o}lkowski \& Zytkow 1969). The final disruption of the planet then takes place in a few orbits (e.g.\ Rasio 1994). In the opposite case mass loss is dynamically stable. The loss rate through inner Lagrangian point (L1) is then far slower, governed by the time-scale of orbital decay through tidal interaction with the host star. In both cases, the material lost through the L1 either forms a disc, or hits the stellar surface and is accreted on to host star directly (as in the case of many Algol type binaries). Which is the case depends on the mass ratio $M_{\rm p}/M_*\equiv q$ of the planet's mass $M_{\rm p}$ to the star's mass $M_*$, and on the radius $R_*$ of the host star.

\label{summary} We have investigated the mass transfer from a planet to its host star, with emphasis on the conditions under which Roche lobe overflow becomes dynamically unstable. The main factors are the adiabatic mass-radius relation of the planet and the processes redistributing angular momentum. Previous analyses of mass transfer were based on the analogy with mass-transferring binary systems (CVs and X-ray binaries) where the moment of inertia of the host star can be neglected, and the angular momentum of the transferred mass is returned to the orbit by tidal interaction with an accretion disc\footnote{ In the \textsc{mesa} code, the default value `Ritter' for the parameter `mdot\_scheme' assumes this, as in Ritter (1988).}. We have argued that this is not an appropriate assumption in the case of Roche lobe overflow of a planet on a MS (or larger) star. The orbital angular momentum of a planet is far too small to affect the rotation of such stars, which can effectively act as an arbitrary sink of angular momentum. The consequence is that planets are much more likely to go through dynamically unstable Roche lobe overflow than predicted by the standard description that applies to interacting binaries. The more massive planets touch the host star surface before overflowing their Roche lobes. These planets consequently spiral into their host on a short time-scale (we call this the `direct merger' case). This is especially likely to happen to planets (if any) that formed close to their host during their PMS life (see Figs.\ \ref{solargiant}-\ref{solargiants7} and \ref{tin1}, \ref{tin2}). At some what lower planetary masses or larger radius of host stars, the mass-transferring stream intersects the surface of the star before completing an orbit. In this simple case (which we call the `stream impact' case), the specific angular momentum of the mass lost by the planet through the L1 is a bit less than that of the planet's centre of mass. This would cause a slight expansion of the orbit. In section \ref{discplanet} we have argued that even if the stream does not impact directly, the truncation radius of the prospective accretion disc is probably only marginally outside the stellar surface of a MS host star. Instead of forming a dics, we find it more likely that the hydrodynamic interaction with the slowly rotating stellar envelope is effective enough to absorb most the accreting angular momentum instead of leading to a spreading dics. We have called this the `minimal assumption', but have compared its consequences with different assumptions. Whether Roche lobe overflow of a planet is dynamically stable and unstable depends on its (adiabatic) mass-radius relation, which differs strongly for gas giant planets and rocky planets (see Figs.\ \ref{zamsgiant}, \ref{zamssilicate}). The areas in the $M_{\rm p}-M_*$ plane differ correspondingly. For low mass rocky planets, the mass-radius index is close to the value $1/3$ for a nearly constant mean density. Since this is also the approximate mass-radius index of the Roche lobe of a low-mass companion, Roche lobe overflow is near marginal stability for the `minimal assumption' (see Fig.\ \ref{silicams}). The more massive rocks are slightly compressible. As a result, their mass transfer is initially unstable, but settles to stable overflow when the mass has decreased to less than $1\,M_{\bigoplus}$ (Fig.\ \ref{mtsilicate}). The higher likelihood of dynamically unstable Roche lobe overflow in our `minimal assumption' increases the change of observing a planet in the process of rapid mass transfer. The low transfer rates driven only by tidal interaction in stable overflow may be hard to recognize observationally.

1607.03919

1607

1607.08526_arXiv.txt

This paper presents the first major data release and survey description for the ANU WiFeS SuperNovA Program (AWSNAP). AWSNAP is an ongoing supernova spectroscopy campaign utilising the Wide Field Spectrograph (\wifes) on the Australian National University (ANU) 2.3m telescope. The first and primary data release of this program (AWSNAP-DR1) releases \nspec\ spectra of \nobj\ unique objects collected over \nnights\ equivalent full nights of observing from July 2012 to August 2015. These spectra have been made publicly available via the WISEREP supernova spectroscopy repository. We analyse the AWSNAP sample of Type Ia supernova spectra, including measurements of narrow sodium absorption features afforded by the high spectral resolution of the \wifes\ instrument. In some cases we were able to use the integral-field nature of the \wifes\ instrument to measure the rotation velocity of the SN host galaxy near the SN location in order to obtain precision sodium absorption velocities. We also present an extensive time series of SN~2012dn, including a near-nebular spectrum which both confirms its ``super-Chadrasekhar'' status and enables measurement of the sub-solar host metallicity at the SN site.

In the last decade, wide-field extragalactic transient surveys -- such as the Palomar Transient Factory \citep[PTF;][]{rau09, law09}, the Panoramic Survey Telescope and Rapid Response System \citep[PanSTARRS;][]{panstarrs}, the Catalina Real-time Transient Survey \citep[CRTS;][]{crts}, and the Texas Supernova Search \citep{quimbythesis, yuanthesis} -- have revolutionised our understanding of the myriad ways in which stars explode through the discovery of new classes of exotic transients. Simultaneously, these surveys have discovered hundreds of supernovae (SNe) of ``traditional'' types \citep[see][for a review]{filippenko97}, enabling statistical analyses of the properties of these SNe. While imaging surveys have provided discovery and light curves for this wealth of new transients, complementary spectroscopy surveys have provided the critical insight into the physical origins of these events. Numerous supernova spectroscopy surveys have released thousands of high quality spectra of nearby SNe into the public domain \citep{matheson08, blondin12, bsnip1, folatelli13, modjaz14}. These surveys have frequently been dedicated to the spectroscopic followup of Type Ia supernovae (\sneia) which, due to their rates and luminosities, dominate any magnitude-limited imaging survey. Such surveys have revealed that photometrically similar SNe can still exhibit diversity of spectroscopic behaviour, indicating spectra remain a critical tool for revealing the full nature of the supernova progenitors (particularly for \sneia). Additionally, spectra remain critical for supernova classification -- particularly at early phases when the full photometric evolution has yet to be revealed. Such early classifications then inform the use of additional SN followup facilities, including those operating outside the optical window. Recently the Public ESO Spectroscopic Survey for Transient Objects \citep[PESSTO;][]{pessto} began a multi-year program on the NTT 3.6m telescope in Chile, with the goal of obtaining high quality spectral time series for roughly 100 transients (of all kinds) to be released to the public. This survey has already released hundreds of spectra in its first two annual data releases, and continues to release all SN classificaion spectra within typically 1 day from observation. Other ongoing SN spectroscopy programs, such as the Asiago Supernova Program \citep{tomasella14}, also make important contributions to the transient community through timely SN classification and spectroscopy releases. Here we describe our ongoing spectroscopy program AWSNAP -- the {\bf A}NU {\bf W}iFeS {\bf S}uper{\bf N}ov{\bf A} {\bf P}rogram -- which uses the Wide Field Spectrograph \citep[WiFeS;][]{dopita07, dopita10} on the Australian National University (ANU) 2.3m telescope at Siding Spring Observatory in Australia. In this paper we describe the data processing procedures for this ongoing program, and describe the first AWSNAP data release (AWSNAP DR1) comprising \nspec\ spectra of \nobj\ supernova of various types obtained during \nnights\ classically-scheduled observing nights over a 3 year period from July 2012 to August 2015. Most of these spectra have been released publically via the Weizmann Interactive Supernova data REPository \citep[WISeREP\footnote{http://wiserep.weizmann.ac.il} --][]{wiserep}, with the remainder set to be released within the next year as part of forthcoming PESSTO papers. This program will continue to observe SNe of interest and classify SN discoveries from transient searches such as the new SkyMapper Transients Survey \citep{keller07}. We aim to release future SN classification spectra from AWSNAP publicly via WISeREP in parallel with any classification announcements. This paper is organised as follows. Section~\ref{sec:data} describes the \wifes\ data processing and SN spectrum extraction procedures. Section~\ref{sec:sn_sample} presents general properties of our SN sample and compares AWSNAP DR1 to other public SN spectra releases. In Section~\ref{sec:results} we present some analysis of the properties of the \sneia\ in our sample, including measurement of narrow sodium absorption features afforded by the intermediate resolution of the \wifes\ spectrograph. Some concluding remarks follow in Section~\ref{sec:conclusions}.

\label{sec:conclusions} This work marks the primary data release for the ANU WiFeS SuperNovA Program (AWSNAP), comprising \nspec\ distinct spectra of \nobj\ unique supernovae. These data were collected using the Wide Field Spectrograph (WiFeS) on the ANU 2.3m telescope during \nnights\ nights of observing over a three-year period from mid-2012 to mid-2015. The AWSNAP spectroscopy sample is comparable in size to other SN spectra data releases, and its composition of SN types is roughly in line with expectations for a magnitude-limited SN search. The phase coverage of the AWSNAP \snia\ sample is comparable to other published \snia\ spectroscopy datasets (for \sneia\ with multiple epochs of observation), with the inclusion of more \sneia\ with a single observation (i.e., classification spectra only). We presented some analyses of the AWSNAP \snia\ sample, including some results uniquely enabled by the fine wavelength resolution available with \wifes. We measured broad absorption features in \snia\ spectra at maximum light, including the ratio of silicon absorption features $R_{Si}$ and the strength of high velocity features $R_{HVF}$. Additionally, we measured the strength and velocity of narrow sodium absorption features, including some cases where the integral-field nature of the instrument allowed us to measure the local systemic velocity within the SN host galaxy. Some expected feature trends, such as a correlation between $R_{HVF}$ and $R_{Si}$, were recovered in our data set. The nature of sodium absorption in our sample was limited by small number statistics. Finally, we presented our observations of the candidate super-Chandrasekhar \snia\ SN~2012dn, and used narrow host galaxy emission features to show the SN site exhibits sub-solar metallicity. The \wifes\ instrument presents several unique advantages for the study of transients, particularly owing to its comparatively narrow velocity resolution ($\sigma_v \sim 45$~\kms). It has previously been employed in the study of SNe with strong narrow emission features such as SN~2009ip \citep{fraser13, fraser15}, SN~2012ca \citep{inserra14, inserra16}, and SN~2013fc \citep{kangas16}. Here we also demonstrated that the fine velocity resolution allows for the measurement of narrow {\em absorption} features, particularly sodium absorption in \sneia. The higher resolution of \wifes\ also frequently revealed narrow host galaxy emission features at the site of the SN, which can at times be used to determine a SN site metallicity (as we showed for SN~2012dn). Finally, the integral field nature of \wifes\ allowed us to measure the local host galaxy rotational velocity at the site of several SNe, even when there was no emission directly at the SN location. Thus \wifes\ is an instrument well-suited for not only standard SN spectroscopic observations, but also a unique suite of capabilities not commonly found in transient followup instruments.

1607.08526

1607

1607.04708_arXiv.txt

Dust trapping accelerates the coagulation of dust particles, and thus it represents an initial step toward the formation of planetesimals. We report $H$-band (1.6 $\mu$m) linear polarimetric observations and 0.87 mm interferometric continuum observations toward a transitional disk around LkH$\alpha$ 330. As results, a pair of spiral arms were detected in the $H$-band emission and an asymmetric (potentially arm-like) structure was detected in the 0.87 mm continuum emission. We discuss the origin of the spiral arm and the asymmetric structure, and suggest that a massive unseen planet is the most plausible explanation. The possibility of dust trapping and grain growth causing the asymmetric structure was also investigated through the opacity index ($\beta$) by plotting the observed SED slope between 0.87 mm from our SMA observation and 1.3 mm from literature. The results imply that grains are indistinguishable from ISM-like dust in the east side ($\beta = 2.0\pm0.5$), but much smaller in the west side $\beta = 0.7^{+0.5}_{-0.4}$, indicating differential dust size distribution between the two sides of the disk. Combining the results of near-infrared and submillimeter observations, we conjecture that the spiral arms exist at the upper surface and an asymmetric structure resides in the disk interior. Future observations at centimeter wavelengths and differential polarization imaging in other bands (Y to K) with extreme AO imagers are required to understand how large dust grains form and to further explore the dust distribution in the disk.

Dust grains are the raw material for planetesimals, which over time form rocky planet cores. The widely accepted core accretion scenario implies that coagulation of micron-sized dust grains into larger bodies is a first step for planet formation. Thus, it is important to investigate the grain size and its distribution in the birthplaces of planets: protoplanetary disks. If the emission is optically thin, the opacity index ($\beta$) can be used as an observational indication of grain size (possibly grain growth), where $\kappa_d$ $(\propto \lambda^{-\beta})$ is the dust opacity and $\lambda$ is the wavelength. A number of observations reveal that the value of $\beta$ is generally low for circumstellar disks around low and intermediate-mass young stellar objects (YSOs) \citep{beckwith91,dalessio01,andrews05,draine06,ricci10a,ricci10b,guilloteau11,menu14}. The exploration of radial and azimuthal dust distribution in various sizes of dust is essential for revealing where and how dust grain coagulation takes place. Since the $\beta$ offers suggestive evidence for the amount of dust grains whose size is approximately comparable to the wavelength of observation, the analysis of $\beta$ and its distribution help understand the initial condition for the subsequent formation of planetesimals. We have conducted $H$-band linear polarimetric observations that trace (sub)micron-sized dust at the disk scattering surface and 0.87 mm dust continuum observations that detect the thermal emission from approximately millimeter-sized grains around the disk mid-plane. LkH$\alpha$ 330 is a G3 star \citep{cohen79} of 2.5 M$_\sun$ \citep{salyk09} and 15 L$_\sun$ \citep{andrews11} located in the Perseus molecular clouds, which is about 250 pc away from the Sun \citep{enoch06}. The disk of LkH$\alpha$ 330 is one of the 4 transitional disks found in the Spitzer Space Telescope Cores to Disks (c2d) Legacy Program, where they took $5-35$ $\mu$m spectra with the Infrared Spectrograph (IRS) for 100 stellar systems \citep{brown07}. Later, \citet{brown08,brown09} constructed the first resolved submillimeter image of the disk with an angular resolution of 0\farcs3 using the Submillimeter Array (SMA) at 340 GHz. This observation directly confirmed the cavity in the dust disk. Then, a two-dimensional Monte Carlo radiative transfer code \citep[RADMC;][]{dullemond04} was used to model the disk. In this model, the disk is flared, density is a power law of radius, and the cavity is represented by the inner and outer gap radius and a density reduction factor. In addition to the best-fitting model of the spectral energy distribution (SED), they indicated that the radius of the cavity was 47 AU and that hot gas resided within 0.8 AU. \citet{andrews11} re-calibrated the data from \citet{brown08,brown09} and combined it with SMA compact (COM) data. They treated the surface density by the similarity solution with an exponential tail in their analysis and derived the cavity radius to be 68 AU. \citet{isella13} also studied the asymmetric component in the LkH$\alpha$ 330 disk with SMA archival data and CARMA data at $\lambda = 1.3$ mm. As results of the visibility fitting with parametric gaussian disk models, they found that the asymmetric component traces a narrow circular arc, which extends in the azimuthal direction by about 90$\arcdeg$ and accounts for about 1/3 of the total disk flux at $\lambda = 1.3$ mm. They discuss possible mechanisms of the arc formation including planets, Rossby waves instabilities (RWI), baroclinic instabilities, disk warping, and disk shadowing. Observations of molecular tracers might help to further confirm the origin of the cavity in the disk. In this paper, the details of the $H$-band linear polarimetric observations and 0.87 mm dust continuum observations of LkH$\alpha$ 330 are provided in section \ref{obs}, and the obtained images which show a spiral arm and asymmetric structures in the LkH$\alpha$ 330 disk are presented in section \ref{result}. In section \ref{discussion}, micron and millimeter-sized dust distributions and the origin of the spiral arm as well as the asymmetric structure are discussed. Finally, we examine the possibility of dust trapping by estimating the dust opacity index using our observations and the data from the literature. Section \ref{summary} summarizes our findings.

\label{discussion} \subsection{Dust distribution and implications of a massive planet} \label{dist} Figure \ref{fig:fig3}a shows the radial surface brightness profiles along the directions of P. A. = 94$\arcdeg$ and 235$\arcdeg$ crossing the scattered emission peaks and P. A. = 258$\arcdeg$ crossing the continuum peak in the radial direction as denoted by red lines in Figure \ref{fig:fig1}c. Both of the scattering emission peaks are found at $r \sim 0\farcs25$ ($\sim$ 62.5 AU) from the central star indicated by dashed circle in Figure \ref{fig:fig1}c. Note that the polarized emission peaks at P. A. = 94$\arcdeg$ and 235$\arcdeg$ are almost located along the direction of semi-major axis (the position angle of LkH$\alpha$ 330 is 80$\arcdeg$ \citep{andrews11} measured from the north in counter-clockwise). While surface brightness asymmetries along the semi-minor axis can be strongly affected by anisotropic scattering, asymmetries along the semi-major axis are unaffected by the scattering efficiency and thus reflect intrinsic variations in the disk structure. Thus, the disk inclination of 35$\arcdeg$ \citep{isella13} does not affect the results very much. Figure \ref{fig:fig3}b represents the azimuthal surface brightness profile at $r = 0\farcs25$ and shows a non-axisymmetric torus in the disk. The surface brightnesses in the west and east sides are approximately 3 and 6 times higher than the base levels found in the north and south regions. As shown in Figure \ref{fig:fig1}c, the two continuum peaks at the east and west sides of the disk match with the peaks seen in the $H$-band PI image, indicating that small-sized dust (0.1 - 1 $\mu$m) has a similar distribution to the millimeter sized dust when phase function including polarization property is the same everywhere. The radial profiles of scattering and continuum emissions also support that they similarly distribute in the disk because their own peaks are seen at the same radial distance of 0$\farcs25$ as shown in Figure \ref{fig:fig3}a. It is important to know whether the disk is optically thick or not at $\lambda = 0.87$ mm continuum when they are compared with other emissions at different wavelength. It can be roughly checked by comparing the brightness temperature ($T_{\rm b}$) with the radiative equilibrium temperature. As a result, $T_{\rm b}$ at the continuum peak was 9.6 K, which is much lower than the radiative equilibrium temperature of 98 K at $r = 0\farcs25$, indicating that the continuum emission is optically thin. Our results suggest that dust grains between (sub)micron and millimeter in size accumulate at roughly the same radius, but not necessarily from an annular structure of equal thickness. Both the $H$-band PI image and the 0.87 mm continuum image shown in Figure \ref{fig:fig1}c indicate that there is a dust concentration in the disk. Recent theoretical studies have proposed several mechanisms to trap dust particles: snow lines, dead zones generated by magneto-rotational instability \citep[MRI;][]{balbus91,dzyurkevich10}, anticyclonic vortices that generate high pressure region triggered by RWI \citep{meheut12}, and pressure bumps generated by massive planets \citep{{paardekooper04},zhu12,pinilla12a,pinilla12b,pinilla15}. Among such possibilities, LkH$\alpha$ 330 implies the possible presence of a massive planet because at least spiral structure is associated in the disk. Spiral arms have been discovered in other transitional disks, such as SAO 206462, MWC 758, and HD 100453 and theoretical simulations reproduce the observed spiral features by embedded planet(s), indicating that spiral arms can be induced by planet-disk interaction \citep{muto12,garufi13,grady13,benisty15,dong15,dong16,wagner15}. A pair of arm-like features in the east and west sides of the LkH$\alpha$ 330 disk, extending from a continuum peak toward the southeast and northwest, may signify the presence of a massive planet probably located in the south east or north west direction as an extension of the spiral arms and which can generate a pressure bump that locally accumulates a large amount of dust. \citet{dong15} produced the $H$-band polarized intensity images by the newly developed Athena $+ +$ MHD code (J. M. Stone et al. 2016, in preparation) showed spiral structure with $m = 2$ mode (see Figure 2 in \citet{dong15}) by a single massive planet at outer disk region. They also predicts that the brightness of the observed arm is expected to be enhanced by 100 to 300 \% from the emission free background, which is consistent with our result of 380 \% enhancement. \citet{isella13} showed that the interaction between the disk and planet is a possible mechanism by utilizing hydrodynamic simulation code FARGO \citep{masset00} as a means to investigate the origin of the asymmetric disk structure. They explored disk structure during 100 to 2000 disk rotations with viscosity of 0.002 and 0.02 and the number of planets with different mass between 3 and 10 $M_J$, and could reproduce the observed azimuthal asymmetries at $\lambda$ = 1.3 mm emission. \citet{birnstiel13} also demonstrated the formation of azimuthal asymmetries in the surface density by disk-planet interactions during their orbital rotation. Therefore, large grains are probably trapped or piled up by a pressure bump generated by a massive planet that depletes gas surface density significantly while interacting with the gas disk \citep{pinilla12a, pinilla12b}. Small grains, on the other hand, trace the gas features because they are mostly coupled to the gas. They further study that such asymmetry possibly leads to azimuthal variations in dust opacity at the millimeter wavelength because large dust grains are trapped at pressure maxima due to the local gas density enhancement. They also argue that the variation of dust temperature in the azimuthal direction might be responsible for generating azimuthal asymmetry. The incident angle of the stellar radiation to the disk is not uniform if LkH$\alpha$ 330 disk is warped such as HD 142527 disk \citep{casassus15a,marino15}, the disk is differently flared in the azimuthal direction, or the shadowing variation behind the asymmetric outer gap wall in the inner disk, causing azimuthal variation in the dust temperature. Therefore, the $H$-band emission depends on internal disk structure, while sub-millimeter is not. All of the possibilities mentioned above can explain the origin of the anti-correlation observed between $H$-band and 0.87 mm continuum emissions. The key questions whether or not a massive planet can form the spirals in scattered light and the sub-millimeter asymmetry will arise. \citet{bae16} simultaneously reproduced the observed spiral arms in NIR and asymmetry in sub-millimeter continuum by a single planet-mass companion at the outer disk region in the case of SAO 206462 transitional disk, and thus their works support in explaining the LkHa330 structure similar to SAO 206462. If this is the case, a massive planet orbiting at the outer disk could be a more plausible explanation for current architecture seen in the LkHa330. They also predicted that brightest peak and the spiral arms will show observable displacement over the next few years. The follow-up observations in NIR and sub-millimeter continuum are demanded for revealing the origin of the structure. Details of the arm structures will be the subject of later studies in this series. Here, we simply stress that LkH$\alpha$ 330 has a spiral arm structure, a supportive evidence for the presence of a massive planet. In the next section, we discuss the possibility of grain growth due to a planet-induced pressure bump. \subsection{Implications of grain growth due to pressure bump} \label{growth} Grain growth can be found observationally from the slope of the SED, expressed as $F_\nu \propto \nu^\alpha$ where $\alpha = \beta + 2$, at mm/sub-mm wavelengths. We thus use this law to estimate the dust opacity in the LkH$\alpha$ 330 disk. $\beta$ is related to composition, shape, and size distribution of the dust that accounts for thermal emission at mm/sub-mm wavelengths. Generally, $\beta$ lies between 0 and 1 for disks around low-mass T Tauri stars and almost every disk falls in the range of $\beta < 2$ \citep{beckwith91}. Note that the $\beta$ of the interstellar medium (ISM) is about 1.7 \citep{li01}. It is widely accepted that the explanation for $\beta_{\rm disk} < \beta_{\rm ISM}$ is primarily a grain growth effect in the disk, and other contributing factors such as the geometry of dust particles or an optically thick disk are less effective and don't alter $\beta$ to the same extent \citep{draine06}. \citet{draine06} also shows that the dust grows from $\mu$m to mm sized grains when $\beta$ decreases below unity. Here we estimate $\beta$ with our 0.87 mm continuum observation and a previous result from a 1.3 mm continuum image taken by CARMA \citep{isella13}. The integrated intensity in the $\lambda = 0.87$ mm and $1.3$ mm continuum observations in west side of the disk is 101 $\pm$ 14.1 and 35 $\pm$ 1 mJy, respectively, but 87 $\pm$ 12.7 and 18 $\pm$ 1 mJy in the east side of the disk, respectively. As a result, the estimated $\beta$ becomes 0.7 $^{+0.5}_{-0.4}$ in the west side and 2.0 $\pm$ 0.5 in the east side, respectively. $\beta$ in the west side is below unity and indicates that grain growth probably takes place as suggested from the work of \citet{draine06}. It is theoretically inferred that the dust grains grow to around 500 $\mu$m or even larger, because such sizes are required for $\beta$ to be less than unity for a power-law distribution of grain sizes ($a$), $n(a) \propto a^{-p}$, where n(a) represents the abundance of grains with a particular size and $p$ is a power index of the distribution \citep{dalessio01,natta04a,natta04b}. $\beta$ in the east side, on the other hand, shows a typical number of $\sim$ 2, indicating that grain growth is not so active than the west side. Note that \citet{dalessio01} also show that $\beta$ will become approximately 1.7 when the maximum size of dust grains is less than 30 $\mu$m. A pressure bump is considered as one possible explanation to promote grain growth since it locally traps dust material in the azimuthal direction and as a result accelerates dust particles growth \citep{barge95,birnstiel13}. If this is the case, the larger grains subsequently stay near the pressure bump while the smaller grains get uniformly distributed in the entire disk as seen in Oph IRS 48 \citep{marel13,marel15}. We note that grains are likely growing everywhere in the disk, but transport is concentrating large grains at the pressure bump, resulting in asymmetric structures. LkH$\alpha$ 330 has at least one spiral arm feature and an asymmetric structure about east and west side of the disk. Although there are several explanations for these complex structures, a large unseen planet naturally explains simultaneously both of the structures. Observations at longer wavelengths available in the future with ALMA or VLA will be greatly helpful for further investigations of grain growth and the connection between large sized grains and rocky cores of planets.

1607.04708

1607

1607.07420_arXiv.txt

High-resolution X-ray spectroscopy with \hitomi\ was expected to resolve the origin of the faint unidentified $E\approx 3.5$ keV emission line reported in several low-resolution studies of various massive systems, such as galaxies and clusters, including the Perseus cluster. We have analyzed the \hitomi\ first-light observation of the Perseus cluster. The emission line expected for Perseus based on the \xmmnewton\ signal from the large cluster sample under the dark matter decay scenario is too faint to be detectable in the \hitomi\ data. However, the previously reported 3.5 keV flux from Perseus was anomalously high compared to the sample-based prediction. We find no unidentified line at the reported high flux level. Taking into account the \xmm\ measurement uncertainties for this region, the inconsistency with \hitomi\ is at a 99\% significance for a broad dark-matter line and at 99.7\% for a narrow line from the gas. We do not find anomalously high fluxes of the nearby faint K line or the Ar satellite line that were proposed as explanations for the earlier 3.5 keV detections. We do find a hint of a broad excess near the energies of high-$n$ transitions of S {\sc xvi} ($E\simeq 3.44$ keV rest-frame) --- a possible signature of charge exchange in the molecular nebula and another proposed explanation for the unidentified line. While its energy is consistent with \xmm\ pn detections, it is unlikely to explain the MOS signal. A confirmation of this interesting feature has to wait for a more sensitive observation with a future calorimeter experiment.

\label{sec:intro} The nature of dark matter (DM) is one of the fundamental unsolved problems in physics and astronomy. Direct particle searches in laboratories as well as searches for electromagnetic signal from celestial objects have been conducted with no unambiguous detection so far. X-ray observations of DM concentrations, such as galaxies and clusters, provide a probe for a particular DM candidate, a sterile neutrino, which is predicted to decay and emit an X-ray line (Dodelson \& Widrow 1994; Abazajian et al.\ 2001). Early searches that provided upper limits on line flux (and thus the particle decay rate) as a function of line energy (which gives the particle mass) are reviewed, e.g., in Abazajian et al.\ (2012) and Boyarsky et al.\ (2012). A possible detection was reported by Bulbul et al.\ (2014, hereafter B14), who found an unidentified line at $E\approx 3.55$~keV in the stacked spectrum of a large sample of galaxy clusters using \xmmnewton\ EPIC MOS and pn. Within their sample was the Perseus cluster (its central region), whose signal was particularly strong. B14 also reported a detection from Perseus with \chandra\ at the same energy. Boyarsky et al.\ (2014) reported an \xmm\ detection in the outer region of Perseus. Urban et al.\ (2015) and Franse et al.\ (2016, hereafter F16) detected the line in several regions of Perseus with \suzaku; however, Tamura et al.\ (2015) did not detect it in the same \suzaku\ data. The 3.5 keV line was also reported from other objects, such as the Galactic Center (Boyarsky et al.\ 2015) and M31 (Boyarsky et al.\ 2014). Other sensitive searches did not detect a significant line signal (e.g., from the Milky Way halo, Sekiya et al.\ 2016; Draco dwarf, Ruchayskiy et al.\ 2016; stacked \suzaku\ clusters, Bulbul et al.\ 2016). Some of the nondetections were inconsistent with other detections under the decaying DM hypothesis (in which the line flux must be proportional to the projected DM mass), most significantly, in a sample of galaxies (Anderson et al.\ 2015). We also note here that the signal from Perseus reported by \xmm, \chandra\ and \suzaku\ was higher than expected given the signal from the rest of the cluster sample (B14). Astrophysical explanations of the reported line, in addition to those considered by B14, have also been proposed; a critical review can be found in F16. An extensive review of the recent observations is given, e.g., by Iakubovskyi (2015). As recognized in all previous studies, the above line detections were near the capability for CCD detectors --- for a $\sim$100 eV resolution, the line reported from clusters with a $\sim$1 eV equivalent width (EW) is a 1\% bump above the continuum, easily affected by errors in modeling the nearby atomic lines and in instrument calibration. A confirmation with a much better spectral resolution was considered essential. \hitomi, launched in February 2016 and lost in March (Takahashi et al.\ 2014, 2016) after having returned a groundbreaking spectrum of the Perseus cluster (\hitomi\ Collaboration 2016, hereafter H16), offered us such a possibility. We present results from this dataset below. We use $h=0.7$, $\Omega_\mathrm{m} = 0.3$ and $\Omega_\mathrm{\Lambda} = 0.7$ cosmology. The cluster heliocentric redshift (average for member galaxies) is 0.0179 (Strubble \& Rood 1999) and the redshift in the CMB frame is 0.01737, which gives $d_L=75.4$ Mpc and a scale of 21.2 kpc per 1\am. We use the 68\% ($1\sigma$) confidence level for errors unless stated otherwise.

\label{sec:disc} Our analysis of the \hitomi\ spectrum of the Perseus cluster core reveals no unidentified emission line around the energy reported by B14. It is inconsistent with the presence of a line at the flux reported by B14 using \xmm\ MOS (as rederived for the approximate SXS FOV). Taking into account the uncertainties of the \xmm\ MOS measurement in this region, which itself is only $3\sigma$ significant, the inconsistency with \hitomi\ for a broad line (that would be emitted by DM) is at the 99\% confidence level, and 99.7\% for a narrow line from the ICM. The broad line exclusion level is 97\% if we force the SXS line flux to be zero, assuming in effect that the mild ``dip'' in the residuals (\S\ref{sec:broad}) is not statistical as we concluded, but some instrumental artifact present only in Perseus and not in other SXS data. We note here that F16, using \suzaku\ data for a similar Perseus region, reported a line flux and its uncertainty similar to that from the \xmm\ SXS-FOV measurement, but given the lack of consensus between different \suzaku\ analyses (cf.\ Tamura et al.), we leave a comparison with \suzaku\ for a later work. We can exclude one of the 3.5 keV line astrophysical explanations proposed by B14 --- namely, anomalously bright K {\sc xviii} He$\alpha$ or Ar {\sc xvii} He$\beta$ satellite lines. These lines are not significantly detected in the SXS spectrum, their fluxes are consistent with expectations and below the MOS 3.5 keV flux. If we consider a slightly wider energy range (Fig.\ \ref{fig:limit}), there is a hint of a broad excess emission feature of the right amplitude (though at very low statistical significance) at $E\approx 3.44$ keV rest-frame, where charge exchange on S {\sc xvi} has been predicted (Gu et al.\ 2015). However, the energy of this feature is $2.6\sigma$ (100 eV) away from the best-fit energy for the MOS SXS-FOV detection, and even more inconsistent with the MOS stacked-cluster sample, though it is consistent with the pn detections (B14). If confirmed with better statistics, it is an interesting feature in itself. Given \hitomi's much greater spectral resolution, it is likely that the inconsistency with \xmm\ that we reported here is attributable to a systematic error in the \xmm\ result. Possible causes will be examined in a future work, using the new accurate knowledge of the fluxes of all the nearby atomic lines from \hitomi, as well as \suzaku\ and \chandra\ spectra and models. One possible reason, mentioned among the Caveats in B14, is that with a CCD resolution, a spurious $\sim$1\% dip in the effective area curve is all that is needed to produce a false line-like residual of the observed amplitude (see Fig.\ 7 in B14). This is an obvious problem for detections in a single object or in local objects, even when different instruments with similar low-resolution detectors are used. Such systematic effects can be minimized by stacking objects at different redshifts. In the cluster sample of B14, the 3.5 keV rest-frame energy spans a 1.2 keV interval of detector energies, which should smear out any such instrument features. Thus this systematic error will be much smaller in the stacked-sample signal. As noted in B14 and subsequent works, the reported line in Perseus, and especially in its core, was much brighter than expected from the signal in the larger cluster sample, scaled by mass under the decaying-DM hypothesis. Assuming that the high Perseus line flux is an artifact but the stacked-sample signal is real, we can evaluate the corresponding expected flux from the SXS FOV. To estimate the projected dark matter mass within this region, we use a total mass profile from Simionescu et al.\ (2011) and one from the Vikhlinin et al.\ (2006) $M-T$\/ scaling relation (the former was used in Urban et al.\ and the latter in B14), correcting them for the 14\% baryon fraction. The projected DM mass within the SXS FOV is $(6-8)\times 10^{12}$\msun. For the sterile neutrino decay rate derived in B14 for the full cluster sample ($\Gamma\approx 2\times 10^{-28}$ s$^{-1}$), we expect a 3.5 keV line with $f=(2.4-3.1)\times 10^{-7}$ \photcms\ (see, e.g., B14 for the equations), 30 times below the flux we ruled out above. This flux, shown by blue dashed line in Fig.\ \ref{fig:limit}, is below the statistical noise in the current observation. The right vertical axis in Fig.\ \ref{fig:limit} shows the sterile neutrino decay rate that corresponds to the line flux on the left axis, using the median projected DM mass estimate. The \hitomi\ $3\sigma$ upper limits on $\Gamma$ are, unfortunately, much higher than many earlier constraints (see, e.g., B14). This is because of the high X-ray brightness of the ICM in the Perseus core, the short exposure (combined with the GV attenuation), and the small SXS FOV. Our results from this relatively short observation illustrate the dramatic improvement in sensitivity for narrow features from that of CCD detectors. However, as expected, the improvement for a putative cluster DM line, which would have a width of 30--35 eV (FWHM), is less significant. The short \hitomi\ observation excluded the anomalously bright signal reported from the Perseus core. However, to test the much weaker stacked-sample detection (provided it withstands the reevaluation of the systematic uncertainties after the \hitomi\ result) will require the photon statistics comparable to that of the CCD stacking studies, or looking at objects where the line is easier to detect. Among clusters, such objects would be non-cool-core systems, in which the line EW should be an order of magnitude higher for the same line flux because of the lower ICM background. A DM line would be narrower, giving a calorimeter a greater leverage, in systems with low velocity dispersion, such as dwarf spheroidals and the Milky Way. Of course, distinguishing a DM line from an astrophysical one would require resolving the line, which only a calorimeter can do.

1607.07420

1607

1607.05563_arXiv.txt

{This paper describes an automatic isophotal fitting procedure that succeeds, without the support of any visual inspection of neither the images nor the ellipticity/position-angle radial profiles, at extracting a fairly pure sample of barred late-type galaxies (LTGs) among thousands of optical images from the Sloan Digital Sky Survey (SDSS). The procedure relies on the methods described in Consolandi et al. (2016) to robustly extract the photometrical properties of a large sample of local SDSS galaxies and is tailored to extract bars on the basis of their well-known peculiarities in their P.A. and ellipticity profiles. It has been run on a sample of 5853 galaxies in the Coma and Local supercluster. The procedure extracted for each galaxy a color, an ellipticity and a position angle radial profile of the ellipses fitted to the isophotes. Examining automatically the profiles of 922 face-on late-type galaxies (B/A$>0.7$) the procedure found that $\sim 36 \%$ are barred. The local bar fraction strongly increases with stellar mass. The sample of barred galaxies is used to construct a set of template radial color profiles in order to test the impact of the barred galaxy population on the average color profiles shown by Consolandi et al. (2016) and to test the bar-quenching scenario proposed in Gavazzi et al. (2015). The radial color profile of barred galaxy shows that bars are on average redder than their surrounding disk producing an outside-in gradient toward red in correspondence of their corotation radius. The distribution of the extension of the deprojected length of the bar suggests that bars have strong impacts on the gradients of averaged color profiles. The dependence of the profiles on the mass is consistent with the bar-quenching scenario, i.e. more massive barred galaxies have redder colors (hence older stellar population and suppressed star formation) inside their corotation radius with respect to their lower mass counterparts. }

Hydrodynamical simulations have made clear that bars have a major impact on the secular evolution of galaxies \citep{atha02,sell14}. It is well known that barred potentials exert non axisymmetric forces onto the gaseous component of the galaxy: the gas within the corotational radius is rapidly funneled to the center of the galaxy \citep[within the Inner Lindblad Resonance, see][]{KK04,K13,sell14, fanali15} while the gas outside is confined to the outer disk \citep{sanders76,shlo89,atha92,beren98, regan04,kim12,cole14}. While it is not clear whether this phenomenon can trigger AGN activity or not \citep{emsell15}, it is out of question that the high density reached by the gas dragged in the center of the galaxy triggers a burst of star formation that rapidly depletes it, turning into gas-poor the region inside the corotational radius \citep{krum05,krum09,daddi10,genzel10}. From an observational point of view, starting from the seminal work of E. Hubble who dedicated half of its world-wide famous tuning fork to barred disk galaxies \citep{hub36}, bars have increasingly captured the interest for the understanding of galaxy secular evolution. As a matter of fact, throughout the years observations have enlighten the physical effects of bars on galaxies \citep{saka99,KK04,jo05,sh05,K13} and, above all, the extremely high frequency of barred galaxies among spirals. Indeed, among local bright disk galaxies, $\sim 60\%$ \citep{kna99,esk00,delmestre07,marinova07} are barred if observed in the near-infrared and about $\sim 40\%$ in the optical bands \citep{esk00,marinova07}, hinting at bars as fundamental drivers of the evolution of late type galaxies (LTGs). If and how the bar fraction evolves across the cosmic time is still under debate \citep{jo04,sh08} as well as the exact determination of the dependence of the bar frequency on stellar mass (especially at the faint end of the mass function) and on galaxy environment \citep{thomp81,marinova12,skibba12,lans14,alonso14}. For example, \citet{masters12}, \cite{skibba12}, \citet{abreu12}, \citet{pg15} all consistently report a bar fraction that increases with increasing mass. Nevertheless, \citet{bara08} recover a strong-bar fraction that increases with decreasing mass while \citet{nairbfrac} find a strong-bar fraction that decreases from $\sim 10^{9}$M$_\odot$ to $\sim 10^{10}$M$_\odot$ and increases again from $\sim 10^{10}$M$_\odot$ to $\sim 10^{11}$M$_\odot$. The work by \citet{cheung13} and \citet{pg15} underline the crucial importance of determining the exact dependence of the fraction of strongly barred galaxies on total stellar mass. Namely, an increasing bar fraction with increasing mass would explain the central quenching of the star formation (SF) in high mass main sequence galaxies that bends the local star formation rate (SFR) vs stellar mass relation at $\sim 10^{9.5}$M$_\odot$. Moreover, \citet{sanja10} report that below $\sim 10^{9.5}$M$_\odot$ disks are systematically thicker, making it difficult to develop bars. \begin{figure} \begin{centering} \includegraphics[scale=0.65]{Fig01.eps} \caption{Wedge diagram of galaxies belonging to the sample studied in this work: the Coma supercluster (c$z>4000$kms$^{-1}$) and the Local supercluster (c$z<3000$kms$^{-1}$). Blue dots represent late-type galaxies wile red dots stand for early-type galaxies. } \label{wedge} \end{centering} \end{figure} \begin{figure} \begin{centering} \includegraphics[scale=0.4]{Fig02.eps} \caption{Distribution of the $i$-band magnitudes of the whole sample and separately for the Local (green line) and the Coma (blue line) supercluster sample.} \label{magdistr} \end{centering} \end{figure} Determining robustly the real bar fraction below $10^9$M$_\odot$ demands statistics as well as sensitivity and resolution. High redshift determinations lack of both aspects and the best environment to determine the optical bar fraction in such a wide range of mass is therefore the local Universe, taking advantage of the publicly available data of the Sloan Digital Sky Survey \citep[SDSS,][]{sdss}. The SDSS fully covered the area of two nearby structures such as the Virgo and the Coma superclusters that give the opportunity to study thousands of galaxies down to a limiting mass as low as $10^7$M$_\odot$ with a physical resolution of $\sim 600$ pc at the distance of Coma. However, the task is non-trivial, as stellar bars among the population of dwarf late type galaxies are often low surface brightness features that are easily confused with or hidden by poorly resolved patches of SF. Purely visual inspection drives the classification of barred galaxies of most of the catalogs of galaxies such as the cases of the RC3 \citep{rc3}, VCC \citep{vcc} and \citet{naircat} classifications. Nevertheless, because of the subjective nature of visual inspection classification, it is habit to average the votes of as many as possible classifiers. A clear example is given by the recent works by \citet{masters11}, \citet{masters12}, \citet{skibba12} and \citet{mel14} which are based on the classification of the Galaxy Zoo project \citep[{\it www.galaxyzoo.org}, see][]{zoo1,masters11}. This classification is made on the basis of thousands of votes assigned by citizens that are asked to visually inspect SDSS galaxies and to answer questions about the morphology of the object.\\ On the other hand, a different approach is adopted for example by \citet{woz95}, \citet{jo04}, \citet{bara08}, \citet{bars09} and \citet{abreu12} who support the visual inspection by performing isophotal fitting analysis and looking at some precise features in the ellipticity and position angle (P.A.) radial profiles. In this paper is tested a different method to produce an objective classification that does not rely on visual inspection of neither the images nor the ellipticity/P.A. radial profiles of galaxies and extract a fairly pure sample of barred LTGs. An automatic procedure based on the sample and on the methods explained in \citet{C16} for the isophotal fitting (explained in section \ref{tilt}) ran on $\sim6000$ SDSS galaxies in the Local and Coma superclusters. Throughout section \ref{vote} is presented the automatic bar-extraction criteria based on the ellipticity and P.A. profiles extracted. Limitations, pureness and completeness of the automatic classification are discussed in section \ref{check} by comparing the final selection to other classifications found in the literature. In section \ref{bfrac} the resulting bar fraction vs mass relation is compared to other published results. Finally, the results are discussed in section \ref{sunto} and the bar-quenching scenario described by \citet{pg15} is tested by creating a set of templates of color profiles of barred galaxies in different bins of mass. These will be also compared to the template color profiles produced in \citet{C16} (from now on C16).

\label{sunto} Holding the selection onto late-type galaxies and excluding S0s, I constructed a template $(g-i)$ color profile in three different bins of mass: i) below $M<10^{9.75}$M$_{\odot}$ ii) $M>10^{9.75}$M$_{\odot}$ and $M<10^{10.25}$M$_{\odot}$ iii) $M>10^{10.25}$M$_{\odot}$. The bins were selected to guarantee that more than 50 galaxies contribute to each template profile. Color profiles are good tracers of the specific SFR and works such as \citet{mac04,mcd11}, C16 have highlighted the good correspondence between average age of the stellar population as a function of radius and color radial profiles. Moreover, using the technique of template profiles, C16 have shown that massive spiral galaxies develop a red and dead component, the importance of which increases with mass. Using the same technique on the tilted color profiles of the subsample of barred galaxies of C16 extracted in this work, we test if the central red and dead component is consistent with the presence of a bar-like structure in the center of galaxies. The template profiles for barred galaxies are displayed in Fig.\ref{tmpl} with radius normalized to the radius of the selected ellipticity peak (R$_{ellpeak}$), a good proxy for the bar length and corotation radius \citep{lauri10}, with a radial step of 0.1 $\rm R/R_{ellpeak}$. \begin{figure*} \begin{centering} \includegraphics[scale=0.45]{Fig07a.eps}\includegraphics[scale=0.45]{Fig07b.eps} \caption{(left) Distribution of the ratio between the deprojected semi major axis of the bar and the semi major axis of the galaxy. The dashed line indicates the average value of the distribution equal to 0.38. (right) The distribution of the ellipticity of the bars extracted in this work.} \label{barextent} \end{centering} \end{figure*} From low to high mass, the template profiles evolve significantly. In the lowest mass bin the color profile is blue over all radii with possibly only a mild gradient toward red inside the bar radius. Things change clearly in the intermediate mass template profile: the outer disk (R/Rpeak > 1) is blue and in the outermost regions overlaps with the lowest mass bin profile, while inside the corotation radius (R/Rpeak <1) the color profile is red as an elliptical of the same mass. The bar has already reached the red sequence and, on the contrary, the disk is still on the blue cloud. In the highest mass bin the red component has once again the typical color of the red sequence in the same range of mass. The disk is still bluer in the outer region but displays an average color of the disk that is redder than the respective average blue cloud values of the same mass. Moreover in Fig.\ref{mkn} we show the template radial color profiles of barred galaxies for two samples up and below the threshold mass indicated by \citet{pg15} to be the mass above which bars have the region inside the corotation radius quenched and therefore red. % In order to quantify the average extension of the region that under the bar influence we correct the radius of the ellipticity peak for projection effect using the measured P.A with respect to its galaxy that is considered to have the P.A. of the last fitted isophote. The distribution of the ratio of the deprojected bar semi major axis and the galaxy semi major axis is plotted in Fig \ref{barextent}. This distribution peaks at 0.3 consistently with others results such as the one published in \citet{marinova07} and \citet{bara08}. The bars that we extracted are primarily strong bars \citep[e$> 0.4$,][]{lauri10} and weak bars represent only the $\sim 9\%$ of the bars extracted which is again consistent with the proportion observed in \citet{marinova07}. Looking at Figure \ref{tmpl} we can deduce that bars are on average redder structures if compared to their associated disks. Fig \ref{bfr} reveals that especially at high mass bars are extremely common and will likely have a big impact on the average photometric properties of the galaxy population. Therefore bars will be strong contributors of the trends in color shown in the template color profiles of C16 especially for high mass objects. Moreover a further clue comes from the distribution of the ratio between a$_{bar}$ and a$_{gal}$ (Fig. \ref{barextent}) indicating that the average optical extension of the bar is $\sim 0.3$ a$_{gal}$ which is consistent with the extension of the intermediate/internal zone identified in the average color radial profiles of C16 that is on average redder than the outer disk zone. A possible explanation for such a correspondence between the presence of the bar and the color of the galaxy, is the one proposed by \citet{pg15} who finds that the sSFR of main sequence local galaxies have a downturn at high mass indicating that massive disks have suppressed sSFR with respect of their lower mass counterparts. These authors suggest that the torque exerted onto the gaseous component by the bar, funnels the gas inside the corotation radius to the very center of the galaxy where it's rapidly consumed by a burst of star formation. The region within the bar extent is therefore gas depleted and grows redder with time and this phenomenon can occur earlier in more massive disks which are dynamically colder. On the contrary, the gas outside the corotation radius is hold in place and keeps feeding the star formation maintaining the disk blue. Nevertheless the bar fraction versus mass relation, along with the well known color-mass relation \citep{C16}, implies that there is an higher fraction of bars among more massive galaxies with redder total colors, although we stress that these are still star forming spiral galaxies. This is still consistent with previous works such as the one of \citet{masters11}, \citet{alonso13}, \citet{alonso14} who consistently find an increasing barfraction in redder galaxies. I found a difference between the bar fraction evaluated among all LTGs (ty>1) and the one that embraces lenticulars. As a matter of fact when lenticulars are taken into account the bar fraction decreases at all masses. At this point, a note of caution is required: the morphology selection relies only on visual morphological classification \citep{vcc} which cannot disentangle the small population \citep[$\sim13\%$ of ETGs,][]{a3d,emsell11} of slow rotators (pure ellipticals) from the much wider population of fast rotators (disks that should have been taken into account when calculating the bar fractions). Therefore this estimate of the bar fraction of the joint population of LTGs and S0/dS0s could be biased by the morphological classification. Nevertheless, the proportion that stands between fast and slow rotators \citep[respectively $\sim87\%$and $\sim13\%$ of ETGs,][]{emsell11} implies that the bar fraction could be even lower, as many fast rotators have been likely classified as ellipticals and therefore excluded from the bar fraction determination. The lower bar fraction can possibly arise from two different scenario: i) S0s are older systems with respect of other disk galaxies and have already undergone buckling instability that weakened the bar; ii) given the well known density-morphology relation \citep{dress80}, S0s populate dense environments which prevents from growing bars because of tidal interactions or fast encounters. Nevertheless it is worth stressing that other results suggest that intermediate/high density environments such as groups can indeed enhance the possibility of growing a galactic bar \citep{skibba12}. As a final note, I would like to highlight that in the highest mass bin the bar fraction is lower. This feature has a low statistical significance but it can be seen also in other works that show the bar fraction as function of mass such as \citet{nairbfrac}, \citet{abreu12}. Although its low significance, such a decrease is possibly consistent with other two possible scenarios: i) more massive disks develops a bar at earlier times with respect to their lower mass counterparts and therefore undergo buckling instability earlier and dismantle the bar at earlier epochs. ii) More massive disks have a different merger history with respect of the low mass population and this may induce a different bar fraction. In the future it would be possible to investigate these hypothesis with the advent of new cosmological simulations at sufficient resolution. Summarizing, in this paper I have developed an IDL-based bar finder that performs isophotal fitting on SDSS images and on the basis of the extracted radial ellipticity an P.A. profiles and recognizes barred galaxies avoiding visual inspection of neither the images nor the profiles. This procedure makes use of the tasks described in C16 and has been tested over the same sample in order to evaluate the bar fraction in the Local and Coma supercluster and quantify th influence of barred galaxies on the average properties of color profiles of LTGs shown in C16. i) The procedure have extracted a fairly pure sample of barred galaxies among face-on LTGs and led to the calculation of a bar fraction of $\sim36\%$ consistent with other literature results \citep{jo04,marinova07,nairbfrac}. ii) The bar fraction shows a strong mass dependency obtained also by previous works in the local volume and at higher redshifts \citep{marinova07,nairbfrac,skibba12,abreu12}. iii) The bars that we extracted typically occupy the central $\sim 30-40 \%$ of the host galaxy and is typically strong ($\sim90 \%$ of times), consistent with the proportions observed by \citet{marinova07}. iv) I constructed color average profiles of barred galaxies in different bins of mass and compared it to the template profiles of C16. ii) and iii) imply that bars likely have a strong impact on the average color profiles created by C16 who observed in the template profiles of LTGs the growth of a red and dead component in an intermediate zone inside 0.3 Petrosian radii whose importance increases with mass. From iv) I was able to assess that bars are redder structures with respect of their disks and can indeed reproduce the upturn toward red of the templates profiles of C16. Moreover this further links the presence of a bar to a decrease of the SFR in a disk galaxies as proposed by \citet{cheung13,pg15}.

1607.05563

1607

1607.05080_arXiv.txt

We report the results of our radio, optical and infra-red studies of a peculiar radio source 4C~35.06, an extended radio-loud AGN at the center of galaxy cluster Abell 407 ($z=0.047$). The central region of this cluster hosts a remarkably tight ensemble of nine galaxies, the spectra of which resemble those of passive red ellipticals, embedded within a diffuse stellar halo of $\sim$1~arcmin size. This system (named the `Zwicky's Nonet') provides unique and compelling evidence for a multiple-nucleus cD galaxy precursor. Multifrequency radio observations of 4C~35.06 with the Giant Meterwave Radio Telescope (GMRT) at 610, 235 and 150 MHz reveal a system of 400~kpc scale helically twisted and kinked radio jets and outer diffuse lobes. The outer extremities of jets contain extremely steep spectrum (spectral index -1.7 to -2.5) relic/fossil radio plasma with a spectral age of a few$\,\times (10^7 - 10^8)$ yr. Such ultra-steep spectrum relic radio lobes without definitive hot-spots are rare, and they provide an opportunity to understand the life-cycle of relativistic jets and physics of black hole mergers in dense environments. We interpret our observations of this radio source in the context of the growth of its central black hole, triggering of its AGN activity and jet precession, all possibly caused by galaxy mergers in this dense galactic system. A slow conical precession of the jet axis due to gravitational perturbation between interacting black holes is invoked to explain the unusual jet morphology.

Galaxy mergers, which take place more frequently in the dense environments of galaxy clusters play a pivotal role in the evolution of galaxies in the Universe across cosmic time. The merger dynamics is influenced by various factors such as size, mass, impact parameter, relative velocity, gas content and the relative inclination of the participating galaxies. Mergers have profound effects on the properties of galaxies on various physical scales. On galactic scales ($\sim10 -100$ kpc) mergers may result in ram pressure stripping of gas, the formation of long tidal tails and enhanced star formation. On smaller scales ($< 1$ pc), the growth of black holes (BHs) and possible triggering of AGN activity with occasional relativistic jet ejection may occur in mergers, which transport matter and energy from the galaxy interiors to the surroundings through AGN feedback processes \citep{MN12}. On the smallest scales ($<<$pc), the final inspiral stage of merging BHs results in powerful gravitational wave emission. However, the formation and merger rates of galactic BHs are still much uncertain. Finally, on very large scales ($\sim100 -1000$ kpc), the shocks and turbulence inducted into the intra-cluster medium (ICM) during mergers may inject large amounts of nonthermal energy and subsequent shock heating of the ICM to X-ray temperatures, with disruption of cooling cores \citep{kandu2006,b64,spaul_01}. Therefore cluster centers are fascinating laboratories to study galaxy formation and evolution. \begin{center} \begin{table*} \caption{Galaxy properties in central region of Abell 407 cluster.} \begin{tabular}{|c|c|c|c|c|c|c|c|c|c|c|} \hline {\hskip -4.50cm} Galaxy$^a$ & {\hskip -5.0cm}Coordinates (J2000)& Redshift$^b$ & \multicolumn {5}{c|}{SDSS Magnitudes} & \multicolumn{3}{|c}{SDSS Colours} \\\cline{4-8} \cline{9-11} & {\hskip -6.0cm}(from SDSS) & &$ u$ &$g$ & $r$ &$i$ &$z$ & $g-r$ & $r-i$ & $i-z$\\ \hline \hline {\hskip -4.50cm}G1 (B)& {\hskip -5.0cm}03h 01m 51.5s +35d 50m 30s & 0.0483 & 18.23 & 15.84 & 14.72 & 14.23 &13.84 &1.12 &0.49 &0.39\\ {\hskip -4.50cm}G2 (A)& {\hskip -5.0cm}03h 01m 51.2s +35d 50m 22s & 0.0454 & 21.09 &19.04 & 18.46 & 17.95 & 17.54 &0.58 &0.51 &0.41\\ {\hskip -4.50cm}G3 (D)& {\hskip -5.0cm}03h 01m 51.8s +35d 50m 20s & 0.0471 & 18.11 &16.00 & 15.00 & 14.44 &13.97 &1.00&0.56 &0.47 \\ {\hskip -4.50cm}G4 (C) & {\hskip -5.0cm}03h 01m 51.7s +35d 50m 12s & 0.0501 & 20.91 &18.81 & 17.76 & 17.37 &16.90 &1.05 &0.39 &0.47 \\ {\hskip -4.50cm}G5 (F) & {\hskip -5.0cm}03h 01m 51.5s +35d 50m 12s & 0.0478 & 20.02 &17.93 & 16.86 & 16.31 &15.85 &1.07 &0.55 &0.46 \\ {\hskip -4.50cm}G6 (E) & {\hskip -5.0cm}03h 01m 52.4s +35d 50m 29s & 0.0444 & 21.33 &19.93 & 19.11 & 18.70 &18.18 &0.82 &0.41 &0.52 \\ {\hskip -4.50cm}G7 (G) & {\hskip -5.0cm}03h 01m 53.2s +35d 50m 26s & 0.0471 & 18.37 &16.34 & 15.31 & 14.80 &14.39 &1.03 &0.51 &0.41 \\ {\hskip -4.50cm}G8 (H) & {\hskip -5.0cm}03h 01m 53.7s +35d 50m 28s & - & - &- & - & - &- &- & -&\\ {\hskip -4.50cm}G9 (I) & {\hskip -5.0cm}03h 01m 54.5s +35d 50m 18s & 0.0451 & 20.27 &18.68 & 17.59 & 17.23 &16.68 &1.09 &0.36 &0.55\\ \hline $^a$ Zwicky's original nomenclature is shown within brackets.\\ $^b$ From \cite{b26}. \end{tabular} \label{tab1} \end{table*} \end{center} Rich galaxy clusters are usually dominated by massive, luminous central elliptical galaxies known as the Brightest Cluster Galaxies (BCG). There are some interesting examples, where extremely massive ($> 10^{12} M_{\odot}$) and brighter ellipticals called the cD galaxies form in the densest regions near the spatial and kinematical center of their host clusters \citep{tov_01,b62}. The distinguishing property of cD galaxies is the presence of a diffuse, faint stellar halo that may extend up to 100s of kpc, well into the intracluster medium \citep{b62,tov_01}. These facts seem to suggest that the formation of cD galaxies is unique to the cluster environment and is linked closely to its dynamical history. It is still far from clear how cD galaxies originate in such a galactic environment, because very few clusters with cD galaxies in the process of formation have been identified. \begin{figure} \centering \includegraphics[width=3.5in,height=4in,keepaspectratio]{A407_3.jpg} \caption{Colour image ($5.6^{\prime} \times 4.7^{\prime}$) of the galaxy cluster Abell~407 taken from the Sloan Digital Sky Survey (SDSS). The central region is host to a striking group of nine close packed galaxy like condensations, embedded within a diffuse stellar halo of intra-cluster light (the `Zwicky's Nonet'). This system possibly represents an exceptionally rare site of a multiple-nucleus cD galaxy precursor assembling in a rich galaxy cluster environment. } \label{fig1} \end{figure} \begin{figure*} \centering \includegraphics[width=3.0in,height=3.5in,keepaspectratio]{JHK_composite_MOD.png}\quad \includegraphics[width=2.9in,height=3.5in,keepaspectratio]{3C3506_150MHz+RASS.png} \caption{Left panel: A pseudo-color near infra-red image ($2.1^{\prime} \times 2.1^{\prime}$) of the compact group `Zwicky's Nonet' obtained by combining the J, H and K-band images available in the 2MASS 6X deep survey. The nine central galaxies of the nonet are marked on the image while Table~\ref{tab1} and Table~\ref{tabk} show their positions, magnitudes, colours and estimated central black hole masses. Right Panel: GMRT 150 MHz radio image ($11^{\prime} \times 11^{\prime}$) of 4C 35.06 is shown superposed on the ROSAT 0.5 - 2.4 keV band smoothed X-ray data shown with contours.} \label{2mass_rass+150} \end{figure*} cD galaxies are almost always radio loud and often eject powerful radio jets from accreting supermassive black holes \citep{bagchi_94,b64}. Galaxy interactions and major mergers remove significant amounts of angular momentum by gravitational torque and drive a part of the constituent gas towards nuclear supermassive black holes (SMBHs), thereby triggering the activity of the central engine \citep{b72}. As a result, cD galaxies are more likely to be radio-loud above radio luminosity $\sim 10^{24.5}$ W Hz$^{-1}$ at 1.4 GHz compared with other galaxies in a cluster \citep{bagchi_94}. The multiple galactic nuclei observed in cDs and BCGs provide evidence supporting the merger scenario. In addition, the SMBHs of these galaxies are believed to grow by multiple galactic mergers \citep{volo_2003,kul_2012,b51}. Thus, gravitational perturbations in the accretion disc in the presence of closely spaced massive black hole pairs may result in the occurrence of distorted radio jets. The inversion symmetry found in `S' or `Z' -shaped radio sources is ascribed to the precession of a spinning black hole \citep{b51,b76} and the associated tilting/perturbation in the accretion disk \citep{b77,b78}. The precessing radio jets in this situation are likely to trace a characteristic helical pattern on the sky \citep{b96}. One prominent object showing this rare phenomenon is PKS 2149-158, a dual radio-loud elliptical galaxy pair in the center of cluster Abell 2382, forming a pair of twisted jet systems \citep{b1}. The well known `C'-shaped wide-angle tail (WAT) source 3C~75 is another striking example of twin AGN producing two pairs of jets showing oscillations and interactions \citep{b2}. In this article, we report radio, optical and near-IR studies of the extraordinary radio source 4C 35.06 (B2 0258+35B), located near the center of the cluster Abell 407, which, interestingly, also harbors a remarkably compact group of nine galaxies embedded inside a diffuse stellar halo of faint intra-cluster light. We make a detailed study of this remarkable system, which possibly provides a unique and compelling evidence for an ongoing formation of a giant cD galaxy at the cluster center and connects to the evolution of its central black hole by mergers. The radio source 4C 35.06 clearly shows a very peculiar twisted jet system on 100 kpc scale, which we investigate further with multi-frequency radio data. This paper is organised as follows; In Section 2 we discuss the optical, near-IR, radio and X-ray properties of the system. In Section 3 we discuss the observations and the data reduction procedure. In Section 4 the main results obtained in the present study are discussed in seven subsections. In the last concluding section, section 5, we summarise our findings. Throughout this article, a Hubble constant $ H_{0} = 73\,$ $\rm km s^{ -1}\, Mpc^{ -1},$ and cosmological parameters $\Omega_{M} = 0.27$ and $\Omega_{\Lambda} = 0.73$ were used. For redshift $z = 0.047$ it implies the linear scale of 0.885 kpc arcsec$^{ -1}$ and a luminosity distance of 200 Mpc \citep{b68}. We define the synchrotron emission spectral index $\alpha$ by $S(\nu) \propto \nu^{\alpha}$, where $S(\nu)$ is the flux density at frequency $\nu$. \section[]{`Zwicky's Nonet': Previous Optical, X-ray and Radio Observations} Abell~407 is a rich galaxy cluster of Bautz-Morgan class II, at a redshift of 0.047 \citep{b56}. The optical image of the central region of this cluster from Sloan Digital Sky Survey (SDSS) shows a complex ensemble of at least nine galaxy-like condensations $\sim 1 \arcmin$ across ($\sim50$ kpc), embedded in a low surface brightness, diffuse stellar halo, which is reminiscent of a giant cD galaxy (Figure~\ref{fig1}). Table~\ref{tab1} lists the SDSS optical magnitudes, redshifts and colours of these nine galaxies. The $g-r$ colour index of $\sim1$ indicates they are passive, early-type red galaxies. Historically, Fritz Zwicky first noticed this extraordinary galactic configuration (V Zw 311) in 1971 \citep{zwicky71}; which was later studied in more detail by \cite{b26} using the Palomar 200-inch and 60-inch telescopes. \cite{b26} described it as the ``most nightmarish known multiple-nucleus system" and concluded that this puzzling galactic system possibly represents an extremely rare and unique snapshot of a giant cD galaxy caught in its formative stages. There were no further investigations of this highly unusual galactic system to understand its nature. Here we propose to name this extraordinary galaxy group of nine galaxies as `Zwicky's Nonet', honoring Fritz Zwicky who first noticed this galaxy group. To our knowledge, this is the most compact and rich system of multiple galaxies known to date. Some other famous compact groups with multiple members are the `Stephen's Quintet' and `Seyfert's Sextet', and the lesser known `Zwicky's Triplet' (Arp 103). In the Uppsala General Catalogue, this multi-galactic system is confusingly listed as a single galaxy UGC~2489 (G1 in our nomenclature) positioned at $03^h01^m51.5^s,$ $+35^d50^m30^s$ \citep{b57}. The SDSS optical and 2MASS near-IR images of the central region of Abell~407 are shown in Figs ~\ref{fig1} and \ref{2mass_rass+150} (Left panel), and a zoomed version is shown in Fig. ~\ref{gmrt610} (Right panel), where we have also labeled the nine galactic condensations by letters G1 to G9 (see Table~\ref{tab1}). The first column in Table~\ref{tab1} also gives Zwicky's original nomenclature (within brackets) for these nine galaxies. \begin{figure*} \centering \includegraphics[width=5.5in,height=5in,keepaspectratio]{610_SDSS_i.png} \caption{Left panel: 4C~35.06 at 610 MHz mapped ($5.8^{\prime} \times 5.4^{\prime}$) with GMRT at $5\arcsec$ resolution, with the positions of central galaxies G1 (the brightest member), and G2, G5 and G6 (possible radio sources) marked on it. Right panel: the SDSS zoomed i-band image ($1.4^{\prime} \times 1.3^{\prime}$) of the central region together with the GMRT 610 MHz radio contours plotted on it. In this figure, all the nine galaxies comprising `Zwicky's Nonet' have been marked. North is up, and east is to the left.} \label{gmrt610} \end{figure*} \subsection{X-ray detection} The Abell 407 cluster is detected in the ROSAT All Sky Survey (RASS), with estimated X-ray luminosity and gas temperature of $5 \times10^{44}$ ergs$^{-1}$ (0.1-2.4 keV band) and 2.5 keV respectively \citep{ebel_98}. In MCXC (Meta Catalogue of X-ray Galaxy Clusters) it is listed as MCXC J0301.8+3550 \citep{Piffa2011} with an estimated cluster mass $M_{500} = 9.16 \times 10^{13}$ M$_{\odot}$, where $M_{500}$ is the total mass enclosed inside a radius $R_{500} = 675.4$ kpc, within which the mean over-density of the cluster is 500 times the critical density at the cluster redshift. From archival X-ray data in RASS hard band (0.5 -2.4keV) we obtained surface brightness contours that are overlaid on GMRT 150 MHz image (Figure.~\ref{2mass_rass+150}; right panel). Here the projected separation of the optical center (taken as G1, the brightest member) and the brighter X-ray peak to the southwest is $\sim 1.7\arcmin$ or $\sim 90$ kpc. Similar offsets in the optical and X-ray emission peaks have been observed in systems with ongoing mergers or in dynamically active clusters \citep{ry_08,man_12}. \begin{figure} \centering \includegraphics[width=3.5in,height=4in,keepaspectratio]{gmrt_610_color.png} \caption{Optical i-band image ($5.7^{\prime} \times 2.9^{\prime}$) of the galaxy cluster Abell~407 taken from the Sloan Digital Sky Survey (SDSS). The white contours show the 610 MHz radio emission morphology of 4C~35.06 as imaged with GMRT at $5\arcsec$ resolution. Both the radio and optical images have been rotated for convenience.} % \label{610_color} \end{figure} \subsection{Previous radio observations of 4C~35.06} One of the earliest detections of this source is at 1.4 GHz frequency with the Cambridge one-mile telescope with fairly poor resolution \citep{Riley75}. Subsequently, this source has been studied using VLA at 1.4 GHz and 5 GHz by \cite{b27}, revealing a bi-lobed structure at $\sim5\arcsec$ resolution. At 1.4 GHz the total flux density of 4C~35.06 is 728 mJy, a core flux density of 10 mJy, an eastern lobe of 305 mJy, and a western lobe of 416 mJy. At 5 GHz, the total flux density was 170 mJy, with core flux of 4 mJy, an eastern lobe of 55 mJy and a western lobe of 114 mJy. A 5 GHz VLBA observation at $4.87 \times 2.23$ mas$^{2}$ resolution detected a compact radio core of 2.6 mJy peak flux (3.5 mJy integrated) associated with the galaxy G3, which is the second brightest member of the nonet \citep{b28}. However, the VLBA scale is only about 3~pc across, while the large-scale radio structure extends over 200 kpc, leaving a vast gap in our understanding of the connection between the compact AGN and larger radio morphology. Moreover, the previous high frequency VLA maps all miss the large-scale jet structure and ultra-steep spectrum outer regions of this source. Recently, using high resolution ($5\arcsec$ FWHM) and high sensitivity (70 $\mu$Jy/beam rms) GMRT observations at 610 MHz frequency, \cite{b65} drew attention to the complete radio structure of an unusual, helically twisted, kinked-jet system of the radio source 4C~35.06. \cite{b97} have also studied this source at a very low radio frequency of 62 MHz with LOFAR at angular resolution of $50\arcsec$ FWHM. We discuss these observations along with our new GMRT observations in the sections below. \section[]{Observations and Data Reduction} \subsection{Multi-wavelength GMRT radio observations} We observed 4C~35.06 with GMRT at three frequencies; 610, 235 and 150 MHz (Project codes $21_{-}066$ and $26_{-}037$ \footnote[1]{https://naps.ncra.tifr.res.in/goa/mt/search/basicSearch}). Table~\ref{tab2} shows the log of radio observations. For flux and bandpass calibrations, 3C 48 was observed at the beginning and end of the observation runs for 10 minutes. The 30-minute scans on target source were alternated by 5 minute scans on the phase calibrator. The data were reduced using NRAO AIPS software package. AIPS tasks SETJY and GETJY were used for flux density calibrations \citep{b30}. The visibility data were flagged for Radio Frequency Interference (RFI) using AIPS tasks. The clean and calibrated solutions for the flux calibrator were used to calibrate the phase calibrator. The bandpass solutions were computed using the flux calibrator. These bandpass solutions were applied to the data and the frequency channels were averaged to increase the signal to noise ratio. The collapsed channel data were recalibrated for phase and amplitude solutions and later applied to the target source. AIPS task IMAGR was used for imaging in 3 dimensions, correcting for the W-term effects at low frequencies. For imaging, Briggs ROBUST weighting parameter was adjusted to detect low surface brightness diffuse emission regions better. Before the final imaging, several rounds of phase self-calibrations and one round of amplitude self-calibration were applied to the data. At 610 and 235 MHz, rms noise levels of 70 $\mu $Jy/beam and 0.90 mJy/beam were achieved, respectively. The rms noise measured in 150 MHz image was $\sim1$ mJy/beam. In Fig.~\ref{610_color}, the GMRT 610 MHz radio image of 4C~35.06 is shown overlaid on the SDSS i-band optical image. From the VLA archives we also created high frequency radio maps using data in D and C scaled arrays, observed at 6cm (4.8 GHz) and 20cm (1.4 GHz) wavelengths, respectively. Standard routines in AIPS were used for calibration and imaging. For spectral index mapping, we imaged both data sets with identical $15\arcsec$ (FWHM) angular resolutions. \begin{center} \begin{table*} \centering \caption{Details of radio observations.} \begin{tabular}{@{}llllll@{}} \hline\hline Telescope&Observed & Band &Obs. time & Beam & Map \\ &frequency & width & & (arc sec)& rms \\ \hline GMRT&610 MHz&32 MHz & 9hrs & 5.83 $\times $ 4.78 & 0.07 mJy/b \\ GMRT&235 MHz&6 MHz & 9hrs & 20.86 $\times $16.68 & 0.9 mJy/b \\ GMRT&150 MHz&16 MHz & 10hrs &19.87 $\times $ 15.77 & 1 mJy/b \\ VLA(NVSS)$^a$&1.4 GHz &100 MHz &survey data &45 $\times $ 45 & 0.4 mJy/b \\ VLA(VLSS)$^b$&74 MHz&1.56 MHz &survey data &80 $\times $ 80 & 100 mJy/b \\ \hline $^a$ \citep{b69} & $^b$ \citep{b70} \end{tabular} \label{tab2} \end{table*} \end{center} \begin{figure*} \includegraphics[width=6.5in,height=6in,keepaspectratio]{A407_GMRT.pdf} \caption{Colour scale GMRT images (Top Panel, left to right): at 610 MHz at resolution $5.83\arcsec \times 4.78\arcsec$, 235 MHz at resolution $20.86\arcsec \times 16.68\arcsec$, and 150 MHz (contour plot) with levels [1, 2, 4, 8, 16,......] $ \times 4 \;m$Jy/beam and resolution $23.9\arcsec \times 19.36\arcsec$. Bottom Panel: Colour scale images showing 150 MHz GMRT contours plotted over 1400 MHz NVSS image (Bottom left ) and 74 MHz VLSS image (Bottom right).} \label{gmrt} \end{figure*} \subsection{Spectroscopic Observations} We attempted to obtain good quality spectra for all nine galaxies comprising `Zwicky's Nonet'. The optical spectroscopic observations of the brighter galaxy members G1, G3, G4, G7 and G9 were taken with the IUCAA Girawali Observatory (IGO) 2m telescope and fainter galaxies G2, G5 and G6 with the Palomar 200-inch telescope. The aim was to characterize their AGN and star forming activity, and to estimate the central velocity dispersions ($\sigma$) and black hole masses ($M_{BH}$) using the well known $M_{BH} - \sigma$ correlations \citep{b3,b4}. Optical long-slit spectroscopic data were taken on 20-22 November 2011 on the IUCAA 2m telescope (IGO). The spectra were obtained using the IUCAA Faint Object Spectrograph and Camera (IFOSC) \footnote[2]{http://igo.iucaa.in}. We used two IFOSC grisms; grism no.7 and grism no.8 in combination with a 1.5-arcsec slit. These grisms provide a wavelength coverage of 3800-–6840 \AA\, and 5800–-8350 \AA. Standard stars were observed during the same nights for flux calibration. Wavelength calibration was done using standard Helium-Neon lamp spectra. Palomar 200-inch observations were carried out on 2014 January 23. The data were taken covering the wavelength range 3800 \AA \; to 9500 \AA, using the double spectrograph (blue and red arms). Wavelength calibration was done using standard Fe-Ar arc lamp spectra. Standard routines in the Image Reduction and Analysis Facility (IRAF) were used for data reduction and one-dimensional spectra were extracted using the {\em doslit} task in IRAF. The analysed data are added as supplementary material. \section[]{Results and Discussion} \subsection{The GMRT images of 4C 35.06: steep spectrum emission and source parameters} Figure~\ref{gmrt} upper row presents the GMRT low frequency radio images at 610, 235 and 150 MHz. The regions marked A1, B1 and A2, B2 represent the features on jets while C denotes the core region. The regions D1 and D2 indicate the outermost diffuse structures of the source. The forthcoming sections give a detailed discussion of these regions. We also show the radio images from the NRAO VLA Sky Survey (NVSS, 1420 MHz) and VLA Low-frequency Sky Survey (VLSS, 74 MHz) in the bottom panel of Figure~\ref{gmrt}. The contours of the GMRT 150 MHz image are plotted over the NVSS and VLSS images for size comparison. The highest resolution (5\arcsec FWHM) and currently the deepest yet GMRT image at 610 MHz shows a bright, complex core region and associated double sided twisted/helical jet structure, while the NVSS 1.4 GHz image does not resolve these structures due to its poor resolution (45\arcsec FWHM). The VLA 1.4 GHz image at $15\arcsec$ resolution detects the extended, twisting-turning jet structure (Fig. ~\ref{gmrt}, bottom left panel). The low frequency 235 and 150 MHz GMRT maps show extended, steep spectrum 'relic' plasma emission features (see Figure~\ref{gmrt}) at the extremities of the helically twisted jet structure. These features were also detected in 62 MHz LOFAR map [\cite{b97}]. At 610 MHz the source is found to have a total flux density of $1.7\pm 0.12$ Jy. The flux densities along the western and eastern jets are observed to be 580 mJy and 193 mJy respectively. The angular size of the source is $260\arcsec$ or linear size of $\sim220$ kpc. The maximum extent of the source at 235 MHz is $430\arcsec$ or $\sim380$ kpc. At 150 MHz the source is found to have the largest size of $460\arcsec$ or $\sim400$ kpc linear size. This implies that the linear extent of the source grows larger with the lowering of frequency, which is suggestive of steep-spectrum emission regions present at the extremities. At these frequencies, the western jet is brighter than the eastern jet (2.37 Jy and 1.18 Jy at 150 MHz, and 1.72 Jy and 800 mJy at 235 MHz respectively). The total source flux densities at 150 MHz and 235 MHz are 6.0$ \pm$0.18 Jy and 4.7$ \pm$0.13 Jy, respectively. The excellent quality GMRT images enabled us to make high resolution ($\sim 25 \arcsec$) spectral index maps down to 150 MHz. These maps are discussed further in the following sections. \subsection{Where is the AGN radio core?} On the $5\arcsec$ resolution GMRT 610 MHz map of 4C 35.06, there are three radio peaks near the center, of which two are shifted south from the jet direction (Figure.\ref{gmrt610}). A previous VLBA observation \citep{b28} has detected a compact radio core in the galaxy G3 on parsec scale. The two GMRT 610 MHz radio peaks are centered near optical galaxies G5 and G2, while the third radio peak on the jet axis is close to the faint galaxy G6. Figure~\ref{gmrt610} (left panel) shows the positions of G2, G5, G6 and the brightest galaxy G1 with `+' signs. \cite{b65} suggested that the probable host AGN emitting the bipolar jet could be galaxy G6 rather than G3, which is clearly offset southward from the principal jet direction. The galaxy G6 is very faint both in optical and infrared light (Table~\ref{tab1} and~\ref{tabk}). From spectroscopy, we obtained a velocity dispersion of (143$\pm$40) Km~s$^{-1}$ for G6, which yielded a relatively small black hole mass of $(0.52\pm0.65)\times 10^8 M_{\odot}$ (see section 4.7). However, the error margin is high due to the low SNR of the spectrum. Even though it is improbable (but not impossible) for a faint galaxy like G6 to produce such a large scale radio jet, it is possible that this now faint galaxy has been stripped off the majority of its outer halo stars in multiple tidal encounters, while still retaining a dense stellar core and black hole at the centre. Possibly this is also reflected in even smaller black hole mass, $(1.5\pm0.07)\times 10^7 M_{\odot}$ derived from its faint K-band luminosity of $M_K = -21.16$ (Table~\ref{tabk}). In an alternative scenario, \cite{b97} have suggested that the rapid movement of galaxy G3 and its episodic AGN radio activity is the reason for the observed peculiar radio morphology of 4C~35.06. In their interpretation, the large scale jet morphology is due to the earlier phase of activity of G3, which had switched off its radio emission and then restarted while it was moving to its current position, resulting in the offset inner double-lobed morphology and steep-spectrum larger jet structure to the north, similar to dying radio galaxies \citep{b98}. Thus, at present, we are observing an aged FRI-like large scale structure with an embedded restarted radio source. \cite{b97} also discuss that this AGN core is less likely to be G6 as argued by \cite{b65}, because of the lower mass of this galaxy and hence the lesser likelihood of it containing an SMBH. However, they have not considered the possibility of stripping of stars from the outer halo during tidal encounters in a dense environment, still retaining the SMBH at its center. In our opinion, the canonical black hole mass- IR bulge luminosity correlation applicable for normal ellipticals in relaxed environments may not hold good in the case of galaxies subjected to violent mergers. We suggest that the observation of the close coincidence of the optical positions of G2 and G5 with the radio peaks and the location of G6 near the center of symmetry of the large scale jets provide ample evidence to argue that the conclusion by \cite{b65} might be justifiable. However, much higher resolution radio or X-ray data are still necessary to identify the compact AGN core and the progenitor galaxy of the large scale jet structure firmly. \begin{figure} \centering \includegraphics[width=3.5in,height=5.0in,keepaspectratio]{Int_spec_RVII_01.pdf} \caption{The integrated radio spectrum of 4C 35.06. A power law fit is shown by a red line. Red points are the data taken from literature and green points the GMRT observations. A pair of spectra shown with blue and pink data points represent the diffuse, relic regions D1 and D2 detected in GMRT low frequency images (Figure ~\ref{gmrt}).\vspace{-0.5cm}} \label{int_spect} \end{figure} \subsection{Spectral Index Maps and Spectral Ageing Results} The integrated spectrum and spectral index maps are derived from the data available in the literature and our present GMRT observations. The broad-band spectrum shows an overall steep power-law from 26 MHz up to 4.9 GHz (Figure.~\ref{int_spect}). The GMRT data points clearly fit with the power-law spectral index $\alpha = -0.99$, thereby indicating overall steep spectrum nature of the source, as compared to typical radio galaxies with jets and lobes. \begin{figure*} \centering \includegraphics[width=6in,height=5in,keepaspectratio]{spix_a407.png} \caption{The spectral index maps obtained from 235 MHz vs. 610 MHz (left panel) and 235 MHz vs. 150 MHz GMRT images (right panel) using matched resolutions. The spectral index values are shown with a color bar on the right edge. In both plots, the spectral index errors in the central region ($\sim 0.02$ and $\sim 0.1$ ) are much lower than that near the jet extremities ($\sim 0.2$ and $\sim 1$ ). } \label{spix} \end{figure*} To understand the energetics of this radio source better, we created spectral index maps at low and high frequency ends, using radio maps convolved to the same resolution. Figure~\ref{spix} shows dual-frequency spectral index maps of the 235 vs. 610 MHz and 150 vs. 235 MHz bands. The spectral index maps clearly show that the radio emission in the central core region has flatter spectra in the range $\alpha = -0.5$ to $ -0.8$, whereas ultra-steep spectrum emission dominate at the outer extremities, with $\alpha \sim -2$ for 235-610 MHz and 150-235 MHz maps. The spectral indices for the diffuse, outermost relic regions marked D1 and D2 are estimated to be $-1.79$ and $-2.10$ respectively, indicating their ultra steep spectral nature (Figure \ref{gmrt} and \ref{int_spect}). This suggests that the radio emission in region D1 and D2 originates from an ageing radio plasma subjected to heavy energy losses, possibly resulting from a previous phase of energy injection in the region. The LOFAR observations at 62 MHz \citep{b97} also detect these ultra-steep spectral regions, but less clearly in comparison with the higher sensitivity GMRT images. The central region shows a flat spectrum, possibly due to the superposition of emission from a few radio-loud AGNs, while the steep spectrum towards the extremities can be attributed to the lack of fresh injection of accelerated particles and radiative energy losses. A high frequency spectral index map (Figure~\ref{vla}) was created by combining VLA D and C array maps at 6 cm (4.8 GHz) and 20 cm (1.4 GHz) wavelengths, both imaged at the same $15\arcsec$ angular resolution. This figure shows a flat spectral index central region with $\alpha \approx -0.5$ (dark shades) which steepens progressively away from the center along the twisting radio jets, with $\alpha > -2$ (light shades) at the extremities of the jets. Although the angular resolution of VLA maps is lower than that of GMRT, the spectral index trend is similar to that with GMRT maps. \begin{figure} \centering \includegraphics[width=3.5in,height=5in,keepaspectratio]{A407_spix-new.pdf} \caption{The spectral index map of 4C 35.06 shown in grayscale, obtained by combining D and C scaled-array VLA data at 6cm and 20cm wavelengths at $15\arcsec$ resolution. The contours are the VLA 20cm (1.4 GHz) radio contours at the same resolution, and contour levels in mJy/beam are shown at the bottom of the image. The optical galaxies in the center are denoted with `+' symbols.} \label{vla} \end{figure} \begin{figure} \centering \includegraphics[width=3in,height=3in,keepaspectratio]{D1t1.pdf}\\ \includegraphics[width=3in,height=3in,keepaspectratio]{D2t1.pdf} \caption{ The second order polynomial spectral fit (blue line) for the outermost regions D1 (Top panel) and D2 (bottom panel). The flux values are taken at frequencies 610, 235, 150 (GMRT) and 74 MHz (VLASS). The dashed red lines represent the tangents drawn at higher (610 MHz) and lower frequency (150 MHz) ends of spectra and the intersection point denotes the estimated break frequency (green dotted line). \vspace{-0.5cm} } \label{breaks} \end{figure} The radiative age of the non-thermal plasma in regions D1 and D2 was obtained from the spectral breaks in the integrated radio spectra of the regions shown in Figure~\ref{breaks}. The spectra are fitted with second order polynomials and tangents are drawn at 610 MHz and 150 MHz frequencies. The intersection point of these tangents gives the break frequency $\nu_{b}$, obtained as 308 MHz and 302 MHz for regions D1 and D2 respectively. Beyond the breaks, a steepening or a possible cut-off in the spectra is suggested, consistent with the scenario of radio emission emanating from a rapidly cooling electron population. The electron spectral age $t_{sp}$ (or cooling time scale) is then estimated from the synchrotron radio spectrum using the formula \citep{b98} \begin{equation} t _{sp} = 1.59 \times 10^{9} \left[\dfrac{B^{1/2}}{\left[ B^2+B_{IC}^2\right] \left[\nu _{b}(1+z)\right]^{1/2}}\right] \; yr \end{equation} This formula is obtained for a uniform magnetic field, neglecting expansion losses over the radiative age. Here $B$ is the magnetic field in $\mu$G, $z$ is the redshift, $ B_{IC}= 3.25 (1+z)^2$ $\mu$G is the inverse Compton equivalent magnetic field and $\nu_{b}$ is the cooling break frequency in GHz. An independent estimate of the magnetic field in the relic regions is needed for a robust estimate of $t_{sp}$. Assuming minimum energy condition, we calculated the energy density and magnetic field in the regions A1, A2 (the inner loop structures), and D1, D2 (the outer relic regions) (see Figure~\ref{gmrt}). The minimum energy density $u_{min}$ is given by: \begin{equation} u_{min} = \xi(\alpha,\nu_{1},\nu_{2}) (1+k)^{4/7} (\nu_o)^{4\alpha/7} (1+z)^{(12+4\alpha)/7} {(I_o/ d)}^{4/7} \end{equation} where k is the ratio of the energy of the relativistic protons to that of electrons, $\alpha$ is the spectral index, $\nu_{1}$ and $\nu_{2}$ are lower and higher limits of frequency, $\nu_o$ is the frequency at which surface brightness $I_o$ is measured, and function $\xi(\alpha,\nu_{1},\nu_{2})$ is tabulated in \cite{GF04}. Here we assume the filling factor to be 1, $\nu_{1}$ = 10 MHz, $\nu_{2}$ = 10 GHz, $\nu_o$ = 150 MHz and z = 0.047. The equipartition magnetic field can be expressed by: \begin{equation} B_{eq} = (\frac{24\pi}{7} \, u_{min})^{1/2} \end{equation} This way the magnetic fields for the relic regions are obtained as $\approx 5\mu $G and $\approx 16\mu $G for $k =1$ and $k=100$, respectively. The corresponding elapsed times are $\sim 170$ and $\sim 40$ Myr, respectively. We point out that \cite{b97} calculated the spectral age by estimating the time taken by the galaxy G3 (putative central AGN) to translate from the former location (at the center of the east-west jet direction), by approximating the radial velocity difference between the galaxy and the stellar envelope \citep{b26} as the velocity of the source in the sky plane. It was assumed that the source was shut off before the translation and restarted once it reached the new position. This approximation enabled them to adopt the translation time of G3 to be the shutdown period ($t_{off}$) of the AGN. An estimate of spectral age was made assuming equal time scales for active ($t_{on}$) and quiescent ($t_{off}$)phases of the AGN and adjudging the sum (($t_{on}$) + ( $t_{off}$)= 70 Myrs) as the age of the radio plasma. Assuming lowest the Lorentz factor values, they obtained a magnetic field of $10 \mu $G and a corresponding break frequency of $\sim 380$ MHz, which is very close to the break frequency ($\sim 300$ MHz) we obtained from the GMRT data points. Even though the restarted AGN activity model and the present model yield almost same values for break frequency, many assumptions had to be invoked to substantiate the former model.\\ \subsection{Twists and kinks in the jet: Dynamical signatures of a perturbed AGN ?} The GMRT 610 MHz maps (Figures~\ref{gmrt610} and ~\ref{gmrt}) reveal that the bipolar radio jet undergoes helical twists, with inversion symmetry, on either side of core C at the points marked A1 and A2 in Figures.~\ref{gmrt} and ~\ref{ss433}. Moreover, the north-west arm of the jet is observed to be brighter, and bends to form a prominent loop/arc starting from A2 up to point B2. A similarly twisted feature between points A1 and B1 is observed in the south-east arm as well, but fainter and more diffuse compared with the north-western counterpart. The region D2 in the western arm shows a peculiar, upward bent jet-like structure in the 150 MHz GMRT map, capped by a `mushroom' like feature at the top, the nature of which is not clear at the moment (Figure.~\ref{gmrt} upper panel). The symmetric counterpart of this knot/mushroom structure could be the downward bent feature D1 in the south-eastern section of the source. The extremely steep high-frequency radio spectra of the mushroom like feature at D2 (and also D1), located $\sim$ 200 kpc from the center, indicate significant energy losses, suggesting an absence of freshly injected particles. One possibility is that they are buoyant non-thermal plasma bubbles rising into the hot intra-cluster medium, inflated by the radio jet in the past, as observed around some BCG/cD galaxies residing in centers of clusters or groups \citep{Bagchi09,MN12}. % In addition to these interesting features there are a few sharp kinks or steps in the western arm of the jet at points marked C1, C2 and C3 (Figure.~\ref{ss433}). On the eastern side of the jet these kinks/steps are not discernible, possibly due to the projection effects or their absence. In our present work, lacking detailed modeling, it is difficult to decipher what these kinks in the jet flow represent physically, but they have been discussed further below ( in Section 4.5 $\&$ 4.6). \begin{figure} \centering \includegraphics[width=3.5in,keepaspectratio]{composite1.pdf} \caption{The high resolution (5$^{\prime\prime}$) grayscale image of the source 4C~35.06 at 610 MHz (GMRT) showing the twisted, helical jet structure. Different regions of the source are marked, and the end extensions D1 and D2 detected at 235 MHz and 150 MHz, are indicated by dotted lines. The colour image shows, for comparison, the cork-screw shaped precessing jets observed in galactic XRB `microquasar' SS433, which is the total intensity image at 4.85 GHz \citep{b52}. Note the linear size of SS 433 jet system is only 0.26 pc while that of 4C 35.06 at 610 MHz is 230 kpc. } \label{ss433} \vspace{-0.5cm} \end{figure} \subsection{Precessing radio jet structure: comparison with galactic microquasar SS~433} The deciphering of the peculiar radio jet morphology of 4C 35.06 gains added impetus when it is compared with a precessing radio jet structure observed in the galactic `microquasar' SS~433 (see Figure~\ref{ss433}). SS~433 is an X-ray binary (XRB) system in the center of supernova remnant W50, consisting of a stellar mass black hole or neutron star accreting matter from an A-type supergiant donor star \citep{b31,b32,b52}. The most unusual aspect of this object, modelled through radial velocity measurements of `shifting' $H_{\alpha}$ lines and high resolution radio imaging, is that the accretion disk around the compact object precesses with a regular period of $\sim 164$ days \citep{b31,b32}. Consequently, the axis of the jet-ejection nozzle also precesses with the same period and the ejected radio plasma traces out a dynamically changing `corkscrew' pattern (see Figure~\ref{ss433} and \cite{b52}). Even though a detailed modelling of the precession geometry of 4C 35.06 is beyond the scope of the present study, it is noticeable that its large-scale jet structure is analogous to that of SS 433, if we ignore the kinks (C1, C2 and C3) for the time being. The loop portions A2 - B2 and A1 - B1 are suggestive of the corkscrew pattern resulting from continuous change of the jet axis, possibly due to a precessing motion. Further, flat spectrum terminal hot spots are absent at the jet extremities of 4C 35.06 as well as SS~433, which is another indication of the continuous shifting of the jet direction. Moreover, the brightness of the western arm of the jet system in 4C~35.06 is nearly double that of the eastern arm. Depending on the inclination of the jet axis with our line of sight, the precession cone angle and the plasma bulk velocity, the projected radio morphology and brightness can appear to be quite different on the approaching and receding sides. This has been shown by \cite{b96} in numerical simulation of relativistic effects in precessing jets. The differences in the observed jet morphology on the two sides can be attributed as partly due to this. The precession of radio jets may be attributable to two mechanisms: first, the presence of a binary black hole system \citep{b76,b51}, where the torque exerted by the companion black hole can precess the accretion disk of the first object, leading to jet precession and secondly, the Lense-Thirring frame dragging effect \citep{bp75}; if the angular momentum vector of the accretion disk is misaligned with that of a fast spinning Kerr black hole, the black hole will try to frame drag the inner accretion disk so as to align it with its spin vector. This will lead to the precession of the accretion disk and of radio jets orthogonal to it. If this is the reason for the helically twisted jets in 4C~35.06, an interesting corollary is that the mass accreting supermassive black hole must be spinning. Moreover, the observed resemblance of the morphology of 4C 35.06 to the precessing relativistic jet system of X-ray binary SS~433 supports the fundamental paradigm that, in spite of a vast difference in involved black hole masses, length and time-scales, almost all relativistic disk-jet coupled phenomena happen in a scale-invariant manner in radio-loud AGNs and the galactic microquasars. \subsection{Jet energetics and interaction of jets with the ambient intracluster medium} In Figure~\ref{2mass_rass+150} (right panel), the GMRT 150 MHz radio image of 4C 35.06 is shown superposed on the ROSAT X-ray map in the 0.5 - 2.4 keV band. Here we observe that the energetic jet system of 4C 35.06 has expanded preferentially in the direction of lower gas pressure in the dense intra-cluster medium. This might affect the ambient X-ray medium by uplifting the gas along the mean jet flow direction. This radio jet feedback effect could be analogous to the distorted, extended lobes of supernova remnant W50 (the `Manatee' nebula), which are shaped by the interaction of powerful jets in microquasar SS~433 with the ambient ISM \citep{Dubner98}. We need much deeper X-ray observations of the A407 cluster to investigate the AGN feedback signatures better. Observationally, the kinetic power of a jet ($\bar{Q}$) is a key descriptor of the state of an accreting SMBH system: its mass, spin and the magnetic field of the accretion disk. Correlation of low-frequency ( $\nu \sim150$~MHz) radiative power of radio sources against their jet power shows that the radio luminosity of the jet constitutes only a small fraction ($<1\%$) of the total kinetic power \citep{b66,Daly2012}. Using the 150 MHz radio flux density of $6.0 \pm$0.18 Jy from the GMRT map, the time averaged kinetic power of jets in 4C 35.06 is computed as $\bar{Q} \approx 3 \times 10^{43}$erg~s$^{-1}$ \citep{b66}. We have not corrected for the (unknown) loss of energy in the outer radio lobes and thus $\bar{Q}$ is likely to be a lower limit. This power is below the transition value $5.0 \times 10^{43}$ erg s$^{-1}$ between FR I and II classes. However, if the jet continues to operate between $10^{7} - 10^{8}$ yr, the injected mechanical energy is $\sim10^{58} - 10^{59}$ erg, which is large enough to affect or even quench any cooling flow strongly and to drive large-scale outflows that redistribute and heat the gas on cluster-wide scales. If we further assume that this jet power is derived from accretion flow onto a black hole at the rate $\dot M$ and $\bar{Q} \sim 0.1 {\dot M} c^{2}$, we obtain ${\dot M} \sim 5.3 \times 10^{-3}$ M$_{\odot}$ yr$^{-1}$. This number is only representative, but it suggests accretion at sub-Eddington rate $\lambda = \bar{Q}/L_{edd} = 0.0024 \times (10^{8}/M_{BH})$, where $L_{edd}$ is eddington luminosity and $M_{BH}$ is the mass of the black hole. This low accretion rate signifies a radiatively inefficient accretion flow (RIAF) in a low-luminosity active galactic nuclei (so called LLAGN or LINER). The optical spectra of the galaxies in Zwicky's Nonet (refer Section. 4.7) also confirms this nature. However a complete picture of the launch of radio jet in 4C~35.06 and its energetic feedback effects requires much deeper and higher resolution X-ray and radio data. Systems hosting helically modulated symmetric jets are the most promising sites for finding close black hole systems (triple or binary)\citep{b51,b76}. In Zwicky's Nonet, it is observed that seven galaxy pairs are separated by distances $\; \sim$ 10 kpc in projection. Their redshift values are also close, with a mean $\bar z = 0.0469$ and standard deviation $\sigma_{z} = 0.00176$, or $v \sim 520$ km s$^{-1}$ (refer Table~\ref{tab1}). The number of binary or triple black hole systems discovered with projected separations less than a few kpc are very low \citep{b51}. In the present dense system of nine galaxies, all packed within a radius of only 25 kpc, the smallest projected separation between galaxy pair combinations is about 5 kpc (between G7 and G8). So the extreme closeness of the galactic members coupled with the helically twisted, large scale jet structure highlights the prime importance of this galaxy group in the search of multiple supermassive black hole systems and their gravitational and electromagnetic merger signatures. Moreover, merging SMBHs would also lead to an enhanced rate of tidal disruption of stars and possible gravitational wave recoil (slingshot) ejection of black holes from galaxies at speeds in excess of $1000$ km~s$^{-1}$. Even though the symmetric helical pattern observed in the jet structure might be explained with precession model, there are a few anomalies like the presence of distinct kinks or steps denoted by C1, C2 and C3 in the north-western arm that pose a challenge. The precession model alone may be insufficient to explain these features. We note that C-shaped twisted paired jet system in radio source 3C~75 is associated with a binary black hole pair separated by only 7~kpc \citep{b2}. Both the jets in 3C~75 show prominent wiggled and kinked structures, which have been modelled as due to the combined linear and orbital motion of the bound binary black hole pair \citep{Yokosawa85}. In 4C~35.05 observed kinks could arise from a similar system, where a black hole with radio jets is orbiting another one at high speed and with large orbital eccentricity \citep{Yokosawa85}. However, in this model, we can not easily explain why these kinks or steps are absent in the south-eastern arm of the jet. \subsection{Optical spectroscopic results: AGN signature and black hole mass estimation} Previously, \cite{b26} measured the redshifts and stellar velocity dispersions ($\sigma$) for a few of galaxies in Zwicky's nonet covering a wavelength range from $3700$ to $5250 \AA$. In our present study, we have obtained good S/N spectra over the wavelength range of $3800 $ to $8500 \AA$ for eight out of the nine galaxies comprising Zwicky's Nonet (The spectra are included as supplementary material). % However, attempt to obtain a fair spectrum of galaxy G8 failed, due to it being very faint. Our main aim was to search for the signs of AGN or star forming activity in the optical spectra of these galaxies and to estimate their black hole masses from the stellar velocity dispersion. The spectra of these galaxies resemble those of passive, early type red ellipticals, devoid of any major emission lines. This is not unusual as optical emission lines are found to be absent in many AGNs showing radio emission and large scale radio jets. It has been observed that many FRI radio sources in galaxy clusters are hosted by galaxies showing very weak or no optical emission lines \citep{b37,b67}. \begin{center} \begin{table} \centering \caption[caption]{K band absolute magnitudes and masses of the \\\hspace{\textwidth} SMBHs associated with the nine galaxies of `Zwicky's Nonet'. } \begin{tabular}{@{}|c|c|c|@{}} \hline Source & K Band absolute & Mass of the \\ & magnitude & SMBH. \\ & & ($ 10^8 M_{\odot} $) \\ \hline \hline G1 &-23.46$\pm$0.039 &1.13 $\pm$0.30\\ G2 &-21.99$\pm$0.085 &0.31$\pm$0.12\\ G3 &-23.49$\pm$0.034 &1.16$\pm$0.30 \\ G4 &-22.60$\pm$0.028 & 0.53$\pm$0.17\\ G5 &-22.73$\pm$0.031 &0.60$\pm$0.18 \\ G6 & -21.16$\pm$0.089 &0.15$\pm$0.07 \\ G7 &-23.56$\pm$0.031 &1.23$\pm$0.30 \\ G8 & -20.06$\pm$0.113 &0.06$\pm$0.03 \\ G9 &--22.50$\pm$0.039 &0.49$\pm$0.16 \\ \hline \end{tabular} \label{tabk} \end{table} \end{center} Internal properties of a galaxy, such as mass and accretion rate of a SMBH are better estimated from the nuclear emission lines \citep{b39}. The observed spectra show that all the suspected radio loud galaxies (G2, G3, G5 and G6) belong to the class of low excitation radio galaxies (LERGs) \citep{b41,b43}. LERGs are mostly found to be hosted by BCGs having extended cD like light profiles \citep{b67}, similar to what we find in Zwicky's Nonet. The well-known tight correlation which connects the mass of the central black hole $M_{BH}$ to the galaxy's bulge stellar velocity dispersion $\sigma$ \citep{b3,b4} is given by \begin{equation} \log_{10}\left(\dfrac{M_{BH}}{M_{\odot}}\right) = \alpha + \beta \log_{10}\left(\frac{\sigma}{ \text{200 km\,s}^{-1}} \right) \end{equation} where $\sigma$ is expressed in km s$^{-1}$. Here we have used $\alpha = 8.38$ and $\beta= 4.53$, as derived in \cite{b95}. The estimated SMBH masses are tabulated in Table~\ref{tab0}. The black hole masses were also calculated using the slightly different $\alpha $ and $\beta $ values taken from \cite{b5} and \cite{b6}. These masses are consistent within one sigma limits with the numbers given in Table~\ref{tab0}. These results show that galaxies G1, G3, G5, G7, and G9 all host supermassive black holes of mass ($M_{BH} \approx few \times 10^{8} \, M_{\odot}$). For the other three galaxies G2, G4 and G6, the estimate of $M_{BH}$ has large errors. Interestingly, the most massive black hole of mass $M_{BH} \approx 10^{9} \, M_{\odot}$ resides in the galaxy G3, which showed a radio loud AGN core in previous VLBA observations \citep{b28}. The galaxy black hole masses are also calculated from their K-band magnitudes using 6 times deeper data on cluster A407 available from 2MASS survey. The equation connecting K-band absolute magnitude ($M_K$), of bulge component to the central black hole mass given by \cite{graham_2007} is, \begin{equation} \log_{10}\left(\dfrac{M_{BH}}{M_{\odot}}\right) = -0.38(\pm 0.06)\left(M_K+24\right) +8.26(\pm 0.11) \end{equation} Table~\ref{tabk} lists the black hole masses estimated with this method for the nine galaxies. The following caveats are worth mentioning here: It is unclear whether the canonical $M_{BH}$-$\sigma$ relation will suffice for galaxies in such a hostile environment, undergoing violent mergers and stripping of stars in multiple tidal encounters. This is clearly evidenced by the formation of a large-scale stellar halo of stripped matter in Zwicky's Nonet. The same concern applies if one were to obtain $M_{BH}$ from the K-band magnitude of bulge using the $M_{BH}$-$M_{K}$ correlation. Moreover, effect of the gravitational potential of the background stellar halo (which is highly dark matter dominated; \citep{b26}) and close merging galaxies on the bulge stellar velocity dispersion of a galaxy are also possible factors that need to be accounted for in black hole mass calculations. In this article, we have not attempted to do so. However, for checking this issue, the last column of Table~\ref{tab0} shows the ratio of black hole mass, from the $M_{BH}$-$\sigma$ correlation ($M_{BH,\sigma}$) to that obtained from $M_{BH}$-$M_{K}$ method ($M_{BH,K}$). The ratio $M_{BH,\sigma}/M_{BH,K}$ is $> 2 $ for galaxies with well determined black hole masses, which suggests that possibly $M_{BH}$-$M_{K}$ method gives smaller black hole masses because of the truncation of the outer envelope of galaxies, which reduces their K-band luminosity. Alternatively, the black hole masses from the $M_{BH}$-$\sigma$ relation are overestimated. The stripped away matter from the presently observed nuclei must provide a large fraction of the total luminosity of the observed stellar halo. The main parameters of the stellar halo, which is detectable up to the r-band surface brightness limit of $\sim 24$ mag arcsec$^{-2}$ (and possibly beyond), quoted by \cite{b26} are as follows; central mass density $\rho(0) = 0.63\pm 0.25 M_{\odot}$ pc$^{-3}$, mass-to-light ratio in r band $M/L = 90\pm35$, and halo radial velocity dispersion $\sigma_{halo} = 610 \pm 200$ km s$^{-1}$. From this value of $\sigma_{halo}$ and taking halo radius $r \approx 30\arcsec$ ($\sim 26.5$ kpc), we obtain the total dynamical mass of halo as $2.2 \times 10^{12}$ M$_{\odot}$, which interestingly is of the same order as that of a super giant cD galaxy. \begin{center} \begin{table} \centering \caption[caption]{The redshifts and masses of the SMBHs associated \\\hspace{\textwidth} with the galaxy like condensations in `Zwicky's Nonet'.\\\hspace{\textwidth} The last column shows the ratio of the SMBH black hole masses \\\hspace{\textwidth} obtained from stellar velocity dispersions and K band magnitudes.} \begin{tabular}{@{}|c|c|c|c|c|@{}} \hline Galaxy &{\hskip -0.5cm} Red & Velocity & Mass of the &SMBH\\ &{\hskip -0.50cm} shift& dispersion & SMBH & mass ratio\\ & & $(km s^{-1})$ & ($ 10^8 M_{\odot} $)&$M_{BH,\sigma} /M_{BH,K}$ \\ \hline \hline G1 &{\hskip -0.50cm}0.0473 &222 $\pm$16 &3.88$\pm$ 1.23&2.76\\ G2 &{\hskip -0.50cm}0.0476 &143 $\pm$27&0.53 $\pm$ 0.44&1.71\\ G3 &{\hskip -0.50cm}0.0470 &273$\pm$18& $9.83\pm 2.96 $&8.47 \\ G4 &{\hskip -0.50cm}0.0503 & 135$\pm$33 & 0.40 $\pm$ 0.45&0.76\\ G5 &{\hskip -0.50cm}0.0476 &230$\pm$8 &4.52$\pm$ 0.74&7.5\\ G6 &{\hskip -0.50cm}0.0454 &143$\pm$40 &0.52$\pm$ 0.65&3.47\\ G7 &{\hskip -0.5cm}0.0468 &211$\pm$16 &3.08$\pm$1.00&2.50\\ G9 &{\hskip -0.50cm}0.0445 &176$\pm$20 &1.34$\pm$0.68 &2.74\\ \hline \end{tabular} \label{tab0} \end{table} \end{center}

We have presented the results of our radio, optical and infra-red observations of the radio source 4C~35.06, located in the central region of the galaxy cluster Abell 407. The cluster center hosts a compact ensemble of nine passive, red elliptical galaxies embedded within a faint, diffuse stellar halo. We proposed to name this galactic system `Zwicky's Nonet'. GMRT observations at 150, 235 and 610 MHz clearly reveal the complete radio structure of 4C~35.06, with a complex central core region and helically twisted and kinked bipolar radio jets extending up to $\sim 400$~kpc. The radio jets terminate into diffuse, ultra-steep spectrum `relic/fossil' plasma lobes D1 and D2. In D2, a peculiar, very steep spectrum ($\alpha < -2$) mushroom like feature is discovered from GMRT 150 MHz map. In regions D1 and D2 of 4C~35.06, the average minimum energy magnetic field is $B \sim 5$~$\mu$G for $k =1$ and $B \sim 16$~$\mu$G for $k =100$. The corresponding spectral ages of electrons are obtained as $ 170 \times 10^6$ and $ 40 \times 10^6$ yr respectively. The time averaged kinetic power of jets is estimated to be $\approx 3 \times 10^{43}$erg~s$^{-1}$, indicating that the source is a FR I type radio galaxy. The unique helical jet system and the very compact configuration of nine galactic nuclei point to the possibility of precessional and orbital motion of the AGN. This also suggests possible gravitational perturbation effects of multiple black holes residing in the extremely dense central region of the cluster. In such an environment, orbital decay assisted by dynamical friction causes the central binary black holes of galaxies to merge, while gravitational torque in the binary phase may cause the accretion disk of AGN to precess, resulting in a helical jet pattern. The morphological similarity of this jet system with that of the galactic microquasar SS 433 also supports a precessional scenario. The absence of terminal hot spots and presence of ultra-steep spectrum regions on both ends of the jet strongly suggest the continuous shifting of the jet direction, further supporting the precessional model. Our study points to the possibility of the fainter member (G6) of the Zwicky's Nonet hosting large-scale radio jets. The faintness of this galaxy is attributed to the stripping of its major stellar envelope due to the tidal interactions in galactic mergers, retaining the SMBH at the center. The high ratio ($> 2 $) of black hole masses from stellar velocity dispersion and K-band luminosity, i.e. $M_{BH,\sigma}/M_{BH,K}$, for galaxies with well determined black hole masses corroborates the diminution of the stellar envelope. The observation of a diffuse stellar halo of stripped matter in the system supports this scenario. The optical spectra of eight galaxies in Zwicky's Nonet fail to show any prominent emission lines, indicating a radiatively inefficient accretion flow onto the central black holes at sub-Eddington rates. No strong star-formation/star-burst activity detected in any of the galaxy spectra. Further high sensitivity and higher resolution radio observations are needed to provide a complete spectral analysis and to obtain the detailed resolved central morphology of this complex source. A deep X-ray observation of the hot intra-cluster gas around the cD galaxy precursor, and detection of the AGNs and their X-ray spectra would be very beneficial in deciphering the nature of this puzzling radio galaxy.

1607.05080

1607

1607.05049_arXiv.txt

We present the analysis of 35.5 square degrees of images in the 1\,--\,0\,S(1) line of \htwo\ from the UK Widefield Infrared Survey for \htwo\ (UWISH2) towards Cassiopeia and Auriga. We have identified 98 Molecular Hydrogen emission-line Objects (MHOs) driven by Young Stellar Objects, 60\,\% of which are bipolar outflows and all are new discoveries. We estimate that the UWISH2 extended emission object catalogue contains fewer than 2\,\% false positives and is complete at the 95\,\% level for jets and outflows brighter than the UWISH2 detection limit. We identified reliable driving source candidates for three quarters of the detected outflows, 40\,\% of which are associated with groups and clusters of stars. The driving source candidates are 20\,\% protostars, the remainder are CTTSs. We also identified 15 new star cluster candidates near MHOs in the survey area. We find that the typical outflow identified in the sample has the following characteristics: the position angles are randomly orientated; bipolar outflows are straight within a few degrees; the two lobes are slightly asymmetrical in length and brightness; the length and brightness of the lobes are not correlated; typical time gaps between major ejections of material are 1\,--\,3\,kyr, hence FU-Ori or EX-Ori eruptions are most likely not the cause of these, but we suggest MNors as a possible source. Furthermore, we find that outflow lobe length distributions are statistically different from the widely used total length distributions. There are a larger than expected number of bright outflows indicating that the flux distribution does not follow a power law.

The formation of stars via disk accretion of material is inevitably related to mass ejection into jets and outflows along the rotational axis of these objects. The outflows from protostars and Young Stellar Objects (YSOs) were first correctly recognised as such by \citet{1980ApJ...239L..17S}. Since then numerous studies of these jets and outflows have been conducted to investigate the details of the excitation mechanism, the mass ejection rates, the jet launching, acceleration and collimation, as well as the outflow energetics and timescales (see reviews of e.g. \citet{1996ARA&A..34..111B}, \citet{2000prpl.conf..867R}, \citet{2007prpl.conf..215B}, \citet{2014prpl.conf..451F}). However, there are still a number of unsolved questions about the outflow phenomenon. Which factors statistically determine the properties (length, luminosity, formation of the main \htwo\ emission features -- knots) of the outflows? Are the properties of the central source (final mass, age) responsible for these or has the environment (density structure in low mass vs. high mass star forming regions) a significant influence? Is the energy and momentum feedback from the outflows significant enough to explain the local turbulent energy near the forming stars and are the jets and outflows able to terminate the star formation process locally? In order to answer these and similar questions, we need to observationally establish the number of jets and outflows from young stars in the Galactic plane and to determine their average properties. The first truly large scale work to identify all jets and outflows in a star forming region via an unbiased survey was done in Orion\,A by \citet{Stanke2002}. They utilised the molecular hydrogen ro-vibrational 1\,--\,0\,S(1) transition at 2.122\,$\mu$m. This line is a proven excellent tracer of hot (T$\sim$2000\,K) and dense (n$\ge$10$^3$\,cm$^{-3}$) gas excited by the fast shocks ($10 - 100$\,km\,s$^{-1}$) caused by the interactions of jets and outflows with the surrounding interstellar medium. As the line is in the K-band, it is less influenced by extinction compared to other tracers of these shocks such as optical H$\alpha$ or [SII] lines, which are the historically used tracers for Herbig-Haro objects. It is usually also stronger than the near infrared (NIR) [FeII] lines, except in strong shocks or purely atomic environments. Observations in other molecular outflow tracers, such as CO or SiO, on large (molecular cloud) scales typically lack the combination of spatial resolution and depth to identify the fainter outflows, which are detectable in the NIR, especially in complex regions along the Galactic plane. Furthermore, the 1\,--\,0\,S(1) line flux is proportional to the total outflow luminosity for a range of excitation conditions \citep{CarattioGaratti2006}. Note that despite the shock velocity and gas density limitations for the excitation of the 1\,--\,0\,S(1) transition, at least some parts of the vast majority of jets and outflows from YSOs are detectable in this line. After the pioneering work by \citet{Stanke2002}, further searches for jets and outflows in star forming regions have been conducted e.g. by \citet{Walawender2005}, \citet{Hatchell2007}, \citet{Davis2009} and \citet{Khanzadyan2012}. In order to establish a truly unbiased sample of jets and outflows from young stars in the Galaxy, not restricted to nearby, mostly low-mass star forming regions, the UKIRT Widefield Infrared Survey for \htwo\ (UWISH2) was conducted by \citet{Froebrich2011}. In this series of papers we are analysing in detail the identified jets and outflows from young stars in this survey as Molecular Hydrogen emission-line Objects (MHOs) defined in \citet{2010A&A...511A..24D}. Previous works based on UWISH2 data by \citet{Ioannidis2012,2012MNRAS.425.1380I} have concentrated on the Serpens and Aquila region along the plane, while \citet{2012ApJS..200....2L,2013ApJS..208...23L} have investigated \htwo\ ouflows from Spitzer-detected extended green objects. This paper is structured as follows: In Sect.\,\ref{dataandanalysis} we briefly discuss the data and our analysis procedures such as the identification of the MHOs and and their most likely driving sources. The results, outflow and driving source properties, are then discussed in detail in Sect.\,\ref{results}.

We have analysed the population of jets and outflows from young stars in a 35.5 square degree region along the Galactic Plane in Cassiopeia and Auriga utilising \htwo\ imaging data from the UWISH2 survey \citep{Froebrich2011} and the catalogue of potential jet/outflow features from \citet{2015MNRAS.454.2586F}. In the investigated area we have identified 98 Molecular Hydrogen emission-line Objects, i.e. potential jets and outflows from young stars, all of which are new discoveries. When scaled up to the entire UWISH2 survey, we thus expect a total sample of about 1500 MHOs for analysis. The detected MHOs are classified as bipolar outflows (60\,\%), single-sided outflows (20\,\%) and individual or small groups of \htwo\ knots (20\,\%). Comparing the number of identified MHOs to the automatically generated and classified catalogue of UWISH2-detected \htwo\ features, we can conclude that the catalogue from \citet{2015MNRAS.454.2586F} contains much less than 10\,\% of false positives amongst the \htwo\ features classified as jets and outflows from young stars. Most of these false positives can be attributed to fluorescently excited edges of molecular clouds. The entire \htwo\ feature catalogue contains only 2\,\% of false positives, i.e. features that are not caused by emission from \htwo. These are mostly variable stars and image artefacts. Finally, the catalogue of jet/outflow features in \citet{2015MNRAS.454.2586F} is complete at the 95\,\% level for objects above the UWISH2 detection limit. Only a small number of low surface brightness features detectable in the UWISH2 data have not been included in the automatically generated catalogue. We could identify driving source candidates for about 75\,\% of the MHOs and about 40\,\% of these objects are associated with groups or clusters of stars, while the remaining 60\,\% seem to be more isolated, indicating that clustered star formation could inhibit the formation of detectable larger outflows. We have identified 15 new star cluster candidates near the MHOs in the survey area. Of the WISE detected driving source candidates, about 20\,\% have positive slopes of their SED, i.e. are protostellar source candidates, while the remaining 80\,\% are most likely CTTSs. This is also supported by their NIR colours. 70\,\% of the driving sources with multiple NIR K-band detections (UKIDSS GPS and 2MASS) show a variability of more than 0.1\,mag and about 20\,\% of the sources vary by more than 0.5\,mag over a timescale of several years. The position angles of the identified outflows have a 99.99\,\% probability to be drawn from a homogeneous distribution. For the typical bipolar outflow the two lobes have position angles within 5\degr. About 10\,\% of the MHOs form X-shaped outflows which could be originating from binary sources \citep{2016ApJ...820L...2L}. The length ratio (short over long) of the lobes of the bipolar outflows have typical values between 0.6 and 1.0, with a median of 0.72. Only about 20\,\% of objects are highly asymmetric with length ratios of less than 0.5. The flux ratios (short over long) typically show a much wider spread (between 0.2 and 5.0) and are not correlated with the length ratio. We measured the length of all outflow lobes and investigated their distributions. We find that the length measurements typically used in the literature (end-to-end for bipolar outflows and source-to-end for single-sided flows) should be avoided. This is due to the typical asymmetry of the bipolar outflows. It causes the lobe length distributions and total length distributions to be different with a probability of 99.95\,\%. There is no apparent difference in the lobe length distributions of outflows from younger and older driving source candidates. The dynamical timescales of the outflows are up to 10\,kyr, while the typical timescales associated with the gaps of large \htwo\ knots in the lobes correspond to 1\,--\,3\,kyr. This indicates that neither FU-Ori or EX-Ori style outbursts are likely to be responsible for the formation of the larger \htwo\ knots seen in the typical outflows. Potentially the population of recently identified NIR eruptive variables, or MNors, with properties in-between FU-Ori and EX-Ori objects, could be the cause for the \htwo\ knot formation, if their as-yet unknown occurrence rate has the correct timescale. The flux distributions of \htwo\ knots and outflow lobes is generally better fit by an exponential than by a power law. This is caused by a increased number of bright knots and lobes (above $30 \times 10^{-18}$\,W\,m$^{-2}$) compared to the expectation for a power law distribution. There are no differences in the flux distribtions for outflows from younger or older driving sources. The completeness limit for the detection of \htwo\ knots in the outflows is estimated as $10^{-18}$\,W\,m$^{-2}$. The number of \htwo\ bright ($> 10^{-3}$\,L$_\odot$ in the 1\,--\,0\,S(1) line) outflows per square parsec in Aquila is comparable to investigations in the inner Galactic Plane. In a future paper we will investigate these trends for a much larger sample of MHOs detected in the Cygnus\,X region where we expect a four- to five-fold increase in the number of MHOs. This will also be used to investigate if any of the trends are caused by environmental effects.

1607.05049

1607

1607.07750_arXiv.txt

The largest geomagnetic storm so far in the solar cycle 24 was produced by a fast coronal mass ejection (CME) originating on 2015 March 15. It was an initially west-oriented CME and expected to only cause a weak geomagnetic disturbance. Why did this CME finally cause such a large geomagnetic storm? We try to find some clues by investigating its propagation from the Sun to 1 AU. First, we reconstruct the CME's kinematic properties in the corona from the SOHO and SDO imaging data with the aid of the graduated cylindrical shell (GCS) model. It is suggested that the CME propagated to the west $\sim$$33^\circ$$\pm$$10^\circ$ away from the Sun-Earth line with a speed of about 817 km s$^{-1}$ before leaving the field of view of the SOHO/LASCO C3 camera. A magnetic cloud (MC) corresponding to this CME was measured in-situ by the Wind spacecraft two days after the CME left LASCO's field of view. By applying two MC reconstruction methods, we infer the configuration of the MC as well as some kinematic information, which implies that the CME possibly experienced an eastward deflection on its way to 1 AU. However, due to the lack of observations from the STEREO spacecraft, the CME's kinematic evolution in interplanetary space is not clear. In order to fill this gap, we utilize numerical MHD simulation, drag-based CME propagation model (DBM) and the model for CME deflection in interplanetary space (DIPS) to recover the propagation process, especially the trajectory, of the CME from $30 R_S$ to 1 AU under the constraints of the derived CME's kinematics near the Sun and at 1 AU. It is suggested that the trajectory of the CME was deflected toward the Earth by about $12^\circ$, consistent with the implication from the MC reconstruction at 1 AU. This eastward deflection probably contributed to the CME's unexpected geoeffectiveness by pushing the center of the initially west-oriented CME closer to the Earth.

As the most important driver of severe space weather, coronal mass ejections (CMEs) and their geoeffectiveness have been studied intensively. Previous statistical studies have shown that not all the front-side halo CMEs are geoeffective~\citep[e.g.,][]{Webb_etal_2001, Wang_etal_2002a, Zhao_Webb_2003, Yermolaev_etal_2005}, and not all non-recurrent geomagnetic storms can be tracked back to a CME~\citep[e.g.,][]{Cane_etal_2000, Cane_Richardson_2003, Yermolaev_etal_2005, Zhang_etal_2007}. These phenomena may cause some failed predictions of the geoeffectiveness of CMEs. The recent notable event exhibiting such a failure was on 2015 March 15 when a fast CME originated from the west hemisphere. Space Weather Prediction Center (SWPC) of NOAA initially forecasted that the CME would at most cause a very minor geomagnetic disturbance labeled as G1, a scale used by SWPC to measure the intensity of geomagnetic storms. However, the CME produced the largest geomagnetic storm so far, at G4 level with the provisional $Dst$ value of $-223$ nT, in the current solar cycle 24~\citep{Kataoka_etal_2015, Wood_etal_2016}. The major geomagnetic storm was called the ``2015 St. Patrick's Day" event as its main phase and peak occurred on March 17, and the surprising CME was selected as a campaign event by International Study of Earth-affecting Solar Transients (ISEST)\footnote{\url{http://solar.gmu.edu/heliophysics/index.php/03/17/2015_04:00:00_UTC}}, a program under SCOSTEP, and also by Coupling, Energetics and Dynamics of Atmospheric Regions (CEDAR)\footnote{\url{http://cedarweb.vsp.ucar.edu/wiki/index.php/2015_Workshop:The_March_17_2015_great_storm}}. \begin{figure*}[b] \centering \includegraphics[width=\hsize]{lasco-gcs} \caption{(a)--(b) SOHO/LASCO C3 difference images showing the fast CME (blue) as well as the preceding slow CME (red) with the GCS fitting meshes superimposed. (c)--(e) Longitudes, latitudes and heights of the leading edges of the two CMEs obtained from the GCS fitting. The line in Panel (g) is the linear fit to the CME height assuming a reasonable uncertainty of $\pm1 R_S$.}\label{fg_gcs} \end{figure*} Such an unexpected phenomenon naturally raises the first question for the forecasting of the geoeffectiveness of a CME, i.e., whether or not a CME will hit the Earth even though we know the source location and initial kinematic properties of the CME. A full understanding of the propagation trajectory of a CME from the Sun to 1 AU is the key to this question. Of course, it is not the only factor determining the geoeffectiveness of a CME. The magnetic field strength and the orientation of the CME flux rope, which directly affect the strength and duration of the interval of the south-component of the magnetic field, is also important for determining its geoeffectiveness. It has been well accepted that the CME's trajectory can be deflected in the corona (within a few tens of solar radii, $R_S$). \citet{Wang_etal_2011} illustrated that such deflections can be classified into three types: asymmetrical expansion, non-radial ejection and deflected propagation. In a statistical sense, CMEs tend to be deflected toward the equator during solar minimum~\citep[e.g.,][]{MacQueen_etal_1986, Cremades_Bothmer_2004, Wang_etal_2011} or deflected away from coronal holes~\citep[e.g.,][]{Gopalswamy_etal_2003, Gopalswamy_etal_2009, Cremades_etal_2006}. The physics behind these deflections is that the gradient of the magnetic energy density may cause the CME to move toward the place where the magnetic energy density reaches the minimum, usually the location of the heliospheric current sheet~\citep{Shen_etal_2011, Gui_etal_2011, Zuccarello_etal_2012, Isavnin_etal_2013, Kay_etal_2013}. Although the CME's deflection in the corona could be tens of degrees and may change the geoeffectiveness of a CME, it still can be monitored by coronagraphs~\citep[e.g.,][]{Mostl_etal_2015}. Thus, the possible deflection of a CME in interplanetary space rather than the deflection in the corona is one of the major sources of uncertainty in the prediction of the CME impact at the Earth. The possibility of the CME deflection in interplanetary space was first proposed by~\citet{Wang_etal_2004b}. They suggested that, different from the deflection in the corona, the CME's trajectory in interplanetary space could be deflected due to the velocity difference between the CME and the ambient solar wind. For a fast CME, the solar wind plasma and interplanetary magnetic field will be piled up from the west and ahead of it, leading to a net deflection force toward the east; for a slow CME, the picture is the opposite. A kinematic model (called DIPS, {\it Deflection in InterPlanetary Space}, hereafter) was therefore developed~\citep{Wang_etal_2004b}. Such deflections are thought to be gradual and much slower than that in the corona, but the total amount of the deflection angle is comparable to that in the corona as it takes place over a much longer distance. Such evidence can also be found in previous studies~\citep[e.g.,][]{Wang_etal_2002a, Wang_etal_2004b, Wang_etal_2006a, Kilpua_etal_2009, Lugaz_etal_2010, Isavnin_etal_2014, Kay_Opher_2015}. One of the most comprehensive analysis of the CME's trajectory in interplanetary space was done by~\citet{Wang_etal_2014} for a slow CME, which was proven to experience a westward deflection all the way from the corona to 1 AU with a total deflection angle of more than 20 degrees. For the 2015 March 15 CME, there were no STEREO~\citep[Solar TErrestrial RElations Observatory,][]{Kaiser_etal_2008} data as the twin spacecraft were behind the Sun and not taking images. All the information of the CME came from the remote-sensing data provided by the Solar and Heliospheric Observatory (SOHO) and the in-situ data by Wind (or ACE) spacecraft at 1 AU. The interplanetary space between the corona to 1 AU thus had an observational gap, and therefore the propagation of the CME from the Sun to 1 AU is unclear. In this paper, we try to recover the kinematic evolution of the CME from the limited observations and fill the gap with the aid of models. We particularly focus on the trajectory of the CME to demonstrate how the CME behavior in interplanetary space favors its strong geoeffectiveness.

Here, we study the fast CME originating on 2015 March 15, which caused the largest geomagnetic storm so far in solar cycle 24. With the aid of forward modeling, the kinematic properties of the CME in the corona are obtained based on the solar imaging data. Within 30 $R_S$ the CME propagated west of the Sun-Earth line along the longitude of about $30^\circ$ and the latitude of about $15^\circ$. The propagation speed is about 817 km s$^{-1}$. Its angular width is wide enough to make its east flank overlap the Sun-Earth line. On the other hand, the in-situ data at 1 AU suggests the flank of a MC arriving at the Earth. Two different models are applied to the MC to infer the configuration of the MC. It is found that the handedness is consistent but the orientation of the flux rope is quite different between the models. As suggested by \citet{Riley_etal_2004}, such a large deviation in the orientation is probably due to the spacecraft being too far away from the MC's axis. Currently, it is still difficult to evaluate the reliability of the model results. Due to the lack of the STEREO observations, the propagation of the 2015 March CME in the heliosphere is unclear. We then try to recover the interplanetary evolution process of the CME from the information at the two ends: near the Sun and at 1 AU. The VFR model results based on the in-situ data suggest that the CME experienced a significant eastward deflection with the ratio of $\frac{v_y}{v_x}$ of about $-0.1$, implying a $12^\circ$-deflection toward the Earth. The detailed trajectory of the CME between the two ends are further reconstructed by using the numerical simulation (for background solar wind), DBM model (for the CME propagation speed) and DIPS model (for the CME trajectory) under the constraints of the CME's kinematics obtained from the solar and in-situ observations. The reconstructed trajectory is bent toward the Earth, quite consistent with the deflection implied by the VFR model. This eastward deflection pushed the center of the initially west-oriented CME closer to the Earth and probably contributed to the unexpected strong geoeffectiveness of the CME. However, the lack of interplanetary observations causes that the above inference for this case cannot be fully validated though it sounds reasonable, and the origin of this strong geomagnetic storm is still somewhat mysterious. Some models applied in this study can be run and tested online. One can go to \url{http://space.ustc.edu.cn/dreams/} for the DIPS model and the VFR model, and to \url{http://oh.geof.unizg.hr/DBM/dbm.php} for the DBM model.

1607.07750

1607

1607.06756_arXiv.txt

Two of the most widely observed and yet most puzzling features of the Sun's magnetic field are coronal loops that are smooth and laminar and prominences/filaments that are strongly sheared. These two features would seem to be quite unrelated in that the loops are near their minimum-energy current-free state, whereas filaments are regions of high magnetic stress and intense electric currents. We argue that, in fact, these two features are inextricably linked in that both are due to a single process: the injection of magnetic helicity into the corona by photospheric motions and the subsequent evolution of this helicity by coronal reconnection. In this paper, we present numerical simulations of the response of a \citet{Parker72} corona to photospheric driving motions that have varying degrees of helicity preference. We obtain four main

\label{sec:intro} A well-known feature of the solar magnetic field is the observation of filament channels at photospheric polarity inversion lines (PILs). These magnetic structures, situated in the upper chromosphere and lower corona, underlie and support the cool plasma that comprises prominences and filaments \citep{Martin98,Gaizauskas00}. Filament channels are regions of highly sheared magnetic field, containing large amounts of free energy that ultimately is converted into kinetic and thermal energy of the plasma, as well as nonthermal particle energies when filament channels erupt and drive coronal mass ejections (CMEs). The shear inherent in the filament channels is a form of magnetic helicity, and filament channels are known as dextral if they have negative helicity and sinistral if they have positive helicity. Observations indicate that dextral (sinistral) filament channels dominate in the northern (southern) hemisphere \citep[e.g.][]{Martin92, Rust94b, Zirker97, Pevtsov03}. This hemispheric helicity rule has also been observed in quiet-Sun magnetic fields \citep{Pevtsov01b}, sigmoids \citep{Rust96}, active-region magnetic fields \citep{Seehafer90}, coronal mass ejections (CMEs), and sunspot whorls \citep{Pevtsov14}. The strength of the preference ranges from about 55\% in active-region filaments \citep{Martin94}, to over 80\% in quiescent filaments \citep{Pevtsov03} and does not seem to change with solar cycle \citep{Hale27, Martin94,Hagino02}.\par A second, seemingly unrelated, feature of the solar magnetic field is the observation of loops in the closed-field corona that appear to be near their minimum energy, current-free state. These loops have been observed at high resolution in \emph{Transition Region and Coronal Explorer (TRACE)} XUV and X-ray images, such as the one in Figure \ref{fig:obs}, where they are seen to be very smooth and laminar, with little to no tangling \citep{Schrijver99}. In other words, there appears to be very little magnetic helicity associated with the topology of these loops. The picture of the corona that emerges, therefore, is one in which magnetic helicity manifests itself at specific locations, namely above PILs, while leaving the rest of the corona generally smooth and quasi-potential.\par Since the corona has very high Lundquist number, its magnetic helicity is believed to originate by injection from the photosphere, either by flux emergence or footpoint motions after emergence, and to be lost to the heliosphere by flux opening in CMEs and streamer blowouts. \emph{Solar Dynamics Observatory (SDO)} measurements of helicity injection into the coronal field indicate that shearing and twisting motions by the photosphere dominate the flux emergence \citep{Liu12}, meaning that the helicity budget of the corona is primarily due to the jostling of existing flux at the surface, rather than due to new flux emerging into the corona. Observations of photospheric convection show that these footpoint motions are highly complex \citep{Schmieder14}, with convective cells appearing randomly throughout the photosphere, and occur over a broad range of scales, of order minutes for granules and days for supergranules \citep{Hirzberger08}. From the standpoint of helicity injection, however, the important flows are those that twist the field. Compression of the field caused by converging flows is not expected to impart any helicity into the field, so these flows cannot be responsible for the shear observed in filament channels. The flows that twist up the field, in contrast, do inject a net helicity into the corona. Therefore, we conclude that it is sufficient to model the helicity injection into the corona with simple twisting motions as has been done by many authors \citep[e.g.][]{WilmotSmith10, Rappazzo13}. Such flows have routinely been observed in helioseismic measurements \citep{Duvall00,Gizon03,Komm07,Seligman14}. Vortical flows on the scale of granules \citep[e.g.][]{Bonet08, Bonet10, VD11, VD15} and supergranules \citep{Brandt88,Attie09} have also been observed.\par These considerations make it very challenging to understand the simultaneous presence in the corona of both filament channels and coronal loops. Magnetic stress is, apparently, injected throughout the solar photosphere, yet is almost nowhere to be found in the corona, except in filament channels. \citet{Antiochos13} presented a new model for the formation of filament channels, magnetic helicity condensation, based on the well-known inverse cascade of magnetic helicity in turbulent systems. In the helicity condensation model, photospheric convection imparts helicity into the coronal field, and this helicity is then transported throughout the corona by magnetic reconnection, which is well-known to conserve helicity \citep{Woltjer58,Taylor74,Taylor86,Berger84b}. Surface convection imparts the same sense of twist to adjacent flux tubes, which are then able to undergo component reconnection at their contact point. This component reconnection produces a single flux tube with an axial flux equal to the sum of the two original axial fluxes, but encircled by the same twist field present on each of the two original flux tubes. In this way the helicity, in the form of twist, inverse-cascades to larger and larger scales. The PIL forms a natural boundary of the flux system, so that when the twist reaches this boundary, it cannot proceed further, since all of the flux has already reconnected. The end result of this process is a mostly axial (untwisted) internal field, and a highly sheared (twisted) field at the PIL, precisely what is observed as a filament channel. At the same time, the untwisted internal field corresponds to the laminar coronal loops. In this way, the helicity condensation model provides a natural mechanism for the simultaneous formation of both highly sheared filament channels and relatively untwisted coronal loops. In this model, these two seemingly unrelated features of the solar atmosphere are actually created by the same process \citep{Antiochos13}.\par The helicity condensation model was initially simulated by \citet{Zhao15}, who injected magnetic helicity into a plane-parallel Parker corona \citep{Parker72}. These authors found that photospheric motions that inject the same helicity everywhere form filament channels at the PIL. Furthermore, randomizing the photospheric motions while keeping the same helicity injection rate did not qualitatively affect the accumulation of twist flux at the PILs. \citet{Zhao15} also tested the effect of injecting helicity of opposite signs on adjacent flux tubes, and found that their fields could not reconnect due to the twist components being co-aligned.\par In subsequent work \citep[][hereafter KAD15]{Knizhnik15}, we rigorously tested the helicity condensation model. We found that it not only qualitatively produced results consistent with the properties of filament channels, but that the inverse cascade of magnetic helicity due to reconnection produces a twist flux at the PIL that agrees quantitatively with the predictions of the helicity condensation model. Based on this result, we estimated that with the helicity preference observed on the Sun, filament channels will form in about a day or so, in agreement with observations of filament channel formation \citep{Martin98, Gaizauskas00}. We showed that helicity condensation agreed both qualitatively and quantitatively with observed properties of filament channels, and that the process produced relatively untwisted coronal loops everywhere except at the PIL. These results, however, were obtained for a $100\%$ helicity rule, meaning that all of the helicity injected into the corona was of the same sign. An obvious question to be raised is: what happens if a fraction of the injected helicity has the opposite sign? Indeed, this is a more realistic scenario since, as described above, the corona has a hemispheric helicity preference, rather than a rule, so that some helicity of the non-preferred sign is injected into the corona at all times.\par It is reasonable to expect that injecting helicity of the opposite sign into the corona would simply slow down the helicity condensation process. However, this result is not as straightforward as it may seem. The simulations of \citet{Zhao15} demonstrated that adjacent flux tubes have difficulty reconnecting if they are twisted in opposite senses. As a result, the twist was unable to inverse-cascade to larger scales, as is required to form the sheared filament channels and smooth coronal loops. Even if reconnection between the adjacent flux tubes is eventually achieved as a result of some instability -- such as ideal kinking -- driving the interaction, the twist-flux cancellation is expected to be far from perfect. Substantial residual twist could remain in what otherwise would have been smooth, untwisted coronal loops. Therefore, it is important to test whether sheared filament channels and smooth coronal loops form when both signs of magnetic helicity are injected into the corona.\par In this paper, we investigate the effect of varying helicity preference on the structure of the closed-field corona. We performed helicity-conserving numerical simulations that inject helicity into a plane-parallel Parker corona, as we did in KAD15. Extending our previous work, the fraction of helicity of each sign that is injected into the corona was varied. We report on three cases: 1) $100\%$ of the injected helicity is positive; 2) $75\%$/$25$\% of the injected helicity is positive/negative; and 3) $50$\%/$50$\% of the injected helicity is positive/negative. To make the simulations as realistic as possible, we also randomized the pattern of helicity injection. We compare the simulations with a fixed pattern of helicity injection to those with a randomized pattern of helicity injection.\par The paper is organized as follows. In \S \ref{sec:model} we discuss the setup and initialization of our numerical simulations. In \S \ref{sec:Hinjection} we describe how magnetic helicity is injected into the domain, and how the helicity preference is employed. In \S \ref{sec:Results} we discuss the results of our simulations, exploring the formation of filament channels and the smoothness of coronal loops for various helicity preferences, and compare the simulations with fixed and randomized patterns of helicity injection. We discuss the implications for understanding coronal magnetic structure in \S \ref{sec:implications}. \par

1607.06756

1607

1607.04279_arXiv.txt

We study how outflows of gas launched from a central galaxy undergoing repeated starbursts propagate through the circumgalactic medium (CGM), using the simulation code {\sc Ramses}. We assume that the outflow from the disk can be modelled as a rapidly moving bubble of hot gas at $\mathrm{\sim1\;kpc}$ above disk, then ask what happens as it moves out further into the halo around the galaxy on $\mathrm{\sim 100\;kpc}$ scales. { To do this we run 60 two-dimensional simulations scanning over parameters of the outflow. Each of these is repeated with and without radiative cooling, assuming a primordial gas composition to give a lower bound on the importance of cooling.} In a large fraction of { radiative-cooling} cases we are able to form rapidly outflowing cool gas from in situ cooling of the flow. We show that the amount of cool gas formed depends strongly on the `burstiness' of energy injection; sharper, stronger bursts typically lead to a larger fraction of cool gas forming in the outflow. The abundance ratio of ions in the CGM may therefore change in response to the detailed historical pattern of star formation. For instance, outflows generated by star formation with short, intense bursts contain up to 60 per cent of their gas mass at temperatures $<5 \times 10^4\,\mathrm{K}$; for near-continuous star formation the figure is $\lsim$ 5 per cent. Further study of cosmological simulations, { and of idealised simulations with e.g., metal-cooling, magnetic fields and/or thermal conduction}, will help to understand the precise signature of bursty outflows on observed ion abundances.

Star-forming galaxies are surrounded by gas (known as the circumgalactic medium, CGM) with a comparable total mass to their stellar mass \citep{2014ApJ...792....8W}. This gas is enriched by metals that were almost certainly ejected from the galaxy; outflows, carrying a mass comparable to the mass of star-forming regions, are ubiquitously observed in star forming galaxies in the local universe \citep{2004ApJ...606..829S,2015Natur.523..169E}, at intermediate redshifts $\lowercase { {z \sim 0.5}}$ \citep{Rubin2010, Nielsen2013, Bordoloi2014} and at high redshifts of $\lowercase{z \sim 6}$ \citep{2002ApJ...576L..25A,2005ApJ...621..227M}. The interrelationship between inflow and outflow is critical to the behaviour of galaxies as a whole, as it reshapes quantities such as the stellar mass function \citep{1986ApJ...303...39D,2010MNRAS.406.2325O} and mass-metallicity relation \citep{2008MNRAS.385.2181F}. Baryon cycling through the CGM likely also plays a role in controlling the distribution of dark matter \citep{pontzen,2015MNRAS.451.1366P}. The mechanism behind outflows is uncertain and may relate to some combination of supernova feedback \citep{2010Natur.463..203G}, winds of high-mass stars \citep{1999ApJ...513..156M, 2008MNRAS.387.1431D} and active galactic nuclei (AGN) \citep{2008A&A...491..407N}. Dwarf galaxies are of particular interest since their small black hole masses makes AGN feedback ineffective, so that the outflows are almost certainly linked directly to star formation feedback. Additionally the higher gas-to-stellar-mass fraction combined with the shallow gravitational potential allow outflows to easily release material to the CGM \citep{2014ApJ...786...54P}. Observations in this low-mass regime have shown evidence of C {\sevensize IV} absorption in the CGM out to 100 kpc. Therefore, these galaxies are useful case studies of the connection between the CGM and the host galaxy. Galactic outflows seem to possess a multi-phase nature, spanning several orders of magnitude in temperature \citep{2014ApJ...792....8W}. There is a particular puzzle over how cold gas material could survive in galactic outflows if they are entrained in a hot flow \citep[e.g.,][]{2005ARA&A..43..769V,2015ApJ...805..158S,2015arXiv150701951Z}. Perhaps this implies that cold gas can be directly accelerated using radiation pressure \citep{2005ApJ...618..569M,2012MNRAS.421.3522H}. { The lifetime of cold clouds is, however, dependent on physical assumptions and simulation methods meaning that robust conclusions are difficult to draw. } Another scenario is that the cooler phases of the outflow are actually formed in situ by radiative cooling \cite[e.g.,][]{2000MNRAS.317..697E,2015ApJ...803....6M}; recently, detailed analytic discussions in support of this possibility have been given by \cite{2016MNRAS.455.1830T} and \cite{2015arXiv150907130B}. In this paper, we study the formation of cool gas in outflows in a different way. Instead of developing analytic solutions to the outflow problem we inject hot gas into a idealised galaxy halo with a fixed potential. The numerical setup is similar to that of \cite{1999ApJ...513..142M} but introduces a simple prescription to deliver the outflow in discrete bursts. While our setup is highly simplified, we use it to argue that the balance between hot and cooler gas will be affected by the duty cycle of star formation. Such a link introduces a new dimension to the relationship between outflows and the evolution of a galaxy, raising the possibility that the relative abundance of different observed ions reflects information about the detailed star formation history. Since it is specifically {\it bursty} star formation that can have a profound impact on dark matter and stellar dynamics \citep{pontzen}, cross-checking typical star formation patterns in the observed universe would be a valuable additional benefit to studies of the CGM. { The simulations are two-dimensional, allowing us to run a much larger parameter study than would be possible in a three-dimensional study. } This paper is organised as follows. In Sec.~\ref{simulations} we explain the initial and boundary conditions to simulate outflows in galaxies using {\sc Ramses} \citep{Teyssier:2002aa}. We discuss the results in Sec.~\ref{results}. Finally, we summarise in Sec.~\ref{conclusions}.

\label{conclusions} In this paper we have considered the possibility that the cool gas material observed in outflows in the circumgalactic medium is a consequence of in situ cooling \citep{2016MNRAS.455.1830T}, using the {\sc Ramses} code. We do not simulate the disk in our galaxies; instead we manually inject gas moving into the base of a box representing the halo \citep[see also][]{1999ApJ...513..142M}. We started by finding and testing equilibrium inflows to be sure that the effects we observe are a result of outflows rather than the choice of initial conditions; see Eqs. \eqref{eqn:u} -- \eqref{eqn:Temp}. We used a fixed potential corresponding to a $50\;\km\,\s^{-1}$ virial velocity dwarf galaxy. The cooling function $\dot Q_{\mathrm{cool}}$ we adopt is suitable for primordial gas as implemented by {\sc Ramses}, which underestimates the true cooling rates and so leads us to conservative conclusions. We modified {\sc Ramses} to inject a time-varying flow into the bottom of the computational domain, according to a set of parameters summarised in Table \ref{on_off}. Our particular focus was on the role that varying star formation rates in the galaxy could have on the evolution of outflows as they traverse the halo. From the complete set of parameters characterising outflows we therefore varied two: the overall star formation cycle length, $\tcyc$, and the fraction of that time spent pumping gas into the CGM, $\fcyc$. We varied these while keeping the total energy injection and mass loading constant. We found a close connection between the parameters and the overall nature of the outflows' traversal of the CGM (Fig.~\ref{pattern}). This in turn has a strong effect on the multiphase nature of the outflows. The amount of cooler $T<5 \times 10^4\,\K$ gas present in the outflow varies strongly over the course of a cycle (Fig.~\ref{coldgas}). Cool material is typically able to form as a shock propagates outwards provided that no hot material is being injected behind the shock. This leads to the time-averaged cool mass fraction depending on both $\tcyc$ and $\fcyc$ (Fig.~\ref{dutyvscycledots}). There are two regimes in which we obtain large cool mass fractions: the first has a small $\fcyc$, corresponding to a rapidly fluctuating star formation rate. Provided $\tcyc$ is greater than a few hundred $\Myr$, the successive shocks do not join up into a coherent flow and the strong time variability triggers waves of effective cooling that travel through the CGM. The second approach is to leave a long period $\gsim 1\;\Gyr$ between successive star formation epochs; in this case cool gas is able to form in the turbulent halo left behind when ejection from the disk shuts off. Our results suggest that steady flow or single-burst solutions with cooling \citep{2016MNRAS.455.1830T} form a lower bound on the amount of in situ cooling to be expected in realistic galaxies with time-varying feedback. Our interest in cool gas is motivated by observational results that show the presence of a cool phase even at large distances from galactic centres \citep[e.g.,][]{Nielsen2013}, which is hard to explain in entrainment scenarios \citep{2015arXiv150701951Z}. However it would be premature to compare our highly idealised study directly to observations. { To enable our large parameter study, we had to restrict ourselves to two-dimensional solutions; the detailed behaviour in three dimensions could differ significantly. Other neglected aspects of the problem include the enhanced cooling rates from metal enrichment and the effects of thermal conduction and magnetic fields. Furthermore a realistic cosmological environment is far more complex than the uniform inflow solution that forms our initial conditions. In terms of the cooling rates, neglecting metals leads to a conservative bound; i.e., more realistic simulations may be able to form cold clouds more easily than our work suggests.} We hope to use our results to interpret the CGM around ab initio cosmological galaxy formation simulations. Some feedback algorithms enforce relatively steady-state star formation \cite[e.g.,][]{2008MNRAS.387.1431D,2008MNRAS.387..577O,2014MNRAS.444.1518V,2015MNRAS.446..521S} whereas others lead to strong bursts \cite[e.g.,][]{2014MNRAS.445..581H,2015MNRAS.453.3499K} and the importance of this distinction for galactic dynamics has already been established \citep{pontzen}. In future work we will study what role in situ cooling plays in these different scenarios, and make the link to observational constraints on the rich phenomenology of the CGM.

1607.04279

1607

1607.03459_arXiv.txt

Locating sources on the sky is one of the largest challenges in gravitational wave astronomy, owing to the omni-directional nature of gravitational wave detection techniques, and the often intrinsically weak signals being observed. Ground-based detectors can address the pointing problem by observing with a network of detectors, effectively triangulating signal locations by observing the arrival times across the network. Space-based detectors will observe long-lived sources that persist while the detector moves relative to their location on the sky, using Doppler shifts of the signal to locate the sky position. While these methods improve the pointing capability of a detector or network, the angular resolution is still coarse compared to the standards one expects from electromagnetic astronomy. Another technique that can be used for sky localization is null-stream pointing. In the case where multiple independent data streams exist, a single astrophysical source of gravitational waves will appear in each of the data streams. Taking the signals from multiple detectors in linear combination with each other, one finds there is a two parameter family of coefficients that effectively null the gravitational wave signal; those two parameters are the angles that define the sky location of the source. This technique has been demonstrated for a network of ground-based interferometric observatories, and for 6-link space interferometers. This paper derives and extends the null-stream pointing method to the unique case of pulsar timing residuals. The basic method is derived and demonstrated, and the necessity of using the method with multiple sub-arrays of pulsars in the pulsar timing array network is considered.

The gravitational wave spectrum covers many decades in frequency space, just like the electromagnetic spectrum. The particular waves that are radiated in any given band of the spectrum reflect the astrophysical phenomena that generated the waves, and in particular the astrophysical timescales that dominate the movement of mass in the system. In the very-low frequency band of the spectrum, $10^{-9} Hz \lesssim f_{gw} \lesssim 10^{-6} Hz$, the primary detection technique is known as pulsar timing. Pulsar timing uses the stable rotation of distant pulsars as clocks. It was first described by Detweiler \cite{Detweiler1979}, and proceeds as follows. The arrival time of pulses from a pulsar are monitored on Earth, and compared against a model for the expected arrival times (the ``ephemeris''). Using simple models for pulsar spindown over time, models for pulse arrival times can be built with precisions at the level of fractions of a microsecond when pulsar monitoring spans several years. The fundamental signature of a gravitational wave passing between the Earth and the distant pulsar is a change in the time of flight for individual pulses, advancing or retarding their time of arrival compared to the model ephemeris. The difference between the arrival time of the pulses and the model are known as ``timing residuals,'' and they constitute the basic data stream in pulsar timing searches for gravitational waves. While the measurement can be made with long-term observations of a single pulsar, the implementation of this method as a viable detection technique has been realized through the development of \textit{pulsar timing arrays}, where many pulsars in many parts of the sky are monitored over long periods of time, combining timing residuals from multiple pulsars to search for gravitational waves. Many efforts have been launched on this front, including the North American Nanohertz Observatory for Gravitational Waves (NANOGrav)\cite{nanograv}, the Parkes Pulsar Timing Array \cite{PPTA}, the European Pulsar Timing Array \cite{EPTA}, and the International Pulsar Timing Array \cite{IPTA}. The expected astrophysical sources of very-low frequency gravitational waves include supermassive black hole binaries (SMBHB), stochastic backgrounds of SMBHBs, as well as a variety of other stochastic sources such as phase transitions in the early Universe or relic radiation from the Big Bang. See \cite{Burke-Spolaor:2015xpf} for a recent review of PTA sources. % Like most gravitational wave detection techniques, pulsar timing arrays are omnidirectional, in the sense that they are sensitive to gravitational waves from any location on the sky. As a general rule of thumb, the sensitivity of any particular pulsar to incident gravitational waves is a function of the angle between the line of sight to the pulsar and the line of sight to the gravitational wave source. This relationship is distinctly evident in the Hellings and Downs curve \cite{Hellings:1983fr}, which relates the correlation of residual signals from two pulsars to their angular separation in the sky. This work develops a data combination technique known as ``pulsar null streams'' to provide a good estimate of the location of a gravitational wave source on the sky. This is an explicit analysis technique for determining the sky location of a putative gravitational wave source without engaging in a full parameter search. Source location knowledge absent other parameter information is useful for counterpart searches, as well as restricting the search space of computationally intensive signal searches. Null stream mapping of gravitational wave sources has been described for interferometric detectors \cite{TintoGursel,ZSS1,ZSS2,NullNetwork}, and relies on the fact that there are correlated gravitational wave signals between detectors in a network. For the case of a pulsar timing array, one has the same situation --- a gravitational wavefront will produce a correlated response in the timing of every pulsar in the array. This correlation may be exploited to create a null stream by taking advantage of the geometrical properties of a pulsar's response to incident gravitational waves. The paper is organized as follows. In section \ref{sec.pulsarTiming} we review a basic signal model for pulsar timing residuals, and express it in a form conducive to build a pulsar null stream. In section \ref{sec.pns} the pulsar null stream is described and written out. Section \ref{sec.demonstration} shows how the pulsar null stream works for an array of three pulsars. Section \ref{sec.PulsarSep} discusses the errors inherent in one sub-array of pulsars. Section \ref{sec.MultSubArray} demonstrates how multiple sub-arrays of pulsars strengthen the null signal technique. Section \ref{sec.noise} examines the efficacy and overall pointing ability of the method in the presence of noise. Section \ref{sec.discussion} summarizes the key results and discusses future directions for this work.

\label{sec.discussion} Null stream mapping of gravitational wave sources relies on the fact that there are correlated gravitational wave signals between detectors. The underlying premise of the null stream construction is that for a collection of pulsars observing the same source, the gravitational wave signal is common to all pulsars in the array, but modified by geometric factors related to the relative position of the source on the sky. We have shown how a linear combination of three pulsar timing streams gives a signal that is minimized at the correct sky location of the gravitational wave source, though not as localized as one might need for electromagnetic counterpart searches. Further we have characterized the error and localization ability of one null signal. Though there are significant errors with one null signal when multiple signals are combined as products the errors decrease and the localization increases dramatically. The techniques here have focused on analysis in the frequency domain specialized for looking at stable sinusoidal signals expected from super massive black hole mergers, however, future work will consider how these sky localization techniques will work for burst type sources in the time domain.

1607.03459

1607

1607.03203_arXiv.txt

The effect of quantum gravity can bring a tiny light speed variation which is detectable through energetic photons propagating from gamma ray bursts (GRBs) to an observer such as the space observatory. Through an analysis of the energetic photon data of the GRBs observed by the Fermi Gamma-ray Space Telescope (FGST), we reveal a surprising regularity of the observed time lags between photons of different energies with respect to the Lorentz violation factor due to the light speed energy dependence. Such regularity suggests a linear form correction of the light speed $v(E)=c(1-E/E_{\rm LV})$, where $E$ is the photon energy and $E_{\rm LV}=(3.60 \pm 0.26) \times 10^{17}~ \rm GeV$ is the Lorentz violation scale measured by the energetic photon data of GRBs. The results support an energy dependence of the light speed in cosmological space.

1607.03203

1607

1607.00161.txt

\label{0. abstract} Recent observations suggest that some type Ia supernovae (SNe Ia) originate from the merging of two carbon-oxygen white dwarfs (CO WDs). Meanwhile, recent hydrodynamical simulations have indicated that the accretion-induced collapse may be avoided under certain conditions when double WDs merge violently. However, the properties of SNe Ia from this violent merger scenario are highly dependent on a particular mass-accretion stage, the so-called WD + He subgiant channel, during which the primary WD is able to increase its mass by accreting He-rich material from a He subgiant before the systems evolves into a double WD system. In this article, we aim to study this particular evolutionary stage systematically and give the properties of violent WD mergers. By employing the Eggleton stellar evolution code, we followed a large number of binary calculations and obtained the regions in parameter space for producing violent mergers based on the WD + He subgiant channel. According to these simulations, we found that the primary WDs can increase their mass by $\sim0.10-0.45\,\rm M_{\odot}$ during the mass-accretion stage. We then conducted a series of binary population synthesis calculations and found that the Galactic SN Ia birthrate from this channel is about $0.01-0.4\times10^{\rm -3}\rm yr^{\rm -1}$. This suggests that the violent WD mergers from this channel may only contribute to $\sim0.3\%-10\%$ of all SNe Ia in our Galaxy. The delay times of violent WD mergers from this channel are $\ge 1.7\rm Gyr$, contributing to the SNe Ia in old populations. We also found that the WD + He subgiant channel is the dominent way for producing violent WD mergers that may be able to eventually explode as SNe Ia.

\label{1. Introduction} Type Ia supernovae (SNe Ia), which are defined as SN explosions without H and He lines but with strong SiII absorption lines in their spectra, have been successfully used as standard distance indicators in cosmological studies of dark energy (e.g. Riess et al. 1998; Perlmutter et al. 1999). There is a theoretical consensus that SNe Ia result from thermonuclear explosions of carbon-oxygen white dwarfs (CO WDs) in close binaries (Hoyle \& Fowler 1960). According to the nature of the mass donor, two widely accepted models have been proposed, i.e. the single-degenerate (SD) model (e.g. Whelan \&Iben 1973; Hachisu, Kato \& Nomoto 1996; Li \& van den Heuvel 1997; Langer et al. 2000; Han \& Podsiadlowski 2004; Meng et al. 2009; Wang, Li \& Han 2010) and the double-degenerate (DD) model (e.g. Webbink 1984; Iben \& Tutukov 1984; Nelemans et al. 2001; Toonen, Nelemans \& Portegies 2012). The difference between these two models is whether the mass donor is a non-degenerate star (a main-sequence star, a red-giant star or a helium star; the SD model) or another WD merging with the primary WD (the DD model). However, there is still no conclusive evidence to support any progenitor models of SNe Ia. Recent studies suggested that more than one progenitor models may be required to reproduce the observational diversities of SNe Ia (see Podsiadlowski et al. 2008; Howell et al. 2011; Wang \& Han 2012; Maoz, Mannucci \& Nelemans 2014). Some recent observations seem to favor the DD model: for example, the lack of H and He lines in the nebular spectra of most SNe Ia (e.g. Leonard 2007; Ganeshalingam, Li \& Filippenko 2011), no conclusive proof for the existence of surviving companions (e.g. Badenes et al. 2007; Graham et al. 2015), the lack of radio emission (e.g. Hancock, Gaensler \& Murphy 2011; Horesh et al. 2012) and the observed evidence for the existence of some super-luminous events (e.g. Howell et al. 2006; Hichen et al. 2007; Scalzo et al. 2010).\footnote{Some of these observed clues may also be explained by the SD model after considering the spin-up/spin-down processes of WDs (e.g. Justham 2011; Di Stefano, Voss \& Claeys 2011; Hachisu, Kato \& Nomoto 2012; Wang et al. 2014).} Meanwhile, the delay times of SNe Ia, defined as the time interval from the star formation to the thermonuclear explosion, provide an important observational constraint for SN Ia progenitor models. Recent observations have suggested that the delay time distributions (DTDs) of SNe Ia follow a power-law distribution $\sim t^{-1}$ (Totani et al. 2008; Maoz et al. 2011; Graur et al. 2011; Barbary et al. 2012; Sand et al. 2012), which is consistent with the results predicted by the DD model (e.g. Ruiter, Belczynski \& Fryer 2009; Mennekens et al. 2010). In addition, Dan et al. (2015) compared the nucleosynthetic yields of thermonuclear explosions from WD mergers with the observations and found that some of their models are good condidates for type Ia events. Moreover, many double WDs have been proposed as possible progenitor candidates of SNe Ia, e.g. Henize 2$-$428, KPD 1930 + 2752, WD 2020-425, V458 Vulpeculae, SBS 1150 + 599A and GD687, etc (Maxted, Marsh \& North 2000; Geier et al. 2007, 2010; Napiwotzki et al. 2007; Rodr\'iguez-Gil et al. 2010; Tovmassian et al. 2010; Santander-Garc\'ia et al. 2015). Especially, Henize 2-428 is a planetary nebula with double degenerate core that has a total mass of $\sim 1.76\,\rm M_{\odot}$, mass ratio $q\sim1$ and orbital period $\sim4.2\,\rm h$, which suggests that Henize 2-428 is a good candidate for super-chandrasekhar mass SNe Ia (e.g. Santander-Garc\'ia et al. 2015). %, the lack of X-ray emissions in elliptical galaxies (e.g. Gilfanov \& Bogd\'an 2010) However, previous studies reveal that the outcome of double WD mergers may be a neutron star resulting from an accretion-induced collapse, rather than a thermonuclear explosion (e.g. Nomoto \& Iben 1985; Saio \& Nomoto 1985; Timmes, Woosley \& Taam 1994). These arguments are based on the assumption that the merging remnant consists of a hot envelope or a thick disc, or even both upon the primary WD (e.g. Kashyap et al. 2015). In that case, the accretion rate from the envelope or disc may be relatively high, leading to the formation of an oxygen-neon (ONe) WD that would collapse into a neutron star when it approaches the Chandrasekhar limit (Nomoto \& Iben 1985; Saio \& Nomoto 1998). We note that Yoon, Podsiadlowski \& Rosswog (2007) argued that the accretion-induced collapse can be avoided for a certain range of parameters and found that they would explode as a SN Ia some $10^5\,\rm yr$ after the initial dynamical merger when considering the rotation of the WDs. Recently, Pakmor et al. (2010) proposed a new explosion scenario for SNe Ia that are produced by the merging of double WDs referred to as the violent merger scenario. In this scenario, a prompt detonation is triggered while the merger is still ongoing, giving rise to a SN Ia explosion (see also Pakmor 2011, 2012). Pakmor et al. (2010) found that the violent merger of double WDs with almost equal masses of $0.9\,\rm M_{\odot}$ can provide an explanation for the formation of sub-luminous 1991bg-like events. Pakmor et al. (2011) suggested that the minimum critical mass ratio for double WD mergers to produce SNe Ia is about 0.8. R\"{o}pke et al. (2012) recently argued that the violent merger model can reproduce the observational properties of SN 2011fe. Additionally, this scenario may also explain the formation of some super-Chandrasekhar mass SNe Ia (e.g. Moll et al. 2014; Cody et al. 2014). For a series of recent theoretical and observational studies of the violent merger scenario see Taubenberger et al. (2013), Kromer et al. (2013), Fesen, H\"oflich \& Hamilton (2015), Seitenzahl et al. (2015), Tanikawa et al. (2015), Chakraborti, Childs \& Soderberg (2015) and Bulla et al. (2016). %Following a three-dimensional simulation for the violent merger of double WDs with masses of 0.9 and $1.1\,\rm M_{\odot}$, Pakmor et al. (2012) suggested that the violent merger scenario may also contribute to normal SNe Ia. Ruiter et al. (2013) recently investigated the distribution of the SN Ia brightness based on the violent merger scenario and argued that the theoretical peak-magnitude distribution from their calculations can roughly reproduce their observed properties. The distribution they obtained depends critically on a particular binary evolutionary stage called the WD + He subgiant channel, during which the primary WD is able to grow in mass by accreting He-rich material from their He companion before eventually evolving into a double WD system. However, this mass-accretion process is still poorly studied, which may influence the binary population synthesis (BPS) results of SNe Ia based on the violent merger scenario (see Ruiter et al. 2013). %Ruiter et al. (2013) found that this mass-accretion process is predicted to occur for about 43\% of all the binaries which can eventually produce SNe Ia based on the violent WD merger scanario. However, this mass-accretion process is poorly studied in their work. Moreover, they employed a relatively weak criteria for the mass ratio of the WD binaries to give the BPS results (see the formula\,1 of Ruiter et al. 2013). In this paper, we systematically study the WD + He subgiant channel for producing SNe Ia via the violent WD merger scenario and obtain the parameter space for the progenitors of SNe Ia. We then present a series of BPS simulations using this parameter space. In Sect.\,2, we describe our methods for the binary evolution calculations and give the main results. The methods and results of our BPS simulations are presented in Sect.\,3. A detailed discussion is provided in Sect.\,4 and finally a summary is given in Sect.\,5.

\label{Discussion} %The particular mass-accretion process from the slightly evolved He star onto the surface of the primary WD is crucial for producing SNe Ia based on the violent merger scenario. For the violent merger scenario, the mass-accretion process during which the primary WD accretes material from a He subgiant is crucial for the production of SNe Ia. By considering an identical value of $\alpha_{\rm CE}$ as Ruiter et al.\ (2013), i.e. $\alpha_{\rm CE}\lambda=0.5$ in this article, we found that such a particular mass-accretion stage can be expected to occur in about 43.7\% of all the binary systems which would eventually contribute to SNe Ia; this is consistent with the fraction given by Ruiter et al.\ (2013). For other cases with larger $\alpha_{\rm CE}\lambda$, this fraction would be much higher in our simulations: 91.8\% for the case of $\alpha_{\rm CE}\lambda=1.0$ and 85.7\% for the case of $\alpha_{\rm CE}\lambda=1.5$. Thus, we suggest that the WD + He subgiant channel may be the dominant way for producing violent mergers of double WDs that can form SNe Ia. During the mass-accretion stage, the primary WDs can increase their mass by about $0.1-0.45\,\rm M_{\odot}$ from the He donors based on our calculations, which is slightly wider than the mass range ($\sim0.15-0.35\,\rm M_{\odot}$) presented by Ruiter et al.\ (2013). This is due to the fact that we employed full stellar evolution calculations, while Ruiter et al.\ (2013) adopted a simple analytical fitting formula for estimating the mass-accretion rate. We compared our prescriptions with Ruiter et al.\ (2013) by calculating the evolution of a WD + He star system provided in Sect.\,2.3 of Ruiter et al.\ (2013). This system has a $0.84\,\rm M_{\odot}$ WD and a $1.25\,\rm M_{\odot}$ He star separated by $1.73\,\rm R_{\odot}$. Our calculation shows that $M_{\rm WD}^{\rm f}=1.2\,\rm M_{\odot}$ and $M_{\rm 2}^{\rm f}=0.84\,\rm M_{\odot}$, compared to $M_{\rm WD}^{\rm f}=1.19\,\rm M_{\odot}$ and $M_{\rm 2}^{\rm f}=0.77\,\rm M_{\odot}$ in Ruiter et al.\ (2013). When this double WDs is formed, their separation is $1.74\,\rm R_{\odot}$ for our calculation and $1.92\,\rm R_{\odot}$ for Ruiter et al.\ (2013). %Similar to the evolution described in Sect.\,2.3, this WD + He star system will eventually evolve to a double WD system. For example, for a specific WD + He star system with ($M_{\rm WD}^{\rm i}$, $M_{\rm 2}^{\rm i}$, $\log\,P^{\rm i})=(0.84, 1.27, -0.74)$, our detailed simulation shows that the final double WDs is ($M_{\rm WD}^{\rm f}$, $M_{\rm 2}^{\rm f}$, $\log\,P^{\rm f})=(1.20, 0.84, -0.73)$, while the fitting method from Ruiter et al. (2013) gives ($M_{\rm WD}^{\rm f}$, $M_{\rm 2}^{\rm f}$, $\log\,P^{\rm f})=(1.19, 0.77, -0.66)$ (see Fig.\,2 of Ruiter et al. 2013). In addition, Ruiter et al.\ (2013) claimed that the DTDs from violent mergers agree rather well with observations and that the violent merger scenario may be the dominant contributor to SNe Ia in field galaxies. They argued that violent mergers from the WD + He subgiant channel may contribute to SNe Ia with delay times $>150\,\rm Myr$, a conclusion that is quite different from our results. This is caused by the different criteria for the critical mass ratio for double WDs between Ruiter et al. (2013) and this work. We note that Ruiter et al. (2013) used a relatively weak criterion with $\eta=1.5$ for the critical mass ratio ($q_{\rm cr}$) of the WD binaries to calculate the SN Ia birthrate and delay time distribution (see formula\,1 in their paper). From Fig.\,5 of Ruiter et al.\ (2013), one can see that the difference between their $q$-cut (with $\eta=1.5$) and no $q$-cut at all is very minor, especially for the massive WDs of double WDs at the high-mass end, which means that the double WDs with nearly all mass ratios are thought to be able to produce SNe Ia in the Ruiter et al.\ (2013) estimates. %Ruiter et al. (2013) used a relatively weak criteria with $\eta=1.5$ for the mass ratio of the WD binaries q (see the Formula\,1 of Ruiter et al. 2013) to calculate the SN Ia birthrate and delay time distribution. From the Fig.\,5 of that paper, we can see that the criteria with q-cut $\eta=1.5$ and with no q-cut are almost the same, which means that the double WDs with almost all mass ratios are thought to be able to produce SN Ia explosions. Thus, the theoretical SN Ia birthrate they gave based on violent merger scenario, which is comparable with the observations, may be too high. The brightness of SN Ia explosions originating from violent mergers are mainly determined by the final mass of the massive WDs ($M_{\rm WD1}$). The reason is that the less massive WD is totally destroyed in the merging process and that almost all of the iron-group elements (including $^{\rm56}\rm Ni$) are produced in the thermonuclear explosion inside the massive WD. For the violent mergers with low-mass WDs ($M_{\rm WD1}<1.1\,\rm M_{\odot}$), their low density causes incomplete silicon burning during the explosion, resulting in a lower yield of $^{\rm56}\rm Ni$ compared with normal SNe Ia, corresponding to sub-luminous SNe Ia (e.g. Pakmor et al.\ 2010, 2011). However, for the violent mergers with massive WDs ($M_{\rm WD1}>1.1\,\rm M_{\odot}$), the yields of these massive mergers are still quite uncertain, and they are potential candidates for explaining super-Chandrasekhar mass SNe Ia (Moll et al. 2014; Cody et al. 2014). %We can make a plausible hypothesis that the former mergers produce sub-luminous explosions and the latter explode as normal SNe Ia. Our calculations show that the proportion of violent mergers which can produce sub-luminous SNe Ia is about 46\%, and normal SNe Ia account for 54\% (see Fig\,8). %However, it is still not fully confirmed about the criteria for the occurrence of the prompt detonations to produce SNe Ia from the violent merger scenario, e.g. the minimum mass of the primary WD and the minimum mass ratio for the WD binaries, etc. In this work, the minimum mass of the primary WD is set to be $0.8\,\rm M_{\odot}$, which is comparable to this value in Pakmor et al. (2010) and Ruiter et al. (2013). For the mass ratio of the double WDs, $q\ge0.8$ is adopted here, which is consistent with Pakmor et al. (2011). However, the minimum critical mass ratio of double WDs to produce prompt detonations and SNe Ia through the violent merger scenario is still quite uncertain (e.g. Pakmor et al.\ 2011; Sato et al.\ 2016). In this work, we adopted $q_{\rm cr}\ge0.8$, which is consistent with Pakmor et al.\ (2011), but somewhat optimistic. If a larger critical mass ratio (e.g. $q_{\rm cr}\ge0.9$) for the double WDs is adopted, the contribution of violent mergers to the Galactic birthrate of SNe Ia would be $\sim$$0.1$$-$$1.1\times10^{\rm -4}\,\rm yr^{\rm -1}$, accounting for only 0.2\% to 3\% of all SNe Ia in the Galaxy, and the delay times of SNe Ia would be larger than 3.5 Gyr. Furthermore, Sato et al.\ (2015) argued that a SN Ia would be produced during the merging process if the masses of both WDs are in the range of $0.9\,\rm M_{\odot} \le M_{\rm WD} \le 1.1\,\rm M_{\odot}$, whereas the thermonuclear explosion can be expected in the stationary rotating merger remnant stage if the massive WDs in double WD systems are in the mass range of $0.7\,\rm M_{\odot} \le M_{\rm WD} \le 0.9\,\rm M_{\odot}$ and the total masses are larger than the Chandrasekhar mass (see also Moll et al.\ 2014; Cody et al.\ 2014). Chen et al.\ (2012) also claimed that the SN Ia birthrate from the DD model may decrease significantly when considering different constraints for double WDs. We note that recent studies suggested that a thin He shell on the surface of CO WDs could potentially produce SNe Ia more easily in WD mergers (e.g. Dan et al.\ 2012; Pakmor et al.\ 2013). Furthermore, the merging of CO WD + He WD systems may also contribute to the birthrates of SNe Ia based on the double-detonation scenario (see Dan et al.\ 2012; Pakmor et al.\ 2013). In this scenario, mass transfer is relatively unstable and occurs within the direct impact regime, in which Kelvin-Helmholz instabilities may trigger an explosion of the He shell on the surface of the CO WD (see Guillochon et al.\ 2010). The shock compression in the CO core caused by the He explosion could potentially lead to the explosion of the whole WD. Less $^{\rm 56}$Ni would be produced during the thermonuclear explosions as the accretors always have sub-Chandrasekhar masses; the SNe Ia resulting from such mergers may have lower luminosities than normal SN Ia explosions, i.e they may contribute to the class of sub-luminous SNe Ia (e.g. Sim et al. 2010; Dan et al. 2012). Considering the possibility of CO WD + He WD systems for producing sub-luminous SNe Ia, future BPS studies are needed to explore the properties of SNe Ia through the CO WD + He WD scenario. %Indeed, Dan et al. (2012) also showed that a thermonuclear runaway can be expected for the merging of a CO WD with a He-rich WD under certain conditions.

1607.00161

1607

1607.06262.txt

We make a comparison for ten typical, popular dark energy models according to their capabilities of fitting the current observational data. The observational data we use in this work include the JLA sample of type Ia supernovae observation, the Planck 2015 distance priors of cosmic microwave background observation, the baryon acoustic oscillations measurements, and the direct measurement of the Hubble constant. Since the models have different numbers of parameters, in order to make a fair comparison, we employ the Akaike and Bayesian information criteria to assess the worth of the models. The analysis results show that, according to the capability of explaining observations, the cosmological constant model is still the best one among all the dark energy models. The generalized Chaplygin gas model, the constant $w$ model, and the $\alpha$ dark energy model are worse than the cosmological constant model, but still are good models compared to others. The holographic dark energy model, the new generalized Chaplygin gas model, and the Chevalliear-Polarski-Linder model can still fit the current observations well, but from an economically feasible perspective, they are not so good. The new agegraphic dark energy model, the Dvali-Gabadadze-Porrati model, and the Ricci dark energy model are excluded by the current observations.

\label{sec:intro} The current astronomical observations have indicated that the universe is undergoing an accelerated expansion~\cite{Riess:1998cb,Perlmutter:1998np,Tegmark:2003ud,Eisenstein:2005su,Spergel:2003cb}, for which a natural explanation is that the universe is currently dominated by dark energy (DE) that has negative pressure. The study of the nature of dark energy has become one of the most important issues in the field of fundamental physics~\cite{Sahni:1999gb,Padmanabhan:2002ji,Bean:2005ru,Copeland:2006wr,Sahni:2006rde,Kamionkowski:2007wv,Frieman:2008sn,Li:2011sd,Bamba:2012cp}. But, hitherto, we still know little about the physical nature of dark energy. The simplest candidate for dark energy is the Einstein's cosmological constant, $\Lambda$, which is physically equivalent to the quantum vacuum energy. For $\Lambda$, one has the equation of state $p_\Lambda=-\rho_\Lambda$. The cosmological model with $\Lambda$ and cold dark matter (CDM) is usually called the $\Lambda$CDM model, which can explain the current various astronomical observations quite well. But the cosmological constant has always been facing the severe theoretical challenges, such as the fine-tuning and coincidence problems. There also exist many other possible theoretical candidates for dark energy. For example, a spatially homogeneous, slowly rolling scalar field can also provide a negative pressure, driving the cosmic acceleration. Such a light scalar field is usually called ``quintessence'' \cite{Zlatev:1998tr,Steinhardt:1999nw,Zhang:2005rg,Zhang:2005rj}, which provides a possible mechanism for dynamical dark energy. More generally, one can phenomenologically characterize the property of dynamical dark energy through parametrizing $w$ of its equation of state (EoS) $p_{\rm de}=w\rho_{\rm de}$, where $w$ is usually called the EoS parameter of dark energy. For example, the simplest parametrization model corresponds to the case of $w={\rm constant}$, and this cosmological model is sometimes called the $w$CDM model. A more physical and realistic situation is that $w$ is time variable, which is often probed by the so-called Chevalliear-Polarski-Linder (CPL) parametrization \cite{MD:2001,EV:2003}, $w(a)=w_0+w_a(1-a)$. For other popular parametrizations, see, e.g., \cite{Huterer:2000mj,Wetterich:2004pv,Jassal:2004ej,Upadhye:2004hh,Xia:2004rw,Linder:2006sv,Lazkoz:2010gz,Ma:2011nc,Li:2011dr,Li:2012vn,Li:2012via}. Some dynamical dark energy models are built based on deep theoretical considerations. For example, the holographic dark energy (HDE) model has a quantum gravity origin, which is constructed by considering the holographic principle of quantum gravity theory in a quantum effective field theory~\cite{A.G:1999wv,M.Li:2004wv}. The HDE model can naturally explain the fine-tuning and coincidence problems \cite{M.Li:2004wv} and can also fit the observational data well \cite{hde1,hde2,hde3,hde4,hde5,hde6,hde7,Li:2009zs,Li:2009bn,Wang:2012uf,Wang:2013zca,Zhang:2015rha,Feng:2016djj,He:2016rvp}. Its theoretical variants, the new agegraphic dark energy (NADE) model \cite{HWRG:2008} and the Ricci dark energy (RDE) model \cite{CGFQ:2009}, have also attracted lots of attention. In addition, the Chaplygin gas model \cite{Kamenshchik:2001cp} is motivated by braneworld scenario, which is claimed to be a scheme for unifying dark energy and dark matter. To fit the observational data in a better way, its theoretical variants, the generalized Chaplygin gas (GCG) model \cite{MCB:2002} and the new generalized Chaplygin gas (NGCG) model \cite{Zhang:2004gc}, have also been put forward. Moreover, actually, the cosmic acceleration can also be explained by the modified gravity (MG) theory, i.e., the theory in which the gravity rule deviates from the Einstein general relativity (GR) on the cosmological scales. The MG theory can yield ``effective dark energy'' models mimicking the real dark energy at the background cosmology level.\footnote{Usually, the growth of linear matter perturbations in the MG models is distinctly different from that in the DE models within GR.} Thus, if we omit the issue of growth of structure, we may also consider such effective dark energy models. A typical example of this type is the Dvali-Gabadadze-Porrati (DGP) model \cite{a:2000}, which arises from a class of braneworld theories in which the gravity leaks out into the bulk at large distances, leading to the accelerated expansion of the universe. Also, its theoretical variant, the $\alpha$DE model \cite{c:2003}, can fit the observational data much better. Facing so many competing dark energy models, the most important mission is to find which one on earth is the right dark energy model. But this is too difficult. A more realistic mission is to select which ones are better than others in explaining the various observational data. Undoubtedly, the right dark energy model can certainly fit all the astronomical observations well. The Planck satellite mission has released the most accurate data of cosmic microwave background (CMB) anisotropies, which, combining with other astrophysical observations, favor the base $\Lambda$CDM model \cite{Ade:2013zuv,Ade:2015xua}. But it is still necessary to make a comparison for the various typical dark energy models by using the Planck 2015 data and other astronomical data to select which ones are good models in fitting the current data. Such a comparison can also help us to discriminate which models are actually excluded by the current observations. We use the $\chi^2$ statistic to do the cosmological fits, but we cannot fairly compare different models by comparing their $\chi_{\rm min}^2$ values because they have different numbers of parameters. It is obvious that a model with more free parameters would tend to have a lower $\chi_{\rm min}^2$. Therefore, in this paper, we use the information criteria (IC) including the Akaike information criterion (AIC)~\cite{H.Akaike:1974} and the Bayesian information criterion (BIC)~\cite{G.schwarz:1978} to make a comparison for different dark energy models. The IC method has sufficiently taken the factor of number of parameters into account. Of course, we will use the uniform data combination of various astronomical observations in the model comparison. In this work, we choose ten typical, popular dark energy models to make a uniform, fair comparison. We will find that, compared to the early study~\cite{Zhang.xin:2010}, in the post-Planck era we are now truly capable of discriminating different dark energy models. The paper is organized as follows. In Sect. \ref{sec.2} we introduce the method of information criteria and how it works in comparing competing models. In Sect. \ref{sec.3} we describe the current observational data used in this paper. In Sect. \ref{sec.4} we describe the ten typical, popular dark energy models chosen in this work and give their fitting results. We discuss the results of model comparison and give the conclusion in Sect. \ref{sec.5}.

\label{sec.5} \begin{figure*} \includegraphics[width=10cm]{aicandbic.pdf}\\ \caption{\label{fig11} Graphical representation of the model comparison result. The order of models from left to right is arranged according to the values of $\Delta{\rm BIC}$, i.e., in order of increasing $\Delta{\rm BIC}$.} \end{figure*} We have considered ten typical, popular dark energy models in this paper, which are the $\Lambda$CDM, $w$CDM, CPL, GCG, NGCG, HDE, NADE, RDE, DGP, and $\alpha$DE models. To investigate the capability of fitting observational data of these models, we first constrain these models using the current observations and then make a comparison for them using the information criteria. The current observations used in this paper include the JLA sample of SN Ia observation, the Planck 2015 distance priors of CMB observation, the BAO measurements, and the $H_0$ direct measurement. The models have different numbers of parameters. We take the $\Lambda$CDM model as a reference. The NADE and DGP models have the same number of parameters as $\Lambda$CDM. The $w$CDM, GCG, HDE, RDE, and $\alpha$DE models have one more parameter than $\Lambda$CDM. The CPL and NGCG models have two more parameters than $\Lambda$CDM. To make a fair comparison for these models, we employ AIC and BIC as model-comparison tools. The results of observational constraints for these models are given in Table \ref{table1} and the results of the model comparison using the information criteria are summarized in Table \ref{table2}. To visually display the model-comparison result, we also show the results of $\Delta {\rm AIC}$ and $\Delta {\rm BIC}$ of these model in Fig. \ref{fig11}. In Table \ref{table2} and Fig. \ref{fig11}, the values of $\Delta {\rm AIC}$ and $\Delta {\rm BIC}$ are given by taking $\Lambda$CDM as a reference. The order of these models in Table \ref{table2} and Fig. \ref{fig11} is arranged according to the values of $\Delta {\rm BIC}$. These results show that, according to the capability of fitting the current observational data, the $\Lambda$CDM model is still the best one among all the dark energy models. The GCG, $w$CDM, and $\alpha$DE models are still relatively good models in the sense of explaining observations. The HDE, NGCG, and CPL models are relatively not good from the perspective of fitting the current observational data in an economical way. We can confirm that, in the sense of explaining observations, according to our analysis results, the NADE, DGP, and RDE models are excluded by current observations. In the models considered in this paper, only the HDE, NADE, RDE, and DGP models cannot reduce to $\Lambda$CDM, and among these models the HDE model is still the best one. Compared to the previous study \cite{Zhang.xin:2010}, the basic conclusion is not changed; the only subtle difference comes from the concrete orders of models in each group of the above three groups. In conclusion, according to the capability of explaining the current observations, the $\Lambda$CDM model is still the best one among all the dark energy models. The GCG, $w$CDM, and $\alpha$DE models are worse than $\Lambda$CDM, but still are good models compared to others. The HDE, NGCG, and CPL models can still fit the current observations well, but from the perspective of providing an economically feasible way, they are not so good. The NADE, DGP, and RDE models are excluded by the current observations.

1607.06262

1607

1607.02339_arXiv.txt

{With only two functional reaction wheels, \kepler\ cannot maintain stable pointing at its original target field and entered a new mode of observation called K2.} {We describe a new pipeline to reduce K2 Pixel Files into light curves that are later searched for transit like features.} {Our method is based on many years of experience in planet hunting for the CoRoT mission. Due to the unstable pointing, K2 light curves present systematics that are correlated with the target position in the \textsc{ccd}. Therefore, our pipeline also includes a decorrelation of this systematic noise. Our pipeline is optimised for bright stars for which spectroscopic follow-up is possible. We achieve a maximum precision on 6 hours of 6 ppm. The decorrelated light curves are searched for transits with an adapted version of the CoRoT alarm pipeline. } {We present 172 planetary candidates and 327 eclipsing binary candidates from campaigns 1, 2, 3, 4, 5 and 6 of K2. Both the planetary candidates and eclipsing binary candidates lists are made public to promote follow-up studies. The light curves will also be available to the community.} {}

Since the launch of \kepler\ in 2009 \citep{Borucki2010}, the number of confirmed exoplanet has grown exponentially, reaching 2933 known transiting exoplanets today and a few thousand unconfirmed candidates. The great diversity of discovered planetary systems is bringing a number of fundamental clues about the processes of planet formation and evolution. Moreover, \kepler\ has also revealed exoplanets around a diversity of hosts from M-dwarfs \citep{Dressing2015} to giant stars \citep{Quinn2015} and binaries \citep{Doyle2011}. The exceptional accuracy of the \kepler\ light curves reaching 15 parts per million (ppm) in 6 hours was in part due to its highly stabilised pointing. The failure of two out of four of the reaction wheels of the \kepler\ satellite put an end to prime \kepler\ mission since the pointing stability could not be maintained at the original target field. Fortunately, clever engineering allowed to give a second life to the \kepler\ satellite through a mission named K2 \citep{Howell2014}. K2 is balanced against solar radiation pressure in an unstable equilibrium and it needs to fire thrusters every 6 hours to maintain the pointing. K2 observes 4 fields a year close to the Ecliptic with a typical duration of 80 days. This important diminution of the time coverage imposed by the new pointing capabilities of the satellite is compensated by the possibility offered to the community to observe different regions of the Milky Way and thus different stellar populations. For example, K2 observes many more M dwarfs than its predecessor \citep{Crossfield2015, Petigura2015}, but also supernovae, clusters and a full campaign (\#9) will be dedicated to microlensing. Furthermore, the targets observed by K2 are globally brighter facilitating the confirmation and characterisation of the detected planetary systems. Initially K2 only delivered Pixels files and not light curves. However, since campaign 3, the K2 mission has produced light curves using its PDC pipeline. The initial lack of light curves triggered the development of many K2 pipelines by many groups. The challenge was to correct for the systematics introduced by the degraded pointing stability coupled with a mis-calibrated pixel response. To correct these systematics in the K2 data, several methods have been developed all presenting two main steps: a photometric extraction with a variety of aperture shapes and positions and a "correction of systematics". \citet{Vanderburg2014} and later on \citet{Armstrong2015b} used normal aperture photometry and corrected the systematics decorrelating the flux and the position variations of the target on the \textsc{ccd}. However, while \citet{Vanderburg2014} used the fact that the main motion of the line of sight was along 1 direction (the roll direction) which reduced the decorrelation to 1D, \citet{Armstrong2015b} opted for a 2D decorrelation. \citet{Aigrain2015} also used aperture photometry but coupled with gaussian processes to model the systematics at the same time as the stellar intrinsic variability. \citet{Huang2015} used the astrometric solution for the position of the targets to extract the light curves and 3 algorithms for the decorrelation of the position related systematics: external parameter decorrelation, trend filtering algorithm (\textsc{tfa}) \citep{Kovacs2005} and semi-periodic stellar oscillations via cosine-filtering. \citet{Foreman-Mackey2015} and \citet{Angus2016} do not decorrelate the systematics but fit the systematics together with their signal of interest respectively transits and periodic signals. They modelled the systematics by identifying common trends in the light curves similarly to the \textsc{tfa} method. Until now the only pipeline using optimised aperture is the one by \citet{Lund2015}. The optimised aperture is calculated with a data clustering algorithm called Density-based spatial clustering of applications with noise (\textsc{dbscan}) and their reduction is optimised for asteroseismology studies. Several groups also preformed planet search in the K2 data but only a couple have published lists of planetary candidates: \citep{Foreman-Mackey2015,Vanderburg2016} for campaign 0 to campaign 3. In this paper we present our K2 custom built pipeline and give planetary and binary candidates for stars brighter than $\textrm{K}_{\textrm{p}}\ \textrm{mag} = 14.7$ from campaigns 1 to 6. In the section~\ref{sec:LC_Extraction}, we describe the pipeline we developed to extract the K2 light curves and to correct the systematics. In section~\ref{sec:PhotPerformances}, we discuss the performances of our pipeline. In section~\ref{sec:TransitDetection}, we present our method to search for transits in the light curves, our vetting procedure and the results of our eclipse signals hunt for bright stars in the first six campaigns. We finish with a summary of our results in section~\ref{sec:summary}.

\label{sec:summary} We provide decorrelated light curves for all long cadence targets of K2 from C1 to C6 (discarding superstamps). The particularity of our pipeline relative to previously published ones is the determination of an optimal aperture and the precision of the centroid determination. Our apertures are in general not circular and follow nicely the \textsc{psf} of the stars. We show that our light curves have precision similar to the light curves from the nominal \kepler\ mission and they will be made public. Using these light curves we searched for eclipse signals on targets brighter than 14.7 magnitudes. This analysis results in a list of 172 planetary and 327 EB candidates. Among the 172 planetary candidates, 129 are new while this is the first release of eclipsing binary candidates from K2 data. All these products will be made public. Other teams have presented candidates for the K2 data till campaign 3. The different methods lead to common and non common candidates since the pipelines have slightly different performances for specific targets. The comparison between the methods will allow to fine tune the methods themselves but most importantly to build a more robust candidate list. There are some possible improvements to the pipeline like, for example, application to the short cadence data , and/or a more robust transit search, an automatic candidate validation instead of the current eyeballing that could result in larger number of candidates. However we think that making candidate lists available is urgent in order to optimise the follow-up of candidates and prevent waste of resources.

1607.02339

1607

1607.05343_arXiv.txt

In a previous study of the L1157 B1 shocked cavity, a comparison between NH$_3$(1$_0$-$0_0$) and H$_2$O(1$_{\rm 10}$--1$_{\rm 01}$) transitions showed a striking difference in the profiles, with H$_2$O emitting at definitely higher velocities. This behaviour was explained as a result of the high-temperature gas-phase chemistry occurring in the postshock gas in the B1 cavity of this outflow. If the differences in behaviour between ammonia and water are indeed a consequence of the high gas temperatures reached during the passage of a shock, then one should find such differences to be ubiquitous among chemically rich outflows. In order to determine whether the difference in profiles observed between NH$_3$ and H$_2$O is unique to L1157 or a common characteristic of chemically rich outflows, we have performed Herschel-HIFI observations of the NH$_3$(1$_0$-0$_0$) line at 572.5 GHz in a sample of 8 bright low-mass outflow spots already observed in the H$_2$O(1$_{\rm 10}$--1$_{\rm 01}$) line within the WISH KP. We detected the ammonia emission at high-velocities at most of the outflows positions. In all cases, the water emission reaches higher velocities than NH$_3$, proving that this behaviour is not exclusive of the L1157-B1 position. Comparisons with a gas-grain chemical and shock model confirms, for this larger sample, that the behaviour of ammonia is determined principally by the temperature of the gas.

A newborn protostar generates a fast and well collimated jet, possibly surrounded by a wider angle wind. In turn, the ejected material drives (bow-)shocks travelling through the surrounding high-density medium and traced by H$_2$ ro-vibrational lines at excitation temperatures of around 2000 K. Consequently, slower and cold (10--20 K) molecular outflows are formed by swept-up material, usually traced by CO. Shocks heat the gas and trigger several processes such as endothermic chemical reactions and ice grain mantle sublimation or sputtering. Several molecules, such as H$_2$O, NH$_3$, CH$_3$OH, H$_2$CO, undergo spectacular enhancements by orders of magnitude in their abundances \citep{vandis98}, as observed at mm-wavelengths in a number of outflows \citep{Garay98,Bachiller97,Jorgensen07}. The link between the gas components at $\sim$ 10 K and the hot 2000 K shocked component is crucial to understand how the protostellar wind transfers momentum and energy back to the ambient medium. In this context, studies of the chemical composition of typical molecules in bow-shocks are essential because they represent a very powerful diagnostic tool for probing their physical conditions. Such studies are also paramount to get a comprehensive understanding of chemistry throughout the various phases of the interstellar medium. \begin{figure}% \centering \includegraphics[bb=82 27 537 450,angle=-90,width=9.0cm]{iras4a_map.ps} \caption{NGC 1333 IRAS4A outflow and corresponding observed positions. The background image represent the H$_2$O emission at 179 $\mu$m from the WISH program \citep{Santangelo14}. The circles show the vertical (V) and horizontal (H) polarization HPBW at the observed positions, whose centers are indicated by the triangles. The stars mark the position of the continuum sources \citep[A and B:][]{Looney00}. } \label{map-iras4} \end{figure} \begin{figure}% \centering \includegraphics[bb=72 65 532 428,angle=-90,width=9.0cm]{l1157_map.ps} \caption{L1157 outflow and corresponding observed positions. The background image represent the H$_2$O emission at 179 $\mu$m from the WISH program \citep{Nisini10b}. The blue contours show the SiO emission from \citet{Bachiller2001}. The circles show the vertical (V) and horizontal (H) polarization HPBW at the observed positions, whose centers are indicated by the triangles. The star marks the position of the central source (L1157-mm). } \label{map-l1157} \end{figure} As part of the Herschel Key Program CHESS \citep[Chemical Herschel Surveys of Star forming regions:][]{Cecc10}, the bow-shock L1157-B1 has been investigated with a spectral survey using the HIFI instrument. From the comparison between NH$_3$(1$_0$--0$_0$) and H$_2$O(1$_{\rm 10}$--1$_{\rm 01}$) profiles, a straightforward estimate of the relative abundance ratios of the gas at different velocities was obtained \citep{Codella10}. As a notable example, the NH$_3$/H$_2$O intensity ratio decreases by a factor of $\sim$ 5 moving towards higher velocities suggesting, in case of optically thin emission along the wings, a similar decrease in the abundance ratios. Other tracers of shocked material such as CH$_3$OH and H$_2$CO show the same profile as that of NH$_3$. In \citet{Codella10} we propose that the difference between the H$_2$O and other species reflects different formation mechanisms: for example, while the bulk of NH$_3$ is released from the grain mantles, H$_2$O is enhanced by the release of the icy mantles {\it as well as} by endothermic reactions occurring in the warm ($\geq 220$ K) shocked gas, which convert all gaseous atomic oxygen into water \citep[e.g.,][and references therein]{Isaskun08}. However, a model by \citet{Viti11} made especially for these data set suggests that the differences observed in the profile of the different molecular tracers are due mainly to the temperature of the gas: if the latter undergoes a period at a temperature close to 4000 K, then NH$_3$ is easily destroyed by the reaction with hydrogen which leads to NH$_2$ $+$ H$_2$ (this reaction has a high barrier of $\sim$5000 K), while H$_2$O remains high in abundance. Such scenario can be explained by the presence of a C-type shock whose pre-shock density and velocity are such that the maximum temperature of the shock reaches 4000 K along the B1 shock of L1157. These findings called for observations of more molecular shocked regions associated with protostellar outflows to investigate whether the difference in profiles between H$_2$O and other species are unique to L1157 or whether it is an ubiquitous characteristic of chemically rich outflows. \begin{figure}% \centering \includegraphics[bb=152 67 511 342,angle=-90,width=9.0cm]{l1448_map.ps} \caption{L1448 outflow and corresponding observed positions. The background image represent the H$_2$O emission at 179 $\mu$m from the WISH program \citep{Nisini13}. The circles show the vertical (V) and horizontal (H) polarization HPBW at the observed positions, whose centers are indicated by the triangles. The green crosses indicate the positions of the millimeter continuum sources \citep[N, S, A, B, and W:][]{Kwon06}. } \label{map-l1448} \end{figure} In this article we present observations of the $J_K=$1$_{\rm 0}$--0$_{\rm 0}$ transition of ortho-NH$_3$ at 572.5 GHz in a number of outflow spots already observed in the ortho-H$_2$O(1$_{\rm 10}$--1$_{\rm 01}$) line as part of the Herschel Key Program WISH \citep[Water In Star-forming regions with Herschel:][]{WISH} and reported by \citet{Tafalla13}. In Sect. 2 the target selection and Herschel observations with HIFI are described, in Sect. 3 we report the line profiles obtained, in Sect. 4 and 5 we develop the analysis of the data, and in Sect. 6 we present the summary and conclusions. \begin{figure}% \centering \includegraphics[angle=90,width=8.5cm]{iras2a-key.ps} \caption{NGC1333 IRAS2A outflow and corresponding observed positions \citep[image adapted from][]{Bachiller98}. The contours show the CH$_3$OH ($2_k-1_k$) PdBI map. The circles show the vertical (V) and horizontal (H) polarization HPBW at the observed positions, whose centers are indicated by the triangles. The star indicates the position of the central source (IRAS2A). } \label{map-iras2} \end{figure}

In the following we summarize the main results and conclusions from our {\it Herschel}/HIFI observations of the ammonia emission from protostellar outflows: \begin{enumerate} \item We detected the NH$_3$ emission from all eight outflow positions we have observed. In all the cases, the ammonia emission reaches terminal velocities ($V_{\rm ter}$) that are lower than H$_2$O, proving that this behaviour is not exclusive of the L1157-B1 position. In addition to ammonia, all the bonus lines (due to CS, H$_2$CO, and CH$_3$OH) were detected in only IRAS4A-B and IRAS4A-R positions, confirming the chemical richness of these regions. \item Comparisons with chemical modelling confirms that the behaviour of ammonia is determined principally by the temperature of the gas. \item While a quantitative comparison between models and observations is not feasible without a proper line radiative transfer model, we constrain the pre-shock density and/or shock velocity for each object based on a comparison of abundance trends. We find that, while several models show agreement with the profiles of the different objects, the best matching model for L1157-B2 has a very low pre-shock density (10$^3$ cm$^{-3}$ and velocity (10 km s$^{-1}$), while IRAS2A-B abundances are best reproduced by a gas that has undergone a relatively high velocity shock (45 km s$^{-1}$) with a pre-shock density of $\sim$ 10$^5$ cm$^{-3}$. L1448-B2, IRAS4A-R and IRAS2A-R are matched by models where ammonia is heavily destroyed at high velocities, as explained above due to the short period when the temperature of the gas is high, at 4000 K. We are not able to constrain the pre-shock density for these objects as it can range from as low as 1000 cm$^{-3}$ to as high as 10$^{6}$ {\it as long as} the maximum temperature of the shock is 4000 K, which can be achieved for a shock velocity of $\sim$ 40km s$^{-1}$. The best matching models also indicate a low level of depletion in the cold phase prior to the passage of the shock, hence it is likely that the pre-shock density is in fact towards the lower limit. Finally, L1157-R, L1448-R4 and IRAS4A-B seem to resemble very closely the abundance profile of L1157-B1. They are likely therefore to have a pre-shock density of 10$^5$-10$^6$ cm$^{-3}$ and a shock velocity of the order of 40 km s$^{-1}$, although we can not exclude a faster shock with a lower pre-shock density: in other words, the behaviour of the H$_2$O/NH$_3$ is again determined by the high temperature the gas can attain and the latter can be achieved by more than one combination of shock parameters. In terms of theoretical abundances, a high H$_2$O/NH$_3$ for much of the dissipation length is only reached within a small range of maximum shock temperatures; in terms of profiles, on the other hand, half of our sample show a high H$_2$O/NH$_3$ ratio: not all of them however require a maximum shock temperature to be close to 4000 K as group 3 models indicate. It is also important to point out that within the observed beam it is unlikely that we are seeing one episodic shock or a group of shocks, all at the same velocities; hence with the present observations it is not possible to draw any statistically meaningful conclusion on the type of shock that is prevalent in outflows around low mass stars. In conclusion, the H$_2$O/NH$_3$ as a function of velocity can be used to determine the most likely combination of 'pre-shock density and shock velocity', although it is not sufficient in itself to be able to constrain each individual parameter. \end{enumerate}

1607.05343

1607

1607.03753_arXiv.txt

In this paper, we use the restricted three body problem in the binary stellar systems, taking photogravitational effects of both the stars. The aim of this study is to investigate the motion of the infinitesimal mass in the vicinity of the Lagrangian points. We have computed semi-analytical expressions for the locations of the collinear points with the help of the perturbation technique. The stability of the triangular points is studied in stellar binary systems Kepler-34, Kepler-35, Kepler-413 and Kepler-16. To investigate the stability of the triangular points, we have obtained the expressions for critical mass which depends on the radiation of both primaries. Fourier-series method is applied to obtain periodic orbits of the infinitesimal mass around triangular points in binary stellar systems. We have obtained Fourier expansions of the periodic orbits around triangular points upto third order terms. A comparison is made between periodic orbits obtained by Fourier-series method and with Runge-Kutta integration of fourth order.

\label{sec:3sec1} In a binary stellar system, each star is exerting a gravitational force on any object with mass in its vicinity. The study of binary stellar systems is very important because of more than 60 percent of the stars in the solar neighborhood build such systems which are interesting systems from the dynamical point of view. For such system the restricted three body problem is an appropriate model. Restricted three body problem is a special case of the three body problem \citep{barrow1999poincar,roy2005orbital,marchal2012three,Mia2016MNRAS.457.1089M} which is very important in many scientific fields such as celestial mechanics, galactic dynamics, molecular theory etc. and it is the most important problem in exoplanetary systems. In the circular restricted three body problem two primaries move in circular orbits around their common center of mass and the third body i.e., small body moves in the same plane and it is influenced by the gravitational forces from the primaries but it does not influence the motion of the two primaries. So many authors have used this problem to find the stability of planet, asteroids or small bodies like meteoroids. Many researchers worked on the study of the existence and stability of libration points in the framework of restricted three body problem considering additional perturbations from radiation pressure \citep{alvarez2014stability,eapen2014study}, the stokes drag \citep{Jain2015Ap&SS.358...51J}, the solar wind drag \citep{pal2015geometry}, oblate and triaxial bodies \citep{singh2014effects}, the variable masses \citep{Abouelmagd2015Ap&SS.357...58A}, the potential from a disk \citep{Kishor2013MNRAS.436.1741K}, and the Yarkovsky effect \citep{ershkov2012yarkovsky}. There are several authors who have studied the significance of radiation from stars in binary systems \citep{Zhou1988Ap&SS.141..257Z,Papadakis2006Ap&SS.305...57P,roman2007photogravitational, das2008effect}. Also, some authors studied the effect of radiation from Sun on the motion of small particle in our solar system \citep{Robertson1937MNRAS..97..423R,Kushvah2008Ap&SS.315..231K,Abouelmagd2013Ap&SS.344..321A, Kumari2013Ap&SS.344..347K,tiwary2015computation}. \cite{Todoran1993Ap&SS.201..281T} studied the effects of the radiation pressure in the restricted three-body problem and the existence of the `out-of-plane' equilibrium points. They found that within the framework of the stellar stability, the five Lagrangian points are the only equilibrium points, at least as far as the force of the radiation pressure is taken into account. The linear stability of the triangular equilibrium points in the photogravitational elliptic restricted problem have examined by \cite{Markellos1992Ap&SS.194..207M}. They determined the stability regions in the space of the parameters of mass, eccentricity, and radiation pressure. They also found that radiation pressure of the larger body for solar system case exerts only a small quantitative influence on the stability regions. \cite{schwarz2012stability} studied the stability of the Lagrangian point L4 in the spatial restricted three-body problem and the possibility of inclined Trojan-like objects in exoplanetary systems (single and binary star systems). They investigated stability computing stability maps by numerical methods. In the case of circular motion of the primary bodies, they have shown that there are stable orbits up to an inclination $i = 61$ degree of the test particle. \cite{abouelmagd2013stability} studied the existence of triangular points and their linear stability when the primaries are oblate spheroid and sources of radiation in the restricted three-body problem. They observed that the locations of the triangular points are affected by the oblateness of the primaries and solar radiation pressure. They also shown that these points are stable for $0\leq\mu\leq\mu_c$ and unstable for $\mu_c\leq\mu\leq\frac{1}{2}$, where $\mu_c$ is the critical value of mass. Recently, \cite{Bosanac2015CeMDA.122...27B} investigated periodic motions near a large mass ratio binary system within the context of the circular restricted three-body problem. They used stability analysis to explore the effect of the mass ratio on the structure of families of periodic orbits. But most of their works includes the use of $q_i=1-\beta_i, i=1,2$, where $\beta_{i}$ is the ratio of the radiation pressure force of binary components to gravitational force of binary components. The luminosity of some star of binary's is unknown so in this present work using relation between luminosity and mass of a star, we have calculated the values of luminosity for different binary systems. Also using the realistic relation among $\beta$, luminosity and mass of the binary components, we have found the values of $q$ for different binary system. After that we study the motion of a small particle around a binary star system in the presence of radiation from both the stars i.e., primaries. Moreover, we have found periodic orbits around triangular points with the aid of Fourier series. Periodic orbits with the help of Fourier series of the classical RTBP have found by \cite{pedersen1935fourier}. Whereas, we determine periodic orbits of the RTBP with radiating primaries in binary systems. Also for the purpose of matching the results we have determined periodic orbits in Earth-Moon system and the result found agree with that of \cite{pedersen1935fourier}. The purpose of this present study is to describe and examine the motion and to find the orbits of the infinitesimal body in stellar binary system using the model of circular restricted three body problem taking into account of the photo gravitational effects of both the stars along with the gravitational forces. We analyze the stability of triangular points in the binary stellar system (Kepler-34, Kepler-35, Kepler-413 and Kepler-16) by the method described in \cite{moulton2012introduction,Szebehely1967torp.book.....S,Murray2000ssd..book.....M}. We obtained periodic orbits around triangular points with the help of Fourier series. The present paper is organized as follows. In Section \ref{sec:3sec2}, we introduce the equations of motion. Numerical values of physical parameters are discussed in Section \ref{sec:3sec3}. In Section \ref{sec:3sec4}, we found location of Lagrangian points of binary stellar systems. Stability of triangular points are discussed in Section \ref{sec:3sec5}. Fourier series in the vicinity of triangular points are discussed in Section \ref{sec:3sec7}. Finally, Section \ref{sec:3sec8} is devoted to conclusions.

\label{sec:3sec8} \label{sec:con} We have studied the motion of a infinitesimal mass in the context of the binary stellar systems Kepler-34, Kepler-35, Kepler-413 and Kepler-16. We have applied restricted three body problem as the model, considering gravitational and radiation effects on the particle from both the stars. With the help of the perturbation technique, semi-analytical expressions for the location of collinear points have been obtained. A comparison is presented in Table \ref{tab:2a} between the analytical and numerical solution of collinear points which shows an excellent level of agreement. We examined the linear stability of triangular points by obtaining the expressions for critical mass. We have found that critical mass depends on the radiation of both primaries. It is observed that for the stellar binary systems the roots of the characteristic equations are complex conjugate with one root has positive real part. Hence the motion around triangular points are unstable. Further, we have obtained Fourier expansions of the periodic orbits around triangular points in the CR3BP with radiation pressure from binaries. We observed that the periodic orbits are expanding for decreasing values of $q_1$ and $q_2$. We have also observed in Sun-Earth system that as we decrease the values of $q_1$, periodic orbits are shifting towards the origin and expanding. Moreover, periodic orbits obtained by Fourier series method have been compared with that of RK4 numerical methods. Also using the actual values of radiation parameters, we have computed periodic orbits in four binary systems. Moreover, since detection of binary systems are increasing in number, studies of such systems provide an important contribution for future observations. Furthermore, as the higher order approximation produce the closer orbit, the work would be extended by considering higher order terms in the governing equations of motion and in the Fourier-series.

1607.03753

1607

1607.01610_arXiv.txt

\baselineskip18pt It has been a puzzle whether quarks may exist in the interior of massive neutron stars, since the hadron-quark phase transition softens the equation of state (EOS) and reduce the neutron star (NS) maximum mass very significantly. In this work, we consider the light U-boson that increases the NS maximum mass appreciably through its weak coupling to fermions. The inclusion of the U-boson may thus allow the existence of the quark degrees of freedom in the interior of large mass neutron stars. Unlike the consequence of the U-boson in hadronic matter, the stiffening role of the U-boson in the hybrid EOS is not sensitive to the choice of the hadron phase models. In addition, we have also investigated the effect of the effective QCD correction on the hybrid EOS. This correction may reduce the coupling strength of the U-boson that is needed to satisfy NS maximum mass constraint. While the inclusion of the U-boson also increases the NS radius significantly, we find that appropriate in-medium effects of the U-boson may reduce the NS radii significantly, satisfying both the NS radius and mass constraints well.

The equation of state (EOS) of isospin asymmetric nuclear matter is of prime importance for the investigation of nuclear structure~\cite{Hor01,jia10,jia10a,wang14}, heavy-ion reaction dynamics~\cite{Li01,Li08}, and many issues in astrophysics~\cite{Li08,lat00,lat01,lat04,lat07,Ste05}. However, the EOS of asymmetric matter at supra-normal densities are still poorly known~\cite{Li08,li15}, though the constraint on the symmetric part of the EOS at supra-normal densities can be extracted from the collective flow data of high energy heavy-ion reactions~\cite{da02}. Different models and approaches can produce rather different high-density behaviors~\cite{br00,chen05,fu06,lie7}, while the complexity arises for the EOS of asymmetric matter when the phase transitions, such as hyperon productions, meson condensations, and quark deconfinement, take place at high densities. Normally, phase transitions can soften the EOS and reduce the neutron star (NS) maximum mass significantly. Recently, massive NS's, for instance, the pulsars PSR J1614-2230 with $M=1.97\pm0.04M_\odot$~\cite{Demorest2010}, and PSR J0348+0432, with $M=2.01\pm0.04M_\odot$~\cite{Antoniadis2013} were identified with the high-precision measurements. The larger NS mass means the stiffer EOS of NS matter at high densities. Considering the phase transitions in the NS core and the EOS constraint from the collective flow data at high densities, it is difficult for the nuclear models to reproduce NS's as massive as $2M_\odot$. This indeed proposes a challenge to the nuclear EOS at supra-normal densities. One may doubt whether there are new degrees of freedom besides nucleons in NS's~\cite{Demorest2010,Trumper04,Ozel06}. On the other hand, one may contemplate what interactions can allow the new degrees of freedom in massive NS's\cite{Bednarek11,jiang12,Ruster04, Horvath04, Alford2005, Alford2007,Ippolito08, Fischer10, Kurkela10a,Kurkela10b,Weissenborn11,Bonanno12,Yasutake14,Xia14}. The mixture of low-density hadronic matter and high-density quark matter may form in the NS core after the hadron-quark transition occurs with the increase of density. In order to evade the potential confliction between the softened EOS due to the appearance of the quark phase and the observed massive NS's, it is necessary to stiffen the quark EOS, for instance, by considering the strong coupling and/or color superconductivity~\cite{Ruster04, Horvath04, Alford2005, Alford2007,Ippolito08, Fischer10, Weissenborn11}. Recently, a new repulsion, provided by the U-boson, was introduced for nucleons in NS's~\cite{Kr09,Wen09,zhang11}. The light U-boson, first proposed by Fayet~\cite{fa80}, might be regarded as the mediator of the putative fifth force~\cite{fa86,fi99,ad03}. Recently, this light U-boson has been considered as the interaction mediator of the MeV dark matter to account for the bright 511 keV $\gamma$-ray from the galactic bulge~\cite{Bo04,Zhu07,Fa07,Je03}. The coupling of the U-boson with the nucleons is very weak, but can increase appreciably the NS maximum mass~\cite{Kr09,Wen09,zhang11} and influence the transition density at the inner edge separating the liquid core from the solid crust, effectively~\cite{Zheng12}. In particular, the weak coupling to baryons plays a striking role in stabilizing neutron stars in the presence of the super-soft symmetry energy~\cite{Wen09} that is extracted from the FOPI/GSI data on the $\pi^{-}/\pi^{+}$ ratio in relativistic heavy-ion collisions~\cite{Xiao08}. In this work, we consider the effect of the U-boson on the softened EOS due to the hadron-quark phase transition. We adopt the relativistic mean-field (RMF) theory, which achieved great success in the past few decades~\cite{Wal74,Bog77,Chin77,Ser86,Rei89,Ring96,Ser97,Ben03,Meng06,Ji07}, to describe hadronic matter, and the MIT bag model for quark matter~\cite{ch74,he86}. For the hadron-quark phase transition, we use Gibbs construction~\cite{Glendenning92,Glendenning01} to determine the mixed phase of hadronic and quark matter. As the stiffening role of the U-boson depends on the softness of the models~\cite{zhang11}, we will examine how the U-boson stiffens the hybrid EOS's initiated with different hadronic models. Since quarks in the bag model are free of interaction, it is also interesting to investigate briefly the effect of the effective correction from the perturbative QCD~\cite{Alford2005,Alcock86,Haensel86,Fraga01,xu15}. Eventually, we will investigate the properties of hybrid stars with various EOS's and discuss how the mass and radius constraints of the NS observations can be satisfied. The paper is organized as follows. In Sec.~\ref{rmf}, we present briefly the formalism. In Sec.~\ref{results}, numerical results and discussions are presented. At last, a summary is given in Sec.~\ref{summary}.

\label{results} For the hadronic phase in hybrid stars, we consider the simple compositions: neutrons, protons, electrons and muons. We do not include hyperons in this work. The appearance of hyperons can largely soften the equation of state, thus reducing the mass of neutron stars greatly. Due to the accurate mass measurement of large-mass neutron stars, people even conclude that the hyperon EOS has to be ruled out~\cite{Demorest2010}. Indeed, the onset densities of hyperons are very model-dependent. With the RMF parameter set NL3, the $\Ld$ hyperon appears at about $0.28fm^{-3}$ ~\cite{Jiang06,Yang08}, while it is at about $0.48fm^{-3}$ with the extended MDI interaction~\cite{Xu10}. The superfluidity of hyperons renders the $\Ld$ hyperons to appear at around $0.64fm^{-3}$~\cite{Takatsuka02}, which just affects the properties of neutron star slightly. When the hyperonic degrees of freedom are included, it is found that the hadron-quark transition density can not change continuously with the emergence of hyperons. This means that the appearance of hyperons disfavors the hadron-quark phase transition in the RMF models. Similar effect of the hyperons is found in the literature~\cite{Chen11}. As implied by the large-mass NS's, the in-medium hyperon potential should be density-dependent~\cite{jiang12}. We may further explore the density-dependent interaction of hyperons in the future and it is beyond the scope of the present work. Among a number of nonlinear RMF models, we select two typical best-fit parameter sets, NL3~\cite{La97} and FSUGold~\cite{Pie05}, to describe the hadron phase of hybrid star matter. Parameters and saturation properties of these two parameter sets are listed in Table~\ref{t:t1}. Usually, the nonlinear RMF models include the nonlinear self-interacting meson terms to simulate appropriate in-medium effect of the strong interaction. The parameter set NL3 includes the nonlinear self-interaction of the $\sigma$ meson, while in addition to the latter, the nonlinear self-interaction of the $\omega$ meson is also included in FSUGold. The resulting EOS of the FSUGold is much softer than that of the NL3 at high densities. It is known before that the vector U-boson can stiffen the EOS, and the stiffening role is much more significant for the soft EOS in pure hadronic matter~\cite{zhang11}. It is thus interesting to see whether we can observe a similar phenomenon in hybrid star matter, if hadronic phase is described with these two different RMF models. \begin{table}[htb] \caption{Parameters and saturation properties for the hadron phase models NL3 and FSUGold. Meson masses, incompressibility $\kappa$ and symmetry energy are in units of MeV, and the density is in unit of $fm^{-3}$. \label{t:t1}} \begin{center} \begin{tabular}{ c c c c c c c c c c c c c c c} \hline\hline &$g_\sg$ & $g_\om$& $g_\r$& $m_\sg$ & $m_\om$ & $m_\rho$ & $g_2$ & $g_3$& $c_3$ & $\Ld_V$ & $\rho_0$ & $\kappa$ & $M^*/M$ & $E_{sym}$ \\ \hline NL3 &10.217 &12.868 &4.474 &508.194 &782.501 &763 &10.431 &-28.890 &- &- &0.148 &271.8 &0.60 &37.4\\ FSUGold &10.592 &14.302 &5.884 &491.500 &782.500 &763 &4.277 &49.934 &418.39 &0.03 &0.148 &230.0 &0.61 &32.5\\ \hline\hline \end{tabular} \end{center} \end{table} While it is still an open question to determine the density at which the hadron-quark phase transition occurs, we investigate the onset of the phase transition in two manners: one is to fix the transition density $\rho_c$ by adjusting the bag constant, and the other to fix the bag constant. Given the bag constant, we also obtain the transition density which is now dependent on the hadron phase models. Once the transition density is determined, we can construct the mixed phase with the Gibbs conditions and then obtain the EOS of hybrid star matter. In this work, we choose the bag constant to be $B^{1/4}=180\sim220$ MeV which is a reasonable range between those values by fitting the light-hadron spectra~\cite{DeGrand75,Bartelski84,Chanowitz83,Carlson83} and those (e.g., 250 MeV) used in the hydrodynamical evolution of the quark gluon plasma~\cite{Ko89,He00}. The present range is also close to that used in the literature~\cite{Jiang13,Xu10,Weissenborn11}. For the choice of the transition density, we are referred to the literature where it is around $1.5\sim4\rho_0$ ~\cite{Alford2005,Sharma07,Shao11,Cavagnoli11}. In addition to the fact that the EOS of hybrid star matter with a transition density $4\rho_0$ or even higher just has a minor effect on the star maximum mass, we choose the range ($2\sim3\rho_0$) for the transition density in this work. As an example, we calculate here the various phase boundaries with $B^{1/4}=180MeV$ and with the transition densities $2\rho_0$ and $3\rho_0$. For the parameter set NL3, the transition density with $B^{1/4}=180MeV$ from the hadronic phase to the mixed phase is about $0.20fm^{-3}$, and $0.77fm^{-3}$ from the mixed phase to pure quark phase. The extent of the mixed phase is about $0.57 fm^{-3}$. For the parameter set FSUGold, the extent of the mixed phase is $1.47fm^{-3}$ with the onset transition density $0.28fm^{-3}$. For the case with the transition density fixed at $2\rho_0$, the extent of the mixed phase is about $0.95fm^{-3}$ with the parameter set NL3, while it is $1.7fm^{-3}$ with the parameter set FSUGold. For the transition density at $3\rho_0$, the extents of the mixed phase are $1.99fm^{-3}$ and $2.34fm^{-3}$ with the NL3 and FSUGold, respectively. \begin{figure}[thb] \begin{center} \vspace*{-5mm} \includegraphics[height=8.0cm,width=10.0cm]{eospt} \end{center} \vspace*{-5mm}\caption{(Color online) Equation of state of NS matter with and without quarks. The curves with the U-boson contribution are also shown for comparisons. The value of the ratio parameter $c_u=(g_u/m_u)^2$ is in unit of $GeV^{-2}$. The hadronic matter EOS is obtained with the NL3 (left panel)and FSUGold (right panel). \label{eospt}} \end{figure} The phase transition usually tends to soften the EOS, as is consistent with the requirement of spontaneous stability in natural processes. This is also true for the hadron-quark phase transition in isospin asymmetric nuclear matter. In Fig.~\ref{eospt}, we depict the EOS of hybrid star matter for various cases with the onset density $3\rho_0$ for hadron-quark phase transition. In comparison to the EOS without phase transition as shown in Fig.~\ref{eospt}, we see that the EOS of isospin asymmetric matter is greatly softened by the hadron-quark phase transition. It is striking to see that the stiff EOS with the NL3, which is not favored by the constraint from the flow data of the heavy-ion reactions~\cite{da02}, even becomes much softer than the soft FSUGold EOS due to the hadron-quark transition. With the inclusion of the U-boson, the soft EOS is stiffened greatly. Because the EOS of hybrid star matter initiated with the NL3 is now much softer than that with the FSUGold, the stiffening effect turns out to be much more appreciable for the EOS initiated with the NL3. The similar stiffening role of the U-boson can also be clearly seen in the case without phase transition, as we compare curves with the NL3 and FSUGold. Since the softening of the EOS due to the hadron-quark phase transition reduces largely the maximum mass of NS's, we will see below that the inclusion of the U-boson plays an important role in satisfying the maximum mass constraint for NS's. \begin{figure}[thb] \begin{center} \vspace*{-5mm} \includegraphics[height=12.0cm,width=12.0cm]{eosnf} \end{center} \vspace*{-5mm}\caption{(Color online) Equation of state of hybrid star matter. The hadronic phase is given by two RMF parameter sets, NL3 (lower panels) and FSUGold (upper panels). For quark phase, we present the results with the fixed bag constant (left panels ) and fixed transition density (right panels). In each panel, the various curves are obtained with different ratio parameters $c_u$. \label{eosnf}} \end{figure} To see the role of the U-boson specifically, we depict in Fig.~\ref{eosnf} the EOS's of hybrid star matter for a set of the ratio parameter $c_u=(g_u/m_u)^2$. Similar to that shown in Fig.~\ref{eospt}, the EOS of hybrid star matter is stiffened significantly due to the repulsion provided by the U-boson. We see that the EOS with the soft (FSUGold) and stiff (NL3) hadron phase models acquires similar stiffening especially for the case with the fixed bag constant. This is different from the case in pure hadronic matter where a much more significant stiffening effect is produced by the soft model~\cite{zhang11}. While for the case with the fixed transition density, the stiffening role of U-boson is more significant for the stiff parameter set NL3. The reason for these to occur lies in the following facts. In pure hadronic matter with RMF models, there is the cancelation between the repulsion provided by the vector meson and the attraction provided by the scalar meson. Thus, more significant cancelation in the soft model sharpens the importance of the repulsion provided by the U-boson. While quarks are modeled by the MIT bag model with the same bag constant, the repulsion provided by the U-boson plays the same stiffening role in the quark phase EOS, after the hadron-quark phase transition occurs. In the case with the fixed transition density, the quark phase EOS connected to the hadronic phase with the NL3 is softer than that with the FSUGold, because of the larger bag constant. Meanwhile, due to the phase equilibrium in the mixed phase, the EOS initiated with the NL3 is softer than that with the FSUGold with increasing the density. Thus, the U-boson provides a more significant stiffening role in the EOS initiated with the stiff NL3 in the hadron phase. \begin{figure}[thb] \begin{center} \vspace*{-5mm} \includegraphics[height=9.0cm,width=8.0cm]{MODE} \end{center} \vspace*{-5mm}\caption{(Color online) The influence of the U-boson on the EOS of hybrid star matter for different hadron phase models at the given bag constant $B=(200MeV)^4$ and transition density $\r_c=3\r_0$, respectively. \label{MODE}} \end{figure} Since in pure hadronic matter the stiffening effect of the U-boson is relevant to the extent of softness of the EOS's, it is interesting to examine whether such a dependence exists in the EOS of hybrid star matter with quark degrees of freedom. Shown in Fig.~\ref{MODE} is the EOS's of hybrid star matter with and without the contribution of the U-boson for two RMF parameter sets. We can see that after the occurrence of the hadron-quark phase transition, the difference in EOS's with the fixed transition density seems to be more apparent than that with the fixed bag constant. This is understandable, because the different bag constants are used to obtain the same transition density. However, compared to the large difference in the EOS's in pure hadronic matter, as shown in Fig.~\ref{eospt}, we see that the phase transition reduces the difference largely in EOS's at high densities. Compared to the evolution of EOS's with the NL3 and FSUGold, we see that the stiffer EOS undergoes a more appreciable softening of the EOS in the mixed phase, as the quark phase is given by the same MIT bag model either with the fixed bag constant or with a given transition density. We have also examined results for the bag constants ranging from 180 MeV to 220 MeV and transition densities from $3\r_0$ to $4\r_0$. The increase of the bag constant or transition density gives rise to a larger extent of the mixed phase, consistent with that found in Ref.~\cite{prak97}. For instance, the transition density with $B^{1/4}=200$ MeV is 0.28$fm^{-3}$ with the parameter set NL3, the extent of the mixed phase increases from 0.57 $fm^{-3}$ (with $B^{1/4}=180$ MeV) to 1.1 $fm^{-3}$. With the parameter set FSUGold, the extent of the mixed phase gets a more apparent rise, and it is 2.8 $fm^{-3}$ starting from the transition density 0.47 $fm^{-3}$. In any case, the appreciable stiffening role of the U-boson in the soft EOS's can be clearly observed as in Fig.~\ref{MODE}, similar to that in Fig.~\ref{eosnf}. In this work, we make comparative study with the stiff and soft EOS's. It is worthy to note the parameter $\Lambda_v$ may also affect the EOS of asymmetric matter. The parameter $\Ld_v$ does not appear in the parameter set NL3. In the parameter set FSUGold, we may change this parameter (also change $g_\rho$) by keeping the symmetry energy at saturation density unchanged. Here, we choose three parameter sets: FSUGw15 ($\Ld_v=0.015$), FSUGlod ($\Ld_v=0.03$), FSUGw45 ($\Ld_v=0.045$), all of which can fit the ground-state properties of $^{208}$Pb~\cite{jia10a}. The transition density increases moderately with the rise of the $\Ld_v$, since the stiffness of the symmetry energy that is controlled by the $\Ld_v$ can affect the chemical equilibrium. With the bag constant $B^{1/4}=180MeV$, we obtain following transition densities: 0.252, 0.282, and 0.297$fm^-3$ with the parameter sets FSUGw15, FSUGlod and FSUGw45, respectively. \begin{figure}[thb] \begin{center} \vspace*{-5mm} \includegraphics[height=12.0cm,width=12.0cm]{pqcd1} \end{center} \vspace*{-5mm}\caption{(Color online) The EOS of hybrid star matter with the inclusion of the QCD correction ($c=0.3$) with and without the U-boson. \label{pqcd1}} \end{figure} In above, the quarks with the bag model are free of interactions. As the QCD coupling is not negligible at densities of interest for compact star physics, it is here necessary to discuss briefly the QCD correction. It is reasonable to take the value of $c$ in Eqs.(\ref{eq:ou1}),(\ref{eq:ou2}) to be 0.3~\cite{Alford2005}. Shown in Fig.~\ref{pqcd1} is the EOS with the QCD correction. We see from Fig.~\ref{pqcd1} that the QCD correction can stiffen moderately the EOS of hybrid star matter at high densities. The QCD correction may affect the phase transition and increase the extent of the mixed phase due to the fact that the quark phase pressure is now modified by the QCD correction. For instance, with the bag constant $B^{1/4}=180MeV$, the onset density of the mixed phase with the parameter set NL3 is then about $0.28fm^{-3}$ with the extent of the mixed phase being about $0.96fm^{-3}$. With the parameter set FSUGold, the onset density of the mixed phase is about $0.46fm^{-3}$, and the extent of the mixed phase grows dramatically up to $6.49fm^{-3}$. As observed in Fig.~\ref{pqcd1}, the QCD correction to the EOS is moderately dependent on the hadron phase models. Specifically, the QCD correction results in an extent of the mixed phase with the soft FSUGold being larger than that with the stiff NL3. We should say that the specific dependence of the QCD correction on the hadron phase models is due to the connection built upon the phase equilibrium conditions. For comparisons, we also plot the results with and without the U-boson. As shown in Fig.~\ref{pqcd1}, we see that the QCD correction is just moderate, compared with the contribution of the U-boson. The stiffening role of the U-boson is similar in cases with and without the QCD correction. \begin{figure}[thb] \begin{center} \vspace*{-5mm} \includegraphics[height=12.0cm,width=12.0cm]{MR} \end{center} \vspace*{-5mm}\caption{(Color online) Mass-radius trajectories of hybrid stars with two hadron phase models. The U boson is included with various ratio parameters. \label{MR}} \end{figure} Now, we turn to the investigations of the hybrid star properties with these EOS's discussed above. Shown in Fig.~\ref{MR} is the relationship between the NS mass and radius for various EOS's with the inclusion of the U-boson. For the stiff hadron phase EOS, the hadron-quark transition at a high density does not reduce the maximum mass of hybrid stars significantly, as shown in the lower right panel in Fig.~\ref{MR}. While in most cases, the phase transition may results in very significant reduction of the maximum mass, which should not be consistent with the observation of large mass NS's~\cite{Demorest2010,Antoniadis2013}. We see that the inclusion of the U boson can remedy the mass eclipse very efficiently with the appropriately chosen ratio parameter. The rise of the maximum mass is consistent with the corresponding stiffening of the high-density EOS, shown in Fig.~\ref{eosnf}. We can see that the role of the U-boson in increasing the maximum mass is more significant for softer EOS's. Though shown in Fig.~\ref{MR} are only the results with the particular bag constant and transition density, it is sufficient to reveal the role of the U-boson since the change of these parameters yields rather similar results. It is, however, interesting to examine the consequences in the internal structure of the particular NS by varying these parameters. Let's first make comparison to the results with $B^{1/4}=180$ and 200 MeV. With $B^{1/4}=180 MeV$, the parameter set NL3 gives a NS maximum mass $1.39 M_\odot$. The size of the quark core is 5.76 km, the mixed phase spreads from 5.76 to 7.62 km, and the hadronic phase starts from 7.62 to 9.68 km. With $B^{1/4}=200 MeV$, the maximum mass is $1.86 M_\odot$. For convenient comparison, we also examine the internal structure of the $1.39 M_\odot$ NS with $B^{1/4}=200 MeV$. The size of the quark core is now 4.42 km, the mixed phase spreads from 4.42 to 6.38 km, and the hadronic phase starts from 6.38 to 10.59 km. This indicates that the smaller bag constant that gives rise to a smaller transition density features a larger quark core in a particular NS. For the case with the smaller transition density, the conclusion is similar since the smaller transition density is given by the smaller bag constant. It is also interesting to examine the effect of the U-boson on the NS structure, although the inclusion of the U-boson does not change the chemical and mechanical equilibriums, the transition densities, and the extents of each phases as well. As an example, here we examine the size of the quark core and of the mixed phase for the maximum mass configuration with the parameters $B^{1/4}=180 MeV$ and $(g_u/m_u)^2=25GeV^{-2}$. With the inclusion of the U-boson, the NS maximum mass increases from $1.39 M_\odot$ to $1.71 M_\odot$. It is found in the NS with the maximum mass configuration that the radius of pure quark core is 5.3km, which is smaller than 5.76 km obtained without the inclusion of the U-boson. The shell of the mixed phase extends from 5.3 to 8.6km, while without the inclusion of the U-boson, the shell starts from 5.76 to 7.6 km. These results indicate that the U-boson may change the ratio of various phases in the particular NS though it does not take part in the phase transition. \begin{figure}[thb] \begin{center} \vspace*{-5mm} \includegraphics[height=8.0cm,width=8.0cm]{gudd2} \end{center} \vspace*{-5mm}\caption{(Color online) Mass-radius trajectories of hybrid stars with the hybrid EOS initiated with the FSUGold. The transition density is set to be $3\r_0$. The dotted curve is the result with the density-dependent coupling constant of the U-boson, depicted in the inset. To obtain the density-dependent coupling, the mass of the U-boson is taken as $1MeV$. \label{gudd}} \end{figure} As shown in Fig.~\ref{MR}, the NS radius is significantly increased by the U-boson. It was pointed out in the literature~\cite{lat00,Li08} that the NS radius is primarily determined by the EOS in the lower density region of $1\rho_0$ to $2\rho_0$. Since the inclusion of the U-boson also increases the pressure in the lower density region appreciably, a large extent of the NS radius is obtained accordingly with various $c_u$'s of interest. This is similar to the previous works in Ref.~\cite{Wen09,zhang11}. It is known that the extraction of NS radii from observations still suffers from large systematic uncertainties~\cite{mil13} involved in the distance measurements and theoretical analyses of the light spectrum~\cite{lat01,ha01,zh07,su11}. There is yet no consensus on the extracted NS radii to date. For instance, using the thermal spectra from quiescent low-mass X-ray binaries (qLMXBs) Guillot and collaborators extracted NS radii of $R_{\rm NS} = 9.4\pm1.2$ km~\cite{gu14}, while a relevant study of spectroscopic radius measurements also suggest small radii $10.8^{+0.5}_{-0.4}$ km for a 1.5 $M_\odot$ NS~\cite{oz15}. There are also larger extracted radii: a 3-$\sigma$ lower limit of 11.1 km on the radius of the PSR J0437-4715~\cite{bog13} and a lower limit of 13 km for 4U 1608-52~\cite{pou14}. If the small radii of NS's is established, we need to reconsider the large NS radii produced by the U-boson. A possible solution of the NS radius suppression is to invoke the appropriate in-medium effects~\cite{ji15}. Considering that the NS radii are mainly determined by the low-density EOS, it is possible to reduce the NS radii by constructing the density-dependent coupling constant for the U-boson in the low-density regime. In the high-density regime, we neglect the in-medium effect for the U-boson, as one knows that the in-medium effect at high densities is suppressed greatly by the Pauli blocking~\cite{br92,ma94,ji15}. As an example, we perform the calculation with the EOS whose hadron phase part is obtained with the FSUGold and find that the appropriate density-dependent coupling constant of the U-boson can reduce the NS radii significantly, as shown in Fig.~\ref{gudd}. The density-dependent coupling constant, depicted in the inset of Fig.~\ref{gudd}, is designed to little change the pressure in the low-density regime~\cite{ji15}. While the energy density still acquires a significant increase from the U-boson, the resulting NS radius is even less than the one obtained with the pure hadron phase EOS. Since the NS maximum mass is determined mainly by the high-density EOS, the present form of the density dependence just gives rise to a moderate reduction of the NS maximum mass. \begin{figure}[thb] \begin{center} \vspace*{-5mm} \includegraphics[height=8.0cm,width=9.0cm]{MR2} \end{center} \vspace*{-5mm}\caption{(Color online) Mass-radius trajectories of hybrid stars with and without the QCD correction. To obtain the 2$M_\odot$ stars, the U-boson is also needed with the parameter denoted in the legend. \label{fmr2}} \end{figure} In Figs.~\ref{MR} and \ref{gudd}, we have not included the contribution of the QCD correction. The influence of the QCD correction on the NS maximum mass is rather sensitive to the transition density. Given large transition densities, the QCD correction just plays limited role in enhancing the NS maximum mass, since the maximum mass in this case is dominated by the hadron phase EOS. As the transition density decreases, the QCD correction brings more significant enhancement for the NS maximum mass. This findings are consistent with those in Ref.~\cite{Weissenborn11}. On the other hand, as we fix the bag constant, the situation for the QCD correction is different because the transition density itself in this case is increased by the correction through the phase equilibrium conditions. As shown in Fig.~\ref{fmr2}, the QCD correction with the fixed bag constant gives rise to a significant enhancement of the NS maximum mass, which is mainly attributed to the appreciable rise of the transition density. In this case, since the hadron phase EOS with the NL3 is much stiffer than that with the FSUGold, the large transition density due to the QCD correction leads to a more appreciable enhancement of the NS maximum mass with the NL3. With the inclusion of the U-boson, the NS maximum mass can further increase to satisfy the $2M_\odot$ constraint~\cite{Demorest2010,Antoniadis2013}. It is interesting to see that we just need a smaller ratio parameter of the U-boson to meet the maximum mass constraint in this case. This is more apparent for the EOS with the stiff NL3: the ratio parameter is largely reduced by the QCD correction, as shown in Fig.~\ref{fmr2} Since the magnitude and in-medium behavior of the U-boson ratio parameter are important for the NS properties, it remains significant to discuss the parameters for the U-boson. To satisfy the NS maximum mass constraint, the ratio parameter $c_u$ in this work is estimated to be around $0\sim 50GeV^{-2}$, depending on the hadron phase models chosen. To avoid the violation of the low-energy nuclear constraints for finite nuclei, we may limit the weak interaction strength of the U-boson with the mass being of order $1MeV$~\cite{zhang11,ji15}. For instance, if $g_u$ is 0.01, then the mass of the U-boson would be below $1.4MeV$, responsible for a long-range interaction. The weak interaction strength of the U-boson does not compromise the success of nuclear models in reproducing the properties of finite nuclei~\cite{xu13}. Interestingly, the inclusion of the QCD correction may reduce the ratio parameter significantly. We note that these estimated orders of magnitude for the U-boson parameters can be compatible with parameter regions allowed by a few experimental constraints~\cite{Kr09}. Moreover, the density-dependent parameter of the U-boson is found to be consistent with the usual predictions on the NS radius, while the density dependence would originate from the in-medium effect in the nuclear many-body system~\cite{ji15,br92,ma94}.

1607.01610

1607

1607.04035_arXiv.txt

{ Observations of the Galactic Center (GC) region in very-high-energy (VHE, $>\unit[100]{GeV}$) $\gamma$-rays, conducted with the High Energy Stereoscopic System (H.E.S.S.), led to the detection of an extended region of diffuse $\gamma$-ray emission in 2006. To date, the exact origin of this emission has remained unclear, % although a tight spatial correlation between the density distribution of the molecular material of the Central Molecular Zone (CMZ) and the morphology of the observed $\gamma$-ray excess points towards a hadronic production scenario.\newline% In this proceeding, we present a numerical study of the propagation of high-energy cosmic rays (CRs) through a turbulent environment such as the GC region. In our analysis, we derive an energy-dependent parametrization for the diffusion coefficient which we use for our simulation of the diffuse $\gamma$-ray emission at the GC. Assuming that hadronic CRs have been released by a single impulsive event at the center of our Galaxy, we probe the question whether or not the interaction processes of the diffusing hadrons with ambient matter can explain the observed diffuse $\gamma$-ray excess. Our results disfavor this scenario, as our analysis indicates that the diffusion process is, on timescales compared to the typical proton lifetime at the GC region, too slow to explain the extension of the observed emission. } \begin{document}

The GC region, harboring the supermassive black hole Sagittarius A* (Sgr A*) at its center, surrounded by massive clouds of dense molecular material, provides a unique opportunity for the study of non-thermal processes like the acceleration and the transport of highly energetic particles in our direct proximity. Besides the detection of the two point-sources HESS~J1745-290 and G0.9+0.1 \cite{Aharonian2004,Aharonian2005}, ongoing observations of this region with the H.E.S.S.\ instrument led to the discovery of an extended region of diffuse TeV $\gamma$-ray emission, spanning a range of \textasciitilde$2^{\circ}$ in Galactic longitude.~A spatial correlation between the density distribution of the molecular material of the CMZ and the morphology of this diffuse emission points to a hadronic production scenario: locally accelerated hadronic CRs diffusing out of the GC might create the observed diffuse $\gamma$-ray flux through interaction processes with the material in the clouds leading to subsequent $\pi^{0}\rightarrow \gamma\gamma$ decays. In the discovery paper \cite{2006Nature}, the H.E.S.S.\ collaboration argued that, assuming a mean diffusion coefficient of \textasciitilde $\unit[10^{30}]{cm^{2}/s}$, a single supernova explosion around $10^4$ years ago could have delivered enough energy to produce such a local population of highly energetic CRs, diffusing fast enough to produce a widely extended emission region comparable to the observed one. Further studies followed up on this idea, deriving diffusion coefficients by fitting the output of high-energy CR propagation models to the characteristics of the observed emission \cite{Dimi2009,Busching2007}, with results in agreement with the H.E.S.S. proposal. A fundamental different approach to the problem is given in \cite{Wommer2008}. Tracking individual particles under influence of the Lorentz force in turbulent magnetic fields, the outcome of this study suggests that particle diffusion is too slow to produce an emission region with an extension of more than a fraction of a degree at the GC, implying that diffusion coefficients should be much smaller than derived in previous studies. Hence, the conclusion of these authors is that only stochastic CR acceleration by magnetic turbulence throughout the GC region could produce the diffuse $\gamma$-ray excess - a scenario pursued in more detail e.g. in \cite{Fatuzzo2012}. Yet another approach assumes that a powerful wind, maintained by regular supernova explosions at the GC, advects particles out of the inner region, leading to a superposition of the diffusion-driven propagation with a ballistic component \cite{Crocker2011}. A general overview of the most recent $\gamma$-ray observations of the GC region and related topics, including a detailed discussion of the detected diffuse TeV $\gamma$-ray emission can be found in \cite{vanEldik}. In summary, there is still no agreement on the question of how highly energetic particles get injected in the GC region and which of the above discussed mechanisms explains best the diffuse radiation. In this proceeding, we focus on the scenario of a single impulsive injection of hadronic CRs at the center of our Galaxy, pursuing a preferably consistent treatment of the problem. From the tracking of ensembles of particles in turbulent magnetic fields, we derive a parametrization of the energy-dependence of the diffusion coefficient. The derived result is used as input parameter for our simulation of the diffuse $\gamma$-ray emission at the GC. Discretizing the diffusion equation with the help of finite differences, our simulation works on a discrete three-dimensional spatial grid, tracking proton distributions of defined energy in discrete time steps. Embedded in this environment is a three-dimensional map of the molecular material of the CMZ, which allows us to take interaction processes into account, % leading to TeV emission that we compare to the H.E.S.S.\ observations.

In this proceeding, we have presented our studies on the scenario of a production of the diffuse VHE $\gamma$-ray emission at the GC via interaction processes of diffusing protons with the ambient molecular matter of the CMZ. Within this scenario, we have assumed that a local population of highly energetic CRs had been released by a single impulsive ejection of a single power source located at the GC.\newline To this end, we have numerically analyzed the motion of protons in turbulent magnetic fields, applying the formalism introduced in \cite{Giacalone1994}, adjusted to the environmental conditions at the GC region. Tracking individual particles in a purely turbulent magnetic field by solving the Lorentz force equation, we derive an energy-dependent parametrization for the diffusion coefficient. In agreement with the analysis presented in \cite{Fatuzzo2010}, we observe two different dynamical regimes in the explored energy range. Modeling both regimes independently according to the usual assumption $D(E)\propto E^\delta$, we derive a value for $D_{10}$ below (but in agreement with) the typically assumed range of $D_{10}\,$\textasciitilde$\ \unit[10^{26}-10^{28}]{cm^2/s}$. The values we derive for the spectral index are $\delta=0.42$ in the low energy range, and $\delta=0.88$ in the high energy regime. We use this parametrization of $D(E)$ to simulate the diffuse $\gamma$-ray emission according to the described scenario. For a total propagation time of $10^4$~years, the excess region exhibits a spatial extension much smaller than the observed one, due to too slow diffusion of the underlying particle distribution. For significantly larger timescales, particle losses due to interaction processes during propagation lead to very high estimates of the total source energy (\textasciitilde$\unit[10^{55}]{erg}$), required to match the measured flux of the diffuse $\gamma$-ray excess.\newline In summary, we have found that diffusion of hadronic CRs away from a single point source at the GC appears to be too slow to explain the diffuse $\gamma$-ray excess measured by H.E.S.S., extending about $2^\circ$ in Galactic longitude. This result is in agreement with the analysis presented in \cite{Wommer2008}. Possibly, the $\gamma$-ray emission of the central source HESS J1745-290 itself might be associated with the addressed scenario. Besides the requirement of further refined simulation studies, the planned Cherenkov Telescope Array (see e.g.\ \cite{CTA2011}) will help to improve our understanding of the question which of the possible remaining scenarios, including stochastic particle acceleration as well as particle advection by a galactic wind outblow, might be the correct one to explain the current VHE $\gamma$-ray observations of the GC and its related particle physics. \newpage

1607.04035

1607

1607.03109_arXiv.txt

We present a configuration-space model of the large-scale galaxy 3-point correlation function (3PCF) based on leading-order perturbation theory and including redshift space distortions (RSD). This model should be useful in extracting distance-scale information from the 3PCF via the Baryon Acoustic Oscillation (BAO) method. We include the first redshift-space treatment of biasing by the baryon-dark matter relative velocity. Overall, on large scales the effect of RSD is primarily a renormalization of the 3PCF that is roughly independent of both physical scale and triangle opening angle; for our adopted $\Omega_{\rm m}$ and bias values, the rescaling is a factor of $\sim 1.8$. We also present an efficient scheme for computing 3PCF predictions from our model, important for allowing fast exploration of the space of cosmological parameters in future analyses.

The clustering of pairs of galaxies, quantified by the 2-point correlation function (2PCF), has a distinctive peak at roughly $100\Mpch$, an imprint of the Baryon Acoustic Oscillations (BAO) occurring prior to recombination (Sakharov 1966; Peebles \& Yu 1970; Sunyaev \& Zel'dovich 1970; Bond \& Efstathiou 1984, 1987; Holtzmann 1989; Hu \& Sugiyama 1996; Eisenstein \& Hu 1998; Eisenstein, Seo \& White 2007). The acoustic peak has been successfully used as a standard ruler to measure the relative size of the Universe at different redshifts as well as the absolute expansion rate (by comparison to the Cosmic Microwave Background), and these measurements in turn constrain the cosmological parameters, particularly dark energy (Eisenstein, Hu \& Tegmark 1998b; Blake \& Glazebrook 2003; Hu \& Haiman 2003; Linder 2003; Seo \& Eisenstein 2003; Weinberg et al. 2012, for a review; Cuesta et al. 2015 for the most recent measurement). Theoretical modeling indicates that the 3-point correlation function (3PCF) of galaxies should also contain BAO features (Sefusatti et al. 2006; Gil-Mar\'in et al. 2012; Slepian \& Eisenstein 2015a, Slepian et al. 2016a; hereafter SE15a and S16a), and indeed hints have been detected in Gazta\~naga et al. (2009) and S16a. If detected at high significance, BAO features in the 3PCF could be used as a standard ruler exactly as is already done with the 2PCF. Given the large spectroscopic datasets presently extant, such as BOSS (Eisenstein et al. 2011; Dawson et al. 2013), 6dFGS (Jones et al. 2009), and WiggleZ (Drinkwater et al. 2010), as well as the even-larger ones planned for the coming decade, such as those from eBOSS, Dark Energy Spectroscopic Instrument (DESI; Levi et al. 2013) and Wide-Field Infrared Survey Telescope (WFIRST; Spergel et al. 2013), it will be desirable to exploit the 3PCF as an additional tool for constraining the cosmic expansion history. However, the true positions of galaxies along our line-of-sight are unknown, and the redshift is used as a proxy, converted to a distance by assuming that the galaxies are comoving with the background Universe's expansion and have no peculiar velocities. This assumption is not perfectly correct. On small scales, galaxies have thermal peculiar velocities due to virialization within clusters. On larger scales, the peculiar velocities are coherent and generated by the growth of structure. Consequently, the positions of galaxies inferred from redshifts will be subject to redshift-space distortions (RSD; Hamilton 1998, for a review) unless these peculiar velocities are accounted for. In the 2PCF, RSD are modeled using the Kaiser formula (Kaiser 1987; Hamilton 1992), which rescales the 2PCF monopole by $1+2\beta/3 + \beta^2/5$, where $\beta = f/b_1$, with $f\approx \Omega_{\rm m}^{0.55}$ the logarithmic derivative of the linear growth rate with respect to scale factor and $b_1$ the linear bias. The Kaiser formula assumes a constant, single line of sight to the entire survey, which appears to be sufficiently accurate for the Sloan Digital Sky Survey (SDSS) volume (see Slepian \& Eisenstein 2015e and references therein for further discussion). Because the Kaiser factor is scale-independent, it does not shift the BAO scale in the 2PCF used to measure the expansion rate. The small-scale thermal velocities additionally alter the 2PCF but can be modeled as a Gaussian smoothing of the power spectrum, which also does not shift the BAO scale. For the 3PCF, there has as yet been no configuration space model from perturbation theory. Scoccimarro, Frieman \& Couchman (1999; hereafter SCF99) use Standard Perturbation Theory (SPT) working to second-order in the linear density field $\delta$ to obtain a model of the bispectrum (Fourier analog of the 3PCF). Rampf \& Wong (2012) show that the tree-level Lagrangian perturbation theory result, from using the Zel'dovich approximation, agrees with the SPT result of SCF99. There have been a number of models for the redshift-space bispectrum, both fully analytic (e.g. Hivon et al. 1995; Verde, Heavens \& Matarrese 1998; Smith, Sheth \& Scoccimarro 2008) and incorporating numerical simulations (e.g. Gil-Mar\'in et al. 2014). However only very few works consider the BAO. Sefusatti et al. (2006) focuses on joint analysis of the power spectrum and bispectrum and notes that BAO can break degeneracies (see their Figures 7, 9, and 10). Gil-Mar\'in et al. (2012) give a fitting formula for the dark matter bispectrum including BAO, and Gil-Mar\'in et al. (2014) includes RSD. The purpose of the present work is to develop a model of the 3PCF in configuration space in a form suitable for fitting the 3PCF of a large-scale redshift survey. First, we will convert the bispectrum model of SCF99 to configuration space. We will find that RSD essentially rescale the no-RSD 3PCF in a way that is roughly independent of both physical scale and triangle opening angle. This conclusion develops ideas first advanced in S16a and helps explain why the configuration-space model without RSD in that work was able to obtain a reasonable fit to the data. In the present work, we also develop a redshift-space model of the baryon-dark matter relative velocity effect, developing previous work on this term's signature in the 3PCF in real-space (SE15a). As a second goal of this paper, we will present a fast scheme for computing 3PCF predictions in the multipole basis first proposed in Szapudi (2004) and further developed in Slepian \& Eisenstein (2015a, b, c; hereafter SE15a, b, c). Typically perturbation theory expressions for the 3PCF $\zeta$ are written as cyclic sums over functions of pairs of sides and their enclosed angle, for instance in the form $\zeta\sim \xi(r_1)\xi(r_2) +{\rm cyc.}$, with $\xi$ the 2PCF (Groth \& Peebles' 1977 ``hierarchical ansatz''; see also Fry \& Peebles 1978; Davis \& Peebles 1977; Ma \& Fry 2000). Each term in the cyclic sum of such an expression corresponds to a different galaxy's contributing a particular bias term to the expectation value $\left< \delta_{\rm g}\delta_{\rm g}\delta_{\rm g}\right>$, as we further explain in \S\ref{sec:cyc_sum}. In reality, it is unknown which galaxy contributes which bias term, and one must cyclically sum so that all galaxies have a chance to contribute all the bias terms relevant at a given order in perturbation theory. Given two sides $r_1$ and $r_2$ and the cosine of their enclosed angle, $\hr_1\cdot \hr_2$, cyclic summing requires computing the third side and the two additional angles. This side and these angles depend on non-separable functions of $r_1,\;r_2$, and $\hr_1\cdot\hr_2$ and so their calculation scales as the number of grid points used for each side, $N_r$, times the number of grid points in angle cosine, $N_{\mu}$---that is, $N_r^2N_{\mu}$. Yet in the end we wish to bin the predictions in side lengths to a relatively modest number of bins, $N_{\rm bins}$, and also project the angular dependence onto Legendre polynomials. In this work we show how to do these operations first, meaning that the cyclic summing can be made to scale as $N_{\rm bins}^2$ for each multipole, for a total scaling as $N_{\rm bins}^2\ell_{\rm max}$ with $\ell_{\rm max}$ the maximal multipole. Computing the 3PCF predictions in the multipole basis using this scheme is consequently significantly more efficient. This efficiency will be important as the 3PCF becomes a standard tool for large-scale structure analyses and it becomes desirable to run a large grid of cosmological parameters through a prediction pipeline. The paper is laid out as follows. In \S\ref{sec:RSD_model}, we present the redshift-space bispectrum model of SCF99 and show how to cast it to configuration space. We then incorporate add tidal tensor biasing and briefly discuss other possible refinements to our model. In \S\ref{sec:cyc_sum}, we present the more efficient cyclic summing scheme summarized above. \S\ref{sec:disc} discusses our results after cyclic summing, and \S\ref{sec:rv_contrib} shows how to add relative velocity biasing in redshift space. \S\ref{sec:concs} concludes. Two Appendices showing mathematical results used in the main text follow \S\ref{sec:concs}. For all of the results displayed in this work, we have used transfer functions output from \textsc{CAMB} (Lewis 2000) with a geometrically flat $\Lambda{\rm CDM}$ cosmology with the following parameters: $\Omega_{\rm b}h^2 = 0.0220453,\;\Omega_{\rm c}h^2 = 0.119006,\;T_{\rm CMB} = 2.7255\;{\rm K}$, $n_{\rm s} = 0.9611$. These parameters match those used in S16a and do not differ substantially from the Planck values (Planck Paper XIII, 2015). Our $\sigma_8(z=0) = 0.8288$, and we rescale $\sigma_8$ by the ratio of the linear growth factor at the survey redshift to the linear growth factor at redshift zero. We take the survey redshift to be $z_{\rm survey} = 0.565$ so that our results are comparable to the CMASS galaxy sample discussed in S16a.

\label{sec:concs} In this work, we have shown how to transform the redshift-space bispectrum model of SCF99 into a configuration-space model suitable for fitting the 3PCF of current and upcoming large-scale structure redshift surveys. Prior to cyclic summing, our configuration model space model can be written entirely in terms of simple 1-D transforms of the linear theory matter power spectrum. We have also shown how to incorporate tidal tensor biasing in this framework, and developed a redshift-space model of the effect on the 3PCF of biasing by the baryon-dark matter relative velocity. Overall, a remarkably simple lesson emerges from this work: that the redshift-space 3PCF is to a fairly good approximation (of order $10\%$) simply a rescaling of the real-space 3PCF, where the rescaling factor is roughly independent of both triangle side length and multipole and depends on $\beta$ as well as the bias coefficients' values. In principle, the model presented in this work should allow constraints to be placed not only on the bias values $b_1,\;b_2,\;b_t$, and $b_v$, but on $\beta$ itself and thence the matter density parameter $\Omega_{\rm m}$, since $\beta = f/b_1 \approx \Omega_{\rm m}^{0.55}/b_1$. However, in practice $\beta$ is substantially degenerate with the linear bias. Future work will explore what precision the 3PCF can offer on $\beta$ as well as on the other bias parameters. In this regard it may be worthwhile to compute higher moments of the 3PCF with respect to the line of sight, as these moments are $\oO(\beta)$ at lowest order and consequently may offer more sensitivity on $\beta$. The mathematical techniques developed in SE15a and in this work are likely sufficient to transform the SCF99 expressions for the quadrupole moment of the 3PCF with respect to the line of sight; so doing is an avenue of possible future work. An addtional avenue of future work would be joint fitting of the 2PCF multipoles and the 3PCF. In particular, the square of the 2PCF monopole scales as $\sigma_8^4 b_1^4[1+4/3\beta+\oO(\beta^2)]$, while at leading order in $\beta$ the 3PCF scales as $\sigma_8^4 b_1^3 [1+4/3\beta+\oO(\beta^2)]$, ignoring the non-linear bias, which enters with a slightly different $\beta$ dependence but empirically is found to be small for e.g. the CMASS sample in SDSS DR12 (Slepian et al. 2016a). Thus the ratio of the 2PCF monopole's square to the 3PCF reveals $b_1[1+\oO(\beta^2)]$. This ratio offers a good starting point for measuring $b_1$ using the 2PCF and 3PCF by largely eliminating the leading-order $\beta$-dependence. This work also shows that the most significant BAO feature appears in the $\ell=1$ post-cyclic multipole of the 3PCF. There is additional BAO information in the $\ell=2$ multipole, but this BAO information comes from the pre-cyclic $L=1$ multipole's contribution to the post-cyclic $\ell=2$. The multipoles $\ell > 3$ look rather similar to $\ell=3$, but $\ell = 0,\;1,\;2$ and $3$ all look distinctively different from each other. There is also interest in removing a possible systematic shift in the BAO scale as measured from the 2PCF caused by galaxy biasing including a term in the baryon-dark matter relative velocity (Dalal, Pen \& Seljak 2010; Yoo, Dalal \& Seljak 2011; SE15a; Beutler et al. 2015; BMH16). Here we have for the first time presented the relative velocity's contribution to the redshift-space 3PCF, and shown that it is likely non-degenerate with the other terms entering at $\ell=1$. This work should enable the 3PCF's use in constraining the relative velocity bias and removing any systematic shift it may induce of the BAO bump in the 2PCF. We have already briefly discussed the treatments of RSD in other works modeling the 3PCF or bispectrum. As regards the purely analytic, SCF99 and Rampf \& Wong (2012) are the current state of the art; fortunately though they use different flavors of perturbation theory (SPT vs. 2LPT) as Rampf \& Wong show they agree at leading order. Rampf \& Wong (2012) also compute the one-loop correction to the redshift-space bispectrum (their equation (4.16)); it involves an exponential of the sum of Legendre polynomials of the cosine between the wave-vectors and the line of sight. We leave for future work consideration of whether the mathematical techniques developed here can be used to convert this expression into a configuration-space model in terms of simple 1-D integral transforms of the power spectrum. Recent observational works on the large-scale 3PCF have adopted different treatments of RSD, discussed in more detail in S16a \S5.2. S16a used the largest sample to date for a 3PCF measurement ($\sim 800,000$ LRGs in the SDSS DR12 CMASS sample), adopting a real-space bias model. The biases measured within this model were interpreted as effective quantitites reflecting the rescalinS16g RSD induce as well as the intrinsic biasing. Gazta\~naga et al. (2009), the only other observational work to access the BAO scale, measured the 3PCF divided by the Peebles \& Groth (1977) hierarchical ansatz, and noted that on large scales RSD cancel out of the 3PCF. Consequently they did not incorporate any additional modeling of RSD. Gil-Mar\'in et al. (2015) measured the bispectrum of a sample similar in size to that of S16a ($\sim 600,000$ LRGs in the SDSS DR11 sample). The starting point of that work was the bispectrum model of SCF99, but as further discussed in \S\ref{subsec:further_model}, they added additional parameters to reflect non-linear structure formation. These parameters are calibrated empirically from N-body simulations. Gil-Mar\'in et al. (2015) also incorporated tidal tensor biasing but assumed the tidal tensor bias follows the theoretical relation for local Lagrangian biasing and consequently is fully determined by $b_1$; thus they did not independently fit for $b_t$. Overall, there has not yet been a simple configuration-space model for the redshift-space 3PCF. The present work fills the gap, and we believe the template presented here will be of considerable utility in extracting BAO information from the 3PCF in current and future redshift surveys. Furthermore, we have developed a model for the relative velocity's contribution to the redshift-space 3PCF that allows a constraint on the relative velocity bias, important for ensuring the 2PCF remains an accurate avenue for measuring the cosmic expansion history. In a companion work (Slepian et al. 2016b) to this paper, we apply the templates presented here to the SDSS DR12 CMASS sample already used for the 3PCF measurement of S16a. We make a highly precise measurement of the linear bias as well as place constraints on the other bias parameters. Furthermore, we make the first high-significance detection of the BAO in the 3PCF and use it to measure the cosmic distance scale to redshift $z=0.57$ with $1.7\%$ precision. In a second companion paper (Slepian et al. 2016c), we use the relative velocity template presented here to constrain the relative velocity bias with $1\%$ precision, sufficient to imply that the BAO scale as measured from the SDSS DR12 2PCF is not systematically shifted. Future work might explore the utility of these 3PCF templates for even larger spectroscopic datasets, such as DESI (Levi et al. 2013).

1607.03109

1607

1607.06815_arXiv.txt

We calculate for the first time the phonon excitation rate in the outer crust of a neutron star due to scattering from light dark matter (LDM) particles gravitationally boosted into the star. We consider dark matter particles in the sub-GeV mass range scattering off a periodic array of nuclei through an effective scalar-vector interaction with nucleons. We find that LDM effects cause a modification of the net number of phonons in the lattice as compared to the standard thermal result. In addition, we estimate the contribution of LDM to the ion-ion thermal conductivity in the outer crust and find that it can be significantly enhanced at large densities. Our results imply that for magnetized neutron stars the LDM-enhanced global conductivity in the outer crust will tend to reduce the anisotropic heat conduction between perpendicular and parallel directions to the magnetic field.

Dark matter constitutes the most abundant type of matter in our universe and its density is now experimentally well-determined $\Omega_{CDM}h^2= 0.1199 \pm 0.0027$ \cite{cdm}. Worldwide efforts to constrain its nature and interactions have led the community to a puzzling situation where null results coexist with direct detection experiments that find high significance excesses \cite{dama}. In particular, in the low mass region of DM candidates i.e. $m_{\chi}<1\; \rm GeV/c^2$, cosmological, astrophysical and collider constraints seem to be the most important, see for example a discussion in \cite{lin}. Direct detection searches of thermalized galactic DM are mostly based on nuclear recoils on selected targets. In this scenario, LDM particles with masses much smaller than that of the nucleon, $m_\chi \ll m_N$, can only provide energies $\sim$ eV which are below the $\sim$keV threshold for conventional terrestrial searches. If one, instead, considers LDM scattering off bound electrons, energy transfer can cause excitation or even ionization and thus seems promising for exploring the phase space in a complementary way in the near future \cite{ele}. DM hitting terrestrial targets is expected to have low velocities $v_\chi \sim 10^{-3}$ (we use $c=\hbar=1$ units) as the gravitational boost is small for the Earth i.e. its Lorentz factor $\gamma_{\varoplus} =1/\sqrt{1-\beta^2_\varoplus} \sim 1$ with $\beta_\varoplus=\sqrt{\frac{2GM_\varoplus}{R_\varoplus}}$. However, for compact objects, such as neutron stars (NSs) with masses $M_{NS}\simeq 1.5 M_\odot$ and radius $R_{NS}\simeq 12$ km, $\gamma_{NS}\sim 1.26$ or $v_\chi \sim 0.7$ and thus provides a mechanism to boost particles to higher velocities or, accordingly, test the same length scales with smaller projectile masses. In particular, the outer crusts in NSs are formed by periodically arranged nuclei with typical densities ranging from $\rho \simeq 2 \,10^6-4\,10^{11}$ $\rm g/cm^3$. In the single-nucleus description \cite{chamel}, a series of nuclei with increasing baryonic number, $A$, from Fe to Kr form a lattice before neutrons start to leak out of nuclei. At these high densities, electrons form a degenerate Fermi sea. At even larger densities and up to nuclear saturation density, around $\rho_0\simeq n_0 m_N \simeq 2.4\,10^{14}$ $\rm g/cm^3$, a number of different nuclear structures called {\it pasta} phases appear \cite{pasta}. In this work we study the effect of LDM scattering in the production of quantized lattice vibrations (phonons) in the outer NS crust. Later, we will discuss how this result can impact subsequent quantities of interest, such as the ion thermal conductivity, that are relevant for computing the cooling behavior of NSs. Phonons are quantized vibrational modes characterized by a momentum $\vec{k}$ and polarization vector $\vec{\epsilon}_{\lambda}$ appearing in a nuclear periodic system \cite{ziman}. They can have a number of different sources. They can be excited due to non-zero temperature $T$ in the medium. The Debye temperature allows us to evaluate the importance of the ion motion quantization. For a bcc lattice \cite{carr}, for example, $T_D\simeq 0.45 T_p$, being $T_p=\omega_p /k_B=\sqrt{\frac{4 \pi n_AZ^2e^2}{k_B^2 m_A}}$ the plasma temperature associated to a medium of ions with number density $n_A$, baryonic number $A$, electric charge $Ze$ and mass $m_A$. $k_B$ is the Boltzmann constant. At low temperatures $T<T_D$, the quantization becomes increasingly important and the thermal phonons produced are typically acoustic modes, following a linear dispersion relation $\omega_{{k},\lambda}=c_{l, \lambda}|{\vec k}|$, where $c_l=\frac{\omega_p/3}{ (6\pi^2 n_A)^{1/3}}$ is the sound speed. In addition, phonon production can be caused by an external scattering agent, for example, standard model neutrinos. In this respect, weak probes such as cosmological neutrinos with densities $n_\nu \sim 116$ $\rm cm^{-3}$ per flavor have been shown to provide small phonon production rates in a crystal target \cite{ferreras}. Due to the tiny mass of the neutrino, the experimental signature of this effect seems however hard to confirm. The main interest in the astrophysical context we discuss in this contribution follows as phonon excitation in a periodic system, such as the outer crust of a NS, can affect the thermal transport coefficients in the star. The potential modification of transport properties of heat/energy in the external layers in NSs is crucial to possibly identifying relevant distortions in the cooling behavior of these astrophysical objects in rich LDM environments. The structure of this contribution is as follows. In section II, we present the effective field theory Lagrangian model using dark matter-nucleon contact interactions via scalar and vector couplings in a relativistic framework and compute the single phonon excitation rate, discussing sources of uncertainty. Later, in section III we compute the thermal conductivity in the outer crust with LDM contributions comparing the results to the standard thermal value and discussing possible astrophysical consequences. Finally, in Section IV we give our conclusions.

In conclusion, we have derived for the first time the single phonon excitation rate in the outer NS crust for relativistic LDM particles in the sub-GeV mass range. We have found that this rate is constant with the phonon momentum and much larger than for cosmological neutrinos at finite $|\vec{k}|$. A non-negligible correction to the local phonon excitation rate of $\sim 20\%$ is obtained when full relativistic phase distribution functions are considered for the incoming $\chi$ particles with respect to a monochromatic approximation, that under-predicts the result. As an astrophysical consequence of the previous, we have calculated the ion thermal conductivity in the dense and hot outer envelope founding that it can be largely enhanced at LDM densities in the maximum of the NS galactic distribution $n_\chi\sim 100 n_{0,\chi}$ due to a net modification of the acoustic phonon population. This effect is non negligible at densities beyond $\sim 3.5\, 10^{11}$ $\rm g/cm^3$ in the base of the outer crust at the level of standard ion-electron or thermal effects \cite{chugunov, pot}. We do not expect the degenerate electron contribution to largely modify this result as this would mildly screen nuclear charge in the lattice however it remains to be further studied. Although a detailed study of the quantitative effect in the surface temperature pattern remains to be undertaken, it is expected that for magnetized NSs the LDM-enhanced global enhancement of the perpendicular thermal conductivity allows a reduction of the difference of heat transport among parallel and perpendicular directions to the magnetid field. Based on previous works only including standard thermal contributions we expect that, as a natural consequence, the surface temperature profile would be more isotropic yielding flatter profiles for intermediate latitudes and remains to be calculated in a future contribution.

1607.06815

1607

1607.06481_arXiv.txt

We report the discovery of 14 Lyman-$\alpha$ blobs (LABs) at $z\sim0.3$, existing at least $4-7$ billion years later in the Universe than all other LABs known. Their optical diameters are $20-70$\,kpc, and \textit{GALEX} data imply \Lya luminosities of $(0.4-6.3)\times10^{43}$\,\ergs. Contrary to high-z LABs, they live in low-density areas. They are ionized by AGN, suggesting that cold accretion streams as a power source must deplete between $z=2$ and $z=0.3$. We also show that transient AGN naturally explain the ionization deficits observed in many LABs: Their \Lya and X-ray fluxes decorrelate below $\lesssim10^6$ years because of the delayed escape of resonantly scattering \Lya photons. High \Lya luminosities do not require \textit{currently} powerful AGN, independent of obscuration. \textit{Chandra} X-ray data reveal intrinsically weak AGN, confirming the luminous optical nebulae as impressive \textit{ionization echoes}. For the first time, we also report mid-infrared \textit{thermal echoes} from the dusty tori. We conclude that the AGN have faded by $3-4$ orders of magnitude within the last $10^{4-5}$ years, leaving fossil UV, optical and thermal radiation behind. The host galaxies belong to the group of previously discovered Green Bean galaxies (GBs). \textit{Gemini} optical imaging reveals smooth spheres, mergers, spectacular outflows and ionization cones. Because of their proximity and high flux densities, GBs are perfect targets to study AGN feedback, mode switching and the \Lya escape. The fully calibrated, coadded optical FITS images are publicly available.

{\label{intro} Lyman-$\alpha$ blobs (LABs) are extended \Lya nebulae with luminosities of $\llya=10^{42-44}$\,\ergs, populating the Universe at $z\gtrsim2$. They are often selected using optical narrow-band filters that isolate redshifted \Lya emission \citep[e.g.][]{ktk03,myh04,myh11,dbs05,ooe09,yzt09,yze10}. LABs are $20-200$\,kpc in size and show a bewildering range of properties. They can be associated with Lyman Break Galaxies \citep{sph95}, visible and obscured AGN, starburst sub-mm galaxies and passively evolving red galaxies \citep[e.g.][]{fwc01,csw04,myh04,gal09,wyh09}. LABs are landmarks of ongoing massive galaxy formation \citep{myh06,pkd08}, yet the ionizing sources in many of them remain mysterious. Our understanding of these processes would greatly benefit from studying the physical conditions in LABs. However, this is difficult as (1) cosmological surface brightness dimming reduces the flux densities by factors $\gtrsim100$, (2) the \Lya line is resonant, and (3) non-resonant optical lines are redshifted into and beyond the near-infrared atmospheric passbands. The resonant character of \Lya causes two main problems with the interpretation of \Lya data. First, \Lya photons scatter efficiently in space and frequency when propagating through a moving medium. Three-dimensional radiative transfer calculations \citep[e.g.][]{mer02,vsm06,kzd10} reveal a great variety of double-peaked \Lya line profiles emerging for various static and kinematic source/halo configurations \citep[for a one-dimensional analytic description in a static medium see][]{neu90}. It is difficult at best to infer the gas kinematics and the \Lya production sites from \Lya imaging and spectroscopy alone. A multi-wavelength perspective is required, including optically thin lines such as [\ion{O}{III}] and \Ha \citep[e.g.][]{sso08,wbg10,yzj11,yzj14,mcm14,znf15,svs15}. The second problem is that we need to understand the processes that govern how many \Lya photons manage to escape, so that they become observable at all. Dust, neutral hydrogen, metallicity and gas outflows control this escape fraction. The latter can range from less than 1 per cent to more than 50 per cent \citep[][and references therein]{ymg15}. Intrinsically, the \Han/\Lya line ratio is fixed for a photo-ionized nebula in equilibrium; the \Ha line can then be used to estimate the escape fraction and the total amount of \Lya produced. However, it is only for a small redshift window of $z=1.9-2.4$ that \textit{both} lines are observable from the ground. Worse, AGN variability may change the escape fraction further, by orders of magnitude, due to delayed \Lya escape \citep{rsf10,xwf11}. Probably the largest mystery with LABs is their frequent lack of ionizing sources; some LABs show no continuum counterparts at all. The lack of accessible diagnostic lines has prevented consistent conclusions on many occasions, and various processes have been suggested that could power LABs. For example, LABs are preferentially found in denser areas and filaments \citep{sso08,yze10,ebs11,myh11}, where LABs easily accrete cold neutral hydrogen from the cosmic web \citep{hsq00,dil09,gds10}. This is a requirement by $\Lambda$CDM galaxy formation models. The \Lya emission arises because of collisional excitation of hydrogen (virial temperature of $10^{4-5}$\,K) when it sinks into the dark matter haloes. It can contribute to an LAB's ionization over $10-30$ per cent of the Hubble time \citep{dil09}. Some calculations including self-shielding and realistic gas phases indicate that cold accretion alone could be insufficient to explain the \Lya fluxes of luminous LABs \citep{fkd10}. On the other hand, \citet{rob12} and \citet{cez13} find cold streams to be rather powerful, reproducing the size--luminosity function of observed LABs. The challenges in modelling the cold streams are mirrored on the observational side. One such gravitationally powered LAB \citep{nfm06} is questioned by \citet{pmb15}, who discovered an embedded obscured AGN and argue that the original data speak \textit{against} cold accretion. Alternatively, LABs can be shock-ionized by starburst-driven superwinds \citep{tas00}, and photoionized by obscured AGN or starbursts \citep[e.g.][]{cls2001,gal09}. Starbursts alone cannot explain \Lya equivalent widths (EW) higher than about 240\,\angstromblank \citep{mar02,sso08}; LABs often exceed this value. Another possibility is centrally produced \Lyan, resonantly scattered by neutral hydrogen in the circum-galactic medium \citep{las07,sbs11}. This leads to a characteristic polarization signal and can thus be distinguished from photo-ionization and shock heating which produce \Lya \textit{in situ} \citep{hss11,hvv13}. However, \citet{tvb14} find that similar polarization signals may arise in cold streams as well. Evidence for obscured AGN in some LABs has been found in infrared and sub-mm data \citep{bzs04,gal09,myh11,ond13,pmb15}, and for other LABs they have been postulated. For example, \citet{myh04} have found 35 LABs, some of which likely powered by superwinds and others by cooling flows \citep{myh06}. About 30 per cent lack UV continuum counterparts. Associated visible AGN are uncommon in this sample, and obscured star formation has been ruled out \citep{tmi13}. \citet{gal09} have shown that 24 out of 29 of these LABs remain undetected in a 400\,ks \textit{Chandra} exposure even after statistical stacking. They have suggested buried AGN and star-bursts as power sources instead of cold accretion. This, however, requires particular combinations of geometrical and radiative transfer effects to explain the substantial escape of \Lyan, while simultaneously preserving the thick obscuration along the line-of-sight \citep[see also][]{sas00}. While this certainly holds for some of these LABs, it seems unlikely to be the case for all of them. In this paper we investigate the effects of episodic AGN duty cycles (\textit{flickering}) on the UV, optical and mid-infrared (MIR) properties of LABs, forming optical \textit{ionization echoes} and MIR \textit{thermal echoes}. Ionization echoes have been reported before, mostly at lower redshift and in smaller and less luminous nebulae \citep{sev10,kls12,kcb12,sdh13,ssk13,kmb15}. We show that transient AGN naturally explain the ionization deficits in LABs. Our analysis is based on the \textit{Green Bean} galaxies \citep[GBs;][hereafter S13]{sdh13} at $z\sim0.3$, hosting luminous extended emission line regions (EELRs). We show that these EELRs are indeed LABs, and that they also host recently faded AGN. GBs form the most impressive ionization and thermal echoes currently known. This paper is structured as follows. In Section 2 we present an overview of the GBs, our \textit{Chandra} X-ray data, archival \textit{GALEX} data, and our Gemini/GMOS optical observations. The data are analysed in Sections 3, 4 and 5, respectively. In Section 6 we discuss the evidence for AGN flickering, and its effect on the UV, MIR and optical properties of LABs. In Section 7 we discuss the LAB size--luminosity function and the evolution of the LAB comoving density. Our summary and conclusions are presented in Section 8. Details about individual targets are given in Appendix \ref{targetnotes}. We assume a flat cosmology with $\Omega_m=0.27$, $\Omega_\Lambda=0.73$ and $H_0=70\,$\kms$\,{\rm Mpc}^{-1}$.

Green Bean galaxies (GBs, $z\sim0.3$) are spectacular and the most powerful emission-line objects known in the nearby Universe. Their extended emission-line regions (EELRs) measure between $20-70$\,kpc, and they are ionized by radio-weak type-2 quasars. GBs are extremely rare with a surface density of $1.1\times10^{-3}$ deg$^{-2}$; only 17 of them are known. They were selected photometrically from the 14500 deg$^2$ footprint of SDSS-DR8. In \citetalias{sdh13} we have suspected that the EELRs are ionization echoes, i.e. the AGN have faded recently and more quickly than the EELRs' typical light crossing times. GBs have not been investigated further apart from the single case study of \citet{dst15}, In this paper we have presented multi-wavelength observations to understand the unusual nature of GBs and to view them in the context of galaxy evolution. With \textit{Chandra} we have probed the current activity of the AGN and the intrinsic obscuration along the line of sight toward 10 GBs. With \textit{GALEX} archival images we have estimated the far-UV properties of 15 GBs, and we have obtained high resolution optical images in $gri$ filters for all objects. \subsection{Main results} \begin{enumerate} \item{\textit{Chandra} has revealed low counts, moderate hardness ratios and weak or absent K$\alpha$ lines. This implies that these AGN are intrinsically weak rather than Compton-thick. The EELRs' high [\ion{O}{III}] luminosities require recent and rapid fading of the AGN, confirming them as ionization echoes (Section \ref{fadingagn}).} \item{Strongly variable AGN do not follow the MIR X-ray relation. The MIR response of a dusty torus can be delayed by up to $\sim10^3$ years \citep{hok11}, forming a thermal echo. Accordingly, the AGN in the GBs must have faded by several factors $10-100$ over the last $100-1000$ years (Section \ref{argmirx}). Combining the thermal and ionization echoes, the AGN must have faded by $3-4$ orders of magnitude within the last $10^{4-5}$ years. This rate is similar to that observed in the \textit{changing look} quasars \citep[][also attributed to change of accretion rates]{lcm15,rcr16}; however, in GBs it could be sustained over much longer periods of time.} \item{\textit{GALEX} FUV images require that at least 85 per cent of the GBs have \Lya luminosities in excess of $10^{43}$\,\ergs. They form Lyman-$\alpha$ blobs or LABs (Section \ref{GALEXresults} and \ref{discussion1}). We have proven that LABs still exist in the Universe $4-7$ billion years later than previously known. Ultimately, the \Lya emission has to be confirmed with FUV spectroscopy using \textit{HST} or \textit{Astrosat} \citep{hut14}.} \item{We propose rapid duty cycles (AGN flickering) as a natural explanation for the mysterious ionization deficits observed in LABs. Resonant \Lya photons are efficiently stored in LABs and only gradually released, \textit{decorrelating} from AGN variability on scales of up to $10^6$ years depending on the optical depth. An AGN may undergo several duty cycles before \Lya escapes; a luminous LAB does not require a \textit{currently} powerful AGN, independent of obscuration (Section \ref{discSecLYA}). This does not mean that we can relinquish e.g. cold accretion, star formation and shocks as ionization sources. It means that multi-wavelength observations are required to identify the ionizing source(s) of individual LABs.} \item{Low-z LABs live mostly in isolation or in low density environments (Section \ref{longslit}), whereas high-z LABs are preferentially found in massive structures. Sometime between $z=2$ and $z=0.3$ these structures must lose their ability to form and sustain LABs, probably because cold accretion streams have depleted. AGN survive, and may continue powering LABs at low redshift (Section \ref{agnoutflows}).} \item{Our comoving volume is 3.9 Gpc$^3$, $100-1000$ times larger than other LAB surveys, second only to \cite{bbb13}. The density at $z\sim0.3$ is $\rho_c=3.3\pm0.9$\,Gpc$^{-3}$ for LABs with $\llya\geq10^{43}$\,\ergs\; and $D\geq20$\,kpc. The density evolves with $\rho_c\propto(1+z)^{3-4}$ for both clusters and in the field between $z=0.3$ and $z=2$. During this time, $99-99.9$ per cent of the LABs disappear. A more accurate determination of the evolution requires (1) better sampling between $z=0.5-1.5$, and (2) volumes of $0.1-1$\,Gpc$^3$ to overcome cosmic variance and to be independent of environment. Otherwise, densities are not directly comparable (Section \ref{clusterdensity}).} \item{LABs with different ionization sources evolve differently. Gravitationally powered LABs do not survive redshift evolution as the cold accretion streams are depleting. We find only one LAB, J1155$-$0147 ($z=0.306$), that could still be powered by cold streams in addition to an AGN. The density of LABs powered solely by AGN is evolving much slower, if at all, depending on the AGN's intrinsic X-ray luminosity (Sects. \ref{diffevolution} and \ref{bridge}).} \end{enumerate} \subsection{Conclusion and future observations} LABs should be considered as efficient \Lya photon stores, albeit with a badly maintained inventory. They trap \Lya photons for a time much longer than their light crossing time. When the thermalised photons are eventually released in a gradual manner, all memory about the location, spectrum and time variability of their source has been lost. Typical storage times are short compared to the life times of cooling flows, star bursts, and shock ionization populating the store with photons. However, the storage times become long compared to episodic AGN bursts and their transition from high to low states (and vice versa). Even if optical and MIR observations point to the most powerful AGN, the absence of the latter in X-rays does not necessarily mean Compton-thick obscuration. The AGN timescale is much shorter than the typical galaxy time-scale and the galaxies response time to AGN activity; care must be taken when seeking or applying relations between AGN and other ``galactic'' observables, in particular if the AGN are suspected to be transient \citep[see also][]{hma14}. The GBs and the associated LABs published in this paper are \textit{much} easier to study than their high-z counterparts, as low fluxes, redshift incompleteness, and physical resolution are not a problem. They are perfectly suited to study AGN feedback, large-scale outflows, quasar duty cycles, mode switching, and the \Lya escape fraction, the latter being controlled by an interplay of geometry and velocity field, metallicity, hydrogen density, dust obscuration and AGN variability. Our own \textit{NUSTAR} observations have commenced in cycle 2 for the two X-ray brightest GBs, J1155$-$0147 and J0113+0106, to accurately determine the obscuration and the actual shut-down ``depth'' of their AGN. Time has also been awarded for an initial survey with the \textit{Hubble Space Telescope}, to determine the FUV properties of the same two GBs, and J2240$-$0927 \citep{dst15}, using ACS/SBC imaging and spectroscopy.

1607.06481

1607

1607.04167_arXiv.txt

Timing results for the black-widow pulsar J2051$-$0827 are presented, using a 21-year dataset from four European Pulsar Timing Array telescopes and the Parkes radio telescope. This dataset, which is the longest published to date for a black-widow system, allows for an improved analysis that addresses previously unknown biases. While secular variations, as identified in previous analyses, are recovered, short-term variations are detected for the first time. Concurrently, a significant decrease of $\sim2.5 \times 10^{-3}\ \rm{cm^{-3}pc}$ in the dispersion measure associated with \psr{} is measured for the first time and improvements are also made to estimates of the proper motion. Finally, \psr{} is shown to have entered a relatively stable state suggesting the possibility of its eventual inclusion in pulsar timing arrays.

\label{sec:intro} Of the $\sim$2600 pulsars known today, roughly 10\% appear to have rotation-periods of the order of a few milliseconds and are known as millisecond pulsars (MSPs). Within the MSP population there exist a variety of configurations, however, most MSPs are found in binary systems. Among these, about 10\% are in tight, eclipsing binaries. Such systems are further classified into the black-widow systems, with very light companions of mass ($\dot{m}_{\rm c}\lesssim 0.05\ M_{\odot}$) and redback systems, with heavier companions \citep[ $ 0.1\ M_{\odot}\ \lesssim \dot{m}_{\rm c} \lesssim 0.5\ M_{\odot}$;][]{r13,CCT+13}. \psr{} is the second black-widow system that was discovered \citep{sbb96}. Its companion is expected to be a $\sim$0.02-0.06 $M_{\odot}$ star, whose exact nature is yet to be determined \citep[see discussions in][]{svbk01,lvt+11}. Pulsar timing relies on making highly precise measurements of the time at which the radio-beam from a rotating pulsar crosses a radio telescope. These measured times are then compared to a theoretical prediction of these crossing events to derive various properties of the pulsar. A more extensive discussion on pulsar timing and the benefits of MSPs for pulsar timing can be found in \cite{lk05} and other reviews of pulsar timing. MSPs are particularly well-suited for this because of their inherent stability and short rotation periods. Even though the pulsars in black-widow systems are MSPs, they are typically excluded from high-precision pulsar timing experiments since several of them have been observed to display variability in their orbital parameters, in particular the orbital period. This variability may be due to many reasons like the interaction of the pulsar with the companion, the presence of excess gas around the companion's orbit or the companion's mass loss. However, only a limited number of studies so far have tried to identify if the variability of such pulsars can be modelled by introducing new parameters into the pre-existing timing models or by defining new timing models for such systems. Given the recent increase in the number of MSPs detected, in large part from surveys of Fermi-LAT sources \citep{2FGLCat}, and the rapid growth in the sensitivity and bandwidth of modern digital receiver systems for pulsar timing making it possible to detect variations in much greater detail, it is pertinent to address this long-standing question. \psr{} has been continuously timed since its discovery in 1995 \citep{sbb96} and therefore the dataset presented in the following analysis represents the longest timing baseline currently published for eclipsing black-widow systems. Given this long time-baseline and other favourable properties discussed in the following sections, this dataset offers an ideal opportunity to attempt such an exercise. Previous pulsar timing analyses of \psr{} have shown that the orbital period, $P_{\rm{b}}$, and projected semi-major axis, $x$, undergo secular variations \citep{sbm+98, dlk+01, lvt+11}. These variations are possibly linked to the variations of the gravitational quadrupole moment of the companion and induced by variations of the mass quadrupole of the companion as its oblateness varies due to rotational effects \citep{lvt+11}. These variations may arise due to a differential rotation of the outer layers of the companion \citep{as94} or due to variations in the activity of the magnetic field of the companion as in the \cite{lr01} model. Similar variations have been measured for a few other pulsars in BW systems like PSR J1959+2048 \citep[PSR B1957+20;][]{fst88}, PSRs J0024$-$7204J and PSRs J0024$-$7204O \citep[47 Tuc J and O;][]{fck+03}, PSR J1807$-$2459A \citep[NGC 6544A;][]{LFR+12} and PSR J1731$-$1847 \citep{NBB+14}. The binary system containing \psr{} has also been recently detected in Fermi-LAT and {\it Chandra}/ACIS data \citep{wkh+12}. The $\gamma$-ray luminosity is $7.66\times10^{32}\ \rm{erg\ s^{-1}}$. The inferred spin-down power, ${\dot{E}}$, from radio observations is $\sim 5.49 \times 10^{33}\ \rm{erg\ s^{-1}}$. The $\gamma$-ray luminosity, therefore, represents $\sim 15\%$ of the total spin-down power, which is consistent with other MSPs for which such a detection has been made. The $\gamma$-ray emission from the system appears to be well fit by a model of emission in the `outer gap accelerator', as discussed in \cite{tct12}. Using the new ephemerides presented here, it may be possible to detect the orbital dependence of pulsed emission from \psr{}. The $X$-ray luminosity is $1.01\times10^{30}\ \rm{erg\ s^{-1}}$ \citep{wkh+12} and the data do not present any evidence for bursts, which suggests that the companion is stable and does not undergo sudden deformations. The flux values fit well for a model with emission from the intra-binary shock, the polar caps and synchrotron emission from the pulsar magnetosphere \citep{wkh+12}. This work provides an update on the timing of \psr{} and presents an improved analysis. Two complementary timing models for \psr{} are provided, one capable of handling small eccentricities and another, utilising orbital-frequency derivatives. A new method for measuring the variations in the orbital period, $\Delta P_{\rm{b}}$, by measuring the change in the epoch of ascending node, $T_{\rm{asc}}$ is also presented.

1607.04167

1607

1607.04398_arXiv.txt

The properties of $\isotope[12]{C}$, $\isotope[16]{O}$, and $\isotope[20]{Ne}$ nuclei in strong magnetic fields $B\simeq 10^{17}\,$G are studied in the context of strongly magnetized neutron stars and white dwarfs. The \sky code is extended to incorporate the interaction of nucleons with the magnetic field and is utilized to solve the time-independent Hartree-Fock equations with a Skyrme interaction on a Cartesian three-dimensional grid. The numerical solutions demonstrate a number of phenomena, which include a splitting of the energy levels of spin-up and -down nucleons, spontaneous rearrangment of energy levels in $\isotope[16]{O}$ at a critical field, which leads to jump-like increases of magnetization and proton current in this nucleus, and evolution of the intrinsically deformed $\isotope[20]{Ne}$ nucleus towards a more spherical shape under increasing field strength. Many of the numerical features can be understood within a simple analytical model based on the occupation by the nucleons of the lowest states of the harmonic oscillator in a magnetic field.

\label{sec:intro} The studies of nuclei and bulk nuclear matter in strong magnetic fields are motivated by the astrophysics of strongly magnetized neutron stars (magnetars) and white dwarfs. The surface fields of magnetars have been inferred to be from observations in the range of $10^{15}$~G. The interior fields of magnetars can be several orders of magnitude larger than their surface fields~\cite{2015RPPh...78k6901T}. The electromagnetic energy of the interaction of baryons with the $B$ field becomes of the order of the nuclear scale $\sim$ MeV for fields $10^{16}-10^{17}$~G and can arise from current-field (for electrically charged particles) and spin-field (for charge neutral particle) interaction. Such interaction can affect the properties of nuclei, including their shell structure, binding energies, and rms radii. This, in turn, can affect the structure and composition of the interiors of neutron stars and white dwarfs where nuclei are predicted to exist, as well as the transport and weak interaction (neutrino emission and absorption) processes due to the changes in charged particle dynamics and transition matrix elements. The equation of state and composition of inhomogeneous nuclear matter featuring nuclei in strong magnetic fields have been studied using various methods including modification of the Thomas-Fermi model ~\cite{1989ApJ...342..958F}, liquid drop model~\cite{1991ApJ...383..745L}, nuclear shell Nilsson model~\cite{2000PhRvL..84.1086K,2001ApJ...546.1137K}, relativistic density functional theory~\cite{2011PhRvC..84d5806P,2015PhRvC..92c5802B}, and non-relativistic Skyrme functional theory~\cite{2012PhRvC..86e5804C}. It has been shown that bulk properties of nuclei and shell structure as well as their shape can be significantly affected if the magnetic field is of the order of $10^{17}$~G and larger. These studies were carried out in the context of neutron star crusts and have concentrated on heavy nuclei beyond (and including) $^{56}$Fe. In this work we consider the properties of carbon-oxygen-neon mass nuclei in strong magnetic fields. Our motivation for doing so is threefold. Firstly, some isolated neutron stars are known to have atmospheres composed of $\isotope[12]{C}$, as is well established in the case of the compact object in Cas A~\cite{2013ApJ...779..186P}. Carbon plays also an important role in the physics of accreting neutron stars, for example, in the superbursts which are associated by unstable ignition of carbon at depth characterized by the density $\rho\simeq 10^9$ g\, cm$^{-3}$~\cite{2014ApJ...791..106S}. Apart from $\isotope[12]{C}$, there are also substantial fractions of oxygen and neon mass nuclei produced in the crusts of accreting neutron stars through nuclear reaction networks~(Ref.~\cite{2014ApJ...791..106S} Table I and Figs. 4 and 5). The composition of the surfaces of magnetars and their properties under accretion are not known. However, one can anticipate the role played by these nuclei in the physics of magnetars, by extrapolating from the physics of low-field neutron stars. Secondly, white dwarfs models composed of $\isotope[12]{C}$, $\isotope[16]{O}$, or $\isotope[20]{Ne}$ are standard in the physics of these compact objects in the non-magnetized regime. The superluminous type-I supernovae were suggested recently to originate from supermassive strongly magnetized white dwarfs~\cite{2013IJMPD..2242004D} with magnetic fields in the range $10^{17}$~G. If such objects exist (for a discussion see~\cite{2014PhRvD..90d3002C}) they would provide the environment where light nuclei would be subjected to intense $B$ fields. Thirdly, besides the astrophysical motivation, there is a technical aspect to our study, as it is a first attempt to include magnetic fields in the widely used \sky code~\cite{2014CoPhC.185.2195M}. Therefore, another motivation of our study is to provide a benchmark for future studies that will include magnetic fields on simple enough systems that are easily tractable, and $\isotope[12]{C}$, $\isotope[16]{O}$, or $\isotope[20]{Ne}$ nuclei are optimal in this respect. As we show below, the splitting of the levels induced by the $B$ field in these nuclei can be easily understood on the basis of a harmonic oscillator model; because of the small number of levels the comparision between the numerical and analytical results becomes possible. (Indeed, in heavy nuclei the number of levels can become very large, which would make such comparison impractical). The density-functional-based Hartree-Fock (HF) theory provides an accurate and flexible framework to study a variety of low-energy nuclear phenomena. The public domain \sky code~\cite{2014CoPhC.185.2195M}, solves the HF equations on a 3D grid (i.e. without any assumptions about the underlying symmetries of the nuclei) and is based on the Skyrme density functional. It has already been utilized to study a broad range of problems ranging from low-energy heavy-ion collisions to nuclear structure to exotic shapes in crusts of neutron stars (for references see Ref.~\cite{2014CoPhC.185.2195M}). In this work we report on the first implementation of a strong magnetic field in the \sky code via extension of the Hamiltonian (and the associated density functional) to include all relevant terms reflecting the interaction of the magnetic field with nucleons. We concentrate on the static solutions, which requires the solution of time-independent Hartree-Fock equations. As an initial application we report convergent studies of the carbon-oxygen-neon mass range nuclei in strong fields. This paper is organized as follows. In Sec.~\ref{sec:theory} we describe briefly the underlying theory and the modifications to the numerical code needed to include $B$ fields. We present our results in Sec.~\ref{sec:results} where we first set up a simple analytical model which is then compared with numerical results. A number of observables such as energy levels, spin- and current-densities, as well as deformations are discussed. Our conclusions are summarized in Sec.~\ref{sec:conclusions}.

\label{sec:conclusions} We have performed numerical computations of $\isotope[12]{C}$, $\isotope[16]{O}$, and $\isotope[20]{Ne}$ nuclei in strong magnetic fields using an extension of the \sky code, which solves Hartree-Fock equations on a three-dimensional grid in a strong magnetic field. The code is based on the Skyrme density functional. Common features found for all three nuclei are (i) the splitting of energy states, which are on the order of MeV for fields $B\sim 10^{17}\,$G; (ii) the increase in spin polarization along the magnetic field as the field is increased, which is characterized by a transition from a regime where $l$-$s$ coupling is dominant to a regime where $l$ and $s$ couple directly to the magnetic field; and (iii) an increase in the flow-velocity in the plane orthogonal to the field with increasing magnetic fields. A number of features are peculiar to specific nuclei and are listed below: \begin{itemize} \item A rearrangement of energy levels in $\isotope[16]{O}$ nucleus is observed at a critical field $4.1\times 10^{17}\,$G, which is accompanied by an abrupt increase in the magnetization of the nucleus and an increase in the velocity flow. This also leads to deformation of the nucleus from its original spherical shape at vanishing value of the field. \item The $\isotope[12]{C}$ nucleus does not change its shape in the magnetic field and there are no energy level rearrangements as seen in $\isotope[16]{O}$. It is found to be more easily polarizable than the heavier nuclei. \item The $\isotope[20]{Ne}$ nucleus is deformed in the ground state. It undergoes significant continuous change in its shape as the magnetic field is increased. The deformation is diminished by a factor of 2 for fields $B\simeq 4.1\times 10^{17}\,$G. \item We have shown that a simple analytical model which fills in the harmonic oscillator states in the magnetic field accounts well for the energy splitting of $\isotope[12]{C}$ and $\isotope[16]{O}$ nuclei as well as their angular momenta and spin projections. In the case of $\isotope[20]{Ne}$ the analytical model is less reliable; it can reproduce qualitatively features obtained with the \sky code; however, because of the nuclear deformation, the magnetic field needs to be directed along the $x$-axis instead of the $z$-axis. \end{itemize} Phenomenologically the most important aspect of these findings is the splitting of the levels in nuclei as a function of the $B$ field. When this splitting is on the order of the temperature it will have an important impact on all transport processes and on neutrino and photon emission and absorption, as well as on the reaction rates. Looking ahead, we would like to extend these studies to nuclei with larger mass numbers and beyond the stability valley in the direction of neutron-rich nuclei that occur in nonaccreting neutron stars~\cite{PhysRevLett.110.041101}. The possibility of non-spherical and extended nuclei (pasta phases~\cite{2016PhRvC..93e5801S}) can also be considered in this context.

1607.04398

1607

1607.04351_arXiv.txt

{ The nearby red dwarf binary GJ65\,AB (UV+BL\,Ceti, M5.5Ve+M6Ve) is a cornerstone system to probe the physics of very low-mass stars. The radii of the two stars are currently known only from indirect photometric estimates, however, and this prevents us from using GJ65\,AB as calibrators for the mass-radius (M--R) relation. We present new interferometric measurements of the angular diameters of the two components of GJ65 with the VLTI/PIONIER instrument in the near-infrared $H$ band: $\theta_\mathrm{UD}(\mathrm{A}) = \thetaA \pm \thetaAerr$\,mas and $\theta_\mathrm{UD}(\mathrm{B}) = \thetaB \pm \thetaBerr$\,mas. They translate into limb-darkened angular diameters of $\theta_\mathrm{LD}(\mathrm{A}) = \thetaLDA \pm \thetaLDAerr$\,mas and $\theta_\mathrm{LD}(\mathrm{B}) = \thetaLDB \pm \thetaLDBerr$\,mas. Based on the known parallax, the linear radii are $R({\rm A}) = \radA \pm \radAerr\ R_\odot$ and $R({\rm B}) = \radB \pm \radBerr\ R_\odot$ ($\sigma(R)/R = 4\%$). We searched for the signature of flares and faint companions in the interferometric visibilities and closure phases, but we did not identify any significant signal. We also observed GJ65 with the VLT/NACO adaptive optics and refined the orbital parameters and infrared magnitudes of the system. We derived masses for the two components of $m(\mathrm{A}) = \MA \pm \MAerr\,M_\odot$ and $m(\mathrm{B}) = \MB \pm \MBerr\,M_\odot$ ($\sigma(m)/m = 4\%$). To derive the radial and rotational velocities of the two stars as well as their relative metallicity with respect to Proxima, we also present new individual UVES high-resolution spectra of the two components. Placing GJ65\,A and B in the mass-radius diagram shows that their radii exceed expectations from recent models by $14 \pm 4\%$ and $12 \pm 4\%$ , respectively. Following previous theories, we propose that this discrepancy is caused by the inhibition of convective energy transport by a strong internal magnetic field generated by dynamo effect in these two fast-rotating stars. A comparison with the almost identical twin Proxima, which is rotating slowly, strengthens this hypothesis because the radius of Proxima does not appear to be inflated compared to models.}

Stellar structure models have traditionally been found to systematically underestimate the radii of very low-mass stars (VLMS) of a given mass by 5 to 15\% \citepads{2002ApJ...567.1140T, 2010ApJ...718..502M}. To explain the discrepancy between models and observations, \citetads{2007A&A...472L..17C} proposed that a decrease in convection efficiency induced by strong magnetic fields can inflate the stellar radius by amounts that qualitatively match the observed differences. Recent progress has been reported by \citetads{2014ApJ...789...53F} using the Dartmouth models \citepads{2008ApJS..178...89D}. The differences between models and eclipsing binary observations are now lower than in the past,~within a few percent for \object{CM Dra} \citepads{2014A&A...571A..70F}, for instance, and the average discrepancy is around 3\% \citepads{2013ApJ...776...87S}. However, there are still only very few test stars for these models around $0.1\,M_\odot$, at the limit of the brown dwarf regime. A thorough review of the field of low-mass star evolution modeling can be found in \citetads{2015ASPC..496..137F}. To validate the VLMS structure models, we thus need radius measurements at the low-mass end of the stellar mass-radius (M--R) function. As shown by \citetads{2003A&A...397L...5S}, \citetads{2009A&A...505..205D}, and \citetads{2012ApJ...757..112B}, for example, long-baseline interferometry is well suited for this task because it provides a simple way to measure the angular diameter of non-eclipsing stars. Unfortunately, the number of VLMS accessible to interferometric measurements is limited by their very small physical radius. With the current optical and infrared interferometry instrumentation, the largest distance at which red dwarfs can be resolved angularly is only about 3\,pc. For this reason, eclipsing VLMS binaries are also used as calibrators for the M--R relation, but as a result of the proximity of the stars, the gravitational and magnetic interactions can affect their physical properties. \citetads{2013ApJ...776...87S} found a comparable discrepancy level on average of observed vs.~model radii for single and binary VLMS, but the components of short-period binaries were found to be the most deviant. Observations of single stars and well-separated binaries is therefore very desirable to calibrate the M--R relation. Its wide physical separation (larger than 4\,AU for most of the orbit) and proximity ($d = 2.7$\,pc) make the \object{Gliese 65} system (\object{GJ65} AB, \object{Luyten 726-8}, \object{BL Cet}+\object{UV Cet}, \object{WDS J01388-1758AB}, \object{2MASS J01390120-1757026}, \object{LDS 838}) a cornerstone system on which to calibrate the M--R relation. This binary consists of two main-sequence red dwarfs of spectral types M5.5Ve and M6Ve \citepads{1994AJ....108.1437H}. Its relatively wide separation is sufficient to neglect gravitational and magnetic interactions, but the two stars are close enough so that the period is relatively short ($P=26.3$\,years). Accurate orbital parameters and masses can therefore be determined in a reasonable time frame. These ideal properties make GJ65\,AB a rare and very favorable configuration among the known VLMS systems. As a side remark, GJ65\,AB will pass within one\,light-year of $\epsilon$\,Eri in about 30\,000\,years \citepads{2010arXiv1004.1557P}, possibly interacting with the hypothetical Oort cloud of this star. We present our new observations in Sect.~\ref{observations}. We used the VLTI/PIONIER instrument to measure the angular diameters of the two components of GJ65 (Sect.~\ref{pionier_obs}), and we combined them with the well-known parallax of the system to derive their linear radius. We also observed GJ65 using NACO (Sect.~\ref{naco_obs}), from which we confirm that a revision of the orbital parameters is necessary. In Sect.~\ref{uves_data} we present high-resolution spectra of the two stars. The determination of the parameters of GJ65 A and B from our observations is described in Sect.~\ref{analysis}. We derive an improved orbital solution in Sect.~\ref{orbital-solution}, including a revised estimate of the mass of the system. We finally compare GJ65\,AB with \object{Proxima} in Sect.~\ref{proxima}. The fundamental parameters of these three stars are very similar, and we discuss the possible reasons for the discrepancy in their measured radii.

We presented the first interferometric measurements of the angular diameters of the two components of the nearby red dwarf binary GJ65. We also obtained new high-accuracy adaptive optics astrometry and infrared photometry, as well as separate high-resolution spectra of the two stars. The latter allowed us to derive the differential radial velocity of GJ65 A and B, estimate their rotational velocity $v \sin i,$ and their metallicity with respect to Proxima taken as fiducial. Based on our new observations, we present refined orbital elements, an accurate value of the total and individual masses of the two components, and of their linear radii. The positions of GJ65 A \& B in the mass-radius diagram confirms that their radii are underestimated by the current stellar structure models by approximately $13 \pm 4\%$. Following \citetads{2007A&A...472L..17C}, we propose that the enlargement of their radii is caused by the inhibition of convection by their magnetic fields, generated through dynamo effects by their fast rotation. This radius inflation is not observed for Proxima, which has almost identical fundamental parameters (in particular the mass) and very similar X-ray activity, but exhibits a slow rotational velocity. Encouragingly, \citetads{2012ApJ...757...42F} and \citetads{2013ApJ...776...87S} showed that the current generation of VLMS models better agree with observational mass and radius determinations than in the past. Further progress might be achieved by an improved modeling of the internal magnetic field in fast-rotating fully convective stars. The complexity of the corresponding simulations represents a very difficult challenge, however. The availability of three VLMS with mostly identical physical properties (GJ65\,AB and Proxima) and differing only in their rotational velocity will be potentially of extremely high value to test their predictions.

1607.04351

1607

1607.07487_arXiv.txt

{Recent timing observation constrained the braking index of the X-ray pulsar PSR J1640-4631 to be $n=3.15\pm0.03$, which is the highest value of all pulsars with measured braking indices so far. In this Letter, we investigate whether pulsar braking by combined between the magnetic dipole emission and the gravitational radiation might have a braking index greater than three. For conventional neutron star and low mass quark star candidates, the inferred ellipticities derived by the observed braking index are obviously much larger than the theoretical estimated maximum value. If PSR J1640-4631 is a low-mass neutron star with a mass of $0.1~ \rm M_{\odot}$, the inferred ellipticity can be approximately equal to the theoretical estimated maximum value. Because of the radio-quiet nature of this source, we employ the vacuum gap model developed by Ruderman and Sutherland to constrain the inclination angle to be $87.2-90^{\circ}$. Based on this, we propose that a low-mass neutron star with a large inclination angle can interpret the high braking index and the radio-quiet nature of this source. Future observations such as gravitational wave detection and long-term timing for this source are required to confirm or confute our scenario. }

The source PSR J1640-4631 is an X-ray pulsar discovered by \cite{gott14} in a NuSTAR survey of the Norma region in the Galactic plane. Recently, the timing data offer some important parameters of this source such as the spin frequency $\nu=4.84~\rm s^{-1}$, the frequency derivative $\dot{\nu}=-2.28\times 10^{-11}~\rm s^{-2}$, and the second frequency derivative $\ddot{\nu}=3.38\times 10^{-22}~\rm s^{-3}$ (Archibald et al. 2016). It is generally thought that the rotational kinetic energy of pulsars is converted into radiation energy through magnetic dipole radiation. From the general power-law of pulsar spin-down $\nu=-K\nu^{n}$ (where $K$ is a constant depending on the momentum of inertia, the magnetic field, and the radius of pulsars), we can obtain the so-called braking index $n=\nu\ddot{\nu}/\dot{\nu}^{2}$. Based on the timing data, \cite{arch16} derived the braking index of PSR J1640-4631 to be $n=3.15\pm0.03$. The authors performed a series of simulations to identify whether this measurement of the braking index arises from the timing noise. Only 0.01 \% of these simulations yielded a braking index greater than three. Therefore, \cite{arch16} thought that the timing noise indicates a very low level and is very probably not responsible for the measured high braking index. Until 2015, eight pulsars were known with relatively reliable measurements for the braking indices (Lyne et al. 2015). All of these pulsars have a braking index lower than three, showing that the spin-down mechanism is not pure magnetic dipole radiation. As a new member of the braking indices population, PSR J1640-4631 challenges the radiation and braking mechanisms of pulsars. The magnetic field decay can result in a braking index higher than three (Blandford \& Romani 1988; Gourgouliatos \& Cumming 2015). The change of inclination angle $\alpha$ between the magnetic axis and the spin axis may also lead to a braking index that is different from three. A long-term observation of the Crab pulsar has revealed that its inclination angle $\alpha$ may be slowly changing at a rate $\dot{\alpha}\sim 1^{\circ}\rm ~century^{-1}$ (\cite{lyne13}). Based on the inclination angle change model, $\alpha$ of PSR J1640-4631 was constrained to be $18.5\pm3$ degrees, and is decreasing at a rate of $\dot{\alpha}=- (0.23\pm0.05)^{\circ}\rm ~century^{-1}$ (\cite{eksi16}). According to the wind braking model, \cite{kou16} found that the alignment of the inclination angle can cause the braking index to first increase and then decrease, corresponding to a braking index first higher and then smaller than three. In addition, either a mass quadrupole or a magnetic quadrupole would produce gravitational radiation (\cite{blan88}). The pulsars braking by the magnetic dipole emission and the gravitational radiation could give rise to a braking index of between three to five (Archibald et al. 2016; de Araujo et al. 2016).. Employing the derived ellipticity by the braking index, \cite{de16} concluded that aLIGO would not detect the gravitational wave from PSR J1640-4631, while the planned Einstein Telescope would be able to do so. We here constrain the properties of PSR J1640-4631. In Sect. 2 the braking torques caused by the magnetic dipole emission and the gravitational radiation are applied to investigate the spin-down of this source. In Sect. 3 we constrain the properties of this source by the inferred ellipticity and the vacuum gap model. Finally, we summarize the results with a brief conclusion in Sect. 4.

Assuming that the gravitational radiation of a pulsar with elastic deformation can be responsible for the high braking index of PSR J1640-4631, we here constrained its properties. To produce the observed high braking index, the ellipticity of the pulsar is $\epsilon \approx 5\times 10^{-3}$ if the pulsar was a CNS or LMNS. According to the theoretical estimation, it is impossible to possess such a high ellipticity for CNS, while this ellipticity is approximately consistent with the estimated maximum for LMNS. For an LMQS candidate, an ultra-high ellipticity $\epsilon \simeq 0.2$ is required. Based on the vacuum gap model and LMNS candidate, an inclination angle in the range of $\alpha= 87.2-90^{\circ}$ might account for the radio-quiet nature of this source. Based on the magnetic dipole model with corotating plasma, \cite{eksi16} obtained two possible inclination angles for PSR J1640-4631: $18.5\pm3^{\circ}$ and $56\pm 4^{\circ}$. However, according to the vacuum gap model, CNS can emit in the radio. If their model is correct, the observer is located exactly in the forbidden zone of the radio emission beam. According to the alignment torque equation given by \cite{phil14} \begin{equation} I\frac{{\rm d}\alpha}{{\rm d}t}=-\frac{2B^{2}R^{6}\Omega^{2}}{3c^{3}}{\rm sin}\alpha {\rm cos}\alpha, \end{equation} the LMNS model predicts a change rate of the inclination angle $\dot{\alpha}\approx -0.03^{\circ}~\rm century^{-1}$, which is approximately one order of magnitude lower than the value ($-0.23\pm0.05^{\circ}\rm ~century^{-1}$) predicted by \cite{eksi16}. This difference is expected to be explained in the future timing observations. Perhaps a serious problem in this work is how to yield such an LMNS. In classical stellar evolution theory, it is impossible to form such an LMNS directly through a supernova explosion. An LMNS might be produced through fragmentation during protoneutron star formation or neutron star collisions (Popov et al. 2007; Horowitz 2010). The deformation of an LMNS should be relatively easy to understand because it experiences a violent event. However, the physical mechanism triggering fragmentation is still a puzzle. To summarize, PSR J1640-4631 cannot be a CNS or LMQS if the high braking index originates from gravitational radiation. However, it is difficult to confirm that this source is an LMNS. One possibility is to try and detect the gravitational wave signals from this source. The characteristic gravitational wave strain equation reads (Abbott et al. 2007) \begin{equation} h_{0}=\frac{16\pi^{2} G}{c^{4}}\frac{\epsilon I \nu^{2}}{d}, \end{equation} where $d$ is the distance of the source. Insertingly, for the observed parameters and the derived ellipticity of PSR J1640-4631 we have $h_{0}=4.2\times 10^{-26}$. In the gravitational wave frequency $\nu_{\rm gw}=9.68 ~\rm Hz$, this signal is lower than the strain sensitivity of aLIGO (de Araujo et al. 2016). Therefore, we expect that the planned Einstein Telescope will investigate this source in the future. According to our scenario, low-mass neutron stars with a braking index greater than three are important gravitational wave sources. If no gravitational wave from this source is detected, the reasons might be as follows. First, the high braking index of PSR J1640-4631 might be caused by other mechanisms such as the inclination angle change (\cite{eksi16}), the quadrupolar field structure (\cite{petr15}), and the magnetic field decrease. Second, the high braking index may arise from the contaminant of an anomalous $\ddot{\nu}$ after glitches (Alpar \& Baykal 2006). Some pulsars were observed to have an exponential recovery in $\dot{\nu}$ after glitches (see also the panel of PSR J1531-5610 in Fig. 7 of Yu et al. 2013). Such an exponential recovery process in $\dot{\nu}$ could result in a high measured value of $\ddot{\nu}$ (Johnston \& Galloway 1999; Hobbs et al. 2010), which does not originate from a secular braking torque. Such a glitch might naturally produce an anomalous braking index. Therefore, we expect further observations aimed to detect gravitational waves and long-term timing of this source to confirm or confute our scenario.

1607.07487