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1607

1607.08082_arXiv.txt

The Andromeda Galaxy recurrent nova \novak\ had been observed in eruption ten times, including yearly eruptions from 2008--2014. With a measured recurrence period of $P_{\mathrm{rec}}=351\pm13$\,days (we believe the true value to be half of this) and a white dwarf very close to the Chandrasekhar limit, \novak\ has become the leading pre-explosion supernova type Ia progenitor candidate. Following multi-wavelength follow-up observations of the 2013 and 2014 eruptions, we initiated a campaign to ensure early detection of the predicted 2015 eruption, which triggered ambitious ground and space-based follow-up programs. In this paper we present the 2015 detection; visible to near-infrared photometry and visible spectroscopy; and ultraviolet and X-ray observations from the \swift observatory. The LCOGT 2\,m (Hawaii) discovered the 2015 eruption, estimated to have commenced at Aug.\ $28.28\pm0.12$ UT. The 2013--2015 eruptions are remarkably similar at all wavelengths. New early spectroscopic observations reveal short-lived emission from material with velocities $\sim13000$\,km\,s$^{-1}$, possibly collimated outflows. Photometric and spectroscopic observations of the eruption provide strong evidence supporting a red giant donor. An apparently stochastic variability during the early super-soft X-ray phase was comparable in amplitude and duration to past eruptions, but the 2013 and 2015 eruptions show evidence of a brief flux dip during this phase. The multi-eruption \swiftk/XRT spectra show tentative evidence of high-ionization emission lines above a high-temperature continuum. Following \citet{2015A&A...582L...8H}, the updated recurrence period based on all known eruptions is $P_\mathrm{rec}=174\pm10$\,d, and we expect the next eruption of \nova to occur around mid-Sep.\ 2016.

Novae are the powerful eruptions resulting from a brief thermonuclear runaway (TNR) occurring at the base of the surface layer of an accreting white dwarf (WD; see \citealp{1949AnAp...12..281S,1951AnAp...14..294S,1959ApJ...130..916C,1957IAUS....3...77G,1972ApJ...176..169S}, and \citealp{Sta08,2016PASP..128e1001S,JS08,Jos16}, for recent reviews). Belonging to the group of cataclysmic variables \citep{1949ApJ...109...81S,1954ApJ...120..377J,1964ApJ...139..457K}, the companion star in these interacting close-binary systems transfers hydrogen-rich material to the WD usually via an accretion disk around the WD. The TNR powers an explosive ejection of the accreted material, with a rapidly expanding pseudo-photosphere initially increasing the visible luminosity of the system by up to eight orders of magnitude \citep[see][for recent reviews]{2008clno.book.....B,2010AN....331..160B,2014ASPC..490.....W}. Following the TNR the nuclear fusion enters a period of short-lived, approximately steady-state, burning until the accreted fuel is exhausted, partly because it has been ejected and partly as that remaining has been burned to helium \citep{1978A&A....62..339P}. As the optical depth of the expanding ejecta becomes progressively smaller, the pseudo-photosphere begins to recede back toward the WD surface, subsequently shifting the peak of the emission back to higher energies until ultimately a supersoft X-ray source (SSS) may emerge \citep[see, for example,][]{2006ApJS..167...59H,2008ASPC..401..139K,2015JHEAp...7..117O}. The `turn-off' of the SSS indicates the end of the nuclear burning, after which the system eventually returns to its quiescent state. All nova eruptions are inherently recurrent, with the WD and companion surviving each eruption, and accretion reestablishing or continuing shortly afterwards. By definition, the Classical Novae (CNe) have had a single {\it observed} eruption, whereas Recurrent Novae (RNe) have been detected in eruption at least twice. Observed intervals between eruptions range from $\sim1$\,yr \citep[for \novak]{2014A&A...563L...9D} up to 98\,yrs \citep[for V2487~Ophiuchi]{2010ApJS..187..275S}, with the shortest predicted recurrence period -- albeit derived from incomplete observational data -- being just six months \citep{2015A&A...582L...8H}. The theoretical limits on the recurrence period of all novae may be as short as 50\,days \citep{2015MNRAS.446.1924H} or even 25\,days \citep{2016ApJ...824...22H}\footnote{In both the \citet{2015MNRAS.446.1924H} and \citet{2016ApJ...824...22H} studies, accretion is assumed to completely stop during the eruption period.}, and as high as mega-years \citep[see, for example,][]{1985ApJ...291..136S,1994ApJ...424..319K,2005ApJ...623..398Y}. The shorter recurrence periods are driven by a combination of a high mass WD and a high mass accretion rate. Such high accretion rates are typically driven by an evolved companion star, such as a Roche lobe overflowing sub-giant star (SG-novae; also the U~Scorpii type of RNe), or the stellar wind from a red giant companion \citep[RG-novae; symbiotic novae; or the RS~Ophiuchi type RNe; see][for recent reviews]{2012ApJ...746...61D,2014ASPC..490...49D}. With the most luminous novae reaching peak visible magnitudes $M_{V}<-10$ \citep{2009ApJ...690.1148S,2016WillLN}, novae are readily observable out to the distance of the Virgo Cluster and beyond \citep[see, for example,][]{2015ApJ...811...34C,2016arXiv160200758S}. But it is the nearby Andromeda Galaxy (\m31), with an annual nova rate of $65^{+16}_{-15}$\,yr$^{-1}$ \citep{2006MNRAS.369..257D}, that provides the leading laboratory for the study of galaxy-wide nova populations \citep[see, for example,][]{1987ApJ...318..520C,1990ApJ...356..472C,2001ApJ...563..749S,2004MNRAS.353..571D,2006MNRAS.369..257D,2008A&A...477...67H,2010A&A...523A..89H,2011A&A...533A..52H,2014A&A...563A...2H,2011ApJ...727...50S,2011ApJ...734...12S,2015ApJS..216...34S,2014ApJS..213...10W,2016ApJ...817..143W}. Since the discovery of the first \m31 nova by \citet[also spectroscopically confirmed]{1917PASP...29..210R} and the pioneering work of \citet{1929ApJ....69..103H} more than 1000 nova candidates have been discovered \citep[see][and their on-line database\footnote{\url{http://www.mpe.mpg.de/~m31novae/opt/m31/index.php}}]{2007A&A...465..375P,2010AN....331..187P} with over 100 now spectroscopically confirmed \citep[see, for example,][]{2011ApJ...734...12S}. Recently, pioneering X-ray surveys with \xmm and \chandra have revealed that novae are the major class of SSSs in \m31 \citep[][]{2005A&A...442..879P,2007A&A...465..375P}. A dedicated multi-year follow-up program with the same telescopes studied the multi-wavelength population properties of \m31 novae in detail \citep[][]{2010A&A...523A..89H,2011A&A...533A..52H,2014A&A...563A...2H}. A major result of this work was the discovery of strong correlations between various observable parameters: indicating that novae with a faster visible decline tend to show a shorter SSS phase with a higher temperature \citep{2014A&A...563A...2H}. This is consistent with the trends seen in Galactic novae \citep[see][]{2011ApJS..197...31S}. Theoretical models indicate that a shorter SSS phase corresponds to a higher mass WD \citep[e.g.][]{2006ApJS..167...59H,2010ApJ...709..680H,2013ApJ...777..136W}. Thus, the \m31 nova population provides a unique framework within which to understand the properties of individual novae and their ultimate fate. Supernovae Type Ia (SNe\,Ia) are the outcome of a thermonuclear explosion of a carbon--oxygen (CO) WD as it reaches and surpasses the \citet{1931ApJ....74...81C} mass limit \citep[see, for example,][]{1973ApJ...186.1007W,1999ApJ...522..487H,1999ApJ...519..314H,2000ARA&A..38..191H}. An accreting oxygen--neon (ONe) WD however, is predicted to undergo electron capture and subsequent neutron star formation \citep[see, for example,][]{1996ApJ...459..701G}. It seems increasingly likely that there is not a single progenitor pathway producing all observed SNe\,Ia; but a combination of different double-degenerate (WD--WD) and single-degenerate (SD; WD--donor) binary systems, with the metallicity and age of the parent stellar population possibly determining the weighting of those pathways \cite[see, for example,][]{2014ARA&A..52..107M}. Novae, in particular the RNe with their already high mass WDs, are potentially a leading SD pathway. Recent studies have indicated that the mass of the WD can indeed grow over time in RN systems \citep[see, for example,][]{2008NewAR..52..386H,2012BASI...40..419S,2016ApJ...819..168H}. A number of questions remain over the size of their contribution to the SN\,Ia rate; including the composition of the WD in the RN systems, the feasibility of growing a CO WD from their formation mass to the Chandrasekhar limit, and the size of the population of high mass WD novae. Of course, the lack of observational signatures of hydrogen following {\it most} SN\,Ia explosions still provides a significant hurdle for the SD scenario \citep[see, for example,][]{2012NewAR..56..122W,2014ARA&A..52..107M}. But the unmistakable presence of hydrogen in PTF\,11kx \citep{2012Sci...337..942D} and the possible presence of hydrogen in SN\,2013ct \citep{2016MNRAS.457.3254M} supports the view that at least some SN\,Ia arise in SD systems. At the time of writing, there have been around 450 detected eruptions of nova candidates in the Milky Way \citep{2014ASPC..490...49D} from which just ten confirmed RN systems are known \citep{2010ApJS..187..275S} accounting for $\sim3\%$ of known Galactic nova systems or $\sim9\%$ of detected Galactic eruptions. A number of recent detailed studies of archival observations have uncovered new results relating to the RN populations of both the Milky Way and \m31, these are summarized below: \citet{2014ApJ...788..164P} used a combination of three different methods to estimate that the RN nova population (essentially $10\le P_{\mathrm{rec}}\le100$\,yrs; A.~Pagnotta, priv.\ comm.) of the Milky Way is $25\pm10\%$ of the Galactic nova population. However, the range of methodologies employed predicted a wide range of contributions, from $9 - 38\%$, with the authors themselves indicating that the statistical errors were ``likely being much too small'' \citep{2014ApJ...788..164P}. \citet{2015ApJS..216...34S} uncovered multiple eruptions of 16 RN systems in \m31, the subsequent analysis predicted an historic \m31 RN discovery efficiency of just 10\% and that as many as 33\% of \m31 nova eruptions may arise from RN systems ($P_{\mathrm{rec}}\le100$\,yrs). \citet{2014ApJS..213...10W,2016ApJ...817..143W} employed a different approach, by recovering the progenitor systems of 11 \m31 RG-novae, they determined that $30^{+13}_{-10}\%$ of all \m31 nova eruptions occur in RG-nova systems, a sub-population that also appears strongly associated with the \m31 disk. Additionally, other recent results for the Milky Way \citep{2016arXiv160602358S}, Magellanic Clouds \citep{2016ApJS..222....9M}, \m31 \citep{2016MNRAS.458.2916C,2016MNRAS.455..668S}, and M\,87 \citep{2016arXiv160200758S} all indicate that the luminosity specific nova rate \citep[see, for example,][]{1990AJ.....99.1079C} may be much higher than previously thought. Together, all these results boost the size of the available `pool' of novae that may contribute to the SN\,Ia population by a factor of $>5$.

\label{sec:conclusions} The 2015 eruption of \novak\ was discovered independently by dedicated monitoring programs utilizing the \swift orbiting observatory and the LCOGT 2\,m (Hawaii) on 2015 Aug 28.41 UT and 28.425, respectively, with pre-eruption non-detections constraining the time of the eruption to 2015 Aug $28.28\pm0.12$ UT. Following detection, a pre-planned pan-chromatic follow-up campaign was initiated which involved ten ground-based telescopes around the globe, but was spearheaded by \swiftk, the LT, and the LCOGT. The eruption light curves spanning the electromagnetic-spectrum from the super-soft X-rays to the $I/i'$-band show remarkable similarity between the 2013, 2014, and 2015 eruptions. The combined visible spectrum from the 2012, 2014, and 2015 eruptions shows tentative evidence for high-excitation coronal lines of [Fe\,{\sc vii}], [Fe\,{\sc x}], and [Fe\,{\sc xiv}], often observed during high temperature shocks, and also hints at the presence of Raman scattered O\,{\sc vi} emission, as seen in spectra of symbiotic stars and novae with red giant companions. The visible spectra from the 2012--2015 eruptions show a consistent decrease in line width. Between days 1 and 4, post eruption, this deceleration is consistent with a power-law decline of the ejection velocity ($v\propto t^{-1/3}$). This deceleration is consistent with consistent with the adiabatic Phase~II shock development as the ejecta interact with significant preexisting circumbinary material. These observations, backed-up by the color--magnitude behavior, point to the donor being a red giant in a long orbital period system. Below we summarize a number of our conclusions. \begin{enumerate} \item The color--magnitude evolution in the visible appears more consistent with the behavior of RS~Oph and V745 Sco (both harboring red giant donors) than that of the sub-giant (e.g.\ U~Sco) or main sequence (T~Pyx) donor RNe. \item There is no evidence at visible wavelengths for optically thick photospheric emission during the early evolution of the eruption. This points to a large {\it minimum} temperature of the expanding photosphere, with photospheric emission therefore peaking in the FUV or EUV. \item The evolving SED of the eruption points to optically thick free-free emission being the dominant process (in the NIR--NUV) throughout the evolution from $t=0.7$\,d to $t=10$\,d. Although significant contribution to the SED from emission lines cannot be ruled beyond day four. \item The $V$- and $r'$-band trends in the SED leads to a prediction of the nebular phase beginning as early as day 5 post-eruption. \item Emission from extremely high velocity ($\mathrm{FWHM}\simeq13000\,\mathrm{km}\,\mathrm{s}^{-1}$) material seen only in the early spectra ($t\la1$\,d) is indicative of outflows along the polar direction -- possible highly collimated outflows or jets. \item We obtained an unprecedentedly detailed UV light curve with \swift UVOT, observing for the first time the rise to maximum and fast decline with subsequent plateaus. The UV peak clearly precedes the visible peak. \item The X-ray light curve of the 2015 eruption was consistent with the last two years in its time scales, $\eton = 5.6\pm0.7$~d and $\etoff = 18.6\pm0.7$~d, as well as in the properties of the early SSS variability and its cessation around day 13. \item The 2015 X-ray light curve also showed evidence of a peculiar dip around day 11 which might have been present in the 2013 light curve as well. \item Merged X-ray spectra tentatively suggest the presence of high-ionization emission lines superimposed on a photospheric continuum that reaches black body temperatures of around 120\,eV. \end{enumerate} The next eruption of \novak\ is predicted for mid Sep.\ 2016 with a $1\sigma$ uncertainty of about 1 month. This prediction holds for both the 1 year and the 6 month recurrence scenarios. In the case of the 6 month period, we expect the subsequent eruption in Feb.\ -- Apr.\ 2017.

1607.08082

1607

1607.08561_arXiv.txt

We investigate how baryogenesis can occur by the presence of an $f(T)$-related gravitational term. We study various cases of $f(T)$ gravity and we discuss in detail the effect of the novel terms on the baryon-to-entropy ratio. Additionally, we study the constraints imposed by the observational values of the baryon-to-entropy ratio and we discuss how more generalized cosmologies can contribute successfully, in a viable and consistent way, in the gravitational baryogenesis mechanism.

One of the main mysteries of the standard cosmological paradigm is the explanation of the excess of matter over antimatter, which is verified by Cosmic Microwave Background observations \cite{Bennett:2003bz}, as well as from Big Bang nucleosynthesis predictions \cite{Burles:2000ju}. Gravitational baryogenesis is one of the mechanisms that have been proposed for the generation of such baryon-anti-baryon asymmetry \cite{Davoudiasl:2004gf,Lambiase:2006dq,Lambiase:2013haa,Lambiase:2006ft,Li:2004hh, Pizza:2015epa,Odintsov:2016hgc}. Moreover, the gravitational baryogenesis has some crucial effects on singular inflation (see for example \cite{oikonomou}). This mechanism for baryon asymmetry incorporates one of Sakharov's criteria \cite{sakharov}, and the baryon-anti-baryon asymmetry is obtained by the presence of a $\mathcal{C}\mathcal{P}$-violating interaction term of the form \begin{equation} \label{baryonassterm} \frac{1}{M_*^2}\int \mathrm{d}^4x\sqrt{-g}(\partial_{\mu} R) J^{\mu}\, . \end{equation} Such a term could be acquired from higher-order interactions in the fundamental gravitational theory \cite{Davoudiasl:2004gf}. In particular, $M_*$ is the parameter that denotes the cutoff scale of the underlying effective theory, $J^{\mu}$ is the baryonic matter current, and $g$ and $R$ are respectively the metric determinant and the Ricci scalar. If one applies the above in the case of a flat Friedmann-Robertson-Walker (FRW) geometry, then the baryon-to-entropy ratio $\eta_B/s$ is proportional to $\dot{R}$, and especially in the case where the matter fluid corresponds to relativistic matter with equation-of-state parameter $w=1/3$ then the net baryon asymmetry generated by the term (\ref{baryonassterm}) is zero. In the present work we are interested in investigating the gravitational baryogenesis mechanism in the framework of $f(T)$ gravity, which is a gravitational modification based on the torsional (teleparallel) formulation of gravity. In particular, in the Teleparallel Equivalent of General Relativity (TEGR) \cite{ein28,Hayashi79,Pereira.book,Maluf:2013gaa} the gravitational Lagrangian is the torsion scalar $T$, and hence one can construct torsional modified gravity by extending it to $f(T)$ \cite{Bengochea:2008gz,Linder:2010py} (see \cite{Cai:2015emx} for a review). The interesting point is that although TEGR is completely equivalent with general relativity at the equation level, $f(T)$ gravity corresponds to different gravitational modification than $f(R)$ one, and therefore its cosmological implications bring novel features \cite{Dent:2011zz,Geng:2011aj,Bamba:2013jqa,F-T-Inf}. Particularly, we shall examine in detail the effects of various gravitational baryogenesis terms which are proportional to $\partial_{\mu}T$ or $\partial_{\mu}f(T)$. As we will show, for the simplest choice of $f(T)$ gravity, the resulting baryon-to-entropy ratio can be compatible to observations, only if some parameters are chosen to be abnormally large. Furthermore, we will constrain the functional form of more general $f(T)$ gravities which can realize a radiation dominated Universe, and finally we shall discuss how more general cosmologies can be contribute successfully to the gravitational baryogenesis scenario. This paper is organized as follows: In section \ref{modela} we briefly review the fundamental properties of $f(T)$ gravity. In section \ref{BaryogenesisF} we discuss various gravitational baryogenesis scenarios in the context of $f(T)$ gravity, and we examine the qualitative implications on the baryon-to-entropy ratio which we calculate in detail for each case under study. Finally, the conclusions follow in the end of the paper.

\label{Conclusions} In this paper we studied the gravitational baryogenesis scenario, generated by an $f(T)$ theory of gravity. In the context of $f(T)$ baryogenesis, the baryon-to-entropy ratio depends on $\dot{T}$, and we discussed two cases of $f(T)$ theories of gravity, the case $f(T)=T$ and also $f(T)\sim (-T) ^n$. In the first case, the resulting picture is not so appealing since in order for the predicted baryon-to-entropy ratio to be compatible to the observational value, some of the parameters must given abnormally small values. The case $f(T)\sim (-T)^n$ is more interesting and we investigated which values should the parameter $n$ take in order to have compatibility with the data. As we showed, the variable $n$ plays a crucial role in the calculation of the baryon-to-entropy ratio. In both cases the interesting new feature is that the baryon-to-entropy ratio is non-zero for a radiation dominated Universe, in contrast to the Einstein-Hilbert gravitational baryogenesis scenario. Finally, we investigated how more general cosmologies affect the baryon to entropy ratio, and when the gravitational baryogenesis term is of the form $\partial_{\mu} T$, inconsistencies may occur in the theory. As we showed, the remedy to this issue is to modify the gravitational baryogenesis term, so that the baryon current is coupled to $\partial_{\mu} f(T)$. In this way more general cosmological evolutions can be considered and the resulting baryon-to-entropy ratio is compatible to the observational data.

1607.08561

1607

1607.03735_arXiv.txt

Establishing the field content during inflation is a fundamental challenge of primordial cosmology. Minimal inflationary models have two massless fields: the Goldstone boson of broken time translations,\footnote{Strictly speaking, $\pi$ is only massless in the decoupling limit $\Mp \to \infty$. However, for adiabatic fluctuations, $\pi$ is directly related to the comoving curvature perturbation, $\zeta=-H\pi +{\cal O}(\pi^2)$, which is the true massless degree of freedom even away from the decoupling limit.} $\pi$, and the graviton, $\gamma_{ij}$. While at present there is no evidence for additional degrees of freedom~\cite{Ade:2015ava}, the imprints of extra particles can be subtle, so it remains important to fully characterize their effects and compare them to observations. Moreover, massive particles are important probes of the ultraviolet completion of inflation. For example, in string theory, massive particles in the low-energy effective theory encode physics at the string and Kaluza-Klein scales~\cite{Baumann:2014nda}. If these scales aren't too far from the inflationary Hubble scale, then their influence may be observable (although the experimental challenge could be enormous). \vskip 4pt \begin{figure}[h!] \centering \includegraphics[scale=0.45]{figures/3pt_intro1}\qquad \includegraphics[scale=0.45]{figures/3pt_intro2} \caption{Diagrams contributing to $\langle\zeta\zeta\zeta\rangle$ and $\langle\gamma\zeta\zeta\rangle$. The solid, dashed, and wavy lines represent the curvature perturbation $\zeta$, a massive spin-$s$ field $\sigma_{i_1\cdots i_s}$, and the graviton $\gamma_{ij}$, respectively. \label{fig:diagrams0} } \end{figure} Since massive particles decay outside of the horizon during inflation, they cannot be observed directly in late-time correlation functions. Instead, the presence of massive particles has to be inferred from their indirect effects on the correlation functions of $\zeta = -H\pi$ and $\gamma_{ij}$ (see Fig.~\ref{fig:diagrams0}). Some of these effects can be mimicked by adding a local vertex in the low-energy effective Lagrangian, which is the result of integrating out the heavy fields. On the other hand, massive particles may spontaneously be created in an expanding spacetime~\cite{parker1968particle, parker1969quantized,parker1971quantized}, an effect which cannot be represented by adding a local vertex to the effective Lagrangian~\cite{Arkani-Hamed:2015bza}. The role of these non-local effects as a means of detecting massive particles during inflation was recently highlighted by Arkani-Hamed and Maldacena (AHM)~\cite{Arkani-Hamed:2015bza}: the spontaneous particle creation allows us to probe massive fields during inflation, even though we are only observing the late-time expectation values of light fields. The rate of particle production in de Sitter space is exponentially suppressed as a function of mass, $e^{-m/T_{\rm dS}}$, with $T_{\rm dS} \equiv H/2\pi$, so their imprints will only be detectable if their masses are not too far above the Hubble rate $H$.\footnote{If the extra fields have strongly time-dependent masses, whose Fourier transforms have support at a frequency~$\hat \omega$, then non-adiabatic particle production occurs at a rate proportional to $e^{- m/\hat \omega}$~\cite{Flauger:2016idt}. The scale $\hat \omega$ may be as large as $\dot \phi^{1/2} = 58 \hskip 1pt H$ without spoiling the slow-roll dynamics. In models with these types of time-dependent couplings, the detectable range of particle masses is somewhat enlarged.} Since the inflationary scale may be as high as $10^{14}\hskip 1pt{\rm GeV}$, this nevertheless provides an opportunity to probe massive particles far beyond the reach of conventional particle colliders. \vskip 4pt Nonlinearities in the decay of the massive particles lead to a non-Gaussianity in the late-time correlation functions of $\zeta$ and $\gamma_{ij}$. The form of this non-Gaussianity will depend on the masses and the spins of the extra particles. The effects of additional scalar fields during inflation have been explored in many previous works, e.g.~in the context of quasi-single-field inflation~\cite{Chen:2009zp, Baumann:2011nk, Noumi:2012vr}. A characteristic signature of these fields are non-analytic scalings in the soft momentum limits of the non-Gaussian correlation functions. These soft limits are particularly clean detection channels, since in single-field inflation their momentum scalings are fixed by the symmetries of the inflationary background~\cite{Maldacena:2002vr, Creminelli:2004yq}. The most straightforward interpretation of such non-analyticity in the correlation functions is therefore the presence of extra particles. Scalar fields with masses less than $\frac{3}{2}H$ give rise to monotonic scalings in the squeezed limit~\cite{Chen:2009zp, Baumann:2011nk}, while those with masses greater than $\frac{3}{2}H$ lead to oscillatory behavior~\cite{Noumi:2012vr, Arkani-Hamed:2015bza, Mirbabayi:2015hva, Chen:2015lza}. The effects of extra massive particles with spin have not been studied in as much detail. Such particles can naturally arise as massive Kaluza-Klein modes or as part of the tower of higher-spin states from string theory~\cite{Rindani:1985pi, Aragone:1988yx}. As was shown by AHM, the spins of new particles lead to a distinctive angular dependence of the soft limits of the non-Gaussian correlators. The analysis of AHM was restricted to the squeezed limit of the bispectrum and interactions that maintained the approximate conformal invariance of the inflationary background. While this assumption made their analysis particularly well controlled, it also implied that the amplitude of the signal is highly suppressed and only observable in the most optimistic and futuristic scenarios. \vskip 4pt We will drop some of the restrictions of the analysis of AHM in our analysis. In particular, we will allow for a large breaking of conformal invariance within the framework of the effective field theory (EFT) of inflation~\cite{Cheung:2007st}. We will find that the signal due to massive spinning particles can be observable within the regime of validity of the EFT. At the same time, the main spectroscopic features of particles with spin during inflation do not rely on conformal invariance and therefore still apply. On the other hand, couplings to particles with odd spins, which are disallowed in the conformally-invariant case, are permitted in the generic effective theory. We also consider the breaking of special conformal invariance by giving the Goldstone fluctuations a nontrivial sound speed. In that case, we find a reduced exponential suppression in the particle production rate, and thus an enhanced level of non-Gaussianity. Finally, we also study the coupling to an external graviton~$\gamma_{ij}$. We demonstrate that the soft graviton limit of the correlator $\langle \gamma \zeta \zeta\rangle$ provides an interesting detection channel for extra particles. Like in the case of massive scalar fields, there will be non-analytic scalings of non-Gaussianities close to the soft momentum limit, but this time only for particles with spin greater than or equal to two. \paragraph{Outline} In this paper, we analyze the allowed couplings of massive particles with spin to the Goldstone boson of broken time translations and the graviton, and discuss their observational signatures. In Section~\ref{sec:dS}, we first collect the equations of motion for massive fields with spin in de Sitter space, whose solutions are presented in Appendix~\ref{app:spindS}. In Section~\ref{sec:EFT}, we then construct the effective action for the leading interactions between the Goldstone boson $\pi$, the graviton $\gamma_{ij}$, and massive spinning fields $\sigma_{\mu_1 \ldots \mu_s}$. We analyze under what conditions the theory is under perturbative control and discuss various constraints on the sizes of the couplings. In Section~\ref{sec:correlators}, we compute the correlation functions associated with the interactions of Section~\ref{sec:EFT}. We estimate the maximal amount of non-Gaussianity that is consistent with the constraints on the couplings of the effective theory. Details of the in-in computation are relegated to Appendix~\ref{app:inin}, and analytic results for soft limits are given in Appendix~\ref{app:squeezed}. Our conclusions are presented in Section~\ref{sec:conclusions}. \paragraph{Notation and conventions} We will use natural units, $c=\hbar=1$, with reduced Planck mass $\Mp^2=1/8\pi G$. Our metric signature is ($-++\hskip 1pt +$). We will use Greek letters for spacetime indices, $\mu, \nu, \ldots =0,1,2,3$, and Latin letters for spatial indices, $i,j,\ldots=1,2,3$. Three-dimensional vectors are written in boldface, $\k$, and unit vectors are hatted, $\hat \k$. A shorthand for the symmetrization of tensor indices is $a_{(\mu} b_{\nu)} \equiv \frac{1}{2} (a_\mu b_\nu + a_\nu b_\mu)$. Overdots and primes will denote derivatives with respect to physical time $t$ and conformal time $\eta$, respectively. The letter $\pi$ will refer both to $3.141\ldots$ and the Goldstone boson of broken time translations. The dimensionless power spectrum of a Fourier mode $f_\k$ is defined as \beq \Delta_f^2(k) \equiv \frac{k^3}{2\pi^2} \langle f_\k f_{-\k}\rangle' \, , \eeq where the prime on the expectation value indicates that the overall momentum-conserving delta function has been dropped.

\label{sec:conclusions} In this paper, we have studied the imprints of massive particles with spin on cosmological correlators using the framework of the effective field theory of inflation \cite{Cheung:2007st}. This generalizes the work of Arkani-Hamed and Maldacena (AHM)~\cite{Arkani-Hamed:2015bza} to cases where conformal symmetry is strongly broken. Let us summarize our results and contrast them with the conclusions of AHM: \begin{itemize} \item In AHM's more conservative analysis, the overall size of non-Gaussianity was too small to be observable even in the most optimistic experimental scenarios. Our results are cautiously more optimistic. Within the regime of validity of the effective field theory, we can accommodate observable non-Gaussianity as long as the masses of the new particles aren't too far above the Hubble scale during inflation. \item The key spectroscopic features of massive particles with spin do not rely on conformal invariance and therefore continue to hold in our analysis. As explained in~\cite{Arkani-Hamed:2015bza}, the masses and spins of extra particles during inflation can be extracted by measuring the momentum dependence in the squeezed limit. \item Our systematic effective field theory treatment of massive spinning particles during inflation allows for a complete characterization of their effects on non-Gaussian cosmological correlators, including their imprints beyond the squeezed limit. We showed that the characteristic angular dependence resulting from the presence of particles with spin persists even for more general momentum configurations. Having access to the complete correlation functions will be valuable for future data analysis. \item We also studied the effects of an explicit breaking of special conformal symmetry by introducing a sound speed $\cs$ for the Goldstone fluctuations. We found that, for $\cs <\mu_s^{-1}$, the exponential suppression in the production of the massive particles, $e^{-\pi\mu_s}$, is changed to $e^{-\pi\mu_s/2}$. For a given mass, the size of non-Gaussianity is therefore enhanced (or less suppressed) for small~$\cs$. \item Finally, we showed that particles with spin greater than or equal to two lead to a signature in the squeezed limit of $\langle\gamma\zeta\zeta\rangle$. This signal may be observable in the $\langle BTT\rangle$ correlator of CMB anisotropies~\cite{Meerburg:2016ecv}. \end{itemize} Figure~\ref{fig:fNLconstraints} is a schematic illustration of current and future constraints on (scale-invariant) primordial non-Gaussianities. We see that the perturbatively interesting regime spans about seven orders of magnitude in $f_{\rm NL}$. Of this regime, three orders of magnitude have been ruled out by current CMB observations, leaving a window of opportunity of about four orders of magnitude. Accessing these low levels of non-Gaussianity will be challenging. Even optimistic projections for future CMB observations won't reduce the constraints by more than an order of magnitude. Digging deeper will require new cosmological probes, such as observations of the large-scale structure (LSS) of the universe~\cite{Alvarez:2014vva} and the tomography of the 21cm transition of neutral hydrogen gas~\cite{Loeb:2003ya}. Our results, together with~\cite{Chen:2009zp, Baumann:2011nk, Noumi:2012vr,Arkani-Hamed:2015bza,Flauger:2016idt}, will help to find optimal observational strategies for extracting the subtle imprints of extra particles during the inflationary era. \begin{figure}[t!] \centering \quad \includegraphics[scale=0.9]{figures/fNL-constraints} \caption{\label{fig:fNLconstraints} Schematic illustration of current and future constraints on (scale-invariant) primordial non-Gaussianity. The ``gravitational floor" denotes the minimal level of non-Gaussianity created by purely gravitational interactions during inflation~\cite{Maldacena:2002vr}. } \end{figure} \paragraph{Acknowledgements} We thank Paolo Creminelli, Garrett Goon, Dan Green, Juan Maldacena, Daan Meerburg, Mehrdad Mirbabayi, Enrico Pajer, Rafael Porto, Eva Silverstein, and Marko Simonovi\'c for helpful discussions, and Maldacena and Meerburg for comments on a draft. H.L.~thanks the Institute of Physics at the University of Amsterdam for its hospitality. H.L~acknowledges support from the EPSRC and the Cambridge Overseas Trust. D.B.~and G.P.~acknowledge support from a Starting Grant of the European Research Council (ERC STG Grant 279617). \newpage \appendix

1607.03735

1607

1607.02443_arXiv.txt

TARGET~5 is a new application-specific integrated circuit (ASIC) of the TARGET family, designed for the readout of signals from photosensors in the cameras of imaging atmospheric Cherenkov telescopes (IACTs) for ground-based gamma-ray astronomy. TARGET~5 combines sampling and digitization on 16 signal channels with the formation of trigger signals based on the analog sum of groups of four channels. {\rev We describe the ASIC architecture and performance.} TARGET~5 improves over the performance of the first-generation TARGET ASIC, achieving: tunable sampling frequency from {\rev $<0.4$~GSa/s to {\mod $>1$~GSa/s}}; a dynamic range on the data path of 1.2 V with {\mod effective dynamic range of 11}~bits and DC noise of ${\sim}0.6$~mV; 3-dB bandwidth of 500 MHz; {\rev crosstalk between adjacent channels $<1.3\%$}; {\mod charge resolution improving from 40\% to $<4\%$ between 3 photoelectrons (p.e.) and $>100$~p.e.} (assuming 4 mV per p.e.); and minimum stable trigger threshold of 20 mV (5 p.e.) with trigger noise of 5 mV (1.2 p.e.), {\rev which is} mostly limited by {\mod interference between trigger and sampling operations}. {\mod TARGET~5 is the first ASIC of the TARGET family} used in {\modakira an} IACT {\mod prototype}, {\rev providing one development path for readout electronics in the forthcoming Cherenkov Telescope Array (CTA)}.

In the past three decades, imaging atmospheric Cherenkov telescopes (IACTs) have greatly advanced observations of very-high-energy gamma-ray emission from the Universe, with numerous implications for astrophysics, particle physics, and cosmology \cite[e.g.,][]{2009ARA&A..47..523H}. This field is {\mod soon going to be revolutionized} with the advent of the Cherenkov Telescope Array (CTA) \cite{2011ExA....32..193A}, which is going to increase the source sensitivity by an order of magnitude at energies from 100~GeV to 10~TeV and to extend observations to the ranges well below 100 GeV and above 100 TeV. {\rev The performance requirements of CTA drive innovation to improve performance and lower cost. One innovative design is the Schwarzschild--Couder telescope \cite{2007APh....28...10V}, which features dual-mirror optics for excellent optical performance (focusing of Cherenkov photons) and a reduced camera plate scale compared to the traditional single-mirror (Davies-Cotton) design used so far for IACTs. The reduced camera size enables compact, inexpensive, densely pixelated photodetectors such as silicon photomultipliers. The optical performance combined with dense pixellation provides improved field of view, angular resolution, and hadronic background rejection capabilities~\cite{Wood201611}.} TARGET is an application-specific integrated circuit (ASIC) series that has been designed for the processing of the photodetector signals in such telescopes. {\mod The goals in the inception of TARGET were to keep the costs low and integrate several functionalities in a compact design.} We have described the concept of TARGET~1, the first generation of ASICs of the TARGET family, and characterized its performance in \cite{2012APh....36..156B}. Several improvements drove the development of TARGET~5 (after a few design iterations), which is described in this paper. {\rev TARGET~5 is the first chip of the TARGET family to be used in a telescope prototype, namely a prototype {\mod of} the Gamma-ray Cherenkov Telescope (GCT) \cite{2015arXiv150806472M}, a Schwarzschild--Couder small-sized telescope proposed in the framework of the CTA project. TARGET~5 is used in the first prototype of the GCT camera, {\modakira also known as Compact High Energy Camera with MAPMTs (CHEC-M)} \cite{2013arXiv1307.2807D,2015arXiv150901480D}. TARGET is also planned to be used in a medium-sized telescope proposed in the framework of the CTA project, namely the Schwarzschild-Couder Telescope (SCT) \cite{SCTICRC2015}.} Key features of TARGET are: \begin{itemize} \item a compact design that combines signal sampling and digitization, as well as triggering, for 16 channels in a single chip, which lowers the cost\footnote{${\sim}\$35$ per channel for the realization discussed in this paper, estimated $<$~\$20 per channel in a large production on the scale required for several dozen CTA telescopes.}, {\mod improves on reliability further reducing maintenance costs}, and enables the use with compact photodetectors such as multi-anode photomultiplier tubes (MAPMTs) or silicon photomultipliers in a compact camera design \item a sampling frequency tunable up to $>1$~GSa/s, ideally suited for the measurement of the $\gtrsim 5$~ns pulses from Cherenkov flashes \item a deep buffer (16,384 samples in TARGET~5) for large trigger latency tolerance between distant (${\sim}1$~km) telescopes\footnote{\rev Within CTA, using a hardware coincidence trigger between telescopes is not planned for GCTs, but it is foreseen for SCTs.} \item dynamic range $>10$~bits \item {\rev moderate power consumption, for the applications described in this paper $\lesssim20$~mW per channel} \end{itemize} {\mod TARGET ASICs are implemented for IACTs into} front-end electronics modules that combine all the functions described above to read out 64 photodetector pixels using six or fewer printed circuit boards, four ASICs, and a companion field-programmable gate array (FPGA) \cite{2012APh....36..156B}. The low number of components supports affordability and reliability. The structure of this paper is as follows. Section~\ref{sec:architecture} describes the architecture of the TARGET~5 ASIC. Section~\ref{sec:performance} presents the characterization of its performance, including sampling and digitization in \ref{sec:datapath}, as well as triggering in~\ref{sec:trigpath}. Section~\ref{sec:cameramod} briefly outlines how TARGET~5 is implemented into front-end electronics modules for CHEC-M, and Section~\ref{sec:conclusions} presents the conclusions and outlook.

\label{sec:conclusions} We have developed a new ASIC of the TARGET family designed to read out signals from the photosensors in cameras of very-high-energy gamma-ray telescopes exploiting {\mod time-resolved} imaging of Cherenkov light from air showers. TARGET~5 processes signals from 16 photodetector pixels in parallel both for sampling and digitization and for trigger formation. Key aspects of the TARGET~5 performance are: \begin{itemize} \item sampling frequency tunable between {\rev $< 0.4$~GSa/s and $>1$~GSa/s} (Fig.~\ref{sampling-frequency}) \item a dynamic range on the data path of 1.2 V with {\mod effective dynamic range 11}~bits and DC noise ${\sim}0.6$~mV (Fig.~\ref{temperature}) \item 3-dB bandwidth of 500 MHz (Fig.~\ref{fig:bandwidth}) \item {\rev crosstalk between neighboring channels $<1.3\%$ (Fig.~\ref{crosstalk})} \item {\mod charge resolution improving from 40\% to $<4\%$ as a function of input charge between 3 p.e. and $>100$~p.e (assuming 4 mV per p.e.)} (Fig.~\ref{charge-linearity} and \ref{charge-resolution}) \item minimum stable trigger threshold of 20 mV (5 p.e.) with trigger noise of 5 mV (1.2 p.e.), mostly limited by {\mod interference between sampling and trigger operations} (Fig.~\ref{trigger-Scurve-normal}) \item minimum stable trigger threshold of $<5$~mV (1.2 p.e.) with trigger noise of 0.5 mV (0.12 p.e.) with sampling disabled (Fig.~\ref{figures/trigger-Scurve-samploff.pdf}) \end{itemize} TARGET~5 is part of the front-end electronics system of the first GCT camera prototype, also known as CHEC-M \cite{2013arXiv1307.2807D,2015arXiv150901480D}, and is the first ASIC in the TARGET family to be used in an IACT {\mod prototype} which is proposed in the framework of the CTA project. To meet the performance desired for CTA, further developments are ongoing that are briefly outlined in \cite{tibaldoICRC2015}.

1607.02443

1607

1607.07785_arXiv.txt

We present digital tables for the radiative terms that appear in the energy and momentum equations used to simulate the accretion onto supermassive black holes (SMBHs) in the center of galaxies. Cooling and heating rates and radiative accelerations are calculated with two different Spectral Energy Distributions (SEDs). One SED is composed of an {\tt accretion disk + [X-ray]-powerlaw}, while the other is made of an {\tt accretion disk + [Corona]-bremsstrahlung} with $T_X=1.16 \times 10^8$ K, where precomputed conditions of adiabatic expansion are included. Quantification of different physical mechanisms at operation are presented, showing discrepancies and similarities between both SEDs in different ranges of fundamental physical parameters (i.e., ionization parameter, density, and temperature). With the recent discovery of outflows originating at sub-parsec scales, these tables may provide a useful tool to model gas accretion processes onto a SMBH.

Most of our present knowledge of the cosmos has come from application of the principles of quantum mechanics and atomic physics. For instance, the evolution of spectroscopy in every band of the electromagnetic spectrum from radio to $\gamma$-rays has allowed the study of the supernova remnants, the solar winds, the accretion onto supermassive black holes (SMBHs), and the large-scale structure of the Universe in a way that we have never imagined to be possible a century ago. Astrophysical processes that involve radiative energy transfer are calculated by the balance between heating and cooling. Analytical prescriptions for the heating and cooling rates in complex environments are only possible under certain limits. Moreover, it is well-known that they depend on the SED used \citep{kallman1982a} and that stability curves also show a dependence on the SED in active galactic nuclei \citep[eg.,][]{chakravorty2009a,chakravorty2012a}. However, the increasing computer power available today has allowed to model complex astrophysical scenarios efficiently and at a relatively low cost, including the dynamical update of the microphysics and chemistry. Non-equilibrium thermodynamics, ionization, molecular states, level populations, and kinetic temperatures of low densities environments are some of the ingredients that have no analytical counterparts and that can be calculated with highly efficient numerical algorithms. Among the several publicly available codes for the calculation of astrophysical environments, \cly \sp \citep{cloudy1303} and {\sc xstar} \citep{kallman2001a} have become the most popular because they treat the atomic physics at an {\it ab-initio} level. In addition, they have the ability to correctly handle a wide variety of scenarios, while predicting the spectrum of different gas geometries, including the Ultraviolet (UV) and the Infrared (IR) as well as a broad range of densities up to $\sim 10^{15}$ \cmd\sp and temperatures from the cosmic microwave background (CMB) to $10^{10}$ K. The electronic structure of atoms, the photoionization cross-sections, the recombination rates, and the grains and molecules are also treated in great detail. In particular, the modeling in \cly \sp includes: i) photoionization/recombination, ii) collisional ionization/3-body recombination to all levels, and iii) collisional and radiative processes between atomic levels so that the plasma behaves correctly in the low density limit and converges naturally to local thermodynamic equilibrium (LTE) either at high densities or when exposed to ``quasi-real" blackbody radiation fields \citep{ferland1998a}. Moreover, collisions, line trapping, continuum lowering, and absorption of photons by continuum opacities are all included as very general processes \citep{rees1989a}. Inner-shell processes are also considered, including the radiative one (i.e., line emission after the removal of an electron) \citep{ferland1998a}. On the other hand, analytical formulas for the heating and cooling rates have been widely used. For instance, previous work on accretion onto SMBHs in the center of galaxies (active galactic nuclei, AGNs) by \cite{psk00}, \cite{pk2004}, \cite{proga2007a}, and \cite{barai2011a} have made use of \cite{blondin1994a} analytical formulas for the heating and cooling rates, which are limited to temperatures in the range $10^4\lesssim T \lesssim 10^8$ K and ionization parameters ($\xi=L/[n_H r^2]$) in the interval $1\lesssim \log(\xi) \lesssim 5$. In this paper, we develop a methodology and present tabulated values that account for highly detailed photoionization calculations together with the underlying microphysics to provide a platform for use in existing radiation hydrodynamics codes based either on Smoothed Particle Hydrodynamics (SPH) or Eulerian methods. Using the Cinvestav-{\sc abacus} supercomputing facilities, we have run a very extensive grid of photoionization models using the most up-to-date version of \cly~({\bf v} 13.03), which allows us to pre-visualize physical conditions for a wide range of distances, from four Schwarzschild radii ($\approx 4 r_{\rm Sch}$) to $r \lesssim 34,000 r_{\rm Sch}$ ($r_{\rm Sch}=\frac{2GM_{BH}}{c^2}$), densities ($10^{-2}\lesssim n_H \lesssim 10^9$ \cmd), and temperatures ($10^2\lesssim T \lesssim 10^9$ K) around SMBHs in AGNs. It is well-known that accretion processes onto compact objects may influence the nearby ambient around SMBHs in the center of galaxies \citep[e.g.,][]{salpeter1964a,fabian1999a,barai2008a,germain2009a}. Together with the outflow phenomena, they are believed to play a major role in the feedback process invoked by modern cosmological models (i.e., $\Lambda$-Cold Dark Matter) to explain the possible relationship between the SMBH and the host galaxy \citep[e.g.,][]{magorrian1998a,gebhardt2000a} as well as in the self-regulating growth of the SMBH. The problem of accretion onto a SMBH can be studied via hydrodynamical simulations \citep[e.g.,][]{ciotti2001a,li2007a,ostriker2010a,novak2011a}. In numerical studies of galaxy formation, spatial resolution permits resolving scales from the kpc to the pc, while subparsec scales are not resolved. This is why a prescribed sub-grid is employed to solve this lack of resolution. With sufficiently high X-ray luminosities, the falling material will have the correct opacity, developing outflows that originate at sub-parsec scales. Therefore, the calculation of the present tables provides a tool to solve the problem of accretion onto SMBHs in the center of galaxies at sub-parsec scales. In addition, two SEDs and three ways of breaking up the luminosity between the disk and the X-ray components are presented. On average, these runs take about 200 minutes using $\approx 4000$ cores ($\approx 13.3$k CPU hours) of the Cinvestav-{\sc abacus} supercomputer. There are several radiation hydrodynamics codes that invoke \cly \sp for spectral synthesis. These codes are used to simulate processes subject to strong irradiation such as the formation and evolution of HII regions, photoevaporation of the circumstellar disks, and cosmological minihaloes. For example, \cite{salz2015a} combine a SPH-based magnetohydrodynamics (MHD) code with \cly \sp for the simulation of the photoevaporation of the hot-Jupiter atmospheres. Moreover, \cite{niederwanger2014a} and \cite{ottl2014a} combine a finite-volume MHD code with \cly \sp to simulate planetary nebulae. The paper is structured as follows: in Section \ref{radicool}, we describe the SEDs used and how they break up between UV and X-ray components. Details of the comparison between photoionization calculations using our two SED bases are also provided. The calculation of the radiative acceleration as included in the momentum equations is described in Section \ref{radiaccel}, while Section \ref{mytables} contains details of the structure of the tables along with the meaning, units, and location in the Internet for public use. The discussion of the results and the conclusions are given in Section \ref{diss1}. Two appendices are added for the description of the Sakura \& Sunyaev disk model and the calculation of the ionic fractions. The symbols appearing through the manuscript have the standard meaning: $G\equiv$ Newtonian gravitational constant, $c\equiv$ speed of light, $m_e\equiv$ electron mass, $M_{BH}\equiv$ black hole mass, $h\equiv$ Planck's constant, $\sigma_T\equiv$ Thompson scattering cross-section, and $T\equiv$ temperature.

\label{diss1} The contribution of the microphysics to the heating and cooling rates is displayed in Fig. \ref{heatAG1}. For instance, at $0.32$ pc we may see from Fig. \ref{heatAG1}(a) that the main contributor to the heating over the temperature range $400\lesssim T \lesssim 8000$ K is the Unresolved Transition Array (UTA, \cite{behar2001b,netzer2004a} and also see \cite{ramirez2008b} for an observational point of view), which accounts for $\approx 12-16$\% of the heating rate. In the interval $10^4\lesssim T \lesssim 8\times 10^4$ K, photoionization heating of O$^{+7}$ becomes the major contributor, providing from $\sim 12$ to 16\% of the heating. At temperatures of $10^5\lesssim T \lesssim 3.2\times 10^5$ K, Fe$^{+18}$ contributes with $\approx 12-17$\%. In these plots, the solid line labels the main contributors, while the dotted and dashed lines depict the second and third contributors, respectively. In Figs. \ref{heatAG1}(a)--(d), we see that in the temperature range $6.3\times 10^6\lesssim T \lesssim 3.2\times 10^9$ K, heating by Compton processes dominate the heating with contributions that rise up to $\approx 100$\% close to the upper extreme of the temperature range. In general, we find a rather complex interplay between the different heating agents, where low-ionization species contribute mostly at low-to-intermediate temperatures, while highly ionized species of heavy metals (e.g., Fe$^{+17}$-Fe$^{+24}$) and intermediate heavy metals in the form of H- and He-like (e.g., O$^{+7}$-O$^{+8}$, C$^{+4}$-C$^{+5}$) become important at temperatures in the range $10^4\lesssim T \lesssim 10^6$ K. At $T\gtrsim 10^6$ K, Compton heating becomes the dominant mechanism. The dashed-dotted lines in Figs. \ref{heatAG1}(a)--(d) depict the contribution of the 100 main heating agents, which clearly account for $\approx 100$\% of the total heating rate. A similar analysis can be done for the cooling rate. In Fig. \ref{coolAG1}, we depict the cooling rates as a function of the temperature at different distances from the source. We see that radiative recombination cooling by H contributes to $\approx 25-74$\% of the total cooling rate in the temperature range $200\lesssim T \lesssim 4\times 10^4$ K. Cooling by H lines dominates at higher temperatures in the interval $5\times 10^4\lesssim T \lesssim 4\times 10^5$ K with a contribution to total cooling of $\approx 25-74$\%. Free-free cooling contributes with up to $\approx 86$\% in the temperature interval between $\sim 1.3\times 10^5$ and $\sim 1.3\times 10^8$ K, with its contribution decreasing to $\approx 72$\% at $T\sim 10^9$ K. The second and third contributors are represented by the dotted and dashed lines, respectively, while the dashed-dotted lines depict the contribution of the 100 main cooling agents. Although the analytical formulas given by \cite{blondin1994a} are useful to study cooling and heating in high-mass X-ray binary systems, they become different for simulations of gas accretion onto SMBHs in the center of galaxies, if an accretion disk emission component and an expansion model are adopted. For instance, \cite{barai2011a} discuss in detail three-dimensional SPH simulations of accretion onto a SMBH, using the heating and cooling rates proposed by \cite{blondin1994a} \citep[see also][for an Eulerian simulation]{mproga2013a}. Some of their runs take longer to reach a steady state compared to the Bondi accretion. When analyzing radiative properties in the $T-\xi$ plane, they find many particles following the equilibrium temperature ($\mathcal{L}=0$) and discuss where and when artificial viscosity plays a dominant role over radiative heating. Below $T=10^4$ K, SED1 and SED2 cooling rates differ by factors of a few. In fact, neutral-to-middle ionized gas contributes mostly to the total cooling below $T=10^4$ K and this can be important in the outflows ($\sim 1400$ km/s) of {\sc n~iii}/{\sc n~iii}$^*$-{\sc s~iii}/{\sc s~iii}$^*$ found at $\sim$ 840~pc \citep{chamberlain2015a}, and also in closer Iron low-ionization broad absorption lines (FeLoBAL) flows ($v\sim$ few thousands \kms) at $\sim 7-70$ pc \citep{mcgraw2015a}. Moreover, at $\log_{10}(\xi)\sim 2$ the SED2 computations may overestimate the equilibrium temperature up to factors of $\sim 20$ in the range $10^5-10^8$ K. This may be used as a discrimination feature for simulations of SEDs in AGNs. For the heating case, however, the differences only reach factors of $\sim 2$ to $\sim 10$. We note that Compton and Coulomb heating could well be operating at temperatures between $10^8$ and $10^{10}$ K \cite[for instance, in pre/post shocked winds in AGNs,][]{FaucherGiguere2012}. In addition, pure 10 keV bremsstrahlung {\sc nowind} heating and cooling differ less from our calculations as the gas approaches the BH. We further note that \cite{vignali2015a} found a gas of high velocity ($\sim 0.14$c) through the identification of highly ionized species of Iron (e.g Fe~{\sc xxv} and Fe~{\sc xxvi}) in a luminous quasar at $z \sim 1.6$ located at distances of $\sim 10^{15}$--$10^{16}$ cm. In fact, through observed high-energy features, \cite{tombesi2015a} relate low- (by molecules) and high-velocity (highly ionized gas) with the predicted energy conserved wind \citep{FaucherGiguere2012}, and locate this gas at $\sim$ 900 $r_{\rm Sch}$, where more precise estimates of the heating and cooling are required. It is therefore clear that a quantitative analysis of the heating and cooling agents operating on these kinds of astrophysical environments are key aspects towards the understanding of the radiation hydrodynamical processes governing the accretion onto SMBHs. We have provided the files {\tt my1Part\_OUT.het} and {\tt my1Part\_OUT.col} as part of the tables, where the default $\approx 10$ agents are given by \cly. The interested reader may request the modified 100 agent files to the corresponding author. These tables have the potential and the flexibility to include other physical effects, like dust and/or molecules. In fact, OH 119 $\mu m$ lines have been found in ultraluminous infrared galaxies (ULIRGs) using the Herschel/PACS telescope at velocities of $\sim 1000$ \kms \sp \citep{veilleux13a}. Also, far-ultraviolet features may be present in Mrk 231 (found with the HST), with velocities of $\sim 7000$ \kms \sp \citep{veilleux16a}. They permit to make more extensive exploration about the influence of the SED on photoionization calculations \citep[see][]{chakravorty2009a,chakravorty2012a} and their impact on the energy and velocity distribution on hydrodynamical accretion processes onto SMBH. Another branch of SED to be explored are those including the reflected spectrum from the accretion disk, a rich mix of radiative recombination continua, absorption edges and fluorescent lines \citep{gar10,gar11,gar13a}. Additionally, if produced close to the black hole, this component suffers alterations due to relativistic effects \citep{dau13,gar14}. These types of SEDs may influence cooling and heating rates as they are very sensitive to the values of the ionization parameter, temperature, and density. In astrophysical ambients like the center of AGNs, they may play an important role in high velocity winds and evolutionary stages of the host galaxies \citep[as they may expel the cold gas reservoirs within $10^6-10^8$ years,][]{sturm11a}. We also are in capacity to include them in tables of radiative acceleration for SPH codes, and will be the subject of a future study. A strict comparison between theoretical models and simulations is beyond the scope of the tables presented here. At present, such simulations are under preparation.

1607.07785

1607

1607.07266_arXiv.txt

We clarify topology of $^3P_2$ superfluids which are expected to be realized in the inner cores of neutron stars and cubic odd-parity superconductors. $^3P_2$ phases include uniaxial/biaxial nematic phases and nonunitary ferromagnetic and cyclic phases. We here show that all the phases are accompanied by different types of topologically protected gapless fermions: Surface Majorana fermions in nematic phases and a quartet of (single) itinerant Majorana fermions in the cyclic (ferromagnetic) phase. Using the superfluid Fermi liquid theory, we also demonstrate that dihedral-two and -four biaxial nematic phases are thermodynamically favored in the weak coupling limit under a magnetic field. It is shown that the tricritical point exists on the phase boundary between these two phases and may be realized in the core of realistic magnetars. We unveil the intertwining of symmetry and topology behind mass acquisition of surface Majorana fermions in nematic phases.

1607.07266

1607

1607.03579_arXiv.txt

This review gives an introduction to spectrometers and discusses their use within radio astronomy. While a variety of technologies are introduced, particular emphasis is given to digital systems. Three different types of digital spectrometers are discussed: autocorrelation spectrometers, Fourier transform spectrometers, and polyphase filterbank spectrometers. Given their growing ubiquity and significant advantages, polyphase filterbanks are detailed at length. The relative advantages and disadvantages of different spectrometer technologies are compared and contrasted, and implementation considerations are presented.

\label{sec:intro} A \emph{spectrometer} is a device used to record and measure the spectral content of signals, such as radio waves received from astronomical sources. Specifically, a spectrometer measures the power spectral density (PSD, measured in units of $\rm{WHz}^{-1}$) of a signal. Analysis of spectral content can reveal details of radio sources, as well as properties of the intervening medium. For example, spectral line emission from simple molecules such as neutral hydrogen gives rise to narrowband radio signals (Fig.~1), while continuum emission from active galactic nuclei gives rise to wideband signals. There are two main ways in which the PSD --- commonly known as \emph{power spectrum} --- of a signal may be computed. The power spectrum, $S_{xx}$, of a waveform and its autocorrelation function, $r_{xx}$, are related by the Wiener-Khinchin theorem. This theorem states that the relationship between a stationary (mean and variance do not change over time), ergodic (well-behaved over time) signal $x(t)$, its PSD, and its autocorrelation is given by \begin{equation} S_{xx}(\nu)=\int_{-\infty}^{\infty}r_{xx}(\tau)e^{-2\pi i\nu\tau}d\tau. \label{eq:psd} \end{equation} where $\nu$ represents frequency, and $\tau$ represents a time delay or `lag'. The autocorrelation function is \begin{equation} r_{xx}(\tau)=\left\langle x(t)x(t-\tau)\right\rangle, \end{equation} where angled brackets refer to averaging over time. \begin{figure}[t] \centering \includegraphics[width=\textwidth]{./figures/hydrogen.pdf} \label{fig:hydrogen} \caption{A galactic hydrogen 21-cm line emission profile, as measured using the Breakthrough Listen digital spectrometer system on the Robert C. Byrd Green Bank Telescope in West Virginia.} \end{figure} Eq.~\ref{eq:psd} shows that that the autocorrelation function is related to the PSD by a Fourier transform. In the discrete case, the relationship becomes \begin{equation} S_{xx}(k)=\sum_{k=-\infty}^{\infty}\left\langle x(n)x(n-k)\right\rangle e^{-2\pi ink},\label{eq:discrete-wiener} \end{equation} which may be recognized as a discrete convolution. It follows from the convolution theorem that \begin{equation} S_{xx}(k)=\left\langle \left|X(k)\right|^{2}\right\rangle ,\label{eq:discrete-pow} \end{equation} where $X(k)$ denotes the Discrete Fourier Transform (DFT) of $x(n)$: \begin{equation} X(k)=\sum_{n=0}^{N-1}x(n)e^{-2\pi ink/N}\label{eq:dft1} \end{equation} with $N \rightarrow \infty$. There are therefore two distinct classes of spectrometers: ones that approximate $S_{xx}(k)$ by firstly forming the autocorrelation, then taking a Fourier transform à~la Eq.~\ref{eq:discrete-wiener}, and those that first convert into the frequency domain to form $X(k)$ before evaluating Eq.~\ref{eq:discrete-pow}. These two routes are shown diagrammatically in Fig.~\ref{fig:wiener}. We will refer to these as autocorrelation spectrometers (ACS, Sec.~\ref{sub:acs}), and Fourier transform filterbanks (FTF, Sec.~\ref{sub:ftf}), respectively. Polyphase filterbank spectrometers (PFB, Sec.~\ref{sub:pfb}) can be thought of as an FTF with enhanced filter response. Note that because the DFT is an \emph{approximation} to the continuous Fourier Transform, ACS and FTF systems have different characteristics. \begin{figure}[t] \centering \includegraphics[width=0.9\textwidth]{./figures/wiener} \caption{The two methods used to compute the PSD of a signal. The top path corresponds to an ACS system while the bottom corresponds to an FTF system. The two approaches are related by the Wiener-Khinchin theorem. \label{fig:wiener} } \end{figure} \subsection{Analysis and synthesis filterbanks} It is important to note the relationship between spectrometers, filters, and filterbanks. A \emph{filterbank} is simply an array of band-pass filters, designed to split an input signal into multiple components, or similarly, to combine multiple components. These are referred to as \emph{analysis} and \emph{synthesis} filterbanks, respectively. When applied to streaming data, a DFT can be considered an analysis filterbank, and an inverse DFT to be a synthesis filterbank. From this viewpoint, a spectrometer is simply an analysis filterbank, where the output of each filter is squared and averaged. \subsection{Polarimetry} Polarization is a key measurement within radio astronomy.\citet{BookTinbergenPolarim} Although most astrophysical radio emission is inherently unpolarized, a number of radio sources --- such as pulsars and masers --- do emit polarized radiation, and effects such as Faraday rotation by a galactic magnetic fields can yield polarized signals. A spectrometer that also measures polarization is known as a \emph{polarimeter} (or spectropolarimeter). \subsubsection{Stokes parameters} The Stokes parameters are a set of four quantities which fully describe the polarization state of an electromagnetic wave; this is what a polarimeter must measure. The four Stokes parameters, $I$, $Q$, $U$ and $V$, are related to the amplitudes of perpendicular components of the electric field: \begin{eqnarray} E_{x} & = & e_{x}(t)cos(\omega t+\delta_{x})\\ E_{y} & = & e_{y}(t)cos(\omega t+\delta_{y}) \end{eqnarray} by time averages of the electric field parameters: \begin{eqnarray} I & = & \left\langle E_{x}E_{x}^{*}+E_{y}E_{y}^{*}\right\rangle \\ Q & = & \left\langle E_{x}E_{x}^{*}-E_{y}E_{y}^{*}\right\rangle \\ U & = & \left\langle E_{x}E_{y}^{*}+E_{y}E_{x}^{*}\right\rangle \\ V & = & i\left\langle E_{x}E_{y}^{*}-E_{y}E_{x}^{*}\right\rangle \end{eqnarray} where $*$ represents conjugation. The parameter $I$ is a measure of the total power in the wave, $Q$ and $U$ represent the linearly polarized components, and $V$ represents the circularly polarized component. The Stokes parameters have the dimensions of flux density, and they combine additively for independent waves. \subsubsection{Measuring polarization products} In order to compute polarization products, a spectrometer must be presented with two voltage signals, $x(n)$ and $y(n)$, from a dual-polarization feed (i.e. a set of orthogonal antennas). With analogy to Eq.~\ref{eq:discrete-pow}, we may form \begin{eqnarray} S_{xx}(k) & = \langle X(k)X^*(k)\rangle & = \langle |X(k)|^2\rangle \label{eq:sxx1} \\ S_{yy}(k) & = \langle Y(k)Y^*(k)\rangle & = \langle |Y(k)|^2\rangle \\ S_{xy}(k) & = \langle X(k)Y^*(k)\rangle & \\ S_{yx}(k) & = \langle Y(k)X^*(k)\rangle & \label{eq:sxx4} \end{eqnarray} where in addition to measuring the PSD of $x(n)$ and $y(n)$, we also compute their cross correlations. Note that while $S_{xx}$ and $S_{yy}$ are real valued, $S_{xy}$ and $S_{yx}$ are complex valued. The four terms $\langle E_x E_x^*\rangle$, $\langle E_y E_y^* \rangle$, $\langle E_x E_y^* \rangle$, and $\langle E_y E_x^* \rangle$ are linearly related (by calibration factors) to the quantities in Eq.~\ref{eq:sxx1}-\ref{eq:sxx4} above. Combining these therefore allows for Stokes $I$, $Q$, $U$ and $V$ to be determined. In order to focus on the fundamental characteristics of spectrometers, the remainder of this chapter details single-polarization systems that compute only $S_{xx}$. Nevertheless, the techniques and characterization approaches are broadly applicable to polarimetry systems. \subsection{Performance characteristics} \begin{figure} \centering \includegraphics[width=\textwidth]{./figures/fb_comparison} \label{fig:pfb_response} \caption{Comparison of the channel response of an ACS (dotted line), FTF (dashed line) and an 8-tap, Hann-windowed PFB (solid line).\label{fig:leakage}} \end{figure} Spectrometers operate over a finite bandwidth $B$, over which $N$ channels with bandwidth $\Delta\nu = B/N$ are computed. With digital systems, channels may be evenly spaced with identical filter shapes. \subsubsection{Spectral leakage} Ideally, each channel would have unitary response over $\nu_c \pm \frac{\Delta\nu}{2}$, where $\nu_c$ is the center frequency, with zero response outside this passband. In practice, this cannot be achieved; each channel has a non-zero response over all frequencies. As such, a signal will `leak' between neighboring channels, known as spectral leakage. Fig.~\ref{fig:leakage} compares the normalized filter response for ACS, FTF and PFB implementations. In the presence of strong narrowband signals, such as radio interference (RFI), spectral leakage is a major concern. \subsubsection{Scalloping loss} A related concern is that a channel's non-ideal shape will cause narrowband signals at channel edges to be attenuated, an effect known as scalloping loss (Fig.~\ref{fig:scalloping}). Spectrometers are often designed such that neighboring channels overlap at their full-width at half-maximum points (FWHM), in which case the signal will be spread evenly over both channels. Wideband signals are not affected by scalloping. \begin{figure} \centering \includegraphics[width=\textwidth]{./figures/pfb_scalloping} \caption{Example of scalloping loss between spectrometer channels. The dashed lines show the response of individual channels, while the black line shows the overall response.\label{fig:scalloping}} \end{figure} \subsubsection{Time resolution}\label{sub:time-res} Time resolution refers to the minimum integration time over which a spectrometer averages over time. For a spectrometer with $N$ channels over a bandwidth $B$, the time resolution is $t_{\rm{res}} = 2B/(RN)$, where $R$ is the length of the averaging window. Detection of transient phenomena, such as fast radio bursts % and pulsars, require $t_{\rm{res}}$ to be as short as a microsecond, whereas integration lengths of several seconds, often averaged even further in post-processing, are common when observing faint sources. % \subsubsection{Dynamic range} Dynamic range refers to the span of input powers over which a spectrometer can operate nominally. The presence of RFI and the input bandwidth are the main drivers for dynamic range; see Sec.~\ref{sub:dynamic-range}.

1607.03579

1607

1607.04445_arXiv.txt

{Mass-loss rate is one of the most important stellar parameters. Mass loss via stellar winds may influence stellar evolution and modifies stellar spectrum. Stellar winds of subluminous hot stars, especially subdwarfs, have not been studied thoroughly.} {We aim to provide mass-loss rates as a function of subdwarf parameters and to apply the formula for individual subdwarfs, to predict the wind terminal velocities, to estimate the influence of the magnetic field and X-ray ionization on the stellar wind, and to study the interaction of subdwarf wind with mass loss from Be and cool companions.} {We used our kinetic equilibrium (NLTE) wind models with the radiative force determined from the radiative transfer equation in the comoving frame (CMF) to predict the wind structure of subluminous hot stars. Our models solve stationary hydrodynamical equations, that is the equation of continuity, equation of motion, and energy equation and predict basic wind parameters.} {We predicted the wind mass-loss rate as a function of stellar parameters, namely the stellar luminosity, effective temperature, and metallicity. The derived wind parameters (mass-loss rates and terminal velocities) agree with the values derived from the observations. The radiative force is not able to accelerate the homogeneous wind for stars with low effective temperatures and high surface gravities. We discussed the properties of winds of individual subdwarfs. The X-ray irradiation may inhibit the flow in binaries with compact components. In binaries with Be components, the winds interact with the disk of the Be star.} {Stellar winds exist in subluminous stars with low gravities or high effective temperatures. Despite their low mass-loss rates, they are detectable in the ultraviolet spectrum and cause X-ray emission. Subdwarf stars may lose a significant part of their mass during the evolution. The angular momentum loss in magnetic subdwarfs with wind may explain their low rotational velocities. Stellar winds are especially important in binaries, where they may be accreted on a compact or cool companion.} \keywords {stars: winds, outflows -- stars: mass-loss -- stars: early-type -- subdwarfs -- hydrodynamics} \titlerunning{Stellar wind models of subluminous hot stars} \authorrunning{J.~Krti\v{c}ka et al.}

Mass loss via stellar winds may influence the evolution of stars and determine their interaction with interstellar environment. Stellar wind also modifies the emergent spectrum and is therefore important for the diagnostics of stars. Radiatively driven stellar winds exist in many types of hot stars \citep[for a review]{pulvina} particularly in hot subluminous stars. The subluminous stars are in the late phases of their evolution and their luminosities are lower than those of corresponding main sequence stars \citep[e.g.,][]{sam1}. Hot subdwarfs are typical subluminous stars, which consist of a bare helium burning stellar core stripped of its envelope during the previous evolution \citep{durman}. It is not clear how a star may end up in such an evolutionary phase. There are more possible evolutionary channels that lead to different types of subluminous objects. Helium low-luminosity stars may originate as a merger of two white dwarfs \citep{iba,saje,zhaff} or in a late thermal pulse \citep{icko,argentinci}. Subluminous stars may be also products of red giants, which were stripped off their envelopes possibly during binary evolution \citep[e.g.,][]{jinanvoxfordu}. Hot subdwarfs are frequently members of binaries. This may be connected with their evolutionary state. Subdwarfs are frequently accompanied by various objects, including white dwarfs, late type stars or substellar objects and, in a rare cases, Be stars \citep[e.g.,][]{dvoj4,dvoj3}. There is growing observational interest in winds of subluminous stars. Ultraviolet (UV) wind line profiles of central stars of planetary nebulae may be used together with other observables to determine the stellar parameters \citep{btpau}. Also subdwarf O (sdO) stars show signatures of wind in the ultraviolet spectral region \citep{jefham}. The X-ray emission of subdwarf stars is likewise connected with their winds and follows a similar trend as the X-ray emission of O stars \citep{bufacek}. The wind may be accreted on a compact companion leading to X-ray sources similar to high-mass X-ray binaries \citep{dvoj18}. In binaries consisting of a subdwarf star and a compact object, the missing X-ray emission may provide an upper limit for the wind mass-loss rate \citep{merne}. \citet{vinca} predicted the mass-loss rates for hot subdwarfs and discussed evolutionary and spectroscopic consequences of these winds. However, these models did not include the hottest subdwarfs and did not predict the terminal velocities. \citet{un} provides independent predictions of mass-loss rates and discussed the role of the wind in the radiative diffusion. Although these models covered a broader range of stellar parameters, the calculations were based on line force multipliers neglecting, for example, the finite disk factor i.e., they assumed the star is a point source of radiation. The physics of the wind of hot subluminous stars is generally complex. These winds are, to some extent, similar to the stellar wind of main-sequence B~stars, where the effects of multicomponent flow and inefficient shock cooling may be important \citep{krtzatmeni,votzameni}. To improve the theoretical description of stellar winds of subluminous hot stars we here provide their wind models, predicting the basic wind parameters. Moreover, we study the effects that have not yet been discussed in the context of stellar winds of subluminous stars. For example, the X-ray irradiation in binaries with a compact component that may affect the wind accretion, or the influence of magnetic fields that may lead to rotational braking. We also discuss the properties of the winds of individual stars that were not available in the literature.

1607.04445

1607

1607.01056_arXiv.txt

We present late-time {\it Hubble Space Telescope} ({\it HST}) images of the site of supernova (SN) 2009ip taken almost 3 yr after its bright 2012 luminosity peak. SN~2009ip is now slightly fainter in broad filters than the progenitor candidate detected by {\it HST} in 1999. The current source continues to be dominated by ongoing late-time CSM interaction that produces strong H$\alpha$ emission and a weak pseudo-continuum, as found previously for 1-2 yr after explosion. The intent of these observations was to search for evidence of recent star formation in the local ($\sim$1kpc; 10{\arcsec}) environment around SN~2009ip, in the remote outskirts of its host spiral galaxy NGC~7259. We can rule out the presence of any massive star-forming complexes like 30 Dor or the Carina Nebula at the SN site or within a few kpc. If the progenitor of SN~2009ip was really a 50-80 $M_{\odot}$ star as archival {\it HST} images suggested, then it is strange that there is no sign of this type of massive star formation anywhere in the vicinity. A possible explanation is that the progenitor was the product of a merger or binary mass transfer, rejuvenated after a lifetime that was much longer than 4-5 Myr, allowing its natal H~{\sc ii} region to have faded. A smaller region like the Orion Nebula would be an unresolved but easily detected point source. This is ruled out within $\sim$1.5 kpc around SN~2009ip, but a small H~{\sc ii} region could be hiding in the glare of SN~2009ip itself. Later images after a few more years have passed are needed to confirm that the progenitor candidate is truly gone and to test for the presence of a small H~{\sc ii} region or cluster at the SN position.

The class of Type~IIn supernovae (SNe~IIn hereafter), whose narrow H lines indicate strong interaction with dense circumstellar material (CSM), have challenged our understanding of stellar evolution and death. Their dense H-rich CSM, progenitor instability, and high initial masses inferred from various clues suggest a link to the class of luminous blue variables (LBVs), which are not supposed to be anywhere near core collapse in the standard scenario of massive single-star evolution (see \citealt{smith14} for a general review). Among well-studied examples of SNe~IIn, the explosion of SN~2009ip in mid-2012 (note that its discovery in 2009 was deemed to be a SN impostor) is one of the most interesting, with by far the best observational characterization of a directly detected progenitor among any SN in history (even SN~1987A). It had a (presumably) quiescent progenitor star candidate detected in archival {\it Hubble Space Telescope} ({\it HST}) images, with a luminosity that implied a very high initial mass of at least 50-60 $M_{\odot}$ \citep{smith10,foley11}. This source also showed a series of outbursts in the few years before the SN \citep{smith10,pastorello13} that were reminiscent of both S~Dor outbursts and giant eruptions of LBVs. Unlike any progenitor source so far, high-quality spectra of these precursor outbursts were obtained, with a detailed analysis before the SN indicating a strong similarity to LBVs \citep{smith10,foley11}.\footnote{Lower-quality, low-resolution photographic spectra were identified for SN~1987A \citep{walborn89}, but so far this is the only other SN with a progenitor spectrum.} With a quiescent and very luminous progenitor, an S~Dor outburst, several bright but brief SN impostor eruptions, and progenitor spectra resembling LBVs, SN~2009ip provides a strong link between LBVs and SNe~IIn. The repeating variable source at the position of SN~2009ip began to brighten again in mid-2012, but this time things were different. Spectra of the fainter 2012a peak showed very broad P Cygni profiles with velocities of 13,000 km s$^{-1}$, suggesting that the event was a core-collapse SN and not another LBV outburst \citep{sm12,mauerhan13}. The subsequent and brighter 2012b event showed a high peak luminosity and a spectrum typical of SNe~IIn with strong CSM interaction. The 2012 SN-like event has already been discussed extensively in the literature \citep{mauerhan13,mauerhan14,pastorello13,prieto13, fraser13,fraser15,ofek13,raf14,smith13,smp14,graham14}. In these publications and in discourse at meetings, there was some uncertainty and controversy about whether the 2012 event was a true core-collapse SN, since (1) CSM interaction can provide bright transients even from relatively low-energy explosions, (2) the initial SN was somewhat fainter than standard SNe~II-P, and (3) the rich observational dataset for the progenitor presented mysteries that were not easily explained by any existing model. These are, however, expressions of the challenge in understanding SNe~IIn and CSM interaction, rather than arguments against a core-collapse event. While it is difficult to prove definitively that the event was a core collapse because of the masking of CSM interaction, a SN is the most straightforward explanation of the data. \cite{smp14} showed that all available evidence was consistent with the core collapse SN explosion of a blue supergiant that encountered strong CSM interaction. Moreover, both line-profile evolution \citep{smp14} and spectrapolarimetry \citep{mauerhan14} show that the CSM interaction was highly aspherical and probably disk-like, forcing the kinetic energy budget of the event to be $\sim$10$^{51}$ ergs. \citet{emily14} also argued for a disk-like CSM based on narrow line ratios. \citet{fraser15} showed that the source at +2 yr was consistent with steady ongoing CSM interaction with no additional outbursts, adding further evidence in favor of a core-collapse event. SN~2009ip provides our clearest example of pre-SN instability that leads to eruptive pre-SN mass loss in the few years before explosions, which may be associated with the final nuclear burning sequences in the last years of a massive star's life \citep{qs12,sa14}. Alternative non-terminal models involving binary mergers and accretion were also proposed for the 2012 event \citep{soker13,kashi13}, but these cannot supply the required 10$^{51}$ ergs of kinetic energy. In this paper, we are mainly concerned with the host galaxy environment around SN~2009ip. A fundamental interesting mystery was that while progenitor detections pointed to a very massive unstable star, the location of SN~2009ip was in the remote outskirts of its spiral host, far away from obvious signs of recent star formation and young stellar populations \citep{smith10,foley11,mauerhan13,raf14}. It was located about 5 kpc from the center of its relatively small host spiral galaxy NGC~7259. By extrapolating the apparent metallicity gradient measured in the inner $\sim$1.5 kpc out to the 5 kpc radius of SN~2009ip, \citet{raf14} infer a mildly subsolar metallicity at the SN site of $0.4 < Z/Z_{\odot} < 0.9$. SN~2009ip's progenitor can therefore be compared with populations of massive stars observed in the Milky Way and Large Magellanic Cloud (LMC). An interesting result is that despite their high luminosities and high inferred initial masses, LBVs in the Milky Way and LMC appear to be relatively isolated compared to expectations for their presumed role in stellar evolution. \citet{st15} demonstrated that LBVs selectively avoid clusters of O-type stars, especially early O-types that are their presumed progenitors. More importantly, LBVs are more dispersed on the sky than WR stars; this rules out the standard picture wherein LBVs are a transitional phase between massive O-type stars and WR stars. Instead, \citet{st15} suggested that most LBVs may be the result of interacting binary evolution, getting rejuvenated by either mass transfer or mergers. This would make them stand out as anomolously young compared to their surrounding populations. In other words, {\it they are evolved massive blue stragglers}. They may become even more isolated upon receiving a kick from their companion star's SN, although it is not yet clear if a kick is required to explain their environments. Also relevant to this story is that SNe~IIn in general appear to be less correlated with bright H$\alpha$ in their host galaxies than other types of SNe \citep{anderson12,habergham14}. While this has been interpreted as signifying lower initial masses by those authors, that interpretation has been a topic of debate \citep{crowther13,st15}. As described below, SN~2009ip seems to follow a similar trend of not having bright H$\alpha$ nearby, despite having a detection of a very massive and luminous progenitor star. \begin{table} \begin{center}\begin{minipage}{3.3in} \caption{New {\it HST} WFC-UVIS Images of SN~2009ip, including ST magnitudes for SN~2009ip and background 3$\sigma$ upper limits}\scriptsize \begin{tabular}{@{}lccccc}\hline\hline Date &Filter &Exp.(s) &Mag &1$\sigma$ &3$\sigma$ U.L. (mag) \\ \hline 2015 May 25 &F275W &2900 &22.53 &0.05 &25.0 \\ 2015 May 23 &F555W &1650 &21.91 &0.01 &26.8 \\ 2015 May 25 &F657N &5911 &18.97 &0.003 &25.2 \\ 2015 May 23 &F814W &1100 &22.91 &0.01 &27.0 \\ \hline \end{tabular}\label{tab:phot} \end{minipage}\end{center} \end{table} \begin{figure} \includegraphics[width=3.2in]{fig2.eps} \caption{The photometry of the F606W progenitor candidate in 1999 (orange square) and the photometry of the source detected in new {\it HST} images in 2015 (pink circles). These are compared to Starburst99 models as discussed in the text for 5, 10, and 20 Myr (gray, blue, and black, respectively). Models with these ages and no reddening (solid) are scaled to the F275W point, and the same models with reddening (dashed/dotted) are scaled to match the F275W and F814W points. See Table~\ref{tab:s99} and text.} \label{fig:sed} \end{figure}

1607.01056

1607

1607.04997_arXiv.txt

It is believed that, in radio-loud active galactic nuclei (AGN), the core radio flux density can be normalized to the flux density of the extended lobe emission to infer the orientation of a radio source. However very little is known about the reliability and precision of this method, and we are unaware of any robust conversion recipe to infer the inclination from the core dominance. Investigating whether or not the radio core dominance parameter R separates the quasars from the radio-galaxies in the $z \ge$~1 3CRR catalog, we found excellent agreement of R with optical type, infrared flux ratios and optical polarization. This indicates that probably both R and optical classification are very good orientation indicators, and the unified model is strongly predictive for these objects. The relative number densities indicate half-opening angles close to 60$^\circ$, as expected from large surveys. The separations of optical types according to radio core dominance as well as NIR/MIR ratios, which are essentially perfect, means that there can be only a small dispersion of torus half-opening angles. Also, even though torus dust is thought to be clumpy, there is an almost zero probability to see a type-1 source at high inclination. Finally, using only the Copernican Principle, i,e, the assumption that solid angle is filled uniformly with source axis orientations, we estimated a semi-empirical relation between core dominance and AGN inclination. This makes it possible to use R to infer the inclination of a source to an accuracy of $\sim$ 10 degrees or less, at least for this type of object.

\label{Intro} The nuclear inclination of active galactic nuclei (AGN) is a fundamental parameter that still remains tricky to directly measure. In radio-quiet objects, proxies for measuring the orientation of AGN usually rely on degenerate models and fitting procedures that sometimes give inconsistent results \citep{Marin2014,Marin2016}. However, in the case of radio-loud AGN, in which there is beamed emission from pc-scale jets and unbeamed emission from diffuse radio lobes, we have an attractive method since the ratio of the former to the latter (the radio core dominance R) should be statistically a function of inclination with respect to the line of sight \citep{Orr1982}. Operationally, it works well because on arcsec-scale maps, double radio sources generally show an unresolved and fairly isolated point source, usually with a flat synchrotron spectrum implying extreme compactness (pc scales). Furthermore selection by the nearly-isotropic lobe emission assures that a sample has a random orientation distribution, filling all solid angles as seen from the nucleus. This core dominance parameter is usually estimated in the radio because, in a general way, the jet and lobe radiation are linked physically. But (for quasars only), the continuum optical flux density \citep{Wills1995} has also been used to normalize the beamed radio core flux density, instead of the flux density of the extended radio lobes. This core dominance parameter is called R$_{\rm V}$. Its usefulness has been confirmed by \citet{Barthel2000}. Narrow line luminosities have also been tried \citep{Rawlings1991}. Very little is known about the reliability and precision of R as a statistical indicator of inclination. One way to test this is to see to what extent it separates radio galaxies (RG) from quasars, since we know that among the high-$z$ 3CR radio galaxies, all RG host hidden quasars, and that they generally have smaller inclinations. We don't know a priori whether or not a lot of noise will be added by a distribution of torus half-opening angles, "holes" in the torus, or dispersion in the radio core emission polar diagrams. But if core dominance correlates very well with object type, then the most natural conclusion is that the sources are roughly generic in nature, core dominance indicates inclination well, and there are very few quasars at high-inclinations. This is in fact what we will show\footnote{A clean separation of optical types by core dominance was reported by \citet{Wilkes2013} in a sample similar to ours.}. The unified model in its simplest form, the "Straw Person Model" of \citet{Antonucci1993}, seems to be correct to first order in this parameter space, as predicted by \citet{Barthel1989}. That is, the model shows very strong predictive power. The same conclusion follows from a detailed study of the X-ray properties of the sample by \citet{Wilkes2013}. High core dominance is known to be associated with low inclination and a flat IR SED; it's also known that radio galaxies have lower core dominance than quasars; we will show that limited optical polarization is also consistent with this idea when available. The purpose of this paper is to note and exploit the fact that there is essentially zero overlap in R$_{\rm 5GHz}$ between RGs and quasars, using that fact and the special sample properties to derive a quantitative formula for converting core dominance to inclination, estimating that it is good to plus or minus 10 degrees. To obtain intelligible results, it is fundamental to select a sample were all RG have hidden quasars, and to select by an approximately isotropic luminosity (the radio lobes luminosity here). The high redshift (i.e. $z \ge$~1) radio-loud AGN from the complete flux-limited, 178~MHz selected, ``3CRR'' catalog of \citet{Laing1983} meet this requirement very well. We know this for several reasons. First, most these objects have been observed polarimetrically and all show substantial optical polarization, oriented roughly perpendicular to the radio axes \citep{Tadhunter2005}. They also show broad emission lines in polarized flux when that information is available\footnote{The case of 3C~368 is controversial; \citet{Dey1999} finds low aperture polarization, but a convincing image of centro-symmetric polarization is shown by \citet{Scarrott1990}.}. Also, all of these radio galaxies for which data is available show strong mid-infrared reprocessing bumps \citep{Honig2011} and high ionization emission lines \citep{Leipski2010}. It is important to note that this sample comprises many of the most radio-luminous objects in the universe, and at lower radio luminosities many radio galaxies lack energetically significant hidden quasars \citep{Ogle2006}; the evidence for that statement from all wavebands is reviewed in detail in \citet{Antonucci2012}.

\label{Conclusion} We have demonstrated that the radio core dominance parameter R$_{\rm 5GHz}$ separates the radio-galaxies and quasars almost perfectly in the $z \ge$~1 3CRR sample. In other words, the agreement of R$_{\rm 5GHz}$ with optical type is not a coincidence, and probably both R$_{\rm 5GHz}$ and optical type are reliable orientation indicators. Since the 3CRR catalog is a complete sample, it must fill the solid angle uniformly, except for small number statistics, just from the Copernican Principle. Therefore we were able to derive an empirical core dominance formula, where R$_{\rm 5GHz}$ is a function of inclination $i$. The essentially perfect separation of the optical types by radio core dominance and infrared (and X-ray columns) is most simply interpreted as meaning that there is only a small dispersion of opening angle, and no holes in the circumnuclear region which would let us see a type-1 AGN at high inclination (contrary to the hypothesis of \citealt{Obied2016}). At lower luminosities, such predictions are less clear as the sample gets less clean. It is possible to extend the catalog to lower redshifts but it would be necessary to exclude nonthermal RG. Weakly-accreting radio-galaxies mostly radiate through kinetic energy in the form of synchrotron jets and lack highly ionized line emission and strong IR reprocessing bumps, which indicates that they lack energetically significant hidden quasars (\citealt{Antonucci2012} and references therein). It is then essential to remove them to have a correct sample. At higher redshifts another problem may arise due to the fact that the ratio of lobe emission to jet power is sensitive to environment \citep{Barthel1996}. The density of the intergalactic medium (IGM) scales with redshift as (1 + $z$)$^3$ \citep{Macquart2013}. For a $z$ = 5 quasar, the IGM density is 27 times higher than compared to a $z$ = 1 quasar, and these should strongly affect the morphology, lobe flux ratios, and lobe distance.

1607.04997

1607

1607.03115_arXiv.txt

The so-called ``dipper" stars host circumstellar disks and have optical and infrared light curves that exhibit quasi-periodic or aperiodic dimming events consistent with extinction by transiting dusty structures orbiting in the inner disk. Most of the proposed mechanisms explaining the dips---i.e., occulting disk warps, vortices, and forming planetesimals---assume nearly edge-on viewing geometries. However, our analysis of the three known dippers with publicly available resolved sub-mm data reveals disks with a range of inclinations, most notably the face-on transition disk J1604-2130 (EPIC~204638512). This suggests that nearly edge-on viewing geometries are {\em not} a defining characteristic of the dippers and that additional models should be explored. If confirmed by further observations of more dippers, this would point to inner disk processes that regularly produce dusty structures far above the outer disk midplane in regions relevant to planet formation.

Understanding how planets form is one of the most compelling problems in astronomy. Close-in planets appear to be common \citep{2010Sci...330..653H,2013ApJ...770...69P,2015ApJ...799..180S} and we can study their nascent systems (i.e., protoplanetary disks) in young stellar associations. However, observing planet formation at $\lesssim1$~AU is complicated by small angular scales and faint disk emission compared to the host star; for the nearest star-forming regions, the best achievable angular resolution for optical/infrared scattered light and sub-mm images is only a few AU. However, new probes of the inner disk may be the so-called ``dipper" stars, whose optical and infrared light curves exhibit episodic drops in flux consistent with extinction by transiting dusty structures orbiting with Keplarian periods of down to a few days. Dipper stars were identified with the CoRoT and {\it Spitzer} missions in the young ($\sim$2--3~Myr) Orion Nebula Cluster \citep{2011ApJ...733...50M} and NGC~2264 region \citep{2010A&A...519A..88A,2014AJ....147...82C}. These studies found that the depth, duration, and periodicity of the dips are consistent with extinction by dust orbiting near the star-disk co-rotation radius. \cite{2015A&A...577A..11M} proposed that the dippers in NGC~2264 could be explained by occulting inner disk warps driven by accretion streams onto the host star, as reported for AA Tau \citep{1999A&A...349..619B}. Unfortunately, the significant distances to these clusters ($\sim$400-750~pc) limited follow-up observations. \cite{2016ApJ...816...69A} identified $\sim$25 dippers in the young ($\lesssim10$~Myr), nearby ($\sim$120--145~Myr) Upper Sco and $\rho$ Oph star-forming regions using high-precision optical photometry from the K2 mission \citep{2014PASP..126..398H}. Follow-up observations showed that the K2 dippers are often weakly accreting late-type stars hosting moderately evolved primordial disks, challenging the disk warp scenario described above. Thus \cite{2016ApJ...816...69A} proposed alternative mechanisms to explain the dips, namely occulting vortices at the inner disk edge produced by the Rossby Wave Instability and transiting clumps of circumstellar material related to planetesimal formation. \cite{2016arXiv160503985B} also found that magnetospheric truncation of weakly accreting disks with misaligned magnetic fields could form occulting accretion streams that produce dippers with high to moderate inclinations. The proposed mechanisms for explaining the dips therefore assume geometries that are closer to edge-on than face-on. The occurrence rate of dippers in co-eval clusters (e.g., $\sim$20--30\% of classical T Tauri stars in NGC~2264; \citealt{2010A&A...519A..88A}, \citealt{2014AJ....147...82C}) and the moderate optical extinction towards these objects \cite[$A_{\rm V}\lesssim1$;][]{2016ApJ...816...69A} suggest that we are not seeing the disks directly edge-on ($i=90$\degrees{}), but rather viewing the dippers at nearly edge-on inclinations (e.g., $i=70$\degrees{}) and thus observing transits of dusty structures lifted above the disk midplane. However, these geometric assumptions have not yet been compared to observations. In this Letter, we present the three known dippers whose circumstellar disks have been resolved in archival sub-mm data such that their inclinations can be constrained. Surprisingly, we find disks with a range of inclinations, most notably the face-on transition disk (TD) J1604-2130 \citep{2012ApJ...753...59M,2014ApJ...791...42Z}. This indicates that nearly edge-on disk inclinations are {\em not} a defining characteristic of the dippers, and motivates a re-examination of the dipper mechanisms so that we can properly interpret these objects in the context of planet formation. \begin{figure*} \begin{centering} \includegraphics[width=17cm]{Fig2.pdf} \caption{\small Top panels show ALMA sub-mm continuum images (5\as{}$\times$5\as{}) with fitted disk inclinations and beam sizes (Section~\ref{sec-alma} \& \ref{sec-inclinations}). Contours are 10$\sigma$ and 100$\sigma$ for EPIC~204638512 and 5$\sigma$, 20$\sigma$, and 50$\sigma$ for EPIC~205151387 and EPIC~203850058. Bottom panels show the real part of the visibilities as a function of projected baseline length.} \label{fig-alma} \end{centering} \end{figure*}

\label{sec-discussion} \subsection{Disk Inclinations\label{sec-inclinations}} For EPIC~204638512, we adopt the disk inclination of $i=6$\degrees{}$\pm$1.5\degrees{} derived by \cite{2012ApJ...753...59M}, who placed strong constraints on disk geometry using their $^{12}$CO first-moment map and assumptions of Keplerian rotation. For EPIC~205151387 and EPIC~203850058, we derive disk inclinations from their sub-mm continuum data (Section~\ref{sec-alma}) using standard routines from the {\it Common Astronomy Software Applications} (CASA) package \citep{2007ASPC..376..127M}; although their existing CO data were insufficient to derive precise disk inclinations, the first-moment maps can be used as rough checks on our continuum results. We derived the disk inclination of EPIC~205151387 using the CASA routine \texttt{uvmodelfit}, which fits simple analytic source component models (point-source, Gaussian, or disk) directly to the visibility data. We assumed an elliptical Gaussian model, which has six free parameters: integrated flux density ($F$), FWHM along the major axis ($a$), aspect ratio of the axes ($r$), position angle (PA), right ascension offset from the phase center ($\Delta\alpha$), and declination offset from the phase center ($\Delta\delta$). We found $F=49.5\pm0.4$~mJy and PA$=-22$\degrees{}$\pm$3\degrees{}, then derived the inclination from $r$ assuming circular disk structure, finding $i=53$\degrees{}$\pm$2\degrees{}. To check our results, we analyzed the first-moment $^{12}$CO map, finding a similar position angle (${\rm PA}\approx-30$\degrees{}) and inclination ($i\approx50$\degrees{}) when assuming Keplerian rotation. EPIC~203850058 is thought to have a high disk inclination based on the detection of an optical jet \citep{2005Natur.435..652W} and the nearly symmetric morphology of a bipolar outflow in molecular CO \citep{2008ApJ...689L.141P}, as previously noted by \cite{2012ApJ...761L..20R}. To derive an inclination, we again used \texttt{uvmodelfit} to fit an elliptical Gaussian model to the continuum visibility data, but with an initial guess of ${\rm PA}\approx20$\degrees{} based on the first-moment $^{12}$CO map in \cite{2012ApJ...761L..20R}. We found $F=4.0\pm0.1$~mJy and PA$=15$\degrees{}$\pm$1\degrees{}, consistent with \citealt{2012ApJ...761L..20R}. We then derived the inclination from $r$ assuming circular disk structure, finding $i\approx90$\degrees{}, but with very large errors. We therefore checked our results with the CASA routine \texttt{imfit}, which fits an elliptical Gaussian to the source in its image plane, then uses the clean beam to return de-convolved fit results; we found a nearly edge-on disk with $i=73$\degrees{}$\pm$23\degrees{}, $F=4.2\pm0.1$~mJy, and ${\rm PA}=$10\degrees{}$\pm$15\degrees{}. Although the small size and faint emission of this source complicates the analysis, the overall picture appears to point to a nearly edge-on viewing geometry for EPIC~203850058. Note that these sub-mm observations have resolutions of $\sim$20--50~AU in radius (Section~\ref{sec-alma}), thus the estimated inclinations reflect bulk disk geometry and assume a uniform inclination angle throughout the disk. Moreover, the uncertainties do not include systematic errors (e.g., we assume the observed disks of EPIC 205151387 and EPIC 203850058 are adequately represented by elliptical Gaussians). \subsection{A Call For Re-thinking Dipper Mechanisms\label{sec-mechanisms}} The proposed mechanisms explaining the dipper phenomenon favor geometries that are nearly edge-on. The disks are likely not seen completely edge-on ($i=90$\degrees{}), however, due to the moderate extinction towards the dippers \cite[$A_{\rm V}\lesssim1$;][]{2016ApJ...816...69A}. Rather, it is thought that we are viewing the dippers at nearly edge-on inclinations (e.g., $i=70$\degrees{}) and thus observing transits of occulting material lifted above the disk midplane by some process \cite[e.g., the breakdown of Rossby waves into vortices;][]{2016ApJ...816...69A}. Notably, none of the proposed dipper mechanisms can account for obscurations from face-on disks. Thus the surprising range of dipper disk inclinations presented in this work, in particular the face-on geometry of EPIC~204638512 (J1604-2130), suggests that nearly edge-on viewing geometries are {\em not} a defining characteristic of the dippers and motivates the exploration of alternative models (or combinations of models) that can explain a range of disk inclinations. For example, occulting accretion streams \cite[e.g.,][]{2015A&A...577A..11M,2016arXiv160503985B} could possibly account for even face-on outer disks if they act in concert with other mechanisms warping the inner disk, such as dynamical interactions with (proto-) planets or low-mass stellar companions \cite[e.g.,][]{2014MNRAS.442.3700F,2015ApJ...798L..44M}. Populations of scattered planetesimals from migrating (proto-) planets may also explain low dipper disk inclinations \citep{2011A&A...531A..80K}, but more work is needed to explore nearly polar orbits. \subsection{EPIC~204638512 (J1604-2130) \label{sec-epic512}} EPIC~204638512 is a particularly interesting case. This source hosts a face-on disk ($i=6$\degrees{}$\pm$1.5\degrees{}; \citealt{2012ApJ...753...59M}) with a large sub-mm dust cavity ($\sim$80~AU in radius; \citealt{2014ApJ...791...42Z}), which seemingly makes it an unlikely dipper. Yet, EPIC~204638512 exhibits the deepest flux dips among the known K2 dippers (up to $\sim$60\%; Figure~\ref{fig-lcs}). How can these characteristics be reconciled? The dipper activity may be related to an inclined and variable inner dust disk, as implied from its infrared emission. The object's {\it Spitzer} IRAC photometry shows no excess \citep{2012ApJ...753...59M}, while its {\it Spitzer} IRS spectrum and WISE photometry reveal excesses consistent with dust at small ($\lesssim$0.1~AU) orbital radii \citep{2010AJ....140.1444D,2014ApJ...791...42Z}. A factor of four variability in mid-infrared flux was also seen over several weeks, indicating a rapidly changing inner dust disk \citep{2009AJ....137.4024D}. Moreover, \cite{2014ApJ...795...71T} used near-infrared imaging polarimetry to identify intensity nulls in the outer disk annulus, which could be self-shadowing from a misaligned inner disk. An inclined transient inner disk has been proposed for HD~142527, which also hosts a face-on transition disk with a large inner dust gap \citep{2006ApJ...636L.153F} and exhibits intensity nulls along the outer disk annulus in its infrared scattered light images \citep{2012ApJ...754L..31C}. HD~142527 has a known inner disk, thought to be a transient feature of accretion from the outer disk \citep{2011A&A...528A..91V,2013Natur.493..191C}. \cite{2015ApJ...798L..44M} modeled the system, finding a relative inclination of $\sim$70\degrees{} between the inner and outer disks, possibly due to dynamical interactions with a low-mass stellar companion orbiting inside the dust gap \citep{2012ApJ...753L..38B,2014ApJ...791L..37R,2014ApJ...781L..30C}. A similar scenario for EPIC~204638512 would reconcile its dips and face-on outer disk. One indication of an inclined inner disk is EPIC~204638512's weak rotational signal (Section~\ref{sec-k2}), which suggests the star is pole-on and thus aligned with the outer disk. The dust cavity of EPIC~204638512 is also thought to have been cleared by giant planet(s) orbiting inside the dust gap \citep{2012ApJ...753...59M,2014ApJ...791...42Z,2015A&A...579A.106V}, which could drive an inner disk warp. However, giant planets alone likely cannot account for all the dippers, as dipper occurrence rates (e.g., $\sim$20--30\% in NGC~2264; \citealt{2010A&A...519A..88A}, \citealt{2014AJ....147...82C}) are much larger than giant planet occurrence rates around late-type stars \cite[e.g., a few percent; see Figure 8 in][]{2013ApJ...771...18G}.

1607.03115

1607

1607.00320_arXiv.txt

We determine the effect of intergalactic magnetic fields on the distribution of high energy gamma rays by performing three-dimensional Monte Carlo simulations of the development of gamma-ray-induced electromagnetic cascades in the magnetized intergalactic medium. We employ the so-called ``Large Sphere Observer'' method to efficiently simulate blazar gamma ray halos. We study magnetic fields with a Batchelor spectrum and with maximal left- and right-handed helicities. We also consider the case of sources whose jets are tilted with respect to the line of sight. We verify the formation of extended gamma ray halos around the source direction, and observe spiral-like patterns if the magnetic field is helical. We apply the $Q$-statistics to the simulated halos to extract their spiral nature and also propose an alternative method, the $S$-statistics. Both methods provide a quantative way to infer the helicity of the intervening magnetic fields from the morphology of individual blazar halos for magnetic field strengths $B \gtrsim 10^{-15}\,{\rm G}$ and magnetic coherence lengths $L_{\rm c} \gtrsim 100\,{\rm Mpc}$. We show that the $S$-statistics has a better performance than the $Q$-statistics when assessing magnetic helicity from the simulated halos.

The origin, strength and structure of intergalactic magnetic fields (IGMF) remain a mystery up to the present day. Possible mechanisms to explain cosmic magnetogenesis may be divided into two main categories: cosmological scenarios predict that magnetic fields were generated through processes taking place in the early universe, such as inflation \cite{PhysRevD.37.2743,Ratra:1991bn,Byrnes:2011aa,Ferreira:2014hma}, electroweak \cite{Vachaspati:1991nm,Enqvist:1993np,PhysRevD.53.662,Grasso:1997nx} or QCD phase transitions \cite{Hogan:1983zz,1989ApJ...344L..49Q,PhysRevD.55.4582,Tevzadze:2012kk}, and leptogenesis \cite{Long:2013tha}, among others; in astrophysical scenarios the fields would be created during the later stages of evolution of the universe, for example during structure formation~\cite{Kulsrud:1996km} or even thereafter \cite{1981ApJ...248...13L}. Measurements of IGMF are rather difficult due to their low magnitude. Common methods to estimate the strength of IGMF are indirect and include the well-known Faraday rotation measurements which yield upper limits of the order of a few nG~\cite{PhysRevD.80.123012}. Lower bounds, $B \gtrsim 10^{-17}\, \text{G}$, have been obtained by several authors using gamma-ray-induced electromagnetic cascades in the intergalactic space \cite{JETPLett.85.10.473,PhysRevD.80.123012,Essey:2010nd,2010ApJ...722L..39A,2010MNRAS.406L..70T,Taylor:2011bn,Takahashi:2013lba,2014arXiv1410.7717C}. These lower bounds are controversial because of the claims~\cite{0004-637X-752-1-22,0004-637X-758-2-102,Schlickeiser:2013eca,SavelievIGM,Chang:2014cta} that the development of the cascade is suppressed by plasma instabilities that arise from interactions with the intergalactic medium. On the other hand, recent direct observations of cascades~\cite{Chen11072015} suggest that plasma instabilities are not operative and that the original bounds hold. We expect that future analyses will clarify the role, if any, of plasma instabilities in the development of the electromagnetic cascade. Magnetic fields can carry helicity ($\mathcal{H}$), which is defined as \begin{equation} \mathcal{H} = \int \mathbf{A} \cdot \mathbf{B} \, {\rm d}^{3}r \,, \end{equation} where $\mathbf{A}$ is the magnetic vector potential and $\mathbf{B}=\nabla \times \mathbf{A}$ is the magnetic field. Since magnetic helicity affects the dynamical evolution of magnetic fields, an indirect way to measure magnetic helicity is to measure the magnetic field power spectrum and compare it with the evolution seen in magnetohydrodynamical (MHD) simulations \cite{Sigl:2002kt,0004-637X-640-1-335,Saveliev:2013uva}. Ther are also some proposals to {\em directly} measure magnetic helicity based on the propagation of cosmic rays~\cite{Kahniashvili:2005yp}. More recently, it has been proposed that helicity can leave characteristic parity-odd imprints on the arrival directions of gamma rays that are the result of gamma-ray-induced electromagnetic cascades \cite{PhysRevD.87.123527,Tashiro:2013ita,Tashiro:2014gfa,Long:2015bda,Chen11072015}. In particular, Long \& Vachaspati~\cite{Long:2015bda} have carried out a thorough analysis of the morphology of the arrival directions of gamma rays using a semi-analytical approach, but without including the stochasticity of the magnetic field or the cascade process. Hence, a full Monte Carlo approach and three-dimensional simulations are needed in order to confirm or refute their findings and provide a solid basis for further analyses. The observation of helical primordial magnetic fields has profound implications for particle physics and the early universe. Scenarios in which the cosmological matter-antimatter asymmetry is generated dynamically are found to also produce helical magnetic fields~\cite{Vachaspati:2001nb}. The handedness of the field is related to details of the matter-genesis scenario~\cite{Vachaspati:2001nb,Long:2013tha}. If the observed magnetic fields are coherent on very large scales, they may have been produced at the initial epoch, perhaps during an inflation \cite{PhysRevD.37.2743,Ratra:1991bn}. Helicity on these scales would indicate the presence of certain parity violating interactions in the fundamental Lagrangian~\cite{Caprini:2014mja}. In the present work we perform simulations of the propagation of gamma rays in both helical and non-helical IGMF. This paper is structured as follows: first, we discuss the theory and implementation of simulations of electromagnetic cascades in Sec.~\ref{sec:EMSim}; in Sec.~\ref{sec:Results} we apply our approach to different magnetic field configurations, focusing in particular on the role of magnetic helicity (Sec.~\ref{sec:HelMag1} - \ref{sec:HelMag3}); in Sec.~\ref{sec:Discussion} we discuss the results, draw our conclusions and give a short outlook.

\label{sec:Discussion} We have performed three-dimensional Monte Carlo studies of the development of gamma-ray-induced electromagnetic cascades in the intergalactic medium in the presence of magnetic fields. We have used the ``Large Sphere Observer'' method for improved computational performance. In this case all cascade photons hitting the surface of the sphere are detected by the the observer. With a standard three-dimensional Monte Carlo simulation most cascade photons would not reach Earth, resulting in wasted computation and very low statistics. A simplification made in our treatment is that the magnetic field evolves adiabatically with redshift as $B(z) = B(z=0) (1+z)^{2}$. This is justified because the cascade development we have discussed occurs in cosmic voids where MHD amplification and contamination by sources is minimal. Also, the sources are at redshifts $z \lesssim 1$. We first compared our computational setup with analytical approximations and then validated it in simple scenarios containing a uniform magnetic field oriented parallelly and perpendicularly to the line of sight of the blazar jet. As expected, for a magnetic field parallel to the direction of the jet of half-opening angle $\Psi$, assumed to be pointing toward Earth, effects of the field were not observed. For a magnetic field perpendicular to the direction of the jet, deflections were non-zero and in the expected direction. Similar results were obtained for stronger and weaker magnetic fields and other orientations. These results are in accordance with Ref.~\cite{Long:2015bda} and also with the predictions of Eq.~(\ref{thetaEBD}). We have also studied the particular case of a magnetic field with a Batchelor power spectrum with and without helicity. The effects of helicity can be clearly seen in Fig.~\ref{fig:test5g}, where arrival directions follow right- or left-handed spirals, depending on the sign of the helicity. For stochastic fields, in general, the results tend to converge toward the case of a uniform magnetic field in the limit of large coherence lengths. We have considered only large values of correlation length ($L_{\rm c} \simeq 120\,{\rm Mpc}$) since for much smaller coherence lengths, with the other parameters being held fixed, no clear signature of helicity can be seen, as shown in Fig.~\ref{fig:diffLc}. Nevertheless, one should bear in mind that the current upper limits of coherence length of magnetic fields in voids range between a few and hundreds of Mpc~\cite{DuNe}, placing the chosen value of 120 Mpc well within the allowed bounds. We have deployed the so-called $Q$-statistics, a powerful analysis tool that makes it possible to determine the properties of magnetic helicity directly from the observables of gamma rays measured at Earth. In this work we for the first time applied $Q$-statistics to realistic three-dimensional simulations of electromagnetic cascades. Our results for $Q$ are shown in Fig.~\ref{fig:Qtest5g}. The plots do not show a strong correlation between $Q$ and the existence and sign of the helicity. At the moment we cannot clearly state whether averaging over several objects will show a stronger correlation. We plan on investigating this issue in a future work. It is important to stress the fact that $Q$-statistics might not be the final method to quantify magnetic helicity, however it is a good initial approach and has been used in several works (Refs.~\cite{PhysRevD.87.123527,Tashiro:2013ita,Long:2015bda}) with satisfactory results. In this work we have, for the first time, introduced the $S$-statistics, which is a direct measure of the handedness of a pattern with respect to the line of sight. We have shown that the orientation, represented by the sign of $S$, is directly correlated with the sign of helicity. This shows that the $S$ measure is also a powerful tool to be used in the analysis of helicity of IGMF. Backgrounds at the $\sim\,$10-100 GeV energy range are expected due to secondary photons from AGN halos whose jet opening angles do not encompass the Earth. Other astrophysical sources of photons in this energy range also exist and have to be taken into account. In this first work we have neglected these backgrounds, which will be included in future studies. We found that it is probably necessary to analyze various sources in order to make a definite statement about the sign of the helicity, since a clear signature cannot always be seen. In the future we will extend our simulations to the case of multiple sources and diffuse gamma rays. We expect to be able to reproduce actual detections and consequently retrieve more precise information about IGMF, which can be used to infer their origin and evolution. In addition, we will extend the analysis by further exploring the parameter space as varying quantities such as the magnetic field strength $B_{\rm rms}$, the magnetic correlation length $L_{\rm c}$ and source parameters such as its distance from the observer, its energy spectrum or its cutoff energy, as they may be important in order to obtain a complete picture of their influence as discussed above and to explain actual observations.

1607.00320

1607

1607.05733_arXiv.txt

Delensing is an increasingly important technique to reverse the gravitational lensing of the cosmic microwave background (CMB) and thus reveal primordial signals the lensing may obscure. We present a first demonstration of delensing on Planck temperature maps using the cosmic infrared background (CIB). Reversing the lensing deflections in Planck CMB temperature maps using a linear combination of the 545 and 857\,GHz maps as a lensing tracer, we find that the lensing effects in the temperature power spectrum are reduced in a manner consistent with theoretical expectations. In particular, the characteristic sharpening of the acoustic peaks of the temperature power spectrum resulting from successful delensing is detected at a significance of $16\,\sigma$, with an amplitude of $A_{\mathrm{delens}}=1.12\pm 0.07$ relative to the expected value of unity. This first demonstration on data of CIB delensing, and of delensing techniques in general, is significant because lensing removal will soon be essential for achieving high-precision constraints on inflationary B-mode polarization.

The gravitational lensing of the cosmic microwave background (CMB) is rapidly becoming a powerful probe of cosmology and fundamental physics \citep{smith07,act1, spt1,plancklens,polarbear,biceplens}. However, the lensing deflections are not only a source of information, but they can also obscure signals in the primordial CMB. In particular, gravitational lensing converts E-mode into B-mode polarization; these lensing-induced B-modes act as a source of noise~\cite{Lewis2002,Hu2002} that limits constraints on the primordial B-mode power from inflationary gravitational waves \citep{KKS97,ZS98}. Gravitational lensing also smooths the acoustic peaks in the temperature power spectrum, which weakens constraints on cosmological parameters probed by the peak positions, such as the number of relativistic particle species. With the rapid increase in sensitivity expected from forthcoming CMB experiments, even near-future CMB B-mode searches will be completely lensing-limited. Delensing, the process of reversing, or removing, the lensing effects in maps, is therefore a method that will be essential to realizing powerful constraints on inflationary B-modes and other parameters~\cite{KnoxSong2002,KCK2002,seljakhirataBmodes}. However, despite their importance for progress in key areas of CMB research and despite much theoretical work, delensing methods have thus far not been successfully demonstrated on data. In this Letter, we show a first demonstration of delensing with CMB temperature data from Planck. The method generally considered to delens is to take a field that traces the dark matter distribution and filter it to obtain a weighted proxy for the true lensing map. This is then used to reverse the lensing in CMB maps, or, equivalently in the case of B-mode delensing, construct a linearized estimate of the lens-induced modes that is subtracted off (see, e.g., Refs.~\cite{HuOkamoto2002,KnoxSong2002,KCK2002,seljakhirataBmodes,Smith2009,delensingexternal,delensinginflation,SKAdelensing}). Though often CMB-internal delensing (i.e., using a lensing map reconstructed from the CMB itself) is considered, Ref.~\cite{sherwin15} discuss in detail the use of the cosmic infrared background (CIB) for this purpose; see also~\cite{delensinginflation}. The CIB is mostly made up of unresolved emission from dusty star-forming galaxies at high redshifts; as the CMB lensing arises from similar high redshifts, the CIB should be strongly correlated with the lensing convergence. This correlation has been measured to be up to around $80\%$, suggesting that one could remove more than half of the lensing~\citep{sherwin15}, assuming the CIB can be accurately separated from Galactic foregrounds. In this paper, we use the CIB estimated from Planck maps at 545 and 857\,GHz to delens the CMB temperature anisotropies measured by Planck at lower frequencies. In particular, we measure the CMB temperature angular power spectrum after delensing and show that the lensing-induced peak-smoothing is significantly reduced in agreement with expectations. Aside from being of general interest as a first demonstration of delensing methods applied to data, this is useful for additional reasons. First, it provides confirmation that CIB delensing is possible despite challenges such as accurately separating the CIB and Galactic dust emission. Second, our demonstration that the change in the power spectrum by delensing is consistent with expectations is a model-independent test that lensing affects the temperature power spectrum in the expected manner.

\label{sec:summary} We have presented a first demonstration of CIB delensing on the temperature anisotropies of the CMB. We demonstrate the expected sharpening of the acoustic peaks at high significance, reporting a $16\,\sigma$ detection of delensing effects in the power spectrum. We note that this sharpening of the peaks is a robust effect that is not easily mimicked by instrumental or astrophysical systematics, which would generally add uncorrelated lensing effects. This is illustrated by the fact that null-tests involving delensing with uncorrelated fields result in a smoothing rather than sharpening of the peaks (i.e., lensing-like rather than delensing effects). The reduction in the peak-smoothing of the temperature spectrum is found to be in agreement with the theoretical expectation (based on the correlation coefficient of our CIB maps with the Planck reconstruction of the lensing potential). In particular, the delensing strength is found to be $A_{\mathrm{lens}}= 1.12 \pm 0.07$ relative to a fiducial expectation of unity. This agreement of data and theory shows consistency between the CMB power spectrum analysis and CIB-derived lensing potential maps. We note that, beyond a consistency test, CIB delensing of the temperature and, particularly, E-mode polarization may allow for improved constraints on parameters such as the effective number of relativistic degrees of freedom ($N_{\text{eff}}$), but we defer such considerations to future work. More broadly, our results demonstrate the viability of CIB delensing as a method. However, while our simple multi-frequency method for Galactic dust removal is sufficient for our current purposes, it results in a correlation coefficient of the CIB with the lensing potential of only around 40\,\% on the degree scales relevant for delensing the temperature anisotropies. The correlation is stronger on the (smaller) scales that are relevant for delensing B-modes, but ultimately more sophisticated methods, e.g., the approach of Ref.~\cite{2016arXiv160509387P} or $\text{H}_{\text{\textsc{I}}}$-based dust cleaning, will be needed to achieve delensing performance closer to that forecasted in Ref.~\cite{sherwin15}. With upcoming CMB experiments, delensing methods will be crucial for revealing small inflationary B-mode polarization signals; our work on CIB-delensing of the CMB temperature anisotropies is an early step in this important research area.

1607.05733

1607

1607.02749_arXiv.txt

{Almost 200 different species have been detected in the interstellar medium (ISM) during the last decades, revealing not only simple species but complex molecules with more than 6 atoms. Other exotic compounds, like the weakly-bound dimer (H$_2$)$_2$, have also been detected in astronomical sources like Jupiter.} {We aim at detecting for the first time the CO-H$_2$ van der Waals complex in the ISM, which if detected can be a sensitive indicator for low temperatures.} {We use the IRAM\,30m telescope, located in Pico Veleta (Spain), to search for the CO-H$_2$ complex in a cold, dense core in TMC-1C (with a temperature of $\sim 10$~K). All the brightest CO-H$_2$ transitions in the 3~mm (80-110~GHz) band have been observed with a spectral resolution of 0.5--0.7~km~s$^{-1}$, reaching a rms noise level of $\sim 2$~mK. The simultaneous observation of a broad frequency band, 16~GHz, has allowed us to conduct a serendipitous spectral line survey.} {No lines belonging to the CO-H$_2$ complex have been detected. We have set up a new, more stringent upper limit for its abundance to be [CO-H$_2$]/[CO]~$\sim 5\times10^{-6}$, while we expect the abundance of the complex to be in the range $\sim 10^{-8}$--$10^{-3}$. The spectral line survey has allowed us to detect 75 lines associated with 41 different species (including isotopologues). We detect a number of complex organic species, e.g.\ methyl cyanide (CH$_3$CN), methanol (CH$_3$OH), propyne (CH$_3$CCH) and ketene (CH$_2$CO), associated with cold gas (excitation temperatures $\sim 7$~K), confirming the presence of these complex species not only in warm objects but also in cold regimes.} {}

} During the last decades, the number of atomic and molecular species detected in the interstellar medium (ISM) has increased considerably thanks to (i) the improved sensitivity of facilities like the IRAM\,30m telescope in Spain or the Atacama Large Millimeter/submillimeter array (ALMA) in Chile, and (ii) new laboratory measurements of transitions of new species included in catalogues such as the Cologne Database for Molecular Spectroscopy (CDMS). Almost 200 different species have been found in Galactic/extragalactic environments such as cold dense cores, hot molecular cores, circumstellar disks, evolved stars or large diffuse molecular clouds. These 200 species do not consist only of simple molecules like the most abundant H$_2$ and CO, but also include complex species usually defined as molecules with 6 or more atoms (\citealt{HvD2009}; see the CDMS database\footnote{http://www.astro.uni-koeln.de/cdms/molecules} for a summary of detected species in space). Molecular hydrogen, H$_2$, is by far the most abundant molecule in the Universe, followed by carbon monoxide, CO. Therefore, the intermolecular forces between these two species are of fundamental interest. If the CO-H$_2$ van der Waals complex\footnote{We note that this complex does not correspond to the formaldehyde molecule, H$_2$CO.} exists in measurable amounts in the ISM, it could be a sensitive indicator for low temperatures. The binding energy of the complex is so small --- typically 20~cm$^{-1}$ or 30~K --- that the relative abundance of the complex in the gas phase is expected to increase at lower temperatures. There is an open debate about the feasibility of observing such weakly-bound species because their formation rates at the very low densities of interstellar molecular clouds (below 10$^{7}$~cm$^{-3}$) are low, due to the small probability of three-body collisions, which is the main formation route of van der Waals complexes in the laboratory. On the other hand, the large timescale on which these processes occur in interstellar space makes radiative association, which is usually a slow process, quite feasible (e.g.\ \citealt{Klem06}). Also non-equilibrium conditions in the ISM may strongly favor the formation and concentration of the CO-H$_2$ complex over time on the surfaces of dust grains in shielded regions at low temperatures, with release to the gas-phase occurring by localized heating processes such as turbulence or jets/outflows (e.g.\ \citealt{Allen97}). However, one also has to consider that CO tends to be frozen out onto dust grains in very cold, dense regions, and it seems difficult to release the CO-H$_2$ complex from grains without destroying it. On the other hand, this is completely unchartered territory, and even sensitive upper limits are useful. A detection of this complex would challenge many beliefs we have about the chemistry of dense molecular clouds. There have been several attempts to observe complexes containing CO and H$_2$ molecules. The detection of the H$_2$ dimer, (H$_2$)$_2$, in the atmosphere of Jupiter has been reported by \citet{Mck88}, while searches for the CO dimer, (CO)$_2$ \citep{Van79}, and the CO-paraH$_2$ complex \citep{Allen97} in Galactic molecular clouds were not successful thus far. Laboratory data have clarified later a spectroscopic problem of these unsuccessful searches. The extensive millimeter-wave (MMW) studies of the CO dimer \citep{Su07} have shown that the radio astronomical search of this complex was based on frequencies which cannot be unambiguously attributed to (CO)$_2$. In the case of the CO-paraH$_2$ complex, the interstellar search was outside the correct frequency position of the most promising R(0) line, as later identified by the first MMW study of CO-paraH$_2$ \citep{Pak99}, and only the weaker Q(1) line was correctly tuned. Recent laboratory studies of the CO-H$_2$ complex have provided precise MMW frequencies with uncertainties of about 50~kHz for the complex in different spin modifications and for its deuterated isotopologues: CO-paraH$_2$ \citep{Po09ApJ}, CO-orthoH$_2$ \citep{Jan13}, CO-orthoD$_2$ \citep{Po09OptSp} and CO-HD \citep{Po15}. Therefore, the availability of precise rest frequencies and modern astronomical receivers (with a sensitivity several times better than the old receivers used 20 years ago), combine in a great opportunity to detect for the first time a van der Waals complex in the ISM. In this paper we present IRAM\,30m observations of a cold region in the Taurus molecular cloud in the search for the CO-H$_2$ van der Waals complex. In Sect.~\ref{s:obs} we describe the observations. In Sect.~\ref{s:res} we present the main results. Unfortunately, we do not have a detection of the CO-H$_2$ complex but we can set a new limit that can be used in future chemical modelling. In addition to the search for the CO-H$_2$ complex, the IRAM\,30m observations allowed us to conduct a spectral line survey of a very cold region ($\sim 10$~K), and we report the detection of complex organic molecules (COMs) as well as first time tentative detections of species in this object. In Sect.~\ref{s:disc} we discuss our results, and we end the paper with a summary of the main results in Sect.~\ref{s:summary}. \begin{figure}[t] \begin{center} \begin{tabular}[b]{c} \includegraphics[width=0.9\columnwidth]{CO-paraH2.pdf} \\ \includegraphics[width=0.9\columnwidth]{CO-orthoH2.pdf} \\ \end{tabular} \end{center} \caption{Energy level diagram for the CO-paraH$_2$ (top panel) and CO-orthoH$_2$ (bottom panel) van der Waals complex. \textit{Top panel}: The energy levels are labeled by quantum numbers $J$, $j_\mathrm{CO}$ and $l$ where $J$ is the total angular momentum, $j_\mathrm{CO}$ refers to the rotation of the CO sub-unit and $l$ refers to the end-over-end rotation of the complex. $K$ corresponds to the projection of $J$ onto the intermolecular axis. The labels $e$ and $f$ indicate the parity of the $J$ levels within a given stack. The parity of an even-$J$ level is `$+$' for stacks labeled by $e$ and `$-$' for $f$, while the parity of an odd-$J$ level is `$-$' for $e$ and `$+$' for stacks labeled $f$. The insert shows the approximate geometrical structure of the CO-H$_2$ complex (see \citealt{Po09ApJ} for details). \textit{Bottom panel}: The energy levels are labeled by quantum numbers $J$, parity $P$ and $n_J$,$P$, a consecutive number of the state for the given values of $J$ and $P$ (see \citealt{Jan13} for details). In both panels, the targeted transitions are indicated by arrows.} \label{f:energylevels} \end{figure} \begin{table} \caption{Frequencies of the brightest CO-H$_2$ targeted lines} \label{t:transitions} \centering \begin{tabular}{l c c} \hline\hline &Transition &Frequency (MHz) \\ \hline CO-paraH$_2$ &(1,1,0)--(0,0,0) & 108480.857 \\ &(1,1,1)--(1,0,1) & \phn91012.364 \\ CO-orthoH$_2$ &(2,f,2)--(1,f,1) & 101907.919 \\ &(3,e,2)--(2,e,1) & \phn93433.726 \\ CO-orthoD$_2$ &(1,1,0)--(0,0,0) & 102791.612 \\ &(1,1,1)--(1,0,1) & \phn89483.510 \\ CO-HD &(1,1,0)--(0,0,0) & 105636.\phnnn \\ \hline \end{tabular} \tablefoot{The uncertainty in the frequency measurements is about 50~kHz for all the transitions, except for CO-HD with a few tens of MHz. The labelling of the quantum numbers of the transitions are explained in detail in Fig.~\ref{f:energylevels}.} \end{table} \begin{figure*}[h!] \begin{center} \begin{tabular}[b]{c} \includegraphics[width=0.95\textwidth]{plot_COH2.pdf} \\ \end{tabular} \caption{\textit{Top panels}: Full spectrum observed with the IRAM\,30m telescope towards the cold dense core in TMC-1C. The mean rms noise level is $\sim 2$~mK. Most of the detected lines emit only in one channel (channel width 0.5--0.7~km~s$^{-1}$), suggesting that the linewidth of the different lines is $\le 0.7$~km~s$^{-1}$ (see Sect.~\ref{s:molecules}). \textit{Bottom panels}: Close-up views of the frequency ranges around the brightest transitions of the CO-H$_2$ van der Waals complex. The corresponding frequencies are listed in Table~\ref{t:transitions}, and are shown in the panels with a vertical dotted line. The expected linewidth is $\approx0.3$~km~s$^{-1}$, as measured in higher spectral resolution observations (e.g.\ \citealt{Spez13}).} \label{f:specCOH2} \end{center} \end{figure*}

} \subsection{Detectability of the CO-H$_2$ complex\label{s:coh2Disc}} In Sect.~\ref{s:coh2} we report an upper limit of $\sim 6$~mK for the CO-H$_2$ lines. We follow the same approach as in \citet{Allen97} to establish an upper limit for the column density of the CO-paraH$_2$ complex. We use the dipole moment of CO and the total partition function of CO-paraH$_2$ calculated from the now known energy level scheme of the complex (see Fig.~\ref{f:energylevels}). Our calculations result in a value of $\sim 3\times10^{12}$~cm$^{-2}$ for the column density of the complex and a fractional abundance of the CO-paraH$_2$ dimer relative to CO of $\sim 5\times10^{-6}$, assuming the CO column density to be $6\times10^{17}$~cm$^{-2}$ (derived from the $^{13}$C$^{18}$O column density listed in Table~\ref{t:molecules}, and assuming standard $^{12}$C/$^{13}$C and $^{16}$O/$^{18}$O ratios of 60 and 500, respectively. In the following, we estimate what number density of the CO-H$_2$ molecular complex do we expect under the ISM conditions, and compare it to the new upper limit. All the reaction rates used in the following are generic rates, which have not specifically been measured or calculated, and are taken from the review paper by \citet{vD14}. The given reactions are the basic types of reactions in space. Following \citet{vD14}, there are two basic processes by which molecular bonds can be formed in the interstellar molecular clouds: radiative association and formation on grain surface with subsequent release to the gas phase. In the radiative association process, the binding energy of a new molecule or molecular complex is carried out through the emission of a photon, and can be described as: \begin{equation} \mathrm{H}_2 + \mathrm{CO} \rightarrow \mathrm{CO}\mathrm{-}\mathrm{H}_2 + h\nu \end{equation} and proceeds at the rate of a radiative association reaction $k_1\approx10^{-17}$--$10^{-14}$~cm$^3$~s$^{-1}$. For the case of the formation on grain surfaces, a dust particle accommodates the released energy, and the process can be described as \begin{equation} \mathrm{H}_2 + \mathrm{CO}\mathrm{-}\mathrm{grain} \rightarrow \mathrm{CO}\mathrm{-}\mathrm{H}_2 + \mathrm{grain}\nonumber \end{equation} which proceeds at a rate of $k_2\approx10^{-17}$~cm$^3$~s$^{-1}$. On the other side, there are three processes for the destruction of the complex: photodissociation, collisional dissociation and neutral-neutral bond rearrangement. The first one can be described by \begin{equation} \mathrm{CO}\mathrm{-}\mathrm{H}_2 + h\nu \rightarrow \mathrm{products} \end{equation} with a reaction rate of $k_3\approx10^{-10}$--$10^{-8}$~cm$^3$~s$^{-1}$. The second and third ones can be given by \begin{equation} \mathrm{CO}\mathrm{-}\mathrm{H}_2 + \mathrm{M} \rightarrow \mathrm{products} \end{equation} where M being a reaction partner, with rates for collisional dissociation of $k_4\approx10^{-26}$~cm$^3$~s$^{-1}$ and for bond rearrangement of $k_5\approx10^{-11}$--$10^{-9}$~cm$^3$~s$^{-1}$. We consider a dense condensation with an H$_2$ density of [H$_2$] = $4\times10^{4}$~cm$^{-3}$, the CO density given by [CO] = [CO-grain] = 10$^{-4}$ [H$_2$], and assume [M] = [H$_2$]. Under these conditions, {the formation is dominated by radiative association, while the destruction mainly occurs by the bond rearrangement. As it is stated by \citet{vD14}, collisional dissociation of molecules is only important in regions of very high temperature ($>3000$~K) and density. Thus,} we determine the CO-H$_2$ abundance in the equilibrium as [CO-H$_2$] = ($k_1$[H$_2$][CO])/($k_5$[M]) = $4\times10^{-8}$--$4\times10^{-3}$~cm$^{-3}$ and [CO-H$_2$]/[CO] $\sim 10^{-8}$--$10^{-3}$. The obtained range for a possible abundance of CO-H$_2$ is quite wide. From the comparison of our estimated [CO-H$_2$]/[CO] abundance to the upper detection limit of [CO-H$_2$]/[CO]$\sim5\times10^{-6}$, we can conclude that the complex might be detected by observations with one or two orders higher sensitivity. \subsection{Molecular inventory in cold regions\label{s:moleculesDisc}} Table~\ref{t:molecules} and Figure~\ref{f:molecules} reveal a relatively rich chemistry in the cold dense core TMC-1C. Despite the average excitation temperature being of only 7~K, we are able to detect a number of species with 6 or more atoms: CH$_3$CN, CH$_3$OH and CH$_3$CCH. The column densities for these species are in the range $10^{11}$--$10^{13}$~cm$^{-2}$, which results in abundances of $10^{-12}$--$10^{-10}$ assuming a H$_2$ column density of 10$^{22}$~cm$^{-2}$ (e.g.\ \citealt{Schnee2005}). These abundances are about two orders of magnitude lower than the typical abundances found toward more massive hot molecular cores. We have searched for more complex species, such as methyl formate (CH$_3$OCHO) or dimethyl ether (CH$_3$OCH$_3$), but we have not detected them with an upper limit on the column density of about $10^{12}$~cm$^{-2}$, assuming an excitation temperature of 7~K. Similarly to the cold core L1689B studied by \citet{Bac12} we also detect ketene (CH$_2$CO), with a column density of $\sim 3\times10^{12}$~cm$^{-2}$ in complete agreement with the column densities determined for L1689B. In addition to the main isotopologues of the detected species, we also detect transitions of the deuterated counterparts CH$_3$CCD and CH$_2$DOH. The deuteration level is estimated to be about 0.045 for CH$_3$CCH and 0.055 for CH$_3$OH, however, this deuteration fractions should be better constraint with future observations of other transitions and with higher spectral resolution (necessary to resolve the lines). The uncertainty of the column density listed in Table~\ref{t:molecules} does not includes the uncertainty in the linewidth, which can not be measured in our coarse spectral resolution observations. The column densities can differ by 30\% if the linewidth is increased/decreased by 0.1~km~s$^{-1}$, or by 50\% if the variation is 0.2~km~s$^{-1}$. Therefore, the column densities reported in Table~\ref{t:molecules} have to be considered with caution. High-spectral resolution observations are necessary to improve the determination of the excitation temperature and column density. Another source of uncertainty in the column density determination is the excitation temperature: Observations of more transitions for the different molecules are required to better constraint the column density and to search for non-LTE effects. In general, a number of deuterated compounds have been detected: DCS$^{+}$, HDCS, NH$_2$D, c-C$_3$HD, c-C$_3$D$_2$, CH$_2$DOH and CH$_3$CCD. The deuteration fraction is 0.2 for H$_2$CS, 0.07 for c-C$_2$HD, and about 0.05 for CH$_3$CCH and CH$_3$OH. It is worth noting that the column density measured for c-C$_3$HD and c-C$_3$D$_2$ is in agreement with the recent measurements of \citet{Spez13}. Finally, in addition to the COMs discussed above, we highlight the detection of some species: (\textit{a}) HCS$^{+}$ has been observed in previous surveys towards Taurus molecular cores (e.g.\ \citealt{Oh98, Kai04}). Here, we present for the first time, a tentative detection of the deuterated counterpart DCS$^{+}$. A detailed study of different deuterated species may help to better understand the routes of deuteration, in particular for those more complex species, and to compare with similar studies conducted in high-mass star forming regions (e.g.\ \citealt{Fontani2011, Fontani2015}); (\textit{b}) Similarly, we report for the first time a tentative detection of HOCO$^{+}$ in this source, for which we determine a column density of $\sim 2\times10^{11}$~cm$^{-2}$; and (\textit{c}) The detection of HCO is common in photon-dominated regions (PDRs; e.g.\ \citealt{schilke2001}), where the chemistry is dominated by the presence of large amounts of far-UV photons. The Taurus molecular cloud is a low-mass star forming complex, and therefore there are no high-mass stars in the region able to produce enough UV photons. In this survey we report the detection of HCO in a cold, dense core, not associated with a PDR, with a column density of $\sim 10^{12}$~cm$^{-2}$. \citet{Bacmann2016} studied HCO in a number of cold prestellar cores, and related its abundance with that of other species such as H$_2$CO, CH$_3$O and finally CH$_3$OH. The authors determine the abundance ratios between the different species to be HCO:H$_2$CO:CH$_3$O:CH$_3$OH $\sim 10:100:1:100$, when the formation of methanol occurs via hydrogenation of CO on cold grain surfaces. The observed abundances of the intermediate species HCO and CH$_3$O suggest they are gas-phase products of the precursors H$_2$CO and CH$_3$OH, respectively. We measure an abundance ratio of HCO:CH$_3$OH $\sim 1:10$ for our cold, dense core (see Table~\ref{t:molecules}), consistent with the results reported by \citet{Bacmann2016}. } We have used the IRAM\,30m telescope to conduct sensitive observations of a cold, dense core in TMC-1C, with the goal of detecting the CO-H$_2$ van der Waals complex. We have not detected any transition of the CO-paraH$_2$ and CO-orthoH$_2$ compounds with a rms noise level of $\sim 2$~mK for a spectral resolution of 0.7~km~s$^{-1}$. This sets a new strong upper limit for the abundance of the complex: [CO-H$_2$]/[CO]~$\sim 5\times10^{-6}$. We estimate that the expected abundance of the complex, with respect to CO, in the ISM can be $\sim 10^{-8}$--$10^{-3}$, which suggest that more sensitive observations would be required to search for and detect for the first time the CO-H$_2$ complex in the ISM. Our sensitive spectral line survey have revealed the detection of 75 different spectral lines coming from 41 different species (including isotopologues). The excitation temperature is $\sim 7$~K, consistent with previous estimates. We detect a number of complex organic molecules such as CH$_3$CN, CH$_3$OH, CH$_3$CCH and deuterated isotopologues. The detection of these species in a cold object is consistent with the similar findings in other objects (e.g.\ L1689B, \citealt{Bac12}). Future studies of these complex species to better constraint the physical parameters, as well as the study of more rare isotopologues, can help to improve the current understanding of the formation of complex species in the cold ISM.

1607.02749

1607

1607.00044_arXiv.txt

{} { We present a custom support vector machine classification package for photometric redshift estimation, including comparisons with other methods. We also explore the efficacy of including galaxy shape information in redshift estimation. Support vector machines, a type of machine learning, utilize optimization theory and supervised learning algorithms to construct predictive models based on the information content of data in a way that can treat different input features symmetrically, which can be a useful estimator of the information contained in additional features beyond photometry, such as those describing the morphology of galaxies. } {The custom support vector machine package we have developed is designated SPIDERz and made available to the community. As test data for evaluating performance and comparison with other methods, we apply SPIDERz to four distinct data sets: 1) the publicly available portion of the PHAT-1 catalog based on the GOODS-N field with spectroscopic redshifts in the range $z < 3.6$, 2) 14365 galaxies from the COSMOS bright survey with photometric band magnitudes, morphology, and spectroscopic redshifts inside $z < 1.4$, 3) 3048 galaxies from the overlap of COSMOS photometry and morphology with 3D-HST spectroscopy extending to $z < 3.9$, and 4) 2612 galaxies with five-band photometric magnitudes and morphology from the All-wavelength Extended Groth Strip International Survey and $z < 1.57$. } {We find that SPIDER-z achieves results competitive with other empirical packages on the PHAT-1 data, and performs quite well in estimating redshifts with the COSMOS and AEGIS data, including in the cases of a large redshift range ($0 < z < 3.9$). We also determine from analyses with both the COSMOS and AEGIS data that the inclusion of morphological information does not have a statistically significant benefit for photometric redshift estimation with the techniques employed here. } {}

\label{intro} An important challenge for the current and coming era of large multi-band extragalactic surveys is obtaining sufficiently accurate photometric redshift estimates and understanding the error properties of these estimates (see e.g. \citet{Huterer06} for a review). Unlike time consuming spectroscopic redshift determination, photometric redshift estimation (photo-z) is subject to significant systematic errors and confusion because the spectral information of a galaxy is limited to the magnitude or flux in a number of wavelength bands. When photo-zs are used, science goals such as using weak lensing for cosmology are strongly affected by the number of outliers --- those objects whose estimated photo-zs are far from the actual redshifts \citep[e.g.][]{Hearin10}. In general, data sets with bands extending into those observed by infrared telescopes (e.g. J, H, and K bands) have more accurate photo-z estimation and fewer outliers. However, most upcoming large surveys, such as the Large Synoptic Survey Telescope \citep[LSST,][]{LSSTover}, will have optical and near-infrared data only. Reducing the number of, and potentially having a method of identifying, the potential outliers in photo-z estimation is an important goal for these projects. Photo-z estimation techniques have traditionally been divided into two main classifications. So-called ``Template fitting'' methods, such as the {\it Lephare} package as described in \citet{lephare} and \citet{Arnouts99}, and {\it Bayesian Photometric Redshift (BPZ)} as described in \citet{bpz}, involve correlating the observed band photometry with model galaxy spectra and redshift, and possibly other model properties. So-called ``Empirical'' or ``Training set'' methods, such as artificial neural networks \citep[e.g. {\it ANNz},][]{annz}, boosted decision trees \citep[e.g. {\it ArborZ},][]{Gerdes10}, regression trees / random forests \citep[e.g.][]{Carliles10,KB13}, support vector machines \citep[e.g.][]{W04}, polynomial mapping \citep[e.g.][]{B05,LY08}, and others develop a mapping from input parameters to redshift with a training set of data in which the actual spectroscopic redshifts are known, then apply the mappings to data for which the redshifts are to be estimated. Both have their drawbacks --- template fitting methods require assumptions about intrinsic galaxy spectra or their redshift evolution, and empirical methods require the training set to be ``complete'' in the sense that it is representative of the target evaluation population in bulk in all characteristics. In a previous work \citep{NNP} we reported on a custom artificial neural network algorithm for photometric redshift determination. Because of the larger frequency of mergers at higher redshifts and the general evolutionary trend from spiral to elliptical shapes among galaxies, it is a reasonable hypothesis that galaxy morphology and redshift are correlated in such a way that the addition of morphological information could improve photo-z estimation. The inclusion of morphological parameters in photo-z estimation has been studied using Sloan Digital Sky Survey (SDSS) data by \citet{Tag03} with an artificial neural network determination, and by \citet{VC2006} and \citet{WS2006} with other methods. \citet{Tag03} find possible modest improvement with the inclusion of shape information, although they restrict their analysis to quite low redshift ($z \leq$ 0.7) galaxies. \citet{WS2006} consider several empirical methods and show marginal improvement for some methods with the addition of morphological information. \citet{VC2006} claim an improvement of between 1 and 3 percent in the RMS error in photo-z determination, however it is not noted whether this result is significant and the method of photo-z estimation is not discussed. \citet{NNP} considered one of the galaxy test data sets used here which includes a more thorough sample of higher redshift galaxies extending to $z \sim 1.57$ and found that including galaxy shape information did not result in a statistically significant benefit in the context of a neural network method. In this work, we evaluate the performance of a custom support vector machine (SVM) package for photo-z determination with comparisons to other photo-z algorithms, and also explore the efficacy of integrating parameters describing the morphological information of galaxies. The SVM package used in this analysis is developed for the IDL environment by one of the authors (EJ), based largely on algorithmic procedures outlined in \citet{LIBSVM}, and has been named SPIDERz (SuPport vector classification for IDEntifying Redshifts).\footnote{available from http://spiderz.sourceforge.net with usage documentation provided there.} It can include additional parameters beyond photometry, such as morphological information, on an equal footing. Here we follow convention \citep[e.g.][]{Hildebrandt10} and define ``outliers'' as those galaxies where \begin{eqnarray} Outliers: {{\vert z_{phot}-z_{spec} \vert} \over {1+z_{spec}}} > .15, \label{erroreq} \end{eqnarray} where $z_{phot}$ and $z_{spec}$ are the estimated photo-z and actual (spectroscopically determined) redshift of the object. The RMS photo-z error in a realization is given by a standard definition \begin{eqnarray} \sigma_{\Delta z/(1+z)} \equiv \sqrt { {{1} \over {n_{gals}}} \Sigma_{gals} \left( {{ z_{phot}-z_{spec} } \over {1+z_{spec}}} \right) ^2 }, \label{RMSeq} \end{eqnarray} where $n_{gals}$ is the number of galaxies in the evaluation set and $\Sigma_{gals}$ represents a sum over those galaxies. For comparison with other works, we also calculate in certain determinations the RMS error without the inclusion of outlier galaxies, referring to this quantity as the ``reduced'' RMS or R-RMS. The paper is organized as follows. In \S \ref{method} we discuss the SVM model implemented in SPIDERz. In the following sections we perform analyses with four separate test data sets and make comparisons to other photo-z methods when possible. In \S \ref{PHAT} we discuss the results of testing SPIDERz on the publicly available portion of the PHoto-z Accuracy Testing real galaxy data catalog (PHAT-1) which spans the redshift range $z < 3.6$ and provides a useful comparison with other photo-z estimation methods. In \S \ref{COSMOS} we discuss the results of testing SPIDERz on the Cosmic Evolution Survey (COSMOS) bright catalog with 14365 galaxies spanning $z < 1.4$ with ten-band photometry, spectroscopic redshifts, and seven morphological measurements, highlighting results obtained with differing numbers of bands and the inclusion of morphological information. We perform a similar analysis in \S \ref{COS-HST} with a catalog of 3048 galaxies spanning a wider redshift range ($z < 3.9$), which were obtained by combining COSMOS photometry and morphology with 3D-HST spectroscopy. In \S \ref{AEGIS} we discuss the results of testing SPIDERz on a catalog of 2612 galaxies with spectroscopic redshifts, five-band photometry, and morphological information from the All-wavelength Extended Groth Strip International Survey (AEGIS) in the range $z < 1.57$, and make a comparison to a neural network determination on this data. We present a summary in \S \ref{disc}.

\label{disc} We have developed a custom SVM classification algorithm for photometric redshift estimation in the IDL environment. The package, SPIDERz, is available to the community. It outputs for each galaxy both an effective distribution of probabilities (in bins of redshift) for each galaxy's photo-z and a single-valued, most likely predicted photo-z bin, with the bin size chosen by the user. In this analysis for practicality of evaluating metrics such as outliers and RMS errors we use only the single-valued photo-z prediction. We compare results obtained with SPIDERz to those previously reported for other codes with the PHAT-1 catalog in \S \ref{PHATresults}. We see that SPIDERz performs comparable to other empirical methods on this catalog, noting that as discussed in \S \ref{PHATdata} far fewer training galaxies were available for SPIDERz. On this latter point, having such a reduced number of training galaxies available is a significant disadvantage for SPIDERz or any empirical method in a dataset such as PHAT-1 where proportionally few of the galaxies are higher redshift. In the context of SVC as discussed in \S \ref{alg}, as the training data in certain redshift bins becomes sparse, there may be insufficient support vectors available for those bins in the binary classifications for effective hyperplane solutions. SPIDERz can naturally include additional parameters beyond band magnitudes. We note that an SVM is in a sense an unbiased way of determining the relative strength of the correlations of a set of input parameters with the output parameter, and that this SVM algorithm is designed to treat all input parameters on an equal footing, thus providing for a convenient method for investigating the inclusion of additional inputs beyond photometry. In order to both explore a much larger dataset and examine the effects of the inclusion of morphological parameters on photo-z estimation, we form a data set of 14365 galaxies with photometry, morphological parameters, and spectroscopic redshifts from the COSMOS field as discussed in \S \ref{COSMOS}. We find that while SPIDERz performs relatively well on this data, that the inclusion of morphological information in the form of morphological parameters does not improve the number of outliers or RMS errors in photo-z estimation over the case of five optical photometric bands only. As expected, using ten photometric bands reduces the number of outliers and RMS errors considerably relative to five bands. To study the performance of SPIDERz and the effects of including morphological parameters over a larger redshift range we form a data set of 3048 galaxies from the overlap of COSMOS photometry and morphology with 3D-HST spectra as discussed in \S \ref{COS-HST}. We find SPIDERz performs quite well on this data, including on higher redshift galaxies, even though these galaxies form a small fraction of the training set. We again find that the inclusion of morphological information in the form of morphological parameters does not improve the number of outliers or RMS errors in photo-z estimation over the case of five optical photometric bands only, and again using ten photometric bands reduces the number of outliers and RMS errors considerably relative to five bands. For comparison with a previous determination with a neural network algorithm, we also estimate photo-z for a data set which consists of 2612 galaxies with five optical band magnitudes, reliable spectroscopic redshifts, and principal components of eight morphological parameters, discussed in \S \ref{AEGIS}. We again find that the inclusion of morphological information does not significantly decrease the number of outliers or RMS error in photo-z estimation. Previously \citet{NNP} found compatible results with the neural network estimation method. We do note that the inclusion of multiple principal components resulted in a diminished performance in photo-z estimation with the neural network algorithm due to the addition of noise while this had less of an effect on the performance of SPIDERz, indicating that perhaps SVMs are less susceptible to the negative effects of adding noise in this situation. The particular findings regarding the inclusion of shape information thus evidence some robustness in regard to different data sets, consideration of morphological parameters or principal components, and consideration of an SVM or neural network method. We conclude, therefore, that these results are likely applicable to all empirical methods with this redshift range and photometry restricted to the visible and near infrared bands. Any gain that may arise in an empirical photo-z determination from correlations between morphology and redshift is overwhelmed by the additional noise introduced. It is likely that any correlations between the morphological parameters and the galaxy type are degenerate to some extent with the correlations between galaxy type and galaxy colors. The possibility remains that under a less input-blind and more complicated scenario where, for example, the regularization parameter for misclassification (see equation \ref{costfn}) varies with the redshift, the additional noise contained in morphological information could be made less consequential. A potentially more straightforward possibility is that morphological information along with probability distribution considerations could be used only to flag galaxies more likely to be catastrophic outliers, which we will explore in a future work.

1607.00044

1607

1607.02088_arXiv.txt

Up-to-date isochrones, zero-age horizontal branch (ZAHB) loci, and evolutionary tracks for core He-burning stars are applied to the color-magnitude diagrams (CMDs) of M\,3, M\,15, and M\,92, focusing in particular on their RR Lyrae populations. Periods for the $ab$- and $c$-type variables are calculated using the latest theoretical calibrations of $\log\,P_{ab}$ and $\log\,P_c$ as a function of luminosity, mass, effective temperature ($\teff$), and metallicity. Our models are generally able to reproduce the measured periods to well within the uncertainties implied by the stellar properties on which pulsation periods depend, as well as the mean periods and cluster-to-cluster differences in \pab\ and \pc, on the assumption of well-supported values of $E(B-V)$, $(m-M)_V$, and [Fe/H]. While many of RR Lyrae in M\,3 lie close to the same ZAHB that fits the faintest HB stars at bluer or redder colors, the M\,92 variables are all significantly evolved stars from ZAHB locations on the blue side of the instability strip. M\,15 appears to contain a similar population of HB stars as M\,92, along with additional helium-enhanced populations not present in the latter which comprise most of its RR Lyrae stars. The large number of variables in M\,15 and the similarity of the observed values of \pab\ and \pc\ in M\,15 and M\,92 can be explained by HB models that allow for variations in $Y$. Similar ages ($\sim 12.5$ Gyr) are found for all three clusters, making them significantly younger than the field halo subgiant HD\,140283. Our analysis suggests a preference for stellar models that take diffusive processes into account.

\label{sec:intro} The globular clusters (GCs) M\,15 (NGC\,7078) and M\,92 (NGC\,6341) are generally thought to have very close to the same metallicity (see the spectroscopic surveys by e.g., \citealt{ki03}, \citealt{cbg09a}) and age (e.g., \citealt{san82}, \citealt{van00}, \citealt{dsa10}). The strongest argument in support of coevality is that color-magnitude diagram (CMD) studies have shown that the difference in magnitude between the turnoff (TO) and the horizontal branch (HB) is nearly identical for these two systems (e.g., \citealt{dh93}). Originally, this so-called ``\delv\ parameter" was measured at the color of the TO (\citealt{ir84}), but the uncertainty of $V_{TO}$ can easily be as high as $\pm 0.1$ mag, implying $\delta$(age) $\sim\pm 1$ Gyr, because of the difficulty of determining the magnitude of the bluest point in a sequence of stars that is, by definition, vertical at the TO. Much more precise ages can be derived by fitting isochrones to the arc of stars from $\sim 1$ mag below the TO through to a point on the subgiant branch (SGB) that is $\sim 0.05$ mag redder than the TO, in conjunction with fits of zero-age horizontal branch (ZAHB) models to the cluster HB populations (\citealt[hereafter V13]{vbl13}; \citealt{lvm13}). Using this technique, which builds on the approaches advocated by \citet{cdk96} and \citet{bcp98}, V13 found that M\,15 and M\,92 have the same age to within $\pm 0.25$ Gyr. [The shapes of modern isochrones in the vicinity of the TO, in particular, appear to be quite a robust prediction and, in fact, stellar models are able to reproduce the turnoff portions of observed CMDs rather well when up-to-date color--$\teff$ relations (e.g., \citealt[hereafter CV14]{cv14}) are employed; see V13.] However, this result is not yet ironclad --- primarily because the two GCs have very different HB morphologies. Indeed, M\,15 is not at all like the majority of clusters with [Fe/H] $< -2$, including M\,92, whose HB populations are located predominately to the blue of the instability strip (IS), and their RR Lyrae stars constitute just a small fraction of the total number of core helium-burning stars. In M\,92-like HBs, both the paucity of variables and their high pulsation periods, relative to those determined for RR Lyrae in more metal-rich clusters (like M\,3), can be plausibly explained if these variables evolved into the IS from ZAHB locations on the blue side of the IS, where most of the HB stars are found (\citealt{ren83}, \citealt{ldz90}, \citealt{psc02}, \citealt{scf14}). Curiously, M\,15 has a horizontal branch that spans a much wider range in color than is typical of extremely metal-deficient GCs, and it is so rich in RR Lyrae that a large fraction of its variables must have evolved from ZAHB structures inside the IS (\citealt{rtr84}, \citealt{bcd84}, \citealt{bcf85}, \citealt{rf88}). Yet, the mean period of its $ab$-type RR Lyrae stars agrees very well with the values of \pab\ that have been derived for other Oosterhoff type II (hereafter, Oo II) systems (\citealt[\citealt{oo44}]{oo39}), including M\,92; see \citet[his Table 2]{cat09}. This suggests that, at the same intrinsic color, M\,15 and M\,92 variables have similar luminosities; and therefore that (in the mean at least) M\,15 RR Lyrae lie above the extension into the IS of the same ZAHB which provides a good fit to the main non-variable, blue HB population of M\,92 (as well as its counterpart in M\,15). There is another important difference between M\,15 and other GCs of very low metallicity in that it is the only one which has been found to have a signficant dispersion in the abundances of heavy neutron-capture elements (e.g., \citealt{coh11}, \citealt{whs13}). This may be (probably is) connected with the fact that M\,15 is one of the most luminous, and thus most massive, clusters in the Galaxy. Indeed, other systems with integrated $M_V < -9$ (see the latest version of the \citealt{har96} catalogue\footnote{www.physics.mcmaster.ca/$\sim$harris/mwgc.dat}) generally exhibit the largest chemical abundance anomalies; see, for instance, recent investigations of 47 Tuc (\citealt{mpd12}, \citealt{gls13}), NGC\,2808 (\citealt{car15}, \citealt{mmp15}), NGC\,2419 (\citealt{ck12}, \citealt{mbi12}), NGC\,6441 (\citealt{bpm13}) and M\,2 (\citealt{yrg14}). Moreover, as discussed in, e.g., the studies of 47 Tuc by \citet{dvd10} and of NGC\,2808 by \citet{dsf11} and \citet{mmp14}, consequences of the observed (or inferred, in the case of helium) abundance variations for their HB populations can often be identified. To be specific, Di Criscienzo et al.~found that the best match to the observed HB morphology of 47 Tuc is obtained if synthetic HB populations are generated on the assumption of $\Delta Y \approx 0.03$ for the initial He abundances (also see \citealt{scp16}). This is approximately the dispersion in $Y$ that has been inferred from the width of the cluster MS by \citet{apk09}. Similarly, D'Alessandro et al.~and Marino et al.~have found that the very unusual HB of NGC\,2808 can be explained if it consists of sub-populations of stars with low, intermediate, and high helium abundances that are consistent with the values of $Y$ implied by the cluster's triple MS (see \citealt{pba07}). Hence, it may turn out that the HB of M\,15 cannot be satisfactorily explained except as a superposition of multiple stellar populations --- something which has long been suspected (see, e.g., \citealt{bcf85}). Indeed, \citet[also see \citealt{jl15}]{jlj14} have recently speculated that the presence of different {\it generations} of stars, which assumes that resident chemically distinct populations formed at different times (\citealt{gcb12}), may be responsible for the appearance of the observed HBs in {\it most} clusters, as well as their separation into Oosterhoff groups. In their scenario, core helium-burning stars with normal helium abundances ($Y \approx 0.25$) populate a different range in color on the HB than those with slightly higher $Y$, enhanced CNO abundances, and younger ages (by 1--2 Gyr), and (if they exist) still younger stars with much higher $Y$. That is, the spread in color on the HB would be due more to the differences in the ages and the abundances of helium and CNO of the existing subgroups than to a large dispersion in mass at nearly constant $Y$ and [CNO/Fe], which is the canonical explanation (\citealt{ro73}). Since HBs are shifted to the red as the metallicity increases, the stars that are located in the IS could belong mostly to the first, second, or third generation depending on the cluster [Fe/H], possibly producing the observed RR Lyrae period shifts (see Jang et al., their Fig.~1 and the accompanying discussion.). However, although difficult to measure, C$+$N$+$O appears to be constant to within measuring uncertainties in most GCs; see, e.g., the spectroscopic results obtained for M\,4 by \citet{sci05}, for NGC\,6397 and NGC\,6752 by \citet{cgl05}, and for M\,3 and M\,13 by \citet[also see \citealt{ssb96}]{cm05}. To date, there is compelling evidence for large star-to-star [CNO/Fe] variations only in NGC\,1851 (\citealt{ygd09}), though there is some suggestion from photometric data that 47 Tucanae harbors a minor CNO-enhanced population of stars in its core (\citealt{apk09}). As shown by \citet{csp08} in the case of NGC\,1851, large variations in [CNO/Fe] cause the SGB to be broadened, or split, and since this is not commonly seen in GC CMDs (see, e.g., the {\it HST} photometric survey carried out by \citealt{sbc07}), intrinsic spreads in [CNO/Fe] larger than $\sim 0.2$ dex are effectively ruled out (unless the effects of age and [CNO/Fe] variations compensate each other). Indeed, even well-developed O--Na and Mg--Al anti-correlations, such as those derived for stars in M\,13 by \citet{jp12} and \citet{dny13}, respectively, can be reproduced remarkably well by theoretical models if the H-burning occurs at a sufficiently high temperature ($\approx 75\times 10^6$ K) and both C$+$N$+$O and the total number of Mg and Al nuclei are constant (see \citealt{dvh15}). At the present time, supermassive stars (\citealt{dh14}) are the only known nucleosynthesis site that has the required H-burning temperatures to achieve this consistency between theory and observations without requiring large {\it ad hoc} modifications to the rates of relevant nuclear reactions.\footnote{\citet{rdc15} have pointed out some difficulties with the scenario proposed by \citet{dh14}, and we do not disagree that there are valid concerns (also see \citealt{ikp16}). However, they may simply be telling us that we do not yet have the correct understanding of how supermassive stars fit into our picture of the very early evolution of GCs, or whether they are but one of the contributors to the chemical evolution of these systems at early times. Given their considerable success in explaining the observed light-element abundance correlations and anti-correlations --- and the limited success, or outright failure, of other hypotheses to accomplish the same thing (see \citealt{dvh15}) --- we suspect that supermassive stars will turn out to be an important piece of the puzzle. Although it is commonly believed that the chemically distinct populations in GCs arose as a result of successive star formation events, this possibility is still conjecture at the present time. The CN-poor, O-rich, Na-poor, $\ldots$ stars could have formed at essentially the same time as the CN-rich, O-poor, Na-rich, $\ldots$\ stars if such chemical abundance variations within GCs have, e.g., a supermassive star origin.} Thus there are ample reasons to question the variations in CNO and age that underpin the explanation of the Oosterhoff dichotomy suggested by \citet{jlj14} and \citet{jl15}. To properly evaluate the validity of their proposals, one should first examine how well updated models for the evolution of HB stars are able to explain both the morphologies of the observed HBs in GCs and the periods of their RR Lyrae variables. Since the difference in \pab\ between Oo II systems (M\,15, M\,92) and Oo I clusters (e.g., M\,3) is of particular interest, a careful consideration of the M\,3 HB is included in this investigation. After describing our evolutionary computations in \S\ref{sec:models}, fits of isochrones to the cluster TOs and of evolutionary tracks for the core He-burning phase to the observed HBs are presented in \S\ref{sec:obs}, along with comparisons of the predicted and observed periods of their RR Lyrae. The main results of this study are summarized and briefly discussed in \S\ref{sec:sum}.

\label{sec:sum} \begin{table*}[t] \begin{center} \caption{Mean Properties of the M\,3, M\,15, and M\,92 RR Lyrae Stars} \begin{tabular}{rcccccccccccccc} \hline \hline Name & & \multicolumn{1}{c}{\per\tablenotemark{a}} & $\sigma$ & \multicolumn{1}{c}{\per\tablenotemark{b}} & $\sigma$ & \tavg & $\sigma$ & \mavg & $\sigma$ & \bavg & $\sigma$ & \vavg & $\sigma$ & $Z$ \\ [0.5ex] \hline \multicolumn{15}{c}{$ab$-type} \\ [0.5ex] M\,3 & & 0.568 & 0.067 & 0.568 & 0.075 & 3.812 & 0.012 & 0.656 & 0.016 & 0.534 & 0.058 & 0.583 & 0.056 & $7.623\times 10^{-4}$ \\ M\,15 & & 0.654 & 0.060 & 0.654 & 0.068 & 3.813 & 0.013 & 0.706 & 0.051 & 0.257 & 0.055 & 0.326 & 0.058 & $2.466\times 10^{-4}$ \\ M\,92 & & 0.645 & 0.033 & 0.645 & 0.042 & 3.815 & 0.005 & 0.673 & 0.011 & 0.283 & 0.042 & 0.347 & 0.041 & $1.786\times 10^{-4}$ \\ [0.5ex] \multicolumn{15}{c}{$c$-type} \\ [0.5ex] M\,3 & & 0.336 & 0.050 & 0.336 & 0.058 & 3.855 & 0.013 & 0.630 & 0.016 & 0.514 & 0.104 & 0.520 & 0.102 & $7.623\times 10^{-4}$ \\ M\,15 & & 0.365 & 0.038 & 0.365 & 0.046 & 3.848 & 0.014 & 0.712 & 0.038 & 0.363 & 0.060 & 0.395 & 0.052 & $2.466\times 10^{-4}$ \\ M\,92 & & 0.352 & 0.050 & 0.352 & 0.053 & 3.860 & 0.016 & 0.662 & 0.005 & 0.340 & 0.040 & 0.360 & 0.029 & $1.786\times 10^{-4}$ \\ [0.5ex] \hline \multicolumn{15}{l}{$^{\rm a}$Observed mean period (in d) of the samples of RR Lyrae considered in this study.} \\ \multicolumn{15}{l}{$^{\rm b}$Predicted mean period (in d) of the samples of RR Lyrae considered in this study.} \\ \end{tabular} \end{center} \end{table*} Mainly during the last 3 decades of the 20$^{\rm th}$ century, but continuing to the present day, many investigators in the GC, stellar evolution, and variable star communities have tried to understand the \citet[1944]{oo39} dichotomy, particularly as regards M\,3 (Oo type I) and M\,15 (Oo II) because they are so rich in RR Lyrae. No one worked harder to explain the difference in the mean periods of their RR Lyrae populations than Allan Sandage and, in the end, it seems that his solution to this problem, that M\,15 RR Lyrae stars have higher helium abundances than those in M\,3 (see \citealt{sks81}), stands a good chance of being the right answer. [Prior to $\sim 2005$, GCs were considered to be simple stellar populations in which all stars within each cluster were thought to be coeval and essentially chemically homogeneous, aside from the ubiquitous star-to-star variations in CN. It was generally assumed that helium did not vary, given that the application of the R-method (\citealt{ib68}) yielded very similar helium abundances, $Y \approx 0.25$, for most clusters, especially those with red HBs (e.g., see \citealt[and references therein]{src04}). Consequently, everyone viewed the possibility that $Y$ varies inversely with [Fe/H], which also seems counter-intuitive from a chemical evolution perspective, with considerable skepticism. Only recently has it been established that GCs contain multiple, chemically distinct stellar populations that have, or probably have, different helium abundances. As mentioned in \S~\ref{sec:intro}, the most massive clusters have provided the most compelling evidence for such variations.] To be sure, our apparent success in modeling the HBs of M\,3, M\,15, and M\,92 is due in part to the improvements made to both the stellar models over the years and the theoretical relations that describe the dependence of the period on luminosity, mass, $\teff$, and metallicity for the fundamental and the first-overtone pulsators (\citealt{mcb15}). With just minor adjustments (well within the associated uncertainties) to the RR Lyrae $\teff$\ scale (or, equivalently, to the coefficients of $\log\teff$\ in these equations), it is possible to reproduce the observed values of \pab\ and \pc\ quite well. The advances that have been made likely explain why we find $\delta Y$(M\,15 {\it minus} M\,3) $\sim 0.03$, as compared with a difference closer to 0.05, in the same sense, that was derived by \citet[also see \citealt{srt87}]{sks81}. Table~1 lists the observed and predicted values of \pab\ and \pc\ and their standard deviations ($\sigma$), as derived from the individual RR Lyrae stars in each cluster, along with the mean values and standard deviations of the temperatures, masses, absolute bolometric magnitudes, and absolute $V$-band magnitudes of the variables. (Note that the mean periods which are calculated from equations (1) and (2) on the assumption of the quantities given in the sixth, eighth, tenth, and last columns agree very well with the values listed in the fourth column.) Nearly the same value of \tavg\ is obtained for the $ab$-type RR Lyrae ($\approx 3.815$) and $c$-type variables ($\approx 3.855$) in all three clusters. In addition, the variables in M\,15 are predicted to have higher mean masses (and significantly larger mass dispersions) than those in M\,92 and M\,3. Consistent with the plots of the RR Lyrae on various CMDs (see, e.g., Figs.~5, 12, and 18), the magnitudes of the first-overtone pulsators in M\,3 have the largest standard deviations, while the luminosity dispersions of both the $ab$- and $c$-type RR Lyrae are the smallest in M\,92. Not surprisingly, the variables in M\,15 and M\,92 have brighter absolute magnitudes by $\gta 0.2$ mag than those in M\,3. (As one would expect, the tabulated properties have some dependence on the samples of RR Lyrae that are considered. For instance, had we retained the most problematic M\,15 variables in our analysis, we would have obtained \bavg\ $= 0.241 \pm 0.068$ and \vavg\ $= 0.313 \pm 0.065$ for the $ab$-type variables in this cluster, as well as \bavg\ $= 0.354 \pm 0.062$ and \vavg\ $= 0.392 \pm 0.053$ for its first-overtone pulsators.) The importance of taking diffusive processes into account should be appreciated. One of the consequences of diffusion is that, due to the settling of helium in the interiors of stars during their main sequence and subgiant evolution, the envelope helium abundance after the first dredge-up (i.e., after the convective envelope has reached its maximum depth on the lower RGB) is predicted to be {\it less} than the initial helium content by $\delta Y \sim 0.003$ (assuming a $0.8 \msol$ model for [Fe/H] $= 1.55$). If diffusion is ignored, the envelope helium abundance is predicted to be {\it higher} than the initial abundance by $\delta Y \sim 0.01$. Since the luminosity of the HB is a sensitive function of $Y$ (see Figs.~8, 21), ZAHBs will be significantly fainter, implying reduced ZAHB-based distance moduli, if diffusion is treated. (As discussed in \S~\ref{subsec:m3}, the value of $(m-M)_V = 15.04$ that we have derived for M\,3 using ZAHB models satisfies the constraint provided by current calibrations of the RR Lyrae standard candle quite well.) Furthermore, the prominence of blue loops in post-ZAHB evolutionary tracks depends quite strongly on the helium abundance (see, e.g., Fig.~21). This has important ramifications for the intermingling of $ab$-type and $c$-type RR Lyrae. For instance, as discussed in \S~\ref{subsec:m3}, there appears to be very little overlap of the colors of these variables in M\,3 (see Fig.~5), which suggests that blue loops must be small or non-existent if the {\it transition} between fundamental and first-overtone pulsation, or vice versa, depends on the direction in which the core He-burning stars are evolving through the instability strip (the so-called ``hysteresis effect"; see \citealt[and especially the very instructive plots provided by \citealt{cct78} and \citealt{san81}]{vb73}). Our diffusive models for an initial helium abundance of $Y = 0.250$ predict small blue loops (see Fig.~5), though better consistency with the observations would be obtained if they were even smaller, which suggests that a slightly lower value of $Y$ should be adopted (but within the uncertainties of the primordial helium abundance) or that our models underestimate the rate at which settling occurs in stars. Higher $Y$ by $\sim 0.013$, as predicted by non-diffusive stellar models, would result in extended blue loops, which seems incompatible with the fairly sharp boundary between the $ab$- and $c$-type variables in M\,3. (The age of HD\,140283 provides another argument that diffusion physics should not be neglected in stellar models; see \S~\ref{subsubsec:other} and \citealt{vbn14}.) There is no overlap of the colors of these RR Lyrae in M\,92, nor is any expected because the evolution through the instability strip is clearly in the direction from blue to red from ZAHB locations on the blue side of the instability strip (see Figs.~9, 11, 12). It is not clear what to make of the M\,15 variables in this regard (see Fig.~21), partly because the magnitude-weighted colors derived by \citet{cbs08} seem particularly uncertain, and because it is not possible to unambiguously determine the helium abundances of the individual RR Lyrae stars. Further observational work to determine improved estimates of the colors of equivalent static stars of the M\,15 variables would be very helpful, as would an extension, towards lower metallicities and additional bandpasses, of the theoretical studies of \citet{bcs95} on the differences between different types of averages and the static magnitudes and colors. As discussed by \citet{amb15}, a separation of the $ab$- and $c$-type RR Lyrae in terms of their colors is found in most GCs, irrespective of whether they are OoI or OoII systems. Notable exceptions are the OoI cluster NGC\,3201 (\citealt{aac14}) and the OoII clusters M\,15 and $\omega$\ Cen (see \citealt{san81}). When there is a mixture of fundamental and first-overtone pulsators in a restricted color range within the instability strip, some of the variables are expected to be evolving from red to blue, and their periods should be decreasing with time, while others will be evolving in the opposite direction with positive period-change rates, $\Delta P/\Delta t$ (see Fig.~3 by \citealt{cct78}). In principle, it should be possible to use measurements of $\Delta P/\Delta t$ to determine the directions in which individual variables are evolving. In practice, however, this seems to be very difficult. As \citet{cc01} have concluded, period-change rates in M\,3 appear to be due more to ``noise" than to evolutionary effects. In their follow-up of the Corwin \& Carney study, \citet{jhs12} noted that ``positive and negative period-change rates with similar size are equally frequent at any period and brightness". For instance, V1 and V10 have comparable mean magnitudes and colors, but the values of $\Delta P/\Delta t$ tabulated for them by Jurcsik et al.~are $-0.417$~d/Myr and $+0.343$~d/Myr, respectively. These stars should both have increasing periods judging from their CMD locations, which are well above the ZAHB, relative to our evolutionary sequences. The same can be said of even brighter variables, and yet, as reported by Jurcsik et al., none of the four brightest RR Lyrae in M\,3 have $\Delta P/\Delta t > 0.0$. This includes the brightest $ab$-type variable in our sample, V42, which has a period-change rate of $-1.132$~d/Myr. In view of such results, and the possible concerns with the mean magnitudes and colors of M\,15 variables mentioned above, we have not attempted to pursue this line of investigation --- though it may be worthwhile to do so when improved data become available. Being able to explain the RR Lyrae in M\,3 and M\,92 so well provides valuable support for our determinations of their distance moduli and ages. We find no compelling evidence for helium abundance variations in either cluster from our analysis of the variable stars, though star-to-star differences at the level of $\delta Y \lta 0.02$ would be very difficult to detect. Our analysis suggests that the faintest HB stars on the blue side of the instability strip and some of the RR Lyrae in M\,15 represent an M\,92-like population. The fact that the difference in magnitude between these stars and the turnoff is identical to within measuring uncertainties in both clusters leads us to conclude that M\,15 and M\,92 are coeval (as most previous studies have found). However, M\,15 appears to contain additional populations of stars with higher helium abundances, up to at least $Y \sim 0.29$ in the vicinity of the instability strip and possibly to higher values along the extended blue tail of the cluster HB. (Fits of isochrones to the turnoff photometry on the assumption of ZAHB-based distance moduli suggest that M\,3, M\,15, and M\,92 all have ages of $\approx 12.5$ Gyr, depending on the assumed CNO abundances. It seems unlikely, in fact, that GCs are as old as the field halo subgiant HD\,140283 --- which is not implausible given the recent discovery of a galaxy at a redshift $z = 11.1$ that seems to have built-up a stellar mass of $\sim 10^9 \msol$ within just $\sim 400$ Myr after the Big Bang; see \citealt{obv16}.) Our explanation of the M\,15 RR Lyrae does raise an important question: why are the ZAHB stars with higher helium abundances distributed to redder colors than those for $Y \approx 0.25$? Possible answers to this question are (i) mass loss rates vary inversely with $Y$, though we are unaware of any empirical or theoretical support for this suggestion, (ii) the abundances of the CNO elements are higher in stars with increased helium abundances, (iii) the stars with higher $Y$ are somewhat younger than those with normal $Y$, or (iv) some combination of these possibilities. \citet{jlj14} have suggested that second generation stars would have enhanced helium and CNO abundances though, in their proposed explanation of the Oosterhoff dichotomy, the different generations of stars would span different color ranges. We see no evidence that this is the case; indeed, a ZAHB for $Y = 0.25$ appears to provide a good fit to the faintest HB stars at all $(V-I_C)_0 \lta 0.3$, where stars with higher helium are presumably also found. More importantly, there is very little spectroscopic evidence for variations in the total CNO abundance in M\,15. Athough they studied only a few giants, \citet{sks97} found that C$+$N$+$O is constant to within the measurement uncertainties in 5 of the 6 stars in their sample. One giant apparently has much higher CNO, given that the derived nitrogen abundance corresponds to [N/Fe] $\sim +1.6$. However, previous studies of much larger samples of upper RGB stars concluded that there are no real differences in the C and N abundances of M\,15 and M\,92 (\citealt{crl82}, \citealt{tlc83}). Both clusters do show a steep decline of [C/Fe] with [Fe/H] (also see \citealt{bbs01}), but C--N and O--N cycling together with deep mixing (\citealt{dd90}, \citealt{lhs93}, \citealt{dv03}) can explain those observations without requiring star-to-star variations in CNO (\citealt{pi88}, \citealt{skp91}, \citealt{cbs05}). Although C$+$N$+$O seems to be approximately constant, the observed variations of CN in present-day main sequence and subgiant stars, as well as star-to-star differences in Mg and Al at any luminosity (e.g., see the relevant studies of M\,92 and M\,15 by \citealt{ksb98}, \citealt{gvb00}, \citealt{cbs05}, \citealt{cbg09b}), is an entirely different issue because they cannot be produced by evolutionary processes within lower mass stars. Such variations must have arisen during an extended period (or possibly successive epochs) of star formation at early times or, if the stars currently observed within a given GC are coeval, they must have been present in the gas out of which the cluster stars formed. That the helium-enhanced stars in M\,15 appear to populate ZAHBs that extend to much redder colors than the ZAHB which fits the faintest stars to the blue of the instability strip suggests that the spread in stellar ages may be larger in M\,15 than in M\,92 or M\,3. [Regardless of which scenario provides the most correct explanation, there is little doubt that H-burning nucleosynthesis at very high temperatures ($\sim 75 \times 10^6$~K, as predicted for supermassive stars; see \citealt{dh14}) is responsible for the observed abundance correlations and anti-correlations, including ratios of the abundances of magnesium isotopes (\citealt{dvh15}).] The potential importance of rotation for our understanding of the HBs in GCs should be kept in mind as well. In the few studies that have been undertaken during the past few decades to measure the rotation of member stars, unexpectedly high rotational velocities have been determined for blue HB stars in the most metal-poor systems (specifically, M\,92; see \citealt{cm97}) and in higher metallicity GCs with extremely blue HBs, including M\,13 (\citealt{pe83}) and NGC\,288 (\citealt{pe85}). A spread of rotational velocities in the precursor red giants would seem to be the most probable cause of the variations in total mass along GC horizontal branches, and the average mass loss could therefore be related to the average rotational velocity in upper RGB stars. The fact that the decline of the surface carbon abundance with increasing luminosity along the giant branch is much more pronounced in M\,13 (\citealt{sm03}) and most, if not all, of the extremely metal-deficient GCs (\citealt{msb08}) than in metal-rich clusters can hardly be a coincidence. \citet{sm79} have shown that such observations can be explained by rotationally driven deep-mixing, which is expected to become less important as the metallicity increases due to the concomitant increase in the mean molecular weight gradient near the H-burning shell. (Red giants in NGC\,288, which has a higher [Fe/H] than M\,13 by $\sim 0.5$ dex, mainly show a bimodality of CN strengths with only a hint of a decline of C and O with increasing luminosity; see \citealt{sl09}.) \smallskip The next paper in the series will carefully examine the differences in the CMDs of M\,3 and M\,13 to try to explain why these two clusters have such different HB morphologies, despite having nearly identical [Fe/H] values. Synthetic HB populations will be presented in this investigation, which will include an analysis of, among other things, the detailed distribution of mass along the observed HBs of these two systems, which is an important and controversial issue; see, e.g., \citet{cd08} and \citet{vc08}.

1607.02088

1607

1607.02331_arXiv.txt

We present a multi-wavelength analysis of five submillimeter sources ($S_\mathrm{1.1mm}$ = 0.54--2.02 mJy) that were detected during our 1.1-mm-deep continuum survey in the SXDF-UDS-CANDELS field (2 arcmin$^2$, 1$\sigma$ = 0.055 mJy beam$^{-1}$) using the Atacama Large Millimeter/submillimeter Array (ALMA). The two brightest sources correspond to a known single-dish (AzTEC) selected bright submillimeter galaxy (SMG), whereas the remaining three are faint SMGs newly uncovered by ALMA. If we exclude the two brightest sources, the contribution of the ALMA-detected faint SMGs to the infrared extragalactic background light is estimated to be $\sim 4.1^{+5.4}_{-3.0}$ Jy deg$^{-2}$, which corresponds to $\sim 16^{+22}_{-12}\%$ of the infrared extragalactic background light. This suggests that their contribution to the infrared extragalactic background light is as large as that of bright SMGs. We identified multi-wavelength counterparts of the five ALMA sources. One of the sources (SXDF-ALMA3) is extremely faint in the optical to near-infrared region despite its infrared luminosity ($L_\mathrm{IR}\simeq 1\times10^{12} \LO$ or SFR $\simeq$ 100 $\MO$ yr$^{-1}$). By fitting the spectral energy distributions (SEDs) at the optical-to-near-infrared wavelengths of the remaining four ALMA sources, we obtained the photometric redshifts ($z_\mathrm{photo}$) and stellar masses ($M_*$): $z_\mathrm{photo} \simeq 1.3$--$2.5$, $M_* \simeq (3.5$--$9.5) \times10^{10}\ \MO$. We also derived their star formation rates (SFRs) and specific SFRs (sSFRs) as $\simeq 30$--$200 \MO$ yr$^{-1}$ and $\simeq 0.8$--$2$ Gyr$^{-1}$, respectively. These values imply that they are main-sequence star-forming galaxies.

Determining the contributors of dust-obscured galaxies to the cosmic star-formation rate density (cosmic SFRD) is a major goal of deep surveys at far-infrared, millimeter, submillimeter, and radio wavelengths. In fact, deep surveys using the {\it Infrared Space Observatory} ({\it ISO}), {\it AKARI}, and the {\it Herschel Space Observatory} ({\it Herschel}) have revealed that dusty star-forming galaxies largely dominate the cosmic SFRD up to the redshift $z\approx$ 1--3 \citep[e.g.,][]{takeuchi2005, goto2011, brugarella2013}. Over the past decade, a series of wide-area surveys performed at millimeter/submillimeter wavelengths using single-dish telescopes have revealed many bright submillimeter galaxies (SMGs) \citep[e.g.,][and references therein]{samil1997, hughes1998, barger1998, blain2002, greve2004, weiss2009, scott2010, hatsukade2011, casey2013, umehata2014} with observed flux densities larger than a few mJy at millimeter/submillimeter wavelengths. They have large total infrared (IR; {rest-frame 8--1000 $\micron$}) luminosities ($L_{\rm IR}\sim10^{12-13}\ \LO$) powered by dust-obscured star formation \citep[e.g.,][]{alexamder2005, laird2010}, and their redshift distribution peaks are at $z\approx$ 2.2--2.5 \citep[e.g.,][]{capman2005,simpson2014}. Their extreme star formation rates (SFRs $\gtrsim$ a few 100--1000 $\MO$yr$^{-1}$) make them non-negligible contributors to cosmic star formation \citep[e.g.,][]{hughes1998, casey2013, wardlow2011, swinbank2014}. However, the contribution of SMGs detected by single-dish surveys to the infrared extragalactic background light, which is believed to be the integrated infrared emissions from all extragalactic sources along the line of sight, is 20--40\% at 850 $\micron$ \citep[e.g.,][]{eales1999, coppin2006, weiss2009} and 10--20\% at 1.1 mm \citep[e.g.,][]{hatsukade2011, scott2012}. Thus, the bulk of infrared extragalactic background light remains unresolved with single-dish telescopes. By using stacking analysis of $K$-selected galaxies, \citet{greve2010} found that these galaxies contribute $\simeq16.5\%$ to the infrared extragalactic background light at 870 $\micron$, although individual source properties remained unexplored in this stacking analysis. The advent of the Atacama Large Millimeter/submillimeter Array (ALMA), which offers high sensitivity and angular resolution capabilities, has allowed a fainter population of SMGs to be unveiled below the confusion limit of single-dish telescopes. Here, we refer to the fainter population of SMGs with flux densities of $\sim$ 0.1--1 mJy at 1.1--1.3 mm as ``faint SMGs.'' Their estimated contributions to the infrared extragalactic background light are $\simeq$ 50\%--80\% \citep{hatsukade2013, ono2014, oteo2015}. Deeper number counts down to $\sim$ 0.02 mJy have recently obtained by \citet{carniani2015} and \citet{fujimoto2015}, who claimed that $\sim$ 100\% of the infrared extragalactic background light is resolved at 1.2--1.3 mm. These results suggest that faint SMGs can play an important role in the cosmic star-formation activities at high redshifts. However, their contributions to the cosmic SFRD are still unknown because of the lack of redshift information. Single-dish telescopes have been used in attempts to detect faint SMGs with the aid of gravitational magnification by lensing clusters. For example, \citet{knudsen2008} constrained the faint end ($S_\mathrm{850 \micron} \simeq 0.1$ mJy) of the 850 $\micron$ number counts by using cluster magnification. \citet{chen2014} performed follow-up observations of these lensed faint SMGs by using the submillimeter array \citep[SMA;][]{moran1998} and implied that there are many faint SMGs that are faint at optical/near-infrared wavelengths and have been missed in deep optical/near-infrared surveys. In lens surveys, however, the effective sensitivity comes at the cost of a reduced survey volume; the effective (source-plane) area within sufficient magnification for faint SMG detection is only $\sim 0.1$ arcmin$^2$ for a typical rich cluster \citep{knudsen2008}.\footnote{\cite{fujimoto2015} also used gravitational magnification from lensing clusters, and their survey area was $\sim$ 0.5 arcmin$^2$ (source-plane).} This also increases the cosmic variance uncertainty \citep[e.g.,][]{robertson2014}. Therefore, it is still necessary to obtain wide ($> 1$ arcmin$^2$) and deep (1$\sigma\lesssim$ 0.1 mJy) blank/unlensed field surveys at a higher angular resolution to gain a better understanding of faint SMGs and their true contributions to the cosmic SFRD. Another key issue to understand galaxy evolution is the star formation properties of galaxies. Star-forming galaxies have a correlation between their stellar masses and SFRs, which is defined as a main sequence \citep[e.g.,][]{daddi2007,rodighiero2011,rodighiero2014,schriber2015}. Main sequence star-forming galaxies are ``normal'' star-forming galaxies selected by optical/near-infrared colors [e.g., {\it BzK} galaxies, typical specific SFRs (sSFRs) $\sim1$ Gyr$^{-1}$ at $z\sim$ 1.4--2.5; \citet{rodighiero2011}]. However, outliers of the correlation exist with higher sSFRs than those of main sequence star formation galaxies \citep[sSFRs $\gtrsim10^{1}$--$10^2$ Gyr$^{-1}$ at $z\sim$ 1.4--2.5;][]{rodighiero2011}. These outliers are often referred to as starburst galaxies. Many bright SMGs are classified as starburst galaxies or the high-mass end of main sequence galaxies \citep[e.g.,][]{takagi2008, dacunha2015}. However, it is not understood whether faint SMGs are on or above the main sequence because the stellar masses of these faint SMGs have not yet been measured. Understanding the star-forming properties of faint SMGs is also helpful to unveil the evolution of the cosmic SFRD because they are thought to be the main contributor to the infrared extragalactic background light \citep[e.g.,][]{carniani2015,fujimoto2015}. In this paper, we present spectral energy distributions (SEDs) for optical-to-radio counterparts to five submillimeter sources that were detected in our 2 arcmin$^2$ 1.1-mm-deep survey of the Subaru/\textit{XMM-Newton} Deep Survey Field \citep[SXDF;][]{furusawa2008} using ALMA (Project ID: 2012.1.00756.S, PI: K. Kohno) to understand their contribution to the infrared extragalactic background light and cosmic SFRD. We also discuss their multi-wavelength properties. Our survey field was also covered by the UKIRT Infrared Deep Sky Survey Ultra-Deep Survey \citep[UDS;][]{lawrence2007} and Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey \citep[CANDELS;][]{grogin2011,Koekemoer2011}. This paper is structured as follows. Section \ref{multi-image} presents the ALMA observations and multi-wavelength data used in this study. Section \ref{counterpart} presents the results of the multi-wavelength counterpart identification of the ALMA sources. Section \ref{SED} derives the photometric redshift, stellar mass ($M_*$), SFRs, and sSFRs and presents the optical-to-radio SEDs. Sections \ref{discussion} and \ref{summary} are devoted to the discussion and summary, respectively. Throughout this paper, we assume a $\Lambda$ cold dark matter cosmology with $\Omega_M = 0.3$, $\Omega_\Lambda = 0.7$, and $H_0 =$ 70 km s$^{-1}$ Mpc$^{-1}$. All magnitudes are given according to the AB system.

\label{discussion} \subsection{Contribution to the infrared extragalactic background light} Owing to the high sensitivity and high angular resolution observations with ALMA, 50\%--100\% of the infrared extragalactic background light has been claimed to be resolved if we go down to $\sim0.1$--0.02 mJy \citep[e.g.,][]{hatsukade2013, ono2014, carniani2015, fujimoto2015}. From the summation of the 1.1 mm flux densities of all of our ALMA sources and the survey area (2 arcmin$^2$), the contribution of the ALMA sources to the infrared extragalactic background light was estimated to be$\sim$ $10^{+6}_{-4}$ Jy deg$^{-2}$, which corresponds to $\sim 40^{+24}_{-16}\%$ of the infrared extragalactic background light obtained by \citet{fixsen1998} using the COsmic Background Explore (COBE) satellite \citep[25$^{+22}_{-13}$ Jy deg$^{-2}$;][]{carniani2015} or $\sim54^{+32}_{-22}\%$ if we adopt the the COBE measurement made by \citet[][18.5 Jy deg$^{-2}$]{puget1996}. In the subsequent discussion, we adopt the Fixsen et al. value for the infrared extragalactic background light, but we caution that there exist uncertainties (likely systematic) in the COBE measurements. Because of the numbers of our sources, we used the Poisson uncertainty values presented by \citet{gehrels1986}. The completeness in the flux range of ALMA sources was $\sim100\%$ \citep{hatsukade2016}. Note that this value can be overestimated because our observation field was selected to include a single bright SMG (Ikarashi et al.~in preparation) and a chain of HAEs at $z=2.5$ \citep{tadaki2013,tadaki2015}. Given that SXDF-ALMA1 and 2 were identified as an AzTEC/ASTE source (Ikarashi et al.~in preparation) and HAEs \citep{tadaki2013}, it is better to exclude SXDF-ALMA1 and 2 when discussing the real contribution of the ALMA sources to the infrared extragalactic background light. The contributions of SXDF-ALMA3, 4, and 5 to the infrared extragalactic background light were estimated to be $\sim$ $4.1^{+5.4}_{-3.0}$ Jy deg$^{-2}$, which corresponds to $\sim 16^{+22}_{-12}\%$ of the infrared extragalactic background light obtained by the COBE satellite. This suggests that their contribution to the infrared extragalactic background light is as large as that of bright SMGs \citep[$S_\mathrm{1.1mm}\geq$ 1.0 mJy, $\sim$ 2.9 Jy deg$^{-2}$;][]{hatsukade2011}. Although our survey area was small and may have been affected by cosmic variance, these results suggest that bright ($S_\mathrm{1.1mm}\geq$ 1 mJy) sources and faint SMGs with 0.5 mJy $< S_\mathrm{1.1 mm}\lesssim$ 1.0 mJy, which is the flux range of the ALMA source, seem to contribute $\sim 28^{+22}_{-12}\%$ to the infrared extragalactic background light. These results suggest that faint SMGs with $S_\mathrm{1.1 mm}\lesssim$ 0.5 mJy are major contributors to the infrared extragalactic background light. The results of stacking analysis of near-infrared selected galaxies with $S_\mathrm{1.1 mm}\lesssim$ 0.5 mJy will be discussed in our upcoming paper (Wang et al.~in preparation). \subsection{Contribution to the cosmic SFRD} Substantial attempts have been made with {\it Herschel} to resolve the redshift evolution of the contribution of infrared selected galaxies up to the $z=$ 3 \citep{brugarella2013}. \citet{brugarella2013} estimated the cosmic infrared SFRD from the infrared luminosity functions inferred from {\it Herschel} observations and found that the contribution peaks at $z=1.35$, which accounts for 89\% of the total cosmic SFRD of $1.1\times10^{-1}$ $\MO$ yr$^{-1}$ Mpc$^{-3}$ (using Chabrier IMF). However, they used the extrapolate infrared luminosity functions below the confusion limit of {\it Herschel} (e.g., $L_\mathrm{IR}\lesssim10^{12}\ \LO$ at $z\simeq 2$) to estimate the cosmic infrared SFRD. \citet{wardlow2011} derived infrared luminosity functions of SMGs detected by the LABOCA Extended Chandra Deep Field South surveys \citep[LESS;][]{weiss2009}. However, they also did not investigate the luminosity range of $L_\mathrm{IR}\lesssim10^{12}\ \LO$. ALMA sources with (photometric) redshifts allow the constraint on the contribution from a faint ($L_\mathrm{IR}\lesssim10^{12}\LO$) population of star-forming galaxies. In addition, the contribution from galaxies undetected by {\it Herschel} can be estimated. We simply assume that the ALMA sources lie in the redshift interval of $1<z<4$ (co-moving volume: $V_{\rm com}\sim$ 1.9 $\times$ 10$^{4}$ Mpc$^3$) to cover the redshift uncertainties of the ALMA sources. Then, we estimate the contribution from all of the ALMA sources, including SXDF-ALMA1 and 2, on the basis of SFRs simply derived from the 1.1 mm flux densities (section \ref{sfr}). Considering the uncertainty in $T_\mathrm{dust}$ (see section \ref{sfr} and Appendix), the inferred infrared SFRD for $1 < z < 4$ is $\sim$ (0.9--5) $\times10^{-2}$ $\MO$ yr$^{-1}$ Mpc$^{-3}$, which accounts for $\sim$ 10--70\% of the average infrared SFRD at $0.9<z<3.6$ as estimated by {\it Herschel} \citep[using Chabrier IMF, $\sim 7\times10^{-2}$ $\MO$ yr$^{-1}$ Mpc$^{-3}$;][]{brugarella2013}. If we exclude SXDF-ALMA1 and 2 to avoid the contribution of the known AzTEC source, then the infrared SFRD for $1 < z < 4$ is estimated to be $\sim$ (0.3--2) $\times10^{-2}$ $\MO$ yr$^{-1}$ Mpc$^{-3}$. The inferred cosmic infrared SFRD is similar to that of bright SMGs ($S_\mathrm{870\micron}>$ 4 mJy) at $z\simeq$ 2--3 [$\sim$ (1--2) $\times10^{-2}$ $\MO$ yr$^{-1}$ Mpc$^{-3}$ using Chabrier IMF; \cite{wardlow2011}]. These results imply that the ALMA sources play an important role in the cosmic SFRD, even if we exclude the contribution of the known AzTEC source at $1 < z < 4$. Note that our results can be affected by the cosmic variance and clustering because of our small survey area. Therefore, future ALMA large surveys will provide a stronger constraint on the role of faint SMGs in the cosmic SFRD. \citet{chen2014} inferred that there are many submillimeter sources that are difficult to detect in deep optical/near-infrared surveys like SXDF-ALMA3. Indeed, submillimeter sources which have no counterparts at optical/near-infrared wavelengths have been reported \citep[e.g.,][]{wang2007, smolcic2012, simpson2014, dunlop2016}. However, their real contributions to the cosmic infrared SFRD is still uncertain. The contribution of SXDF-ALMA3 to the cosmic infrared SFRD may be $\sim$ 0.1--1 $\times10^{-2}$ $\MO$ yr$^{-1}$ Mpc$^{-3}$ or $\sim$ 1\%--10\% of the average infrared SFRD at $0.9<z<3.6$ as estimated by {\it Herschel} ($\sim 7\times10^{-2}$ $\MO$ yr$^{-1}$ Mpc$^{-3}$) if this object lies somewhere in the redshift interval of $1<z<4$. \subsection{Star formation properties of the ALMA detected sources} Determining the star-forming properties of faint SMGs is important to understand the evolution of the cosmic SFRD because they are main contributors to the cosmic SFRD. We investigated the star formation mode of faint SMG counterparts to check whether or not their star-forming properties are similar to starburst galaxies or not. Figure \ref{m-sfr} plots the total SFRs (SFR$_\mathrm{UV}$ + SFR$_\mathrm{IR}$) of SXDF-ALMA1, 2, 4, and 5 as functions of their stellar mass. We also show the average values of {\it BzK} galaxies derived by the PACS stacking analysis \citep{rodighiero2015}, SMGs identified in ALESS surveys \citep[][]{hodge2013} at $1.3 < z < 2.5$, and faint 1.3 mm sources detected with ALMA \citep{hatsukade2015}. We plotted SFRs of the ALESS sources obtained by \citet{dacunha2015} by fitting SEDs at ultraviolet to radio wavelengths. This figure shows that SXDF-ALMA1, 2, 4, and 5 are located in the main sequence. This means that they are more like ``normal'' star-forming galaxies rather than extremely starburst galaxies. These results are consistent with those of \citet{koprowski2014} and \citet{hatsukade2015}. Note that the total SFR of SXDF-ALMA1 should be treated as an upper limit because we used the 3$\sigma$ upper limit value as its SFR$_\mathrm{UV}$. However, this does not affect our results because the SFR$_\mathrm{UV}$ of SXDF-ALMA1 is negligible compared to its SFR$_\mathrm{IR}$ (see table \ref{table2}). Figure \ref{m-sfr} also shows the constraints of the stellar mass and SFR of SXDF-ALMA3. The results imply that SXDF-ALMA3 is an starburst galaxy with a small stellar mass compared to bright SMGs \citep[$M_*\sim9.0\times10^{10}\ \MO$ using Chabrier IMF;][]{hainline2011}. Submillimeter sources such as SXDF-ALMA3 have been missed in previous deep optical/NIR surveys and submillimeter single-dish surveys. Future spectroscopic identification of such sources using ALMA is highly encouraged. Finally, we compared our results with the theoretical predictions obtained by recent simulations and semi-analytical models. \citet{bethermin2012} empirically predicted the number counts at far-infrared and millimeter wavelengths from mid-infrared and radio number counts and suggested that galaxies with $S_\mathrm{1.1 mm}\lesssim$ 1 mJy are more likely to be associated with main sequence star-forming galaxies by using the SED library based on {\it Herschel} observations. From a theoretical point of view, \citet{hayward2013} predicted the number counts at submillimeter wavelengths from a semi-empirical model with three-dimensional hydrodynamical simulations and three-dimensional dust radiative transfer and also suggested that galaxies with $S_\mathrm{1.1 mm}<$ 1 mJy are more likely to be associated with main sequence star-forming galaxies. These predictions are consistent with our results that two of the three faint SMGs (SXDF-ALMA4 and 5) are main sequence star-forming galaxies, as shown in figure \ref{m-sfr}.

1607.02331

1607

1607.00334_arXiv.txt

{The behaviour of a massive, non-interacting and non-minimally coupled quantised scalar field in an expanding de Sitter background is investigated by solving the field evolution for an arbitrary initial state. In this approach there is no need to choose a vacuum in order to provide a definition for particle states, nor to introduce an explicit ultraviolet regularization. We conclude that the expanding de Sitter space is a stable equilibrium configuration under small perturbations of the initial conditions. Depending on the initial state, the energy density can approach its asymptotic value from above or below, the latter of which implies a violation of the weak energy condition. The backreaction of the quantum corrections can therefore lead to a phase of super-acceleration also in the non-interacting massive case.} \emailAdd{tommi.markkanen@kcl.ac.uk} \emailAdd{a.rajantie@imperial.ac.uk} \begin{document}

An eternally expanding de Sitter space has been shown in the classical context to be a stable attractor solution for a variety of cosmological models with a cosmological constant in the context of cosmic no-hair theorems \cite{Wald:1983ky,Starobinsky:1982mr,Gibbons:1977mu,Hawking:1981fz} (see \cite{Schmidt:2004xv,Brandenberger:2016uzh} for review and references). For quantized theories de Sitter space has been analyzed in terms of back-reaction of cosmological fluctuations in \cite{Tsamis:1996qq2,Tsamis:1996qq3,Geshnizjani:2002wp} quantum gravity in \cite{Tsamis:1996qq,Tsamis:1996qq0,Tsamis:1996qq1,Antoniadis:1985pj} and quantum fields on a classical curved background in \cite{ford,Bousso:2001mw,Polyakov:2007mm,Polyakov:2007mm1, Polyakov:2007mm2,akh,Marolf:2010zp,Marolf:2010nz,Mottola:1984ar,Anderson:2013ila,Anderson:2013ila1,Habib:1999cs, Anderson:2013ila2,Onemli:2004mb,Onemli:2002hr,Dabrowski:2014ica,Antoniadis:2006wq,Anderson:2000wx,Koivisto:2010pj, Albrecht:2014hxa,Glavan:2014uga,Glavan:2015cut,Jiang:2016nok}. In \cite{Anderson:2013ila,Anderson:2013ila1} a thorough investigation of de Sitter space in a global coordinate system was presented and as the main finding it was stated that the global de Sitter space with a quantized scalar field is unstable thus reinforcing the arguments made already in \cite{Mottola:1984ar}. Quite generically a quantum field coupled to a curved background leads to the creation of quanta which is often called cosmological particle production \cite{Parker:1968mv,Parker:1968mv1,Parker:1968mv2,Parker:1968mv3}, or in the case of a black hole, Hawking radiation \cite{Hawking:1974sw,Hawking:1974rv}. Also in flat space there are closely related processes: in Schwinger pair creation particles are produced due to a background electric field \cite{Schwinger:1951nm} (for a recent review, see \cite{Gelis:2015kya}) and in the Unruh effect because of constant acceleration \cite{Unruh:1976db}. A crucial conclusion from Hawking radiation is that black holes may be viewed as thermodynamic objects, which is manifested as the relation between the event horizon area and entropy of a black hole \cite{Bekenstein0,Bekenstein1,Bekenstein2,Bardeen:1973gs,Bekenstein3,Hawking:1976de}. Importantly, analogous thermal properties have also been assigned space-times with boundaries \cite{Gibbons:1977mu}, which can be assigned entropy and temperature, a connection made even more profound by the observation that in many cases it is possible to derive the Einstein equations from thermodynamics \cite{Jacobson:1995ab,Padmanabhan:2002sha,kof,cai,Padmanabhan:2009vy}. However, it is not obvious how far one may take the thermodynamic analogy of de Sitter space as discussed in \cite{Davies:2003me,Birrell:1982ix}. Based on the purely thermodynamic description of \cite{Gibbons:1977mu}, de Sitter space has been investigated for example in \cite{Antoniadis:2006wq,Mottola:1985qt,Sekiwa:2006qj,Padmanabhan:2003gd}. The thermal features in de Sitter space are often studied in the static coordinates with explicitly a horizon and a coordinate singularity much like in the Schwarzschild metric. In such coordinates de Sitter space appears then essentially as an insulated cavity that is internally in equilibrium containing a thermal spectrum of particles \cite{kof,Kaloper:2002cs,Susskind:2003kw}. However, our current view of the Universe is best described as that of a co-moving observer in the flat FLRW coordinates $ds^2=-dt^2+a(t)^2d\mathbf{x}^2$, which warrants a detailed investigation of de Sitter space in this specific coordinate system. In this work we set out to investigate the evolution of a quantized scalar field in expanding de Sitter space and its implications for the backreaction in the flat FLRW coordinates. We study the backreaction in the semi-classical approach with a strictly classical geometry. Our model consists of a quantized, massive, noninteracting and non-minimally coupled scalar field in a classical spacetime curved by vacuum energy. One of our key findings is that quantum energy-momentum tensor of a non-interacting massive scalar field can violate the weak energy condition. A similar result was found earlier for a massless, minimally coupled and self-interacting model in \cite{Onemli:2004mb,Onemli:2002hr}. In that case the origin of such a violation was the generation of a nonperturbative mass via interactions, as described in \cite{Starobinsky:1994bd}, see also the related discussion in \cite{Winitzki:2001fc}. A violation of the weak energy condition in the massive and non-interacting model has different origins and highlights the generality of the effect. Furthermore, for our study we define a new covariant renormalization prescription for de Sitter space. Our method has all the benefits of the popular adiabatic subtraction technique \cite{Parker:1974qw,Parker:1974qw1,Bunch:1980vc}, but it does not require the lengthy expressions needed in adiabatic subtraction and is physically more meaningful. Our units are $\hbar\equiv c\equiv k_B\equiv1$ and we use the (+,+,+) conventions of \cite{Misner:1974qy}.

\label{sec:con} The exponentially expanding de Sitter space has significant relevance for the early and late time Universe. In this work we have studied the evolution of a massive, non-interacting and non-minimaly coupled scalar field in such a spacetime. The Bunch-Davies vacuum state enjoys a special status in de Sitter space and we have shown here that the de Sitter invariance of this state is independent of the details of renormalization and can be understood as a manifestation of covariant conservation. Importantly, although the behaviour of the modes in the Bunch-Davies state can be interpreted to constantly go through a particle creation process as indicated by a nontrivial Bogolubov transformation this is not visible in the semi-classical backreaction: It bears no sign of a density from classical particles and implies strictly $w=-1$ for the equation of state. We also studied the behaviour of the energy-momentum when starting from non-de Sitter invariant initial conditions. It is known that the concept of a particle is non-trivial in curved space and in our approach there is no need to give the definition of a particle. Instead, the evolution is fully determined by the initial condition for the field modes. By effectively using the Bunch-Davies vacuum mode contributions as renormalisation counterterms, we were able to investigate the time dependence of the energy-momentum tensor without the need for ultraviolet regularisation. Via analytic and numerical examples, we saw that although it is possible to obtain a initial period where the equation of state is $w\neq-1$ all initial conditions approach the Bunch-Davies state as an equilibrium configuration on timescales of $\sim 1/H$. This is in full agreement with the conclusions of \cite{Habib:1999cs,Anderson:2000wx,Koivisto:2010pj}. Finally we considered the semi-classical backreaction onto the Hubble rate. Due to the tendency of the scalar field to equilibrate in de Sitter space, we concluded that the expanding patch in the FLRW coordinates is a stable configuration under small perturbations of the initial conditions. As an important result from a fundamental point of view we showed that the weak energy condition can in principle be violated by quantum corrections i.e. that the equation of state can be $w<-1$, even though this effect is quickly exponentially suppressed. This was verified at the late time limit via an analytic calculation as well as a numerical analysis. Thus in support of the analysis of a massless self-interacting scalar field in \cite{Onemli:2004mb,Onemli:2002hr} we obtain that also for a massive non-interacting scalar field quantum corrections may lead to a period of superacceleration with $\dot{H}>0$.

1607.00334

1607

1607.08213_arXiv.txt

In this paper, we study in detail the role of general relativity on the global dynamics of giant pulsar glitches as exemplified by Vela. For this purpose, we carry out numerical simulations of the spin up triggered by the sudden unpinning of superfluid vortices. In particular, we compute the exchange of angular momentum between the core neutron superfluid and the rest of the star within a two-fluid model including both (non-dissipative) entrainment effects and (dissipative) mutual friction forces. Our simulations are based on a quasi-stationary approach using realistic equations of state (EoSs) following \cite{sourie2016numerical}. We show that the evolution of the angular velocities of both fluids can be accurately described by an exponential law. The associated characteristic rise time $\tau_{\text{r}}$, which can be precisely computed from stationary configurations only, has a form similar to that obtained in the Newtonian limit. However, general relativity changes the structure of the star and leads to additional couplings between the fluids due to frame-dragging effects. As a consequence, general relativity can have a large impact on the actual value of $\tau_{\text{r}}$: the errors incurred by using Newtonian gravity are thus found to be as large as $\sim 40 \%$ for the models considered. Values of the rise time are calculated for Vela and compared with current observational limits. Finally, we study the amount of gravitational waves emitted during a glitch. Simple expressions are obtained for the corresponding characteristic amplitudes and frequencies. The detectability of glitches through gravitational wave observatories is briefly discussed.

\label{intro} Pulsars are very compact stars rotating rapidly with exceptionally stable periods spanning from $\sim$ 1.4 milliseconds to a few seconds. Nevertheless, some pulsars exhibit sudden increases in their observed angular velocity $\Omega$, with relative amplitude $\Delta \Omega/ \Omega$ ranging between $\sim 10^{-11}$ and $\sim 10^{-5}$ \citep{wong2001observations, espinoza2011study}. These spin-up events, known as \emph{glitches}, are usually followed by a slow relaxation on time scales up to months or years and are sometimes accompanied by abrupt changes in the pulsar spin-down rate, $\Delta\dot{\Omega}/\dot{\Omega}\sim 10^{-4} - 10^{-2}$ (we use a dot to denote time derivative). Presently, 472 glitches have been detected in 165 pulsars\footnote{http://www.jb.man.ac.uk/pulsar/glitches.html.} \citep{espinoza2011study}, with angular velocities ranging from 0.09 Hz to 327 Hz (see, e.g., the ATNF Pulsar Database\footnote{http://www.atnf.csiro.au/research/pulsar/psrcat.}; \citet{manchester2005australia}). At least two distinct glitching behaviours have been identified \citep{espinoza2011study, yu2013detection}: (i) quasi-periodic giant glitches with a very narrow spread in size around $\Delta\Omega/\Omega \sim 10^{-6}$, and (ii) smaller glitches of various sizes at random intervals of time. The most emblematic pulsar of the first kind is Vela (PSR~B0833--45) with a rotation frequency $f=\Omega /(2\pi) \simeq 11.19$~Hz \citep{dodson2007two} corresponding to a period $P=1/f\simeq 89$~ms. Since its discovery in 1969, 19 glitches have been detected so far every $\sim$ 2-3 years~\citep{espinoza2011study}. The second type of glitching pulsars is exemplified by the Crab (PSR~B0531+21) with a rotation frequency $f \simeq 29.95$~Hz ($P\simeq 33$~ms). Since the first detections of glitches \citep{radhakrishnan1969detection, reichley1969observed}, different mechanisms have been proposed to explain these events (see, e.g., the review by \citet{haskell2015models}). A glitch is nowadays commonly thought as the manifestation of an internal process, except possibly for highly magnetised neutron stars for which some evidence of magnetospheric activity have been found (e.g. \citet{archibald2013anti, keith2013connection, antonopoulou2015unusual}). The interior of neutron stars can thus be probed using observations of pulsar glitches. Glitches were first suggested to arise from crustquakes \citep{ruderman1969neutron, baym1971neutron}. Following this idea, the presence of a solid crust (which crystallized when the star was young and rapidly rotating) prevents readjustments of the stellar shape, as the star spins down due to electromagnetic emission. Crustal stresses thus build up, until the crust cracks and the star suddenly adopts a more spherical shape. The resulting reduction of the moment of inertia leads the pulsar to spin up, assuming conservation of angular momentum. This scenario can account for small glitches, such as those exhibited by the Crab pulsar. However, as pointed out by \citet{ruderman1969neutron}, this mechanism fails to predict the occurrence frequency of giant glitches, as observed in the Vela pulsar. Giant glitches are generally thought to be the manifestation of superfluid matter inside neutron stars, as suggested by the very long time scales observed during post-glitch relaxations \citep{baym1969superfluidity}. From theoretical calculations, the interior of a neutron star is expected to contain an isotropic neutron superfluid in the inner crust, an anisotropic neutron superfluid in the outer core, and possibly other superfluid species in the inner core (see, e.g., \cite{page2013stellar}). In a seminal work, \citet{anderson1975pulsar} proposed that glitches themselves could be triggered by the sudden unpinning of neutron superfluid vortices. The idea is the following. It is well-known from laboratory experiments (see, e.g., \citet{yarmchuk1979observation, abo-shaeer2001observation, zwierlein2005vortices}) that a superfluid can only rotate by forming an array of quantized vortices, each carrying a quantum $\hbar$ of angular momentum, where $\hbar$ is the Planck-Dirac constant. The neutron superfluid present in the core and the inner crust of neutron stars is thus expected to be threaded by a huge number of vortex lines, with a mean surface density given by \begin{equation} \label{surfdens_vortex} n_{v} \, (\text{cm}^{-2}) = \frac{4 m_{\n} \Omega_{\n}}{h} \simeq\frac{10^7}{P(\text{ms})}\, , \end{equation} where $h=2\pi \hbar$ is the Planck constant, $m_{\n}$ is the neutron rest mass and the coarse-grained averaged angular velocity $\Omega_{\n}$ of the neutron superfluid is approximated by that of the star \citep{ginzburg1965}. The neutron superfluid is supposed to be weakly coupled to the rest of the star by so-called mutual friction forces arising from the dissipative forces acting on individual vortices~\citep{alpar1984rapid}. Due to the spin down of the star induced by the electromagnetic torque, vortices tend to move away from the rotation axis. The key assumption of vortex-mediated glitch theories is that vortices can pin to nuclear clusters in the inner crust~\citep{anderson1975pulsar} and/or to quantized magnetic flux tubes in the core if protons form a type II superconductor \citep{baym1969superfluidity,sauls1989superfluidity, ruderman1998neutron}. In such case, the neutron superfluid is decoupled from the rest of the star, and can rotate more rapidly, as schematically illustrated on Fig.~\ref{fig:evol}. The lag $\delta\Omega = \Omega_{\n} - \Omega >0 $ induces a Magnus force on the vortices. The larger the lag, the stronger the force. For some critical value $\delta \Omega_0$ of the lag, vortices will suddenly unpin, the superfluid will spin down and, by conservation of angular momentum, the rest of the star will spin up leading to the observed glitch (see Fig.~\ref{fig:evol}). During the subsequent relaxation, these vortices are thought to progressively repin and a lag can grow anew. This vortex-mediated scenario is supported by laboratory experiments \citep{tsakadze1980properties} and the ability of the vortex creep model to reproduce the post-glitch relaxations in different pulsars \citep{alpar1984vortex, alpar1984vortexII, alpar1993postglitch, alpar1996postglitch, gugercinoglu2014vortex}. \begin{figure} \center \includegraphics[width = \columnwidth]{schema4} \caption{Schematic time evolution of the observed angular velocity $\Omega$ of the pulsar and of the angular velocity $\Omega_{\n}$ of the interior neutron superfluid over two glitch events.} \label{fig:evol} \end{figure} It is noteworthy to mention that the two mechanisms described above are not necessarily independent. Indeed, starquakes can be induced by the presence of superfluids in neutron star interiors, whether superfluid vortices are pinned~\citep{ruderman1991neutron} or not~\citep{carter2000centrifugal, chamel2006effect}. In turn, sudden motions of vortices can be triggered by quakes~\citep{ruderman1991neutron,chau1993correlated, alpar1996postglitch,eichler2010}. Although mesoscopic studies of large collections ($\sim10^2-10^4$) of vortices provide useful insight (e.g. \cite{warszawski2011gross, warszawski2013knock}), simulating pulsar glitches requires to follow the dynamics of all the superfluid vortices contained in the star. Given their huge number, $\sim 10^{17}$ for Vela, the overall transfer of angular momentum between the neutron superfluid and the rest of the star can be studied using a smooth-averaged hydrodynamic approach, still involving microscopic parameters determined by the local dynamics of individual vortices (see, e.g., \cite{bulgac2013strength} and references therein). Whereas the general relativistic framework for describing starquakes was developed a long time ago~\citep{carter1975relativistic}, the general relativistic formulation of the vortex-mediated glitch model is more recent~\citep{langlois1998differential}. As a matter of fact, most global numerical simulations of pulsar glitches have been performed within the Newtonian framework~(e.g., \cite{larson2002simulations, peralta2006transitions, sidery2010dynamics,haskell2012modelling}). Recently, \cite{seveso2012effect} and \cite{antonelli2016axially} have developed a non-relativistic hydrodynamic model for describing the different stages of the glitch phenomenon based on the static structure of the neutron star computed in general relativity. However, general relativity could also play an important role for the global dynamics of glitches. Furthermore, general relativity is essential to determine the amount of gravitational waves associated with glitch events. Observations of gravitational waves are of particular interest since they could potentially provide additional information on the glitch phenomenon (see, e.g., \cite{stopnitzky2014gravitational, haskell2015models} and references therein). In this paper, we present global numerical simulations of vortex-mediated pulsar glitches. We focus on the spin-up stage regardless of the glitch triggering mechanism. On the other hand, we study the glitch dynamics in full general relativity. We also derive the associated gravitational wave characteristic amplitudes and frequencies using the standard quadrupole formula. The paper is organized as follows. We start by presenting, in Section~\ref{typical}, the different assumptions on which our model is based. In Section~\ref{sec:evol_eq}, we introduce the evolution equations governing the transfer of angular momentum in the interior of a pulsar during a glitch. Results of stationary rotating configurations are discussed in Section~\ref{stat_conf}. In Section~\ref{numerical}, we detail the numerical procedure underlying our simulations. Results for the glitch rise time are presented and discussed. We study the emission of gravitational waves in Section~\ref{gws}. Finally, we conclude in Section~\ref{conclusion}.

\label{conclusion} In this paper, we have studied in detail the impact of general relativity on the global dynamics of giant pulsar glitches as observed in Vela. We have carried out numerical simulations of the spin up triggered by the sudden unpinning of quantized vortices. To this end, we have computed the exchange of angular momentum between the neutron superfluid in the core and the rest of the star within a two-fluid model including neutron-proton entrainment effects. Both fluids were assumed to be coupled by mutual friction arising from dissipative forces acting on individual vortices. Since the hydrodynamical time scale is typically much smaller than the glitch rise time, we have described the time evolution of the two fluids by a sequence of quasi-stationary axisymmetric rigidly rotating configurations following \cite{sourie2016numerical}. We have calculated the mutual friction torque considering straight vortices arranged on a regular array, following \cite{langlois1998differential}. In order to get some physical insight, we first solved analytically the dynamical equations by expressing the change in the lag as $\delta \dot{\Omega} /\delta \Omega \approx -1/\tau_{\text{r}}$, where the characteristic spin-up time scale $\tau_{\text{r}}$ can be expressed in a form similar to that obtained in the Newtonian limit (see, e.g., \cite{carter2001relativistic, sidery2010dynamics}). However, general relativity not only changes the structure of the star, but also impacts the fluid dynamics. In particular, frame-dragging effects induce additional fluid couplings of the same form as the entrainment arising solely from neutron-proton interactions. For all these reasons, general relativity can change substantially the glitch rise time. To test the validity of this analytical approach and to assess the importance of general relativity, we have also solved numerically the equations governing the transfer of angular momenta. For this purpose, two different kinds of inputs are needed: macroscopic quantities (the rotation frequency of the star, the glitch amplitude and the neutron star mass) and microscopic properties (the EoS and the stellar-averaged mutual friction coupling $\bar{\mathcal{B}}$). We have explored in detail various stellar configurations, using two different relativistic mean-field EoSs and considering the observed properties of glitching pulsars. The results obtained by numerical simulations are found to be very well reproduced by the analytical approximation. In particular, the glitch rise time $\tau_\text{r}$ can thus be expressed in terms of the moments of inertia of the fluids, the stellar rotation rate and $\bar{\mathcal{B}}$, which can be obtained from stationary configurations. Furthermore, we have studied the effects of general relativity on $\tau_\text{r}$ by using two different polytropic EoSs of the kind previously introduced by \cite{prix2005relativistic}. Both the effects of general relativity on the structure of the star and on the fluid couplings are found to be important and therefore realistic simulations of the global glitch dynamics should be carried out in full general relativity. Depending on the stellar compactness and on the rotation rate, the errors incurred by using Newtonian gravity instead of general relativity are found to be very sensitive to the adopted EoS, and amount to $\sim 20-40 \%$. These errors, however, might not be the dominant source of uncertainties. In particular, neutron superfluid vortices may not be arranged on a regular array parallel to the rotation axis, as assumed here. The dynamics of superfluid vortices and proton flux tubes remain highly uncertain, and warrant further studies. Considering the current upper limit $\tau_{\text{r}} < 30$ s \citep{dodson2007two}, we have found that the mutual friction parameter $\bar{\mathcal{B}}$ should be higher than $\sim 10^{-5}$ to explain Vela glitches. Since $\bar{\mathcal{B}}$ represents the average over the whole star, the mutual friction coupling $\mathcal{B}$ might be locally much stronger ($\mathcal{B}\sim 10^{-4}-10^{-3}$) as discussed for instance by \cite{sedrakian2005type} and \cite{haskell2014new}. In any case, since the actual value of $\tau_{\text{r}}$ is found to be much longer than the hydrodynamical time scale for current estimates of the mutual friction forces, the whole dynamical evolution of star during the spin up can be accurately computed by considering a sequence of stationary configurations only. We have also determined the amount of gravitational radiation emitted by the star during the spin up. For this purpose, we have studied the time variation of the mass quadrupole moment of the star resulting from changes in the fluid angular velocities. Using the quadrupole formula, we have numerically computed the characteristic amplitudes and frequencies associated with glitch events. Although the peak frequencies are found to lie in the sensitivity bands of current interferometers like Advanced LIGO, the corresponding amplitudes are too small for the gravitational waves to be detected. Their observations would require to improve the sensitivity by orders of magnitude. In particular, the characteristic amplitude for Vela is estimated to be at most $\sim10^{-32}$ for the (unrealistic) value $\bar{\mathcal{B}}=1/2$. If existing, the most promising sources would thus be pulsars rotating much more rapidly than Vela and undergoing high amplitude glitches. Although glitches are unlikely to be detected through gravitational waves, the Low Frequency Array (LOFAR) radio telescope \citep{stappers2011observing} and the future Square Kilometer Array (SKA) \citep{watts2015probing} will be able to observe the spin up with unprecedented accuracy. It would thus lead to much more stringent constraints on the characteristic time $\tau_{\text{r}}$ and thereby on the underlying glitch mechanism. This calls for more realistic models of glitching pulsars including the crust magnetoelasticity and superfluidity (whose formalism has been already developed, see, e.g. \cite{carterchachoua2006,carter2006relativistic}), and accounting for the local dynamics of quantized vortices.

1607.08213

1607

1607.04498_arXiv.txt

Two observations of V959 Mon, done using the {\sl Chandra} X-ray gratings during the late outburst phases (2012 September and December), offer extraordinary insight into the physics and chemistry of this \edit2{Galactic} ONe nova. The X-ray flux was 1.7 $\times$ 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$ and \edit2{8.6} $\times$ 10$^{-12}$ erg cm$^{-2}$ s$^{-1}$, respectively at the two epochs. The first result, coupled with electron density diagnostics and compared with published optical and ultraviolet observations, indicates that most likely in 2012 September the X-rays originated from a very small fraction of the ejecta, concentrated in very dense clumps. We obtained a fairly good fit to the September spectrum with a model of plasma in collisional ionization equilibrium (CIE) with two components; one at a temperature of 0.78 keV, associated with flat-topped and asymmetrical emission lines, blueshifted by $\simeq$710-930 km s$^{-1}$; the other at a temperature of 4.5 keV, mostly contributing to the high-energy continuum. However, we cannot rule out a range of plasma temperatures between these two extremes; we also modeled the spectrum as a static cooling flow, but the available models and the data quality are not adequate yet to differentiate between the two-component fit and a smoothly varying temperature structure. In December, the central white dwarf (WD) became visible in X-rays. We estimate an effective temperature of $\simeq$680,000 K, consistent with a WD mass $\geq$1.1 M$_\odot$. The WD flux is modulated with the orbital period, indicating high inclination, and two quasi-periodic modulations with hour timescales were also observed. No hot plasma component with temperature above 0.5 keV was observed in December, and the blue-shifted component cooled to kT$\simeq$0.45 keV. Additionally, new emission lines due to a much cooler plasma appeared, which were not observed two months earlier. We estimate abundances and yields of elements in the nova wind that cannot be measured in the optical spectra and confirm the high Ne abundance previously derived for this nova. We also find high abundance of Al, 230 times the solar value, consistently with the prediction that ONe novae contribute to at least 1/3rd of the Galactic yield of $^{26}$Al.

\label{sec:intro} Novae eruptions occur because of explosive CNO burning at the bottom of the envelope accreted onto WDs in binary systems. Dredged-up and mixed nuclei of C, N and O act as catalysts accelerating the burning and causing the envelope to expand as the electron degeneracy is lifted. Convection brings beta decaying nuclei to the upper layers, heating the envelope and allowing mass ejection. The outflow eventually occurs, mostly, or only, through a radiation driven wind \citep[see][and references therein]{Kelly13}. Each outburst ejects mass into the interstellar medium (ISM), from a few 10$^{-7}$ M$_{_\odot}$ to a few 10$^{-4}$ M$_{_\odot}$, depending on the mass accretion rate and time necessary for the pressure at the bottom of the accreted layer to exceed the gravitational pressure. This pressure is mostly an inverse function of the WD mass \citep[][]{Starrfield12, Wolf13}. Two types of novae are very important for Galactic chemistry. The first \edit2{type} are primordial novae \edit2{in} low metallicity binaries, that should yield Ti and nucleosynthetic end-points around Cu-Zn \citep{Jose07}. The \edit2{second type are the} more common oxygen-neon novae (ONe), attributed to WDs of oxygen and neon. \edit2{These novae} inject a significant amount of dredged up Ne and Mg and many intermediate elements including Al, Ar, Cl, and F into the ISM. Despite the rarity of ONe WDs, if \edit2{these WDs} accrete H-rich material, the outbursts are frequently repeated: there are probably $\approx$15 a year \edit2{such eruptions} in the Galaxy, assuming an observed nova rate of 50 novae of all types each year \citep[see][]{Shafter97, Shafter16}. ONe novae play a very interesting role in Galactic chemistry, and we have been able to cast light on their high metallicities with the work we present in this article. We analyse here in detail two {\sl Chandra} X-ray gratings' observations of the ONe nova V959 Mon, focusing also on the contribution of this nova to the Galactic chemistry. In Section 2 we review the known information about the nova. In Section 3 we describe the observations, and their analysis, including model fits for the ejecta and for the central WDs. Section 4 focuses on the analysis of the ejecta, especially on their chemistry as derived from the ``X-ray harder'' 2012 September spectrum, with a special attention for the problem of the yield of Al. Section 5 describes the timing analysis of the 2012 December observation. Finally, in Section 6 we discuss our findings and derive conclusions.

The {\sl Chandra} spectra of V959 Mon give insight into the nature of the central source and the physical conditions and chemical composition of the ejecta. The observation of the emission lines due to hot ejected plasma in the December 2012 spectrum prompted us to examine the September 2012 archival spectrum. At the earlier epoch, the X-ray high resolution spectrum was still much richer in emission lines due to high ionization potential transitions observed with better S/N than later in December, without superimposition with the WD atmospheric continuum. Both {\sl Chandra} spectra highlight the potential of X-ray high resolution spectra to derive important parameters of the nova physics. With some assumptions, namely that the optically measured abundances of He and C were the same in the plasma observed in the X-rays, and postulating that the O and N abundances were approximately constant in the two epochs of {\sl Chandra} observations, we were able to estimate absolute abundances of the special elements of which a nova on an ONe WD enriches the interstellar medium. The very large Al abundance of the two-component fit, confirmed by the yield calculation made by evaluating the flux in the H-like and He-like emission lines, implies that ONe novae significantly contribute to the Galactic Al abundance, assuming, of course, that the abundances derived in the small X-ray emitting mass are representative of all the ejecta (there is no reason to think they may be very different). Our results are consistent with the expectations that novae on ONe WDs contribute {\it at least} 30\% of the total galactic yield of Al$^{26}$ in the ISM. Future observations of the spectral region of the Al lines in these novae will clarify the origin of the Galactic Al, a problem that bears implications even in order to evaluate the SN II rate. In the future {\it Athena} will allow a very accurate estimation of line fluxes in this region for a number of ONe novae. The lines' ratios in the September spectrum, based on the Mg XI He-like triplet, suggest that the X-rays were emitted by concentration of material of very small total mass (less than 10$^{-10}$ M$_\odot$), concentrated in very small clumps. The spectral shape of the lines suggest a spread in space rather than collimation. The remnant shell of the recurrent nova T Pyx from outbursts previous to 2011 presents a very knotty structure that was flash-ionized again in 2011 \citep[][]{Shara15, Shara97}. \citet[][]{Toraskar13} attribute these knots to the collision of new ejecta with the swept-up, cold, dense shell from the older outbursts, which would drive Richtmyer-Meshkov instabilities in it. An alternative explanation may be that a fraction of the ejecta, moving more slowly than the rest, is {\it from the beginning of the outburst} ejected in dense X-ray emitting clumps, which later cool and constitute the clumpy material that was detected by \citet[][]{Shara97} with HST. This is only a working hypothesis, to be explored in further research, but the concept of small clumps in the ejecta clearly brings to one's mind the image of the old shell of T Pyx. \added{\citet[][]{Williams13} attributed hard X-rays to dense blobs of gas colliding with each other while expanding. In this author's scenario, the WD ejecta are all in such blobs or clumps, and they cools very quickly. The more homogeneous circumbinary gas is due instead to mass loss from the secondary. The small mass of X-ray emitting material at a given time would thus depend on the mass loss rate from the WD and the cooling time of the globules. However, for an X-ray emitting mass of the order of 10$^{-8}$ M$_\odot$ and a mass loss rate of the order of 10$^{-6}$ M$_\odot$ year$^{-1}$, a value that seem reasonable for this nova, the cooling/expanding time must have been of the order of few days, much longer than the minute time scale predicted by \citet[][]{Williams13}. This implies that \citet[][]{Williams13} model may be viable after carefully revisiting and constraining the physical parameters. Incidentally, we note that, if only a portion of the material is ejected by the WD and a significant \edit2{part of it} is instead ejected by the secondary, this makes difficult to estimate the yield of chemical elements from observations at any wavelengths, so it would be crucial to devise ways to estimate the amount of ejected material from the secondary. The comparison between abundances of the same elements obtained in the X-ray range and in the optical range may be the best way to obtain this information. } In the September spectrum the emission clearly originated in at least two regions with distinct plasma temperatures, 4.5 and 0.78 keV, both in CIE. However, we could not rule out a smooth variation of temperature between the two values, or a lower minimum plasma temperature. In our best fit model, 57\% of the unabsorbed flux (but only 18\% of the measured flux) was emitted in blueshifted material that produces slightly asymmetric, flat topped emission lines, probably indicating a large spread in spatial directions and radius-dependent absorption. 43\% of the unabsorbed flux (and 82\% of the measured, absorbed flux) originated in a much hotter region, contributing to almost a third of the flux in the H-like emission lines. The discrepancy in the evaluation of the Ne abundance obtained with two methods, the global fit and the measurement of the flux in the lines (assuming only the N(H) and emission measure values estimated with the global fit), may imply that the temperature structure is more complex and/or more smoothly varying than our two zones model. In the future, we should study a cooling flow type of model, with gradually changing emissivity, including line broadening and other elements appropriate for the nova physics. In addition to working on the models, we also hope to obtain better quality data for future ONe nova outbursts. Only one component of plasma in CIE with kT around half a keV was found for the December spectrum, while new emission lines, that most likely are not associated with the central source, had emerged, probably indicating a much cooler plasma (kT$\simeq$0.1 keV or less), at the lower limit of the temperature of the CIE plasma models in XSPEC or other analysis packages. The main goal we had in mind in proposing the December 2012 LETG observation was to measure the WD temperature and chemical composition. It is quite remarkable that the hot WD, emitting supersoft X-rays, emerged and became detectable even at high inclination and with a considerable amount of interstellar absorption towards the nova. The modulation of the central supersoft X-ray source confirms the high inclination at which the system is viewed. This bears implications \added{for those cases in which the central supersoft X-ray source (SSS) is never detected. \citet[][]{Ness13} have suggested that the detection of the SSS may be impossible in nova system viewed at high inclination, while in these novae the emission line spectrum of the ejecta is better observed and more easily measurable. As a consequence, unless the \edit2{central source} detection is aided by a phenomenon of Thomson scattering \citep[see][]{Orio13}, we would not be able to draw any conclusion from ``missing'' SSS (e.g. whether it has not turned on yet, or it has already turned off) because in most systems at high inclination, because the SSS will never appear. Not knowing the range of inclinations, it would be impossible to obtain statistics from the study of extragalactic novae like the one of \citet[][]{Henze14}. The example of V959 Mon, where the SSS eventually was measurable even at high inclination and without clear evidence of Thomson scattering, indicates that a missing SSS is instead more likely to be due to slow evolution in an ejected shell that is optically thick to supersoft X-rays, rather than to the inclination. There are several examples in the literature of novae that in which an emission lines X-ray spectrum later became an SSS X-ray spectrum when the SSS really emerged, like N LMC 2009,\citep[][]{Orio13b}, so we suggest that the inclination is unlikely to be the determining factor in the SSS detection. } We find that the WD temperature was consistent with a WD of at least 1.1 M$_\odot$, comparing our best-fit temperature T$_{\mathrm eff}$=680,000 K with the current theoretical models \citep{Wolf13}. However, the increase in supersoft luminosity in the following two weeks is consistent with a further increase in T$_{\rm eff}$ \citep[][]{Page13}, probably indicating an even higher mass. In fact, a blueshift of 2260 km s$^{-1}$ in the emission lines indicates residual mass loss, \edit2{implying that} we did not observe the WD after all mass loss had ceased; the WD had not returned to its pre-outburst dimensions yet. However, the detected absorption features are indeed in good agreement with those in the model for a static atmosphere on a WD of mass of at least 1.1 M$_\odot$ (that is, a H-burning WD with peak temperature of, or above, 680,000 K). The detection of two periods of 55 and 102 minutes respectively in the LETG light curve in 2012 December is consistent with several periodicities measured in the supersoft X-ray source after a nova outburst \citep[][]{Orio12, Ness12}. Our final conclusion is that the two {\sl Chandra} grating spectra of V959 Mon highlight the potential of X-ray grating observations to gain insight into the nova physics, \edit2{by} deriving the effective temperature and relative mass of the central WD, \edit2{by} better understanding the dynamics of the mass ejection, and \edit2{by} estimating the chemical yields in the interstellar medium. \begin{table}[t] \centering \caption{Chemical abundances, by number, derived for ONe novae from optical and ultraviolet spectra, compared with the results for V959 Mon. } \label{table:abundances-other} \begin{tabular}{lllll} \hline Nova & Ne/Ne$_\odot$&Mg/Mg$_\odot$& Al/Al$_\odot$ & Reference \\ \hline V382 Vel & 17$\pm$3 & 2.6$\pm$0.1 & 21$\pm$2 & \citet{Shore03} \\ QU Vul & $21.7\pm1.7$ & $10\pm5.1$ & 53.3$\pm$14.7 & \citet{2002S} \\ LMC 1990 & 62$\pm44$ & 16$\pm$7.5 & 257$\pm98$ & \citet{1999V} \\ V693 Cra & 247$\pm144$ & 7.9$\pm7$ & 60$\pm59$ & \citet{1997V} \\ V838 Her & $52.5\pm2.3$ & $1.4\pm0.8$ & 29$\pm$21 & \citet{2007S} \\ V1974 Cyg & 41.5$\pm$17 & 4.6$\pm$3 & ... & \citet{2005V} \\ V959 Mon & 1660$^{+40}_{-160}$ & 230$_{-20}^{+40}$ & 230$^{+70}_{-70}$ & 2012 September (two-component fit) \\ & 587$\pm$169 & 308$\pm$146 & 229$\pm$66 & 2012 September (flux method) \\ & 190$^{+80}_{-120}$ & 220$^{+70}_{-60}$ & 770$^{+380}_{-380}$ & 2012 December \\ & 95 & ... & ... & 2013 March-April \citet{Tarasova14} \\ \hline \end{tabular} \end{table}

1607.04498

1607

1607.07791_arXiv.txt

{We obtain a family of regular static, spherically symmetric solutions in Einstein--Cartan theory with an electromagnetic field and a nonminimally coupled scalar field with the correct sign of kinetic energy density. At different values of its parameters, the solution, being \asflat\ at large values of the radial coordinate, describes (i) twice \asflat\ symmetric \whs, (ii) asymmetric \whs\ with an AdS asymptotic at the ``far end'', (iii) regular black holes with an extremal horizon or two simple horizons, and (iv) black universes with a de Sitter asymptotic at the ``far end''. As in other black universe models, it is a black hole as seen by a distant observer, but beyond its horizon there is a nonsingular expanding universe. In all these cases, both the metric and the torsion are regular in the whole space. } Keywords: {Einstein--Cartan theory, scalar field, nonminimal coupling, wormholes, regular black holes, black universes} PACS number: 04.20.-q, 04.20.Jb, 04.40.-b, 98.80.Jk \bigskip ] % \email 1 {kb20@yandex.ru} \email 2 {agal17@mail.ru}

The origin of the presently observed accelerated expansion of the Universe has become one of the most important problems in modern cosmology and even in theoretical physics as a whole. Among different theoretical models trying to explain it (see, e.g., the recent reviews [1--4] and references therein), two main trends can be distinguished: (1) introduction of a new hypothetic form of matter with large negative pressure, called dark energy (DE), in the framework of general relativity (GR) (the cosmological constant, various kinds of quintessence, phantom matter etc.), and (2) different suggestions in alternative theories of gravity, such as $f(R)$ theories, multidimensional theories and theories involving non-Riemannian geometries, such as the Riemann-Cartan geometry with torsion. The simplest theory of this kind is the Einstein-Cartan theory (ECT) [5--8], also leading to models of accelerated expansion [9--11]. The ECT can be considered as a degenerate version [12--14] of the Poincar\'e gauge theory of gravity (PGTG), in which the gravitational Lagrangian contains invariants quadratic in the curvature and torsion tensors. Unlike that, in the ECT the torsion is not dynamic since its gravitational action reduces to the curvature scalar of Riemann--Cartan space--time, directly generalizing the action of GR. It is nevertheless a viable theory of gravity: its observational predictions agree with the classical tests of GR, but it substantially differs from GR at very high densities of matter \cite{HO07,Trautman,BH11}. Theories with torsion also attract attention since torsion naturally arises in many approaches such as supergravity [17--19] and superstring [20--22] theories. One of the simplest extensions of the ECT is $f(R)$ gravity with torsion \cite{CCSV07,CCSV08}, and, as shown in \cite{CCSV08}, torsion can play the role of DE and cause an accelerated expansion of the Universe. Moreover, the existence of bouncing cosmologies in the ECT \cite{G11} shows that torsion can replace ``exotic'' sources, violating the weak and null energy conditions (WEC and NEC). As is well known, such a violation is in GR a necessary condition for the existence of traversable wormholes \cite{hoh-vis}. Wormholes are a subject of particular interest as possible time machines or shortcuts between different universes or distant parts of the same universe, for reviews see \cite{vis-book, lobo-rev, BR-book} and references therein. Quite a number of \wh\ solutions are known, see, e.g, [29--32] for solutions with minimally coupled scalar fields, \cite{br73, BVis99} for solutions with conformal coupling and \cite{BVis00, br-96, br-gr02} for other couplings. In agreement with the general results \cite{hoh-vis}, minimally coupled scalar fields supporting \whs\ have to be phantom (i.e., have a wrong sign of kinetic energy). In the case of a nonminimal coupling, there are special \wh\ solutions with normal fields, but in all such cases there are always regions where the effective gravitational constant becomes negative, that is, the gravitational field itself becomes a phantom \cite{br-JMP,br-star07}. Moreover, all such configurations, whose existence is connected with the phenomenon of conformal continuation \cite{br-Pol,br-JMP}, turn out to be unstable under radial perturbations \cite{br-gr01,br-gr02,br-gr04}. Some extensions of GR predict the existence of \whs\ without exotic matter, in particular, brane world models \cite{kb-kim1,kb-kim2}, Einstein-Gauss-Bonnet gravity \cite{EGB} and other high-order theories \cite{HOG-13}, the Horndeski theory \cite{Sush-Horn} and others. From the properties of the ECT it is also natural to expect that in this theory wormholes can exist without exotic matter or at least without manifestly phantom fields with a wrong sign of kinetic energy. And indeed, a family of exact \ssph\ wormhole solutions in the ECT was recently found \cite{br-gal15} with a pair of canonical scalar fields as sources of gravity. One of these fields was nonminimally coupled to Riemann-Cartan curvature and provided the effect of torsion on the space-time metric. A shortcoming of these solutions was an infinite value of the torsion scalar at the \wh\ throat. Other kinds of configurations in GR whose existence is connected with NEC and WEC violation are regular black holes which, instead of a singularity at $r=0$, contain, at the ``far end'', flat or (anti-) de Sitter asymptotic regions \cite{BF06, bu-Dehn, BBS12}. Among them of particular interest are the so-called black universes. By definition, a black universe is a nonsingular black hole in which, beyond the event horizon, there is an expanding universe. This class of models provides avoidance of singularities in both black holes and cosmology and combines the properties of a wormhole (absence of a center, a regular minimum of the area function in the case of spherical symmetry) and a black hole (a horizon separating static and cosmological regions of space-time). Moreover, the Kantowski-Sachs cosmology in the T region can be asymptotically isotropic and approach a de Sitter mode of expansion, which makes such models potentially viable for a description of an inflationary Universe or the present accelerated expansion. A number of such solutions of GR have been obtained with different kinds of phantom scalar fields as sources, with and without \elmag\ fields \cite{BF06, BBS12, BDon11, BKor15}. In the present study, we again seek \ssph\ solutions in the ECT but now with a single nonminimally coupled scalar field (being a source of torsion) and an \elmag\ field. Our purposes are (1) to obtain both \wh\ and regular black hole solutions with a normal scalar field, (2) to include electric or magnetic fields into consideration, and (3) to avoid a singular behavior of torsion in the whole space-time. The paper is organized as follows. In Section 2 we present the ECT equations both in the general case and for static, spherically symmetric configurations involving an electromagnetic field and a scalar field nonminimally coupled to space-time curvature. Section 3 is devoted to finding and analyzing the properties of a family of exact solutions, and Section 4 is a discussion.

We have found a family of exact static, spherically symmetric solutions in the Einstein-Cartan theory (ECT) of gravity, with sources in the form of a nonminimally coupled non-phantom scalar field and an electromagnetic field. From this whole family we have selected asymptotically flat solutions with a nonnegative \Scw\ mass. The remaining subfamily depends on four constants: the nonminimal coupling coefficient $\xi > 1/2$, an arbitrary length scale $b > 0$ and two significant integration constants: the dimensionless electromagnetic charge $Q$ and the ``asymmetry factor'' $\alpha$. With different values of these parameters, the solution describes (i) twice \asflat\ (or M-M) symmetric \whs, (ii) asymmetric M-AdS \whs\ with zero or nonzero masses, (iii) regular M-AdS black holes with an extremal horizon or two simple horizons, and (iv) M-AdS black universes. It is important that, in all these solutions with a nonsingular metric, the torsion scalar also remains finite in the whole space, unlike the recently obtained \wh\ solutions in the ECT with two scalars \cite{br-gal15}. In addition, there are a number of singular solutions that remained beyond the scope of this study, for example, those with $\alpha$ outside the range \rf{25a}. It seemed that if $\alpha$ takes a marginal value such that the spherical radius tends to a finite constant $r=r_0$ at the ``far end'', $x\to -\infty$. It turns out, however, that such solutions possess a repulsive singularity at $r=r_0$, and, depending on $Q$, this singularity can be naked or hidden beyond horizons from the viewpoint of a distant observer. It is clear that in our solutions the existence of a minimum of $r(x)$ (a throat) is provided by WEC and NEC violation by the effective SET (\ref {10}) of the scalar--torsion field. A question of interest is whether or not the purely scalar SET respects the energy conditions, so that their violation might be completely ascribed to the contribution of torsion. We notice, however, that this question cannot be asked correctly because only the effective SET $T^{i \rm (eff)}_k [\phi]$ satisfies the conservation law \rf{11}, and the purely scalar contribution cannot be unambiguously separated from that of torsion. Nevertheless, let us try to calculate the quantity $\rho + p_r \equiv T^t_t - T^x_x$ (whose negative value indicates NEC violation) for the purely scalar contribution in the SET \rf{7}, excluding there the terms containing $S$ and its derivatives, which would be a correct expression for the scalar field SET in the absence of torsion. We obtain \beq \rho[\phi] + p_{r} [\phi] = \frac{A(x)\bigl [1 -2\xi+ (1 + 4\xi )x^2 \bigr]}{\kappa \xi b^2 (x^2 + 1)^2}. \eeq This expression is negative at small $x$ (near the throat) since $\xi > 1/2$, but becomes positive far from it. This resembles the concept of a ``trapped ghost'' \cite{BDon11, BS10}, but, in contrast to these papers, here the kinetic term of the scalar field does not change its sign. It is of interest that the effective density $\rho^{\rm (eff)}[\phi] = T^{t \rm (eff)}_t$ can be positive near the throat in our solution, so that NEC violation is caused by a larger negative value of the effective radial pressure $p_r^{\rm (eff)}[\phi] = -T^{x \rm (eff)}_x$ Thus, for symmetric configurations $(\alpha = 0,\ K = 0)$ we obtain at $x=0$ \beq \rho^{\rm (eff)}[\phi] = \frac{1}{\kappa b^2} \Bigr[2 + 2(p^2 -1)B_0 + Q\Bigr]. \eeq For $p = 0.5$ \, $(a = Q = \pi/4)$ we have \beq \rho^{\rm (eff)}[\phi] = \frac{1}{\kappa b^2 } \Bigr[6 + 4\ln 2 - \frac{3\pi }{4}\Bigr ] > 0 . \eeq We conclude that the ECT, like some other extensions of GR, provides the existence of regular configurations without a center (\whs, \bus, and regular black holes with two asymptotic regions) with normal (non-phantom) fields. It means that torsion here replaces (or plays the part of) exotic matter while forming such regular objects. A feature of interest in the present solution is the necessity of the \elmag\ field for obtaining asymptotic flatness, but it is evidently a property of the present family of solutions rather than a general property of the theory. A challenging problem is that of stability of these and other solutions of the ECT, and we hope to deal with it in our further studies. \subsection*

1607.07791

1607

1607.00387_arXiv.txt

{Spatially resolved polarized (sub-)mm emission has been observed for example in the protoplanetary disk around HL Tau. Magnetically aligned grains are commonly interpreted as the source of polarization. However, self-scattering by large dust grains with a high enough albedo is another polarization mechanism, becoming a compelling method independent of the spectral index to constrain the dust grain size in protoplanetary disks.} {We study the dust polarization at mm wavelength in the dust trapping scenario proposed for transition disks, when a giant planet opens a gap in the disk. We investigate the characteristic polarization patterns and their dependence on disk inclination, dust size evolution, planet position, and observing wavelength.} {We combine two-dimensional hydrodynamical simulations of planet-disk interactions with self-consistent dust growth models. These size-dependent dust density distributions are used for follow-up three-dimensional radiative transfer calculations to predict the polarization degree at ALMA bands due to scattered thermal emission.} {Dust self-scattering has been proven to be a viable mechanism for producing polarized mm-wave radiation. We find that the polarization pattern of a disk with a planetary gap after 1\,Myr of dust evolution shows a distinctive three ring structure. Two narrow inner rings are located at the planet gap edges. A third wider ring of polarization is situated in the outer disk beyond 100\,au. For increasing observing wavelengths all three rings slightly change their position, where the innermost and outermost rings move inward. This distance is detectable comparing the results at ALMA bands 3, 6 and 7. Within the highest polarized intensity regions the polarization vectors are oriented in the azimuthal direction. For an inclined disk there is an interplay between polarization originating from a flux gradient and inclination-induced quadrupole polarization. For intermediate inclined transition disks the polarization degree is as high as $\sim 2\%$ at $\lambda=3.1\,$mm (band 3), which is well above the detection limit of future ALMA observations.} {}

\label{sec:introduction} To investigate dust coagulation and ongoing planet formation processes observational constraints on the size of grains embedded in a protoplanetary disk are crucial. Observations at millimeter (mm) wavelengths have probed large (mm sized) grains in the disk midplane, through calculation of a low spectral index of the dust opacity (e.g. \citealt{beckwith1991,testi2001,rodmann2006,ricci2010b,ricci2010,guilloteau2011}). For a continuous disk with a monotonically decreasing radial gas pressure, mm dust particles in the outer disk can experience an excessive radial drift towards the star, in contradiction to observations that revealed the existence of mm sized particles in the outer disk regions \citep{wilner2005,andrews2005,ricci2010,ricci2011}. A particle trap caused by a pressure bump might be a solution to prevent this drift problem and to allow the grains to grow efficiently \citep{whipple1972,klahr1997,fromang2005,johansen2009,pinilla2012,zhu2012}. This bump can for example result from the presence of a massive planet carving a gap in the gas density. Hence, so-called transition disks characterized by large radial gaps are excellent targets to study the impact of planet formation on the disk structure. To explain pronounced dust rings in protoplanetary disks there are certainly also other mechanisms, such as the generation of gaps at the dead-zone edges in magnetized disks without any planet \citep{flock2015}. Particle trapping can be probed by measuring low spectral index values inside the trap, however, small opacity index values could be also explained by optically thick emission from compact regions \citep{ricci2012}. Among other uncertainties, such as the composition and porosity of dust aggregates \citep{henning1996,kataoka2014}, this shows that constraining the grain size only with opacity index measurements is still problematic.\\ Recently, \citet{kataoka2015} introduced an alternative, independent method to constrain the grain size distribution in protoplanetary disks based on dust polarization at mm wavelengths. The classical picture from the optical and near-infrared (NIR), where stellar photons are scattered by small dust in the disk surface layers, is transferred to dust self-scattering\footnote{Self-scattering means that the source of incident radiation is the thermal emission of the dust itself, which is then scattered off large dust grains resulting in polarized mm disk emission.} in the mm wavelength regime. As mentioned above, it is known that dust grains in protoplanetary disks can grow to sizes comparable to mm wavelengths, meaning that they are expected to have a large albedo and thus, can produce scattered light. When the radiation field is anisotropic, the continuum emission is expected to be partially polarized due to self-scattering of dust thermal emission. So far, polarized (sub-)mm emission has been observed in the disks around a few young stellar objects, e.g. IRAS 16293-2422B \citep{tamura1995,rao2014}, HL Tau \citep{tamura1995,stephens2014}, DG Tau \citep{tamura1999, hughes2013}, MWC 480 \citep{hughes2013}, and L1527 \citep{seguracox2015}. Commonly, magnetically aligned grains (see e.g. \citealt{lazarian2007}) are assumed as the source of polarization, but dust self-scattering is another important mechanism that needs to be considered. This idea has been explicitly applied to the protoplanetary disk around HL Tau by \citet{kataoka2016} and \citet{yang2016}, who successfully reproduced the polarization signatures observed by means of scattered mm radiation.\\ In this paper we study the dust trapping scenario when a massive planet is embedded in the disk, and investigate where polarization due to scattering can be detected. We combine 2D hydrodynamical simulations of planet-disk interactions with self-consistent dust growth models (cf. \citealt{pinilla2012b,dejuanovelar2013}). These results are used to perform 3D radiative transfer calculations in order to predict the polarization at ALMA wavelengths due to scattered thermal emission. In difference to \citet{kataoka2015}, we consider the simulated spatial dust density distribution for each grain size from the dust evolution model instead of a parametrized dust density and a simplified power law grain size distribution. We compare our results for different dust evolution timescales and analyze the dependence of the polarization degree on the disk inclination, dust composition, planet's position and observing wavelength. Moreover, it is discussed whether the polarization is detectable with future ALMA observations.\\ This paper is organized as follows. In Sect. \ref{sec:methods} we describe the numerical methods to obtain a transition disk model and the radiative transfer calculations. Section \ref{sec:results} presents our results of the dust growth modeling and the simulated polarization maps. Furthermore, we discuss our findings in terms of disk inclination, dust evolution timescales, dust composition, planet's position and observing wavelength. Finally, our results and the conclusions of this work are summarized in Sect. \ref{sec:conclusions}.

\label{sec:conclusions} We present a new technique to investigate the dust trapping scenario in transition disks hosting a pressure bump by means of mm-wave polarization of scattered thermal emission. More precisely, the continuum emission is polarized due to dust self-scattering of an anisotropic radiation field induced by disk inclination. For the mm-wave polarization to work the scattering grains must have grown to a maximum size of a few hundred microns. Additionally, further populations of either very small micron-sized or mm-/cm-sized grains need to be present in order to account for the large portion of unpolarized continuum emission. Our model predictions at different ALMA bands (bands 3, 6 and 7) are based on self-consistent dust growth models and radiative transfer calculations, which allows us to estimate the polarization degree in disks hosting a massive planet. The planet is considered to be located at 20\,au and 60\,au, respectively. Compared to previous studies a dust size distribution in radial direction is included in our model. Measuring the dust polarization degree is a direct and an unambiguous method to probe the location of large particles (a few hundred micron to mm sized particles, depending on the observing wavelength) when particle trapping occurs in the disk. We emphasize that the polarization technique presented can be applied to any disk with ring-like dust structures segregating the grain sizes, and is not limited to the gap opening scenario by planets. In this paper we focus on transition disks, because they are excellent candidates where planet formation may be ongoing. Thus, polarization observations are especially a unique tool to investigate the dust distribution when particle trapping is triggered by a planet embedded in the disk. The main findings of this paper are the following: \begin{enumerate} \item The polarization pattern of a disk hosting a planetary gap after 1 Myr of dust evolution shows a characteristic three ring structure, where the two inner, narrow rings are located just at the gap edges. Additionally, there is a third polarization ring in the outer disk beyond 100\,au with a larger radial extension. Detecting such a distinctive ring structure with polarization observations may represent the radial size distribution of dust grains and hint to regions with a specific grain size corresponding to the wavelength.\\ \item For an inclined disk there is an interplay between polarization originating from a flux gradient and from an inclination-induced anisotropy. The fraction of scattered radiation polarized due to disk inclination increases with the inclination angle. For intermediate inclined transition disks the polarization degree at $\lambda=3.1\,$mm is as high as $~2\,\%$, which is well above the detection limit of future ALMA polarization observations. A spatial resolution as high as 0\farcs2 is required to certainly resolve the ring structure in polarized intensity.\\ \item The local degree of polarization is very sensitive to the maximum grain size at a certain location in the disk. Large cm and mm grains do not contribute to the polarization, which leads to a significant local reduction of the polarization degree, e.g. at the pressure bump region in our transition disk models. Hence, if the maximum grain size is larger than a few hundred microns, the polarization is not detectable, even if numerous small grains are present and contribute.\\ \item For the face-on disks the polarization vectors are in azimuthal direction within the highest polarized intensity regions. For inclined disks the majority of the polarization vectors in the two inner rings are orientated along the minor axis. In the outer ring the gradient-induced polarization and, therefore, the azimuthal vector orientation still dominates.\\ \item With increasing observing wavelengths the innermost and outermost polarization rings move inwards by a detectable distance, while the middle ring slightly moves radially outside. Hence, the positions of the two rings at the gap edges are shifted in opposite directions. For the outermost ring the moving distance is $\sim 20\,$au comparing the results at 0.87\,mm and 3.1\,mm. Furthermore, without any giant planet embedded in the disk, this third ring reflects the only polarized region, even though the polarization degree is quite low.\\ \item We find that the dust composition has no effect on the overall polarization ring structure. The presence of even a very small fraction of silicate in the dust mixture causes the local polarization degree and net polarization to be very similar to the case of pure silicate species. Silicate grains dominate the refractive index of the dust mixture, and the fractional abundances of carbonaceous material and water ice hardly affect the polarization pattern. A significant change of the local polarization degree with dust species can be led back to either optical depths effects or different flux gradients. \end{enumerate} We showed that polarization due to dust self-scattering is a powerful tool in order to constrain the grain size in protoplanetary disks independent of the spectral index. Nevertheless, there are also other polarization mechanisms such as alignment of dust grains along the magnetic field vectors. A keystone to investigate in future work is how to distinguish between these two mechanisms. Further multi-wave and spatially resolved polarization observations of protoplanetary disks are necessary. The wavelength and grain size dependence of polarization are also the keypoints from the modeling side. Due to the self-scattering mechanism we expect the polarization degree to significantly change with the observing wavelength, while it is thought to be approximately constant due to magnetically aligned grains in a toroidal magnetic field (\citealt{cho2007}).

1607.00387

1607

1607.00949_arXiv.txt

We study sets of oscillators that have high quantum occupancy and that interact by exchanging quanta. It is shown by analytical arguments and numerical simulation that such systems obey classical equations of motion only on time scales of order their relaxation time $\tau$ and not longer than that. The results are relevant to the cosmology of axions and axion-like particles.

1607.00949

1607

1607.07875_arXiv.txt

{ We derive the luminosity function \lf\ and redshift distribution \psiz\ of short Gamma Ray Bursts (SGRBs) using (i) all the available observer--frame constraints (i.e. peak flux, fluence, peak energy and duration distributions) of the large population of \fe\ SGRBs and (ii) the rest--frame properties of a complete sample of SGRBs detected by \sw. We show that a steep $\phi(L)\propto L^{-\alpha}$ with $\alpha\ge2.0$ is excluded if the full set of constraints is considered. We implement a Monte Carlo Markov Chain method to derive the \lf\ and \psiz\ functions assuming intrinsic \yone\ and \ama\ correlations to hold or, alternatively, that the distributions of intrinsic peak energy, luminosity and duration are independent. To make our results independent from assumptions on the progenitor (NS--NS binary mergers or other channels) and from uncertainties on the star formation history, we assume a parametric form for the redshift distribution of the population of SGRBs. We find that a relatively flat luminosity function with slope $\sim 0.5$ below a characteristic break luminosity $\sim 3 \times10^{52}$ erg s$^{-1}$ and a redshift distribution of SGRBs peaking at $z\sim1.5-2$ satisfy all our constraints. These results hold also if no \yone\ and \ama\ correlations are assumed, and they do not depend on the choice of the minimum luminosity of the SGRB population. We estimate that, within $\sim$200 Mpc (i.e. the design aLIGO range for the detection of gravitational waves produced by NS--NS merger events), there should be 0.007--0.03 SGRBs yr$^{-1}$ detectable as $\gamma$--ray events. Assuming current estimates of NS--NS merger rates and that all NS--NS mergers lead to a SGRB event, we derive a conservative estimate of the average opening angle of SGRBs $\langle \theta_{\rm jet}\rangle\sim3^\circ$--$6^\circ$. The luminosity function implies a prompt emission average luminosity $\left\langle L \right\rangle \sim 1.5 \times 10^{52}\,\rm{erg\,s^{-1}}$, higher by nearly two orders of magnitude compared to previous findings in the literature, which greatly enhances the chance of observing SGRB ``orphan'' afterglows. Efforts should go in the direction of finding and identifying such orphan afterglows as counterparts of GW events.}

The population of short Gamma Ray Bursts (SGRBs) is still poorly understood due to the relatively few events with measured redshift \citep[see e.g.][for recent reviews]{2014ARA&A..52...43B,2015JHEAp...7...73D}. Available information is rather sparse, but the low density of the close circumburst medium \citep{2013ApJ...776...18F,Fong:2015fp}, the variety of galaxy morphologies \citep[e.g.][]{2015JHEAp...7...73D}, the lack of any associated supernova in the nearby SGRBs and the possible recent detection of a ``kilonova'' \citep{1989Natur.340..126E,1998ApJ...507L..59L,Yang:2015lr,Yang:2015hl,Jin:2016rr,Jin:2015cr} signature \citep{2013ApJ...774L..23B,2013Natur.500..547T}, all hint to an origin from the merger of two compact objects (e.g. double neutron stars) rather than from a single massive star collapse. However, the prompt $\gamma$--ray emission properties of SGRBs \citep{2009A&A...496..585G,2015JHEAp...7...81G} and the sustained long lasting X--ray emission (despite not ubiquitous in short GRBs - \citealt{Sakamoto:2009lr}) and flaring activity suggest that the central engine and radiation mechanisms are similar to long GRBs. Despite still based on a couple of breaks in the optical light curves, it seems that also SGRBs have a jet: current measures of $\theta_{\rm jet}$ are between 3$^\circ$ and 15$^\circ$ while lower limits seem to suggest a wider distribution \citep[e.g.][]{2014ARA&A..52...43B,2015ApJ...815..102F}. Recently, it has been argued that the customary dividing line at $T_{90}=2\,\rm{s}$ between short and long GRBs provides a correct classification for \textit{Fermi} and \textit{CGRO} GRBs, but it is somewhat long for \textit{Swift} bursts \citep{2013ApJ...764..179B}. A renewed interest in the population of SGRBs is following the recent opening of the gravitational wave (GW) ``window'' by the LIGO--Virgo discovery of GW150914 \citep{Abbott:2016lr} and by the most recent announcement of another event, GW151226, detected within the first data acquisition run \citep{2016arXiv160604856T,Abbott:2016lr}. Despite no electromagnetic (EM) counterpart was identified within the large localisation region of these event, there are encouraging prospects for forthcoming GW discoveries to have an EM--GW association, thanks to the aLIGO--Virgo synergy and world wide efforts for ground and space based follow up observations. If the progenitors are compact object binary (NS--NS or NS--BH - \citealt[e.g.][]{Giacomazzo:2013fk}) mergers, SGRBs are one of the most promising electromagnetic counterparts of GW events detectable by the advanced interferometers. Other EM counterparts are expected in the optical \citep{Metzger:2012fj}, X-ray \citep{Siegel:2016qy, Siegel:2016lq} and radio bands \citep{Hotokezaka:2016uq}. The rate of association of GW events with SGRBs is mainly determined by the rate of SGRBs within the relatively small horizon set by the sensitivity of the updated interferometers aLIGO and Advanced Virgo \citep{2016LRR....19....1A}. However, current estimates of local SGRB rates range from 0.1--0.6 Gpc$^{-3}$ yr$^{-1}$ (e.g. Guetta \& Piran 2005; 2006) to 1--10 Gpc$^{-3}$ yr$^{-1}$ \citep[WP15]{2006A&A...453..823G,2009A&A...498..329G,2012MNRAS.425.2668C,2014MNRAS.437..649S} to even larger values like 40-240 Gpc$^{-3}$ yr$^{-1}$ \citep{2006ApJ...650..281N,2006A&A...453..823G}\footnote{All these rates are not corrected for the collimation angle, i.e. they represent the fraction of bursts whose jets are pointed towards the Earth, which can be detected as $\gamma$--ray prompt GRBs.}. Such rate estimates mainly depend on the luminosity function \lf\ and redshift distribution \psiz\ of SGRBs. These functions are usually derived by fitting the peak flux distribution of SGRBs detected by BATSE \citep{2005A&A...435..421G,2006A&A...453..823G,2006ApJ...650..281N,2006ApJ...643L..91H,2008MNRAS.388L...6S}. Due to the degeneracy in the parameter space (when both \lf\ and \psiz\ are parametric functions), the redshift distribution was compared with that of the few SGRBs with measured $z$. % The luminosity function \lf\ has been typically modelled as a single or broken power law, and in most cases it was found to be similar to that of long GRBs \citep[i.e.\ proportional to $L^{-1}$ and $L^{-2}$ below and above a characteristic break $\sim 10^{51-52}$ erg s$^{-1}$ -][D14 hereafter]{2006A&A...453..823G,2008MNRAS.388L...6S,2011ApJ...727..109V,2014MNRAS.442.2342D} or even steeper \citep[$L^{-2}$ and $L^{-3}$ -][WP15 hereafter]{2015MNRAS.448.3026W}. Aside from the mainstream, \cite{2015MNRAS.451..126S} modelled all the distributions with lognormal functions. The redshift distribution \psiz\ (the number of SGRBs per comoving unit volume and time at redshift $z$) has been always assumed to follow the cosmic star formation rate with a delay which is due to the time necessary for the progenitor binary system to merge. With this assumption, various authors derived the delay time $\tau$ distribution, which could be a single power law $P(\tau)\propto\tau^{-\delta}$ (e.g. with $\delta=1-2$, \citeauthor{2005A&A...435..421G} \citeyear{2005A&A...435..421G}, \citeyear{2006A&A...453..823G}; D14; WP15) with a minimum delay time $\tau_{\rm min}=10-20$ Myr, or a peaked (lognormal) distribution with a considerably large delay (e.g. 2--4 Gyr, \citeauthor{2005AAS...20715803N} \citeyear{2005AAS...20715803N}; WP15). Alternatively, the population could be described by a combination of prompt mergers (small delays) and large delays \citep{2011ApJ...727..109V} or to the combination of two progenitor channels, i.e. binaries formed in the field or dynamically within globular clusters \citep[e.g.][]{2008MNRAS.388L...6S}. Many past works, until the most recent, feature a common approach: parametric forms are assumed for the compact binary merger delay time distribution and for the SGRB luminosity function; free parameters of such functions are then constrained through (1) the small sample of SGRBs with measured redshifts and luminosities and (2) the distribution of the $\gamma$--ray peak fluxes of SGRBs detected by past and/or present GRB detectors. A number of other observer frame properties, though, are available: fluence distribution, duration distribution, observer frame peak energy. The latter have been considered in \cite{2015MNRAS.451..126S} which, however, lacks a comparison with rest--frame properties of SGRBs as done in this article. Another issue was the comparison of the model predictions with small and incomplete samples of SGRBs with measured $z$. Indeed, only recently D14 worked with a flux--limited complete sample of SGRBs detected by \textit{Swift}. The aim of this paper is to determine the redshift distribution \psiz\ and the luminosity function \lf\ of the population of SGRBs, using all the available observational constraints of the large population of bursts detected by the \fe--Gamma Burst Monitor (GBM) instrument. These constraints are: (1) the peak flux, (2) the fluence, (3) the observer frame duration and (4) the observer frame peak energy distributions. In addition we also consider as constraints (5) the redshift distribution, (6) the isotropic energy and (7) the isotropic luminosity of a complete sample of SGRBs detected by \sw\ (D14). This is the first work aimed at deriving \lf\ and \psiz\ of SGRBs which considers constraints 2--4 and 6--7. Moreover, we do not assume any delay time distribution for SGRBs but derive directly, for the first time, their redshift distribution by assuming a parametric form. In \S2 we describe our sample of SGRBs without measured redshifts detected by \fe/GBM, which provides observer--frame constraints 1--4, and the (smaller) complete sample of \sw\ SGRBs of D14, which provides rest--frame constraints 5--7. One of the main results of this paper is that the \lf\ of SGRBs is flatter than claimed before in the literature: by extending standard analytic tools present in the literature, we show (\S3) that a steep \lf\ is excluded when all the available constraints (1--7) are considered. We then employ a Monte Carlo code (\S4) to derive the parameters describing the \lf\ and \psiz\ of SGRBs. In \S5 and \S6 the results on the \lf\ and \psiz\ of SGRBs are presented and discussed, respectively, and in \S7 we compute the local rate of SGRBs, discussing our results in the context of the dawning GW era. We assume standard flat $\Lambda$CDM cosmology with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ and $\Omega_{\rm{m}} = 0.3$ throughout the paper.

We derived the luminosity function \lf, redshift distribution \psiz\ and local rate of SGRBs. Similarly to previous works present in the literature, we fitted the properties of a synthetic SGRB population, described by the parametric \lf\ and \psiz, to a set of observational constraints derived from the population of SGRBs detected by \fe\ and \sw. Any acceptable model of the SGRB population must reproduce their prompt emission properties and their redshift distributions. Our approach features a series of improvements with respect to previous works present in the literature: \begin{itemize} \item (observer frame) constraints: we extend the classical set of observational constraints (peak flux and - for few events - redshift distribution) requiring that our model should reproduce the peak flux $P$, fluence $F$, peak energy \epo\ and duration $T$ distributions of 211 SGRBs with $P_{64}\geq 5\,\rm{ph\,s^{-1}\,cm^{-2}}$ as detected by the GBM instrument on board the \fe\ satellite. The uniform response of the GBM over a wide energy range (10 keV -- few MeV) ensures a good characterisation of the prompt emission spectral properties of the GRB population and, therefore, of the derived quantities, i.e. the peak flux and the fluence; \item (rest frame) constraints: we also require that our model reproduces the distributions of redshift, luminosity and energy of a small sample (11 events) of \sw\ SGRBs with $P_{64}\geq 3.5\,\rm{ph\,s^{-1}\,cm^{-2}}$ (selected by D14). This sample is 70\% complete in redshift and therefore it ensures a less pronounced impact of redshift--selection biases in the results; \item method: we parametrize \psiz\ as in Eq.~12 and derive the redshift distribution of SGRBs independently from their progenitor nature and their cosmic star formation history. Instead, the classical approach depends (i) on the assumption of a specific cosmic star formation history $\psi(z)$ and (ii) on the assumption of a delay time distribution $P(\tau)$; \item method: we derive our results assuming the existence of intrinsic \yone\ and \ama\ correlations in SGRBs (``case (a)''), similarly to what has been observed in the population of long GRBs. However, since evidence of the existence of such correlations in the population of SGRBs is still based on a limited number of bursts, we also explore the case of uncorrelated peak energy, luminosity and energy (``case (c)''). \end{itemize} Our main results are: \begin{enumerate} \item the luminosity function of SGRBs (case (a)), that we model with a broken power law, has a slope $\alpha_1 = 0.53^{+0.47}_{-0.14}$ (68\% confidence interval) below the break luminosity of $L_{\rm b} = 2.8^{+0.6}_{-1.89}\times 10^{52}$ erg s$^{-1}$ and falls steeply above the break with $\alpha_2 = 3.4^{+0.3}_{-1.7}$. This solution is almost independent from the specific assumption of the minimum luminosity of the \lf\ (case (b)). Moreover, it implies an average isotropic equivalent luminosity $\left\langle L \right\rangle \approx 1.5\times 10^{52}\,\rm{erg\,s^{-1}}$ (or $3\times 10^{52}\,\rm{erg\,s^{-1}}$ in case (c)), which is much larger than e.g.\ $\left\langle L \right\rangle \approx 3\times 10^{50}\,\rm{erg\,s^{-1}}$ from D14 or $\left\langle L \right\rangle \approx 4.5\times 10^{50}\,\rm{erg\,s^{-1}}$ from WP15; \item the redshift distribution of SGRBs \psiz\ peaks at $z\sim1.5$ and falls rapidly above the peak. This result is intermediate between those reported in the literature which assume either a constant large delay or a power law distribution favoring small delays. We find that our \psiz\ is consistent with the MD14 SFH retarded with a power law delay time distribution $\propto \tau^{-1}$; \item as a by-product we find that, if SGRBs feature intrinsic \yone\ and \ama\ correlations, they could be slightly steeper than those derived with the current small sample of short bursts with redshift, e.g. \cite{Tsutsui:2013lr}, but still consistent within their 68\% confidence intervals; \item if we assume that there are no correlations between \epo\ and \liso(\eiso) (case (c)), we find similarly that the \lf\ is flat at low luminosities and the formation rate peaks at slightly larger redshift ($z\sim 2$); \item we estimate the rate of SGRBs as a function of $z$ within the explorable volume of advanced LIGO and Virgo for the detection of double NS mergers or NS--BH mergers. Assuming the design aLIGO sensitivity averaged over sky location and over binary orbital plane orientation with respect to the line of sight, NS--NS mergers can be detected up to 200 Mpc (410 Mpc for NS--BH mergers). This is usually referred to as the detection \textit{range} for these binaries. The rate of SGRBs within the corresponding volume is $\sim$7$\times10^{-3}$ yr$^{-1}$ (0.028 yr$^{-1}$ for NS--BH merger distance), assuming the existence of \yone\ and \ama\ correlations for the population of short bursts (model (a)). Rates larger by a factor $\sim 4$ are obtained if no correlation is assumed (model (c)). If binaries producing observable SGRBs are preferentially face--on (which is the case if the GRB jet is preferentially aligned with the orbital angular momentum), then the actual explorable volume extends to a somewhat larger distance \citep[a factor of $\sim 1.5$ larger, see][]{schutz2011}, increasing the rates of coincident SGRB--GWs of about a factor of $3.4$ \citep{schutz2011}; \item we compare our SGRB rates with the rates of NS mergers derived from population synthesis models or from the statistics of Galactic binaries. This enables us to infer an average opening angle of the population of SGRBs of 3$^\circ$--6$^\circ$ (assuming that all SGRBs are produced by the NS--NS mergers) which is consistent with the few bursts with $\theta_{\rm jet}$ measured from the break of their afterglow light curve. \end{enumerate} \noindent Our SGRB rate estimates might seem to compromise the perspective of a joint GW--SGRB observation in the near future. We note, though, that these rates refer to the prompt emission of SGRBs whose jets point towards the Earth. SGRBs not pointing at us can still be seen as ``orphan'' afterglows (i.e.\ afterglows without an associated prompt emission - see e.g. \citealt{Ghirlanda:2015fk,Rhoads1997} for the population of long GRBs) especially if the afterglow emission is poorly collimated or even isotropic \citep[e.g.][]{Ciolfi:2015kx}. The luminosity of the afterglow correlates with the jet kinetic energy, which is thought as proportional to the prompt luminosity. Point 1 above shows that the average luminosity in the prompt emission, as implied by our result, is higher by nearly two orders of magnitude than previous findings. This enhances the chance of observing an orphan afterglow in association to a GW event (e.g. \citealt{Metzger:2015fk}). Efforts should go in the direction of finding and identifying such orphan afterglows as counterparts of GW events.

1607.07875

1607

1607.02511_arXiv.txt

We present a coherent database of spectroscopic observations of far-IR fine-structure lines from the \textit{Herschel}/PACS archive for a sample of 170 local AGN, plus a comparison sample of 20 starburst galaxies and 43 dwarf galaxies. Published \textit{Spitzer}/IRS and \textit{Herschel}/SPIRE line fluxes are included to extend our database to the full $10$--$600\, \rm{\micron}$ spectral range. The observations are compared to a set of \textsc{Cloudy} photoionisation models to estimate the above physical quantities through different diagnostic diagrams. We confirm the presence of a stratification of gas density in the emission regions of the galaxies, which increases with the ionisation potential of the emission lines. The new \oiv$_{25.9\, \rm{\micron}}$/\oiii$_{88\, \rm{\micron}}$ {\it vs} \neiii$_{15.6\, \rm{\micron}}$/\neii$_{12.8\, \rm{\micron}}$ diagram is proposed as the best diagnostic to separate: \textit{i)} AGN activity from any kind of star formation; and \textit{ii)} low-metallicity dwarf galaxies from starburst galaxies. Current stellar atmosphere models fail to reproduce the observed \oiv$_{25.9\, \rm{\micron}}$/\oiii$_{88\, \rm{\micron}}$ ratios, which are much higher when compared to the predicted values. Finally, the (\neiii$_{15.6\, \rm{\micron}}$+\neii$_{12.8\, \rm{\micron}}$)/(\siv$_{10.5\, \rm{\micron}}$+\siii$_{18.7\, \rm{\micron}}$) ratio is proposed as a promising metallicity tracer to be used in obscured objects, where optical lines fail to accurately measure the metallicity. The diagnostic power of mid- to far-infrared spectroscopy shown here for local galaxies will be of crucial importance to study galaxy evolution during the dust-obscured phase at the peak of the star formation and black-hole accretion activity ($1 < z < 4$). This study will be addressed by future deep spectroscopic surveys with present and forthcoming facilities such as \textit{JWST}, ALMA, and \textit{SPICA}.

Rest-frame mid- to far-infrared (IR) spectroscopy is a unique tool to study dust-enshrouded galaxies and, in particular, to disentangle the emission originated by star-forming activity from that originated by accretion onto supermassive black holes in the nuclei of active galaxies \citep[e.g.][]{sm92}. Optical/UV lines provide access only to relatively unobscured gas, hampering our ability to investigate the physics in obscured regions. The extinction is however negligible in the mid- to far-IR range, which contains several lines that provide information of the physical conditions even for heavily obscured gas (see Tables\,\ref{tbl_lines} and \ref{tbl_guide}). Therefore, mid- to far-IR lines are the key to investigate not only ``star formation'' in galaxies, but in general all the processes in the majority of galaxies that occur in a dust-embedded phase, and thus are hidden to optical studies. In the dark side of star formation and active galactic nuclei (AGN) activity we find critical aspects of these phenomena: e.g. deeply buried active nuclei, the onset of the star formation and AGN feeding and feedback through gas inflow and outflow. These processes have a major impact on galaxy evolution but cannot be addressed from an optical/UV perspective. The potential of mid- to far-IR spectroscopy was already proven by the \textit{Infrared Space Observatory} \citep[\textit{ISO};][]{kes96}, using the Short Wavelength Spectrometer \citep[SWS;][]{deg96} and the Long Wavelength Spectrometer \citep[LWS;][]{cle96}. As mid- and far-IR fine-structure emission lines are sensitive to the physical conditions of the interstellar medium (ISM), these transitions can be used to investigate its different phases \citep[e.g.][]{stu00,neg01,spi05} and characterise the primary spectrum of the ionising radiation \citep[e.g.][]{ale99}. \textit{ISO} enabled the development of the first IR diagnostics to separate the AGN, the ISM, and the stellar contributions \citep[e.g.][]{stu02,bra08}, and shed light on the nature of ultra-luminous IR galaxies \citep[ULIRGs;][]{gen98}. The full exploitation of the mid-IR window was possible thanks to the InfraRed Spectrograph \citep[IRS;][]{hou04} on-board the \textit{Spitzer Space Telescope} \citep{wer04}. The study of large samples of galaxies including ULIRGs, Quasars, Seyfert galaxies, radio, and starburst galaxies \citep{arm07,tom10,vei09,bra06,bs09} showed, e.g., the ability of the \nev$_{14.3, 24.3\, \rm{\micron}}$ and the \oiv$_{25.9\, \rm{\micron}}$ lines, associated to high-density and high-excitation photoionised gas ($n_{\rm e} \gtrsim 10^3\, \rm{cm^{-3}}$, I.P.\,$\gtrsim 54\, \rm{eV}$), to investigate: \textit{i)} the inner part of the narrow-line region (NLR) and identify optically-hidden AGN \citep{arm07}; \textit{ii)} the AGN contribution in ULIRGs \citep{vei09}; \textit{iii)} the AGN nature of radio-galaxies \citep{haa05,lei09}; \textit{iv)} the power of the extinction-free density tracers in the $10^2$--$10^4\, \rm{cm^{-3}}$ range based on the \siii$_{18.7, 33.5\, \rm{\micron}}$ and the \nev$_{14.3, 24.3\, \rm{\micron}}$ lines \citep{tom08,tom10}; \textit{v)} star formation tracers based, e.g., on \neii$_{12.8\, \rm{\micron}}$ and \neiii$_{15.6\, \rm{\micron}}$ \citep{ho07}, or polycyclic aromatic hydrocarbon (PAH) features \citep{odo09,sar09}; \textit{vi)} the warm molecular gas traced by H$_2$ lines \citep[e.g.][]{tom08,bau10}; \textit{vii)} the extreme ultraviolet (EUV) spectra of AGN revealed by high-ionisation fine-structure IR lines \citep{mel11}. In the far-IR, the \textit{Herschel Space Observatory} \citep{pil10} provided a large gain in spectroscopic sensitivity when compared to \textit{ISO}, allowing to access this spectral range for a larger number of galaxies in the nearby Universe. The Photoconductor Array Camera and Spectrometer \citep[PACS;][]{pog10} and the Spectral and Photometric Imaging REceiver Fourier-transform spectrometer \citep[SPIRE;][]{gri10,nay10} sampled the $50$--$210\, \rm{\micron}$ and $200$--$670\, \rm{\micron}$ ranges, respectively. This spectral region covers the main cooling lines of photo-dissociation regions (PDR), \oi$_{63, 145\, \rm{\micron}}$ and \cii$_{158\, \rm{\micron}}$, which provide information of the cold neutral gas \citep{th85}, as well as the high-\textit{J} CO transitions \citep{kam14,kam15}, allowing the investigation of the excitation of molecular gas; the fine-structure lines of \oi$_{63\, \rm{\micron}}$ and \nii$_{122\, \rm{\micron}}$ as tracers of the IR luminosity and star formation rate in ULIRGs \citep{far13}; the origin of \cii$_{158\, \rm{\micron}}$ emission, mainly associated to the PDR \citep[][hereafter S15]{s15}; X-ray dissociation regions \citep[XDR;][]{spi12b}. The aim of the present study is to extend for the first time, with a statistical approach, the spectroscopic work done by \textit{Spitzer} in the mid-IR to the longer wavelengths using \textit{Herschel}. Taking advantage of the \textit{Herschel} scientific archive, we extended the previous study presented in S15, limited to a sample of 26 Seyfert galaxies, to a total sample of 170 active galaxies, and a comparison sample of 20 starburst galaxies and 43 dwarf galaxies --\,the latter taken from the Dwarf Galaxy Survey \citep{mad13}. Of particular interest is the development of the diagnostics that will be exploited by future IR observatories, e.g. the \textit{James Webb Space Telescope} \citep[\textit{JWST};][]{gar06} and the \textit{SPace Infrared telescope for Cosmology and Astrophysics} \citep[\textit{SPICA};][]{swi09} for galaxies up to the peak of the star formation rate density (SFRD) at $1 < z < 4$, but also the Atacama Large Millimeter/submillimeter Array (ALMA) for far-IR rest-frame observations of galaxies at high-redshift ($z \gtrsim 3$). The combination of mid- and far-IR lines expands the possible diagnostics to, e.g., the \oiv$_{25.9\, \rm{\micron}}$/\oiii$_{88\, \rm{\micron}}$ ratio (hereafter \oiv$_{25.9}$/\oiii$_{88}$) that has been proposed as a powerful diagnostic to discriminate AGN from star formation activity (S15). The larger statistics in the present work allow us to perform a more robust analysis of the diagnostics already tested in S15, including new line ratios sensitive to metallicity and ionisation \citep{nag11,nag12}. We include in this work the results on the ``compact sample'' of 43 dwarf galaxies presented by \citet{cor15} to investigate how strong star formation activity in low-metallicity environments can be well separated from either AGN and ``normal'' starburst galaxies through mid- and far-IR line ratios. A set of \textsc{Cloudy} photoionisation models were developed using AGN and starburst galaxies as ionisation sources in order to interpret the behaviour of the observed line ratios and their dependence on density, ionisation parameter, and metallicity. The text is organised as follows. The selection of the sample of galaxies is explained in Section\,\ref{sample}, Section\,\ref{obs} presents the data reduction and describes the catalogs taken from the literature and included in our study, the \textsc{Cloudy} photoionisation models are described in Section\,\ref{models}, the results of this work are detailed in Section\,\ref{results}, and the main conclusions are summarised in Section\,\ref{summ}. Due to the large dataset and associated tables needed to show the results of this study, the majority of the latter appear in their complete form only in the online version of this journal.

\label{summ} In this work we have presented the complete database of far-IR fine-structure lines observed by \textit{Herschel}/PACS during its 3.5 years operational life for 170 AGN-classified galaxies in the \citet{ver10} catalog. In order to complete the full mid- to far-IR spectra in the $10$--$600\, \rm{\micron}$ range we collected published fine-structure line fluxes measured with \textit{Spitzer}/IRS and \textit{Herschel}/SPIRE. As a comparison sample, we compiled an equivalent database for starburst galaxies extracted from \citet{bs09} and \citet{ga09}. Additionally, we included 43 dwarf galaxies from the Dwarf Galaxy Survey \citep{mad13,cor15}, in order to probe a more extreme star formation environment at low metallicities. The photoionisation code \textsc{Cloudy} has been used to reproduce the physical conditions of the gas exciting the mid- to far-IR fine-structure lines. We produced four models using different ionising spectra: a pure AGN model ($\alpha = -1.4$; $S_\nu \propto \nu^{\alpha}$), a LINER model ($\alpha = -3.5$), a starburst galaxy model ($20\, \rm{Myr}$ continuum star formation with $Z = \rm{Z_\odot}$), and a dwarf galaxy model ($1\, \rm{Myr}$ instant starburst with $Z = 1/5\, \rm{Z_\odot}$). Calculations for the starburst and the dwarf galaxy models have been extended down to $T = 50\, \rm{K}$ in order to include the PDR emission region. The main results of this study are: \begin{itemize} \item The \oi$_{145/63}$ line ratio, used as temperature tracer in the $100$--$400\, \rm{K}$ range for the neutral gas in the PDR, does not show a clear correlation with the ionisation nor the density of the ionised gas, traced by the \siv$_{10.5}$/\siii$_{18.7}$ and the \siii$_{33.5/18.7}$ ratios, respectively. High \oi$_{145/63}$ ratios of $\gtrsim 0.1$ are found in objects with self-absorption in the \oi$_{63\, \rm{\micron}}$ line profile. \item The \ci$_{609/371}$ line ratio, sensitive to the temperature in the PDR ($20$--$100\, \rm{K}$), show similar median values for AGN ($0.53 \pm 0.21$) and starburst galaxies ($0.45 \pm 0.11$). This is much higher than the ratios predicted from XDR simulations ($\sim 0.15$--$0.19$), suggesting that the neutral carbon lines in our AGN sample are likely dominated by relatively low-density PDR emission originated in the ISM of the host galaxies. \item The density stratification found in S15 is confirmed here with 155 pairs of lines from \nii$_{205/122}$, \siii$_{33.5/18.7}$, \oiii$_{88/52}$, and \nev$_{24.3/14.3}$ line ratios, plus 93 upper limits. Both the least squares fit and the Kaplan-Meier residuals fit are consistent with an increasing gas density --\,measured by the line ratios\,-- with the ionisation potential (IP) of the transition. This suggests that harder radiation fields are traced by the higher density gas found in the innermost part of the NLR. \item The line ratios of S1 and S1h are always very similar, suggesting that both AGN types are indistinguishable from an IR perspective, thus S1h correspond to optically obscured S1 \citep[e.g.][]{tom10}. \item The \cii$_{158}$/\nii$_{122, 205}$ line ratios, sensitive to the relative contributions of the PDR and the low excitation photoionised gas, suggest a major contribution ($\gtrsim 80\%$) of PDR emission to the \cii$_{158\, \rm{\micron}}$ line. The observed ratios \cii$_{158}$/\nii$_{122, 205}$ $\gtrsim 3$ are in agreement with the simulations stopped at $T_{\rm stop} = 50\, \rm{K}$, which include PDR emission. Values lower than those observed are predicted for pure photoionised gas ($T_{\rm stop} = 1000\, \rm{K}$). An additional correlation of \cii$_{158}$/\nii$_{122, 205}$ with ionisation potential (traced by \neiii$_{15.6}$/\neii$_{12.8}$) might also be present, as \nii \ emission decreases in favour of \niii. \item The inclusion in the diagnostic diagrams of dwarf galaxies demonstrates that classical line ratios such as \oiv$_{25.9}$/\neii$_{12.8}$ fail as AGN/starburst tracers at low metallicities. Instead, the \oiv$_{25.9}$/\oiii$_{88}$ ratio is an excellent tracer to discriminate between non-thermal excitation in AGN and thermal ionisation produced by any kind of star formation activity. Since it probes a relatively hard range of the spectrum ($54.94$--$35.12\, \rm{eV}$, $4$--$2.6\, \rm{Ry}$) this line ratio is ideal for composite objects with extreme star formation bursts and mergers involving a chemically unevolved companion, a scenario that seems to be more common with increasing redshift. Alternatively, the \oiv$_{25.9}$/(\neiii$_{15.6}$+\neii$_{12.8}$) ratio can also be used as a proxy for the relative AGN to starburst contribution. \item A new AGN/star formation diagnostic diagram is proposed, based on the \oiv$_{25.9}$/\oiii$_{88}$ and \neiii$_{15.6}$/\neii$_{12.8}$ (or alternatively \siv$_{10.5}$/\siii$_{18.7}$). The \oiv$_{25.9}$/\oiii$_{88}$ ratio is sensitive to the relative contributions from AGN and star formation, while the \neiii$_{15.6}$/\neii$_{12.8}$ ratio (\siv$_{10.5}$/\siii$_{18.7}$) is sensitive to the ionisation, mostly driven by the metallicity in the stellar population. The combination of these line ratios allows to clearly separate the different populations in the diagram, i.e. AGN, dwarf galaxies, and starburst galaxies, and provides an extinction-free diagnostic ideal for future mid- and far-IR spectroscopic surveys. \item Photoionisation by current stellar atmosphere models cannot reproduce the observed \oiv$_{25.9}$/\oiii$_{88}$ line ratios due to the lack of predicted photons above the He\,\textsc{ii} ionisation edge ($\gtrsim 54\, \rm{eV}$, $4\, \rm{Ry}$). Furthermore, both dwarf and starburst galaxies show similar \oiv$_{25.9}$/\oiii$_{88}$ line ratios in the $5 \times 10^{-3}$--$3 \times 10^{-1}$ range, with no apparent dependency on the metallicity of the stellar population. An improved treatment of stellar atmosphere models above $54\, \rm{eV}$ will be necessary in order to explain the observed line ratios, e.g. by including the contribution of shocks and non-thermal emission in the EUV. \item LINERs show intermediate line ratios of \oiv$_{25.9}$/\oiii$_{88}$, \neiii$_{15.6}$/\neii$_{12.8}$, and \siv$_{10.5}$/ \siii$_{18.7}$, between Seyfert and starburst galaxies. These line ratios can be explained by assuming a softer UV spectrum with a slope of $\alpha \approx -3.5$, in line with the receding accretion disc scenario expected at low-luminosities and/or low-accretion efficiencies. \item The (\neiii$_{15.6}$+\neii$_{12.8}$)/(\siv$_{10.5}$+\siii$_{18.7}$) line ratio is sensitive to the gas metallicity for AGN, starburst and dwarf galaxies, this is confirmed by the models including the effect of Sulphur depletion. Thus, this ratio is proposed as a powerful extinction-free metallicity tracer, ideal for dust regions and obscured objects where optical determinations are not possible. Future facilities, such as \textit{JWST} for the Local Universe and \textit{SPICA} along Cosmic history, at the peak of star formation and black hole accretion activity ($1 < z < 4$), will be able to trace the build up of heavy elements during galaxy evolution. \end{itemize}

1607.02511

1607

1607.03517_arXiv.txt

The Mission Accessible Near-Earth Objects Survey (MANOS) aims to physically characterize sub-km Near-Earth Objects (NEOs). We report first photometric results from the survey which began in August, 2013. Photometric observations were performed using 1~m to 4~m class telescopes around the world. We present rotational periods and lightcurve amplitudes for 86 sub-km NEOs, though in some cases, only lower limits are provided. Our main goal is to obtain lightcurves for small NEOs (typically, sub-km objects) and estimate their rotational periods, lightcurve amplitudes, and shapes. These properties are used for statistical study to constrain overall properties of the NEO population. A weak correlation seems to indicate that smaller objects are more spherical than the larger ones. We also report 7 NEOs that are fully characterized (lightcurve and visible spectra) as the most suitable candidates for a future human or robotic mission. Viable mission targets are objects fully characterized, with a $\Delta v$$^{NHATS}$ $\leq$12~km~s$^{-1}$, and a rotational period P$>$1~h. Assuming a similar rate of object characterization as reported in this paper, approximately 1,230 NEOs need to be characterized in order to find 100 viable mission targets.

Near Earth Objects (NEOs) are minor bodies (asteroids, comets, meteoroids) on orbits with perihelia distances q$\textless$1.3~AU. As of April 2016, 14,263 NEOs have been discovered\footnote{Numbers from the Minor Planet Center: \url{http://www.minorplanetcenter.net/}}. About 90$\%$ of NEOs originated in the asteroid belt and have a rocky nature \citep{Jewitt2002, DeMeo2008}. Despite the impressive number of discovered NEOs, physical information for these objects remains limited. Rotational light curves are one tool to constrain the physical evolution of these objects. The rotational states of asteroids provide information about physical properties such as a lower limit to density, internal structure, cohesion, and shape or surface heterogeneity \citep{Pravec2000, Holsapple2001, Holsapple2004}. Large objects (diameter greater than 1~km) have been well-studied with photometric, spectroscopic, and/or radar techniques \citep{Benner2015, Thomas2014, Thomas2011, Warner2009, Pravec2002, Binzel2002}, but small objects are also of interest for a number of reasons. First, objects in the meter to decameter size regime can impact the Earth on human timescales, as opposed to the 10$^{6}$~years impact interval of km-scale objects \citep{Harris2015}. As evidenced in Chelyabinsk, Russia in 2013 \citep{Popova2013} relatively small objects can pose a modest impact hazard. In addition, these small NEOs are the immediate parent bodies of meteorites. To interpret meteorites in an astrophysical context requires that we better understand their source population. In addition, studying these small objects can provide deeper insight into size-dependent evolutionary processes such as the radiative Yarkovsky and Yarkovsky-O'Keefe-Radzievskii-Paddack (YORP) effects \citep{Bottke2006}. Finally, the much greater number of NEOs with sizes of $\sim$100~m compared to km-size objects provides more opportunities for detailed physical study. This includes increased possibilities for a variety of exploration mission scenarios (e.g., \citet{Abell2009}) as well as more frequent near-Earth encounters to study physical changes associated with gravitational perturbation events (e.g., \citet{Scheeres2005}; \citet{Binzel2010}). Because NEOs have their origin in the Main Belt of asteroids, and are the result of multiple collisions, their shape as well as rotational properties are valuable tracers of their evolution. \citet{Binzel2002} suggested that NEOs should be similar in rotation and shape to similar sized Main Belt asteroids, and spin distribution of km-size Main Belt Asteroids and similar sized NEOs supports this assertion \citep{Pravec2008, Polishook2009}. Therefore, the study of small NEOs (sub-km objects) allow us to infer the properties of small Main Belt asteroids, which remain unobservable with current facilities. We present here a study focused on the rotational properties of sub-km NEOs. Our ultimate goal is to obtain the most comprehensive data-set of sub-km NEOs to date, allowing a homogeneous and detailed study of the shape, surface and rotational properties of these objects. This paper is divided into six sections. In the next section, we introduce briefly the Mission Accessible Near-Earth Objects Survey (MANOS). Then, we describe the observations and the data set analyzed. In Section 4, we present our main results regarding rotational period and lightcurve amplitudes of our targets. In Section 5, we discuss our results and compare them to the literature. In Section 6, we put constraints on the internal structure of NEOs. Finally, Section 7 is dedicated to the summary and conclusions of this work.

We present a homogeneous dataset composed of 86 objects (data reduced and analyzed the same way). We report rotational periods and lightcuve amplitude for most of them, but in some cases, we only report constraints on these properties. We report that 70~$\%$ of our sample shows at least one full rotation or a partial lightcurve with a period estimate. We report partial lightcurves for 14~$\%$ of our sample, i.e. objects that do not show a clear rotation period but do show a clear increase or decrease in magnitude. Finally, 16~$\%$ of the lightcurves are flat. Most of the observed obejcts are small and fast rotators, with $\sim$50$\%$ of objects spinning in less than 5~min. MANOS found two ultra-fast rotators: 2014~RC, and 2015~SV$_{6}$ with rotational periods of 15.8~s and 17.6~s, respectively. Discovery of these objects confirmed that MANOS is highly sensitive to the detection of fast spinning objects thanks to the use of large facilities allowing us to use short exposure time. We also highlight fast rotators in the Amor population, confirming again that our survey is sensitive to fast rotating small NEOs. We studied rotational frequency distribution according to size, and dynamical class. We noted an excess of both slow and fast rotators that does not allow us to fit a Maxwellian distribution to the observable distribution. Rotational periods are not significantly different in the Amor, Aten or Apollo groups. Axis ratio corrected from phase angle has been derived for our MANOS sample. No strong correlation between axis ratio and size or axis ratio and period has been found. Among the 30 mission accessible MANOS targets with complete lightcurves, six objects have rotational periods higher than 2~h, whereas three have periods between 1 and 2~h. The rest (i.e. 21 objects) have periods ranging from few seconds up to 1~h: 10 objects are rotating in less than 5~min, 3 objects have period between 5 and 10~min, and 8 with period longer than 10~min. Their sizes range from 3~m to 215~m, i.e. an absolute magnitude of 29.6 to 20.7. In conclusion, 33 of our 86 MANOS targets are mission accessible according to NHATS (i.e. $\Delta v$$^{NHATS}$$\leq$12~km~s$^{-1}$, launch window 2015-2040), 26 of these 33 are fully characterized with lightcurve (partial and flat lightcurves also considered) and visible spectrum. Only 7 objects of the objects presented here meet the NHATS dynamical criteria and the 1~h rotation limit \citep{Abell2009}: 2002~DU$_{3}$, 2010~AF$_{30}$, 2013~NJ, 2014~YD, 2015~CO, 2015~FG$_{36}$, and 2015~OV. Assuming a similar rate of object characterization as reported in this paper, $\sim$1,230 objects (i.e. approximately 10$\%$ of the known NEO population) need to be characterize in order to find 100 viable mission targets. Approximately 400,000 NEOs with diameter between 10~m and 1~km are estimated \citep{Tricarico2016}. To find 100 viable mission targets, $\sim$0.3$\%$ of the estimated population need to be characterized. \citet{Harris2015} estimated a population of $\sim$8$\times$10$^{7}$ objects in the 10~m-1~km size range, and so $\sim$0.002$\%$ of this population has to be characterized in order to find 100 viable targets. This means that $\sim$33,000 NEOs are expected to be mission accesible targets using the \citet{Tricarico2016} estimate, whereas the \citet{Harris2015} value gives us a total of $\sim$6,000,000 objects. As our main goal is to get a large set of fully characterized objects (lightcurve, and visible spectra), it is important to complete the study of some objects that have been partially characterized (incomplete/unknown lightcurve and/or no spectra, Table~\ref{Tab:candidates}). For most of these objects, their next optical windows are within the next ten years. NHATS generates the next optical windows through to the year 2040 (see NHATS webpage for more details). Unfortunately, because of their highly uncertain orbits most of these objects will be lost by then and current/future surveys will have to re-discover them. It is also important to point out that because of their uncertain orbits, their next windows of visibility can be off as well as their visual magnitude.

1607.03517

1607

1607.08605_arXiv.txt

{We report the results of a comprehensive analysis of X-ray ({\it Chandra} and {\it Swift} observatories), optical (Southern African Large Telescope, {\it SALT}) and near-infrared (the {\it VVV} survey) observations of the Be/X-ray binary pulsar \2s. Accurate coordinates for the X-ray source are determined and are used to identify the faint optical/infrared counterpart for the first time. Using {\it VVV} and SALTICAM photometry, we have constructed the spectral energy distribution (SED) for this star and found a moderate NIR excess that is typical for Be stars and arises due to the presence of circumstellar material (disk). A comparison of the SED with those of known Be/X-ray binaries has allowed us to estimate the spectral type of the companion star as B1-2V and the distance to the system as $>15$ kpc. This distance estimation is supported by the X-ray data and makes \2s\ one of the most distant X-ray binaries within the Milky Way, residing on the far side in the Scutum-Centaurus arm or even further.}

Determining the nature of Galactic X-ray sources through their optical counterparts and measurements of their distance is key not only in the study of individual sources, but also in population studies of different classes of objects. In the case of X-ray emission from a compact object (namely an X-ray binary) the former allows one to understand physical processes near the compact object. Only a good knowledge of the optical counterparts to sources of X-ray emission allows testing of different theoretical models of the emission mechanisms and accretion processes. In turn, properties of different populations of X-ray sources in our and other galaxies contain information about fundamental physical mechanisms responsible for their formation and evolution and about properties of the host galaxy itself. Distance estimations can be made relatively straightforwardly for X-ray sources in other galaxies, but is a non-trivial task in our own galaxy. There are at least two problems to contend with: the high density of stars, which complicates the process of optical identification, and the large and variable interstellar absorption, which complicates measurements of different properties of the source, including its distance. These difficulties have led to the situation where a significant number of Galactic X-ray sources are still unclassified or have no reliable distance measurements \citep[][]{liu06,liu07}. In this work we utilize high quality X-ray, optical and near-infrared (NIR) data to localize and identify the poorly studied transient X-ray pulsar \2s. It was discovered in 1975 by the {\it SAS-3} observatory \citep{1976IAUC.2959....2W}. Later a strong coherent variability, with a period of 9.3 s, was found \citep{1982IAUC.3667....3K} in the source light curve. Subsequent observations of \2s revealed strong outbursts in 2007 \citep{2007ATel.1345....1K} and most recently in 2015 \citep{2015ATel.7018....1S}. Observations with the {\it NuSTAR} observatory and {\it Fermi}/GBM monitor during the 2015 outburst led \citet{tsygankov2016} to the discovery of a cyclotron absorption line in the source spectrum at $\simeq23.5$ keV and to improve the accuracy of the binary parameters. Moreover these authors also estimated the distance to the system $d\simeq20$ kpc. The possible nature of \2s\ as a pulsating Be/X-ray binary system was first suggested on the basis on its transient activity \citep{1983ApJ...274..765K} and X-ray spin modulation. However, no optical or infrared counterpart had been directly determined until now. This made \2s\ one of the first transient X-ray binaries to be discovered with a probable Be-companion, though the confirmation of this and the distance to the system were still to be confirmed. Moreover, the X-ray behaviour of the source is quite unusual -- in particular, the source never displays type I outbursts. \citet{oka_neg01} proposed the model of the truncated disc in which such a behaviour is naturally explained for systems with a low eccentricity. Indeed, current measurements of \citet{tsygankov2016} revealed a very low eccentricity in this system $e\simeq0.035$. Based on recent data from the {\it Chandra} and {\it Swift} observatories, Southern African Large Telescope (SALT) and the {\it VVV} survey, we have measured for the first time an accurate position for the X-ray pulsar \2s, leading to the identification of its infrared counterpart and an estimation of its distance.

In this paper we have used {\it Chandra} and {\it Swift}/XRT data to accurately determine the coordinates of the long known, but poorly studied, transient X-ray pulsar \2s\ for the first time. The subsequent optical observations with the \textit{SALT} telescope and a comprehensive analysis of the \textit{VVV} survey data led to the identification of its likely infrared counterpart, whose spectral energy distribution is consistent with a Be-star, as expected from its previous classification as a BeXRB with an accreting X-ray pulsar. Moreover, a comparison of the measured SED with those of known Be/X-ray binaries demonstrates the latter are similar and has allowed us to estimate the spectral type of the \2s\ companion star as B1-2V and the distance to the system as $>15$ kpc. This value agrees well with the independent estimations of $d=20\pm4$ kpc from the temporal properties of the pulsar \citep{tsygankov2016}. These distance estimates place \2s\ in the Scutum-Centaurus arm, on the far side of the Galaxy from the Sun, or even further, depending on the possible spectral type of its companion star. This result can be considered as the first reliable distance determination to \2s\ and, for an outlying Galactic object, makes \2s\ one of the most distant Galactic high-mass X-ray binaries known. We note that estimations of distances to other distant Galactic sources have, as a rule, quite large uncertainties and still require additional confirmations \citep[see, e.g.,][]{tsygankov2005,shaw09,karas10,pel11}. In turn, the determination of distances to high-mass X-ray binaries at the furthest regions of the Galaxy, behind the Galactic Center, is very important for determining the space density of such objects \citep{2013MNRAS.431..327L}. The measurement of the distance to \2s\ ($d>15$ kpc) has allowed us to estimate a lower limit to its unabsorbed luminosity near the outburst maximum as $L_{0.5-10 keV}\simeq3.2\times10^{37}$ \lum\ and in the faint state $L_{0.5-10 keV}\simeq1.1\times10^{34}$ \lum, from the {\it Swift}/XRT and {\it Chandra} data, respectively. Taking into account that the ratio between the X-ray bolometric source flux and its flux in the 0.5-10 keV energy band is about 2 (from the analysis of the source broadband spectrum, obtained by {\it NuSTAR} and our results, which demonstrate that its spectral shape is not dependent on the luminosity) we can estimate a lower limit to the maximum (bolometric) X-ray luminosity of the source $L_{X, max}\simeq6.4\times10^{37}$ \lum. The value $L_{X, max}\simeq1.2\times10^{38}$ \lum, obtained for a distance of 20 kpc, is comparable with that measured for other bright Be/X-ray binary transients such as V\,0332+53 and 4U\,0115+63, and implies that the observed outbursts from \2s\ are Type II in nature. The corresponding source luminosity in quiescence, namely $L_{X, faint}\sim(2.2-4)\times10^{34}$ \lum\, is high enough for testing different models for the low accretion regime and neutron star cooling scenario, should one obtain sufficiently long exposures in the faint X-ray state.

1607.08605

1607

1607.04966_arXiv.txt

Numerical hydrodynamic simulation of inviscid and viscous flows have shown that significant outflows could be produced from the CENtrifugal pressure supported BOundary Layer or CENBOL of an advective disk. However, this barrier is weakened in presence of viscosity, more so, if there are explicit energy dissipations at the boundary layer itself. We study effects of viscosity and energy dissipation theoretically on the outflow rate and show that as the viscosity or energy dissipation (or both) rises, the prospect of formation of outflows is greatly reduced, thereby verifying results obtained through observations and numerical simulations. Indeed, we find that in a dissipative viscous flow, shocks in presence of outflows can be produced only if the Shakura-Sunyaev viscosity parameter $\alpha$ is less than $0.2$. This is a direct consequence of modification of the Rankine-Hugoniot relation across the shock in a viscous flow, when the energy dissipation and mass loss in the form of outflows from the post-shock region are included. If we ignore the effects of mass loss altogether, the standing dissipative shocks in viscous flows may occur only if $\alpha <0.27$. These limits are tighter than the absolute limit of $\alpha=0.3$ valid for a situation when the shock itself neither dissipates energy nor any outflow is formed. We compute typical viscosity parameters required to understand spectral and temporal properties of several black hole candidates such as GX399-4, MAXI J1659-152 and MAXI J1836-194 and find that required $\alpha$ are indeed well within our prescribed limit.

Outflows in most of the astrophysical objects are considered to be associated with accretion phenomena \citep{livio97}. Indeed, they are generally thought to be produced from the boundary layers of the central objects, including black holes (\citealt[][hereafter C96a]{c96a};\citealt{c99,dc99}), paradoxical as it may sound. In a standard Keplerian disk, where the centrifugal force is totally balanced by gravity, no boundary layer could be formed around a black hole. This is not the case when the angular momentum is `not-Keplerian' such as when the flow is transonic which has low and almost constant angular momentum. Here, the outward centrifugal force slows down the infalling matter forming a static or oscillating shock (see, C96a and references therein) or, even a shock free flow with a broad and diffused centrifugal barrier \citep{c97}. The post-shock region is the so-called CENtrifugal pressure supported BOundary Layer or CENBOL. It has been shown that a large region of the parameter space allows such a shock formation \citep[e.g.,][and references therein]{c96b}. When the flow has some viscosity, this parameter space starts to shrink \citep{cd04}. Furthermore, when the thermal energy at the base of the jet is dissipated, the drive required to form outflows is also reduced. So, it is pertinent to ask how the outflow rate depends on both the energy dissipation and the viscosity at the base of the flow, i.e., the CENBOL. Even more important is to know the upper limit of viscosity parameter, when all the flow parameters such as the specific energy and angular momentum at the inner edge of the disk are held constant, which will still allow the formation of shocks in presence of these dissipations. There are estimates of $\alpha$ parameter in accretion disks both from observational and numerical simulations. If these estimates are lower as compared to our limits, it would indicate that CENBOL would form and radiation emitted from it is an integral part of the spectral and timing properties of black holes. \defcitealias{nagchak16}{Paper~I} It has been shown recently (\citet{nagchak16} [\citetalias{nagchak16}]; see also, \citet{kumarchatt13}, with a different viscous stress prescription), that there is an upper limit on the \citet{ss73} viscosity parameter above which three sonic points and therefore standing shocks are not possible in a viscous accretion flow. These results were predicted by \citet{c90a} and \citet{c96a} in the context of isothemal and polytropic flows respectively where it was shown that the topology of solutions change dramatically beyond a critical value of viscosity. In Paper I, the limit is shown to be $\alpha_{sup} \sim 0.3$. In presence of dissipation at the shock and mass outflow, this limit is likely to change. In the present paper, we address this very important issue and show that indeed, shock dissipation tightens the limit to $\alpha_{sup} \sim 0.27$. If, furthermore, outflows are included, then one requires to have $\alpha_{sup} \sim 0.2$ for shock formation. This reduces the available parameter space even further. Nevertheless, as we show below, this limit is not low enough to prohibit formation of the shocks since even the estimated viscosities in flows from observations are much lower than this value. Thus the observational evidences of shocks which decide on the spectral properties of black holes and shock oscillations in explaining quasi-periodic oscillations (QPOs) remain valid even when outflows from dissipative shocks are present. In the next Section, we present the model equations used for our study. In \S 3, we present procedures to solve for the transonic flows properties, including the shock locations. In \S 4, we discuss how the parameter space, spanned by the specific energy at the inner sonic point and the specific angular momentum on the black hole horizon behave in presence of viscosity, energy dissipation and outflow rate. In \S 5, we show a similar analysis for the data from outbursts of different black hole candidates. Finally, in \S 6, we draw our conclusions.

1607.04966

1607

1607.06152_arXiv.txt

We used the newly commissioned 50 cm Binocular Network (50BiN) telescope at Qinghai Station of Purple Mountain Observatory (Chinese Academy of Sciences) to observe the old open cluster NGC 188 in $V$ and $R$ as part of a search for variable objects. Our time-series data span a total of 36 days. Radial-velocity and proper-motion selection resulted in a sample of 532 genuine cluster members. Isochrone fitting was applied to the cleaned cluster sequence, yielding a distance modulus of $(m-M)_V^0=11.35\pm0.10$ mag and a total foreground reddening of $E(V-R)=0.062\pm0.002$ mag. Light-curve solutions were obtained for eight W Ursae Majoris eclipsing-binary systems (W UMas) and their orbital parameters were estimated. Using the latter parameters, we estimate a distance to the W UMas which is independent of the host cluster's physical properties. Based on combined fits to six of the W UMas (EP Cep, EQ Cep, ES Cep, V369 Cep, and---for the first time---V370 Cep and V782 Cep), we obtain an average distance modulus of $(m-M)_V^0=11.31 \pm 0.08$ mag, which is comparable with that resulting from our isochrone fits. These six W UMas exhibit an obvious period--luminosity relation. We derive more accurate physical parameters for the W UMa systems and discuss their initial masses and ages. The former show that these W UMa systems have likely undergone angular-momentum evolution within a convective envelope (W-type evolution). The ages of the W UMa systems agree well with the cluster's age.

W Ursae Majoris (W UMa) variables are low-mass, so-called ``overcontact'' binary systems, where the Roche lobes of both stellar components are filled. W UMas share a common convective envelope. Both components are characterized by rapid rotation, with periods ranging from $P = 0.2$ d to $P = 1.0$ d. It is straightforward to obtain complete and high-quality W UMa light curves in a few nights of observing time on small to moderate-sized telescopes. Since W UMas are very common in both old open clusters (OCs) and the Galactic field, they have significant potential as distance indicators. Approximately 0.1\% of the F-, G-, and K-type dwarfs in the solar neighborhood are W UMas \citep{Duerbeck84}, while in OCs their frequency may be as high as $\sim$0.4\% \citep{Rucinski94}. Although the occurrence frequency of main-sequence contact binaries in old globular clusters is low, the frequency of ``blue straggler''-type contact binary systems is two to three times higher there than that in OCs \citep{Rucinski00}. W UMas can reveal the evolutionary history of their host cluster, since they are thought to result from dynamical interactions in the cluster. Alternatively, these systems may represent a possible final evolutionary phase of primordial binary systems, once these binaries have lost most of their angular momentum. Since W UMas are more than 4 mag fainter than Cepheids or RR Lyrae variables, studies of W UMa distances have only been undertaken for just a few decades. \citet{Rucinski97} used W UMas as distance tracers for 400 objects from the Optical Gravitational Lensing Experiment (OGLE) variable-star catalog. If a reliable (orbital) period--luminosity (PL) relation can be established for W UMa systems, they could potentially play a similarly important role as Cepheids in measuring the distances to old structures in the Milky Way, including those traced by old OCs and the Galactic bulge. Although distances based on individual W UMas are not as accurate as those resulting from Cepheid analysis, their large numbers could potentially overcome this disadvantage. Clusters represent good stellar samples to study distances, because their distances can be estimated in a number of independent ways \citep[e.g.,][]{Chen15}. In our modern understanding, these systems are most likely formed through either nuclear evolution of the most massive component in the detached phase (A subtype) or angular-momentum evolution of the two component stars within a convective envelope (W subtype) \citep{Hilditch1988, Yildiz13}. \citet{Yildiz14} provided a method to calculate the typical evolution timescales of W UMa systems. Armed with accurate distances, significantly improved physical parameters (such as masses and luminosities) can be determined. In turn, such W UMa systems can then be used to constrain the evolution of W UMa systems in general. NGC 188 is an old OC at a distance of $\sim$2 kpc. It contains a large number of W UMa variables. Seven W UMas near its center were first found by \citet{Hoffmeister64} and \citet{Kaluzny87}. \citet{Zhang02, Zhang04} undertook a detailed survey covering 1 deg$^2$ around the cluster center and found 16 W UMas. \citet{Branly96} calculated light-curve solutions for five central W UMas using the Wilson--Devinney (W--D) code and discussed the W UMa distance in relation to the cluster distance. \citet{Liu11} obtained orbital solutions for EQ Cep, ER Cep, and V371 Cep, while \citet{Zhu14} published similar results for three additional W UMas, i.e, EP Cep, ES Cep, and V369 Cep. In this paper, we study all eight previously identified NGC 188 W UMas based on high-cadence observations obtained over a continuous period of more than two months, resulting in accurate, self-consistent, and homogeneous magnitudes and a well-determined average distance, relying on up to 3000 data points for a single W UMa system. We also establish a physical relationship between these eight variables and their host cluster, NGC 188, based on proper motions, radial velocities, and features in the color--magnitude diagram (CMD). Although Cepheid variables are among the most useful objects to establish the distance ladder, the small number of Cepheids in our Galaxy introduces relatively large statistical errors, while their disparate distances introduce comparably large systematic errors in distance modulus, reddening \citep{An07}, and metallicity \citep{Sandage08}. Compared with bright O- and B-type Cepheids, faint W UMa dwarfs are much more plentiful. Our aim is to obtain more accurate cluster distances than available to date based on our new W UMa observations and, consequently, improve the corresponding PL relation. At the same time, an important secondary goal is to obtain significantly improved W UMa stellar parameters, thus allowing us to better constrain the evolution of our cluster W UMa systems as a population. In Section 2, we discuss our observations and the calibration of both the NGC 188 data and the W UMa properties used in this study. The light-curve results, as well as the results from CMD fitting, proper-motion, and radial-velocity selection, and our distance analysis, are covered in Section 3. We discuss the properties of our eight sample W UMas, as well as the feasibility of using W UMas as distance indicators, e.g., based on their period--luminosity--color (PLC) or PL relations, in Section 4. In Section 5, we summarize our main conclusions.

Observations of the old OC NGC 188 were obtained with the recently commissioned 50BiN telescope in $V$ and $R$. We collected 36 nights of time-series data, spanning an unprecedented total of 3000 frames for each star. To select genuine cluster members, we performed a detailed radial-velocity and proper-motion analysis. The radial velocity of NGC 188 is $v_{\rm{RV}}=-42.32 \pm 0.90$ km s$^{-1}$, while its proper motion is ($\mu_{\alpha},\mu_{\delta}$) = ($-5.2 \pm 0.6, -0.3 \pm 0.6$) mas yr$^{-1}$. Of our total sample of 914 stars, 532 stars are probable cluster members. They delineate an obvious cluster sequence down to $V=18$ mag. We use the Dartmouth stellar evolutionary isochrones \citep{Dotter08} to match the cluster members, adopting an age of 6 Gyr and solar metallicity. A distance modulus and reddening of, respectively, $(m-M)_V^0=11.35\pm0.10$ mag and $E(V-R)=0.062\pm0.002$ mag were obtained. Accurate light-curve solutions were obtained for the eight W UMas, and parameters such as their mass ratios and the components' relative radii were estimated. We subsequently estimated the distance moduli for the W UMas, independent of the cluster distance. W UMas can be used to derive distance moduli with an accuracy of often significantly better than 0.2 mag. Object $V_5$ (V371 Cep) is not a genuine W UMa system. $V_3$ (ER Cep) was excluded from the distance-modulus analysis because of its low cluster-membership probability. For the remaining six OC W UMas---EP Cep, EQ Cep, ES Cep, V369 Cep, V370 Cep, and V782 Cep---we obtained a joint best-fitting distance modulus of $(m-M)_V^0=11.31 \pm 0.12$ mag, which is comparable to the result from our isochrone fits, as well as with previous results from the literature. The resulting accuracy is better than that resulting from application of the previously established empirical parametric approximation. To double check our results for NGC 188 and the applicability of W UMas as distance tracer, we applied it to the OC Berkeley 39. Based on four of its W UMas, we derived a distance modulus of $(m-M)_V^0=13.09\pm0.23$ mag, which is also in accordance with literature results. W UMas as distance tracers have significant advantages for poorly studied clusters. The six W Umas in NGC 188 satisfy a tight PL relation. Armed with the latter, W UMas could indeed play an important role in measuring distances and to map Galactic structures on more ambitious scales than done to date. Based on better distances and photometry, more accurate physical parameters for eight W UMa systems were derived. Using the evolutionary W UMa model of \citet{Yildiz13}, the initial masses of seven W UMa systems were estimated. All seven W UMa systems have initial primary masses below 1.8 $M_\odot$, which means that they have evolved along the evolutionar W-type route (i.e., through angular-momentum evolution within a convective envelope). The ages of seven sample W UMa systems were estimated based on their initial masses and the current luminosity-based masses of the secondary components. Six of our cluster W UMa systems have similar ages as the host cluster itself, which provides confirmation of the age-dating method of \citet{Yildiz14}.

1607.06152

1607

1607.03893_arXiv.txt

We report on the search for galaxies in the proximity of two very metal-poor gas clouds at $z\sim 3$ towards the quasar \qso. With a 5-hour MUSE integration in a $\sim 500\times 500$ kpc$^2$ region centred at the quasar position, we achieve a $\ge 80\%$ complete spectroscopic survey of continuum-detected galaxies with $m_{R} \le 25$ mag and Ly$\alpha$ emitters with luminosity $L_{\rm Ly\alpha} \ge 3\times 10^{41} ~\rm erg~s^{-1}$. We do not identify galaxies at the redshift of a $z\sim 3.2$ Lyman limit system (LLS) with $\log Z/Z_\odot = -3.35 \pm 0.05$, placing this gas cloud in the intergalactic medium or circumgalactic medium of a galaxy below our sensitivity limits. Conversely, we detect five Ly$\alpha$ emitters at the redshift of a pristine $z\sim 3.1$ LLS with $\log Z/Z_\odot \le -3.8$, while $\sim 0.4$ sources were expected given the $z\sim3$ \lya\ luminosity function. Both this high detection rate and the fact that at least three emitters appear aligned in projection with the LLS suggest that this pristine cloud is tracing a gas filament that is feeding one or multiple galaxies. Our observations uncover two different environments for metal-poor LLSs, implying a complex link between these absorbers and galaxy halos, which ongoing MUSE surveys will soon explore in detail. Moreover, in agreement with recent MUSE observations, we detected a $\sim 90~\rm kpc$ \lya\ nebula at the quasar redshift and three \lya\ emitters reminiscent of a ``dark galaxy'' population.

The assembly and growth of galaxies throughout cosmic times requires the continuous accretion of substantial amounts of fresh fuel. Observations of molecular or atomic gas in the interstellar medium of both high-redshift systems and nearby galaxies reveal that, without accretion rates that are at least commensurable with the observed star formation rates (SFRs), galaxies would exhaust their gas reservoirs on timescales of $\sim 1-2~\rm Gyr$, much shorter than the Hubble time \citep[e.g.,][]{san08,gen10}. From this consideration, it follows that the accretion of gas from the halo (i.e. the circumgalactic medium or CGM) and, ultimately, from the baryon reservoir present in the intergalactic medium (IGM), needs to be ubiquitous at all redshifts. In support of this argument, modern cosmological simulations predict accretions rates of $\gtrsim 10-20~\rm M_\odot~yr^{-1}$ at $z\sim 2-3$ galaxies \citep[e.g.,][]{dek09,ker09,fau11}. However, despite a general consensus that gas accretion is a dominant process for galaxy evolution, direct observational evidence of cold gas infalling onto galaxies is scarce at $z\sim 1$ \citep{rub12,mar12} and even more tenuous at higher redshifts \citep[e.g.,][]{cri13,bou13,mar15}. While in apparent contradiction with expectations of ubiquitous inflows, the lack of direct detections is often justified in the context of the so-called ``cold stream'' or ``cold flow'' paradigm if a significant fraction of the gas is accreted along narrow and dense filaments of cold gas \citep[e.g.,][]{bir03,ker05,dek06}. Indeed, if the infalling gas covers only a small fraction of the solid angle as seen from a galaxy \citep[e.g.,][]{goe12}, then it is natural that only a limited number of sightlines will intersect these infalling cold streams, showing redshifted absorption lines. Moreover, for typical infall velocities of $\sim 100-200~\rm km~s^{-1}$, the signature of inflows is often masked by interstellar absorption at the systemic redshifts \citep{rub12}. Further, when selecting absorbers via metal lines, some ambiguity remains in separating recycled gas falling back onto galaxies from material that is being accreted from the IGM for the first time. Due to these intrinsic limitations, observers have to resort to other, more indirect, signatures of the presence of cold gas accretion. For instance, simulations predict that accretion in the form of cold flows is a dominant contributor to the cross section of optically-thick gas ($N_{\rm HI} \ge 10^{17.2}~\rm cm^{-2}$) that gives rise to Lyman limit systems (LLSs) near galaxies \citep[e.g.,][]{fauker11,fum11sim,van12}. While cold streams occupy only a small fraction of the solid angle seen from a galaxy, the probability to intersect these filaments in the transverse direction with background sources is higher. Indeed, simulations predict covering factors $f_{\rm c}$ for optically-thick gas within the virial radius in the range of $f_{\rm c} \sim 0.1-0.4$, although with large variations between different models \citep{fauker11,fum11sim,shen13,fum14cf,fau14}. Thus, in principle, statistical comparisons between the observed properties of LLSs near galaxies \citep[e.g.,][]{rud12,pro13} and the predictions of numerical simulations offer interesting constraints for the cold accretion paradigm \citep[see also][]{leh13,fum14cf,coo15,fum16lls}. On an object by object basis, however, the lack of direct kinematic signatures of infall requires that multiple diagnostics are combined to establish whether the gas observed in absorption is potentially being accreted onto galaxies observed in emission at close projected separations. A few examples that rely on low metallicity \citep[e.g.,][]{rib11,cri13,leh13}, rotational signatures \citep[e.g.,][]{bou13,mar15}, or filamentary morphology \citep[e.g.,][]{can12} can be found in the literature. Following this approach, we present in this paper a dedicated search of galaxies around two very metal-poor LLSs with $Z \lesssim 5\times 10^{-4} Z_\odot$. Their extremely-low metallicity is at odds with what expected for gas that has been enriched by outflows. Thus, even without direct kinematic measurements, metal-poor LLSs that reside near galaxies are among the most compelling examples of nearly chemically-pristine gas infalling for the first time inside halos. It is however worth noting that very metal poor LLSs represent only a small fraction of the parent population, with $\lesssim 20\%$ of the LLSs having $Z \lesssim 10^{-3} Z_\odot$ between $z\sim 2.5$ and $3.5$ \citep{fum16lls,leh16}. Our observations target the field of the quasar \qso\ (a.k.a. SDSSJ095852.19$+$120245.0), which lies at a redshift $z_{\rm qso}=3.3088\pm0.0003$ \citep{hew10} and hosts two strong absorption line systems along its line of sight \citep{fum11sci}: a pristine gas cloud at $z_{\rm lls,1}=3.096221 \pm 0.000009$ with \ion{H}{I} column density\footnote{Throughout this work, \HI\ column densities are expressed in units of cm$^{-2}$.} $\log N_{\rm HI} = 17.18 \pm 0.04$ and without discernible metals to a limit of $Z<10^{-3.8}~\rm Z_\odot$; and a second LLS at $z_{\rm lls,2} = 3.223194 \pm 0.000002$, with column density $\log N_{\rm HI} = 17.36 \pm 0.05$ and metallicity $\log Z/Z_\odot = -3.35 \pm 0.05$ \citep{leh16}. A summary of the physical properties measured for these two systems or inferred via photoionization modelling by \citet{fum11sci} and \citet{leh16} is presented in Table \ref{tab:llsprop}. \begin{table*} \caption{Physical properties of the two LLSs, which are measured in absorption or inferred via photoionization modelling.}\label{tab:llsprop} \centering \begin{tabular}{c c c c c c c c} \hline \hline ID & $z_{\rm abs}$ & $\log N_{\rm HI}$ ($\rm cm^{-2}$) & $\log Z/Z_\odot$ & $\log n_{\rm H}$ ($\rm cm^{-3}$) & $\log x_{\rm HI}$ & $\log \ell$ (pc) \\ \hline LLS 1 & $3.096221 \pm 0.000009$ & $17.18 \pm 0.04$ & $<-3.8$ & $<-2.0$ &$<-2.4$ & $>3.1$ \\ LLS 2 & $3.223194 \pm 0.000002$ & $17.36 \pm 0.05$ & $-3.35 \pm 0.05$ & $-3.3$ &$-4.1$ & 6.3 \\ \hline \hline \end{tabular} \flushleft{The columns of the table are: (1) the system ID; (2) the redshift measured in absorption; (3) the neutral hydrogen column density measured in absorption; (4) the inferred metallicity; (5) the inferred neutral fraction; (6) the inferred size. Values are from \citet{fum11sci} and \citet{leh16}.} \end{table*} Details of the imaging and spectroscopic observations are presented in Section \ref{sec:obs}, followed by the analysis of continuum-detected sources and \lya\ emitters in Section \ref{sec:continuum} and Section \ref{sec:lines}. In Section \ref{sec:qso}, we report on the discovery of an extended nebula at the quasar redshift, with discussion and conclusions in Section \ref{sec:end}. Given the technical nature of Sections \ref{sec:obs}-\ref{sec:lines}, readers who are primarily interested in the final results may prefer to continue reading from Section \ref{sec:qso}. Throughout this work, we assume solar abundances from \citet{asp09} with $Z_\odot = 0.0134$, and we use the ``Planck 2013'' cosmology \citep{pla14} for which the Hubble constant is $H_0 = (67.8 \pm 0.8)~\rm km~s^{-1}~Mpc^{-1}$ and the matter density parameter is $\Omega_{\rm m} = 0.308 \pm 0.010$. Magnitudes are expressed in the AB system.

1607.03893

1607

1607.01772_arXiv.txt

% This study represents the most sensitive \Chandra\ X-ray point source catalogue of M31. Using 133 publicly available \Chandra\ ACIS-I/S observations totalling $\sim1$ Ms, we detected \num\ X-ray sources in the bulge, northeast, and southwest fields of M31, covering an area of $\approx$ \area\ deg$^{2}$, to a limiting unabsorbed $0.5-8.0$ keV luminosity of $\sim10^{34}$ \es. In the inner bulge, where exposure is approximately constant, X-ray fluxes represent average values because they were determined from many observations over a long period of time. Similarly, our catalogue is more complete in the bulge fields since monitoring allowed more transient sources to be detected. The catalogue was cross-correlated with a previous \xmmn\ catalogue of M31's $D_{25}$ isophote consisting of 1948 X-ray sources, with only 979 within the field of view of our survey. We found \xcmatch\ (\matchperall\%) of our \Chandra\ sources (\xcmatchuniq\ or \matchper\% unique sources) matched to within 5\arcsec\ of \xcmatchuniq\ \xmmn\ sources. Combining this result with matching done to previous \Chandra\ X-ray sources we detected \absnew\ new sources in our catalogue. We created X-ray luminosity functions (XLFs) in the soft ($0.5-2.0$ keV) and hard ($2.0-8.0$ keV) bands that are the most sensitive for any large galaxy based on our detection limits. Completeness-corrected XLFs show a break around \flatten\ \es, consistent with previous work. As in past surveys, we find the bulge XLFs are flatter than the disk, indicating a lack of bright high-mass X-ray binaries in the disk and an aging population of low-mass X-ray binaries in the bulge.

\label{sec:intro} M31 is the nearest large spiral galaxy to our own, and as such allows for the best possible spatial resolution and sensitivity of any Milky Way sized galaxy. The X-ray population of spiral galaxies can include X-ray binaries (XRBs), both low-mass (LMXBs) and high-mass (HMXBs), supernova remnants, supersoft sources, and massive stars. There is also contamination from background active galactic nuclei (AGN) or galaxies/galaxy clusters and Galactic foreground stars. Unlike the Milky Way, which is difficult to observe in X-rays due to absorption and the fact we are in the disc, M31 has moderate Galactic foreground absorption ($N_{\rm{H}} = 7\times10^{20}$ cm$^{-2}$ \citealt{dickey-90}) and can provide a galaxy-wide survey of the X-ray population, specifically XRBs. With an increased sample size, it is possible to study large-scale properties like the radial distribution across a range of $L_{X}$. The goal of this paper is to create the deepest \Chandra\ X-ray catalogue of regions with archival observations in M31. Specifically, the bulge has $>1$ Ms of data largely from monitoring programs that have not been fully exploited. The X-ray point source population of M31 was first studied by \citep{trinchieri11-91} using $\approx$ 300 ks of \Einstein\ imaging observations. They detected 108 sources of which 16 showed variability. They did not find a significant difference between the luminosity distribution of the bulge and disc population. \citet{primini06-93} completed a survey of the central $\sim$34\arcmin\ of M31 using the \rosat\ High-Resolution Imager (HRI). They found 18 variable sources within $7.5\arcmin$ of the nucleus and 3 probable transients. Also, $>75\%$ of the unresolved X-ray emission in the bulge was either thought to be diffuse of from a new class of X-ray sources. This work was followed-up by deeper \rosat\ observations \citep{supper01-97,supper07-01} detecting 560 sources down to $5\times10^{35}$ \es\ in the $\sim10.7$ deg$^{2}$ view. They associated 55 sources with foreground stars, 33 with globular clusters, 16 with supernova remnants, and 10 with radio sources and galaxies, leaving 80\% of sources without an optical/radio identification. They confirmed the previous result from \Einstein\ that the total luminosity is distributed evenly between the bulge and disc. A comparison with the \Einstein\ results revealed 11 variable sources and 7 transients, while comparison with the first \rosat\ Position Sensitive Proportional Counter (PSPC) survey found 34 variable sources and 8 transients. The \rosat\ surveys also revealed the presence of supersoft sources (a class of white dwarf X-ray binary) in M31 \citep{supper01-97,kahabka04-99}. \citet{trudolyubov12-04} used \xmmn\ and \Chandra\ to detect 43 X-ray sources coincident with globular cluster candidates, finding 31 of the brightest had spectral properties similar to Galactic LMXBs. X-ray monitoring of optical novae in the centre of M31 with \rosat, \xmmn, and \Chandra\ \citep{pietsch11-05,pietsch04-07} showed them to be primarily supersoft sources. \citet{shaw-greening03-09} completed an \xmmn\ spectral survey of 5 fields along the major axis of M31 (excluding the bulge) and detected 335 X-ray sources, which were correlated with earlier X-ray surveys and radio, optical, and infrared catalogues. They classified 18 sources as HMXB candidates by spectral fitting with a power law model with a photon index of $0.8-1.2$, indicating they were magnetically accreting neutron stars. \citet{peacock10-10} used the 2XMMi catalogue of X-ray point sources along with supplemental \Chandra\ and \rosat\ observations to identify 45 globular cluster LMXBs. This study covered 80\% of the known globular clusters in M31 \citep{peacock07-10} and confirmed trends whereby high metallicity, luminosity, and stellar collision rate correlated positively with the likelihood of a cluster hosting an LMXB. \citet{henze03-14} completed monitoring observations with \xmmn\ and \Chandra\ of the bulge of M31 and detected 17 new X-ray counterparts of optical novae, with 24 detected in total. \begin{table*} \caption{Summary of Previous M31 X-ray Surveys\label{tab:sum}} \resizebox{\textwidth}{!}{% \begin{tabular}{ c c c c c} \hline\hline Observatory & Detected Sources & $L_{X}$ (\es) & Region & References \\ \hline \Einstein & 108 & $5\times10^{36} - 10^{38}$ ($0.2-4.0$ keV) & 14 \Einstein\ imaging observations ($\sim4$ deg$^2$) & \citet{trinchieri11-91} \\ \rosat\ (HRI) & 86 & $\gtrsim1.8\times10^{36}$ ($0.2-4.0$ keV) & central $\sim$34\arcmin\ ($\sim0.3$ deg$^2$) & \citet{primini06-93}\\ \rosat\ (PSPC) & 560 & $5\times10^{35} - 5.5\times10^{38}$ ($0.1-2.4$ keV) & whole galaxy ($>D_{25}$ ellipse, $10.7$ deg$^2$) & \citet{supper01-97,supper07-01}\\ \xmmn/\Chandra & 43 & $\sim10^{35} - 10^{39}$ ($0.3-10.0$ keV) & bulge \& major axis ($1.7$ deg$^{2}$) & \citet{trudolyubov12-04}\\ % \xmmn & 335 & $\sim10^{34} - 10^{39}$ ($0.3-10.0$ keV) & 5 fields along major axis ($1$ deg$^{2}$) & \citet{shaw-greening03-09}\\ \xmmn/\Chandra & 45 & $\sim10^{35} - 7\times10^{38}$ ($0.2-12.0$ keV) & whole galaxy ($>D_{25}$ ellipse, $4$ deg$^{2}$) & \citet{peacock10-10}\\ \xmmn & 1897$^{1}$ & $4.4\times10^{34} - 2.7\times10^{38}$ ($0.2-4.5$ keV) & whole galaxy ($>D_{25}$ ellipse, $4$ deg$^{2}$) & \citet{stiele10-11}\\ \xmmn/\Chandra\ & 24 & $\sim10^{35} - 9\times10^{37}$ ($0.2-2.0$ keV) & centre ($\sim0.2$ deg$^{2}$) & \citet{henze03-14}\\ % \hline \end{tabular} } \begin{list}{}{} \item Luminosities have all been corrected to a distance of 776 kpc used in this paper. \Chandra\ surveys are summarized in Table \ref{tab:chandra}. See \citet{stiele10-11} for a more comprehensive list. \item ${^1}$ The \xmmn\ LP total catalogue includes 1948 X-ray sources. \end{list} \end{table*} The most comprehensive X-ray population survey of M31 to date was completed by \citet{stiele10-11} using the \xmmn\ European Photon Imaging Camera. They detected 1897 sources to a limiting luminosity of $4.4\times10^{34}$ \es, including 914 new X-ray sources. Their source classification/identification was based on several methods: X-ray hardness ratios, spatial extent of the sources, long-term X-ray variability, and cross-correlation with X-ray, optical, infrared, and radio catalogues. Confirmed identifications included 25 supernova remnants, 46 LMXBs, 40 foreground stars, and 15 AGN/galaxies. There were many candidates for each of these classes as well, including 2 HMXBs and 30 supersoft sources. Nevertheless, 65\% of their sources had no classification. We summarize a few of the major X-ray surveys of M31 in Table \ref{tab:sum} (for a more detailed list please see \citet{stiele10-11}). \Chandra\ has not observed all of M31 as previous observatories have, but instead mostly monitored the supermassive black hole in the nucleus, with the majority of exposures each being 5 ks. Various groups have used a handful of observations to survey the bulge and create a catalogue of sources with either the Advanced CCD Imaging Spectrometer (ACIS-I/S) or the High-Resolution Camera (HRC-I/S). \citet{kong10-02} compiled the first \Chandra\ catalogue of M31 within the bulge, finding 204 sources above $\gtrsim2\times10^{35}$ \es. Their most important result was finding different X-ray luminosity functions (XLFs) when different regions (inner/outer bulge and disc) were considered separately. The inner bulge showed a break at $10^{36}$ \es\ and this break shifted to higher luminosities when moving outwards from the inner bulge to the disc. In addition, the slopes became steeper, indicating non-uniform star formation history. \citet{kaaret10-02} used the HRC-I to detect 142 sources, which when compared to \rosat\ observations revealed 50\% of the sources $>5\times10^{36}$ \es\ to be variable. No evidence was found for X-ray pulsars, leading to the conclusion that most sources should be LMXBs. \citet{di-stefano05-02} used 3 disc fields in M31 to analyse globular cluster LMXBs while \citet{di-stefano07-042} used these fields with a nuclear pointing to study supersoft/quasi-soft sources. \citet{williams07-04} used HRC-I to study the disc and bulge of M31 with snapshot images, finding variability in 25\% of 166 detected sources. \citet{voss06-07} combined 26 ACIS observations to investigate the X-ray population in the bulge, finding 263 X-ray sources (64 new) down to $10^{35}$ \es. They clearly demonstrated the power of merging observations to obtain deeper exposures and detect the faintest sources, decreasing the (completeness-corrected) XLF limit by a factor of 3. \citet{hofmann07-13} used 64 HRC-I observations totalling 1 Ms to detect 318 X-ray sources. They studied the long-term variability of sources by producing light curves and found 28 new sources, along with classifying 115 as candidate XRBs. A further 14 globular cluster XRB candidates, several new nova candidates, and a new supersoft X-ray source outburst were discovered. Recently, the $\sim$12 yrs of monitoring observations of the nucleus have been utilized in a number of studies to investigate variability and detect transients \citep{barnard09-12,barnard092-12,barnard06-13}. Specifically, \citet{barnard01-14} used 174 \Chandra\ ACIS and HRC observations to detect 528 X-ray sources in the bulge down to $10^{35}$ \es. By studying source variability, they identified 250 XRBs (200 new) with X-ray data alone, a factor of 4 increase. Table \ref{tab:chandra} summarizes previous \Chandra\ M31 X-ray catalogues. At the time of writing, a large \Chandra\ program (350 ks) has been accepted to survey a part of the star-forming disc of M31 (PI: B. Williams). Aside from being able to confirm the first HMXBs in M31, it will completely characterize a large part of the X-ray source population using optical photometry and spectroscopy. \begin{table*} \caption{Summary of Previous \Chandra\ M31 X-ray Catalogues\label{tab:chandra}} \begin{tabular}{c c c c c} \hline\hline Instrument & Detected Sources & $L_{X}$ (\es) & Region & References \\ \hline ACIS -I & 204 & $\gtrsim2\times10^{35}$ ($0.3-7.0$ keV) & central $\sim17\arcmin\times17\arcmin$ (0.08 deg$^{2}$) & \citet{kong10-02} \\ HRC-I & 142 & $2\times10^{35} - 2\times10^{38}$ ($0.1-10.0$ keV) & central $\sim30\arcmin\times30\arcmin$ (0.25 deg$^{2}$) & \citet{kaaret10-02}\\ ACIS -I/S \& HRC-I & 28 & $5\times10^{35} - 3\times10^{38}$ ($0.3-7.0$ keV) & 3 disc fields (0.7 deg$^{2}$) & \citet{di-stefano05-02}\\ ACIS -S3 & 33 & $\sim10^{35} - 10^{38}$ ($0.1-7.0$ keV) & 3 disc fields and nucleus (0.7 deg$^{2}$) & \citet{di-stefano07-042}\\ HRC-I & 166 & $\sim10^{36} - 5\times10^{38}$ ($0.1-10.0$ keV) & disc \& bulge (0.9 deg$^{2}$) & \citet{williams07-04}\\ ACIS -I/S & 263 & $5\times10^{33} - 2\times10^{38}$ ($0.5-8.0$ keV) & $12$\arcmin\ radius from core (0.126 deg$^{2}$) & \citet{voss06-07}\\ HRC-I & 318 & N/A ($0.1-10.0$ keV) & $30$\arcmin\ radius from core ($0.8$ deg$^{2}$) & \citet{hofmann07-13}\\ ACIS -I/S \& HRC-I & 528 & $5\times10^{34} - 5\times10^{38}$ ($0.3-10.0$ keV) & $20$\arcmin\ radius from core ($0.35$ deg$^{2}$) & \citet{barnard01-14}\\ \hline \end{tabular} \begin{list}{}{} \item Luminosities have all been corrected to a distance of 776 kpc used in this paper. \end{list} \end{table*} This paper aims to study the properties of M31 X-ray point sources down to the lowest luminosities for any large galaxy. Specifically, we will report general source catalogue characteristics (e.g. flux and radial distributions), cross-correlate our catalogue with previous \xmmn\ and \Chandra\ surveys, and study the XLF. In addition, high-resolution \HST\ data from the Panchromatic Hubble Andromeda Treasury (PHAT) Survey \citep{dalcanton06-12} allows optical counterpart identifications for X-ray sources, especially AGN. We adopt a distance to M31 of 776 $\pm$ 18 kpc as in \citet{dalcanton06-12}, which corresponds to a linear scale of 3.8 pc arcsecond$^{-1}$.

\label{sec:summary} We have used 133 publicly available \Chandra\ ACIS-I/S observations totalling $\sim1$ Ms to create the deepest X-ray point source catalogue of M31. We detected \num\ X-ray sources within our field of view (\area\ deg$^{2}$) to a limiting unabsorbed $0.5-8.0$ keV luminosity of $\sim10^{34}$ \es. Our 90\% ($0.3-8.0$ keV) completeness limit is $4\times10^{35}$ \es. We detected \blg, \nes, and \sw\ sources in the bulge, northeast, and southwest fields of M31 respectively. In the bulge fields, X-ray fluxes are closer to average values because they are calculated from many observations over a long period of time. Similarly, our catalogue is more complete in the bulge fields since monitoring allows more transient sources to be detected. Cross-correlating our catalogue with a previous \xmmn\ catalogue of 1948 X-ray sources, with only 979 within the field of view of our survey area, we found \xcmatch\ (\matchperall\%) of our \Chandra\ sources (\xcmatchuniq\ or \matchper\% were unique sources) matched to within 5\arcsec\ of \xcmatchuniq\ \xmmn\ sources with a median offset of \medsepx\arcsec. Similarly, we matched our catalogue to a master list of \allprevc\ previously published \Chandra\ sources in M31 and found \catc\ of our sources within 1\arcsec. Collating the matching results from all catalogues we found \absnew\ new sources in our catalogue. We also created XLFs in the soft and hard bands that are the deepest for any large galaxy based on our detection limits. Using published catalogues of AGN and LMXBs we determined the contribution to the XLF from these populations. The observationally identified AGN in M31 are incomplete below $\sim10^{-13}$ \esc\ (hard band) based on data from the \Chandra\ Deep Field South. The completeness-corrected XLFs show a break at \flatten\ \es, which is consistent with previous work in M31. We found that the bulge XLFs are flatter compared to the disk, consistent with other studies. This indicates a lack of bright high-mass X-ray binaries in the disk due to a low star formation rate and an aging population of low-mass X-ray binaries in the bulge. This catalogue is more robust and complete than the latest \Chandra\ source catalogue release due to the stringent processing and requirements we have placed on source detection. In addition, we have published a much more detailed set of source characteristics using {\em ACIS Extract}. Impending \Chandra\ and \nustar\ X-ray surveys in M31 will cover new regions of the galaxy (e.g. PHAT field) that have only been observed by \xmmn. This will result in high spatial resolution $0.3-30$ keV data that will be crucial for classifying and characterising the X-ray source population.

1607.01772

1607

1607.05639.txt

% The next generation interferometric radio telescope, the Square Kilometre Array (SKA), which will be the most sensitive and largest radio telescope ever constructed, could greatly contribute to the detection, survey and characterization of Gamma Ray Bursts (GRBs). By the SKA, it will be possible to perform the follow up of GRBs even for several months. This approach would be extremely useful to extend the Spectrum Energetic Distribution (SED) from the gamma to the to radio band and would increase the number of radio detectable GRBs. In principle, the SKA could help to understand the physics of GRBs by setting constraints on theoretical models. This goal could be achieved by taking into account multiple observations at different wavelengths in order to obtain a deeper insight of the sources. Here, we present an estimation of GRB radio detections, showing that the GRBs can really be observed by the SKA. The approach that we present consists in determining blind detection rates derived by a very large sample consisting of merging several GRB catalogues observed by current missions as \textit{Swift}, \textit{Fermi}, Agile and INTEGRAL and by previous missions as BeppoSAX, CGRO, GRANAT, HETE-2, Ulysses and Wind. The final catalogue counts 7516 distinct sources. We compute the fraction of GRBs that could be observed by the SKA at high and low frequencies, above its observable sky. Considering the planned SKA sensitivity and through an extrapolation based on previous works and observations, we deduce the minimum fluence in the range 15\,-\,150 keV. This is the energy interval where a GRB should emit to be detectable in the radio band by the SKA. Results seem consistent with observational capabilities.

\label{Introduction} % Due to their high fluence, between $10^{-7}$ and $10^{-5}$ erg/cm$^2$, and to their huge isotropic energies, between $\sim 10^{48}$ - $10^{54}$ erg emitted in a very short time, Gamma Ray Bursts (GRBs) are the most violent and energetic astrophysical phenomena currently known in the Universe. \\ Discovered in 1963 but announced only in the 1974 \citep{annuncioGRB}, the study of GRB is continuing and improving. Although GRBs are the brightest sources in the Universe and studied for long time, there are several unclear aspects that has to be understood. For example, several models try to explain the physics of these phenomena, but none can be considered final and self-consistent. One of the principal issues is that GRB spectra are very different from each other and this fact makes them very difficult to study and classify: in general, different spectra correspond to various combinations of parameters. So far, the central engine of these sources is still debated (see, e.g., \cite{FWK00,Da03,Mes06}), because there is no single theoretical model capable of explaining in a comprehensive way all the observations. \\ In general, they are cosmological objects and may occur at any point of the Universe, namely at different redshift, and, for this reason, strong selection effects plays an important rule in the discussions about the physics of GRBs. Currently, the farthest GRBs have been detected at $z = 8.1$ \citep{Salvaterra2009} (spectroscopic redshift) and $z = 9.4$ \citep{GRBrecord} (photometric redshift). Using them as distance indicators is crucial issue in order to probe the Hubble flow up to early epochs. \par % In general, GRBs are peculiar sources and because of the variety of light curves and spectra, many efforts are concentrating on a possible standardization. These attempts are principally focused on high energies (i.e., $\gamma$\,-\,rays, X\,-\,rays), even though some recent studies are extending the frequency range towards the low energies. The general aim is to discover a standard behavior for a specific class of GRBs presenting defined characteristics and features. \\ In other words, it is necessary to observe GRBs in several spectral bands and consider their complete envelopes in a large energy range. Studying in detail a large number of Spectral Energy Distributions (SEDs) could allow to highlight the emission process in a large spectral band. In this sense, a crucial breakthrough has been achieved by the launch of the {\it Swift} satellite in 2004. The \textit{Burst Alert Telescope} (BAT) in 15\,-\,150 keV energy range, the \textit{X\,-\,Ray Telescope} (XRT) in 0.3\,-\,10 keV and the \textit{Ultra-Violet/Optical Telescope} (UVOT) in 170\,-\,650 nm make up the payload of \textit{Swift} and allow a rapid follow-up of the afterglows in different wavelengths. These facilities give a better coverage of the GRB light curve than the previous satellite missions. By the instrumentation for the X and UV counterparts, \textit{Swift} allows a rapid localization of GRBs and several efforts have been dedicated to trace afterglow light curves at different wavelengths \citep{N06,K11,O09}.\\ % By using large catalogues, comparative studies among optical and X\,-\,ray light curves allow to fix constraints on GRB theoretical models (see, e.g., \cite{O09}, \cite{S11} for the standard fireball model) and to carry out correlations among optical and X\,-\,ray properties, e.g. between fluence and brightness \citep{G08, N09, K11}. \par Regarding to the radio band, the GRB light curve could be tracked for hundreds of days after the $\gamma$ \,-\,ray onset, but fluxes are very faint at these frequencies and only a few of current instruments can detect them. This paper, without claiming for completeness, is a discussion on interferometric radio observations since these could improve the detection, the surveys and the characterization of GRBs at frequencies which are not usually investigated. Our aim is to show that these observations are a realistic option. In general, the analysis of the GRB afterglow light curves, at different wavelengths, allows to investigate the physics of these fascinating objects probing how the blast wave generated by the burst propagates in the circumburst medium. At the present state of the art, it is crucial to boost the multi-wavelength study of the afterglow by radio observations and, in particular, by the Square Kilometre Array (\href{http://www.skatelescope.org/}{SKA}\footnote{http://www.skatelescope.org/}). \\ Although satellite observations are the primary step for the GRB detection, they are not and cannot be the only observational channel. Satellite missions are very expensive, have limited lifetimes and something could go amiss (e.g., malfunctioning, mission delays, wrong orbiting, etc). This means that satellite GRB investigations should be supported with ground-based telescopes, promoting, finally, research towards radio frequencies. It is worth noticing that radio band is not affected by radiation extinction, contrary to higher frequencies. An accurate calorimetry in the radio band for well-detectable GRBs would be possible in principle. In addition, radio observations can give substantial indications of the inverse Compton effect, since only radio frequencies can probe the density of the interstellar medium. In other words, redshift measurements of the host galaxies can be obtained observing the hydrogen spin-flip at 1.4 GHz. These observations give fundamental information for GRBs guested in galaxies. \par % The layout paper is as follows: in Sec. \ref{importance of radio} we discuss GRBs in radio band; in Sec. \ref{GRBs above SKA} the SKA telescope is shortly introduced, considering, in particular, its instrumental sensitivities. Specifically, in subsections (Sec. \ref{Collection of GRB catalogues}, \ref{sec: The sky above the SKA and its shadow cone}, \ref{Sec: A suggestion for a Spectral Energy Distribution from radio to gamma band}, \ref{Sec: The SKA sensitivity and the radio GRB-detection probability}, \ref{sec: Serendipitous detection rate probability for the SKA}), we calculate the minimum fluence between 15\,-\,150 keV to detect a GRB by the SKA introducing a first simple SED from gamma range to radio band, give some estimations of the GRB detection rates for the SKA at high and low frequencies. We use a catalogue assembled by collecting several previous catalogues and obtain a large list with 7516 GRBs. Discussions and conclusions are given in Sec. \ref{sec: Discussions and conclusions}. %% %%

\label{sec: Discussions and conclusions} %%% %%% In this paper, we discussed the possibility use the SKA for GRB radio detections and surveys. In the last part of the previous section, we calculated some serendipitous detection rates. However, to detect GRBs in radio band, it is worth noticing that SKA dishes will be used principally to go on-source with specific (and precise) coordinates, namely after a GRB alert. In other words, GRBs will not be detected randomly, but moving antennas on-source after a satellite previous detection (e.g., Robo-AMI, mentioned above). In our opinion, only considering a comprehensive synergy among the radio band detections and the detections in other wavelengths ($\gamma$-ray, X-ray, optical, infrared) will be possible to really understand the light curves of these phenomena. In fact, the SKA will considerably be able to contribute to an increment of GRB radio data, so that the statistics will be improved, and finally the understanding of the GRB physics will be enhanced. \\ On the other hand, a serendipitous detection by the SKA could give an important contribution to the development of GRB science. In fact, another point to highlight concerns \textit{prompt} radio emissions. Even though these emissions might not exist, it is also important to point out that they have never been directly observed by a radio telescope. Currently, this very short emission might be detected randomly only, because moving an antenna on-source takes some minutes (or seconds, in the best case). In collaborations like Robo-AMI and \Swift, after the satellite detections, the radio antennas go on-source in $\sim 4$ minutes but this time is too slow for a radio observation in the GRB \textit{prompt} phase. \par It is important to stress that GRBs are peculiar sources and basically emit covering all the spectrum. The principal investigations are limited around high frequencies (i.e., $\gamma$, X, optical) and often authors try to point out any single aspect of them, to classify these objects in several classes. However, because of the very large emission spectrum, GRBs should be studied with a different approach, that is without focusing only on a narrow range of the electromagnetic spectrum. Analyzing GRBs with a large ``spectral field of view'' would give the chance to see aspects related to each other among more bands, perhaps discovering common features which are not clear if one sees only specific details in narrow electromagnetic bands. GRB studies in radio band could become crucial for discoveries in this sense, therefore observations by appropriate telescopes are necessary. Additional radio data will give the possibility to find possible correlations either within the radio band or between the radio frequencies and the higher ones. Without a lot of radio observations and focused analyses, it is impossible to get precise and self-consistent observations. In particular, we cannot say that every GRB with a fluence greater than an estimated value can emit in radio band (because only observations can confirm or not this issue), but the results shown here about the radio detection rates and GRB radio sensitivities are relevant enough to encourage the study in this direction by the SKA. \par %%% Another important issue is related to cosmology. As it is well known, GRBs are cosmological sources and their studies can be addressed in a cosmological perspective. The most ambitious goal is demonstrating that these objects could be used as \textit{distance indicators}. In order to find intrinsic relations amomg GRB parameters that effectively translates into distances, it is worth noticing that only observational quantities should be used. However, obtaining a standard light curve for GRBs, as Phillips achieved for SNeIa \citep{PhillipsRelation}, is extremely more difficult. Indeed, GRBs spread in a very large range of variables that could be used to fix their fundamental features. If the redshift is known, it is possible to obtain a distance for GRBs. The SKA can help in this goal considering also observations not directly related to GRBs (e.g., the spin-flip emission of the hydrogen at 1.4 GHz from the host galaxy). \par To conclude, precise radio observations and surveys would allow a deeper and complete understanding of these still mysterious objects. The expansion of the SED for GRBs would help to relate different bands to each other and a focused study also in radio frequencies would open a research perspective which has not been explored yet in detail. In this work we have shown that this opportunity is a fact since the SKA will be able to observe and detect GRBs. It is also useful to point out that the radio band is not affected by radiation extinction, contrary to higher frequencies. In general, an accurate calorimetry for radio well-detectable GRBs would be possible. In addition, radio observations can be useful for estimations of the inverse Compton scattering, since only radio frequencies can probe the density of the interstellar medium. However, it is imperative to keep in mind that synergies among ground-based and satellite telescopes, observing at different frequencies, have to be realized in order to understand the fundamental nature of GRBs. %%% %%% % % %% Acknowledgements

1607.05639

1607

1607.00774_arXiv.txt

We present the discovery of three new transiting giant planets, first detected with the WASP telescopes, and establish their planetary nature with follow up spectroscopy and ground-based photometric lightcurves. WASP-92 is an F7 star, with a moderately inflated planet orbiting with a period of 2.17 days, which has $R_p = 1.461 \pm 0.077 R_{\rm J}$ and $M_p = 0.805 \pm 0.068 M_{\rm J}$. WASP-93b orbits its F4 host star every 2.73 days and has $R_p = 1.597 \pm 0.077 R_{\rm J}$ and $M_p = 1.47 \pm 0.029 M_{\rm J}$. WASP-118b also has a hot host star (F6) and is moderately inflated, where $R_p = 1.440 \pm 0.036 R_{\rm J}$ and $M_p = 0.513 \pm 0.041 M_{\rm J}$ and the planet has an orbital period of 4.05 days. They are bright targets (V = 13.18, 10.97 and 11.07 respectively) ideal for further characterisation work, particularly WASP-118b, which is being observed by K2 as part of campaign 8. WASP-93b is expected to be tidally migrating outwards, which is divergent from the tidal behaviour of the majority of hot Jupiters discovered.

The WASP consortium \citep{pollacco2006} has been highly successful at identifying and confirming the existence of transiting hot Jupiters (orbital period, P$<$10 days). Wide angle ground-based surveys are able to find rarer objects with small orbital separations and deep transits found predominantly around stars brighter than V=13. This wealth of discoveries of bright stars hosting planets has enabled detailed follow up observations. The planets discovered by these surveys dominate the targets used for planetary atmosphere characterisation and further understanding of planet formation mechanisms and migration e.g. \citet{sing2016,ford2006,matsumura2010}. Hot Jupiters are rare objects with an occurrence rate estimated to be $\sim1\%$, depending on the stellar population surveyed \citep{wang2015,howard2012,wright2012}. Searches for close in companions to hot Jupiters via transit timing variations (TTV; \citet{steffen2012}) have been unsuccessful. This lack of close in planets points towards the theory that these objects form outside the snow line and migrate towards the orbits they are observed in through high eccentricity migration (HEM) \citep{rasio1996,fabrycky2007,mustill2015}. It was first noted by \citet{winn2010} and expanded upon by \citet{albrecht2012} that the distribution of projected hot Jupiter orbital obliquities appeared correlated with the temperature of the host star. The two groups of obliquities seen were separated at a stellar effective temperature around $6250$ K, with cooler stars more likely to host hot Jupiters with orbits aligned to the stellar rotation axis, and hotter stars with a range of alignments. The range of alignments observed is another indicator that HEM is a strong candidate for the prevailing migration mechanism for hot Jupiters with strong misalignments. Recent statistical work has however shown that not all hot Jupiters can have undergone HEM due to the lack of super-eccentric orbits found \citep{dawson2015}. The K2 mission \citep{howell2014} allows for the detection of the transits of small companion planets for hot Jupiters. This capability was demonstrated with the discovery of 2 further planets in close in orbits around WASP-47 \citep{becker2015}, which leads to the question of whether close in planets can be detected with space-based photometry for other systems with hot Jupiters. In this paper we present the discovery of WASP-92b, WASP-93b and WASP-118b. The latter system will be observed in campaign 8 of the K2 mission. All three of these planetary systems are excellent candidates for spin-orbit alignment follow up. Section \ref{sec:obs} introduces the observational data collected for the systems in this paper. Section \ref{sec:stellar} presents the spectral analyses of the host stars, and Section \ref{sec:analy} describes the methods used to determine the parameters of the newly discovered systems. The discoveries and their tidal evolution are discussed in Section \ref{sec:discu}.

\label{sec:discu} \begin{figure} \includegraphics[width=0.48\textwidth]{evo} \caption{Results of the \textsc{Bagemass} MCMC analysis for WASP-92 (upper plot); WASP-93 (middle), and WASP-118 (lower). For each of the plots, the dotted black line is the ZAMS. The solid blue line is the evolutionary track for the mass found, and the dashed tracks either side are for the 1-$\sigma$ error of the mass. The solid orange line is the isochrone for the stellar age found, with the 1-$\sigma$ error denoted by dashed lines in the same colour. The density of MCMC samples is shown in the colour scale of the posterior distribution plotted.} \label{fig:evo} \end{figure} \subsection{WASP-92 system} WASP-92b is a {\bf$\sim0.81$} $M_{\rm J}$ planet in a 2.17 days orbit around an F7 spectral type star. The \textsc{Bagemass} tool presented in \citet{maxted2015} was used with the observed measurements for {[Fe/H]}, \teff and $\rho_*$ to estimate the age and mass of WASP-92, which were found to be $2.99\pm1.03$ Gyr and $1.21\pm0.06$ $M_\odot$. The posterior distributions of this analysis can be seen in the upper plot of Figure \ref{fig:evo} with the associated stellar evolution tracks and isochrones for the optimal mass and age found by the \textsc{Bagemass} tool. The tracks used by the tool were calculated from the \textsc{GARSTEC} code \citep{weiss2008}. These results are in agreement with the gyrochronological age presented in Section \ref{sec:stellar92}, but much older than the age suggested by the lithium abundance observed in the spectra. When the solution for the system was allowed to include eccentricity, the value for $e$ found was $0.084^{+0.118}_{-0.060}$, which affected each of the reported parameters by $<$5\%. The BIC for the eccentric solution was 19.1, and 5.7 for the circular solution. Given the low value for eccentricity found and a $\Delta$BIC\footnotemark \footnotetext{where $\Delta$BIC = BIC$_{ecc}$ - BIC$_{circ}$} $>$10, which \citet{kass1995} indicates is very strong evidence against additional free parameters, the orbit of WASP-92b can be assumed to be circular. The parameters found in the fit for a circular orbit were used as the final parameters. \subsubsection{Tidal evolution} Using the parameters of the orbital solution, the tidal stability of the system can be investigated. The majority of hot Jupiters observed are in Darwin-unstable orbits \citep{darwin1879}, where the planet is migrating towards the Roche limit of the system and tidal disruption of the planet \citep{matsumura2010}. No stable orbits exist if the total angular momentum of the system, $L_{tot}$ is below a critical angular momentum, $L_{c}$, \begin{equation} L_{c} = 4 \left( \frac{G^{2}}{27} \frac{M_{*}^{3}M_{p}^{3}}{M_{*}+M_{p}} \left( C_{*} + C_{p}\right) \right)^{\frac{1}{4}} \,\, , \end{equation} where $C_{*}$ and $C_{p}$ are the moments of inertia of the star and planet respectively \citep{counselman1973,hut1980}. The total angular momentum is defined as \begin{equation} L_{tot} = L_{orb} + C_{*}\omega_{*} + C_{p}\omega_{p} \, \, , \end{equation} where $L_{orb} = M_{*}M_{p}\sqrt{ \frac{Ga(1-e^{2})}{M_{*} + M_{p}} }$. For WASP-92, $L_{tot} / L_{c} \sim 0.67$, which indicates that the planetary orbit is unstable and is continuing to migrate inwards towards the Roche limit, which is the case for most hot Jupiters. The rate of this migration can be estimated by, \begin{equation} t_{\textrm{remain}} = \frac{2 Q'_{*,0}}{117n} \frac{M_{*}}{M_{p}} \left( \frac{a}{R_{*}} \right)^{5} \,\, , \end{equation} as presented in \citet{brown2011} for slowly rotating stars, where $Q'_{*,0}$ is the current tidal quality factor for the star and $n$ is the orbital frequency. If $Q'_{*,0}$ is set to $10^8$\footnotemark \footnotetext{The use of $Q'_{*,0} = 10^{8}$ is suggested in \citet{penev2011}.}, the spiral in time ($t_{\textrm{remain}}$) is approximately 16 Gyr, which is significantly longer than the remaining lifetime of the host star. For lower values of $Q'_{*}$, the decay timescale remains too large for changes in the orbital period to be observable - a period change of $>60$s will occur after $\sim 20$ Myr for $Q'_{*,0} = 10^{8}$ when solely accounting for the effects of tidal orbital decay. \subsection{WASP-93 system} WASP-93b is a {\bf$\sim1.47$} $M_{\rm J}$ planet in a 2.73 days orbit around an F4 spectral type star. \textsc{Bagemass} was used for the WASP-93 system, resulting in age and mass estimates of $1.61\pm0.48$ Gyr and $1.39\pm0.08$ $M_\odot$. This age estimate is just beyond the upper 1-$\sigma$ uncertainty for the age derived from gyrochronology, and below the estimate of several Gyr from the lithium abundance. When the eccentricity parameters were included in the orbital solution, the value for $e$ found was $0.012^{+0.035}_{-0.008}$, and eccentric BIC was 168.3 and circular BIC was 147.4. The orbit is assumed to be circular, as the value for $e$ is small, and the $\Delta$BIC$>$10. \begin{figure} \includegraphics[width=0.49\textwidth]{tomographywasp93} \includegraphics[width=0.49\textwidth]{w93_rvexpanded} \caption{Upper plot showing the time series of CCF residuals after the subtraction of the average CCF profile with reference to transit phase for each of the two attempts to observe spectra during transit. The vertical dashed lines show the width of the CCF profile centred on the centre-of-mass velocity of the system. The horizontal dashed lines show beginning of transit ingress and egress. The trail of the transit of WASP-93b is visible near the blue-shifted limb of the CCF. Lower plot showing the RV measurements around the transit phase. The colours used correspond to those in the lower plot of Figure \ref{fig:wasp93}, and the green and blue RV points match the CCFs shown in the left and right of the upper plot respectively. The RV and RM effect model overplotted in black is for the case that the planetary orbit and the stellar spin are aligned.} \label{fig:w93tomography} \end{figure} In Section \ref{sec:spec} it was mentioned that two attempts were made to collect a time series of spectra during transits of WASP-93b. Due to uncertainty in the transit ephemeris at the time of these observations, the measurements are not well centred on the mid-transit time. The CCFs for these observations were reduced as described in \citet{cameron2010}, and included in the fit, which now includes the orbital obliquity and FWHM of the planet signal in the CCF as jump parameters. The MCMC code did not converge on a solution when including the CCFs to determine orbital obliquity, which indicates that the SOPHIE spectra were not of a high enough precision to determine the misalignment of the transit of WASP-93b. A value for the orbital obliquity was also not found when the $v \sin i$ was fixed to the value found in the global fit presented in Table \ref{tab:parameters}, $37.0$ km s$^{-1}$. The maximum amplitude of the RV anomaly produced by the RM effect is defined as, \begin{equation} \Delta RV_{RM} = \sqrt{1 - b^{2}}\left( \frac{R_{p}}{R_{\star}} \right) ^{2} v \sin i \,\, , \end{equation} which equates to $\sim 0.18$ km s$^{-1}$ for WASP-93. The average 1-$\sigma$ uncertainty for the first series of spectra is 0.13 km s$^{-1}$ and 0.11 km s$^{-1}$ for the second, which will make the effect difficult to detect above the noise in the CCFs. The lower plot in Figure \ref{fig:w93tomography} shows the RV measurements calculated for the CCFs shown in the upper plot with reference to transit phase. The plotted RM effect model is for an aligned orbit, for the data shown, it is unclear what the best fit shape of the RM effect curve should be, which further indicates that higher quality data is required to fully determine the orbital obliquity of the WASP-93 system. The upper plot in Figure \ref{fig:w93tomography} shows the time series of the residuals of the CCFs after the subtraction of the average CCF shape, in which a signature of planetary transit is visible near the blue shifted limb of the CCFs, particularly in the first set of data. The effect is predominantly visible during the centre of the transit, where the most starlight is occulted by the planet. Given that none of the red shifted limb of the CCF is occulted by the observed planet signal, it is expected that WASP-93b has an orbit almost entirely mutually misaligned with the stellar spin axis. \subsubsection{Blend scenario} Given the shape of the transit appears to be almost v-shaped, it is important to investigate whether a background blended eclipsing binary or hierarchical triple system could be producing the eclipse signal observed in the lightcurve of WASP-93 and the RV variation detected. The AO image shows a blended stellar companion separated by $0.69\pm0.01$ arcsec, which is fainter by more than 3 magnitudes, so would not be able to produce the transit depth observed, even if the star was fully occulted by a non-emitting body. The AO imaging also shows no other stellar bodies within 0.3 arcsec of the stellar core of WASP-93 which would be bright enough to produce an eclipse mimicking a planet transiting WASP-93. Whilst the transit signal could be created by a chance-aligned background eclipsing binary which is not resolved in the AO images, this scenario is extremely unlikely given the small region of space where a bright enough stellar binary system would have to exist in to not be resolved in the AO image. The most likely blend scenario would be that WASP-93 is a hierarchical triple system, but this would produce a blended line profile in the observed CCFs. Even with CCFs calculated to $\pm 100$ km s$^{-1}$ from the systemic RV, there is no evidence of an additional line profile. The modulation of the CCF shape during the transit also confirms the planetary nature of the signal, rather than that of a hierarchical triple system. \subsubsection{Tidal evolution} \begin{figure} \includegraphics[width=0.49\textwidth]{w93_tidal3} \caption{Plot showing the tidal equilibrium curves for WASP-93. In the upper plot, the blue line shows the total angular momentum of the system when dual-synchronised for the range of semi-major axis; the green line shows the current total angular momentum with an assumption of spin-orbit alignment, and the red line shows the current separation of the star and planet. Each of the angular momenta are scaled with the current total orbital angular momentum, L$_{orb,0} = 2.986\times 10^{42}$ kg m$^{2}$/s. The lower plot shows the curves for the total energy in the system, where the blue line shows the total orbital and rotational energy for the system when dual synchronised for the range of orbital separation; the green line shows the total energy when angular momentum is conserved, and the red line shows the current separation. Each of the energies are scaled with the current orbital energy, E$_{orb,0} = -3.976\times 10^{37}$ kg m$^{2}$/s$^{2}$. The 1-$\sigma$ uncertainties plotted are calculated from the output chains of the global MCMC analysis of the system.} \label{fig:w93tidal} \end{figure} Unlike most hot Jupiters, WASP-93 has a value of $L_{tot} > L_{c}$, which indicates that an orbit exists for this system where tidal equilibrium can be reached -- in this case $L_{tot} / L_{c} = 1.57 \pm{0.16}$. Tidal equilibrium is characterised by the minimisation of total orbital and spin energy, constrained by the conservation of angular momentum. Figure \ref{fig:w93tidal} shows where the current angular momentum of the system (in green) intersects with the angular momentum curve for the dual synchronous state (in blue), which is where stable orbits exist. The current orbital separation of the star and planet (in red) is between the inner and outer equilibrium states. The inner equilibrium is unstable, which can be seen in the lower plot demonstrating the total energy of the system, where that equilibrium denotes a maximum in energy. The energy in the plot has been defined as \begin{equation} E_{tot} = -\frac{G M_{*} M_{p} }{2a} + \frac{1}{2} C_{*} \omega_{*}^{2} + \frac{1}{2} C_{p} \omega_{p}^{2} \, \, , \end{equation} where the first term is the orbital energy, and the remaining terms are the rotational energy of the star and planet respectively. In the case of pseudo-synchronisation, the orbital frequency, $n = \omega_{p}$, and the value for $\omega_{*}$ is determined by angular momentum conservation (green line) or synchronisation with the orbital frequency (blue line). If the planetary orbit and the stellar spin are aligned and since angular momentum is conserved, the WASP-93 system will migrate along the path of constant $L$ until it reaches the outer equilibrium, which is at an energy minimum and is a Darwin stable dual-synchronised orbit. The timescale for this outwards migration can be approximated using a numerical integration of Equation 18 from \citet{matsumura2010}\footnotemark \footnotetext{$\frac{da}{dt} = \frac{9}{Q'_{*,0}}\frac{M_{p}}{M_{*}}\frac{R^{5}_{*}}{a^{4}} \left( \frac{a_{0}}{a} \right)^{3/2} (\omega_{*,0}\cos{\epsilon_{*,0}} - n)$, where $\epsilon_{*,0}$ is the current spin orbit misalignment.}. Using this integration, the timescale until the equilibrium orbit is reached is $\sim 2 \times 10^{7}$ Gyr for $Q'_{*,0}=10^{8}$, which is far longer than the remaining lifetime of the star even for lower values of $Q'_{*,0}$. The timescale for an observable change in the orbital period ($>60$s) is also too long to detect at $\sim 50$ Myr for $Q'_{*,0}=10^{8}$. Since the tomographic analysis indicates that the orbit of the planet is significantly misaligned with the spin axis of the host star, the current total angular momentum of the system in the plane of the stellar rotation is lower than if the system were aligned. If the orbits are misaligned, beyond $\lambda \sim 30^{\circ}$, which is indicated by the tomographic signal, then either $L_{tot} < L_{c}$ or the current separation of the star and the planet would mean that the planet is tidally migrating towards its host star. \subsection{WASP-118 system} \begin{figure} \includegraphics[width=0.45\textwidth]{tomographywasp118} \includegraphics[width=0.45\textwidth]{w118_rvexpanded} \caption{Upper plot showing the time series of CCF residuals after the subtraction of the average CCF shape with reference to transit phase the attempt to observe spectra throughout a transit, where the trail of the transit of WASP-118b is not visible. Lower plot showing the RV measurements around the transit phase. The colours used correspond to those in the lower plot of Figure \ref{fig:wasp118}, and the purple RV points match the CCFs shown in the upper plot.} \label{fig:w118detail} \end{figure} WASP-118b is a {\bf$\sim 0.52$} $M_{\rm J}$ planet in a 4.04 days orbit around an F6 spectral type star. \textsc{Bagemass} was also used to obtain mass and age estimates for WASP-118, which gave $2.38\pm 0.38$ Gyr and $1.40\pm 0.05$ $M_\odot$. The results of the \textsc{Bagemass} modelling are shown in the lower plot in Figure \ref{fig:evo}, which shows that WASP-118 is slightly evolved. This result is in agreement with the spectral analysis for this system in Section \ref{sec:stellar118}. Since the weather on La Palma was poor for the spectral observations on the night beginning 1st October 2015, the uncertainties calculated by the DRS appear lower than would be expected for these conditions. The time series of CCFs collected can be seen in the upper plot of Figure \ref{fig:w118detail}, which highlights the varying observing conditions, and the effect of the 79\% illuminated moon an average of 42$^{\circ}$ away from the star during the observations. The CCFs with the average CCF profile subtracted showed no sign of the planet occulting part of the line profile during transit, which is not surprising given the noisy data. The scatter in the RV measurements calculated can also be clearly seen in the lower plot in Figure \ref{fig:w118detail}, which shows the measured RVs with respect to transit phase. Due to the concern about the reliability of these measurements, which could bias the results, the global fit was completed and compared with the inclusion and exclusion of this data set. With the transit series of HARPS-N RVs excluded when the solution was allowed to be non-circular, the value for $e$ was $0.089^{+0.057}_{-0.047}$. The associated BICs were 20.3 for a circular solution, and 67.0 for an eccentric solution. If the extra RV points are included in the global fit, $e=0.132^{+0.105}_{-0.080}$, and the BIC for a circular solution is 50.3, whereas for a non-circular solution the BIC is 520.4. In each case, the $\Delta$BIC $=$ BIC$_{ecc}$ - BIC$_{circ}$ $>$ 10, which is strong evidence against including the two extra free parameters for eccentricity, so the system is assumed to be circular. The scatter of the transit series of HARPS-N RVs around the fitted model, which is not accounted for with the uncertainties quoted, would not be expected to be caused by stellar activity on the timescale of a few of hours. The data points have been excluded from the fit used for the final parameters quoted, to prevent the underestimated uncertainties from biasing the parameters presented in this work. The distribution of the RV measurements taken during the transit suggest that the orbit of WASP-118b is aligned with respect to the stellar spin axis, which suggests a disc-migration origin for the planet and makes the system a strong target for looking for companion planets, such as those found for WASP-47 \citep{becker2015,neveuvanmalle2015}. Were any companion planets around WASP-118 assumed to have orbits co-planar with the orbit of WASP-118b, planets up to 0.45 AU\footnotemark \footnotetext{Using the relation $a_{t} = {R_*}/{\sin(|90-i|)}$, where $a_{t}$ is the furthest orbital separation where transits will occur for the given inclination of the orbit.} away from the host star would also transit, which will be detectable with K2 depending on the radii of the planets. \subsubsection{Tidal evolution} \begin{figure} \includegraphics[width=0.49\textwidth]{w118_tidal3} \caption{Plot showing the tidal equilibrium curves for WASP-118, which are presented in the same way as Figure \ref{fig:w93tidal}. The upper plot is scaled with orbital angular momentum, L$_{orb,0} = 1.164\times 10^{42}$ kg m$^{2}$/s, and the lower plot is scaled with orbital energy, E$_{orb,0} = -1.046\times 10^{37}$ kg m$^{2}$/s$^{2}$.} \label{fig:w118tidal} \end{figure} WASP-118 also has $L_{tot} > L_{c}$ ($=1.07\pm 0.08$), thus the system very likely has enough angular momentum to exist in a tidal equilibrium with the same assumptions about the system geometry used above. In order to test whether the system is evolving to a stable orbit, the tidal equilibrium curves for WASP-118 are plotted in Figure \ref{fig:w118tidal} as for WASP-93. The red line denoting the current orbital separation is clearly inside the inner equilibrium state, which indicates the system cannot reach a stable orbit. WASP-118b will follow the path of constant $L$ towards the local lowest energy state, which is tidal migration inwards towards the Roche limit of the star. Given the relatively high $\vsini$ of WASP-118, the expression for $t_{\textrm{remain}}$ used for WASP-92b is not an appropriate approximation in this case. In order to estimate the spiral in time for the system, the equation for $\frac{da}{dt}$ used for WASP-93b from \citet{matsumura2010} was numerically integrated from the current orbital separation to the stellar radius. In order to construct this integral, it is approximated that $e$=0, and that the planetary spin is synchronised with the orbital period, which are valid assumptions for a short period giant planet system. The spin and orbit are also assumed to be aligned, which is implied from the in transit RV measurements for WASP-118. The result of the calculation is $t_{\textrm{remain}} \sim 150$ Gyr when $Q'_{*,0}=10^{8}$. As for WASP-92b, the spiral in timescale of the planet is too long to be observable even with a higher tidal efficiency (lower $Q'_{*,0}$.)

1607.00774

1607

1607.05198_arXiv.txt

The co-orbital satellites of Saturn, Janus and Epimetheus, swap radial positions every 4.0 years. Since \textit{Cassini} has been in orbit about Saturn, this has occurred on 21 January in 2006, 2010, and 2014. We describe the effects of this radial migration in the Lindblad resonance locations of Janus within the rings. When the swap occurs such that Janus moves towards Saturn and Epimetheus away, nonlinear interference between now-relocated density waves launches a solitary wave that travels through the rings with a velocity approximately twice that of the local spiral density wave group velocity in the A ring and commensurate with the spiral density wave group velocity in the B ring.

\label{sec::introduction} \textit{Cassini} investigations of the rings of Saturn have revealed a rich collection of phenomena both secular and transient. One of the most important drivers of processes in the rings is the gravitational influence of a subset of Saturn's more than 60 moons. Two of these moons are particularly dynamic: co-orbitals Janus and Epimetheus. Every 4.0 years, they migrate radially and switch their positions relative to the planet. When this migration occurs, so too move the resonance locations between the rings and the two moons. \citet{Tiscareno2006b} have described the effects of this swap on overlapping linear spiral density waves, but stellar occultations observed by \textit{Cassini} at various times during this process have revealed that the nonlinear density waves do not respond in the same way. We investigate that phenomenon here. \subsection{Resonances and spiral density waves} Saturn's numerous moons provide for a large number of inner Lindblad resonances within the A and B rings. \citet{Lissauer1982} give the condition for such a resonance as: \begin{eqnarray} m(\Omega_p - n) = -\kappa, \label{eqn::inner_lindblad_resonance} \end{eqnarray} where $m\Omega_p = mn_s + k\kappa_s$ is an integer multiple of the angular speed $\Omega_p$ with which the gravitational potential of an orbiting, exterior satellite with no inclination rotates. It is given in terms of $n_s$, the satellite's mean orbital angular velocity and $\kappa_s$, the radial, or epicyclic, frequency of its orbit. The numbers $n$ and $\kappa$ represent the same quantities for a test particle within the rings. The coefficients $m$ and $k$ are integers, with $m$ specifying the number of spiral arms the resonance will create. To first order, $n \approx \kappa,~n_s \approx \kappa_s$ and we can rewrite equation~\ref{eqn::inner_lindblad_resonance} as \begin{eqnarray} (m+k)n_s = (m-1)n. \end{eqnarray} This leads to the common practice of labeling resonances as $(m+k):(m-1)$. In order to precisely compute locations of resonances within the rings, however, we must not make approximations like $n_s~\approx~\kappa_s$ and instead compute them separately based on a more complete model of Saturn. This is necessary because of Saturn's oblateness; if the planet were a perfect sphere, then $n_s~=~\kappa_s$. We use the method of \citet{Lissauer1982} to compute precise locations for the resonances. Table~\ref{tab::res_params} lists the parameters used in these computations. \begin{table} \begin{center} \begin{tabular} { c c } Parameter & Value\\ \hline $GM_\text{Saturn}$ & 37931207.7 km$^3$s$^{-1}$\\ $r_\text{Saturn}$ & 60,330 km\\ $J_2$ & 1.629071$\times 10^{-2}$\\ $J_4$ & -9.3583$\times 10^{-4}$\\ $J_6$ & 8.614$\times 10^{-5}$\\ $J_8$ & -1$\times 10^{-5}$\\ $n_\text{Jan, in}$ & 1.04708751$\times 10^{-4}$ s$^{-1}$\\ $n_\text{Jan, out}$ & 1.04687109$\times 10^{-4}$ s$^{-1}$\\ $n_\text{Epi, in}$ & 1.04736844$\times 10^{-4}$ s$^{-1}$\\ $n_\text{Epi, out}$ & 1.04658870$\times 10^{-4}$ s$^{-1}$\\ \end{tabular} \caption{The parameters used to compute the resonance locations specified in this work. All except the mean motions of Janus and Epimetheus are taken from \citet{Jacobson2006}.} \label{tab::res_params} \end{center} \end{table} At many of these resonance locations, features are excited within the rings. \citet{Lin1964} were the first to develop a theory of density waves to explain the spiral structure of many galaxies, thus giving the phenomenon its name. Later, \citet{Goldreich1982} and \citet{Shu1984} were among the first to apply the theory to explain the structure of Saturn's rings. Linear density waves are raised by weaker resonances as small perturbations on the background surface mass density of the rings. Stronger resonances generate larger perturbations, which can be of the same order as the background density. When this occurs, a nonlinear density wave forms, with sharp peaks (regions of highest density) and shallow troughs (regions of lowest density). The predator-prey model of \citet{Esposito2012} identified strongly-perturbed regions such as nonlinear density waves as a likely site of increased aggregation within the rings. \subsection{Surface mass density in the A and B rings} Using images of spiral density waves observed with the \textit{Cassini} Imaging Science Subsystem (ISS), \citet{Tiscareno2007} derive an approximate surface mass density for the A ring of 40 g/cm$^2$. This surface density appears to increase as one moves farther from the planet. \citet{Esposito1983} used a \textit{Voyager} occultation of the Janus 2:1 spiral density wave to estimate a surface mass density for the inner B ring of 70$\pm$10 g/cm$^2$. This is somewhat larger than the results of \citet{Reffet2015}, who use \textit{Cassini} CIRS observations to estimate a value on the order of 40 g/cm$^2$ for the inner B ring, 100 g/cm$^2$ for the middle of the ring and 50 g/cm$^2$ towards the outer boundary. The results of \citet{Hedman2016} bridge this gap by finding that the surface mass density varies within the Janus 2:1 resonance from 69 g/cm$^2$ near resonance to 47 g/cm$^2$ several hundred kilometers exterior. These values are all substantially lower than an estimate that considers the dynamics and aggregation of ring particles by \citet{Robbins2010} of 240-480 g/cm$^2$. The background surface mass density $\sigma_0$ strongly governs the group velocity $v_g$ of spiral density waves excited at resonance locations, given by $v_g = \frac{\pi G \sigma_0}{\kappa}$, for epicyclic frequency $\kappa$ and gravitational constant $G$ \citep{Toomre1969}. For this study, we assume a surface mass density of 40 g/cm$^2$ for the A ring and as a lower limit for the B ring. We take 70 g/cm$^2$ as an upper limit for the B ring in the vicinity of the Janus 2:1 resonance. \subsection{Janus and Epimetheus} Among Saturn's many moons, Janus and Epimetheus are of particular interest because they represent the only pair of co-orbiting satellites in the solar system. \citet{Dermott1981} showed that this co-orbital configuration leads to each moon traversing a horseshoe orbit about the pair's shared mean orbital radius of 151,450 km in a frame of reference rotating with their mean angular velocity. This gives a small relative velocity between the two. Every 4.0 years, the pair approach each other within 15,000 km \citep{Nicholson1992} and exchange orbital angular momentum. This causes a rapid shift, in which the inner and outer bodies switch position in the course of approximately 100 days. Each moon is radially shifted in proportion to its relative mass ($m_E/m_J = 0.278$): 20 km for Janus and 80 km for Epimetheus. During \textit{Cassini's} time at Saturn, this has occurred three times: 21 January 2006 (Janus moves inwards), 21 January 2010 (Janus moves outwards), and 21 January 2014 (Janus moves inwards). Because of this change in the radius of their orbits, special care must be taken when computing the mean motions (and thus the pattern speeds and resonance locations) of these bodies. The semi-major axes given in table~\ref{tab::moon_info} are valid only for times distant from the orbital swap; during the swap, the values evolve continuously. Table~\ref{tab::resonances} lists the computed resonance locations for the first-order Janus and Epimetheus resonances used in this study. Although the focus of this work is the effect that the Janus/Epimetheus orbital swap has on spiral density waves raised within the rings, \citet{ElMoutamid2016} have also observed that changes in the Janus 7:6 resonance affect the shape of the outer edge of the A ring. \begin{table} \begin{center} \begin{tabular}{ c c c c c c } Name & $M$ (kg) $^1$& $a$ (km) $^2$& $i$ ($^\circ$) $^3$& $e$ $^3$\\ \hline Epimetheus & 5.3$\times 10^{17}$ & & 0.351 & 0.0098\\ 2002-2006, 2010-2014 & & 151,410 & & \\ 2006-2010, 2014-2018 & & 151,490 & &\\ Janus & $1.9\times 10^{18}$ & & 0.163 & 0.0068\\ 2002-2006, 2010-2014 & & 151,460 & & \\ 2006-2010, 2014-2018 & & 151,440 & &\\ \end{tabular}\\ $^1$\citet{Thomas2010}, $^2$\citet{Jacobson2008}, $^3$\citet{Spitale2006}\\ \caption{Basic properties of Janus and Epimetheus. Values with specified dates refer to 21 Jan of that year and are not accurate in the immediate vicinity (approximately 100 days) of those end points.} \label{tab::moon_info} \end{center} \end{table} \begin{table} \center{ \begin{tabular}{ c c c c c c } Resonance & Time period & $r_{\text{res}}$ (km) & $n_\text{occ}$ & $n_\text{features}$ & $n_\text{empty}$\\ \hline Epimetheus 6:5 & & & 152 & - & -\\ &2002-2006, 2010-2014 & 134,223 &&&\\ &2006-2010, 2014-2018 & 134,289 &&&\\ Epimetheus 5:4 && & 151&-&-\\ &2002-2006, 2010-2014 & 130,660 &&&\\ &2006-2010, 2014-2018 & 130,724 &&&\\ Epimetheus 4:3 && & 142 &-&-\\ &2002-2006, 2010-2014 & 125,228 &&&\\ &2006-2010, 2014-2018 & 125,290 &&&\\ Epimetheus 3:2 && & 132 &-&-\\ &2002-2006, 2010-2014 & 115,922 &&&\\ &2006-2010, 2014-2018 & 115,979 &&&\\ Epimetheus 2:1 && & 103 &-&-\\ &2002-2006, 2010-2014 & 96,216 &&&\\ &2006-2010, 2014-2018 & 96,263 &&&\\ Janus 6:5 && & 152 & 92 & 50\\ &2002-2006, 2010-2014 & 134,265 &&&\\ &2006-2010, 2014-2018 & 134,247 &&&\\ Janus 5:4 && & 151 & 110 & 33\\ &2002-2006, 2010-2014 & 130,701 &&&\\ &2006-2010, 2014-2018 & 130,683 &&&\\ Janus 4:3 && & 142 & 127 & 19\\ &2002-2006, 2010-2014 & 125,267 &\\ &2006-2010, 2014-2018 & 125,250 &\\ Janus 3:2 && & 132 & 1 & 104\\ &2002-2006, 2010-2014 & 115,959 &&&\\ &2006-2010, 2014-2018 & 115,943 &&&\\ Janus 2:1 && & 103 & 177 & 5\\ &2002-2006, 2010-2014 & 96,246 &&&\\ &2006-2010, 2014-2018 & 96,233 &&&\\ Mimas 5:3 & & 132,298 & 151 & 7 & 139\\ Prometheus 31:30 & & 136,389 & 146& 0 & 130\\ Prometheus 14:13 & & 132,716 & 153& 0 & 141 \\ \end{tabular} \caption{Inner Lindblad resonance locations used for this work. The orbit swap of Janus and Epimetheus means that the resonance location shifts on 21 January of the specified years. The number of examined occultations $n_\text{occ}$ does not equal the sum of observed features $n_\text{features}$ and occultations with no observed features $n_\text{empty}$ because some occultations contain multiple features and some have a signal-to-noise ratio too low to resolve the density wave. Because the Epimetheus resonance regions overlap the (much stronger) Janus ones, they were not independently searched for anomalous features.} \label{tab::resonances} } \end{table}

It is clear that the orbital swap of Janus and Epimetheus has a substantial effect on the rings. \citet{Tiscareno2006b} have demonstrated that it accounts for the unusual morphology of second-order, linear spiral density waves and \citet{ElMoutamid2016} have shown it substantially alters the shape of the outer edge of the A ring. Our investigation reveals that in Saturn's strongest density waves, the effect is even more dramatic. Nonlinear interference between the waves generated at the pre- and post-swap resonance locations results in the formation of a solitary wave. In the A ring, this wave propagates outward at about twice the group velocity of the local spiral density waves and in the B ring the propagation velocity is commensurate to the density wave group velocity. Occultations observed by \textit{Cassini} throughout 2016 and images captured until the end of the mission should reveal the next iteration of this cycle and provide a longer baseline over which to understand the phenomenon. A future model describing the behavior of the solitary wave may also allow for an independent estimation of the ring surface mass density. Finally, these observations have potential applications in understanding the physics of protoplanetary and accretion disks, where the migration of massive bodies has been theorized. Although we may generally lack the capability of detecting these moving bodies, the effects they render onto the neighboring disk of particles and gas may reveal their motion indirectly.

1607.05198

1607

1607.04541_arXiv.txt

{ Electrical conductivity of finite-temperature plasma in neutron star crusts is studied for applications in magneto-hydrodynamical description of compact stars. We solve the Boltzmann kinetic equation in relaxation time approximation taking into account the anisotropy of transport due to the magnetic field, the effects of dynamical screening in the scattering matrix element and correlations among the nuclei. We show that conductivity has a minimum at a non-zero temperature, a low-temperature decrease and a power-law increase with increasing temperature. Selected numerical results are shown for matter composed of carbon, iron, and heavier nuclei present in the outer crusts of neutron star. } \FullConference{The Modern Physics of Compact Stars 2015\\ 30 September 2015 - 3 October 2015\\ Yerevan, Armenia} \begin{document}

Magneto-hydrodynamics (MHD) forms the basis of large-scale description of physics of dense plasma in compact stars. A key quantity in the dissipative formulations of MHD is the conductivity of matter. It determines, for example, the dissipation of currents and therefore the decay of magnetic fields, the dispersion of plasma waves, etc. In turn, magnetic field decay affects the rotational and thermal evolutions of neutron stars and consequently a broad array of their observational manifestations. Transport in compact star plasma was studied traditionally in the cold (essentially zero-tem\-pe\-rature) and dense regime where the constituents form degenerate quantum liquids. This regime is relevant for mature isolated or accreting neutron stars as well as interiors of white dwarfs. The dilute and warm (non-zero temperature) regime is of interest in the context of transient, short-lived states of neutron stars, such as proto-neutron stars newly born in supernova explosions or hypermassive remnants formed in the aftermath of neutron star binary mergers. We start this article with an overview of the transport calculations of electrical conductivity of compact star matter in the density regime corresponding to their outer crusts ($\rho \le 10^{11}$ g cm$^{-3}$). Then we go on to describe our recent effort to calculate the electrical conductivity of non-zero temperature crustal plasma. We focus on sufficiently high temperatures where nuclei form a liquid coexisting with electronic background of arbitrary degeneracy. We close this review with a summary and outlook. Below we use the natural (Gaussian) units with $\hbar= c = k_B = k_e = 1$, $e=\sqrt{\alpha}$, $\alpha=1/137$ and the metric signature $(1,-1,-1,-1)$.

In this contribution we gave an overview of our current work on the conductivity of dense matter in the envelopes of neutron stars at non-zero temperature. One ingredient of our effort is the formulation of the transport in a manner which allows us to include the dynamical screening exactly, provided that the polarization tensor of electrons (or equivalently the self-energies of QED photons) in plasma can be computed to desired accuracy. Here we employed the results based on the hard-thermal-loop approximation and low-frequency expansion appropriate at not very high temperatures. We have shown that for electron-ion scattering the dynamical screening is suppressed parametrically by a factor $M/T$, but we anticipate that its effect would be substantial in the cases (a) of high temperatures and presence of light clusters, (b) of low temperatures, where the Umklapp processes with $e$-$e$ scattering are important, (c) of transport processes where $e$-$e$ scattering may be dominant from the outset, such as the thermal conductivity and shear viscosity. Our numerical results show that the scalar conductivity (no anisotropy due to the $B$-field) has a minimum as a function of temperature, with a power-law decrease at low-temperatures and a power-law increase at higher temperatures. The range of validity of the zero-temperature Drude formula extends from low temperatures up to 0.1$-$1 MeV ($10^9-10^{10}$ K), where the lower of these bounds corresponds to density $\rho \sim 10^{6}$ g cm$^{-3}$ and the upper one to $\rho \sim 10^{10}$ g cm$^{-3}$. The behaviour of the off-diagonal $\sigma_1$ component of the conductivity tensor is similar to the one described above except at low densities, where it remains almost constant. Finally the $\sigma_0$ component shows strongly density dependent behaviour: for large densities (high degeneracy) it behaves analogous to $\sigma$, but shows the inverse trend for low-densities which is associated with the transition from the regime $\omega_c\tau <1 $ to $\omega_c\tau >1 $.

1607.04541

1607

1607.03966_arXiv.txt

We consider the dispersion on the supernova distance-redshift relation due to peculiar velocities and gravitational lensing, and the sensitivity of these effects to the amplitude of the matter power spectrum. We use the MeMo lensing likelihood developed by Quartin et al., which accounts for the characteristic non-Gaussian distribution caused by lensing magnification with measurements of the first four central moments of the distribution of magnitudes. We build on the MeMo likelihood by including the effects of peculiar velocities directly into the model for the moments. In order to measure the moments from sparse numbers of supernovae, we take a new approach using Kernel Density Estimation to estimate the underlying probability density function of the magnitude residuals. We also describe a bootstrap re-sampling approach to estimate the data covariance matrix. We then apply the method to the Joint Light-curve Analysis (JLA) supernova catalogue. When we impose only that the intrinsic dispersion in magnitudes is independent of redshift, we find $\sigma_8=0.44^{+0.63}_{-0.44}$ at the one standard deviation level, although we note that in tests on simulations, this model tends to overestimate the magnitude of the intrinsic dispersion, and underestimate $\sigma_8$. We note that the degeneracy between intrinsic dispersion and the effects of $\sigma_8$ is more pronounced when lensing and velocity effects are considered simultaneously, due to a cancellation of redshift dependence when both effects are included. Keeping the model of the intrinsic dispersion fixed as a Gaussian distribution of width 0.14 mag, we find $\sigma_8 = 1.07^{+0.50}_{-0.76}$.

\label{sec:intro} Understanding the nature of dark energy is one of the key goals of modern cosmology. A recurring theme in understanding dark energy is measuring both the growth of structures and geometry of the Universe \citep[e.g.][]{2008MNRAS.389L..47A,2010PhRvD..81j3510Z,2012PhRvD..85l3546Z,2013MNRAS.429.1514S,2014MNRAS.439.3504S,2015PhRvD..91f3009R}. In this work, we consider using supernovae, typically considered a probe of geometry, to constrain the growth of structure. \subsection{Growth and Geometry} Here, geometry refers to observational probes that are predominantly sensitive to the background expansion of the Universe: observations that constrain the distance-redshift relation. Among geometrical probes, observations are either `standard rulers', such as Baryon Acoustic Oscillations \citep[e.g.][]{2003ApJ...598..720S,2003ApJ...594..665B,2005ApJ...633..560E,2011MNRAS.418.1707B,2012MNRAS.427.3435A,2013A&A...552A..96B,2014MNRAS.441...24A}, or `standard candles', such as supernovae Ia \citep[e.g.][]{2009ApJ...700..331H,2009ApJS..185...32K,2010A&A...523A...7G,2010AJ....139..120F,2011ApJS..192....1C,2014A&A...568A..22B,2016arXiv160303823N}. Supernovae are well known as a key observable probe in establishing the accelerated expansion of the Universe \citep{1998AJ....116.1009R,1999ApJ...517..565P}. These geometrical observables constrain the densities of the constituents of the Universe (such as the matter density, $\Omega_{\rm{m}}$), and the equations of state of these densities. A key aim of these observations is to measure $w$, the equation of state of dark energy. With the accelerating expansion of the Universe now well established \citep[e.g.][]{2008ARA&A..46..385F,2015arXiv150201590P}, and ever more precise measurements of $w$, new dark energy observations are focusing on measuring the growth of the Universe \citep[e.g.][]{2005PhRvD..72d3529L,2012MNRAS.425..405B}. Here, growth refers to the growth of cosmological density perturbations. The motivation for measuring the growth of density perturbations is that theoretical models of dark energy (which must reproduce the observed distance-redshift relation) often predict different expectations for the growth of density perturbations \citep[e.g.][]{2012PhR...513....1C}. Among growth probes, there is a natural division of the observations into two kinds: either relativistic or non-relativistic. Relativistic techniques such as gravitational lensing \citep{2001PhR...340..291B,2012MNRAS.427..146H} and the integrated Sachs Wolfe effect \citep{1967ApJ...147...73S,2003astro.ph..7335S,2005PhRvD..72j3519P,2006PhRvD..74f3520G,2015arXiv150201595P} are sensitive to the path that photons take, and are therefore sensitive to both time-like and space-like perturbations to the metric. On the other hand, non-relativistic observations, such as Redshift Space Distortions, \citep{1987MNRAS.227....1K,2001Natur.410..169P,2008Natur.451..541G,2012MNRAS.426.2719R} focus on the positions and velocities of large scale structure. These observations are sensitive only to time-like perturbations to the metric, the `Newtonian' potential. A key feature of general relativity (GR) is that time-like and space-like perturbations to the metric should be equal \citep[e.g.][]{2007PhRvD..76j4043H,2013PhRvD..87b4015B}, and so measuring both Newtonian and Lensing density fluctuations is a powerful technique to constrain modified gravity alternatives to a cosmological constant dark energy \citep[e.g.][]{2010Natur.464..256R,2013MNRAS.429.2249S,2013A&A...557A..54D,2015JCAP...12..051L,2015arXiv151104457P,2016MNRAS.456.2806B}. The equality of these two potentials is only true in the absence of anisotropic stress, as we would expect for cold dark matter (CDM), so an inequality may also suggest dark matter interactions. We thus have several observational tests that any theory of the dark Universe must pass: the distance-redshift relation, and measurements of Newtonian and Lensing density fluctuations. Supernovae are one of the most established probes of the distance-redshift relation. In this paper, we focus on using supernovae to also constrain Newtonian and Lensing density fluctuations. The main motivation for this work is to contribute to the development of a new method to measure the amplitude of the matter power spectrum with supernovae. $\sigma_8$ is the amplitude of the matter power spectrum at scales of 8 $h^{-1}$Mpc, and a key parameter to measure in order to constrain cosmological density fluctuations, although as we will see, supernova magnitude fluctuations are sensitive to fluctuations on very different physical scales. The best-fitting value for $\sigma_8$ from the Planck measurements of the cosmic microwave background is 0.830$\pm$0.015 \citep{2015arXiv150201589P}. However, it is important to remember that this is a derived value (not a measurement) that depends on the measured value of the amplitude of primordial density fluctuations ($A_s$), and a cosmological model. Current measurements of $\sigma_8$ from weak lensing shear and redshift space distortions find a value that is lower than expectations from the $\Lambda$CDM model with Planck parameters \citep[e.g.][]{2013MNRAS.429.1514S,2013PhRvL.111p1301M,2014MNRAS.444..476R,2016PDU....13...66C,2015arXiv150201590P,2016arXiv160600439G}. Depending on the choice of model (such as imposing flatness, or a cosmological constant), the tension is at the level between two and three standard deviations. While the tension remains only moderate, and unaccounted-for systematic effects in the observables cannot be ruled out, there is the possibility that the tension may represent some of the first hints of physics beyond the $\Lambda$CDM model. New methods to measure $\sigma_8$ are important to determine if the current tension is due to new physics, or systematic effects. \subsection{Signals in the Noise} A more general motivation for this work is to develop a method to measure cosmological parameters from signals that are often considered as noise. \cite{2013JCAP...06..002B,2013PhRvL.110b1301B} modelled the effect of cosmological density fluctuations on luminosity distance measurements. They found that peculiar velocities and gravitational lensing are the dominant sources of the dispersion, and place a fundamental limit on our ability to measure parameters affecting the distance redshift relation, such as $\Omega_{\Lambda}$. In this work, instead of considering the dispersion from lensing and peculiar velocities as noise in the distance-redshift relation, cosmological signals in this distribution are our primary focus. At low redshift, peculiar velocities can be inferred for individual galaxies with an independent distance indicator, such as the Tully-Fisher, Fundamental Plane or supernovae \citep[e.g.][]{2012ApJ...751L..30H,2012MNRAS.423.3430B,2016MNRAS.455..386S}. This can be achieved at higher redshifts if galaxy sizes or magnitudes are cross-correlated with galaxy overdensities \citep{2014MNRAS.443.1900B}. Also at higher redshift, peculiar velocities can be inferred statistically in a galaxy redshift survey, via redshift space distortions \citep[e.g.][]{2011MNRAS.415.2876B,2015arXiv151108083O,2016PhRvD..93b3525S}. \cite{2007PhRvL..99h1301G} analysed the peculiar velocities of 271 low redshift supernovae, and found $\sigma_8=0.79 \pm 0.22$. \cite{2015JCAP...12..033H} analysed the bulk flow of the 100 lowest redshift galaxies in the joint light-curve analysis (JLA) catalogue. Instead of fitting for $\sigma_8$ directly, the analysis fitted the amplitude $A$ of the peculiar velocity covariance matrix, normalized to the expected value in $\Lambda$CDM. The analysis found a value of $A$ consistent with zero peculiar velocities, but with a large uncertainty that also included the expected $\Lambda$CDM value. The magnitudes of standard candles are also affected by gravitational lensing. In rare cases, when the photons pass close to a massive cluster, a supernova can be strongly lensed, or even lensed into multiple images \citep[e.g.][]{2015Sci...347.1123K}. However, all supernovae will be weakly-lensed by cosmological density perturbations. \cite{2005ApJ...633..589S} detected lensing magnification in the cross-correlation of quasar and foreground galaxy correlations. One of the most established techniques for detecting weak gravitational lensing is by measuring coherent distortions in the shapes of galaxies \citep[e.g.][]{2014MNRAS.441.2725F,2015arXiv150705552T}. In contrast, supernova lensing is sensitive to the change in magnitude caused by weak gravitational lensing. Since lensing depends on the integrated path of the photons, the effect is most significant for high redshift supernovae \citep[e.g.][]{2016MNRAS.456.1700S}. \cite{2006ApJ...640..417G} considered correcting for the lensing dispersion in the Hubble diagram by estimating the magnification effect from large scale structure, and \cite{2010MNRAS.402..526J} considered using supernovae magnification to constrain the properties of the lensing dark matter haloes. \cite{2013MNRAS.432..679C} considered modelling the line of sight lensing signal to improve time delay measurements. \cite{2014ApJ...780...24S} tested for lensing magnification with 608 supernovae from the Sloan Digital Sky Survey by cross-correlating the magnitude residuals with the expected lensing signal from foreground galaxies. Although the significance of the cross-correlation was low at the 1.4 standard deviation confidence level, the correlation suggests that lensing provides some contribution to the distances measured with supernovae. \cite{2006PhRvD..74f3515D} proposed using the dispersion in the Hubble diagram due to weak lensing to measure $\sigma_8$, but noted that the Gaussian model used in their work can be biased towards incorrect values of $\sigma_8$ due to the non-Gaussian nature of the lensing dispersion. In a series of papers, \cite{2013PhRvD..88f3004M} and \cite{2014PhRvD..89b3009Q} developed a method to measure $\sigma_8$ from the effect of lensing magnification on the magnitude residuals of supernovae (called Method-of-the-Moments -- MeMo). In \cite{2014MNRAS.443L...6C}, the MeMo technique was applied to the JLA and Union2 supernovae catalogues, finding $\sigma_8=0.84^{+0.28}_{-0.65}$ at the 68\% confidence level, or $\sigma_8 < 1.45$ at the 95\% confidence level. In \cite{2016PDU....13...66C}, the MeMo lensing likelihood was combined with a peculiar velocity likelihood. These two different physical effects were combined by using a peculiar velocity likelihood for supernovae with redshift $z<0.1$, and the lensing-only MeMo likelihood for higher redshifts. This approach has the advantage that correlations between supernovae from large scale bulk flows can be modeled in the velocity likelihood, but does not model the combined effect of the two different kinds of perturbations on the moments. \cite{2016PDU....13...66C} allowed both $\sigma_8$ and the perturbation growth index $\gamma$ to vary, finding the best-fitting $\sigma_8=0.65^{+0.23}_{-0.37}$, with $\gamma=1.35^{+1.7}_{-0.65}$. Keeping $\gamma$ fixed at the expected value in GR of 0.55, the best-fitting value of $\sigma_8$ was $0.40^{+0.21}_{-0.23}$. Our approach to treating lensing and velocities is different from that of \cite{2016PDU....13...66C}. Instead of treating the two effects as independent -- with two different likelihoods -- we use a single likelihood that directly combines predictions for lensing and velocities into the expectations for the moments. The advantage of this approach is that the total expectations for the moments include contributions for both effects, which would otherwise be underestimated. \cite{2006PhRvD..74f3515D} assumed that the intrinsic dispersion in supernova magnitudes can be modeled with a Gaussian distribution with mean zero and a standard deviation that is independent of redshift. \cite{2014MNRAS.443L...6C} relaxed this assumption somewhat by allowing the intrinsic dispersion to be further modeled with intrinsic third and fourth moments (although also constant in redshift). In both papers, the only variation in the distribution of residuals was assumed to be due to perturbations in the metric, as a function of $\sigma_8$. In reality, the intrinsic dispersion in the magnitudes of supernovae may vary with redshift, and may not be Gaussian. For example, Malmquist bias may affect the distribution of fainter residuals, sub-populations of different types of Ia supernovae may skew the intrinsic dispersion, or correlations with host-galaxy evolution may introduce redshift dependence \citep[e.g.][]{2009ApJ...691..661H,2010ApJ...722..566L,2010MNRAS.406..782S,2010ApJ...715..743K,2016MNRAS.457.3470C}. \cite{2016arXiv160408885S} tested the Hubble residuals of recent supernova data with the Kolmogorov Smirnov test, and found the residuals to be consistent with a Gaussian distribution. \cite{2014MNRAS.443L...6C} tested models of the intrinsic dispersion that are both constant and vary with redshift, and found that the Bayesian evidence favoured a model for the intrinsic dispersion that is constant in redshift. In this work, we aim to place limits on cosmological density fluctuations via the effects of peculiar velocities and lensing magnification on moments of the Hubble diagram. We verify the MeMo lensing likelihood on simulated catalogues from the MICE light cone simulation. We include the effects of peculiar velocities directly into the moments likelihood, as opposed to including velocities with a separate likelihood. In Section \ref{sec:method_modelling} we describe the physical model of the moments, based on the effects of peculiar velocities and lensing magnification. In Section \ref{sec:data} we review the JLA catalogue and the simulated realizations of the catalogue that we generate. In Section \ref{sec:measurements} we describe our techniques to measure the moments and estimate the data covariance matrix. In Section \ref{sec:results} we present the results of fitting for the lensing and velocity models to the JLA catalogue, and compare our results to other analyses. In Section \ref{sec:conclusions}, we summarize our main conclusions.

\label{sec:conclusions} We have considered the effects of peculiar velocities and lensing magnification on moments of the distance redshift relation. We have described theoretical models for these effects, and statistical methods for undertaking these measurements on sparse data. We note that at redshift $z\sim0.2$--$0.5$, the effects of peculiar velocities and lensing magnification contribute similarly to the dispersion in the Hubble diagram. The cancellation of redshift dependence of these two effects across this redshift range makes the effects of $\sigma_8$ on the second moment in this redshift range degenerate with the intrinsic supernova dispersion. We thus emphasize the importance of modelling both effects simultaneously, and also the importance of measuring higher moments of the Hubble diagram, in order to break this degeneracy. We confirm that the simulated lensing convergence in the MICE light cone simulation is in excellent agreement with the moments modelled by \cite{2013PhRvD..88f3004M}. We present an extension of the MeMo likelihood of \cite{2014PhRvD..89b3009Q}, which directly includes the effects of peculiar velocities in a single likelihood. We note that standard estimators for the moments of the magnitude residuals underestimate the genuine moments for typical numbers of supernovae. We show that KDE can be used to reduce the bias in estimates of the moments from sparse samples. We then apply the velocity and lensing likelihood and the KDE estimators to the JLA supernovae catalogue. Comparing to other work, we note that using the moments estimator of equation (\ref{eq:moments_def_discrete}), we can reproduce the values of $\sigma_8$ and $\sigma_{\rm{I}}$ from \cite{2014MNRAS.443L...6C}. However, our result with this estimator is likely to be an underestimate of the genuine moments of the JLA survey. There are some limitations to the current analysis: as with other work focusing on the Hubble diagram residuals \citep{2014ApJ...780...24S,2014MNRAS.443L...6C}, we do not account for correlations in the distance moduli in our analysis (either from the light-curve fitting, or density perturbations). This approach frees us from the imposition of Gaussianity (which is implicit in a typical covariance-matrix analysis) in the distribution of residuals, which is the dominant signal for lensing, and the primary focus of this paper. In the case of lensing, since the signal depends on the integrated line of sight density, we expect the magnification effect to be uncorrelated for angular separations greater than a few arcminutes. However, we note that \cite{2016arXiv161101315S} found that the magnitude residuals in the JLA catalogue are consistent with zero correlations from lensing magnification. In the case of peculiar velocities, we do expect large-scale correlations, although we have verified by calculating the full covariance matrix that the correlations are negligible for all but the lowest redshift supernovae. In the case of the light curve parameters, \cite{2016PDU....13...66C} showed that marginalizing over the values slightly increased the uncertainties, but did not significantly bias the results or conclusions. Currently, the main limitation in the analysis is measuring unbiased moments of sparse samples of the residuals. However, as the size of supernova catalogues increases, this issue will become less problematic. With larger catalogues, however, it will become more important to model the intrinsic dispersion in the supernovae, such as the dependence with redshift and correlations with host galaxy type. We emphasize that it is possible to place limits on the amplitude of the matter power spectrum with the supernovae Hubble diagram -- both the background expansion and perturbations with the same observable. By fitting for the effects of lensing and velocities, we can also test for consistency in the Newtonian and lensing potentials. Furthermore, by measuring the moments of the residuals, we test for consistency or evidence for outliers in the supernova population.

1607.03966

1607

1607.08559_arXiv.txt

In anticipation of improved observational data for Jupiter's gravitational field from the \textit{Juno} spacecraft, we predict the static tidal response for a variety of Jupiter interior models based on \textit{ab initio} computer simulations of hydrogen-helium mixtures. We calculate hydrostatic-equilibrium gravity terms using the non-perturbative \textit{concentric Maclaurin Spheroid} (CMS) method that eliminates lengthy expansions used in the theory of figures. Our method captures terms arising from the coupled tidal and rotational perturbations, which we find to be important for a rapidly-rotating planet like Jupiter. Our predicted static tidal Love number $k_2 = 0.5900$ is $\sim$10\% larger than previous estimates. The value is, as expected, highly correlated with the zonal harmonic coefficient $J_2$, and is thus nearly constant when plausible changes are made to interior structure while holding $J_2$ fixed at the observed value. We note that the predicted static $k_2$ might change due to Jupiter's dynamical response to the Galilean moons, and find reasons to argue that the change may be detectable, although we do not present here a theory of dynamical tides for highly oblate Jovian planets. An accurate model of Jupiter's tidal response will be essential for interpreting \textit{Juno} observations and identifying tidal signals from effects of other interior dynamics in Jupiter's gravitational field.

\label{sec:intro} The \textit{Juno} spacecraft began studying Jupiter at close range following its orbital insertion in early July 2016. The unique low-periapse polar orbit and precise Doppler measurements of the spacecraft's acceleration will yield parameters of Jupiter's external gravitational field to unprecedented precision, approaching a relative precision of $\sim 10^{-9}$ \citep{kaspi2010}. In addition to providing important information about the planet's interior mass distribution, the non-spherical components of Jupiter's gravitational field should exhibit a detectable signal from tides induced by the planet's closer large moons, possibly superimposed on signals from mass anomalies induced by large-scale dynamic flows in the planet's interior \citep{cao2015,kaspi2010,kaspi2013}. As a benchmark for comparison with expected \textit{Juno} data, \citet{hubbard2016} constructed static interior models of the present state of Jupiter, using a barotropic pressure-density $P(\rho)$ equation of state for a near-solar mixture of hydrogen and helium, determined from \textit{ab intio} molecular dynamics simulations \citep{militzer2013a,militzer2013b}. In this paper, we extend those models to predict the static tidal response of Jupiter using the three-dimensional concentric Maclaurin spheroid (CMS) method \citep{wahl2016}. The \citet{hubbard2016} preliminary Jupiter model is an axisymmetric, rotating model with a self-consistent gravitational field, shape and interior density profile. It is constructed to fit pre-\textit{Juno} data for the degree-two zonal gravitational harmonic $J_2$ \citep{jacobson2003}. While solutions exist matching pre-\textit{Juno} data for the degree-four harmonic $J_4$, models using the \textit{ab initio} EOS required unphysical compositions with densities lower than that expected for the pure H-He mixture. As a result, the preferred model of \citet{hubbard2016} predicts a $J_4$ with an absolute value above pre-\textit{Juno} error bars. Preliminary Jupiter models consider the effect of a helium rain layer where hydrogen and helium become immiscible \citep{stevenson1977a}. The existence of such a layer has important effects for the interior structure of the planet, since it inhibits convection and mixing between the molecular exterior and metallic interior portions of the H-He envelope. This circumstance provides a physical basis for differences in composition and thermal state between the inner and outer portions of the planet. Adjustments of the heavy element content and entropy of the $P(\rho)$ barotrope allow identification of an interior structure consistent with both pre-\textit{Juno} observational constraints and the \textit{ab initio} material simulations. The preferred preliminary model predicts a dense inner core with $\sim$12 Earth masses and an inner hydrogen-helium rich envelope with $\sim$3$\times$ solar metallicity, using an outer envelope composition matching that measured by the \textit{Galileo} entry probe. Although the \textit{Cassini} Saturn orbiter was not designed for direct measurements of the high degree and order components of Saturn's gravitational field, the first observational determination of Saturn's second degree Love number $k_2$ was recently reported by \citet{lainey2016}. This study used an astrometric dataset for Saturn's co-orbital satellites to fit $k_2$, and identified a value significantly larger than the theoretical prediction of \citet{gavrilov1977}. The non-perturbative CMS method obtains values of $k_2$ within the observational error bars for simple models of Saturn's interior, indicating the high value can be explained completely in terms of static tidal response \citep{wahl2016}. The perturbative method of \citet{gavrilov1977} provides an initial estimate of tidally induced terms in the gravitational potential, but neglects terms on the order of the product of tidal and rotational perturbations. \citet{wahl2016} demonstrated, that for the rapidly-rotating Saturn, these terms are significant and sufficient to explain the observed enhancement of $k_2$.

Our study has predicted the static tidal Love numbers $k_{nm}$ for Jupiter and its three most influential satellites. These results have the following features: (a) They are consistent with the most recent evaluation of Jupiter's $J_2$ gravitational coefficient; (b) They are fully consistent with state of the art interior models \citep{hubbard2016} incorporating DFT-MD equations of state, with a density enhancement across a region of H-He imiscibility \citep{morales2013}; (c) We use the non-perturbative CMS method for the first time to calculate high-order tesseral harmonic coefficients and Love numbers for Jupiter. The combination of the DFT-MD equation of state and observed $J_{2n}$ strongly limit the parameter space of pre-\textit{Juno} models. Within this limited parameter space, the calculated $k_{nm}$ show minimal dependence on details of the interior structure. Despite this, our CMS calculations predict several interesting features of Jupiter's tidal response that the \textit{Juno} gravity science system should be able to detect. In response to the rapid rotation of the planet the $k_2$ tidal Love number is predicted to be much higher than expected for a non-rotating body. Moreover, the rotation causes state mixing between different tesseral harmonics, leading to a dependence of higher order static $k_{nm}$ on both $m$ and the orbital distance of the satellite. An additional, significant dependence on $a/r$ is expected in the dynamic tidal response. We present an estimate of the dynamical correction to our calculations of the static response, but a full analysis of the dynamic theory of tides has yet to be performed.

1607.08559

1607

1607.04294_arXiv.txt

\label{sec:intro} The QCD axion is a Nambu--Goldstone boson which arises in association with spontaneous breakdown of Peccei--Quinn (PQ) symmetry, and it dynamically solves the strong CP problem~\cite{Peccei:1977hh,Weinberg:1977ma,Wilczek:1977pj}. The strength of interactions with the standard model particles as well as the periodicity of the axion potential are determined by the axion decay constant, which is currently constrained as \begin{equation} 10^9\, \mathrm{GeV} \lesssim f_a \lesssim 10^{12}\, \mathrm{GeV}. \label{f-bound} \end{equation} Here the lower bound is set by astrophysical arguments of star cooling~\cite{Raffelt:1996wa}, while the upper bound comes from the requirement that the abundance of the axion be less than or equal to that of cold dark matter (CDM) without fine-tuning the initial misalignment angle~\cite{Preskill:1982cy,Abbott:1982af,Dine:1982ah}. Cosmological considerations further set bounds on a combination of the decay constant~$f_a$ and the inflationary Hubble scale~$H_{\mathrm{inf}}$. Firstly, $f_a$ should be larger than $H_{\mathrm{inf}}$, otherwise the PQ symmetry breaking would happen after inflation and thus would produce domain walls which overclose the universe. An exception to this statement is when the number of degenerate vacua along the bottom of the PQ scalar's Mexican hat potential (the so-called domain wall number) is $N = 1$; then the walls are bounded by cosmic strings without being connected to other walls, and such walls of finite size can annihilate soon after the QCD phase transition~\cite{Vilenkin:1982ks,Sikivie:1982qv,Linde:1990yj,Lyth:1992tx}. Furthermore, if the axion constitutes the CDM, then even tighter bounds on $f_a$ and $H_{\mathrm{inf}}$ are obtained from constraints on CDM isocurvature perturbations~\cite{Seckel:1985tj,Lyth:1989pb}. For instance, with a decay constant of $f_a = 10^{12}\, \mathrm{GeV}$, current constraints on CDM isocurvature from measurements of the cosmic microwave background (CMB)~\cite{Ade:2015lrj} require $H_{\mathrm{inf}} \lesssim 10^7\, \mathrm{GeV}$. The bound on~$H_{\mathrm{inf}}$ is particularly strong for $f_a \lesssim 10^{11}\, \mathrm{GeV}$, due to the nonquadratic form of the axion potential~\cite{Lyth:1991ub}, dubbed the anharmonic effect. The anharmonic enhancement of the axion isocurvature perturbations becomes so strong for small values of~$f_a$ that a decay constant of $f_a \lesssim 10^9\, \mathrm{GeV}$ is excluded~\cite{Kobayashi:2013nva}; thus the lower bound of~(\ref{f-bound}) can also be derived from cosmology when the dark matter consists of axions. These observations indicate that, any detection of primordial gravitational waves from inflation in the near future would exclude the QCD axion as a dark matter candidate. Here it should be noted that the above arguments on $f_a$ and $H_{\mathrm{inf}}$ from domain walls and isocurvature perturbations are modified if the decay constant takes different values during the inflation epoch and today. The works~\cite{Linde:1990yj,Linde:1991km} pointed out that if $f_a$ was larger during inflation, then the axion isocurvature perturbations would actually be smaller than that inferred from the present value of~$f_a$, allowing high-scale inflation without contradicting observations. (See also Refs.~\cite{Higaki:2014ooa,Chun:2014xva,Fairbairn:2014zta,Kearney:2016vqw} for recent works along this line. A similar effect can also be induced by a non-minimal coupling of the axion kinetic term to gravity~\cite{Folkerts:2013tua}.) While this is certainly an intriguing possibility, however, the introduction of a dynamical decay constant not only leads to smaller-than-expected isocurvature perturbations on the CMB scales, but also {\it dynamically enhances} the isocurvature perturbations on smaller length scales. This is understood by viewing the decay constant~$f_a$ and the axion field~$\theta$ respectively as the radial and angular directions of the PQ scalar; if the radial component~$f_a$ is forced to quickly shrink while the angular velocity~$\dot{\theta}$ is nonzero, then the motion along the angular direction would speed up, leading to the growth of the fluctuation~$\delta \theta$ among different regions of space. Moreover, this effect not only works on sub-horizon scales, but further stretches out to scales much larger than the Hubble radius at the time when the decay constant varies.\footnote{A similar effect was invoked in~\cite{Kobayashi:2014sga} for the generation of large-scale magnetic fields in the post-inflationary universe.} In this paper we examine the evolution of the full axion perturbation spectrum under a time-dependent decay constant. We show that, while the dynamical decay constant can help to evade the isocurvature constraints on CMB scales, it can also strongly enhance the axion fluctuations on smaller scales, which may even exceed the periodicity of the axion potential and thus lead to the formation of axionic domain walls. These domain walls are produced due to the enhanced axion fluctuations, although the PQ symmetry is broken before inflation, and thus they are different from the well-studied axionic walls. In particular, the axionic walls from the enhanced fluctuations are not bounded by cosmic strings, and thus would overclose the universe even in the case of $N = 1$~\cite{Preskill:1991kd}. The strong enhancement of the axion fluctuations may also lead to the recovery of the PQ symmetry in the post-inflation epoch, which would upset the cosmological expansion history.\footnote{In this paper we investigate the enhancement of the sub- and super-horizon scale axion fluctuations due to the shrinking of the decay constant. We also note that if the radial component of the PQ scalar oscillates about its minimum, both the axion and the radial component can be produced by parametric resonance. Such sub-horizon fluctuations can lead to a nonthermal symmetry restoration, which has been studied in~\cite{Tkachev:1995md,Kasuya:1996ns,Kasuya:1997ha,Kasuya:1998td,Tkachev:1998dc,Kearney:2016vqw}.} On the other hand, if the axion perturbations are only mildly enhanced, then such small-scale isocurvature could cause a variety of observable consequences, as discussed in e.g.~\cite{Chluba:2013dna,Sekiguchi:2013lma}. Before moving on, we also comment on another possibility to suppress the axionic isocurvature perturbations. If the PQ symmetry is badly broken during inflation, the axion mass may become heavier than or comparable to the Hubble parameter. Then, it does not acquire sizable super-horizon fluctuations, suppressing the isocurvature perturbations~\cite{Dine:2004cq,Higaki:2014ooa,Dine:2014gba,Kawasaki:2015lea,Takahashi:2015waa,Kawasaki:2015lpf,Nomura:2015xil}. In this case, one has to make sure that the explicit PQ symmetry breaking is sufficiently suppressed in the present universe to be consistent with the neutron electric dipole moment experiments. The layout of this paper is as follows: We describe the setup and notation in Section~\ref{sec:setup}. In Section~\ref{sec:dynamics} we give general discussions on the evolution of the axion fluctuations under a dynamical decay constant, then in Section~\ref{sec:DDC} we compute the fluctuations in the post-inflationary universe. In Section~\ref{sec:cosmo}, we discuss the cosmological implications of the enhanced fluctuations, and also carry out case studies. We conclude in Section~\ref{sec:conc}.

\label{sec:conc} While a dynamical decay constant of the QCD axion can help accommodate high-scale inflation, we have shown that it also enhances the small-scale axion isocurvature perturbations. This effect stretches out to super-horizon modes, and thus the enhanced perturbation modes will continue to enter the horizon for some time after the decay constant stops its evolution. The axion fluctuation can become much larger than the potential period, thus unless the fluctuation redshifts away by the time the universe cools down to $T \sim \Lambda_{\mathrm{QCD}}$, it would lead to the formation of axionic domain walls. These walls produced by the enhanced axion fluctuations are not bounded by cosmic strings, and thus would dominate the universe even in the case of $N = 1$. In our analyses we have assumed the decay constant~$f$ to vary as a power-law of the scale factor. However in actual cases the dynamics of~$f$ could be more complicated. If, for instance, $f$ instead oscillates around its final value~$f_0$, then the axion perturbations may get a ``kick'' during the first oscillation, and be enhanced much stronger than in the cases analyzed in this paper. The oscillations of~$f$ could further induce resonant amplifications of the sub-horizon axion perturbations~\cite{Tkachev:1995md,Kasuya:1996ns,Kasuya:1997ha,Kasuya:1998td,Tkachev:1998dc,Kearney:2016vqw}. We also note that, for the evolution of $f$ to take place not too far from the QCD phase transition, a relatively light degree of freedom is required. In a simple case, the mass of the field that sets the axion decay constant~$f$ needs to be much lighter than~$f$ itself. Such a hierarchy in scales is realized in a supersymmetric axion model, where the saxion (corresponding to the radial component~$r$, and thus setting~$f$) acquires a mass only from supersymmetry breaking effects. When the fluctuations of the axion are enhanced by the dynamical $f$, the saxion may also acquire large fluctuations, which results in large spatial fluctuations of $f$. It would be important to study the evolution of perturbations in explicit models of axions with dynamical decay constants. While we have mainly focused on the QCD axion in this paper, similar discussions hold also for axion-like fields with dynamical decay constants, or more generally, for fields that have kinetic terms with time-dependent coefficients. If such a field possesses a periodic potential at energy scales higher than the QCD scale, then even if the evolution of the decay constant happens in the very early times, the enhanced perturbations could still be transformed into domain walls before redshifting away. (However in such cases one may have to take into account the effect of the potential, and in particular the induced mode-mode coupling, during the evolution of the decay constant.) We also note that, depending on when the decay constant evolves, the enhanced axion(-like) isocurvature modes may leave observable signals for upcoming CMB and large scale structure surveys.

1607.04294

1607

1607.03361_arXiv.txt

Cosmic ray accelerators like supernova and hypernova remnants in star forming galaxies are one of the most plausible sources of the IceCube observed diffuse astrophysical neutrinos. The neutrino producing hadronic processes will also produce a diffuse gamma ray flux, constrained by the Fermi-LAT bounds. The fact that point sources like blazars also contribute to the diffuse gamma ray flux implies large gamma opacity of the neutrino sources. Indeed, for these high redshift star forming galaxies the gamma absorption during the intergalactic propagation can be significant. In addition, large gamma attenuation inside these extreme source galaxies can reduce the cascade component of the diffuse flux. Under the current astrophysical uncertainties affecting these absorptions processes, we find the associated diffuse gamma ray flux can remain compatible with the current Fermi-LAT bounds.

\label{sec:intro} The signature of high energy neutrinos in the IceCube (IC) experiment for the first time has opened the window observing neutrinos originating at the cosmic distances~\cite{Aartsen:2013bka,Aartsen:2013jdh,Aartsen:2014gkd,Aartsen:2014muf,Aartsen:2015knd,Aartsen:2015rwa}. The simple explanation of connecting these events to some high energy astrophysical point source remained elusive as the events are found to be compatible with an isotropic distribution~\cite{Aartsen:2014gkd}. Although many astrophysical objects are not compatible with the observed IC flux, the catalog of proposed sources is nevertheless extensive. It ranges from long GRB remnants embedded in molecular clouds~\cite{Dado:2014mea} to active galactic nuclei (AGN)~\cite{Tamborra:2014xia} to intense star forming galaxies~\cite{He:2013cqa,Murase:2013rfa,Liu:2013wia,Chakraborty:2015sta,Senno:2015tra,Xiao:2016rvd}. These IC neutrino events observed at such high energies (TeV to PeV) have another important implication. One of the yet unanswered crucial questions of astroparticle physics is the nature of the interaction of the cosmic rays (CR). The observation of these high energy neutrino events is indirectly pointing towards the hadronic interaction of the CR accelerators as neutrinos are expected to be produced through hadronic processes (via pions from pp or p$\gamma$ interactions). This hadronic origin assumption of the neutrino flux will make the production of gamma rays from neutral pions unavoidable, creating a diffuse background of $\gamma$ ray flux. These multimessenger signatures are of great help to constrain the different astrophysical models compatible with IC high energy neutrino flux. Among these possibilities there is one natural appealing candidate, consisting of stellar remnants (hypernova and supernovae remnants) inside galaxies with large star formation rate, known as star forming galaxies (SFGs). SFGs are distinguished between normal star forming galaxies (NSFGs) and starburst galaxies (SBGs), i.e, galaxies with extraordinary star formation. While hypernovae remnants (HNRs) are able to accelerate protons to hundreds of PeV energies thus generating PeV neutrinos, the more abundant supernovae remnants (SNRs) will dominate the flux in the hundred TeV range. This scenario is not only able to explain the measured flux by IceCube, but also predicts a break on the neutrino spectrum around TeV energies~\cite{Chakraborty:2015sta}. Indeed, the large flux of neutrinos detected in the TeV energies with respect to the flux detected at PeV energies points to different sources of cosmic accelerators in the different energy ranges. The associated gamma ray flux with these sources will also lay in the GeV-TeV range. However, at this energy range, the intergalactic medium is highly opaque as the $\gamma$ rays interaction with the cosmic microwave background (CMB) and the extragalactic background light (EBL) remains significant. Of course, the gamma absorption depends on the distance traveled. SBGs, being distant objects, will result in a very efficient gamma rays absorption. Thus, the highest energy gamma rays cannot reach us from these far objects. However, this flux can reappear in the lower energies. The initial $\gamma \gamma$ collisions in the intergalactic medium also produce a $e+$/$e-$ pairs, which will interact additionally to the EBL photons via the inverse Compton mechanism, resulting in a cascaded low energy $\gamma$ ray flux in the GeV energy range. Thus the total gamma ray flux in the GeV energies is one sensitive measure of all these different interaction processes. Measuring this diffuse $\gamma$ flux is one of the great challenges of the present day physics experiments. One of the main goal of the Fermi-LAT collaboration experiment is to measure this isotropic diffuse gamma ray background (IGRB), in fact the measured extragalactic gamma ray background (EGB) includes the diffuse flux. The known blazar and other $\gamma$ ray sources also contribute to the EGB. Any plausible hadronic model of the IC neutrinos should not overpopulate the bounds of the isotropic diffuse gamma ray background (IGRB) measured by the Fermi-LAT collaboration~\cite{TheFermi-LAT:2015ykq}. For a recent analysis of observed SFG-SBG point sources and their gamma ray-neutrino correlation, see~\cite{Moharana:2016mkl}. The large number of neutrinos detected in the TeV energies have started putting pressure on this multimessenger picture as these larger fluxes demand a greater diffuse $\gamma$ background. Recently, the Fermi collaboration did a survey of the blazar sources in the high energy tail of the observed Fermi EGB spectra~\cite{Ackermann:2014usa}. The stacking of all the blazar point sources gives the limit of diffuse gamma rays in GeV energies. In the energy interval of $50$ GeV to $1$ TeV this limit is found to be very close to the observed diffuse background, leaving no more than the 14\% percent of the observed fluxes to the hadronic channel gamma rays produced with neutrinos. This apparent tension in the picture is also pointing to a class of sources that are opaque in the gamma rays~\cite{Wang:2015mmh, Murase:2015xka} etc. Recently, another study~\cite{Lisanti:2016jub} for these known gamma sources found that these limits can actually be more relaxed, i.e, 32 $\%$ non blazar gamma ray contribution. This is due to the fact that the percentage of the blazar component in EGB still has big uncertainties. However, there is no doubt that they own a significant fraction of the total EGB. Note that, in the case of blazars, the leptonic models are found to be successful in explaining the source gamma spectra~\cite{LAT:2011aa}. There is no co-production of neutrinos in these models. Of course, invoking a hadronic model with subsequent neutrino production for the blazar gammas could relieve this tension. In the following, we focus on the SBGs and calculate the limits coming from the different components of the diffuse gamma and neutrino spectra. The large dimensions of the early galaxies with large number of background photons make sure that the gamma rays with higher energies (1-10 TeV) cannot escape. Thus the cascade component of the gamma rays coming from the high energy part can get greatly attenuated. The lower energy component with ordinary intergalactic absorption can still reach the detectors. However, the flux depends on several free parameters and the estimation of the EBL. In this work, we try to understand the broad picture, independent of the parameter uncertainties. We find that the neutrino flux from this model explaining the IC events can still be accommodated in view of the recent Fermi-LAT non-blazar EGB estimate. Under the present uncertainties of the EBL and considering the intragalactic absorption, the hadronic gammas can remain under the observed non-blazar EGB gamma limit. Thus not only the parameter space describing this model can get narrowed down but also the absorption models of gamma rays can get constrained. \cite{TheFermi-LAT:2015ykq,Bechtol:2015uqb}. The outline of this paper is the following. In section~\ref{sec:flux} we briefly describe the $\nu$ and the $\gamma$ flux production by the stellar remnants (SRs) embedded in the star forming galaxies. In section~\ref{sec:diffga} we describe the details of gamma ray spectrum of SFGs, both the cascading and non cascading component. Then, section~\ref{sec:intra} , describes the absorption of the $\gamma$ rays inside the source galaxies and their effect on the total diffuse gamma ray flux. % In section~\ref{sec:gamma} we discuss the different uncertainties in the estimation of the gamma spectra. Finally, we comment and conclude in section ~\ref{sec:conclus}.

\label{sec:conclus} Stellar remnants in normal star forming galaxies and star burst galaxies have been proven to provide a diffuse neutrino flux compatible with the results obtained in the IceCube experiment. The existence of two different kind of accelerators SNRs and HNRs in those galaxies also give rise to the characteristic break in the neutrino spectrum. Nevertheless, because of the hadronic origin of the neutrino flux associated gamma ray production would be unavoidable. In light of the recent studies of the diffuse gamma spectra of Fermi-LAT and the blazar contribution to this diffuse spectra a tension between the diffuse gamma and neutrino spectra has been suggested. We have studied how the different models for the intergalactic absorption affect the diffuse gamma ray spectrum. Moreover, very importantly we have added the internal absorption of gamma rays present in SBGs to show how this modify the spectrum in the GeV-TeV energies. In particular, the cascading part of the diffuse gamma spectra gets a major reduction from this galactic absorption. Combining these two effects, we have shown that under reasonable assumption of the astrophysical uncertainties this diffuse $\gamma$ ray flux can be compatible with the current experimental limits. The main conclusion from this study is therefore that the diffuse gamma ray flux associated to this SBG models cannot be used to rule out this scenario that stellar remnants embedded in SBGs are the main source of the diffuse IC neutrino spectra. This work is a reassurance of this previously presented models. Therefore, the neutrino diffuse flux spectrum footprint predicting a spectra break around 100 TeV is a better tool to test these models describing the origin of the cosmic neutrinos. This is based on the fact that the weakly interacting neutrinos do not suffer from the above mentioned absorptions and one can avoid the uncertainties introduced by the EBL density and the SBGs internal absorption. Indeed, it would be interesting to investigate the IceCube data accumulation time frame and the statistical significance necessary to confirm or to rule out this spectral break scenario. Nevertheless, as shown in this work, it is always essential to provide a consistent multimessenger picture.

1607.03361

1607

1607.01364_arXiv.txt

The light variations of first-overtone RR Lyrae stars and contact eclipsing binaries can be difficult to distinguish. The Catalina Periodic Variable Star catalog contains several misclassified objects, despite the classification efforts by \citet{Drakeetal2014}. They used metallicity and surface gravity derived from spectroscopic data (from the SDSS database) to rule out binaries. Our aim is to further constrain the catalog using SDSS colours to estimate physical parameters for stars that did not have spectroscopic data.

Briefly, \textsc{Photo-Met} estimates the unknown physical parameters of stars by interpolating the known parameters of other stars that have very similar broad-band colours. The method relies on two numerical algorithms: \\ \noindent - efficient $k$-nearest-neighbour finding in a four-dimensional metric colour-colour space and \\ \noindent - local linear regression. As with any empirical parameter estimation algorithm, the reliability of the entire process depends much more on the training set than on the actual numerical method. For this study, we used an empirical training set based on SDSS PSF magnitudes of approximately 360\,000 stars. The stellar parameters [Fe/H], $T_{\rm eff}$ and $\log g$ are the adopted weighted averages from SSPP (SEGUE Stellar Parameter Pipeline). For further details on the method see \citet{Kerekesetal2013}. The variable list of the Catalina Sky Survey consists of 5467 stars marked as RRc. The cross-match with the SDSS DR10 Cross-ID tool using an 1\farcs2 radius search resulted in 2762 stars. For delivering new results we took the stars with photometric measurements only (1732 objects) and applied the \textsc{Photo-Met} method to estimate the surface gravity, the effective temperature and metallicity.

We could identify several contact binaries that were originally classified as RRc stars in the CSS Periodic Variable Star Catalog. The \textsc{Photo-Met} method, however, clearly placed them outside the RRc domain set by \citet{Drakeetal2014}. We plotted in Fig.~1 four light curves where the asymmetry of the minima can be clearly seen when folded with twice the assumed variation period. However, despite our success in finding new contaminating binary stars in the sample, the large errors in the $\log g$ and [Fe/H] determination (up to $\pm 1.0$ dex) and the low quality of the light curves of faint stars make this method uncertain, leaving a lot of ambiguous objects in the catalog. \begin{figure}[!ht] \includegraphics[width=0.95\textwidth]{Banyai-fig1.eps} \vspace*{-3mm} \caption{Folded light curves of newly identified eclipsing binaries in the CSS RRc sample. Blue and red points are shifted in phase by 0.5 with respect to each other to highlight the differing minima.} \label{authorsurname-fig1} \end{figure} \vspace*{-3mm}

1607.01364

1607

1607.08245_arXiv.txt

We present evidence of bar-induced secular evolution in galactic discs using 3.6 {\mum} images of nearby galaxies from the {\it Spitzer} Survey of Stellar Structure in Galaxies ({\s4g}). We find that among massive galaxies ($M_{\ast}/${\Msun}$> 10^{10}$), longer bars tend to reside in inner discs having a flatter radial profile. Such galaxies show a light deficit in the disc surrounding the bar, within the bar radius and often show a $\Theta$-shaped morphology. We quantify this deficit and find that among all galaxies explored in this study (with $10^{9}<M_{\ast}/${\Msun}$< 10^{11}$), galaxies with a stronger bar (i.e. longer and/or with a higher Bar/T) show a more pronounced deficit. We also examine simulation snapshots to confirm and extend results by Athanassoula and Misiriotis, showing that as bars evolve they become longer, while the light deficit in the disc becomes more pronounced. Theoretical studies have predicted that, as a barred galaxy evolves, the bar captures disc stars in its immediate neighbourhood so as to make the bar longer, stronger and thinner. Hence, we claim that the light deficit in the inner disc is produced by bars, which thus take part in shaping the mass distribution of their host galaxies.

The evolution of galaxies is initially driven by fast and violent processes at early times. Later on, however, as the merger rate drops, slow and secular processes start to become dominant (e.g., \citealt{kormendy_04}). One of the major contributors that drive internal secular evolution in disc galaxies are non-axisymmetric structures, such as bars (\citealt{athanassoula_13_book, kormendy_13_book, sellwood_14_rev} for reviews covering the theoretical and observational perspectives). Bars are common in nearby disc galaxies (a fraction of 50$\sim$ 70 per cent, e.g., \citealt{buta_15} and references therein, a fraction of $\sim$30 per cent when only strong bars are counted. e.g., \citealt{masters_11, buta_10, buta_15} and references therein). Numerical and analytic studies show that once disc galaxies are massive enough and rotation-dominated, the bar instability develops relatively fast, within a few hundred Myr (e.g., \citealt{pfenniger_91, friedli_93, athanassoula_02a}, hereafter AM02; \citealt{martel_13}). However, the bar formation can be delayed if the disc is dispersion-dominated, if the initial dark matter halo is dominant, or if the galaxy initially contains a large amount of gas (\citealt{athanassoula_86}; AM02; \citealt{athanassoula_03, athanassoula_13}). Thus the fraction of barred galaxies provides us statistics on how many galaxies already have sufficiently massive discs dominated by rotation (\citealt{sheth_12}). Bar fractions appear to change with redshift: from $z=0.8$ to $z=0.2$, overall bar fractions increase from 20 to 65 per cent, while strong bar fractions increase from 10 to 30 per cent (\citealt{sheth_08, cameron_10}, see also \citealt{melvin_14}). The lower bar fraction at high redshift is mainly due to the lower mass galaxies not yet having developed bars (\citealt{sheth_12}). At z $>$1, there are tentative bar detections in the near infrared data (\citealt{sheth_03}), and the bar fractions appears to be $\sim 10$ per cent at z=1.5$\sim$2 (\citealt{simmons_14}). Note that finding bars is subject to limits due to image resolution and band-shifting. Cosmological simulations suggest that bars formed during the violent phase (z$>$1) may be easily destroyed or may become too weak to be observed, whereas bars formed in the stellar disc at the secular phase ($z<0.8$) are predicted to be generally robust and long-lived. The fraction of bars increases with time (\citealt{kraljic_12, martig_12}), consistent with the observed trend. It is challenging to gauge when bars are formed, as the age of stellar populations building up a bar does not necessarily refer to the formation epoch of the bar itself. Nevertheless, several attempts have been made (e.g., \citealt{gadotti_05, wozniak_07, perez_07, perez_11, elmegreen_09, sanchez_blazquez_11, delorenzo_caceres_13, james_16}). The age of bars seems to be correlated to the mass and dynamical status of their parent galaxies. Massive, and rotation dominated galaxies form their bars first (\citealt{sheth_08, sheth_12}). Recent studies find that some galaxies have hosted their bars for a long time. An integral field spectrograph study on NGC 4371 reveals that the inner disc and nuclear ring -- thought to be composed of stars formed from gas funnelled by the bar -- are mainly composed of old ($>10$ Gyr) stellar populations, and indicate that the bar was formed at $z\sim1.8$ (\citealt{gadotti_15}). Although not a direct measure of bar age, radial light profiles of bars have been used to infer the dynamical age of bars (\citealt{kim_15a}). Because bars are formed from disc material, the radial profiles of bars at an early evolutionary stage would be similar to those of discs, which are exponential (AM02). However, bars with exponential profiles are not necessarily young (AM02). Whereas, bars that show flat radial profiles are expected to be strong, and therefore dynamically old. Bars in massive and bulge-dominated galaxies are found to show flat radial profiles (\citealt{kim_15a}). Numerical simulations find that once bars are formed, it is difficult to dissolve them (e.g., \citealt{shen_04, athanassoula_05b, athanassoula_13, debattista_06, berentzen_07, villa_vargas_10, kraljic_12, martig_12}). Hence, bars must have been influencing their host galaxies since they are formed, and have an extended impact in the evolution of galaxies (\citealt{gadotti_15}). Therefore, it is important to explore the impact of bar driven secular evolution on their host galaxies. Because of their non-axisymmetry, bars induce large-scale streaming motions (e.g., \citealt{athanassoula_92a}). Observational studies find enhanced central gas concentrations in barred galaxies (\citealt{regan_97, sakamoto_99, sheth_00, zurita_04, jogee_05}). This leads to barred galaxies also showing enhanced star formation in the central region (e.g, \citealt{ho_97, ellison_11}), and pronounced nuclear rings (\citealt{knapen_02, comeron_10, kim_w_12a, seo_13}). Eventually, bars induce the formation of discy pseudo bulges in their host galaxies (\citealt{kormendy_04, sheth_05, athanassoula_05a, debattista_06, cheung_13}). The impact of bars on disc galaxies is not only limited to the central part of the galaxy. Bar torques bring the gas inside the corotation inwards, while pushing the gas between the corotation radius and the outer Lindblad resonance radius outwards (\citealt{combes_85, combes_08_gas, kubryk_13}). Barred galaxies are often accompanied by an outer ring where one of the bar resonances is expected to be located (\citealt{schwarz_81, buta_96, buta_03, buta_15, romero_gomez_06, athanassoula_09a}). Thus bars drive secular evolution in their host galaxies, slowly re-arranging mass and angular momentum distributions throughout the different galactic components. Numerical simulations predict that bar-induced angular momentum redistribution leads discs to show a break in their radial density profile (\citealt{valenzuela_03, debattista_06}) with a shallower inner disc and a steeply decreasing outer disc. Observational data also show that, compared to unbarred galaxies, barred ones show larger global disc scale length and fainter central surface brightness of discs among massive galaxies (\citealt{sanchez_janssen_13, diaz_garcia_16xx}). The majority of disc galaxies are found to show at least one disc break (e.g., \citealt{pohlen_02, pohlen_06, erwin_08, gutierrez_11, maltby_12a, munoz_mateos_13, kim_14, laine_14}; for edge-on galaxies, see also \citealt{comeron_12, martin_navarro_12}), and this has been confirmed by simulations (e.g., \citealt{roskar_08a, athanassoula_16b}). However, it should be noted that not all disc breaks are produced by bars, and even unbarred galaxies may show disc breaks. Apart from the bar-driven one, several other mechanisms responsible for disc breaks have been proposed (e.g., \citealt{vanderkruit_87, tagger_87, kennicutt_89, laurikainen_01, elmegreen_07, younger_07, roskar_08b, minchev_12a, munoz_mateos_13}). If a galaxy has a disc break, there are two different disc scale lengths: the inner and the outer disc scale length. Compared to the global disc scale length which ignored the disc break, the inner and outer disc scale lengths differ by $\sim40$ per cent on average (\citealt{kim_14}). It is therefore important to separate a disc into its inner and outer parts. \cite{laine_14} find that the average disc profile of Type I is similar to the average outer disc profile of Type II, while inner disc profiles of Type II are flatter than disc profiles of Type I, and \citet{laine_14} attribute this difference to the effects of bars. A morphology often seen in barred galaxies is well represented by the so-called $\Theta$-shaped galaxies, which show a light deficit in the disc surrounding the bar at radii smaller than the radius of the inner ring (i.e. the ring that is often present near the ends of the bar), and thus within the bar radius. \citet{gadotti_03} have studied two of such barred galaxies (NGC 4608 and NGC 5701; see also \citealt{laurikainen_05, gadotti_08}). We show more examples in Fig.~\ref{fig:deficit_light}, in which arrows indicate the disc light deficit around the bar within the bar radius (note that the presence of an inner ring is not a necessary condition for the occurrence of the light deficit). Theoretical work also finds light deficits in simulated galaxies (e.g., AM02; \citealt{athanassoula_13}). In particular, a model with a more centrally concentrated halo shows a more prominent light deficit, as well as a stronger, longer and thinner bar, as compared to a model with a less centrally concentrated halo (AM02). Theoretical studies have also predicted that bars give up angular momentum as they evolve, which leads to several changes in bar properties that might also affect the disc. \citet{athanassoula_03}, considering the bar as an ensemble of orbits, presents schematically three possible changes in its orbital structure. Firstly, the bar traps stars which were initially on quasi-circular orbits just outside the bar. The new orbits are elongated and thus the bar becomes longer and/or more massive, to the detriment of the surrounding disc. Secondly, orbits in the bar become more elongated, i.e. the bar becomes thinner. Lastly, the bar slows down by lowering its pattern speed. These three possible processes are closely linked together and occur simultaneously. In summary, simulations predict that as bars evolve disc stars are captured onto bar orbits. With the help of these newly captured stars the bar becomes longer and more massive. Thus, since stars are removed from the disc and added to the bar, the inner part of the disc surrounding the bar (i.e. at galactocentric radii below the bar radius, or $r< R_{\rm bar}$) should in principle become less dense. Hence, we should expect that there will be a deficit of light from the disc surrounding the bar, and, as the galaxy evolves, this light deficit in the inner disc will become more pronounced. However, this effect has not yet been explored with observational datasets and no direct comparisons with simulations have been reported either. In this study, we test this hypothesis by checking whether there is a relation between the bar and the {\it inner} disc properties through detailed structural analysis, quantifying the light deficit in both an observational dataset and a set of simulation snapshots. The paper is organized as follows. In \S 2 we give a brief overview of our data and data analysis. Results on the impact of bar-driven secular evolution on discs are presented in \S 3. We explore how bar properties change with time using the snapshots of a simulation in \S 4. We discuss our results in \S 5, and summarize and conclude in \S 6. \begin{figure} \centering \includegraphics[width=7cm]{fig1.eps} \caption{(Left): Images of NGC 1015 and IC 1438 at 3.6 {\mum}. Red arrows indicate the deficit of light from the inner disc surrounding the bar. Horizontal bars in the bottom left corners of these panels span 30 arcsec. North is up and east is to the left. (Right): Residual images using the models fitted in Paper I highlight the light deficits in the inner discs. } \label{fig:deficit_light} \end{figure}

Bars act as a driving force for the evolution of their host galaxies. With the aim of assessing the impact of bar-driven secular evolution on discs, we used 3.6 {\mum} images of 118 nearby barred galaxies from the {\s4g} with type II (down-bending) radial surface brightness profiles. We investigated how the properties of bars are related to those of the inner parts of their host discs. In particular, we investigated the origin of the light deficits often observed in the part of inner discs surrounding the bar, within the bar radius \citep[see e.g.][for earlier discussions]{gadotti_08}. Our main results can be summarized as follows: \begin{itemize} \item Among massive galaxies with a prominent bar (Bar/T$>0.1$), there is a clear trend that longer bars reside in more flattened inner discs (larger inner disc scale length and lower central surface brightness) than shorter bars do. Such galaxies often show the light deficit around the bar in the inner part of the disc. \item To better understand the relation between the bar and the light deficit in inner discs, we quantify the light deficit. We measure the maximum difference between the surface brightness profiles along the bar major and minor axes, Max($\Delta \mu$). As it measures the light above the disc, Max($\Delta \mu$) is a measure of the bar prominence. Because bars evolve by capturing disc stars, Max($\Delta \mu$) is also a indicator for the light deficit in the inner disc. We find that Max($\Delta \mu$) is strongly related to the bar size and to Bar/T, so that the light deficit is directly proportional to bar size and to how conspicuous the bar is. \item By studying a time sequence of snapshots from the evolution of a simulated barred galaxy, we find that as the bar evolves, it becomes longer and the light deficit in the inner disc becomes more pronounced. This can be understood by the fact that as a barred galaxy evolves, the bar loses angular momentum and becomes longer and more massive by trapping nearby disc stars onto bar orbits. Therefore, the light deficit is produced as a consequence of the capture of disc stars by the bar, which are thus removed from the inner part of discs. \end{itemize} The observed correlations between the light deficit and bar size and Bar/T (Fig.~\ref{fig:max_diff}) are consistent with the picture drawn from the analysis of the evolution of a simulated barred galaxy (Fig.~\ref{fig:gtr116_img}), in that bars grow longer and more conspicuous by capturing nearby disc stars. Based on these results, we therefore propose that the light deficit often observed in the part of the inner discs within the bar radius is produced by bars. This is direct evidence for bar-driven secular evolution in galactic discs, and a strong indication that bars are actively involved in shaping the mass distribution of their host galaxies.

1607.08245

1607

1607.01014_arXiv.txt

Recent observational progress has led to the establishment of the standard $\Lambda$CDM model for cosmology. This development is based on different cosmological probes that are usually combined through their likelihoods at the latest stage in the analysis. We implement here an integrated scheme for cosmological probes, which are combined in a common framework starting at the map level. This treatment is necessary as the probes are generally derived from overlapping maps and are thus not independent. It also allows for a thorough test of the cosmological model and of systematics through the consistency of different physical tracers. As a first application, we combine current measurements of the Cosmic Microwave Background (CMB) from the Planck satellite, and galaxy clustering and weak lensing from SDSS. We consider the spherical harmonic power spectra of these probes including all six auto- and cross-correlations along with the associated full Gaussian covariance matrix. This provides an integrated treatment of different analyses usually performed separately including CMB anisotropies, cosmic shear, galaxy clustering, galaxy-galaxy lensing and the Integrated Sachs-Wolfe (ISW) effect with galaxy and shear tracers. We derive constraints on $\Lambda$CDM parameters that are compatible with existing constraints and highlight tensions between data sets, which become apparent in this integrated treatment. We discuss how this approach provides a complete and powerful integrated framework for probe combination and how it can be extended to include other tracers in the context of current and future wide field cosmological surveys.

Introduction} The past two decades have seen immense progress in observational cosmology that has lead to the establishment of the $\Lambda$CDM model for cosmology. This development is mainly based on the combination of different cosmological probes such as the CMB temperature anisotropies, galaxy clustering, weak gravitational lensing, supernovae and galaxy clusters. Until now, these probes have been, for the most part, measured and analysed separately using different techniques and combined at late stages of the analysis, i.e. when deriving constraints on cosmological parameters. However, this approach is not ideal for current and future surveys such as the Dark Energy Survey (DES\footnote{\tt{http://www.darkenergysurvey.org}.}), the Dark Energy Spectroscopic Instrument (DESI\footnote{\tt{http://desi.lbl.gov}.}), the Large Synoptic Survey Telescope (LSST\footnote{\tt{http://www.lsst.org}.}), Euclid\footnote{\tt{http://sci.esa.int/euclid/}.} and the Wide Field Infrared Survey Telescope (WFIRST\footnote{\tt{http://wfirst.gsfc.nasa.gov}.}) for several reasons. First, these surveys will cover large, overlapping regions of the observable universe and are therefore not statistically independent. In addition, the analysis of these surveys requires tight control of systematic effects, which might be identified by a direct cross-correlation of the probes statistics. Moreover, each probe provides a measurement of the cosmic structures through a different physical field, such as density, velocity, gravitational potentials, and temperature. A promising way to test for new physics, such as modified gravity, is to look directly for deviations from the expected relationships of the statistics of the different fields. The integrated treatment of the probes from the early stages of the analysis will thus provide the cross-checks and the redundancy needed not only to achieve high-precision but also to challenge the different sectors of the cosmological model. Several earlier studies have considered joint analyses of various cosmological probes. \citet{Mandelbaum:2013, Cacciato:2013} and \citet{Kwan:2016} for example derived cosmological constraints from a joint analysis of galaxy-galaxy lensing and galaxy clustering while \citet{Liu:2016} used the cross-correlation between the galaxy shear field and the overdensity field together with the cross-correlation of the galaxy overdensity with CMB lensing to constrain multiplicative bias in the weak lensing shear measurement in CFHTLenS. Recently, \citet{Singh:2016} performed a joint analysis of CMB lensing as well as galaxy clustering and weak lensing. Furthermore, \citet{Eifler:2014} and \citet{Krause:2016} have theoretically investigated joint analyses for photometric galaxy surveys by modelling the full non-Gaussian covariance matrix between cosmic shear, galaxy-galaxy lensing, galaxy clustering, photometric baryon acoustic oscillations (BAO), galaxy cluster number counts and galaxy cluster weak lensing. Extending beyond this, we present and implement an integrated approach to probe combination. In this first implementation we combine data from CMB temperature anisotropies, galaxy overdensities and weak lensing. We use data from Planck 2015 \cite{Planck-Collaboration:2015af} for the CMB, for galaxy clustering we use photometric data from the 8$^{\mathrm{th}}$ data release of the Sloan Digital Sky Survey (SDSS DR 8) \cite{Aihara:2011} and the weak lensing shear data comes from SDSS Stripe 82 \cite{Annis:2014}. We combine these probes into a common framework at the map level by creating projected 2-dimensional maps of CMB temperature, galaxy overdensity and the weak lensing shear field. In order to jointly analyse this set of maps we consider the spherical harmonic power spectra of the probes including their cross-correlations. This leads to a spherical harmonic power spectrum matrix that combines CMB temperature anisotropies, galaxy clustering, cosmic shear, galaxy-galaxy lensing and the ISW \cite{Sachs:1967} effect with galaxy and weak lensing shear tracers. We combine this power spectrum matrix together with the full Gaussian covariance matrix and derive constraints on the parameters of the $\Lambda$CDM cosmological model, marginalising over a constant linear galaxy bias and a parameter accounting for possible multiplicative bias in the weak lensing shear measurement. In this first implementation, we use some conservative and simplifying assumptions. For instance we include a limited range of angular scales for the different probes to reduce our sensitivity to systematics, nuisance parameters and nonlinear corrections. With this, we work under the assumption of Gaussian covariance matrices and with a reduced set of nuisance parameters. This paper is organised as follows. In Section \ref{sec:framework} we describe the framework for integrated probe combination employed in this work. The theoretical modelling of the cosmological observables is summarised in Section \ref{sec:theorypred}. Section \ref{sec:maps} describes the data analysis for each probe, especially the map-making procedure. The computation of the spherical harmonic auto- and cross-power spectra is discussed in Section \ref{sec:cls} and the estimation of the covariance matrix is detailed in Section \ref{sec:covariance}. In Section \ref{sec:results} we present the cosmological constraints derived from the joint analysis and we conclude in Section \ref{sec:conclusions}. More detailed descriptions of data analysis as well as robustness tests are deferred to the Appendix.

Conclusions} To further constrain our cosmological model and gain more information about the dark sector, it will be essential to combine the constraining power of different cosmological probes. This work presents a first implementation of an integrated approach to combine cosmological probes into a common framework at the map level. In our first implementation we combine CMB temperature anisotropies, galaxy clustering and weak lensing shear. We use CMB data from Planck 2015 \cite{Planck-Collaboration:2015af}, photometric galaxy data from the SDSS DR8 \cite{Aihara:2011} and weak lensing data from SDSS Stripe 82 \cite{Annis:2014}. We take into account both the information contained in the separate maps as well as the information contained in the cross-correlation between the maps by measuring their spherical harmonic power spectra. This leads to a power spectrum matrix with associated covariance, which combines CMB temperature anisotropies, galaxy clustering, cosmic shear, galaxy-galaxy lensing and the ISW \cite{Sachs:1967} effect with galaxy and weak lensing shear tracers. From the power spectrum matrix we derive constraints in the framework of a $\Lambda$CDM cosmological model assuming both a Gaussian covariance as well as a Gaussian likelihood. We find that the constraints derived from the combination of all probes are significantly tightened compared to the constraints derived from each of the three separate auto-power spectra. This is due to the complementary information carried by different cosmological probes. We further compare these constraints to existing ones derived by the Planck collaboration and find reasonable agreement, even though the joint analysis slightly prefers lower values of both $\Omega_{\mathrm{m}}$ and $\Omega_{\mathrm{b}}$ and a higher Hubble parameter $h$. For a joint analysis of three cosmological probes, the constraints derived are still relatively weak, which is mainly due to our conservative cuts in angular scales. Nevertheless this analysis already demonstrates the potential of integrated probe combination: the complementarity of different data sets, that alone yield rather weak constraints on the full $\Lambda$CDM parameter space, allows us to obtain robust constraints which are significantly tighter than those obtained from probes taken individually. In addition, our analysis reveals challenges intrinsic to probe combination. Examples are the need for foreground-correction at the map as opposed to the power spectrum level and the need for coordinate-independent bias corrections. In this first implementation we have made simplifying assumptions. We assume a Gaussian covariance matrix for all cosmological probes considered. This is justified for the CMB temperature anisotropies and the galaxy overdensity at large scales. The galaxy shears on the other hand exhibit non-linearities already at large scales and their covariance therefore receives significant non-Gaussian contributions \cite{Sato:2009}. Furthermore, we do not take into account the cosmology-dependence of the covariance matrix \cite{Eifler:2009}. In addition we only include systematic uncertainties from a potential multiplicative bias in the weak lensing shear measurement and neglect effects from other sources. Finally we also used the Limber approximation for the theoretical predictions. We leave these extensions to future work but we do not expect them to have a significant impact on our results since we restrict the analysis to scales where the above effects are minimised. In order to fully exploit the wealth of cosmological information contained in upcoming surveys, it will be essential to investigate ways in which to combine these experiments. It will be thus interesting to extend the framework presented here to include additional cosmological probes, 3-dimensional tomographic information and tests of cosmological models beyond $\Lambda$CDM.

1607.01014

1607

1607.01825_arXiv.txt

The Atacama Cosmology Telescope Polarimeter (ACTPol) is a polarization sensitive upgrade to the Atacama Cosmology Telescope, located at an elevation of 5190 m on Cerro Toco in Chile. ACTPol uses transition edge sensor bolometers coupled to orthomode transducers to measure both the temperature and polarization of the Cosmic Microwave Background (CMB). Calibration of the detector angles is a critical step in producing polarization maps of the CMB. Polarization angle offsets in the detector calibration can cause leakage in polarization from E to B modes and induce a spurious signal in the EB and TB cross correlations, which eliminates our ability to measure potential cosmological sources of EB and TB signals, such as cosmic birefringence. We calibrate the ACTPol detector angles by ray tracing the designed detector angle through the entire optical chain to determine the projection of each detector angle on the sky. The distribution of calibrated detector polarization angles are consistent with a global offset angle from zero when compared to the EB-nulling offset angle, the angle required to null the EB cross-correlation power spectrum. We present the optical modeling process. The detector angles can be cross checked through observations of known polarized sources, whether this be a galactic source or a laboratory reference standard. To cross check the ACTPol detector angles, we use a thin film polarization grid placed in front of the receiver of the telescope, between the receiver and the secondary reflector. Making use of a rapidly rotating half-wave plate (HWP) mount we spin the polarizing grid at a constant speed, polarizing and rotating the incoming atmospheric signal. The resulting sinusoidal signal is used to determine the detector angles. The optical modeling calibration was shown to be consistent with a global offset angle of zero when compared to EB nulling in the first ACTPol results and will continue to be a part of our calibration implementation. The first array of detectors for Advanced ACTPol, the next generation upgrade to ACTPol, will be deployed in 2016. We plan to continue using both techniques and compare them to astrophysical source measurements for the Advanced ACTPol polarization calibration.

\label{sec:intro} The Atacama Cosmology Telescope (ACT) is an off-axis Gregorian telescope constructed in 2007 \cite{2007ApOpt..46.3444F}. The Atacama Cosmology Telescope Polarimeter (ACTPol) is a polarization sensitive receiver upgrade to ACT. Starting in 2013, in a staged deployment, ACTPol began observing the Cosmic Microwave Background (CMB). In the most recent observation season ACTPol used ${\sim}3{,}000$ transition edge sensor bolometers, in two frequency bands across three arrays \cite{2016arXiv160506569T}. The primary elements of the ACTPol optical chain are a 6 m primary mirror, a 2 m secondary mirror and three optics tubes, each of which contains a set of three silicon reimaging lenses. Other elements include the receiver window, several band defining filters and a set of corrugated feedhorns per array. Figure \ref{fig:ray_trace} shows the optical chain with a simple ray trace for one optics tube. \begin{figure} [ht] \begin{center} \begin{tabular}{c} \includegraphics[width=0.5\linewidth]{telescope_raytrace_2_merged.jpg} \end{tabular} \end{center} \caption{\label{fig:ray_trace} Single field ray trace of the ACTPol optics. Shown here is a simple ray trace from CODE V, with a single field point on the sky propagated to the primary mirror, secondary mirror, and through the refractive silicon optics in the ACTPol receiver for a single optics tube. The focal plane for one array is the final element, where the ACTPol detectors lie.} \end{figure} The corrugated feedhorns each couple to an orthomode transducer (OMT) which separates incoming light into orthogonal polarizations. The angle of the two OMT fins are defined by lithography in fabrication. The ACTPol detectors are fabricated on 3-inch wafers which were etched into hexagonal and partial hexagonal shapes (referred to as ``hexes'' and ``semi-hexes'') and tiled in an array of three hexes and three semi-hexes each. The orthogonal antenna probes on each wafer are oriented at 0/90 degrees and 45/135 degrees, such that the full array of six wafers has sets of detectors ranging from 0 to 180 degrees at 15 degree intervals. The initial set of calibration angles are determined by the photolithography of the detector wafers and their placement into the array. We then add an array-specific angle when the array is installed into the receiver cryostat. This angle is constrained mechanically, hereafter referred to as the ``installation angle.'' Finally, the optics chain itself (the reflectors and lenses) causes a position dependent polarization rotation, which we solve for using the optical design software CODE V\footnote[1]{Synopsys Optical Solutions Group -- https://optics.synopsys.com/codev/}. All three of these angles are combined to produce the final polarization angle calibration for ACTPol.

We have outlined the steps taken to calibrate the ACTPol detector polarization angles. Working from the fabricated design angles we use planet observations to determine the installation angle and match these to our model for the optical distortions. Finally, we add the additional polarization rotation determined from optical modeling. This calibration technique has produced detector angles consistent with a global offset angle from zero when compared to the EB-nulling offset angle and is currently used to calibrate the ACTPol detectors. A technique for measuring the polarization rotation due to the reimaging optics is also described. We observe a rotating polarization grid, fit the detector time streams to determine the detector angles and take the difference of this measured angle with the expected detector angle. The measured angles demonstrate the potential for this technique to constrain polarization rotation present in the reimaging optics; however, improvements need to be made with respect to the uniformity of the polarization grid to achieve sub-degree measurements of individual detector angles. \vspace{-0.05in}

1607.01825

1607

1607.06634_arXiv.txt

{There is evidence in the 125-Myr Pleiades cluster, and more recently in the 5-Myr NGC\,2264 cluster, that rotation plays a key role in the Lithium (Li) depletion processes among low-mass stars. Fast rotators appear to be less Li-depleted than equal-mass slow rotators.} {We intend to explore the existence of a Li depletion - rotation connection among the $\beta$ Pictoris members at an age of about 24 Myr, and to use such correlation either to confirm or to improve \rm the age estimate based on the Lithium Depletion Boundary (LDB) modeling.} {We have photometrically monitored all the known members of the $\beta$ Pictoris association with at least one Lithium equivalent width (Li EW) measurement from the literature.} {We measured the rotation periods of 30 members for the first time and retrieved from the literature the rotation periods for other 36 members, building a catalogue of 66 members with measured rotation period and Li EW.} {We find that in the 0.3 $<$ M $<$ 0.8 M$_\odot$ range, there is a strong correlation between rotation and Li EW. For higher mass stars, no significant correlation is found. For very low mass stars in the Li depletion onset, at about 0.1 M$_\odot$, data are too few to infer a significant correlation. The observed Li EWs are compared with those predicted by the Dartmouth stellar evolutionary models that incorporate the effects of magnetic fields. After decorrelating the Li EW from the rotation period, we find that the hot side of the LDB is fitted well by Li EW values corresponding to an age of 25$\pm$3 Myr in good agreement with independent estimates from the literature.}

In recent years much attention has been payed to the 24$\pm$3\,Myr \citep{Bell15} young \object{$\beta$ Pictoris} stellar association. Several studies have allowed to increase significantly the number of confirmed members and to discover many more candidate members. Just to mention the most recent works, we refer the readers to \citet{Lepine09}; \citet{Kiss11}; \citet{Schlieder10, Schlieder12}; \citet{Shkolnik12}; \citet{Malo13, Malo14a, Malo14b}. The first comprehensive search for the rotation periods of the low-mass members of $\beta$ Pictoris was carried out by \citet{Messina10, Messina11} who retrieved a total of 38 low-mass members (i.e., spectral types from late F to M) from the earlier compilations of \citet{Zuckerman04}, \citet{Torres06}, and \citet{Kiss11}. Their study provided the rotation periods of 33 out of 38 members.\\ In the light of the mentioned studies, we have again explored up to the most recent literature and, at this time, we could finally compile a new list of 117 stars \rm among members and candidate members with spectral types later than about F3V. Then, we have started a new rotational study on this enlarged sample. To get the photometric rotation periods of our targets, we used our own observations, archive data, and also we made use of periods from the literature.\\ As result of our photometric investigation, we obtained the rotation periods of 112 out of 117 stars. \rm Specifically, we measured for the first time the rotation periods of 56 stars. \rm For another 27 stars, we could confirm with our analysis of new or archived data the values reported in the literature. For 29 stars we adopted the literature values. For the remaining 5 stars, our periodogram analysis did not provide the rotation period. The results of this investigation are presented in the catalogue of photometric rotation period of the \object{$\beta$ Pictoris} association members \rm (Messina et al. 2016; Paper I) where we describe the photometric observations newly obtained, their reduction and analysis, and a detailed discussion of our results obtained for each individual star.\\ In this paper (Paper II), we focus on a sub-sample consisting of 66 members for which we know the rotation period and have one measurement at least (from the literature) of the Lithium equivalent width (EW). This sub-sample is to date the largest of any known young loose association. We will make use of it to investigate the correlation between rotation and Li depletion \rm and to compare with earlier results \rm the age of the $\beta$ Pictoris association obtained by the modeling of the Lithium Depletion Boundary (LDB), after decorrelating the Li EW from rotation. \rm

We have retrieved from the literature the measured Li EW of 66 members of the young $\beta$ Pictoris association. We have carried out a photometric monitoring of these members that allowed us to measure for the first time the rotation periods of 30 members. For other 16 members we retrieved the rotation periods from \citet{Messina10, Messina11}, and for the remaining members we used different sources in the literature.\\ We have explored the existence of a connection between rotation and Li depletion. After removing the mass dependence of the Li EW, using linear fits to EW versus V$-$K$_s$, we found that for stars with 0.5 $\le$ V$-$K$_s$ $\le$ 3.4\,mag, roughly corresponding to masses M $>$ 0.8\,M$_\odot$, no significant correlation is found between Li EW and rotation period. On the contrary, in the color range 3.4 $\le$ V$-$K$_s$ $\le$ 4.5\,mag, roughly corresponding to masses 0.3 $<$ M $<$ 0.8\,M$_\odot$, we find a strong correlation between the Li EW and the rotation period, where fast rotators are much less Li-depleted than slow rotators. Finally, in the color range 5.4 $<$ V$-$K$_s$ $<$ 5.9\,mag, roughly corresponding to masses M $\sim$ 0.1\,M$_\odot$, we have some hint for an inverted correlation where fast rotators are more depleted than slow rotators. However, this correlation is currently based on only 7 stars and the significance level is high (80\%), but this is due to just one point that could be an outlier for whatever reason.\\ Interestingly, the dispersion in the 3.4 $\le$ V$-$K$_s$ $\le$ 4.5\,mag has a peak-to-peak amplitude that amounts to 180\,m\AA\,\, which is about a factor of 2 larger than that measured in the same mass range in the 5-Myr \object{NGC\,2264} open cluster \citep{Bouvier16}, and about a factor 2 smaller than that measured in the 125-Myr \object{Pleiades} open cluster \citep{Soderblom93}. Therefore, we note that the effect of rotation on the Li depletion is also age dependent and increases with age.\\ A comparison with the \citet{Baraffe15} models, shows a mismatch of about 0.7\,mag in the V$-$K$_s$ color \rm between the observed and the predicted color range where the Li gap falls. Models predict the Li gap at bluer colors than observed. On the contrary, the Dartmouth models that incorporate the effects of magnetic fields provide a good match with the hot side of the Li depletion gap, although some mismatch on the cool side of the Li gap remains. Comparing the Dartmouth models with the hot side of the Li depletion gap, we infer an age of 25$\pm$3\,Myr, which is in very good agreement with the most recent age estimates for the $\beta$ Pictoris association. However, the relatively large values of the reduced chi-squares suggest that either the intrinsic Li EW variability is still underestimated or some other factor, apart from rotation, plays a relevant role in producing the observed scatter among stars with similar mass.\\ \rm \begin{figure}[!h] \begin{minipage}{10cm} \includegraphics[scale = 0.3, trim = 0 0 0 0, clip, angle=90]{chisq.eps} \end{minipage} \caption{\label{chisq} Reduced chi-squares versus age obtained from the residuals from the fit to the hot boundary of the Li depletion gap.} \end{figure} {\it Acknowledgements}. Research on stellar activity at INAF- Catania Astrophysical Observatory is supported by MIUR (Ministero dell'Istruzione, dell'Universit\'a e della Ricerca). This research has made use of the Simbad database, operated at CDS (Strasbourg, France). SM thanks Jerome Bouvier for useful discussion and the anonymous Referee for useful comments that helped us to improve this paper. \begin{table*} \caption{Properties of the 66 members of the $\beta$ Pictoris association studied in this work} \begin{minipage}{20cm} \hspace{-1cm} \scriptsize \begin{tabular}{l l l l l l l c r r r r} \hline Target & \multicolumn{3}{c}{RA } & \multicolumn{3}{c}{DEC} & Sp.T & V\hspace{0.3cm} & V$-$K$_s$ & P \hspace{0.2cm} & Li EW\\ & \multicolumn{3}{c}{(hh, mm, ss)} & \multicolumn{3}{c}{($^{\circ}$, $^{\prime}$, $^{\prime\prime}$)} & & (mag) & (mag) & (d)\hspace{0.2cm} & (m\AA)\\ \hline \object{HIP\,560}& 00& 06& 50.08& $-$23& 06& 27.20& F3V& 6.15 & 0.910 & 0.224 & 87.0 \\ \object{TYC\,1186-0706-1}& 00& 23& 34.66& +20& 14& 28.75& K7.5V+M5& 10.96 & 3.623 & 7.900 & 338.0 \\ \object{GJ\,2006A}& 00& 27& 50.23& $-$32& 33& 06.42& M3.5Ve& 12.87 & 4.858 & 3.990 & 29.0 \\ \object{GJ\,2006B}& 00& 27& 50.35& $-$32& 33& 23.86& M3.5Ve& 13.16 & 5.044 & 4.910 & 26.0 \\ \object{2MASS\,J01112542+1526214}A& 01& 11& 25.42& +15& 26& 21.50& M5+M6& 14.46 & 6.252 & 0.911 & 593.0 \\ \object{2MASS\,J01351393-0712517}& 01& 35& 13.93& $-$07& 12& 51.77& M4.5V& 13.22 & 5.502 & 0.703 & 46.7 \\ \object{TYC\,1208-0468-1}& 01& 37& 39.42& +18& 35& 32.91& K3+K5& 9.85 & 3.114 & 2.803 & 440.0 \\ \object{HIP\,10679}& 02& 17& 24.74& +28& 44& 30.43& G2V& 7.75 & 1.488 & 0.777 & 160.0 \\ \object{HIP\,10680}& 02& 17& 25.28& +28& 44& 42.16& F5V& 6.95 & 1.163 & 0.240 & 140.0 \\ \object{HIP\,11437}A& 02& 27& 29.25& +30& 58& 24.60& K4& 10.12 & 3.040 & 12.500 & 248.0 \\ \object{HIP\,11437}B& 02& 27& 28.05& +30& 58& 40.53& M1& 12.44 & 4.219 & 4.660 & 220.0 \\ \object{HIP\,12545}& 02& 41& 25.90& +05& 59& 18.00& K6Ve& 10.37 & 3.301 & 4.830 & 450.0 \\ \object{GJ3305}& 04& 37& 37.30& $-$02& 29& 28.00& M1+M?& 10.59 & 4.177 & 4.890 & 140.0 \\ \object{2MASS\,J04435686+3723033}& 04& 43& 56.87& +37& 23& 03.30& M3Ve+M5?& 12.98 & 4.179 & 4.288 & 194.0 \\ \object{HIP\,23200}& 04& 59& 34.83& +01& 47& 00.68& M0.5Ve& 10.05 & 3.990 & 4.430 & 270.0 \\ \object{HIP\,23309}& 05& 00& 47.10& $-$57& 15& 25.00& M0Ve& 10.00 & 3.756 & 8.600 & 360.0 \\ \object{HIP\,23418}A& 05& 01& 58.80& +09& 59& 00.00& M3V& 11.45 & 4.780 & 1.220 & 53.0 \\ \object{BD\,-211074A}& 05& 06& 49.90& $-$21& 35& 09.00& M1.5V& 10.29 & 4.350 & 9.300 & 20.0 \\ \object{BD\,-211074B}& 05& 06& 49.90& $-$21& 35& 09.00& M2.5V& 11.67 & 4.643 & 5.400 & 20.0 \\ \object{2MASS\,J05082729-2101444}& 05& 08& 27.30& $-$21& 01& 44.40& M5.6V& 14.70 & 5.867 & 0.280 & 618.0 \\ \object{2MASS\,J05241914-1601153}& 05& 24& 19.15& $-$16& 01& 15.30& M4.5+M5& 13.50 & 5.603 & 0.401 & 217.0 \\ \object{HIP\,25486}& 05& 27& 04.76& $-$11& 54& 03.47& F7V& 6.22 & 1.294 & 0.966 & 191.0 \\ \object{2MASS\,J05335981-0221325}& 05& 33& 59.81& $-$02& 21& 32.50& M2.9V& 12.42 & 4.725 & 7.250 & 49.0 \\ \object{2MASS\,J06131330-2742054}& 06& 13& 13.31& $-$27& 42& 05.50& M3.V:& 12.09 & 5.230 & 16.9 & 28.0 \\ \object{HIP\,29964}& 06& 18& 28.20& $-$72& 02& 41.00& K4Ve& 9.80 & 2.986 & 2.670 & 420.0 \\ \object{TWA\,22}& 10& 17& 26.89& $-$53& 54& 26.50& M5& 13.99 & 6.301 & 0.830 & 510.0 \\ \object{HIP\,76629}& 15& 38& 57.50& $-$57& 42& 27.00& K0V& 7.97 & 2.118 & 4.270 & 292.0 \\ \object{2MASS\,J16430128-1754274}& 16& 43& 01.29& $-$17& 54& 27.50& M0.6& 12.50 & 3.951 & 5.140 & 300.0 \\ \object{HIP\,84586}& 17& 17& 25.50& $-$66& 57& 04.00& G5IV+K5IV& 7.23 & 2.528 & 1.680 & 250.0 \\ \object{HD\,155555C}& 17& 17& 31.29& $-$66& 57& 05.49& M3.5Ve& 12.71 & 5.081 & 4.430 & 20.0 \\ \object{TYC\,872822621}& 17& 29& 55.10& $-$54& 15& 49.00& K1V& 9.55 & 2.186 & 1.830 & 360.0 \\ \object{GSC\,08350-01924}& 17& 29& 20.67& $-$50& 14& 53.00& M3V& 12.86 & 4.766 & 1.982 & 50.0 \\ \object{V4046\,Sgr}& 18& 14& 10.50& $-$32& 47& 33.00& K5+K7& 10.44 & 3.191 & 2.420 & 440.0 \\ \object{2MASS\,J18151564-4927472}& 18& 15& 15.64& $-$49& 27& 47.20& M3V& 12.86 & 4.780 & 0.447 & 46.0 \\ \object{HIP\,89829}& 18& 19& 52.20& $-$29& 16& 33.00& G1V& 8.89 & 1.837 & 0.571 & 290.0 \\ \object{2MASS\,J18202275-1011131}A& 18& 20& 22.74& $-$10& 11& 13.62& K5Ve+K7Ve& 10.63 & 3.350 & 4.650 & 530.0 \\ \object{TYC\,907724891}& 18& 45& 37.02& $-$64& 51& 46.14& K5Ve& 9.30 & 3.204 & 0.345 & 490.0 \\ \object{TYC\,907307621}& 18& 46& 52.60& $-$62& 10& 36.00& M1Ve& 11.80 & 3.946 & 5.370 & 332.0 \\ \object{HD\,173167}& 18& 48& 06.36& $-$62& 13& 47.02& F5V& 7.28 & 1.144 & 0.250 & 107.0 \\ \object{TYC\,740800541}& 18& 50& 44.50& $-$31& 47& 47.00& K8Ve& 11.20 & 3.660 & 1.075 & 492.0 \\ \object{HIP\,92680}& 18& 53& 05.90& $-$50& 10& 50.00& K8Ve& 8.29 & 1.924 & 0.944 & 287.0 \\ \object{TYC\,687210111}& 18& 58& 04.20& $-$29& 53& 05.00& M0Ve& 11.78 & 3.762 & 0.503 & 483.0 \\ \object{2MASS\, J19102820-2319486}& 19& 10& 28.21& $-$23& 19& 48.60& M4V& 13.20 & 4.985 & 3.640 & 55.0 \\ \object{TYC\,687801951}& 19& 11& 44.70& $-$26& 04& 09.00& K4Ve& 10.27 & 2.904 & 5.650 & 320.0 \\ \object{2MASS\,J19233820-4606316}& 19& 23& 38.20& $-$46& 06& 31.60& M0V& 11.87 & 3.598 & 3.242 & 422.0 \\ \object{TYC\,744311021}& 19& 56& 04.37& $-$32& 07& 37.71& M0.0V& 11.80 & 3.954 & 11.300 & 110.0 \\ \object{2MASS\, J19560294-3207186}AB& 19& 56& 02.94& $-$32& 07& 18.70& M4V& 13.23 & 5.116 & 1.569 & 100.0 \\ \object{2MASS\,J20013718-3313139}& 20& 01& 37.18& $-$33& 13& 14.01& M1& 12.25 & 4.056 & 12.700 & 100.0 \\ \object{2MASS\,J20055640-3216591}& 20& 05& 56.41& $-$32& 16& 59.15& M2:& 11.96 & 4.022 & 8.368 & 140.0 \\ \object{HD\,191089}& 20& 09& 05.21& $-$26& 13& 26.52& F5V& 7.18 & 1.104 & 0.488 & 91.0 \\ \object{2MASS\, J20100002-2801410}AB& 20& 10& 00.03& $-$28& 01& 41.10& M2.5+M3.5& 12.80 & 4.640 & 0.470 & 46.0 \\ \object{2MASS\,J20333759-2556521}& 20& 33& 37.59& $-$25& 56& 52.20& M4.5& 14.87 & 5.993 & 0.710 & 504.0 \\ \object{HIP\,102141}A& 20& 41& 51.20& $-$32& 26& 07.00& M4Ve& 10.36 & 5.416 & 1.191 & 20.0 \\ \object{HIP\,102141}B& 20& 41& 51.10& $-$32& 26& 10.00& M4Ve& 10.36 & 5.416 & 0.781 & 20.0 \\ \object{2MASS\,J20434114-2433534}& 20& 43& 41.14& $-$24& 33& 53.19& M3.7+M4.1& 12.73 & 4.971 & 1.610 & 28.0 \\ \object{HIP\,102409}& 20& 45& 09.50& $-$31& 20& 27.00& M1Ve& 8.73 & 4.201 & 4.860 & 80.0 \\ \object{HIP\,103311}& 20& 55& 47.67& $-$17& 06& 51.04& F8V& 7.35 & 1.539 & 0.356 & 147.0 \\ \object{TYC\,634902001}& 20& 56& 02.70& $-$17& 10& 54.00& K6Ve+M2& 10.62 & 3.541 & 3.410 & 420.0 \\ \object{2MASS\,J21100535-1919573}& 21& 10& 05.36& $-$19& 19& 57.40& M2V& 11.54 & 4.344 & 3.710 & 41.0 \\ \object{2MASS\, J21103147-2710578}& 21& 10& 31.48& $-$27& 10& 57.80& M4.5V& 14.90 & 5.600 & 0.650 & 502.0 \\ \object{TYC\,9486-927-1}& 21& 25& 27.49& $-$81& 38& 27.68& M2V& 11.70 & 4.360 & 0.542 & 104.0 \\ \object{TYC\,221113091} & 22 & 00 & 41.59 & +27 & 15 & 13.60 & M0V & 11.39 & 3.666 & 1.109 & 40.0 \\ \object{TYC\,934004371}& 22& 42& 48.90& $-$71& 42& 21.00& K7Ve& 10.60 & 3.706 & 4.460 & 440.0 \\ \object{HIP\,112312}& 22& 44& 58.00& $-$33& 15& 02.00& M4Ve& 12.10 & 5.168 & 2.370 & 30.0 \\ \object{TX Psa}& 22& 45& 00.05& $-$33& 15& 25.80& M4.5Ve& 13.36 & 5.567 & 1.080 & 450.0 \\ \object{TYC\,583206661}& 23& 32& 30.90& $-$12& 15& 52.00& M0Ve& 10.54 & 3.971 & 5.680 & 185.0 \\ \hline \end{tabular} \end{minipage}% \end{table*}

1607.06634

1607

1607.07632_arXiv.txt

{ In this first investigation of the MOCCA database with respect to cataclysmic variables, we found that for models with Kroupa initial distributions, considering the standard value of the efficiency of the common-envelope phase adopted in BSE, no single cataclysmic variable was formed only via binary stellar evolution, i. e., in order to form them, strong dynamical interactions have to take place. Our results also indicate that the population of cataclysmic variables in globular clusters are, mainly, in the last stage of their evolution and observational selection effects can change drastically the expected number and properties of observed cataclysmic variables.

Cataclysmic variables (CVs) are among the most interesting objects in globular clusters (GCs). They are interacting binaries composed of a white dwarf (WD) that accretes matter stably from a main sequence (MS) star or a brown dwarf (BD) \citep[e. g.,][for a comprehensive review]{Knigge_2011_OK}. CVs are subdivided according to their photometric behaviours as well as the WD magnetic field strength, being, mainly, magnetic CVs (where the accretion is partially or directly via magnetic field lines) and non-magnetic CVs (where the accretion is via an accretion disk). Among the non-magnetic CVs, the most prominent subgroup is that composed of dwarf novae (DNe) which exhibit repetitive outbursts due to the thermal instability in the accretion disk. In this initial investigation, we analysed six GC models: three (called S models) with ``Standard'' distributions of the initial binary properties (uniform distribution for the mass ratio, uniform in log or log-normal distribution for the semi-major axis and thermal distribution of eccentricities), and three with the Kroupa initial binary population \citep{Kroupa_INITIAL} (called Kroupa models). In what follows, we will present the main results achieved so far.

The study of CVs in GCs with the MOCCA code has just started and we expect more exciting results in future investigations. With the analysis of only six models, we could already find interesting results. The Kroupa initial binary population might require some adjustments, although more tests with respect to the binary stellar evolution parameters must be done before claiming that. Above all, it seems clearly that the initial binary population can be potentially constrained with specific binary populations (like CVs) in the field and GCs. The search for DNe should be in the direction of deeper observations (optical, H$\alpha$ and FUV), in order to reach the faint CVs during their quiescence. Additionally, more effort should be put in finding secure optical counterparts of faint X-ray sources down to $\sim 10^{30}$ erg/s. Finally, the CV population in GCs are intrinsically old, which indicates that cluster CVs are substantially older in comparison with the observed field CVs. This implies that cluster CVs are intrinsically fainter than observed field CVs. It is worth mentioning that we have been preparing two papers where the results will be presented and discussed with more details.

1607.07632

1607

1607.02221_arXiv.txt

Usually, we assume that there is no inhomogeneity isotropic in terms of our location in our universe. This assumption has not been observationally confirmed yet in sufficient accuracy, and we need to consider the possibility that there are non-negligible large-scale isotropic inhomogeneities in our universe. The existence of large-scale isotropic inhomogeneities affects the determination of the cosmological parameters. In particular, from only the distance-redshift relation, we can not distinguish the inhomogeneous isotropic universe model from the homogeneous isotropic one, because of the ambiguity in the cosmological parameters. In this paper, in order to avoid such ambiguity, we consider three observables, the distance-redshift relation, the fluctuation spectrum of the cosmic microwave background radiation(CMBR) and the scale of the baryon acoustic oscillation(BAO), and compare these observables in two universe models; One is the inhomogeneous isotropic universe model with the cosmological constant and the other is the homogeneous isotropic universe model with the dark energy other than the cosmological constant. We show that these two universe models can not predict the same observational data of all three observables but the same ones of only two of three, as long as the perturbations are adiabatic. In principle, we can distinguish the inhomogeneous isotropic universe from the homogeneous isotropic one through appropriate three observables, if the perturbations are adiabatic.

\label{Sec1} Usually, the modern physical cosmology adopts the Copernican principle which states that we do not live in the privileged domain in the universe. The Copernican principle and the observed high isotropy of the CMBR provide the high homogeneity and isotropy of our universe in the globally averaged sense. By contrast, if we remove the Copernican principle, the isotropy around us does not necessarily imply the homogeneity of our universe. The Copernican principle can not be directly confirmed by observations since in order to do so we have to move to the other clusters of galaxies from our galaxy. Hence there is the possibility that there are large-scale isotropic inhomogeneities in our universe. The existence of isotropic inhomogeneities around us affects the determination of the cosmological parameters. The universe model with large-scale isotropic inhomogeneities has been studied in the context of the scenario to explain the observed distance-redshift relation without introducing dark energy components within the framework of general relativity. There are several severe observational constraints on the scenario without dark energy\cite{Tomita:2000jj,Tomita:2001gh,Celerier:1999hp,Iguchi:2001sq,Yoo:2008su,Bull:2012zx,Clifton:2008hv,Vanderveld:2006rb,Yoo:2010qn}, in particular, these universe models are constrained by observations of the kinetic Sunyaev-Zeldovich effect\cite{Zhang,Zibin:2010a,Zibin:2011ma}; The scenario with adiabatic isotropic inhomogeneities has already been ruled out, even though both growing and decaying modes are assumed to exist. On the other hand, the scenario with non-adiabatic isotropic inhomogeneities has not been ruled out yet, although there is an argument on whether the initial condition is contrived. Even if there are dark energy components, not so large isotropic inhomogeneities may exist and significantly affect observational results\cite{Romano:2010nc,Romano:2011mx,Marra:2010pg,Sinclair:2010sb,Valkenburg:2013qwa,Valkenburg:2012td,deLavallaz:2011tj,Valkenburg:2011ty,Marra:2012pj,Tokutake:2016hod,Negishi:2015oga}. Since our observation is confined on a past light cone, a spatially more distant event we observe occurred a longer time ago. If the universe is homogeneous and isotropic, the temporal evolution of the universe is revealed by observing events of various distances. By contrast, if the universe is inhomogeneous and isotropic, observational data contain the information about not only the temporal evolution of the universe but also its spatial inhomogeneities. An intrinsic degeneracy between temporal evolution and isotropic inhomogeneities around us may cause systematic errors. Denoting the total energy density and the total pressure of dark energy components by $\rho_{\rm d}$ and $p_{\rm d}$, respectively, its equation of state is given by \begin{equation} p_{\rm d}=w\rho_{\rm d} \label{EOS} \end{equation} with $\rho_{\rm d}>0$, $w$ is a function less than $-1/3$. The case of $w=-1$ corresponds to the cosmological constant. We can not distinguish the inhomogeneous isotropic universe model with the cosmological constant from the homogeneous isotropic universe model with dark energy of $w\neq -1$, if we have the observational date of the distance-redshift relation only; isotropic inhomogeneities may cause systematic errors on the amount of dark energy and its equation of state. Some authors studied systematic errors due to isotropic inhomogeneities and have evaluated the magnitude of them\cite{Marra:2010pg,Valkenburg:2012td,Romano:2010nc,Romano:2011mx,Sinclair:2010sb,Valkenburg:2013qwa,deLavallaz:2011tj,Valkenburg:2011ty,Marra:2012pj,Tokutake:2016hod,Negishi:2015oga}. However, nobody has not shown how to remove the systematic errors caused by isotropic inhomogeneities. It is just the main purpose of the present paper. The homogenous isotropic universe model is often called the Friedmann-Lema\^{\i}tre-Robertson-Walker (FLRW) universe model, and hereafter we call so. In this paper, we study whether we can distinguish the inhomogeneous isotropic universe model with the cosmological constant from the FLRW universe model with dark energy of $w\neq -1$ and remove the systematic error due to isotropic inhomogeneities, if we use multiple observables: the distance-redshift relation, the fluctuation spectrum of the CMBR and the scale of the BAO in the distribution of galaxies. In other words, we investigate whether there is an inhomogeneous isotropic universe model whose distance-redshift relation, fluctuation spectrum of the CMBR and BAO scale are identical with those of the FLRW universe model. We assume that the inhomogeneous isotropic universe model is filled with non-relativistic matter which is cold dark matter(CDM) and baryonic matter and a positive cosmological constant. Furthermore, we restrict ourselves to the case that the amplitude of isotropic inhomogeneities is so small that they can be treated by the linear perturbation approximation of the flat FLRW universe model, and the scale of isotropic inhomogeneities are larger than the BAO scale but should be smaller than the present horizon scale. If we can find such an inhomogeneous isotropic universe model, it is impossible to distinguish these two universe models from each other, otherwise we can. The organization of this paper is as follows. In Sec.~II, we derive the basic equations for the inhomogeneous isotropic universe model. In Sec.~III, we derive the null geodesic equations which are used to construct the inhomogeneous isotropic universe model from observables given from the FLRW universe model. In Sec.~IV, we show expressions of observables in the inhomogeneous isotropic universe model. In Sec.~V, we give the observables in the FLRW model and derive conditions to determine the inhomogeneous isotropic universe model. We explain the numerical procedure and the numerical result in Sec.~VI. Finally, Sec.~VII is devoted to the summary and discussion. In this paper, we adopt the sign conventions of the metric and Riemann tensor of Ref.\cite{wald} and the geometrized unit in which the speed of light and Newton's gravitational constant are one.

We studied whether we can distinguish the inhomogeneous isotropic universe model with the cosmological constant from the FLRW universe model with dark energy other than the cosmological constant and remove the systematic error due to isotropic inhomogeneities, by considering multiple observables: the distance-redshift relation, the fluctuation spectrum of the CMBR and the BAO scale. We found that we can do so; There is no inhomogeneous isotropic universe model whose distance-redshift relation, fluctuation spectrum of the CMBR and the BAO scale are identical with those of the FLRW universe model, as long as the density perturbations are adiabatic. It is nontrivial that we can distinguish the inhomogeneous isotropic universe model from the FLRW universe model by using a finite number of observables, since the inhomogeneous isotropic universe model has a functional degree of freedom. Here it should be noted that we have used not only the information about the universe on the past light cone but also that inside the past light cone, since we assumed that the BAO scale at the decoupling time is everywhere constant. However, there is a possibility that the BAO scale at the decoupling time is inhomogeneous, if the ratio between the energy densities of baryonic matter and CDM has been inhomogeneous at the decoupling time. If the ratio between the energy densities of baryonic matter and CDM is inhomogeneous and isotropic, we can not distinguish the inhomogeneous isotropic universe model from the FLRW universe model by only the distance-redshift relation, the fluctuation spectrum of the CMBR and the BAO scale. In the case of $w_{0}=-1.01$ and $w_{1}=0$, the inhomogeneous isotropic universe model can also explain observations, if the fluctuation of the ratio between the energy densities of baryonic matter and CDM at $t=t_0$ is $-0.037408$ (see Appendix A). It is very important to observe the ratio between the energy densities of baryonic matter and CDM in the domain $0<z<2$.

1607.02221

1607

1607.00012_arXiv.txt

{The spin of a planet or brown dwarf is related to the accretion process, and therefore studying spin can help promote our understanding of the formation of such objects. We present the projected rotational velocity of the young sub-stellar companion GQ Lupi b, along with its barycentric radial velocity. The directly imaged exoplanet or brown dwarf companion joins a small but growing ensemble of wide-orbit sub-stellar companions with a spin measurement. The GQ Lupi system was observed at high spectral resolution ($R \sim$\num{100000}), and in the analysis we made use of both spectral and spatial filtering to separate the signal of the companion from that of the host star. We detect both \ce{CO} (S/N=11.6) and \ce{H2O} (S/N=7.7) in the atmosphere of GQ Lupi b by cross-correlating with model spectra, and we find it to be a slow rotator with a projected rotational velocity of $5.3^{+0.9}_{-1.0}$ \si{\km\per\s}. The slow rotation is most likely due to its young age of $< \SI{5}{Myr}$, as it is still in the process of accreting material and angular momentum. We measure the barycentric radial velocity of GQ Lupi b to be \SI{2.0 \pm 0.4}{\km\per\s}, and discuss the allowed orbital configurations and their implications for formation scenarios for GQ Lupi b.}

Measurements of the spin and the orbit of giant extrasolar planets and brown dwarf companions may hold important clues to their origin and evolution. Generally, two formation processes are considered for giant planets: i) core accretion and ii) disk fragmentation. Jupiter and Saturn are commonly accepted to have formed through core accretion, a class of formation models where gas accretes onto solid planetary embryos of several to ten Earth masses which may have formed beyond the iceline by runaway accretion from kilometer-sized planetesimals \citep{Pollack1996, Laughlin2004, Hubickyj2005}. The discovery of extrasolar giant planets led to the reinvigoration of the disk fragmentation hypothesis which states that giant (exo)planets may form as a disk gravitational instability that collapses on itself in the outer protoplanetary disk \citep{Boss1997, Boss2000}. Disk fragmentation is also considered a potential formation scenario for the more massive brown dwarf companions \citep{Chabrier2014}, or alternatively brown dwarf companions can be the result of prestellar core fragmentation during the earliest stages of the cloud collapse \citep{Jumper2013}. Spin is predominantly a result of accretion of angular momentum during the formation, and if core accretion and gravitational instability result in differences in spin angular momentum, it is possible this will show up in studies of spin of sub-stellar companions as function of mass. In the Solar System, the spin angular momenta of those planets not influenced by tidal effects or tidal energy dissipation by a massive satellite follow a clear relationship, spinning faster with increasing mass \citep{Hughes2003}. In particular, gas giants far away from their central star are likely to have primordial spin angular momentum, making the directly imaged sub-stellar companions ideal candidates for exploring the connection between formation and spin. The rotational velocity of an exoplanet was measured for the first time by \citet{Snellen2014}, observing the directly imaged planet $\beta$ Pictoris b with high-dispersion spectroscopy, measuring it to have a projected rotational velocity of $v\sin(i)=$ \SI{25}{\kilo\m\per\s}. Another young directly imaged planet, 2M1207 b, became the first exoplanet to directly have its rotational period measured ($P_{\textrm{rot}}=$ \SI{10.7}{hr}), when \citet{Zhou2016} detected rotational modulations in HST/WFC3 photometric monitoring of the object. Both results are in accordance with an extrapolation of the spin-mass trend observed in the Solar System planets (see Fig. \ref{fig:planet_spin}). Apart from being related to the accretion process, the spin of an exoplanet is also a fundamental observable that affects in particular its atmospheric dynamics and climate as well as e.g. its magnetic fields. On Earth, the Coriolis effect generates large-scale ocean currents which in turn promote cyclones. For fast rotators, including many brown dwarfs, the wind flows are rotation dominated \citep{Showman2013}. On the other hand, exoplanets orbiting close to their parent star are expected to be tidally locked. \citet{Brogi2016} and \citet{Louden2015} both recently made use of high-dispersion spectroscopy to detect a Doppler signature in the transmission spectrum of the hot Jupiter HD 189733 b, consistent with synchronous rotation. Synchronous rotation is the cause of large temperature differences between the day- and night-side which in turn can cause fast winds flowing from the hot day-side to the cold night-side. Another approach to understanding the formation and dynamical evolution of exoplanets is to study their orbits. In the case of directly imaged sub-stellar companions, it is often difficult to constrain the orbits, due to the long time-scales involved \citep[e.g.][]{Pearce2015, Ginski2014A}. However, using a high-dispersion slit spectrograph in combination with adaptive optics, it is possible to extract spatially resolved high-dispersion spectra for the companion \citep{Snellen2014}, and thereby measure even very small Doppler shifts due to the orbital motion of the companion. Thus, even just one radial velocity measurement can in some cases prove to be a powerful orbital constraint \citep{Nielsen2014, Lecavelier2016}. In this paper we present both the spin measurement and the barycentric radial velocity of the widely-separated sub-stellar companion GQ Lupi b from high-dispersion spectroscopy. We introduce the GQ Lupi system in Section \ref{sec:The-GQ-Lupi-system}, and give the details of the observations in Section \ref{sec:Observations}. The data analysis is detailed in Section \ref{sec:Data-analysis} with special emphasis on how the spatially resolved spectrum of the companion is extracted and cleaned from telluric and stellar spectral lines. In Section \ref{sec:Measuring-the-signal} we explain the cross-correlation analysis which is employed to measure the rotational broadening and doppler shifts of the molecular lines in the companion spectrum, and the results from the companion are presented in Section \ref{sec:Results}, along with the host star spin and systemic velocity. We discuss the implications of the spin measurement in Section \ref{subsec:The-slow-spin} and the constraints on orbital elements in Section \ref{subsec:The-orbital-orientation}.

\label{sec:Conclusion} The young GQ Lupi system has a central Classical T Tauri star surrounded by a warm dust disk, and is orbited by a sub-stellar companion GQ Lupi b at $\sim$\SI{100}{\au}, which is either a gas giant or a brown dwarf. We observed the parent star and the companion simultanously in the K-band by careful positioning of the slit of the high-dispersion spectrograph CRIRES, in combination with adaptive optics. We made use of both the spatial and spectral information to separate the spectrum of the companion from that of the host star. We detect both water and CO in the companion spectrum. The molecular lines are rotationally broadened and Doppler shifted due to the orbital motion of the companion. We have measured the projected rotation velocity to be $v\sin(i) = 5.3^{+0.9}_{-1.0} $\si{\km\per\s} and the barycentric radial velocity to be RV$=$\SI{2.0 \pm 0.4}{\km\per\s}. GQ Lupi b is a slow rotator compared to the giant planets in the Solar System, and to the recent spin measurements of the exoplanets $\beta$ Pic b ($v\sin(i)=$\SI{25}{\km\per\s}, \citet{Snellen2014}) and 2M1207 b ($v\sin(i)=$\SI{17}{\km\per\s}, \citet{Zhou2016}). This is in spite of GQ Lupi b being likely more massive than either of these, and thus this new spin measurement does not agree with the spin-mass trend of the others. However, we argue that the slow spin is a manifestation of the young age of GQ Lupi, and that the discrepancy cannot be used to argue for fundamental differences in formation path at this time. We have used the barycentric radial velocity measurement to place new constraints on orbital elements such as the semi-major axis, the eccentricity and the orbital inclination with respect to the observer. This shows the strength of high-dispersion spectroscopy, because it is possible to measure even small radial velocities. The spin and RV measurements of GQ Lupi b demonstrate the potential of the combination of spatial and spectral filtering through the use of high dispersion spectrographs together with adaptive optics.

1607.00012

1607

1607.07404_arXiv.txt

name}{} \renewcommand{\refname}{\sc references} \renewcommand{\figurename}{Fig.} \makeatletter \renewcommand{\@oddhead}{\textit{Advances in Astronomy and Space Physics} \hfil} \renewcommand{\@evenfoot}{\hfil \thepage \hfil} \renewcommand{\@oddfoot}{\hfil \thepage \hfil} \makeatother \let\oldthebibliography=\thebibliography \let\endoldthebibliography=\endthebibliography \renewenvironment{thebibliography}[1]{\begin{oldthebibliography}{#1}\setlength{\parskip}{0ex}\setlength{\itemsep}{0ex}}{\end{oldthebibliography}} \begin{document} \fontsize{11}{11}\selectfont % \title{Evolution of density and velocity profiles of matter \\ in large voids} \author{\textsl{M.~Tsizh\footnote{tsizh@astro.franko.lviv.ua}, B.~Novosyadlyj}} \date{\vspace*{-6ex}} We analyse the evolution of cosmological perturbations which leads to the formation of large voids in the distribution of galaxies. We assume that perturbations are spherical and all components of the Universe — radiation, matter and dark energy - are continuous media with ideal fluid energy-momentum tensors, which interact only gravitationally. Equations of the evolution of perturbations in the comoving to cosmological background reference frame for every component are obtained from equations of conservation and Einstein's ones and are integrated by modified Euler method. Initial conditions are set at the early stage of evolution in the radiation-dominated epoch, when the scale of perturbation is mush larger than the particle horizon. Results show how the profiles of density and velocity of matter in spherical voids with different overdensity shells are formed.\\[1ex] {\bf Key words:} cosmology: dark energy, large-scale structure of Universe

1607.07404

1607

1607.07859_arXiv.txt

{We report three newly discovered exoplanets from the SuperWASP survey. WASP-127b is a heavily inflated super-Neptune of mass $0.18 \pm 0.02~\rm M_J$ and radius $1.37 \pm 0.04~\rm R_J$. This is one of the least massive planets discovered by the WASP project. It orbits a bright host star ($\rm V_{mag} = 10.16$) of spectral type G5 with a period of $4.17$ days. WASP-127b is a low-density planet that has an extended atmosphere with a scale height of $2500 \pm 400~\rm km$, making it an ideal candidate for transmission spectroscopy. WASP-136b and WASP-138b are both hot Jupiters with mass and radii of $1.51 \pm 0.08~\rm M_J$ and $1.38 \pm 0.16~\rm R_J$, and $1.22 \pm 0.08~\rm M_J$ and $1.09 \pm 0.05~\rm R_J$, respectively. WASP-136b is in a $5.22$-day orbit around an F9 subgiant star with a mass of $1.41 \pm 0.07~\rm M_{\odot}$ and a radius of $2.21 \pm 0.22~\rm R_{\odot}$. The discovery of WASP-136b could help constrain the characteristics of the giant planet population around evolved stars. WASP-138b orbits an F7 star with a period of $3.63$ days. Its radius agrees with theoretical values from standard models, suggesting the presence of a heavy element core with a mass of $\sim 10 ~\rm M_{\oplus}$. The discovery of these new planets helps in exploring the diverse compositional range of short-period planets, and will aid our understanding of the physical characteristics of both gas giants and low-density planets.}

\hspace{0.5cm}Over 3000 exoplanets have been discovered as of 2016 July.\footnote{\url{http://exoplanetarchive.ipac.caltech.edu/}} The space-based mission \textit{Kepler} \citep{2010Sci...327..977B} and its successor \textit{K2} \citep{2014PASP..126..398H} have discovered a large number of transiting planets. The results show that small Earth- and Neptune-sized planets are common around solar-like stars \citep{2011ApJ...736...19B}. The {\it Kepler} discoveries have provided a large sample of planetary systems that are very useful for statistical studies of planetary populations (e.g. \citealt{2013ApJ...766...81F,2013ApJ...767...95D,2013PNAS..11019273P}). However, most of the {\it Kepler} planets orbit very faint stars, for which it is quite difficult to obtain precise radial velocity (RV) measurements that are necessary to constrain planetary masses. On the other hand, ground-based systematic surveys (e.g. HATNet: \citealt{2002PASP..114..974B}; SuperWASP: \citealt{2006PASP..118.1407P}; KELT: \citealt{2007PASP..119..923P}; QES: \citealt{2013AcA....63..465A}; HATSouth: \citet{2013PASP..125..154B}; NGTS: \citealt{2013EPSC....8..234W}) usually provide a large number of candidates around stars that are sufficiently bright to enable measuring their RVs, from which strong observational constraints for theoretical studies can be obtained. The upcoming space mission, TESS \citep{2015JATIS...1a4003R}, will target main-sequence dwarf stars that are brighter than $13^{\rm th}$ magnitude and could yield over 1000 planets smaller than Neptune. In addition, all-sky surveys such as MASCARA \citep{2012SPIE.8444E..0IS}, Evryscope, \citep{2015PASP..127..234L} and Fly's Eye Camera \citep{2016PASP..128d5002P} will also be able to provide planet candidates around bright stars for detailed characterisation. Precise measurements and analysis of these systems reveals that planets with similar masses can have very different radii, resulting in very unique physical characteristics. For example, Jupiter-mass planets can range in radii from $0.775 \rm R_J$ (WASP-59b \citealt{2013A&A...549A.134H}) to $1.932 \rm R_J$ (WASP-17b \citealt{2010ApJ...709..159A,2012MNRAS.426.1338S}). A growing number of short-period sub-Saturn and super-Neptune mass planets such as WASP-39b ($\rm M_p = 0.28 M_J$; \citealt{2011A&A...531A..40F}), HAT-P-11b ($\rm M_p = 0.081 M_J$; \citealt{2010ApJ...710.1724B}), and HAT-P-26b ($\rm M_p = 0.059 M_J$; \citealt{2011ApJ...728..138H}) were also found. Many of these planets were found to have radii larger than predicted from standard coreless models (e.g.~\citealt{2007ApJ...659.1661F}). Some theories suggest that the planet radius is correlated with the equilibrium temperature. For example, strong stellar irradiation could heat up the planet, inflating its radius \citep{1996ApJ...459L..35G}. The planetary interior could also be tidally heated as the orbit circularises \citep{2001ApJ...548..466B,2003ApJ...592..555B}. An enhanced atmospheric opacity can hinder the cooling process of the planet such that the planet radius can remain larger for longer \citep{2007ApJ...661..502B}. The interaction between stellar wind and the magnetic field of the planet can lead to Ohmic heating, which could also influence the temperature of the planet \citep{2011ApJ...738....1B}. Low-density planets with extended atmospheres orbiting bright host stars are ideal targets for transmission spectroscopy, which can further reveal the composition of planets. We present here the discovery of three new planets, WASP-127b, WASP-136b, and WASP-138b, discovered by the SuperWASP survey. Section \ref{observation} summarises the observations from the WASP detection, follow-up photometry, and spectroscopic data of each of the planets. In Sect.~\ref{results}, we describe our analysis and present the derived results of the system parameters. Lastly, we discuss the system characteristics and how evolution theories could explain the existence of these planets in Sect.~\ref{discussion}.

Photometry observation of WASP-127, WASP-136, and WASP-138.} \begin{tabular}{lllll} \hline Planet & Date & Instrument & Filter & Comment \\ \hline \multirow{4}{*}{WASP-127b} & \multicolumn{1}{l}{18/03/2014} & \multicolumn{1}{l}{TRAPPIST} & \multicolumn{1}{l}{z} & \multicolumn{1}{l}{partial transit} \\ & \multicolumn{1}{l}{28/04/2014} & \multicolumn{1}{l}{EulerCam} & \multicolumn{1}{l}{Gunn r} & \multicolumn{1}{l}{partial transit} \\ & \multicolumn{1}{l}{13/02/2016} & \multicolumn{1}{l}{LT RISE} & \multicolumn{1}{l}{V + R} & \multicolumn{1}{l}{full transit} \\ & \multicolumn{1}{l}{18/04/2016} & \multicolumn{1}{l}{Zeiss 1.23m} & \multicolumn{1}{l}{Cousins-I} & \multicolumn{1}{l}{full transit} \\ \hline \multirow{3}{*}{WASP-136b} & \multicolumn{1}{l}{21/08/2014} & \multicolumn{1}{l}{EulerCam} & \multicolumn{1}{l}{z} & \multicolumn{1}{l}{partial transit} \\ & \multicolumn{1}{l}{24/11/2014} & \multicolumn{1}{l}{TRAPPIST} & \multicolumn{1}{l}{z} & \multicolumn{1}{l}{partial transit} \\ & \multicolumn{1}{l}{21/08/2015} & \multicolumn{1}{l}{EulerCam} & \multicolumn{1}{l}{z} & \multicolumn{1}{l}{full transit} \\ \hline \multirow{1}{*}{WASP-138b} & \multicolumn{1}{l}{17/12/2015} & \multicolumn{1}{l}{EulerCam} & \multicolumn{1}{l}{NGTS} & \multicolumn{1}{l}{partial transit} \\ \hline \end{tabular} \end{table} \hspace{0.5cm}Multiple follow-up photometry was taken for the three stars to place better constraints on the system parameters. The photometry was obtained with EulerCam at the $1.2$ m Euler-Swiss telescopes \citep{2012A&A...544A..72L} and TRAPPIST \citep{2011Msngr.145....2J,2011EPJWC..1106002G}, which are situated at ESO La Silla Observatory in Chile, the RISE camera on the Liverpool Telescope at the Observatorio del Roque de los Muchachos on La Palma \citep{2008SPIE.7014E..6JS} and the Zeiss $1.23$ m telescope at the German-Spanish Astronomical Center at Calar Alto in Spain. The summary of our follow-up photometric observations is given in Table \ref{PhotometrySummary}, and the phase-folded light curves are shown in Figs.~\ref{photomtery127}, \ref{photomtery136} and \ref{photomtery138}, along with their best-fit transit models derived from our analysis in Sect.~\ref{mcmcanalysis}. \textbf{TRAPPIST:}~ WASP-127 and WASP-136 were both observed with the 0.6 m TRAnsiting Planets and PlanetesImals Small Telescope (TRAPPIST) robotic telescope. The telescope is equipped with a thermoelectrically cooled 2k$\times$2k CCD camera, which has a pixel scale of 0.65'' that translates into a 22'$\times$22' field of view. For details of TRAPPIST, see \citet{2011EPJWC..1106002G} and \citet{2011Msngr.145....2J}. A partial transit of WASP-127b was observed on 2014 March 18 through a Sloan-z’ filter (effective wavelength $=896.3\pm0.8$ nm) with an exposure time of 9 seconds. The same filter was used to observe a partial transit of WASP-136b on 2014 November 24 (effective wavelength $=895.0\pm0.6$ nm), but with an exposure time of 7 seconds. Throughout the observations, the telescope was kept in focus and the positions of the stars on the detector were retained on the same few pixels, thanks to a software guiding system that regularly derives an astrometric solution for the images and sends pointing corrections to the mount when needed. The observation of WASP-136 was made towards the end of night. Therefore the increase in airmass is reflected in the uncertainties of the photometry. Data were reduced as described in \citet{2013A&A...552A..82G}. After a standard pre-reduction (bias, dark, and flat-field correction), the stellar fluxes were extracted from the images using IRAF/DAOPHOT\footnote{IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.} \citep{1987PASP...99..191S}. For each light curve, we tested several sets of reduction parameters and chose the one giving the most precise photometry for the stars of similar brightness as the target. After a careful selection of reference stars, the transit light curves were finally obtained using differential photometry. \textbf{EulerCam:}~ We observed one full transit of WASP-127, one partial and one full transit of WASP-136, and one full transit of WASP-138 with EulerCam \citep{2012A&A...544A..72L}. The full transit of WASP-127 was observed on 2014 April 28 with a Gunn r filter. The telescope was defocused throughout the observation with a FWHM of between 1.6 and 2.5 arcsec. A circular aperture of radius 4.7 arcsec was used along with one reference star for the extraction of the light curve. A partial and a full transit of WASP-136 were obtained on 2014 August 21 and 2015 August 21, respectively. Both observations were taken using a Gunn z filter and an exposure time of 50 seconds. The telescope was substantially defocused throughout both nights. The FWHM of the first night was between 1.5 and 2.3 arcsec. A circular aperture of radius 2.7 arcsec and four reference stars were used for photometry extraction. The FWHM of the second night was between 1.9 and 3.0 arcsec. We used a circular aperture of radius 4.5 arcsec along with five reference stars for the photometry reduction. The observation of WASP-138 was carried out on 2015 December 17 with an NGTS filter (with a custom wavelength of 550 - 900 nm) and exposure times of between 50 and 85 seconds. The telescope was substantially defocused and the FWHM was between 1.3 and 2.5 arcsec. A photometric aperture of 5.6 arcsec radius was used to extract the fluxes, and one reference star was used to generate the relative light curve. See \citet{2012A&A...544A..72L} for further details on EulerCam and its data reduction procedures. \textbf{RISE:}~ A full transit of WASP-127 was observed with RISE \citep{2008SPIE.7014E..6JS}. The camera is equipped with a back-illuminated frame-transfer CCD of $1024\times1024$ pixels. A V+R filter (a custom filter constructed from 3 mm OG515 + 2 mm KG3, with a bandwidth of 500-900 nm) and a $2\times2$ binning of the detector were used for the observation, resulting in a pixel scale of $1.08$ arcsec/pixel. We used an exposure time of 1.5 seconds and defocused the telescope by 0.5 mm for all the observations. Images were automatically bias, dark, and flat corrected by the RISE pipeline. We selected four comparison stars for data reduction. The data were reduced with the standard IRAF apphot routines using a 1.4 pixels (4.86 arcsec) aperture. We attribute the increased scatter around mid-transit to thin clouds (see Fig.~\ref{photomtery127}). \textbf{ZEISS:}~ The Zeiss $1.23$ m telescope has a focal length of $9857.1$ mm and is equipped with the DLR-MKIII camera, which has 4k$\times$4k pixels of size 15 micron. The plate scale is $0.32$ arcsec/pixel and the field of view is $21.5\times21.5$ arcmin. It has previously been successfully used to follow-up many planetary transits (e.g. \citet{2015A&A...579A.136M}). The telescope was defocused and the exposure time was adjusted several times during the night in a range of between 65 and 105 seconds. The guiding camera did not operate correctly on the night of our observations, and the usual precision was not achieved. The CCD was windowed to decrease the readout time and therefore sped up the cadence of the observations. The night was not photometric, and several clouds disturbed the observations. The data were reduced using a revised version of the \textsc{defot} code \citep{2014MNRAS.444..776S}. In brief, the scientific images were calibrated and the photometry was extracted by the standard aperture-photometry technique. The resulting light curve was normalised to zero magnitude by fitting a straight line to the out-of-transit data. \begin{figure}[htbp!] \centering \includegraphics[scale=0.45]{wasp127_mcmcfit.pdf} \caption{\label{photomtery127}Photometry follow-up of WASP-127 observed from EULERCam (red), TRAPPIST(green), RISE(magenta) and Zeiss(cyan). The data are phase-folded with the ephemeris from our analysis. The light curves are assigned an arbitrary offset from the zero magnitute and are binned to a 2-minute cadence for clarity. The best-fit transit model from \citet{2002ApJ...580L.171M} is plotted as a black solid line, and the residuals of the fit are plotted directly below the light curves.} \end{figure} \begin{figure}[htbp!] \centering \includegraphics[width=0.45\textwidth]{wasp136_mcmcfit.pdf} \caption{\label{photomtery136} Photometry follow-up of WASP-136 observed from EULERCam (red) and TRAPPIST(green). The data are phase-folded with the ephemeris from our analysis. The light curves are assigned an arbitrary offset from the zero magnitute and are binned to a 2-minute cadence for clarity. The best-fit transit model from \citet{2002ApJ...580L.171M} is plotted as a black solid line, and the residuals of the fit are plotted directly below the light curves. } \end{figure} \begin{figure}[htbp!] \centering \includegraphics[width=0.45\textwidth]{wasp138_mcmcfit.pdf} \caption{\label{photomtery138}Photometry follow-up of WASP-138 observed from EULERCam (red). The data are phase-folded with the ephemeris from our analysis. The EulerCAM light curve is assigned an arbitrary offset from the zero magnitute for clarity. The best-fit transit model from \citet{2002ApJ...580L.171M} is plotted as a black solid line, and the residuals of the fit are plotted directly below the light curves.} \end{figure} Discussion and conclusion} \subsection{WASP-127b} \hspace{0.5cm}From our best-fit MCMC solution, we obtain a planet with a mass of $0.18 \pm 0.02 ~\rm M_{J}$ and a radius of $1.37 \pm 0.04 ~\rm R_{J}$ ($\rm M_{pl} = 0.16 \pm 0.02 ~M_J$ and $\rm R_{pl} = 1.41 \pm 0.06 ~R_J$ for the case where RVs from both CORALIE and SOPHIE were included for analysis). This means that WASP-127b has a density of $0.07_{-0.01}^{+0.01} ~\rho_{J}$, making it one of the lowest density planets ever discovered. It is also the planet with the second lowest mass discovered by WASP, only more massive than WASP-139b \citep{2016arXiv160404195H}. Compared to the standard coreless model of \citet{2007ApJ...659.1661F}, WASP-127b is over $30\%$ larger than expected for a planet with an orbit at 0.045 au around a 4.5 Gyr solar-type star. From Sect.~\ref{SpectralAnalysis}, WASP-127 is estimated to be much older than the Sun. Hence the theoretical radius of the planet should be even smaller. The anomalously large radius of WASP-127b could be explained by several inflation mechanisms. One such mechanism is tidal heating \citep{2001ApJ...548..466B,2003ApJ...592..555B}, where the planetary interior receives heat energy as the orbit circularises. As with many short-period gas giants, the orbit of WASP-127b may have shrunk and migrated to its current position through planet-planet scattering \citep{2008ApJ...686..621F} or Kozai mechanism \citep{2007ApJ...669.1298F,1962AJ.....67..591K,1962P&SS....9..719L}. The planetary orbit may also have been tidally circularised during the migration process, which results in rapid transfer of energy to its interior, inflating the radius of the planet. This energy transfer process occurs rapidly at the early stages of the evolution history of the system, hence it is unclear how efficient this way of heating up the planetary interior is for such an old system. On the other hand, the atmosphere of WASP-127b could have an enhanced opacity resulting from enhanced metallicity \citep{2007ApJ...661..502B}. This can delay the cooling effect and maintain the inflated radius for a longer time. \citet{2011ApJ...738....1B} suggested that the Ohmic heating mechanism could account for the large planetary radii. The interaction between the planetary magnetic field and the flow of ionised atmospheric heavy elements could induce an electro-motive force. This reaction can drive electrical currents throughout the planet and lead to inflation as the planet heats up. The Ohmic dissipation, however, is limited by the depth of the dissipation. The \citet{2012ApJ...757...47H} model suggests that the convective zone boundary will move deeper if Ohmic dissipation occurs outside the convection zone, which in turn will slow down the efficiency of the planet cooling process. Recently, \citet{2016ApJ...818....4L} have suggested that a planet whose radius was inflated through internal heating could re-inflate again as its host star moves towards the RGB phase because of increased irradiation. Their studies also showed that re-inflation of a planet is most likely for low-mass and short-period planets. WASP-127 is estimated to have a main-sequence lifetime of approximately $\rm t_{MS} = t_{\odot}(M/M_{\odot})^{-2.5} \approx 8$ Gyr, where $\rm t_{\odot}$ is the solar main-sequence lifetime, $\rm M_odot$ is the solar mass and M is the stellar mass. WASP-127 has an isochronal age of $11.41 \pm 1.80$ Gyr, implying that the star is at the end of the main-sequence phase. Therefore, WASP-127b may be undergoing a phase of re-inflation as its host star enters the subgiant branch. WASP-127b is a highly inflated planet that shares many similarities with a number of low-density planets such as WASP-39b \citep{2011A&A...531A..40F}, WASP-113b \citep{2016arXiv160702341B}, HAT-P-8b \citep{2015AJ....150...49B}, HAT-P-11b \citep{2010ApJ...710.1724B}, HAT-P-47b, and HAT-P-48b \citep{2016arXiv160604556B}. Assuming it has an atmosphere identical to that of Jupiter ($\mu = 2.2$u, where the atomic mass unit is u $=1.66\times10^{-27}$ kg), its scale height would be $\rm H \approx 2500 \pm 400 ~km$. WASP-127b also orbits a bright G-type star of $\rm V_{mag} = 10.172$, making it an excellent candidate for transmission spectroscopy. Figure \ref{NeptuneDesert} shows the planetary mass as a function of the orbital period. WASP-127b falls into the short-period Neptune desert \citep{2016A&A...589A..75M}, a region between Jovian and super-Earth planets with a lack of detected planets. The presence of such a region indicates the existence of two unique types of short-period planets, hot Jupiters and super-Earths, which have very different formation and evolution mechanisms. It also suggests that short-period planets with intermediate masses are likely to be destroyed. \begin{figure}[htbp!] \includegraphics[width=0.5\textwidth]{NeptuneDesertplot.pdf} \caption{\label{NeptuneDesert} Planetary mass against orbital period plot. Exoplanet data are taken from the NASA Exoplanet Archive (\protect\url{http://exoplanetarchive.ipac.caltech.edu/}) and are plotted as grey dots. WASP-127b, WASP-136b, and WASP-138b are represented by green, red, and cyan circles, respectively. The black dashed lines are the upper and lower boundaries of the Neptune desert as defined by Mazeh et al. 2016. WASP-127b falls in this Neptune desert between Jovian and super-Earth planets.} \end{figure} \subsection{WASP-136b} \hspace{0.5cm}WASP-136b is an inflated hot-Jupiter transiting a bright F5 host star ($\rm V_{mag} = 9.928$). The best-fit MCMC solution of the system gives a planetary mass of $1.51 \pm 0.08 ~\rm M_{J}$ and a radius of $1.38 \pm 0.16 ~\rm R_{J}$ which yields a planet density of $0.58_{-0.15}^{+0.23} ~\rho_{J}$. The estimated main-sequence lifetime of WASP-136 is approximately $\rm t_{MS} \approx 4$ Gyr. In Sect.~\ref{SpectralAnalysis} we estimated the isochronal age of WASP-136 to be $\sim 3.6 \pm 0.7$ Gyr. This suggests that WASP-136 is at the end of the main-sequence phase. The density and surface gravity of WASP-136 implies that the star is a subgiant, which is consistent with the age estimate. The radius of WASP-136b is approximately $25\%$ larger than predicted from the planet evolution model of \citet{2007ApJ...659.1661F} for a coreless planet. Like many bloated hot-Jupiters (WASP-54b: \citealt{2013A&A...551A..73F}; WASP-78b and WASP-79b: \citealt{2012A&A...547A..61S}; WASP-142b: \citealt{2016arXiv160404195H}), WASP-136b receives a stronger irradiation from its F-type host star compared to a G-type, which can lead to a more inflated planet. Similar to EPIC 211351816.01 \citep{2016arXiv160605818G}, the inflation mechanism of WASP-136b could follow the class I model from \citet{2016ApJ...818....4L}, where WASP-136b might be heated up through the deposit of some of the incident stellar irradiation into the planetary interior. Consequently, when the host of WASP-136b enters the subgiant branch, the planet could receive an increased stellar irradiation that will lead to re-inflation of the planet radius. The existence of short-period hot Jupiters around subgiant stars is rare. The lack of giant planets at short orbital distances may be a consequence of tidal disruption \citep{2013ApJ...772..143S}. A planet residing inside the synchronous orbit could experience strong tidal forces and spiral inwards towards the star. When the planet eventually reaches the Roche limit, the tidal force becomes stronger than the planet's gravitational force, and the planet will be tidally disrupted. WASP-136b has an orbital distance of $0.0661$ au and is approximately 46 times the Roche limit. We followed the derivation of \citet{2010ApJ...725.1995M} and estimated that WASP-136b has a remaining lifetime of $\sim 0.683$ Gyr. Furthermore, during the post-main-sequence phase, the stellar radius expands and will eventually lead to the engulfment of the planet \citep{2009ApJ...705L..81V}. \subsection{WASP-138b} \hspace{0.5cm}Our best-fit solution gives a system with a planetary mass and radius of $1.22 \pm 0.08 ~\rm M_{J}$ and $1.09 \pm 0.05 ~\rm R_{J}$. The density of WASP-138b is $0.92_{-0.15}^{+0.10} ~\rm \rho_J$. It orbits around an F9 star with a metallicity of $\rm [Fe/H]=-0.09 ~dex$, slightly more metal-poor than the Sun. WASP-138b is a hot Jupiter that has similar characteristics to many short-period gas giants with a period of $\sim 3$days, for example, WASP-35b \citep{2011AJ....142...86E} and WASP-141b \citep{2016arXiv160404195H}. With an isochronal age of $3.4 \pm 0.9$ Gyr, the planet evolution model of \citet{2007ApJ...659.1661F} suggests that WASP-138b has a core mass of $\sim 10 ~\rm M_\oplus$ of heavy elements. \hspace{0.5cm}Many interesting large short-period exoplanets are being discovered by ground-based surveys. Discoveries such as that of WASP-127b will provide exceptional targets for future missions such as JWST and for further characterisation. With the diverse range of exoplanets, it is important to understand the mass-radius relation and distinguish the mechanisms responsible for each distinct class of planet. This will ultimately help our understanding of the dynamics of planetary systems.

1607.07859

1607

1607.02145_arXiv.txt

We use the SPARC (Spitzer Photometry \& Accurate Rotation Curves) database to study the relation between the central surface density of stars \sstar\ and dynamical mass \sdyn\ in 135 disk galaxies (S0 to dIrr). We find that \sdyn\ correlates tightly with \sstar\ over 4 dex. This central density relation can be described by a double power law. High surface brightness galaxies are consistent with a 1:1 relation, suggesting that they are self-gravitating and baryon dominated in the inner parts. Low surface brightness galaxies systematically deviate from the 1:1 line, indicating that the dark matter contribution progressively increases but remains tightly coupled to the stellar one. The observed scatter is small ($\sim$0.2 dex) and largely driven by observational uncertainties. The residuals show no correlations with other galaxy properties like stellar mass, size, or gas fraction.

\label{sec:intro} Several lines of evidence suggest that the stellar and total surface densities of disk galaxies are closely linked. The inner rotation curves of low surface brightness (LSB) galaxies rise more slowly than those of high surface brightness (HSB) ones \citep{deBlok1996a, Verheijen1997}, indicating that low stellar densities correspond to low dynamical densities \citep{deBlok1996b}. \citet{Lelli2013} find that the inner slope of the rotation curve $S_0$ (extrapolated for $R\rightarrow0$) correlates with the central surface brightness $\mu_0$ over 4 dex. This scaling relation has been independently confirmed by \citet{ErrozFerrer2016}, who show that other structural parameters (stellar mass, bulge-to-disk ratio, bar strenght) do \textit{not} correlate with $S_0$ as tightly as $\mu_{0}$. The inner slope $S_0$ scales as the square root of the central dynamical surface density \sdyn, hence the $S_0-\mu_0$ relation provides key information on the ratio between baryons and dark matter (DM) in galaxy centers \citep{Lelli2014}. In this Letter we explore another method to estimate \sdyn. \citet{Toomre1963} provides the relation between the central surface density and the rotation curve $V(R)$ of self-gravitating disks. His Equation 16 states \begin{equation}\label{Eq:Toomre} \Sigma_{\rm dyn}(0) = \dfrac{1}{2\pi G} \int_{0}^{\infty} \dfrac{V^{2}(R)}{R^{2}}dR, \end{equation} where $G$ is Newton's constant. This formula holds as long as the baryonic disk is nearly maximal (Sect.\,\ref{sec:SigmaDyn}). \sdyn\ has key advantages over $S_0$: (i) it is independent of fitting procedures or extrapolations for $R \rightarrow 0$, and (ii) the disk thickness can be easily taken into account. We consider 135 galaxies from the SPARC (Spitzer Photometry and Accurate Rotation Curve) database \citep[][hereafter Paper I]{Lelli2016a}. These objects have both high-quality \hi rotation curves and $Spitzer$ [3.6] surface photometry. The availability of [3.6] images is a major improvement over previous studies \citep{Lelli2013} since the near-IR surface brightness provides the best proxy to the central stellar surface density \sstar. We find that \sdyn\ tightly correlates with \sstar\ over 4 dex (Figure\,\ref{fig:Central}), even for LSB galaxies that appear not to be self-gravitating.

In this Letter we establish a scaling relation between the central dynamical density and the central stellar density of disk galaxies. HSB galaxies are consistent with unity, suggesting that they are self-gravitating and baryon dominated in the inner parts. LSB galaxies systematically deviate from the 1:1 line, indicating that the DM contribution progressively increases but remains tightly coupled to the baryonic one. This central density relation represents a key testbed for cosmological models of galaxy formation.

1607.02145

1607

1607.05286_arXiv.txt

We present deep near-infrared photometry and spectroscopy of the globular cluster 2MASS-GC\,03 projected in the Galactic disk using MMIRS on the Clay telescope (Las Campanas Observatory) and VISTA Variables in the Via Lactea survey (VVV) data. Most probable cluster member candidates were identified from near-infrared photometry. Out of ten candidates that were followed-up spectroscopically, five have properties of cluster members, from which we calculate $<$[Fe/H]$> = -0.9 \pm 0.2$ and a radial velocity of $<v_{\rm r}> = -78 \pm 12$\,km/s. A distance of 10.8\,kpc is estimated from 3 likely RR\,Lyrae members. Given that the cluster is currently at a distance of 4.2\,kpc from the Galactic center, the cluster's long survival time of an estimated $11.3\pm 1.2$\,Gyr strengthens the case for its globular-cluster nature. The cluster has a hint of elongation in the direction of the Galactic center.

Galactic globular clusters (GCs) have played a role in diverse areas of study, from the evolution of stellar populations to the hierarchical formation of galaxies. The current population of known GCs in the Milky Way is composed of $\sim 160$ of these systems \citep{Harris2010}, although it has long been suggested \citep[e.g.][]{Ashman1992} that a fraction of the total number of Galactic GCs might be hidden in lines of sight with high extinction (i.e. disk, bulge) or have intrinsic properties that make them difficult to detect. The search for new Galactic GCs has only returned a small number of the expected missing GCs. Some of these newly identified GCs have properties mid-way between a GC and an ultra-faint dwarf galaxy \citep{Carraro2005,Willman2005,Koposov2007,Longmore2011,Munoz2012,Balbinot2013,Laevens2014,Kim2015}. Most of these new detections were possible thanks to the arrival of wide-sky photometric surveys such as Sloan Digital Sky Survey \citep[SDSS,][]{York2000}, the Two Micron All Sky Survey \citep[2MASS,][]{Skrutskie2006} and the VISTA Variables in the Via Lactea (VVV) Public Survey \citep{Minniti2010,Saito2012}. The latter survey reported the discovery of VVV CL\,001 \citep{Minniti2011}, VVV CL\,002 \citep{MoniBidin2011} and several other GC candidates \citep{Borissova2014}. It is well known that the Galactic GC system contains a metal-rich subsystem \citep{Zinn1985} which is mainly composed of GCs associated with the Galactic bulge. Among the 13 GCs found with $|b| < 2^{\circ}$ (8$\%$ of the Galactic population), 11 lie projected on the Galactic bulge, while only two, 2MASS-GC\,03 and GLIMPSE\,01 \citep{Simpson2004,Kobulnicky2005,Davies2011}, are likely associated with the Milky Way disk. However, the nature of GLIMPSE\,01 remains unclear \citep[see][]{Davies2011}. In contrast, numerous young cluster candidates have been discovered in the Galactic plane using infrared photometry \citep[e.g.][]{Froebrich2007a,Borissova2011,Borissova2014,Solin2014}. Therefore, the scarce GCs at low Galactic latitudes represent a peculiar population of stellar systems in a region of the Milky Way dominated by the presence of young stellar clusters. \begin{figure*} \begin{center} \includegraphics[scale=0.37]{cluster_Ks.png} \includegraphics[scale=0.37]{cluster_MIR.png} \caption[Imagenes]{VVV $K_S$ ({\it left}) and WISE false color ({\it right}) images for GC\,03. The red circles show the position of the stars observed spectroscopically. The false color image was constructed using the WISE filters {\it W}1 (blue), {\it W}2 (green) and {\it W}3 (red). The image sizes are 7\,arcmin $\times$ 7\,arcmin.} \label{fcr1735_VVV_WISE} \end{center} \end{figure*} \citet{Froebrich2007a} first identified 2MASS-GC\,03 as a local density enhancement in 2MASS photometry maps, listing it as \#1735 (or FSR\,1735) in their catalog. Follow-up $JHK_s$ photometry of stars associated with that overdensity \citep{Froebrich2007b} allowed the authors to suggest that FSR\,1735 is actually a GC at $d_{\odot} \sim 9$\,kpc, although they point out that a definitive determination of the object's nature would require follow-up observations. A systematic analysis of the cluster was performed by \cite{Kharchenko2013}, as a part of a global investigation on the Galactic star-cluster system. Using the 2\,MASS catalogs and PPMXL Catalogue of Positions and Proper Motions on the ICRS \citep{Roeser2010}, they estimated the metallicity, distance and age of FSR\,1735 as [Fe/H] = -0.6, $d_{\odot} \sim 8.8$\,kpc and $t = 9.9$\,Gyr, respectively. Because of its discovery in 2\,MASS data, FSR\,1735 has the SIMBAD denomination 2\,MASS-GC03 and it also corresponds to the candidate number 71 in the \cite{Solin2014} catalog, which contains new star cluster candidates discovered in the VVV field of view. In this paper, we will use the SIMBAD designation, namely 2\,MASS-GC\,03 or simply GC\,03. As it can be seen in Figure \ref{fcr1735_VVV_WISE}, we observe a clear overdensity of stars in the VVV $K_s$-image, but it is not associated with mid-infrared nebulosity in WISE images (Figure \ref{fcr1735_VVV_WISE}, right panel), confirming the Froebrich et al.\ conclusion of small amounts of dust and gas in that area. In this paper, we present deep MMIRS and VVV photometry and low-resolution near-infrared spectroscopy of several stars in order to confirm the globular nature of the cluster GC\,03 and derive accurate values for the fundamental parameters that describe that system.

We have derived fundamental parameters for the stellar cluster 2MASS-GC\,03 using MMIRS photometry and spectroscopy and VVV data. Out the 10 stars for which we have obtained spectra, 5 were considered as cluster members based on their proper motions derived from VVV photometry. We have estimated metallicities and radial velocities for field and candidate stars and set mean values for the cluster in $<[Fe/H]> = -0.9$ and $<v_{\rm r}> \sim -78$\,km\,s$^{-1}$. With that information, we have established that 2MASS-GC\,03 is an intermediate-age system with $t = 11.3$\,Gyr at $d_{\odot} = 10.8$\,kpc. Its location within the Galaxy and overall properties (summarized in Table~2) confirms its globular nature.

1607.05286

1607

1607.00006_arXiv.txt

\setcounter{footnote}{1} We report the discovery of five new transiting hot Jupiter planets discovered by the HATSouth survey, \hatcurb{31} through \hatcurb{35}. These planets orbit moderately bright stars with $V$ magnitudes within the range \hatcurCCtassmvshort{33}--\hatcurCCtassmvshort{32}\,mag while the planets span a range of masses \hatcurPPmshort{31}--\hatcurPPmshort{35}\,\mjup, and have somewhat inflated radii between \hatcurPPrshort{33}--\hatcurPPrshort{31}\,\rjup. These planets can be classified as typical hot Jupiters, with \hatcurb{31} and \hatcurb{35} being moderately inflated gas giant planets with radii of \hatcurPPr{31}\,\rjup{} and \hatcurPPr{35}\,\rjup, respectively, that can be used to constrain inflation mechanisms. All five systems present a higher Bayesian evidence for a fixed circular orbit model than for an eccentric orbit. The orbital periods range from \hatcurLCP{35}\,day for \hatcurb{35}) to \hatcurLCP{31}\,day for \hatcurb{31}. Additionally, \hatcurb{35} orbits a relatively young \hatcurISOspec{35}\ star with an age of \hatcurISOage{35}\,Gyr. We discuss the analysis to derive the properties of these systems and compare them in the context of the sample of well characterized transiting hot Jupiters known to date. \setcounter{footnote}{0}

\label{sec:introduction} Planets that eclipse their host star during their orbit are key objects for the study of exoplanetary systems. The special geometry of transiting extrasolar planets (TEPs) enables measurements of not only the planet size but other important physical parameters, such as their masses and densities, and the characterization of the alignment between the orbital axis of a planet and the spin axis of its host star through the Rossiter-McLaughlin effect. The majority of well-characterized TEPs have been discovered by wide-field photometric surveys, including Kepler \citep{2010Sci...327..977B}, the Wide Angle Search for Planets \citep[WASP;][]{2006PASP..118.1407P}, the Hungarian-made Automated Telescope Network \citep[HATNet;][]{2004PASP..116..266B,bakos:2013:hatsouth}, COnvection ROtation and planetary Transits \citep[CoRoT;][]{2008A&A...482L..17B}, and the Kilodegree Extremely Little Telescope survey \citep[KELT;][]{2012ApJ...761..123S}. The known sample of exoplanets present a great diversity of orbital and planetary parameters. Extending the sample of close-orbiting TEPs is a key goal of ground-based surveys as they allow for a large array of additional observational measurements, such as information about the chemical composition of the atmospheres of the planets using emission and transmission spectroscopy for sufficiently bright targets. The HATSouth survey \citep{bakos:2013:hatsouth} has been designed to increase the sample of well-characterized TEPs. Some recent examples of planets discovered by HATSouth are HATS-18b \citep{2016arXiv160600848P} and HATS-25b through HATS-30b \citep{2016arXiv160600023E}. A full list of TEPs discovered by the HATSouth survey, along with all discovery and follow-up light curves, can be found at \url{http://hatsouth.org/}. In this paper we present five new transiting planets discovered by the HATSouth network around moderately bright stars: \hatcurb{31} through \hatcurb{35}. In Section~\ref{sec:obs} we describe the photometric transit detection with HATSouth, as well as the data analysis methods and the procedures used to confirm the planetary nature of the transit signal using follow-up spectroscopic and photometric observations. In Section~\ref{sec:analysis} we describe the analysis carried out to rule out false positive scenarios that could mimic a planetary signal, and to ascertain the stellar and planetary parameters. We discuss the implication of our results and compare them with all known transiting hot Jupiters to date in Section~\ref{sec:discussion}.

\label{sec:discussion} \begin{figure*} \centering \begin{tabular}{cc} \includegraphics[width=\columnwidth]{mass-rad} & \includegraphics[width=\columnwidth]{eqt-rad} \end{tabular} \caption{\emph{Left panel}: Mass-radius diagram of all known transiting hot Jupiters, i.e. planets with masses of $0.1M_J<M<5M_J$ and periods $P<10$ days, with precisely measured masses and radii. \hatcurb{31} is shown with a red circle, \hatcurb{32} with a red square, \hatcurb{33} with a red triangle, \hatcurb{34} with a red diamond and \hatcurb{35} with a red star. Isodensity curves with density of 0.1, 0.25, 0.5 and 1 $\rho_\mathrm{J}$ are shown by the dashed lines. \emph{Right panel}: the mass-density diagram \emph{Right panel}: Planet equilibrium temperature versus radius for the same sample of transiting hot Jupiters plotted in the mass-radius diagram. \hatcurb{31} through \hatcurb{35} are represented with the same symbols. } \label{fig:mrd} \end{figure*} \begin{figure} \centering \includegraphics[width=\columnwidth]{dens-mass} \caption{Mass-density diagram of all known transiting hot Jupiters, planets with masses of $0.1M_J<M<5M_J$ and periods $P<10$ days with well-characterized masses and radii, taken from the NASA Exoplanet Archive. Red data points as per Fig.~\ref{fig:mrd}. } \label{fig:dens-mass} \end{figure} We have presented five new transiting hot Jupiters, \hatcurb{31}\ through to \hatcurb{35}, discovered by the HATSouth survey. Our analysis of the combined photometric and spectroscopic data rules out the possibility that these transit detections are blended stellar eclipsing binary systems, and we conclude that these objects are transiting planets. In Figure~\ref{fig:mrd} we show the mass-radius and equilibrium temperature versus radius diagrams of all known transiting hot Jupiters with well determined masses and radii discovered to date retrieved from the NASA Exoplanet Archive\footnote{\url{http://exoplanetarchive.ipac.caltech.edu/}} on 2016 May 30, with \hatcurb{31} through \hatcurb{35} superimposed in red. From the mass-radius diagram, the planets presented in this paper can be classified as typical hot Jupiters. \hatcurb{31}, \hatcurb{32} and \hatcurb{34} are slightly less massive than Jupiter with \hatcurPPmlong{31}\,\mjup, \hatcurPPmlong{32}\,\mjup{} and \hatcurPPmlong{34}\,\mjup, respectively. However the radius values of the five objects, all higher than that of Jupiter, vary between \hatcurPPrlong{33}\,\rjup{} for \hatcurb{33} to \hatcurPPrlong{31}\,\rjup{} for \hatcurb{31}. The planet equilibrium temperature versus radius diagram in the right panel of Figure~\ref{fig:mrd} shows that the planets agree with previously observed general trends. \hatcurb{31} and \hatcurb{35} have higher equilibrium temperature, in the range \hatcurPPteff{31} K to \hatcurPPteff{35} K, compared with the other three objects. From the mass-radius and equilibrium temperature-radius diagrams it can be seen that \hatcurb{31} and \hatcurb{35} reside in a different region than the other three planets. \hatcurb{31} and \hatcurb{35} have a radius of \hatcurPPrlong{31}\,\rjup{} and \hatcurPPrlong{35}\,\rjup{}, respectively, and are therefore moderately inflated planets, while \hatcurb{32}, \hatcurb{33} and \hatcurb{34} have radii from \SIrange{1.2}{1.4}{\rjup}, which is close to the mean radius of known hot Jupiters. This indicates that the inflated radii are linked to the increased irradiation from their parent star. All of the discovered planets have a period below the mean value of transiting hot Jupiters, with the shortest period of the sample being \hatcurLCP{35}\, days for \hatcurb{35}. In Figure~\ref{fig:dens-mass} we show the planet density against mass for \hatcurb{31}--\hatcurb{35} in the context of all known exoplanets with well-characterized densities. \hatcurb{31} is the lowest density planet of the objects presented in this paper, with a mean density of \hatcurPPrho{31}\,\gcmc, while the other objects have typical densities between \hatcurPPrho{35}\,\gcmc{} and \hatcurPPrho{33}\,\gcmc, for objects of their mass and period.

1607.00006

1607

1607.07410_arXiv.txt

We present results of a dark matter search performed with a 0.6\,kg\,d exposure of the DAMIC experiment at the SNOLAB underground laboratory. We measure the energy spectrum of ionization events in the bulk silicon of charge-coupled devices down to a signal of 60\,eV electron equivalent. The data are consistent with radiogenic backgrounds, and constraints on the spin-independent WIMP-nucleon elastic-scattering cross section are accordingly placed. A region of parameter space relevant to the potential signal from the CDMS-II Si experiment is excluded using the same target for the first time. This result obtained with a limited exposure demonstrates the potential to explore the low-mass WIMP region ($<$10\,\gev ) with the upcoming DAMIC100, a 100\,g detector currently being installed in SNOLAB.

Introduction} The DAMIC (dark matter in CCDs) experiment~\cite{Barreto2012264} employs the bulk silicon of scientific-grade charge-coupled devices (CCDs) to detect coherent elastic scattering of weakly interacting massive particles (WIMPs), highly motivated candidates for being the dark matter in the Universe~\cite{Kolb:1990vq, *Griest:2000kj, *Zurek:2013wia}. By virtue of the low readout noise of the CCDs and the relatively low mass of the silicon nucleus, DAMIC is particularly sensitive to low-mass WIMPs in the Galactic halo with masses in the range 1--20\,\gev , which would induce nuclear recoils of keV-scale energies. Throughout 2015, dark matter search data were acquired in the SNOLAB laboratory with 8\,Mpix CCDs (2.9\,g each) in dedicated one- to two-month-long periods. In this paper, we present results from a 0.6\,kg\,d exposure reaching a sensitivity to the spin-independent WIMP-nucleus elastic-scattering cross section $<$\powero{-39}\,cm$^2$ for WIMPs with masses $>$3\,\gev\ and directly probing the signal excess in the CDMS II silicon experiment~\cite{Agnese:2013rvf} with the same nuclear target. This work establishes the calibration and stable performance of the detector, the understanding of backgrounds, and the analysis techniques necessary for the full deployment of the eighteen 16\,Mpix CCDs (5.8\,g each) of DAMIC100.

1607.07410

1607

1607.07626_arXiv.txt

We present large-scale trends in the distribution of star-forming objects revealed by the Hi-GAL survey. As a simple metric probing the prevalence of star formation in Hi-GAL sources, we define the fraction of the total number of Hi-GAL sources with a 70\micron counterpart as the ``star-forming fraction'' or SFF. The mean SFF in the inner galactic disc (3.1\,kpc$\,< R_{\rm GC}\,<$ 8.6\,kpc) is 25\%. Despite an apparent pile-up of source numbers at radii associated with spiral arms, the SFF shows no significant deviations at these radii, indicating that the arms do not affect the star-forming productivity of dense clumps either via physical triggering processes or through the statistical effects of larger source samples associated with the arms. Within this range of Galactocentric radii, we find that the SFF declines with $R_{\rm GC}$ at a rate of $-$0.026$\pm$0.002 per kiloparsec, despite the dense gas mass fraction having been observed to be constant in the inner Galaxy. This suggests that the SFF may be weakly dependent on one or more large-scale physical properties of the Galaxy, such as metallicity, radiation field, pressure or shear, such that the dense sub-structures of molecular clouds acquire some internal properties inherited from their environment.

Molecular gas is a dominant component in the interstellar medium (ISM) and the principal location of star formation. Clouds of molecular gas take on a hierarchical structure throughout the Milky Way, where the densest clumps of gas account for roughly 5-10\% of the total mass in a typical cloud \citep{Battisti2014,Ragan2014}. With the exception of the central molecular zone (CMZ), this appears to be a universal property of molecular clouds on average, regardless of a cloud's proximity to a spiral arm where, globally, molecular gas is concentrated \citep{Eden2012,Eden2013}. The conditions of gas in molecular clouds in the Milky Way is often characterised by CO emission, and large sections of the Galactic plane have now been surveyed in a number of CO transitions and isotopologues. For example, \citet{Roman-Duval2010} use the Galactic Ring Survey (GRS) $^{13}$CO (1-0) data and find a steep decline in the Galactic surface mass density of molecular clouds with Galactocentric radius ($R_\mathrm{GC}$), which extends to the outer Galaxy until a truncation point of the molecular disc at $R_\mathrm{GC}$ = 13.5\,kpc \citep{Heyer1998}. The excitation temperature of CO declines with $R_\mathrm{GC}$, which may link to interplay between the slow decline of the cooling rate (due to lower metallicity) and more rapid decline of the heating rate (attributed to a decrease in star formation rate [SFR]) with $R_\mathrm{GC}$ \citep{Roman-Duval2010}. Thermal emission from interstellar dust grains which follow the gas distribution provide a secondary tracer of ISM structure and properties. \citet{Sodroski1997} derive a dust temperature gradient with $R_\mathrm{GC}$, owing in part to the metallicity gradient in the Galactic disc \citep{Lepine2011b} and to variations in the strength of the interstellar radiation field with $R_\mathrm{GC}$. The links between these trends and star formation, however, remain tenuous. Data from the {\em Herschel} Hi-GAL survey \citep{Molinari2010b} provides a new high-resolution perspective on the distribution of dust which is necessary to distinguish between active and quiescent molecular clouds throughout the Galaxy and to quantify their star formation activity in detail. In this paper, we examine global trends in the properties of Hi-GAL sources with Galactocentric radius. \begin{figure} \includegraphics[width=0.5\textwidth]{hou_spiral.pdf} \caption{\label{f:hou_spiral} A bird's eye schematic view of the Milky Way. The solid coloured spiral features correspond to the \citet{HouHan2014} four-arm spiral arm model based on HII regions, where red is the Norma arm, yellow is the Perseus arm, green is the Sagittarius-Carina arm, blue is the Scutum-Centaurus arm, and black is the local arm. The grey ellipse represents the area influenced by the Galactic bar. The dashed lines show the \citet{Reid2016} loci of the spiral arms, colour-coded in kind as above. The star indicates the position of the Sun. The grey shaded regions show the area bound by the longitude limits of the current catalogue. } \end{figure}

\subsection{What does the star-forming fraction (SFF) mean?} Hi-GAL sources detected at the four longest {\em Herschel} wavebands -- 160, 250, 350 and 500\,\micron -- are equivalent to the submillimetre-continuum sources detected in surveys such as ATLASGAL \citep{ATLASGAL} or the JCMT Plane Survey \citep[JPS;][]{Moore2015}. For the most part, such objects have virial ratios clustered around the critical value \citep{Urquhart2014c} and therefore at least half are potentially star-forming clumps. As explained above, we take it that those that have 70\,\micron emission are already actively star-forming. The SFF may therefore be considered as the fraction of dense clumps with embedded YSOs. If all the IR-dark clumps detected by Hi-GAL were to evolve into IR-bright sources, the SFF would give the relative timescales of the pre-stellar and protostellar stages, but this is not necessarily the case. Those not currently forming stars may either go on to form stars in the future or may dissipate without doing so. The SFF must be somewhat related to the evolutionary state of clumps, as traced by e.g. the ratio of their infrared luminosity ($L_{\rm IR}$) to mass \citep[cf.][]{Molinari2008, Urquhart2013a}, and has some dependence on the average clump mass, since the highest-mass clumps have an undetectably short infrared-dark lifetime \citep{Motte_cygX,Urquhart2014c}. The mean SFF is therefore set by the relative timescale of the IR-bright protostellar stage to that of the pre-stellar stage, multiplied by the average fraction of productive dense cores. The measured mean SFF value of 0.25 happens to be consistent with equal timescales and 50\% of clumps being eventually star-forming \citep[see also][]{Moore2015}. Relative variations in SFF indicate changes (in both, but presumably mainly the latter) and/or variations in the time gradient of the SFR on timescales similar to the clump lifetime ($\sim 10^5$ years; if SFR is increasing, SFF will be low, since there will be more bound starless clumps, and vice versa). The SFF is therefore a quantity related to the current star-formation efficiency (SFE) within dense, potentially star-forming clumps, being the fraction of dense clumps that are forming stars within the timescale set by submillimetre and far-IR detection. The SFF does not, however, tell us the actual conversion efficiency of clump mass into stellar mass, and is thus not a star-forming {\it efficiency}, strictly speaking. A change in SFF with location may indicate a spatial variation in environmental factors influencing the probability that a clump will form stars. Such factors may include the availability of dense molecular gas \citep{Roman-Duval2016}, turbulent pressure \citep{Wolfire2003}, local magnetic-field strength \citep{HeilesTroland2005}, or the presence of a triggering agent such as a wind- or radiation-driven bubble \citep{Bertoldi1989,Bisbas2009}. \subsection{The SFF associated with spiral arms} Our knowledge of the spiral structure of the Galaxy comes from high-resolution surveys of the Milky Way plane, which have informed various efforts to model Galactic structure \citep[e.g.][]{Dame2001,Roman-Duval2010,Vallee2014c,HouHan2014,Reid2014a,Reid2016}. The Milky Way has either two or four arms \citep[depending on the choice of tracer;][]{Robitaille2012} and exhibits spatial offsets between tracers of cold molecular gas and those of active star formation \citep[e.g.][]{Vallee2014a}. While these studies have shown that the spiral arms are undoubtedly where material is concentrated in the Galaxy, studies of SFE metrics across the Galactic plane have found no compelling evidence of variation associated with the spiral arms \citep{Moore2012,Eden2012,Eden2013,Eden2015}, though small number statistics were a limitation to these studies. Nevertheless, we find similar signatures in our results. Figure~\ref{f:Nsources_area_Rgc} shows that in terms of total number of sources (top histogram panels), there are enhancements at the spiral arm radii. In the north, the Scutum tangent at $R_{\rm GC} \sim 4.5$\,kpc, the Sagittarius arm between $5.5 < R_{\rm GC} < 6.5$\,kpc show clear peaks in total distribution. The southern Norma ($R_{\rm GC} \sim 4.7$\,kpc) and Centaurus ($R_{\rm GC} \sim$ 6.5\,kpc) tangents are also peaks in total source surface density. One of the largest discrepancies between competing spiral arm models is the path of the Carina arm. If it lies between 7\,kpc $< R_{\rm GC} <$ 8\,kpc as R16 suggest, is not a peak in overall source surface density or SFF. The HH14 model puts the Carina arm about 1-2\,kpc further from the Galactic centre in \ $R_{\rm GC}$, in which case the current catalogue misses the tangent longitude and the Carina arm is too poorly-sampled for our consideration in this paper. Turning to the lower panels in Figure~\ref{f:Nsources_area_Rgc} (SFF versus $R_{\rm GC}$), we see that the SFF at these radii do not exhibit compelling peaks (i.e. $>3\sigma$ deviation from the mean SFF), with the possible exception of one bin near the northern Sagittarius arm tangent ($R_{\rm GC} \sim$ 6.5\,kpc) where the SFF is $\sim$0.31 ($\sim$2$\sigma$), however since the adjacent bins lack any elevation in SFF, this peak should be taken cautiously. Otherwise and interestingly, if anything, the SFF exhibits weak depressions in SFF at the Perseus, (southern) Centaurus and Carina arm radii. That the Sagittarius arm (at $R_\mathrm{GC} \sim$ 6 -- 6.5\,kpc in the North) may be unremarkable in SFF versus $R_{\rm GC}$ is of particular interest. This arm is prominent in CO\,(3-2) and therefore has abundant molecular gas content at these longitudes \citep{Rigby2016}. It is also a strong feature in the RMS source distribution \citep{Urquhart2014a}, gas temperature \citep{Roman-Duval2010}, and the ratio of infrared luminosity to clump mass \citep[$L_{\rm IR}/M_{\rm clump}$][]{Moore2012,Eden2015} suggestive of enhancement in the SFE, albeit on the kiloparsec scales probed by earlier surveys. This, however, may be a consequence of local variations (e.g. a few high-luminosity sources) which are not captured by the SFF metric, which is based strictly on source count surface densities. The lack of significant change in SFF across the spiral arms indicates that the arms have little effect on the star-forming productivity of dense clumps or on the average evolutionary state of star-forming clumps, and no change in the latter across the spiral arms where the line of sight is along a tangent. The latter might be surprising since a lag between dust/gas-traced and star-traced arms is predicted by the density-wave theory of spiral arms and has been reported several times in qualitative studies of nearby face-on galaxies, but not supported by more recent observational work \citep[see][and references therein]{Foyle2011}. \begin{figure} \includegraphics[width=0.5\textwidth]{sff_dist.pdf} \caption{\label{f:sff_heliodist} The star forming fraction (SFF) as a function of heliocentric distance in kiloparsecs. The result of a linear weighted least squares fit to the trend is shown in the green solid line, the slope of which is $+$0.007$\pm$0.001 per kpc.} \end{figure} \begin{figure} \includegraphics[width=0.5\textwidth]{sff_Rgc_heliodist_chunks.pdf} \caption{\label{f:sff_heliodist_bias} SFF plotted as a function of $R_{\rm GC}$ using sources from 2\,kpc-wide heliocentric distance bins specified in the upper left of each panel. The best weighted least-squares linear fit is shown in the red line, the slope (m) of which is shown in each panel. Further statistics can be found in Table~\ref{t:sff_rgc_chunks}.} \end{figure} \begin{table} \caption{SFF versus $R_{\rm GC}$ in variable heliocentric distance ranges. \label{t:sff_rgc_chunks}} \begin{center} \begin{tabular}{crcc} $D$ range & $N_{\rm tot}$ & mean SFF & slope$^a$ \\ (kpc) & & & (kpc$^{-1}$) \\ \hline 4 $< D <$ 6 & 5865 & 0.226 & $-$0.015$\pm$0.007 \\ 5 $< D <$ 7 & 5832 & 0.227 & $-$0.037$\pm$0.006 \\ 6 $< D <$ 8 & 7937 & 0.214 & $-$0.030$\pm$0.009 \\ 7 $< D <$ 9 & 9802 & 0.217 & $-$0.042$\pm$0.007 \\ 8 $< D <$ 10 & 10845 & 0.247 & $-$0.053$\pm$0.006 \\ 9 $< D <$ 11 & 10529 & 0.264 & $-$0.045$\pm$0.007 \\ 10 $< D <$ 12 & 9981 & 0.274 & $-$0.024$\pm$0.008 \\ \hline \end{tabular} \end{center} $^a$ Slope of the SFF versus $R_{\rm GC}$ relation in 2\,kpc-wide heliocentric distance bins, as shown in Figure~\ref{f:sff_heliodist_bias}, fit between 3.1\,kpc $<$ $R_{\rm GC}$ $<$ 8.6\,kpc. \end{table} \subsection{Potential biases in measuring the SFF} As Hi-GAL provides us with an unprecedented number of uniformly-surveyed sources, we expect that any bias in our findings is distance-related. Our study considers sources out to $D\sim$20\,kpc heliocentric distance, but given the longitude limits, this translates to a much smaller range of $R_{\rm GC}$, such that 96\% of these sources fall within the 3.1\,kpc $< R_{\rm GC} <$ 8.6\,kpc range used in the above analysis. The SFF as a function of heliocentric distance is shown in Figure~\ref{f:sff_heliodist}. There is a shallow but statistically significant slope of $+0.007 \pm 0.001$ which suggests a distance-related bias affecting the sample. There are several possible effects at work here. First, the physical size corresponding to the resolution element increases with distance, such that (assuming a uniform distribution of sources on average) the number of sources overlapping with the beam will increase with distance and also the likelihood that one of those sources is 70\,$\mu$m-bright, tending to increase the SFF with heliocentric distance (see Figure \ref{f:sff_heliodist_bias} and Table \ref{t:sff_rgc_chunks}). Second, the typical spectral energy distributions of both starless and protostellar Hi-GAL sources \citep[e.g.][]{Giannini2012} are intrinsically brighter at wavelengths longer than 160\micron. If a 70\micron counterpart is detected, it is typically a ``weaker'' (i.e. fewer $\sigma$ above rms) detection \citep[see Fig 3 in][]{Molinari2016b}. Thus, at large distances, sources are more readily detected at longer wavelengths, resulting in an increasing fraction of genuine protostellar sources being mis-classified as starless, effectively reducing the SFF with distance. Another potential related bias is the effect of distance on the average observed clump mass and luminosity of sources. At large distances, a higher fraction of the sources will be higher mass, which have shorter infrared-dark lifetimes \citep{Urquhart2014c}. In any case the gradient of this relationship is only one third of that in the relationship of SFF with $R_{\rm GC}$ and cannot be the cause of the latter, especially since the relationship between $R_{\rm GC}$ and $D$ is not a one-to-one correlation. We can get a sense of the impact this shallow distance bias may have on the trend with $R_{\rm GC}$ by looking at the SFF versus $R_{\rm GC}$ using objects confined to narrow heliocentric distance bins. The SFF as a function of $R_{\rm GC}$ using sources within 2\,kpc intervals\footnote{The selection of the 2\,kpc width was made to ensure a good sample size ($N > 5000$) was available.} of heliocentric distance is shown in Figure~\ref{f:sff_heliodist_bias}. As expected from Figure~\ref{f:sff_heliodist}, the mean SFF increases slightly as the distance centre moves outward. We note that not only do all distance intervals show a significant anti-correlation, also the most populated distance intervals (covering distances between 7 and 11\,kpc) exhibit a steeper slope than the full sample value by a factor of $\sim$2, lending credibility to the overall robustness of the trend. \begin{figure} \includegraphics[width=0.5\textwidth]{radial_trends_compare.pdf} \caption{{\bf (a)} The ratio of dense ($^{13}$CO-emitting) gas surface density to the total ($^{12}$CO + $^{13}$CO) gas surface density \citep[from Figure 12 of][]{Roman-Duval2016}. The black line shows the mean, and the light grey shaded area represents the total error budget. {\bf (b)} The median dense gas mass fraction (DGMF, defined as the ratio mass traced by sub-millimetre dust emission to the mass in the parent cloud traced by $^{13}$CO) reported in \citet{Battisti2014} as a function of $R_{\rm GC}$ plotted with the standard error of the DGMF in each 0.5\,kpc-wide bin. {\bf (c)} The SFF gradient with $R_{\rm GC}$ from Figure~\ref{f:Nsources_area_Rgc}. \label{f:radial_trends}} \end{figure} \subsection{What drives the gradient in SFF with $R_{\rm GC}$?} It is far from clear what the physical origin of a gradual decline in SFF with $R_{\rm GC}$ over 5\,kpc might be. Since star formation is observed to be closely correlated with dense gas \citep[e.g.][]{Lada2010}, one might expect the SFF to be greater where the fraction of dense gas is higher. On kiloparsec scales, \citet{Roman-Duval2016} show that the fraction of ``dense'' gas -- defined as the fraction of mass in $^{13}$CO out of the ``total'' molecular mass (traced by $^{12}$CO + $^{13}$CO emission) -- does decline with $R_{\rm GC}$, roughly from 0.9 to 0.6 over the 3\,kpc $< R_{\rm GC} <$ 8\,kpc range (a gradient of $-$0.06\,kpc$^{-1}$, Figure~\ref{f:radial_trends}a). Within individual molecular clouds, however, the fraction of gas at even higher densities -- defined as the ratio of total mass in compact sub-millimetre clumps of dust emission to the total mass of the host cloud traced by $^{13}$CO -- shows no dependence on $R_{\rm GC}$ \citep[Figure~\ref{f:radial_trends}b, see also][]{Eden2013,Battisti2014}. This suggests that once molecular clouds form dense structures (which we observe as sub-millimetre or Hi-GAL clumps), the prevalence of star formation (or SFF) is governed by other internal properties, perhaps inherited from their environment. Below, we focus our discussion on the known large scale radial properties that have been observed in the Galactic disc including metallicity, radiation field, thermal and turbulent pressure and rotational shear. The known negative metallicity gradient in the Galactic disc, when traced by HII regions and OB stars, is in the region of 0.06 -- 0.07 dex kpc$^{-1}$ within the approximate $R_{\rm GC}$ range covered in the present study \citep{Chiappini2001,Lepine2011b}, which translates to a reduction by a factor of 2 over 5\,kpc, while the measured SFF slope of $-$0.026 kpc$^{-1}$ produces only a 13\% decline in 5\,kpc. Reduced metallicity implies lower dust-to-gas ratio and reduced CO/H$_2$ abundance, and so less efficient cooling and turbulent energy dissipation. This might be expected to result in less efficient star formation. However, \citet{GloverClark2012c} predict that, while the fraction of total molecular cloud mass traced by CO may decrease, the star-formation rate within clouds has little sensitivity to the metallicity. \citet{Hocuk2016} also suggest that grain surface chemistry has only a small effect on star formation in molecular clouds. We cautiously note that the SFF traces star formation within clumps and not clouds and is independent of measured clump masses. A radial decrease in radiation-field strength (in both photon intensity and hardness) should offset, to some extent, the reduced shielding that a declining metallicity produces via reduced dust and CO abundance \citep{Sandstrom2013}, so the destruction rate of both these will be less than expected from reduced metallicity alone \citep{GloverClark2012c}. Other potential effects of the radiation field related to star formation include changes in the ionisation fraction, a decrease in which may reduce magnetic-field support of clumps against gravity, and in thermal energy input to the ISM, but both effects are more likely to produce a positive SFF gradient than the observed negative one. \citet{Wolfire2003} estimate the typical thermal pressure in the Galactic plane interval 3\,kpc $<$ $R_{\rm GC}$ $< 18$\,kpc to be $P_{\rm therm}/k \simeq 1.4 \times 10^4\,\exp(-R/5.5\,\mbox{kpc})$\,K\,cm$^{-3}$ in the 3\,kpc $<$ $R_{\rm GC}$ $<$ 18\,kpc range i.e. a shallow declining exponential. They predict a similar but flatter turbulent pressure gradient ($\propto \exp(-R/\mathrm{7.5\,{kpc}})$) between 3 and 10\,kpc. The SFF gradient with $R_{\rm GC}$ is much shallower, however the \citet{Wolfire2003} relation predicts a factor of 3.6 reduction in the pressure between 3\,kpc and 10\,kpc, corresponding to a linear gradient of 0.5$<P>$ per kpc. \citet{Rigby_PhD} find that, while the average thermal pressure in the denser parts of molecular clouds traced by $^{13}$CO ($J = 3\rightarrow 2$) is similar to that of the neutral gas, the turbulent pressures are higher by one or two orders of magnitude. On the other hand, the negative SFF gradient appears inconsistent with the proposition that increased turbulent pressure produces a raised density threshold for star formation \citep{Kruijssen2014}. Rotational shear might be another suspect, contributing to turbulent pressure and the specific angular momentum of clouds and clumps, both of which may affect star-formation productivity. However, shear has been shown to have little effect on SFE within clouds \citep{Dib2012}. While shear is high at inner $R_{\rm GC}$, it decreases rapidly with increasing $R_{\rm GC}$ and is relatively flat and low beyond 3\,kpc, where the SFF decreases steadily. Again, the gradient appears to be in the wrong sense with high SFF where the shear is also higher. As part of their investigation into the low star-formation efficiency in the CMZ, \citet{Kruijssen2014} suggest that the gravitational stability of the Galactic disc is increased inside $\sim$4\,kpc, due to the ratio of Toomre Q parameter to gas surface density. We might therefore expect the SFE (i.e. the conversion of total gas mass to stars) to decrease within this radius, but it is not clear how this might relate to the rate of production of stars in dense clumps measured by the SFF. \citet{Koda2016} show that the molecular gas fraction increases steadily with decreasing $R_{\rm GC}$, but we see in Figure~\ref{f:Nsources_area_Rgc} that the surface density of mass in dense clumps falls rapidly within 4\,kpc \citep{Bronfman1988,Urquhart2014a}. The production of molecular clouds from neutral gas therefore is more efficient at small $R_{\rm GC}$ where the H$_2$/HI ratio is nearly 100\% \citep{Koda2016}. The fraction of molecular gas in the form of dense clumps within these clouds, while more or less steady, on average outside $\sim$4\,kpc, albeit with very large, apparently random variations from cloud to cloud \citep{Eden2012,Eden2013}, falls sharply inside this radius. Of the above mechanisms that might have a connection to the star-formation productivity of these dense clumps traced by Hi-GAL, most should affect the SFF in the opposite sense than is observed. Therefore the connection between the several-kpc-scale consistent gradient in SFF and environmental conditions is obscure, not least because the dense clumps, once formed, might be expected to go on to form stars independent of their environment.

1607.07626

1607

1607.01370_arXiv.txt

The cosmic horseshoe gravitational lens is analyzed using the perturbative approach. The two first order perturbative fields are expanded in Fourier series. The source is reconstructed using a fine adaptive grid. The expansion of the fields at order 2 produces a higher value of the chi-square. Expanding at order 3 provides a very significant improvement, while order 4 does not bring a significant improvement over order 3. The presence of the order 3 terms is not a consequence of limiting the perturbative expansion to the first order. The amplitude and signs of the third order terms are recovered by including the contribution of the other group members. This analysis demonstrates that the fine details of the potential of the lens could be recovered independently of any assumptions by using the perturbative approach.

Strong gravitational lensing offers a unique opportunity to probe the dark halos potential in the vicinity of the Einstein circle. The lensing potential relates directly to the projected matter distribution and as a consequence is a direct measurement of the matter distribution in the lens. However deriving a precise relation between the observations of a gravitational lens and the lensing potential is generally difficult. There are basically two main problems. First the lens may show some degree of complexity and may not be properly described with simple analytical models. And secondly some degree of degeneracy in the modeling of the lens is generally present see for instance, ~\cite{Saha2006}, ~\cite{Wucknitz2002}, ~\cite{Chiba2002}. A solution to the first point is to use a non parametric method for the reconstruction of the potential. For instance potential reconstruction on a grid offers a general model free solution, but the obvious drawback is a dramatic increase in the number of parameters, which in turn aggravates the degeneracy issue. The only solution is to developp a method that offers a direct relation between the arc morphology and the potential. This is precisely what the perturbative approach (~\cite{Alard2007}) achieves. At first order the potential is expanded using two angular functionals, the fields $f_1$ and $\frac{d f_0}{d \theta}$. Each of these fields relates directly to the arc morphology, $f_1$ is related to the mean radial position of the arc, while $\frac{d f_0}{d \theta}$ is related to width of the arc in the radial dimension. This direct relation offers a simple solution to the degeneracy problem. Another aspect is that the reconstruction of the potential is general and does not require any specific assumptions. Some specific examples of reconstruction of arcs systems using the perturbative approach are presented in ~\cite{Alard2009} and ~\cite{Alard2010}. In particular the reconstruction the lens in ~\cite{Alard2009} shows that very complex systems can be handled in this approach. The application of the perturbative method to the cosmic horseshoe gravitational lens offers the possibility to push the reconstruction to a high level of accuracy. The HST data available for this lens offer an excellent resolution and a wealth of details allowing to probe the fine details of the halo dark matter distribution. Let's recall the basic equations of the order one theory by starting from the lens equation, Eq. (~\ref{lens_eq}). \begin{equation} {\bf r_S} = {\bf r} -\nabla \phi \label{lens_eq} \end{equation} Using the equations relating the fields $f_0$ and $f_1$ to the potential, \begin{equation} \left\{ \begin{aligned} \phi(r,\theta) &= \phi_0(r)+\epsilon \psi(r,\theta) \\ \psi(r,\theta) &= f_0(\theta)+f_1(\theta) (r-1) \end{aligned} \right. \label{pot_def} \end{equation} The lens equation Eq. (~\ref{lens_eq}) is expanded to order one in $\epsilon$ (~\cite{Alard2007}): \begin{equation} {\bf r_S} = \left(\kappa_2 \ dr-f_1 \right) {\bf u_r} - \frac{d f_0}{d \theta }\bf {u_{\theta}} \label{pert_0} \end{equation} With: \begin{equation} f_1=\left[\frac{d \psi}{d r} \right]_{r=1} \ \ ; \ \ f_0=\psi(1,\theta) \ \ ; \ \ \kappa_2=1-\frac{d^2 \phi_0}{d r^2} \label{eq_f1_df0} \end{equation} It is useful to introduce the impact parameter of the source, namely ${\bf r_S=\tilde r_S+r_0}$ leading to: \begin{equation} {\bf \tilde r_S} = \left(\kappa_2 \ dr-\tilde f_1 \right) {\bf u_r} - \frac{d \tilde f_0}{d \theta }\bf {u_{\theta}} \label{pert_1} \end{equation} With: $$ \tilde f_i=f_i+x_0 \cos(\theta)+ y_0 \sin(\theta) \ \ \ \ , i=0,1 $$ And the impact paremeter vector ${\bf r_0}$ $$ {\bf r_0}=(x_0,y_0) $$

The perturbative method allowed the reconstruction of the cosmic horseshoe lens without making any particular assumptions. The inclusion of third order terms was dictated only by the necessity to optimize the chi-square. These third order terms were related to the group contribution later in the analysis, but it was not necessary to make any hypothesis about the group contribution when reconstructing the lens. In this case the group is made of quite a large number of galaxies and trying to make an extensive model including each object would require too many parameters. Such models including many parameters are generally plagued with degeneracies issues which is a constant re-occuring problem in conventional gravitational lens analysis. This analysis does not have to include all theses parameters but reduces the lens to a number of fundamental parameters. It is clear that this minimal set of parameters (basically the expansion of the fields to order 3) corresponds to the expectation of many models when the model include more parameters that the fundamental parameters (which would be the case here when modeling all the group). As a consequence it is clear that the perturbative approach is a method of choice for complex systems. The perturbative approach allows a model free, non-degenerate, fast and simple analysis of any gravitational lens system. It is important to note that even in the case of an a priori simple lens system it is useful to apply the perturbative method since this method could reveal unexpected complex contributions. Essentially in the same way that the contribution of the group was discovered without making any initial hypothesis about the presence of the group.

1607.01370

1607

1607.04280_arXiv.txt

We present new analysis of multi-epoch, \textit{H}-band, scattered light images of the AB Aur system. We used a Monte Carlo, radiative transfer code to simultaneously model the system's SED and \textit{H}-band polarized intensity imagery. We find that a disk-dominated model, as opposed to one that is envelope dominated, can plausibly reproduce AB Aur's SED and near-IR imagery. This is consistent with previous modeling attempts presented in the literature and supports the idea that at least a subset of AB Aur's spirals originate within the disk. In light of this, we also analyzed the movement of spiral structures in multi-epoch \textit{H}-band total light and polarized intensity imagery of the disk. We detect no significant rotation or change in spatial location of the spiral structures in these data, which span a 5.8 year baseline. If such structures are caused by disk-planet interactions, the lack of observed rotation constrains the location of the orbit of planetary perturbers to be $>$47 AU.

AB Aur (also known as HD 31293 and SAO 57506, $d=144$ pc) is a young, ($4\pm 1$ Myr) intermediate-mass ($2.4\pm 0.2$ M$_\sun$), Herbig Ae star \citep{van den Ancker,DeWarf2} that is actively accreting material. It is surrounded by a large envelope which extends out to at least 1320 AU and blends into a nearby nebula \citep{Grady}. Within the envelope, a protoplanetary disk (r $\sim 450$ AU; Mannings \& Sargent 1997) surrounds the central star and displays many complex structures \citep{Hashimoto,Fukagawa,Grady}. At 1.6 and 2 $\mu$m the disk has a region of decreased polarized intensity, which is likely due to the scattering geometry of the surface of AB Aur's inclined disk \citep{Oppenheimer,Perrin}. However, more recent HiCIAO \textit{H}-band imagery has shown an additional six regions of decreased polarized intensity (PI) that are not explained by geometric scattering effects \citep{Hashimoto}. These data also revealed an approximately 16 AU wide gap in the disk. Centered at 80 AU from the central star, it is similar to the mid-IR gap inferred by \cite{Honda} from models of their 24.6 $\mu$m imagery, but appears to be different from the gap detected by \cite{Tang}, who find a gap which ends at approximately 110 AU and is about 90 AU wide in 1.3 mm continuum emission. Spiral structures in the disk were first detected in STIS imagery by \cite{Grady} and later imaged in the \textit{H} band by \cite{Fukagawa}, who suggested that they are either maintained by a planet, due to gravitational instabilities within the disk, or the result of the outer envelope replenishing disk material. Subsequent work by \cite{Hashimoto} and \cite{Lin} show evidence of the spiral structures in \textit{H}-band PI imagery, $^{12}$CO (3-2) maps, and at 850 $\mu$m. Both works favor planetary bodies perturbing the disk as an explanation for the formation of the spirals. However, a planet has yet to be detected in the system. \cite{Tang} recently detected four spirals in the CO gas, which are generally not coincident with the spirals detected in the near-IR. For example, \cite{Tang}'s CO S2 spiral appears to share a base with the \textit{H}-band S1 spiral as labeled by \cite{Hashimoto}, but the outer regions of the two spiral arms do not overlap. Similarly, the \textit{H}-band S3 spiral appears to potentially be a continuation of the CO S3 spiral at regions farther from the central star, but there is a gap between the regions where the spirals have been detected; the innermost region detected of the \textit{H}-band S3 spiral and the outermost detected region for the CO S3 spiral do not spatially overlap. Finally, the other two CO spirals, CO S1 and CO S4, appear to have no near-IR counterpart at all. \cite{Tang} recently suggested a different formation mechanism for the spirals whereby a combination of the rotation and infall of material from the envelope allows for the build up of higher density regions along the envelope's bipolar cavities. Due to the system's low inclination angle ($i=22\degree$; Tang et al. 2012), these regions of higher density are projected onto the disk and form the observed spiral structures. They appear as part of the disk, but are actually high density regions of the envelope. However, observational constraints on the density, infall rate, and rotational speed of the envelope do not exist, making it difficult to determine the likelihood of this scenario. Very little is known about AB Aur's envelope, partly because of the difficulty of disentangling the observed contributions of the disk and envelope (e.g. CO lines trace the mid-plane disk structure, but \citealt{Tang} also suggests it traces the envelope morphology). \cite{Pietu} found no evidence for any infall of material in their study of the CO lines. This agrees with the overall conclusions of \cite{Robitaille}, who use two dimensional radiative transfer modeling of the system's SED to place constraints on the mass accretion rate from the envelope. They find that the infall rate might be as high as $10^{-6}$ M$_\sun$ yr$^{-1}$, but their best fit model uses no infall at all, suggesting that the envelope is very optically thin. In this paper, we analyze multi-epoch \textit{H}-band imagery of AB Aur to investigate whether the positions of its spiral arms at these wavelengths have changed with time. A variety of observations of AB Aur have been modeled in the past (including but not limited to, its SED by \citealt{Bouwman} and \citealt{Robitaille}, SED and NIR interferometry by \citealt{tan08}, NIR scattered light imagery by \citealt{Perrin} and \citealt{jan10}, SED and mid-IR imagery by \citealt{Honda}, and mm emission by \citealt{Pietu}). However, self-consistent models of the SED and near-IR imagery of the system, which can be useful in interpreting multi-epoch imagery, have not been extensively explored. Therefore, we first used a three-dimensional, Monte Carlo, radiative transfer code to model the overall behavior of the system's SED and \textit{H}-band imagery. After finding that some of the spiral structures in the system could arise in the disk, as noted in previous works, we compare two sets of archival \textit{H}-band imagery in order to determine if the position of the spirals have changed with time.

We have used the \texttt{HOCHUNK3D} Monte Carlo code to self-consistently model the near-IR imagery and SED of the AB Aur system simultaneously. Our modeling results are consistent with those already present in the literature; a disk-dominated model reproduces many of AB Aur's observed features. This suggests that a spiral formation scenario involving disk material remains a possibility for at least some of the spirals, particularly in the \textit{H} band where the envelope does not significantly contribute to our model's SED or imagery. Given our findings, we analyzed the 2004 and 2009 total light and PI imagery from the CIAO and HiCIAO instruments in the \textit{H} band to measure any potential rotation of the disk's spirals. In the event that the spiral structures are formed due to disk-planet interactions, the spirals' movement, or lack thereof, can constrain the locations of possible planets within the AB Aur system. We find no significant rotation of any of the spiral structures over the 5.8 year baseline between these two datasets. By purposely rotating our data and comparing them to unrotated versions of the data, we find that if the spirals did move, they did so by less than $10\degree$. This suggests that if a planet were responsible for the observed structures, it is in an orbit that is least 47 AU away from the central star.

1607.04280

1607

1607.03375_arXiv.txt

We present deep imaging observations of activated asteroid P/2016 G1 (PANSTARRS) using the 10.4m Gran Telescopio Canarias (GTC) from late April to early June 2016. The images are best interpreted as the result of a relatively short-duration event with onset about $\mathop{350}_{-30}^{+10}$ days before perihelion (i.e., around 10th February, 2016), starting sharply and decreasing with a $\mathop{24}_{-7}^{+10}$ days (Half-width at half-maximum, HWHM). The results of the modeling imply the emission of $\sim$1.7$\times$10$^7$ kg of dust, if composed of particles of 1 micrometer to 1 cm in radius, distributed following a power-law of index --3, and having a geometric albedo of 0.15. A detailed fitting of a conspicuous westward feature in the head of the comet-like object indicates that a significant fraction of the dust was ejected along a privileged direction right at the beginning of the event, which suggests that the parent body has possibly suffered an impact followed by a partial or total disruption. From the limiting magnitude reachable with the instrumental setup, and assuming a geometric albedo of 0.15 for the parent body, an upper limit for the size of possible fragment debris of $\sim$50 m in radius is derived.

P/2016 G1 (PANSTARRS) (hereafter P/2016 G1 for short) was discovered by R. Weryk and R. J. Wainscoat on CCD images acquired on 2016 April 1 UT with the 1.8-m Pan-STARRS1 telescope \citep{Weryk16}. From the derived orbital elements ($a$=2.853 AU, $e$=0.21, $i$=10.97$^\circ$), its Tisserand parameter respect to Jupiter \citep{Kresak82} can be calculated as $T_J$=3.38, so that the object belongs dynamically to the main asteroid belt. The discovery images revealed however a cometary appearance showing clear evidence of a tail extending for approximately 20\arcsec, and a central condensation broader than field stars \citep{Weryk16}. Since the discovery of the object 133P/Elst-Pizarro in 1996 \citep[see e.g.][and references therein]{Hsieh04}, about twenty objects of this class have been discovered, whose activity triggering mechanisms have been proposed to range from impact-induced to rotational disruption, while the activity has been found to last from a few days or less to a few months. In this latter case, sublimation-driven of volatile ices has been invoked as the most likely mechanism of dust production, although gaseous emissions lines have remained undetected to date. For a review of the different objects discovered so far, their orbital stability, and their activation mechanisms, we refer to \cite{Jewitt15}. In this paper we report observations of P/2016 G1 acquired with the 10.4m GTC, and present models of the dust tail brightness evolution from late April to early June 2016. We provide the onset time, the total dust loss, and the duration of the activity, and attempt to identify which physical mechanism is involved in its activation.

1) Asteroid P/2016 G1 was activated 350$^{+10}_{-30}$ days before perihelion, i.e., around 10th February 2016. The activity had a duration of 24$^{+10}_{-30}$ days (HWHM), so that no dust has been produced since our first observation of April 21, 2016. The total dust mass emitted was at least $\sim$2$\times$10$^7$ kg, with a maximum level of activity of $\sim$8 kg s$^{-1}$. These parameters were estimated assuming a power-law size distribution of particles between 1 $\mu$m and 1 cm, with power index of $\kappa$=--3.0, geometric albedo of 0.15, and being emitted isotropically from an otherwise undetected nucleus. The calculated peak and total dust mass are lower limits, as if larger values for the maximum particle size were assumed, these quantities would increase. In addition, if different values for the geometric albedo and/or for the density of the particles were assumed, these quantities would also change accordingly. 2) While the inverted-C feature which is apparent in the out-of-plane images of May 29 and June 8, 2016 is approximately mimicked by the isotropic ejection model, a westward brightness feature cannot be reproduced with that model. However, if some dust mass is ejected from a specified direction right at the time of activation, which turns out to be approximately along the Sun-to-asteroid vector, that feature becomes apparent in the simulations. We speculate that this dust ejection could be associated to an impact, and that the subsequent modeled activity is due to the asteroid partial or total disruption. The impact itself had produced the ejection of some 2.4$\times$10$^5$ kg of dust. 3) The inferred ejection velocities of the dust particles are very small, in the range of 0.015 to 0.14 m s$^{-1}$, with an average value of $\sim$0.08 m s$^{-1}$, corresponding to the escape velocity of an object of 35 m radius and 3000 kg m$^{-3}$ density. An object of that size would have been remained well below the detection limit of the images acquired, so that we cannot assure whether fragments of that size or smaller could exist in the vicinity of the dust cloud. Deeper imaging of the object is clearly needed to assess this fact and to determine the fragment dynamics.

1607.03375

1607

1607.03143_arXiv.txt

We extract cosmological information from the anisotropic power spectrum measurements from the recently completed Baryon Oscillation Spectroscopic Survey (BOSS), extending the concept of clustering wedges to Fourier space. Making use of new FFT-based estimators, we measure the power spectrum clustering wedges of the BOSS sample by filtering out the information of Legendre multipoles $\ell>4$. Our modelling of these measurements is based on novel approaches to describe non-linear evolution, bias, and redshift-space distortions, which we test using synthetic catalogues based on large-volume \Nbody simulations. We are able to include smaller scales than in previous analyses, resulting in tighter cosmological constraints. Using three overlapping redshift bins, we measure the angular diameter distance, the Hubble parameter, and the cosmic growth rate, and explore the cosmological implications of our full shape clustering measurements in combination with CMB and SN Ia data. Assuming a \LambdaCDM cosmology, we constrain the matter density to $\Om = 0.311_{-0.010}^{+0.009}$ and the Hubble parameter to $H_0 = 67.6_{-0.6}^{+0.7} \Unit{km \, s^{-1} \, Mpc^{-1}}$, at a confidence level (CL) of 68 per cent. We also allow for non-standard dark energy models and modifications of the growth rate, finding good agreement with the \LambdaCDM paradigm. For example, we constrain the equation-of-state parameter to $w = -1.019_{-0.039}^{+0.048}$. This paper is part of a set that analyses the final galaxy clustering dataset from BOSS. The measurements and likelihoods presented here are combined with others in \citet{Alam:2016hwk} to produce the final cosmological constraints from BOSS.

\label{sec:intro} Together with observations of the cosmic microwave background (CMB) \changed{and} type-Ia supernova (SN) samples, the analysis of the large-scale structure (LSS) of the Universe based on galaxy redshift surveys has been a prolific source of cosmological information over the past few decades \citep{Davis:1983,Maddox:1990,Tegmark:2004,Cole:2005sx,Eisenstein:2005su,Anderson2012,Anderson:2013oza,Anderson:2013zyy}. \changed{These datasets have helped to stablish the \LambdaCDM model as the current standard cosmological paradigm, and to determine the values of its basic set of parameters with high precision.} The \LambdaCDM model assumes that the energy density of the observable universe is dominated by (pressureless) cold dark matter (CDM) and \changed{a mysterious `Dark Energy' (DE) component that drives the accelerated expansion of the late-time universe, which can be described by a cosmological constant $\Lambda$ or vacuum energy.} Observations of the clustering of galaxies can shed light onto the underlying physical nature of this energy component by probing the growth of structure and the expansion history of the Universe. \changed{Thus, important} recent and ongoing spectroscopic galaxy-redshift surveys, such as the Baryon Oscillation Spectroscopic Survey \citep[BOSS;][]{Dawson:2012va} and its extension eBOSS \citep{Dawson:2015wdb} \changed{are} very valuable probes of the late-time evolution of the Universe. \changed{A major goal of galaxy surveys is to obtain precise measurements of the expansion history of the Universe by means of a feature imprinted into the two-point clustering statistics, the baryonic acoustic oscillations \citep[BAO; for a review see \eg][]{Bassett:2009mm}. The BAO are relics of pressure waves that propagated through the photon-baryon plasma prior to recombination and froze in at the time of last scattering. The interaction between dark and baryonic matter after recombination resulted in a signal of enhanced correlation of density peaks separated by a well defined physical scale, the sound horizon at the drag redshift.} This \changed{scale can be used} as a robust standard ruler for measurements of cosmic distances \citep{Eisenstein:2004an,Seo:2005ys,Angulo:2007fw,Sanchez:2008iw}. \changed{The first detections of the BAO feature \citep{Eisenstein:2005su,Cole:2005sx} relied on angle-averaged clustering statistics. However, separate measurements of the BAO signal along the directions parallel and perpendicular to the line of sight (LOS) can be used to obtain separate constraints on the Hubble parameter $H(z)$ at and the angular diameter distance $D_{\rm A}(z)$ to the mean redshift of the survey by means of the Alcock-Paczynski \citep[AP;][]{AP:1979} test. In this way, anisotropic clustering measurements can break the degeneracy obtained from angle-averaged quantities, which are only sensitive to the average distance $D_{\rm V}(z) \propto(D_{\rm A}(z)^2/H(z))^{1/3}$ \citep*{Hu:2003ti,Wagner:2007in,Shoji:2008xn}.} \changed{The dominant source of anisotropy of the measured clustering signal are the redshift-space distortions (RSD), which are due to the impact of the LOS component of the peculiar velocities of the galaxies on the observed galaxy redshifts. The pattern of RSD provides} additional cosmological information beyond that of the BAO signal. \changed{As, to linear order, peculiar velocities are related to the infall of matter} into gravitational potential wells \citep{Kaiser:1987qv}, the RSD are a probe of the growth of structure. As modifications to general relativity (GR) can change \changed{the growth rate of density fluctuations}, RSD can be used to constrain the theory of gravity \citep[\eg,][]{Guzzo:2008}. However, the galaxy velocity field is highly non-linear even on large scales so that \changed{a} detailed modelling is required \citep[\eg,][]{Scoccimarro:2004tg}. One way to \changed{characterize the anisotropies in the clustering of galaxies} is to use the concept of clustering wedges \changed{introduced by \citet{Kazin:2011xt}, which} correspond to the average the correlation function over wide bins of the LOS parameter, $\mu$, \changed{defined as the cosine of the angle between the total separation vector between two galaxies and the LOS direction.} Anisotropic BAO distance measurements \changed{obtained using clustering wedges were first presented} in \cite{Kazin:2013rxa} as part of the BOSS DR9 CMASS analysis \citep{Anderson:2013oza}, \changed{while \citet{Sanchez:2013uxa,Sanchez:2013tga} performed an analysis of the full shape of the wedges measured from the BOSS DR9 and DR11 galaxy catalogues, respectively.} An alternative \changed{tool to wedges are the Legendre multipole moments} of the two-point statistics \citep{Padmanabhan:2008ag}. The multipoles of the correlation function measured from BOSS DR11 galaxy catalogues were used in several recent galaxy clustering analyses \citep[\eg;][]{Samushia:2013yga,Alam:2015qta,Reid:2014iaa}. In Fourier space, the first anisotropic clustering \changed{studies} \citep[\eg,][]{Blake:2011rj,Beutler:2013yhm} were \changed{performed} on measurements of the Legendre multipoles of the power spectrum \changed{obtained by means of the Yamamoto-Blake estimator \citep{Yamamoto:2005dz,Blake:2011rj}. In this work we extend the concept of clustering wedges to Fourier space and adapt the Yamamoto-Blake estimator to provide a measurement of these statistics.} \changed{We perform an analysis of the full-shape of the Fourier-space clustering wedges measured from the final BOSS galaxy samples \citep{Reid:2015gra}, corresponding to SDSS data release 12 \citep[DR12;][]{Alam:2015mbd}.} In order to make use of new estimators based on fast Fourier transforms \citep[FFT;][]{Bianchi:2015oia, Scoccimarro:2015bla}, we measure the power spectrum clustering wedges of the BOSS sample by filtering out the information of Legendre multipoles $\ell>4$. \changed{Exploiting the signature of BAO and RSD in these measurements, we} derive distance and \changed{growth-of-structure} constraints. We also explore the implications of the full shape of our measurements on the parameters of the standard \LambdaCDM model, as well as its most important extensions, making use also of complementary cosmological information from CMB and \changed{SN} samples. This work is part of a series of papers that analyse the clustering properties of the final BOSS sample. \nocite{Sanchez:2016a} Besides the approach of this work, the analogous full-shape analysis using configuration space wedges is discussed in \citet{Sanchez:2016b}. \changed{Complementary RSD measurements using Fourier and configuration space multipoles are presented in \citet{Beutler:2016arn} and \citet{Satpathy:2016tct}, respectively. Tinker et al. (\Inprep) compares the performance of the different methodologies to extract cosmological information from the full shape of anisotropic clustering measurements. Anisotropic} BAO distance measurements are presented in \citet{Ross:2016gvb} and \citet{Beutler:2016ixs} for configuration and Fourier space, respectively, making use of the linear density-field reconstruction technique \citep{Eisenstein:2006nk,Cuesta:2015mqa}. \changed{\citet{VargasMagana:2016}} investigates the potential sources of theoretical systematics in the anisotropic BAO analysis for the final BOSS galaxy BAO analysis in configuration space. \changed{All final BOSS} analyses are summarised \changed{in \citet{Alam:2016hwk}, where they are combined} into a set of consensus measurements following the methodology described in \citet{Sanchez:2016a}. A different approach is followed in \citet{Salazar-Albornoz:2016psd}, who perform a tomographic analysis by means of angular correlation functions in thin redshift shells. This paper is organised as follows: Section~\ref{sec:boss} describes the \changed{final BOSS DR12 galaxy catalogue and the optimal estimator we use to measure the Fourier-space clustering wedges of this sample, which are the basis for our cosmological constraints.} \changed{This section describes also the methodology we follow to estimate the covariance matrix of our measurements (section~\ref{sec:covariance_matrix}) and to account for the window function of the survey (section~\ref{sec:win_func}).} The model for the Fourier space wedges is discussed in section~\ref{sec:model} \changed{where we} describe the recipe for the non-linear gravitational dynamics, galaxy bias and RSD and analyse the performance of the model using \changed{\Nbody} simulations and synthetic catalogues mimicking the clustering properties of the BOSS galaxy sample. Anisotropic BAO and RSD constraints derived from the full-shape analysis of the DR12 clustering wedges analysis in Fourier space \changed{are} described in section~\ref{sec:BAO_and_RSD_measurements}. In section~\ref{sec:cosmological_implications}, we present the cosmological results from combining the measurements of the Fourier-space wedges with complementary data sets and infer cosmological \changed{constraints for different parameter spaces}. Finally, in section~\ref{sec:conclusions} we conclude our analysis with a summary and discussion of the results.

\label{sec:conclusions} \changed{In this work, we performed a cosmological analysis of the full shape of anisotropic clustering measurements in Fourier space, of the final galaxy samples from BOSS, the DR12 combined sample \citep{Reid:2015gra}, a galaxy catalogue that is unprecedented in its volume. This information can be used to place tight constraints on the expansion history of the Universe and he growth-rate of cosmic structures.} We extended the concept of clustering wedges \citep{Kazin:2011xt} to Fourier space by defining an estimator for this quantity analogous to the Yamamoto-Blake estimator \changed{for the power spectrum multipoles} \citep{Yamamoto:2005dz,Blake:2011rj}. We revised the definitions of the shot noise and optimal-variance weights of the power spectrum estimator to fully account for the observational systematics of BOSS. However, in order to make use of FFT-based estimators \citep{Bianchi:2015oia, Scoccimarro:2015bla}, we approximate the power spectrum wedges of the BOSS sample by filtering out the information of Legendre multipoles $\ell>4$. We obtain the estimate for the covariance matrices associated with our clustering measurements from the \MDPatchy \citep{Kitaura:2015uqa} and \QPM mock catalogues, which were specifically designed to mimic the clustering and observational systematics of the BOSS combined sample. Our modelling of the anisotropic power spectrum relies on novel approaches to describe non-linearities, galaxy bias, and RSD. The full model was validated using synthetic galaxy catalogues obtained from a set of 100 full \Nbody simulations using the theoretical recipe of the covariance matrix of the power spectrum wedges of \citet{Grieb:2015bia}. Further model performance tests were conducted as part of the BOSS RSD `challenge' and using the \MDPatchy mocks that mimic the entire combined sample. These tests show that any systematic biases in the distance and growth measurements introduced by our analysis method are smaller than the statistical errors obtained from the DR12 sample and can be neglected. The BAO distance and the growth rate measurements inferred from our BAO+RSD fits of the Fourier space wedges are in excellent agreement with the configuration-space results of \citet{Sanchez:2016b}, which are based on the same gRPT+RSD model, and are consistent with previous measurements on the BOSS LOWZ and CMASS samples. However, thanks to the optimization of the analysis and the improved modelling, our constraints are significantly more precise than the results obtained from previous analyses. The BAO and RSD \changed{measurements} inferred from BOSS are in good agreement with the \LambdaCDM predictions from the \Planck data at the 1-sigma level. \changed{The results presented here and those of all companion papers in the series analysing the BOSS DR12 combined sample are combined into the final consensus constraints in \citet{Alam:2016hwk}, which are computed using the methodology described in \citet{Sanchez:2016a}.} We also explored the cosmological implications of our clustering measurements by directly comparing them with the predictions obtained for different cosmological models. We combined the information in the full-shape of the clustering wedges with CMB data from the \Planck satellite and the JLA SN sample to infer constraints on the parameters of the standard \LambdaCDM cosmological model and a number of its most important extensions such as modified DE models, non-flat universes, neutrino masses and possible deviations from the predictions of GR. Assuming a \LambdaCDM cosmology, the combined data sets constrain the matter density parameter to $\Om = 0.311_{-0.010}^{+0.009}$ and the Hubble constant to $H_0 = 67.6_{-0.6}^{+0.7} \Unit{km \, s^{-1} \, Mpc^{-1}}$. These values are in good agreement with the results from the \Planck 2013 + DR11 BAO BAO + SN constraints found in \citet{Anderson:2013zyy}. Relaxing the assumption of a cosmological constant and allowing for a \changed{constant} EOS with $w \neq -1$, we find $w = 1.019_{-0.039}^{+0.048}$. In all tested DE models, the \LambdaCDM case is always found to be very well within the $1 \sigma$ confidence intervals. The most extreme case are the constraints using a $w$CDM model and a free $\sum m_\nu$, in which case we find $w = -1$ close to the edge of the 1-sigma interval. Allowing for a modification in the growth rate by varying the exponent $\gamma$ in $f = [\Om(z)]^\gamma$, we measure $\gamma = 0.52 \pm 0.10$ in perfect agreement with GR ($\gamma_\mathrm{GR} = 0.55$) and with an uncertainty reduced by a factor of 1.5 compared to the previous results of \citet{Sanchez:2013tga}. The curvature parameter $\OK$ is found to be completely consistent with zero in the tested cases. Using the \Planck + BOSS measurements for a $K$-\LambdaCDM model, the total density of the Universe today is only allowed to deviate less than $0.3\%$ from the critical density at 68\% CL. The neutrino mass is found to be $\sum m_\nu < 0.260 \Unit{eV} $ (95\% CL), which is consistent with other recent cosmological analyses such as weak lensing based on CFHTLenS \citep[$\sum m_\nu < 0.28 \Unit{eV}$ at 68 per cent CL]{Kitching:2016hvn}. We conclude that \LambdaCDM is the preferred cosmological model among the variations explored in this work and the standard paradigm has thus been further consolidated. Our analysis methodology can easily be applied to the data from other galaxy samples. In the near future, surveys such as the Hobby Eberly Telescope Dark Energy Experiment \citep[HETDEX;][]{Hill:2008mv}, the Dark Energy Spectroscopic Instrument \citep[DESI;][]{Levi:2013gra}, the Subaru Prime Focus Spectrograph \citep[PFS;][]{Ellis:2012rn} and the ESA space mission \emph{Euclid} \citep{Laureijs:2011gra} will provide even more detailed views of the LSS of the Universe, helping to improve our knowledge of the basic cosmological parameters and to further test for possible deviations from the standard \LambdaCDM model.

1607.03143

1607

1607.07140_arXiv.txt

We explore the effects of specific assumptions in the subgrid models of star formation and stellar and AGN feedback on intrinsic alignments of galaxies in cosmological simulations of ``MassiveBlack-II'' family. Using smaller volume simulations, we explored the parameter space of the subgrid star formation and feedback model and found remarkable robustness of the observable statistical measures to the details of subgrid physics. The one observational probe most sensitive to modeling details is the distribution of misalignment angles. We hypothesize that the amount of angular momentum carried away by the galactic wind is the primary physical quantity that controls the orientation of the stellar distribution. Our results are also consistent with a similar study by the EAGLE simulation team.

\label{S:intro} The intrinsic shapes and orientations of galaxies are correlated with each other and the large scale density field. This intrinsic alignment of galaxies is an important astrophysical systematic in weak lensing measurements \citep{{2000MNRAS.319..649H},{2000ApJ...545..561C},{2001MNRAS.320L...7C},{2002MNRAS.335L..89J},{2004PhRvD..70f3526H}} of upcoming surveys such as the Large Synoptic Survey Telescope\footnote{\url{http://www.lsst.org/lsst/}} (LSST; \citealt{LSST09}) and Euclid \footnote{\url{http://sci.esa.int/euclid/}, \url{http://www.euclid-ec.org}} \citep{LAA+11}. Ignoring intrinsic alignments in weak lensing analysis can significantly bias the constraints on cosmological parameters such as the dark energy equation of state parameter \citep{2016MNRAS.456..207K}. Therefore, intrinsic alignments have been studied with analytical models and also cosmological simulations including $N$-body and hydrodynamic simulations which can help in mitigating this contaminant signal. Analytically, intrinsic alignments have been modeled with a linear alignment model \citep{{2001MNRAS.320L...7C},{2004PhRvD..70f3526H}} and modifications of the model which includes the non-linear evolution of the density field \citep{{2007NJPh....9..444B},{2015JCAP...08..015B}}. However, it is difficult to analytically describe the alignments of a galaxy's stellar component by accurately considering the physics of galaxy formation. There are also limitations to the use of $N$-body simulations as one has to populate halos with galaxies by assigning a random orientation \citep{2006MNRAS.371..750H} or employ semi-analytic methods \citep{2013MNRAS.436..819J}. Recently, intrinsic alignments of galaxies in large volume hydrodynamic simulations have been extensively studied with simulations of galaxy formation such as MassiveBlack-II \citep{2015MNRAS.450.1349K}, Horizon-AGN \citep{2014MNRAS.444.1453D}, EAGLE \citep{2015MNRAS.446..521S} and Illustris \citep{{2014Natur.509..177V},{2014MNRAS.444.1518V},{2014MNRAS.445..175G}}. Cosmological hydrodynamic simulations of galaxy formation are an important tool to study intrinsic alignments as it is directly possible to measure the shape and orientation of the stellar component of galaxies in the simulations. In a precursor of this paper, \cite{2015MNRAS.448.3522T} studied the galaxy shapes and two-point statistics in the MassiveBlack-II cosmological hydrodynamic simulation. This study was extended to compare the galaxy alignments based on their morphological type in MassiveBlack-II and Illustris simulations \citep{2015arXiv151007024T}. \cite{2015MNRAS.454.2736C} used the Horizon-AGN simulation, an Adaptive Mesh Refinement (AMR) based hydrodynamic simulation of galaxy formation to study intrinsic alignments of spirals and elliptical galaxies. The redshift and luminosity evolution of alignments in the same simulation was studied in \cite{2016arXiv160208373C}. Recently, \cite{2016arXiv160603216H} studied the mass and redshift dependence of intrinsic alignments in the Illustris simulation and their dependence on stellar mass, luminosity, redshift and photometric type. Qualitatively, the properties of galaxy shapes and alignments have a similar trend with mass across different simulations. However, differences have been noted in the amplitude of galaxy alignments and morphological fraction of disk galaxies in MassiveBlack-II and Illustris \citep{2015arXiv151007024T}, as well as qualitative differences in the comparison of alignments of spirals with the over-density and the redshift dependence of intrinsic alignments in the Horizon-AGN simulation \citep{{2015MNRAS.454.2736C},{2016arXiv160208373C}}. Given the differences in the models of subgrid physics adopted in these simulations and also the numerical implementations of hydrodynamics, it is important to understand the details of the subgrid physics responsible for changes in the galaxy alignments and to explore the robustness of simulation results. In a previous study, \cite{2015MNRAS.453..721V} studied intrinsic alignments using the EAGLE suite of simulations with variations in the strength of feedback. Here, we undertake a parameter space study of the subgrid model adopted in the MassiveBlack-II simulation using a suite of small volume simulations with box size of $25h^{-1}Mpc$ on a side. We vary the free parameters in the feedback models of the simulation and test the robustness of the galaxy shapes, orientations and two-point statistics of shape correlations to variations in these parameters. Since high resolution hydrodynamic simulations of large volume are computationally expensive, we also test the usefulness of using small volume simulations to capture the sensitivity of intrinsic alignment statistics to variations in the feedback parameters. This paper is organized as follows. In Section~\ref{S:simulations}, we describe the simulations used in this study along with a brief overview of the feedback models adopted in the MassiveBlack-II simulation. Section~\ref{S:methods} provides the details of the methods adopted to calculate shapes and intrinsic alignment statistics studied in this paper. In Section~\ref{S:dcmode} we compare the results from the suite of small volume simulations with the fiducial MBII model and different amplitudes of the DC mode with those of the original $100h^{-1}Mpc$ box size MBII simulation. The intrinsic alignment statistics in the small volume runs with different feedback parameters are compared with those from the fiducial model in Section~\ref{S:baryonic}. Finally, we provide a summary of our conclusions in Section~\ref{S:conclusions} \section {Simulations and Feedback Models} \label{S:simulations} In this paper, we use the MassiveBlack-II (MBII) simulation \citep{2015MNRAS.450.1349K}, a high resolution cosmological hydrodynamic simulation performed in a box of volume $(100h^{-1}Mpc)^3$, which includes galaxy formation physics as our base model. We complement MassiveBlack-II with smaller volume simulations of size $25h^{-1}Mpc$, in which we vary the key parameters of the star formation and stellar and AGN feedback model. We denote the smaller volume simulations as MBII-25. The simulations are performed with the TreePM-Smoothed Particle Hydrodynamics (SPH) code, P-Gadget, a modified version of GADGET2 \citep{2005MNRAS.361..776S}. The same version of the code has been used earlier to perform the large volume MBII simulation \citep{2015MNRAS.450.1349K}. The simulations include the wide range of physical effects thought to be crucial for properly modeling galaxy formation, such as multiphase ISM, star formation, supernova and stellar wind feedback, as well as black hole accretion and feedback. Radiative cooling and heating are included as in \cite{1996ApJS..105...19K}, along with photoheating due to an imposed ionizing UV background. Initial conditions are generated at $z=159$ and simulations are evolved to $z=0$ with an equal initial number of gas and dark matter particles. The cosmological parameters are chosen with the WMAP7 cosmology\citep{2011ApJS..192...18K}: $h=0.701$, $\Omega_{m}=0.275$, $\Omega _{b} = 0.046$, $\Omega _{\Lambda} = 0.725$, $\sigma _{8} =0.816$, spectral index, $\eta _{s} = 0.968$ The mass of each dark matter particle is $1.1\times 10^{7}h^{-1}M_{\odot}$. The smaller volume simulations are performed with the same mass and spatial resolution as the original simulation. Accordingly, the initial number of gas and dark matter particles are equal to $2 \times 1792^{3}$ and $2 \times 448^{3}$ in the $100h^{-1}Mpc$ and $25h^{-1}Mpc$ box size simulations respectively. We note that all the small volume simulations have been started with the same initial conditions at $z=159$. The details of the star formation and feedback models of the simulation and the changes adopted in the small volume runs are described below. \subsection{Star formation and Stellar and AGN Feedback } \label{SS:sf} The star formation and feedback model adopted in the simulation is based on an earlier multiphase ISM model of Springel \& Hernquist \cite{2003MNRAS.339..289S}. Specifically, if the local gas density $\rho$ is greater than a critical density threshold $\rho _{th}$, a multiphase ISM consisting of cold clouds in pressure equilibrium with a hot ambient gas is assumed. The effective pressure $P_{eff}$ is defined as $P_{eff} = (\gamma - 1)(\rho _{h} \mu _{h} + \rho _{c} \mu _{c})$ \citep{2003MNRAS.339..289S}, where $\rho _{c}$, $\rho _{h}$ are the local densities of cold and hot phases respectively, $\rho = \rho _{c} + \rho _{h}$, and $\mu _{h}$ and $\mu _{c}$ are specific energies of hot and cold components. The threshold density $\rho _{th}$ is determined self consistently by requiring that the effective pressure is a continuous function of density. Star formation is modeled by spawning individual stellar particles stochastically from the cold clouds. The rate of star formation is given by \begin{equation} \label{eq:starfor} \frac{d\rho _{*}}{dt} = \frac{\rho _{c}}{t_{*}} - \beta \frac{\rho _{c}}{t_{*}} \end{equation} where $\beta = 0.1$ is the mass fraction of short lived stars and $t_{*}$ is the star formation time scale with density dependence given by \begin{equation} t_{*}(\rho) = t_{0}^{*}(\frac{\rho}{\rho_{th}})^{-0.5}, \end{equation} where $t_{0}^{*} = 2.1\mathrm{Gyr}$. The energy released by supernovae heats the ambient gas and the heating rate is set by the energy balance condition \begin{equation} \label{eq:sf_hrate} \frac{d}{dt}(\rho _{h} \mu _{h}) = \beta \frac{\rho _{c}}{t_{*}}(\mu _{SN}). \end{equation} Here $\mu _{SN}=\frac{3}{2}kT_{SN}$ where $T_{SN}$ is the equivalent supernova temperature which is equal to $10^{8}$ K in the fiducial model. \subsection{Wind Feedback} \label{SS:wf} Galactic winds are implemented with the wind velocity given by \begin{equation} v_{w} = \sqrt{\frac{2\beta \chi \mu _{SN}}{\eta (1-\beta)}}, \end{equation} where $\chi = 1.0$ is the fraction of supernova energy carried by the wind and $\eta = 2.0$ is the wind loading factor. For a given time step $\Delta t$, a gas particle is added to the wind probabilistically with the probability \begin{equation} p_{w} = 1 - \exp{[-\frac{\eta (1-\beta)x\Delta t}{t_{*}}]}. \end{equation} \subsection{AGN Feedback} \label{SS:agn} The simulations also include the physics of black hole accretion and feedback, based on the models of \cite{2005MNRAS.361..776S} and \cite{2005Natur.433..604D}. Black holes are treated as collisionless particles introduced into halos of mass greater than $5.0\times10^{10}h^{-1}M_{\odot}$ at regular time intervals, separated by $\Delta \log (a) = \log (1.25)$. The densest particle is converted into a seed black hole of mass $M_{\mathrm{BH,seed}} = 5\times10^{5}h^{-1}M_{\odot}$ which grows in mass by black hole accretion and mergers. The black hole accretion rate is given by the modified Bondi rate formula \begin{equation} \dot{M}_{BH} = \frac{4\pi \alpha G^{2}M^{2}_{\mathrm{BH}}\rho}{(c^{2}_{s} + v^{2}_{\mathrm{BH}})^{3/2}}, \end{equation} where $\rho$ is the local gas density, $c_{s}$ is the local speed of sound, $v$ is the velocity of BH relative to the gas. The accretion rate is limited to $2$ times the Eddington rate, $\dot{M}_{\mathrm{Edd}}$. A dimensionless parameter $\alpha$ is set to 100; that value has been found experimentally to approximately correct for the gas density close to the black hole, which is reduced in the effective sub-resolution model of the ISM. The AGN feedback is modeled by coupling $5\%$ (the value chosen to match the slope in the observed $M_{\mathrm{BH}}-\sigma$ relation \citep{2005MNRAS.361..776S}) of the bolometric luminosity radiated from the BH, \begin{equation} L_{\mathrm{bol}}= \epsilon _{r}\dot{M}_{\mathrm{BH}}c^{2}, \end{equation} with the radiation efficiency $\epsilon _{r}=0.1$. The energy is deposited isotropically to the 64 nearest gas particles within the BH particle kernel. \subsection{Parameters Space Study} \label{SS:feedbackpar} In the simulations analyzed here, we vary the key parameters in the star formation and stellar and AGN feedback models. In particular, we consider the effect of a lower or higher star formation efficiency by increasing and decreasing the star formation timescale $t_{0}^{*}$ by a factor of 3. We also consider the effects of increasing the AGN feedback by increasing the scaling parameter $\alpha$ in the AGN feedback model to $300$, which triples the black hole accretion rate. Similarly, the effect of wind velocity is weakened by decreasing the wind loading factor 10 times to study the effects of wind feedback.

\label{S:conclusions} Our primary goal in this paper is to explore the effects of model parameters in the star formation and feedback models on the galaxy shapes and alignments using small volume simulations of size $25 h^{-1}Mpc$ on a side. As our fiducial model for the simulation, we adopted the same star formation and feedback model as in the MassiveBlack-II hydrodynamic simulation of galaxy formation \citep{2015MNRAS.450.1349K}, which is performed in a box of volume $(100h^{-1}Mpc)^3$. Simulations with significantly (by factors of of 3 - 10) varying feedback show remarkable consistency with the fiducial run. Within the statistical precision we are able to achieve in our small volume runs, most of observational probes are insensitive to the details of subgrid physical modeling, with the exception of misalignment angles. We hypothesize that the angular momentum ejected by galactic winds is the most crucial physical quantity that determines the alignment of stellar shapes, and it remains one of the least robust quantities predicted in modern simulations of galaxy formation. Our conclusions are also in good agreement with similar exploration of the role of subgrid physics on intrinsic alignments by the EAGLE simulation team.

1607.07140

1607

1607.02881_arXiv.txt

We study the X-ray and optical properties of the ultraluminous X-ray source (ULX) X-6 in the nearby galaxy NGC 4258 (M106) based on the archival {\it XMM-Newton}, {\it Chandra}, {\it Swift}, and {\it Hubble Space Telescope} ({\it HST}) observations. The source has a peak luminosity of $L_{\mathrm{X}} \sim 2 \times 10^{39}$ erg s$^{-1}$ in the {\it XMM-Newton} observation of 2004 June. Consideration of the hardness ratios and spectral model parameters shows that the source seems to exhibit possible spectral variations throughout the X-ray observations. In the images from the {\it HST}/Advanced Camera for Surveys (ACS), three optical sources have been identified as counterpart candidates within the 1$\sigma$ error radius of 0$\arcsec$.3. The brightest one has an absolute magnitude of $M_{V} \approx$ $-$7.0 and shows extended structure. The remaining two sources have absolute magnitudes of $M_{V} \approx$ $-$5.8 and $-$5.3 mag. The possible spectral types of the candidates from brightest to dimmest were determined as B6$-$A5, B0$-$A7, and B2$-$A3, respectively. The counterparts of the X-ray source possibly belong to a young star cluster. Neither the standard disk model nor the slim disk model provides firm evidence to determine the spectral characteristics of ULX X-6. We argue that the mass of the compact object lies in the range $10-15M_{\sun}$ indicating that the compact source is most likely a stellar-mass black hole.

Ultraluminous X-ray sources (ULXs) are extragalactic off-nuclear point-like sources with luminosities exceeding the Eddington limit for a 10 $M_{\sun}$ black hole (BH) ($L_{\mathrm{X}}$ $\textgreater$ $10^{39}$ erg s$^{-1}$) \citep{fen11}. If the emission is isotropic, the compact objects in some of the bright ULXs might be intermediate-mass BHs with masses $\sim$ 10$^{2}-$10$^{4}$ $M_{\sun}$ \citep{mil04}. Conversely, some ULXs might contain stellar-mass BHs and their high luminosities may arise from supercritical accretion \citep{sha73,pou07}. Recent studies on ULXs showed that stellar-mass BH scenarios are reliable \citep{liu13,mot13,fab15}. On the other hand, pulsations with an average period of 1.37 s were detected from a ULX in M82 using {\it NuSTAR} data, which indicates that the compact object in this system is a neutron star \citep{bac14}. That result has led to the idea that some ULX systems may harbour neutron stars instead of BHs. The nature of the ULX binary systems is still unclear. Studying the X-ray spectral states and state transitions of ULXs with the help of available multi-epoch data and comparing them with the well-known characteristics of Galactic BH binaries (BHB) are essential tools for understanding the radiative mechanisms of these sources. There are three active states that have been defined for Galactic BHBs: thermal, hard, and steep power law (PL). In the thermal state, a geometrically thin, optically thick accretion disk dominates the emission, while in the hard state the emission is produced by a geometrically thick, optically thin Comptonizing region. The hard state is characterized by non-thermal PL emission with a photon index $1.4<\Gamma<2.1$. However, the steep PL state is defined by a softer spectrum having a photon index of $\Gamma>2.4$ \citep{rem06}. In the steep PL or thermal state, most of the Galactic BHBs have relatively higher luminosities than in the hard state. A similar correlation between luminosity and photon index has been found in ULXs X-1 in NGC 1313 \citep{fen06,dew10} and X37.8+54 in M82 \citep{jin10}, although there are some ULXs that exhibit the opposite behavior (NGC 1313 X-2, \citealt{fen06}; NGC 4736 X-2, \citealt{avd14}). Additionally, distinct spectral state transitions have been observed in some ULXs (e.g. NGC 2403 src 3, \citealt{iso09}; IC 342 X-1, \citealt{mar14}). On the other hand, identification of the optical counterparts of the ULXs may provide valuable information. The optical emission could originate from the donor star and/or the accretion disk via X-ray photoionization \citep{fen11}. The optical counterparts of the several ULXs have been found in nearby galaxies using {\it Hubble Space Telescope} ({\it HST}) data (\citealt{tao11} and references therein; \citealt{gla13}). Broadband {\it HST} photometry of the optical counterparts allows constraints to be placed on the mass and spectral type of the companion star \citep{yan11,gri11,gri12}. These constraints could also be defined by studying the environment of the ULX if the system belongs to a stellar cluster or association \citep{gri11,pou13}. In this work, the X-ray spectral properties of the ULX X-6 \footnote{We adopted the source number from the work of Akyuz et al. (2013). They numbered the detected sources in NGC 4258 as XMM-$n$, where $n$ represented the source number with decreasing EPIC pn count rate. We have shortened their designation to X-6 for convenience.} in NGC 4258 have been studied using archival {\it XMM-Newton}, {\it Chandra} and {\it Swift} observations. Also the optical counterpart of X-6 has been searched on the {\it HST}/Advanced Camera for Surveys (ACS)/WFC archival images. NGC 4258 (M106) is a nearby (7.7 Mpc, \citealt{swa11}) Seyfert-type spiral galaxy. It is well known for its anomalous arms, discovered on the basis of H$\alpha$ imaging \citep{wil01}. The X-6 is located 2$\arcmin$.4 away from the center of the galaxy and its {\it Chandra} coordinate is R.A. $=$ 12$^{\mathrm{h}}$18$^{\mathrm{m}}$43$^{\mathrm{s}}.887$, decl. $=$ +47$^{\circ}$17$\arcmin$31$\arcsec.81$. The source was classified as a ULX by \citet{swa11} with an unabsorbed X-ray luminosity of 1.6$\times 10^{39}$ erg s$^{-1}$ in the 0.3$-$10 keV energy band. X-6 is not positionally coincident with any X-ray point source in the {\it Einstein} and {\it ROSAT} catalogs. \citet{aky13} also studied the X-ray spectrum and the temporal properties of this source. They presented spectral and timing analyses based on the {\it XMM-Newton} observations with the longest exposure available for the non-nuclear X-ray point sources in the $D_{25}$ of NGC 4258. The paper is organized as follows: the observations and data reductions are described in Section 2. The details and results of the analyses are given in Section 3. Discussion of the physical properties of the ULX and a summary are given in Section 4.

We examined the X-ray temporal and spectral properties of X-6 using seven {\it XMM-Newton}, one {\it Chandra}, and 12 {\it Swift} observations available in the archives. Also, the optical counterpart of X-6 was investigated using the archival {\it HST}/ACS/WFC data. With the help of the simultaneous multi-band {\it HST}/ACS/WFC data, the CMDs for optical counterpart candidates and the cluster members have been obtained. As seen in Figure 3(a) and 3(c), X-6 exhibits possible spectral variations. However, the light curves from the {\it Swift} data do not show significant variation and indicate a rather persistent behaviour (see Figure 3(b)). Most notably, X-6 has the lowest HR value in XM5 data. The source has a steeper PL photon index ($\Gamma \sim 2.40$) and the lowest inner disk temperature ($T_{\mathrm{in}} \sim 0.83$ keV) at this epoch. If we assume that the source emits at the Eddington limit in XM6 data (in which the spectrum gives the highest $L_{\mathrm{X}} = 2.2 \times 10^{39}$ erg s$^{-1}$) the mass of the compact object in this system is found to be $M_{\mathrm{BH}}\sim 15 M_{\sun}$. The luminosity of X-6 changes by a factor $\sim$2 during these variations. Nonetheless, this variation in luminosity seems not to correlate with the canonical BHB states. Generally in Galactic BHBs, the luminosities are usually lower in the hard state, higher in the thermal (soft) state, and switch to a very high value in the steep PL state. However, there are some Galactic BHBs and ULXs that do not show similar behavior (e.g. XTE J1550-564, \citealt{rem06}; NGC 1313 X-2, \citealt{fen06}; IC 342 X-1, \citealt{mar14}; NGC 4736 X-2, \citealt{avd14}). Taking into account that the ULXs are different than usual Galactic BHBs, each might harbour a stellar-mass BH with a supercritical accretion disk (SCAD) \citep{fab15} similar to GRS 1915+105 \citep{vie10}. In SCAD scenario the disk is expected to be slim ($H/R \sim 1$) within a spherization radius ($r_{\mathrm{sp}}$) and the temperature is expected to depend on the radius as $R^{1/2}$ \citep{pou07}. We tried to fit the X-ray spectra of X-6 with a $p$-free model in {\scshape xspec} to determine whether X-6 has slim disk properties \citep{abr88,wat01}. The DISKPBB model yields acceptable fits to the spectra of the source in C1, XM3, and XM6 data (see Table 4). The $p$ parameters were found to be consistent with the slim disk model with one exception (C1 data, $p \sim$ 0.75). This may be due to the decrease in the flux of X-6 in C1 data. The inner disk temperatures obtained from DISKPBB fits are consistent with the model \citep{pou07} and similar to some other ULXs (e.g. \citealt{vie06}; \citealt{gla09}; \citealt{sor15}). It is possible to constrain the mass of the compact object using the parameters of the disk models (DISKBB and DISKPBB). For this calculation, we used the technique described by \citet{sor15}. Since the spectrum of X-6 is relatively better modeled with DISKBB model in XM3 data, the DISKBB normalization parameter obtained using that observation was adopted for this calculation. We found an apparent inner disk radius of $r_{\mathrm{in}}\sqrt{\cos i} \approx 66$ km. The apparent radius was corrected to the true value using the equation $R_{\mathrm{in}}=\xi \cdot \kappa^{2} \cdot r_{\mathrm{in}}$ equation, where the correction factor $\xi = 0.412$ and $\kappa$ is a spectral hardening factor (see \citealt{kub98}). Assuming $\kappa = 1.7$ \citep{shi95} and disk inclination $i = 60$, the true inner disk radius was calculated as $R_{\mathrm{in}} \approx 100$ km. Using the relation between inner disk radius and mass \citep{mak00}, we found a BH mass of $M \sim 10M_{\sun}$ for a non-spinning BH. Also, if we consider the normalization parameter of the DISKPBB model in XM3 data, we may calculate another mass value for the compact object. We derived a true inner disk radius of $R_{\mathrm{in}}\sqrt{\cos i} \approx 40$ km using the correction factor $\xi=0.353$ and a spectral hardening factor $\kappa =$ 3 \citep{vie08}. Assuming a moderate disk inclination $i=60$ and taking the mass correction factor as minimum ($\sim$ 1.2) the mass of the compact object in X-6 can be calculated as $M \sim 10M_{\sun}$ for a non-spining BH. This value is consistent with the estimation in the paragraph above. Three optical counterpart candidates were identified after the astrometric correction. We calculated the X-ray to optical flux ratios for the three sources. This ratio is given as log($f_{\mathrm{X}}$/$f_{V}$)$=$log$f_{\mathrm{X}}$ $+$ $m_{V}$/2.5 $+$ 5.37 where $m_{V}$ is the extinction-corrected visual magnitude and $f_{\mathrm{X}}$ is the unabsorbed X-ray flux in the energy band 0.3$-$3.5 keV \citep{mac82}. Simultaneous X-ray and optical observations are not available for X-6. Therefore we calculated log($f_{\mathrm{X}}$/$f_{V}$) using the minimum (XM5) and maximum (XM6) $f_{\mathrm{X}}$ values adopted from PL model parameters. For source 2, log($f_{\mathrm{X}}$/$f_{V}$) was found to be 1.5$-$1.8. Although these values are within the given ratios for active galactic nuclei (AGNs) ($-$1 to 1.7, \citealt{mac88}), they are acceptable values also for a high-mass X-ray binary. For source 1 and 3, the ratios were found to be 2.2$-$2.5 and 2.0$-$2.3, respectively. These ratios are greater than the values for AGNs, normal stars, normal galaxies, and BL Lac objects \citep{mac88,sto91} and are similar to those for the optical counterparts of other ULXs \citep{fen08,yan11,tao11,avd16}. Colors and absolute magnitudes of these sources were obtained (see Table 5). These sources have $M_{V}$ values that are consistent with the optical counterparts of ULXs in Table 4 of \cite{tao11}, which lie in the range $-$7 $<$ $M_{V}$ $<$ $-$3. If we assume that the optical emission of X-6 is dominated by the companion star and we use the Schmidt--Kaler table (\citealt{all82}) of intrinsic colors, then the probable spectral types of sources 1, 2, and 3 can be estimated to be B2$-$A3, B6$-$A5, and B0$-$A7 supergiants, respectively. We also found, by using the {\it HST}/ACS/WFC images, that the counterpart candidates of X-6 possibly belong to a star cluster. After obtaining CMDs for the stars in the cluster and using the Padova isochrones, the ages of sources 1, 2, and 3 have been estimated to be in the ranges 15$-$30, 12$-$14, and 6$-$25 MYr, respectively. The masses of the counterpart candidates are estimated from the PARSEC isochrones by taking into account their ages and absolute magnitudes as 9$-$13 M$_{\sun}$ for source 1, 14$-$16 M$_{\sun}$ for source 2, and 10$-$25 M$_{\sun}$ for source 3. The mass ranges of the candidates are compatible with the donor stars of other ULXs \citep{pat08}. On the other hand, the optical emission could also arise from the accretion disk. When the disk is dominant, this emission is expected to vary significantly, as is found for most ULXs, e.g., M101 ULX-1 (\citealt{tao11}). If the optical variability of the counterpart candidate is taken into account, source 1 is a promising candidate in the present case and the contribution from the disk cannot be ignored. Therefore, the donor star and the accretion disk may give comparable contributions to the optical emission of X-6. In summary, even though the source does not exhibit variations in its X-ray spectrum similar to the Galactic BHBs, there are some spectral changes in both HR and spectral model parameters. Nonetheless, it is hard to say anything conclusively about the accretion regime. Additional X-ray data with better statistical quality can constrain the physical parameters better, shedding more light on the origin of the X-ray emission from X-6. Both broadband photometric and high-resolution spectroscopic observations will help to distinguish the optical counterpart and find the origin of the optical emission.

1607.02881

1607

1607.05373_arXiv.txt

The recent direct detection of gravitational waves (GWs) from binary black hole mergers~\cite{Abbott:2016blz,Abbott:2016nmj} opens up an entirely new non-electromagnetic window into the Universe making it possible to probe physics that has been hidden or dark to electromagnetic observations. In addition to cataclysmic events involving black holes, GWs can be triggered by physical processes and systems involving neutron stars. Properties of neutron stars are largely determined by the equation of state (EOS) of neutron-rich matter, which is the major ingredient in calculating the stellar structure and properties of related phenomena, such as gravitational wave emission from elliptically deformed pulsars and neutron star binaries. Although the EOS of neutron-rich matter is still rather uncertain mainly due to the poorly known density dependence of nuclear symmetry energy at high densities, significant progress has been made recently in constraining the symmetry energy using data from terrestrial nuclear laboratories. These constraints could provide useful information on the limits of GWs expected from neutron stars. Here after briefly reviewing our previous work on constraining gravitational radiation from elliptically deformed pulsars with terrestrial nuclear laboratory data in light of the recent gravitational wave detection, we estimate the maximum gravitational wave strain amplitude, using an optimistic value for the breaking strain of the neutron star crust, for 15 pulsars at distances 0.16 kpc to 0.91 kpc from Earth, and find it to be in the range of $\sim[0.2-31.1]\times 10^{-24}$, depending on the details of the EOS used to compute the neutron star properties. Implications are discussed.

Gravitational waves are tiny disturbances in space-time predicted by General Relativity in 1916. A century after the fundamental predictions of Albert Einstein, the Laser Interferometer Gravitational-Wave Observatory (LIGO) has detected directly gravitational waves from black hole mergers~\cite{Abbott:2016blz,Abbott:2016nmj}. This detection has the potential to transform profoundly our understanding of the Universe, as gravitational wave astronomy opens up the possibility to probe physics that has been hidden to current electromagnetic observations~\cite{Maggiore:2007}. Because gravity interacts extremely weakly with matter, gravitational waves carry much cleaner and detailed picture of their sources as opposed to their electromagnetic counterparts~\cite{Flanagan:2005yc}. Several types of gravitational waves and their corresponding sources have been discussed in the literature: (i) {\it Inspiral gravitational waves} are triggered during the latest stages of compact binary systems where the two objects finally collide and merge. These systems typically consist of two neutron stars, two black holes, as demonstrated by LIGO~\cite{Abbott:2016blz,Abbott:2016nmj}, or a neutron star and a black hole whose orbit decayed to the point that the two masses are about to coalesce. This requires one of the original stars to be massive enough to undergo collapse to a compact object without destroying its companion, and without disrupting the bound orbit~\cite{Riles:2012yw}. (ii) {\it Burst gravitational waves} are generated in sudden cataclysmic events, such as supernova explosions or gamma-ray bursts~\cite{Riles:2012yw}. A spherically symmetric explosion cannot lead to emission of GWs in General Relativity. In order to produce GWs a supernova must exhibit some asymmetry. It is known that many pulsars formed in supernovae have large speeds relative to neighboring stars (high "birth kicks"), and this suggests strongly that certain spernovae do exhibit considerable non-spherical motion~\cite{Ott:2006qp}. (iii) {\it Stochastic gravitational waves} could even have been produced during the very early Universe, well before any stars had been formed, merely as a consequence of the dynamics and expansion of the Cosmos~\cite{Riles:2012yw}. Very similar to the Cosmic Micro-wave Background (CMB), which is an electromagnetic leftover from the Big Bang, these GWs originate from many independent events amounting to a Cosmic Gravitational-wave Background (CGB). Such gravitational waves may carry critical information about the very early Universe, as they would have been stretched out as the Cosmos expanded. (iv) {\it Continuous gravitational waves} are generated by sources with steady and well-defined frequency. Examples of such sources are binary neutron star or black hole systems orbiting the common center of mass, or a (rapidly) spinning neutron star with some long-living axial asymmetry~\cite{Jaranowski:1998qm}. Ground-based gravitational wave observatories, such as LIGO and VIRGO, will also allow studies of a large population of neutron stars. Additionally, in space eLISA(Evolved Laser Interferometric Space Antenna), a planned space mission by the European Space Agency expected to launch in the early 2020s, will provide an unprecedented instrument for GWs search and detection~\cite{AmaroSeoane:2012je}. The eLISA will survey the low-frequency gravitational-wave bandwidth and is expected to detect a wide variety of events and systems in Space. Rotating neutron stars are considered to be one of the main prospects for sources of continuous GWs potentially detectable by the LIGO~\cite{Abbott:2004ig} and VIRGO (e.g. Ref.~\cite{Acernese:2007zzb}) laser interferometric observatories. According to General Relativity, a perfectly symmetric rotating object self-bound by gravity does not generate GWs. To emit GWs over an extended time period, a pulsar needs to exhibit a long-living axial asymmetry, e.g., a "mountain" on its surface~\cite{Jaranowski:1998qm}. In the literature, several different mechanisms causing such asymmetries have been discussed: \begin{enumerate}[label=(\roman*)] \item{Anisotropic stress built up during the crystallization period of the neutron star crust may be able to support long term asymmetries, such as static "mountains" on the neutron star surface~\cite{PPS:1976ApJ}.} \item{In addition, because of its violent formation in supernova the rotational axis of a neutron star may not necessarily be aligned with its principal moment of inertia axis, which results in a neutron star precession~\cite{ZS:1979PRD}. Because of this, even if the pulsar remains symmetric with respect to its rotational axis, it generates GWs~\cite{ZS:1979PRD,Z:1980PRD}.} \item{Also, neutron stars have extremely strong magnetic fields, which could create magnetic pressure and, in turn, deform the pulsar, if the magnetic and rotational axes do not coincide~\cite{BG:1996AA}.} \end{enumerate} These mechanisms generally cause a triaxial pulsar configuration. GWs are characterized by a very small dimensionless strain amplitude, $h_0$. The magnitude of the strain amplitude depends on how much the pulsar is distorted from axial symmetry which depends on details of the EOS of neutron-rich matter. The EOS of nuclear matter under extreme pressure, density and/or isospin asymmetry is still largely uncertain and rather theoretically controversial. A major source of these uncertainties is the rather poorly known density dependence of the nuclear symmetry energy, $E_{sym}(\rho)$ which encodes the information about the energy associated with the neutron-proton asymmetry in neutron-rich matter, see, e.g.~\cite{Li:1997px,ibook01,Lattimer:2004pg,Bar05,Ste05,EPJA}. Fortunately, both nuclear structures and reactions provide useful means to probe the $E_{sym}(\rho)$ in terrestrial nuclear laboratories \cite{Bah14}. For example, several observables in heavy-ion reactions are known to be sensitive to the $E_{sym}(\rho)$ from sub-saturation to supra-saturation densities, see, e.g., Refs.~\cite{Li97,Li00,Li02,Rei16}. For a comprehensive recent review, see articles in Ref.~\cite{EPJA}. Thanks to the hard work of many people in both nuclear physics and astrophysical communities, considerable progress has been achieved in recent years in constraining the symmetry energy especially around and below the nuclear matter saturation density using data from both terrestrial experiments and astrophysical observations, see, e.g., Refs.~\cite{LCK08,Tsa12,LiHan13,Hor14,Bal16}. In this chapter, we first briefly review our earlier work on constraining the GWs expected from elliptically deformed pulsars~\cite{Krastev:2008PLB} in light of the recent direct gravitational wave detections~\cite{Abbott:2016blz,Abbott:2016nmj}. Applying several nucleonic EOSs constrained by laboratory nuclear experiments, we then estimate the GW strain amplitude for fifteen pulsars for selected neutron star configurations. Our focus is on illustrating effects of nuclear symmetry energy on the GW strain amplitude from deformed pulsars. Recent reviews on effects of the density dependence of nuclear symmetry energy on GWs from various neutron star oscillations~\cite{New14}, binary neutron star mergers~\cite{Fattoyev:2013rga}, the pure general relativistic w-mode~\cite{Wen09} as well as the EOS-gravity degeneracy~\cite{He15} can be found in the referred references.

In this work we revisited and updated our estimates on the upper limits on the gravitational wave signals expected from elliptically deformed millisecond pulsars and compared our predictions with the current strain sensitivity of LIGO and VIRGO. More specifically, we provided estimates on the strain amplitude of gravitational waves to be expected from fifteen pulsars at distances from Earth 0.16 kpc to 0.91 kpc and rotational frequencies below 300 Hz. We have applied the MDI EOS constrained by nuclear laboratory data and calculated the upper limit on the gravitational-wave signal to be expected from these pulsars. Depending on the EOS details the {\it maximum} $h_0$ is found to be in the range of $\sim[0.2-31.1]\times 10^{-24}$. Effects of the density dependence of nuclear symmetry energy are examined. In the near future, new experiments especially those using high intnsity/energy radioactive beams in terrestrial nuclear laboratories will provide more accurate information about the EOS of dense neutron-rich matter, especially the high density behavior of nuclear symmetry energy. The EOS of neutron-rich matter will be better constrained in a wide density range and most of the studies presented in this work will be refined further. The better constrained EOS will then help predict more accurately properties of pulsars and the gravitational waves emitted by them. Similarly, rapid progress in astrophysical observations of neutron stars and measurements of gravitational wave signals will help us better understand the structure of pulsars and extract more accurately information about the underlying EOS. Ultimately, combining results from terrestrial laboratories and astrophysical observations will help us reach the shared goal of thoroughly understanding the nature of dense neutron-rich nucleonic matter.

1607.05373

1607

1607.05003_arXiv.txt

For several years, we have been developing vortex phase masks based on sub-wavelength gratings, known as Annular Groove Phase Masks. Etched onto diamond substrates, these AGPMs are currently designed to be used in the thermal infrared (ranging from 3 to 13~$\mu$m). Our AGPMs were first installed on VLT/NACO and VLT/VISIR in 2012, followed by LBT/LMIRCam in 2013 and Keck/NIRC2 in 2015. In this paper, we review the development, commissioning, on-sky performance, and early scientific results of these new coronagraphic modes and report on the lessons learned. We conclude with perspectives for future developments and applications.

\label{sec:intro} % Over the past ten years, direct imaging has become one of the main tools for the detection and characterization of planetary systems, including proto-planetary disks, debris disks, and giant extrasolar planets. Direct imaging enables not only spectral characterization of the detected objects, but also to obtain direct information of planetary systems architecture at various ages (including the proto-planetary phase), which is key to our understanding of how planetary systems form and evolve. In the context of direct planet imaging, the vortex coronagraph\cite{Mawet05,Foo05} has been considered for more than ten years as one of the most promising concepts to reduce the stellar glare\cite{Guyon06}. Yet, vortex coronagraphs have only recently been installed on 10-m class telescopes. Vortex coronagraphs feature a vortex phase mask in their focal plane. The textbook effect of the vortex phase mask is to move the light of an on-axis source outside the geometric image of the input pupil. When followed by a---usually undersized---diaphragm (aka Lyot stop) in a downstream pupil plane to block the diffracted light, the vortex phase mask provides a theoretically perfect starlight cancellation for a clear, circular aperture. One of the possible implementations of the vortex phase mask is the Annular Groove Phase Mask (AGPM)\cite{Mawet05}, made up of a concentric sub-wavelength grating etched onto a diamond substrate\cite{Forsberg13,Delacroix13}. Sub-wavelength gratings produce form birefringence, which can be used to synthesize a helical phase ramp for the two orthogonal polarization states of light separately. By tuning the grating parameters, one can adjust the wavelength-dependence of the birefringent effect induced by the sub-wavelength grating to produce quasi-achromatic behavior across a wide bandwidth\cite{Mawet05b,Delacroix12}. Together with its high throughput and $360^{\circ}$ discovery space, this makes the AGPM an excellent candidate for installation on high-contrast imaging instruments. Here, we review the design, manufacturing, and testing of these AGPMs (Sect.~\ref{sec:design}). We then report on their installation and commissioning on 10-m class telescopes, including operational aspects (Sect.~\ref{sec:comm}), and present their measured on-sky performance (Sect.~\ref{sec:perfo}), as well as a few examples of early scientific results (Sect.~\ref{sec:science}). We conclude with a discussion of the lessons learned from three years of on-sky operations, and of the perspectives for the infrared vortex coronagraph in the coming years (Sect.~\ref{sec:lessons}).

1607.05003

1607

1607.00626_arXiv.txt

Significant undulations appear in the light curve of a recently discovered super-luminous supernova (SLSN) SN 2015bn after the first peak, while the underlying profile of the light curve can be well explained by a continuous energy supply from a central engine, possibly the spin-down of a millisecond magnetar. We propose that these undulations are caused by an intermittent pulsed energy supply, indicating an energetic flare activity of the central engine of the SLSN. Many post-burst flares were discovered during X-ray afterglow observations of Gamma-Ray Bursts (GRBs). We find that the SLSN flares described here approximately obey the empirical correlation between the luminosity and time scale of GRB flares, extrapolated to the relevant longer time scales of SLSN flares. This confirms the possible connection between these two different phenomena as previously suggested.

During the past ten years, a series of modern supernova surveys have discovered an unusual type of supernovae with an absolute magnitude at peak emission of $M_{\rm AB}<-21$, which are more luminous than normal supernovae by a factor of $\sim10-100$ (Benetti et al. 2014; Bersten et al. 2016; Chatzopoulos et al. 2011; Chen et al. 2016a; Chomiuk et al. 2011; Chornock et al. 2013; Gal-Yam et al. 2009, 2012; Howell et al. 2013; Inserra et al. 2013, 2016; Kangas et al. 2016; Leloudas et al. 2012; Lunnan et al. 2013, 2016; McCrum et al. 2014, 2015; Nicholl et al. 2013, 2016; Ofek et al. 2007; Papadopoulos et al. 2015; Quimby et al. 2007, 2011; Smith et al. 2007, 2016; Vreeswijk et al. 2014; Yan et al. 2016). The total radiated energy of a typical superluminous supernova (SLSN) is on the order of $\sim10^{51}$ erg. If this radiation is mainly powered as usual by the radioactive chain $\rm ^{56}Ni\rightarrow^{56}Co\rightarrow^{56}Fe$, then an extremely large amount (several to several tens of solar masses) of radioactive $^{56}$Ni would be required. In principle, such a high mass of $^{56}$Ni could be produced by core-collapse explosions of very massive progenitors with a very large explosion energy (Umeda \& Nomoto 2008; Moriya et al. 2010) or by disruption explosions of very massive progenitors due to pair-production instability (Barkat et al. 1967; Heger \& Woosley 2002; Gal-Yam et al. 2009). In both cases, however, the corresponding high masses of supernova ejecta would lead to a broad slowly-evolving supernova light curve, whereas the observational light curves rise and often decline rapidly. Moreover, for the pair-instability events, there is still controversy whether they occur locally and if they are related to some SLSNe or not (e.g. McCrum et al. 2014; Georgy et al. 2017). Therefore, generally speaking, the radioactivity power scenario is seriously challenged by the very high luminosity of SLSNe. Alternatively, a powerful central engine is believed to play an essential role in driving SLSN explosions and in powering their emission, by an instantaneous and/or a long-lasting energy injection into the explosion-ejected stellar envelope. To be specific, at the initial time of some SLSN explosions, their central engines could impulsively provide a great amount of energy to the supernova ejecta, and lead the ejecta to have a very high initial velocity corresponding to a kinetic energy on the order of $10^{52}$ erg. If these explosions happen in dense, extended circum-stellar material (CSM; e.g., stellar wind and some particular material clusters), then the ejecta can be subsequently heated by the conversion of the kinetic energy through shock interaction between the ejecta and surrounding material (Smith \& McCray 2007; Chevalier \& Irwin 2011; Moriya et al. 2011, 2013; Ginzburg \& Balberg 2012; Inserra et al. 2016). Such interaction-powered supernovae can usually be indicated by narrow Balmer emission lines (Chatzopoulos et al. 2011). On the contrary, for hydrogen-poor SLSNe, shock interaction could usually be negligible, and the supernova emission is probably powered directly by a long-lasting central engine. During the past few years, a remarkable number of light curves of SLSNe, even including several hydrogen-rich ones without narrow lines, have been successfully explained with a continuous energy injection (Dessart et al. 2012; Inserra et al. 2013, 2016; Nicholl et al. 2013, 2016; Howell et al. 2013; McCrum et al. 2014; Wang et al. 2015; Dai et al. 2016; Lunnan et al. 2016; Bersten et al. 2016; Yu et al. 2017). This indicates that the central engine could be the most viable and most common energy source for most SLSNe, which makes these SLSNe very relevant to another engine-driven phenomenon: gamma-ray bursts (GRBs). In principle, different energy sources could coexist in some SLSNe (e.g. Wang et al. 2015a; Lunnan et al. 2016). In the framework of the long-lasting energy injection model, it is convenient to connect SLSN light curves with the temporal behaviors of their central engines. Specifically, while a continuous energy injection determines the basic profile of the SLSN light curves, it can also be expected that some light curve undulations could be caused by flare activity of the central engines. In this case, the modeling of these light curve undulations could provide a special insight in probing the nature of these SLSN engines and even their possible connection with GRB engines.

Recently, it has been suggested that Type I SLSNe and long GRBs could be two branches from an united origin (Metzger et al. 2015; Yu et al. 2017). To be specific, both are produced as outcomes of the formation of a millisecond magnetar from core-collapse of a massive rapidly rotating progenitor star. The primary difference between them is the magnetic field strengths of their magnetars, which leads to various observational differences between these two explosion phenomena. The discovery of the light curve undulations of SN 2015bn, which is very fortunate, indicates that late but energetic flare activity has taken place on its central engine. This strongly reveals the similarity and connection between SLSNe and GRBs from a new perspective, in view of the ubiquitous existence of GRB flares. The possible universal $L_{\rm flare,p}-t_{\rm flare,p}$ relationship further indicates that the central engines of SLSNe and GRBs could have a common nature, e.g., the magnetar nature as supposed. This provides an independent and robust support to the suggestion of Metzger et al. (2015) and Yu et al. (2017). The different energy and time scales of the flares of SLSNe and GRBs could arise from the different magnetic fields of their magnetar engines. In the framework of the magnetar engine model, we tend to believe that the flare activity is associated with reconnections of ultra-high multi-polar magnetic fields in the magnetar, by referring to the magnetic reconnection model previously proposed by Dai et al. (2006) for explaining GRB flares. Although the dipolar field derived above for SN 2015bn is not so high, the discovery of some so-called low-field magnetars in our galaxy (e.g. Rea et al. 2010) indicates that the internal multi-polar fields of a magnetar could be much higher than its surface dipolar field. In Dai et al. (2006), a strong internal toroidal field is considered to form due to differential rotation of the magnetar and subsequently float to the stellar surface to be reconnected. Qualitatively, this scenario could have many similarities with the physical processes of solar flares. This was statistically confirmed by Wang \& Dai (2013), who discovered that the distributions of energies, durations, and waiting times of both GRB flares and solar flares can all be understood within the physical framework of self-organized-criticality avalanche-like processes. Nevertheless, in Dai et al. (2006), the magnetar is assumed to simply consist of a clearly-separated solid crust and core, which is probably invalid for an extremely hot magnetar. An actual evolution of magnetic fields of a newly-born magnetar could be much more complicated than that considered in Dai et al. (2006). For example, fluid instabilities could be involved (e.g. Cheng \& Yu 2014). Therefore, an elaborate model is demanded to quantitatively describe the magnetic reconnections and to account for the properties of the observed SLSN flares as well as the related GRB flares. An elaborate model is also needed to consolidate our flare explanation for the light curve undulations and distinguishing it from other possible scenarios (e.g., the CSM interaction model). First of all, the dynamics of the supernova ejecta and the radiation transfer in it need to be described by a more detailed radiation hydrodynamic code such as the public SuperNova Explosion Code (SNEC; Morozova et al. 2015). Furthermore, in contrast to the 1D case considered here, a 2D simulation of Chen et al. (2016) showed that the interaction between a magnetar wind and a supernova ejecta can in fact lead to many fluid instabilities, which mix the ejecta material and fracture the ejecta into filamentary structure. The consequent inhomogenity and anisotropy of the system could substantially influence the early dynamics and emission of the supernova. It can even be expected that the clumpy structure of the ejecta could also cause some light curve undulations, which is worth to be investigated in future simulations. In any case, it will be crucial and helpful to collect more observational data to exhibit the details of early light curves (in particular during the increasing phase) and to implement some synergic multi-wavelength observations to SLSNe including in the high-energy bands.

1607.00626

1607

1607.03625_arXiv.txt

We present the first detection of a correlation between the \LyA forest and cosmic microwave background (CMB) gravitational lensing. For each \LyA forest in SDSS-III/BOSS DR12, we correlate the one-dimensional power spectrum with the CMB lensing convergence on the same line of sight from Planck. This measurement constitutes a position-dependent power spectrum, or a squeezed bispectrum, and quantifies the non-linear response of the \LyA forest power spectrum to a large-scale overdensity. The signal is measured at 5~$\sigma$ and is consistent with the $\Lambda$CDM expectation. We measure the linear bias of the \LyA forest with respect to the dark matter distribution, and contrain a combination of non-linear terms including the non-linear bias. This new observable provides a consistency check for the \LyA forest as a large-scale structure probe and tests our understanding of the relation between intergalactic gas and dark matter. In the future, it could be used to test hydrodynamical simulations and calibrate the relation between the \LyA forest and dark matter.

Quasars are bright lanterns that illuminate the dark Universe and probe the distribution of gas. The \LyA forest observed in their spectra reveals the presence of intervening neutral hydrogen absorbing light at $1216\ \angstrom$. It has been used to study the thermal history of intergalactic gas and hydrogen reionization \cite{2007ApJ...662...72B,2015MNRAS.447.3402B,Lee:2015cm,2011MNRAS.415.2257M}. Assuming that the flux transmission in the \LyA forest traces the matter density makes it a powerful probe of the large-scale structure of the Universe at intermediate redshifts, on a wide range of scales. The one-dimensional power spectrum along the line of sight probes the matter fluctuations on the smallest scales, constraining the sum of neutrino masses \cite{2015JCAP...11..011P,2015JCAP...02..045P}, models of warm dark matter \cite{2013PhRvD..88d3502V,Seljak:2006hv} and primordial black hole dark matter \cite{Afshordi:2003ic}. Combining different lines of sight enables probing the three-dimensional power spectrum on larger scales \cite{2011MNRAS.415.2257M,2011JCAP...09..001S} and provides a measurement of the baryonic acoustic oscillations (BAO) at high redshift \cite{Slosar:2013bz,2013A&A...552A..96B,2015A&A...574A..59D}. The interpretation of these results rely on the \LyA transmission tracing the underlying matter density field. If the hydrogen is in photoionization equilibrium with a uniform UV background, and there is no other sources of entropy, then the relationship is described analytically through variations of the fluctuating Gunn-Peterson approximation (FGPA; \cite{Croft:1998gi}) and is evaluated numerically using hydrodynamical simulations \cite{2014JCAP...07..005B,Rossi:2014jy}. However, the connection between \LyA transmission and the underlying matter density is complex \cite{2014PhRvD..89h3010P} and non-linear. It is affected by the proximity effect on the largest scales \cite{2014PhRvD..89h3010P}, by thermal broadening, Jeans smoothing and non-linear gravitational evolution on the smallest scales \cite{2015JCAP...12..017A} and by the gas equation of state throughout. For these reasons, and in the light of the tension between BAO measurements from the \LyA forest and galaxies \cite{2015PhRvD..92l3516A}, consistency checks for the link between \LyA transmission and matter density are valuable. Gravitational lensing of the cosmic microwave background (CMB) is sourced by large-scale structures located between the last scattering surface and the observer, and provides a measurement of the projected density of matter. In this paper, we cross-correlate, for the first time, the power spectrum of the flux transmission in the \LyA forest of quasars with CMB lensing to test our understanding of the relation between intergalactic gas and dark matter. Since both the CMB lensing convergence and the mean \LyA transmission probe the mean density on a given line of sight, it is natural to compute their cross-spectrum. However, the mean \LyA transmission is strongly affected by continuum fitting in the quasar spectrum, making this a challenging observable. For this reason, we instead correlate the CMB lensing convergence with the small-scale \LyA power spectrum on the same line of sight. The origin of this signal is more complex, and corresponds to a position-dependent power spectrum \cite{2014JCAP...05..048C,2014PhRvD..90l3523S}, or a squeezed bispectrum of the matter density. Simply put, a positive CMB convergence corresponds to an overdense line of sight; on this line of sight, the matter power spectrum is enhanced on all scales, due to non-linear evolution under gravity \cite{2014JCAP...05..048C,Li:2014hm,2014PhRvD..90l3523S} (see Fig.~\ref{fig:schematic} for a schematic of this idea). This bispectrum would therefore vanish at linear order in the perturbation theory of the density field, where short and long modes are independent. Instead, for a non-linear density field, this signal probes the response of the \LyA power spectrum to a mean overdensity. \begin{figure}[t]\centering\includegraphics[width=1\columnwidth]{figures/cartoon_v2}\caption{Schematic of the correlation: overdense regions (respectively underdense regions), in red on the top panel (blue on the bottom panel) have positive (negative) CMB lensing convergence and are expected to produce more (less) small-scale structures under non-linear gravitational evolution, which is detectable in the amplitude of the \LyA forest power spectrum. The extent of the aforementioned regions is determined by the angular resolution $\theta^{\rm WF}$ of the Wiener-filtered convergence map and depth of the lensing efficiency function. In this analysis, we select \LyA forests in the redshift range $2.1-3.6$.}\label{fig:schematic}\end{figure} This method was proposed in \cite{2001ApJ...551...48Z,2009PhRvL.103i1304V,Vallinotto:2011bn}. In this paper, we present the first detection of this signal, and propose a new theoretical description of it, based on the response of the matter power spectrum to a mean overdensity.

We have presented the first detection of a cross-correlation between the \LyA forest of quasars and the gravitational lensing of the CMB. Our understanding of this correlation is based on the response of small-scale fluctuations in the matter density, measured by the one-dimensional power spectrum of the transmission in the \LyA forest, to large-scale overdensities probed by the convergence of CMB lensing. This signal corresponds to a bispectrum in the squeezed limit configuration where the two small-scale modes are of order $k \sim 0.1 - 1\ \mathrm{Mpc}/h$. This is the first measurement of the ${\mathrm{CMB \ lensing}-\mathrm{Ly}\alpha}$ integrated bispectrum, and it measures the non-linearity in the \LyA forest. Finally, this new observable tests our understanding of the relation between neutral hydrogen and dark matter. We measured the one-dimensional power spectrum and the linear bias of the \LyA forest, finding values consistent with hydrodynamical simulations. Even though the power spectrum is sensitive to a number of systematic effects, these are much less important in a cross-correlation measurement like the integrated bispectrum that we computed. The theoretical bispectrum is the sum of two contributions: the response of the linear matter power spectrum, theoretically well-understood and involving no free parameters, and the response of the bias and non-linear terms, computed up to an effective non-linear bias $b_2^{\rm eff}$ which we have fitted. We believe this model provides a reasonable explanation of the observed signal. However, we notice that our interpretation of the measured bispectrum is limited by theoretical uncertainties mainly related to baryonic physics. That is, the term $D(k, \mu)$, taken from \cite{2015JCAP...12..017A}, encodes a number of effects that are significant at very small scales (of order $k\sim 60\ h/\mathrm{Mpc}$, see \cite{2016JCAP...03..016C}), but the integral of the three-dimensional power spectrum in Eq.~\eqref{eq:p1d_th} gets contributions from $k$-modes greater than 10~$h$/Mpc, and we cannot neglect these effects. Moreover, this term and the redshift-space distorsion term $\beta$ may also respond to large-scale overdensities. Therefore, the effective non-linear bias term $b_2^{\rm eff}$ encompasses several uncertain contributions: it could be compared with simulations, providing both a valuable check for the simulation assumptions while shedding light on the relation between \LyA and dark matter. Another uncertainty arises from the fact that the ionizing UV flux of the quasars reduces the amount of neutral hydrogen around them, a phenomenon known as the proximity effect. Because overdense regions radiate more, the bias of neutral hydrogen $b_{\mathrm{H\textsc{i}}}(k)$ becomes negative at scales larger than ${k\sim0.01\ h/\mathrm{Mpc}}$ \cite{2014PhRvD..89h3010P}. This impacts the H\textsc{i} power spectrum for scales ${k \lesssim 0.1\ h/\mathrm{Mpc}}$ and may also affect our measurement in the lowest $k$-bin. In the future, hydrodynamical simulations could be compared to this new observable and used to model the dependence of the various non-linear terms in the one-dimensional power spectrum on the mean overdensity. In particular, the anisotropic linear bias of the \LyA forest, \emph{i.e.} its dependence on angle $\mu$, could be explored. Meanwhile, this measurement can inform simulations and help contrain the relation between intergalactic gas and dark matter. Other avenues of exploration are to study the dependence of the signal on redshift and perpendicular separation $r_\perp$. The upcoming {SDSS-IV/eBOSS} \cite{2016AJ....151...44D} data, covering a broad redshift range, will improve the signal-to-noise ratio and help measure the redshift dependence. Combined with a high signal-to-noise ratio CMB lensing map, it could allow for a measurement of the angular correlation between the \LyA forest and large-scale overdensities. Finally, because this bispectrum is sensitive to small scales observed in the \LyA forest, it could also provide additional constraints on the total mass of neutrinos and be used as a tool to study alternative models of dark matter predicting small-scale cut-offs.

1607.03625

1607

1607.01416_arXiv.txt

We present an update of a spectroscopic follow-up survey at low-resolution of a large number of RR Lyrae halo overdensity candidates found in the southern sky. The substructure candidates were identified in the RR~Lyrae catalog of Torrealba et al. (2015) using Catalina Real-time Transient Survey (CRTS) data. Radial velocities and mean metallicities have been estimated for target stars in almost half of the original overdensities to assess their potential membership to coherent halo features.

1607.01416

1607

1607.01550_arXiv.txt

We briefly introduce the disadvantages for Type Ia supernovae (SNe Ia) as standard candles to measure the Universe, and suggest Gamma-ray bursts (GRBs) can serve as a powerful tool for probing the properties of high redshift Universe. We use GRBs as distance indicators in constructing the Hubble diagram at redshifts beyond the current reach of SNe Ia observations. Since the progenitors of long GRBs are confirmed to be massive stars, they are deemed as an effective approach to study the cosmic star formation rate (SFR). A detailed representation of how to measure high-$z$ SFR using GRBs is presented. Moreover, first stars can form only in structures that are suitably dense, which can be parameterized by defining the minimum dark matter halo mass $M_{\rm min}$. $M_{\rm min}$ must play a crucial role in star formation. The association of long GRBs with the collapses of massive stars also indicates that the GRB data can be applied to constrain the minimum halo mass $M_{\rm min}$ and to investigate star formation in dark matter halos.

Gamma-ray bursts (GRBs) are the most powerful explosions in the cosmos, which can be divided into long GRBs with duration times $T_{90}>2$ s and short GRBs with $T_{90}<2$ s.\cite{Kouveliotou et al.1993} In theory, it is generally accepted that long GRBs are formed by the core collapses of massive stars,\cite{Woosley1993,Woosley et al.2006} while short GRBs are arised from the mergers of binary compact stars.\cite{Eichler et al.1989,Narayan1992} Because of their high luminosities, GRBs can be discovered out to very high redshifts. To date, the farthest burst detected is at $z=8.2$ (GRB 090423\cite{Tanvir et al.2009}) \footnote{A photometric redshift of $z\sim9.4$ for GRB 090429B was measured by Ref.~\refcite{Cucchiara et al.2011}.}. GRBs are therefore considered as a powerful tool for probing the properties of the early Universe. Type Ia supernovae (SNe Ia) have been treated as an ideal standard candle to measure dark energy and cosmic expansion.\cite{Perlmutter et al.1998,Schmidt et al.1998,Riess et al.1998} However, the highest redshift of SNe Ia so far is $z=1.914$.\cite{Jones et al.2013} It is hard to observe SNe Ia at redshifts $>2$, even with the excellent space-based platforms such as SNAP.\cite{Sholl et al.2004} Since lots of the interesting evolution of the Universe occurred before this epoch, the usage of SNe Ia in cosmology is limiting. In contrast to SNe Ia, GRBs can be detected at higher redshifts. Moreover, gamma-ray photons from GRBs, unlike that the optical photons of supernovae, are immune to dust extinction. The observed gamma-ray flux is an actual measurement of the prompt emission flux. Thus, GRBs are potentially a more promising cosmological probe than SNe Ia at higher redshifts. The possible use of GRBs as ``relative standard candles'' started to become reality after some luminosity relations between energetics and spectral properties were found.\cite{Amati et al.2002,Ghirlanda et al.2004,Yonetoku et al.2004,Liang et al.2005} These GRB luminosity relations have been deemed as distance indicators for cosmology. For example, Ref.~\refcite{Schaefer2003} constructed the first GRB Hubble diagram using two GRB luminosity indicators. With the correlation between the collimation-corrected gamma-ray energy $E_{\rm \gamma}$ and the spectral peak energy $E_{\rm p}$, Ref.~\refcite{Dai et al.2004} obtained tight constraints on cosmological parameters and dark energy. Ref.~\refcite{Ghirlanda et al.2004b} placed much tighter limits on cosmological parameters with the same $E_{\rm p}-E_{\rm \gamma}$ relation and SNe Ia. Ref.~\refcite{Liang et al.2005} constrained cosmological parameters and the transition redshift using a model-independent multivariable GRB luminosity relation. Ref.~\refcite{Schaefer2007} constructed a GRB Hubble diagram with 69 bursts with the help of five luminosity indicators. Ref.~\refcite{Kodama et al.2008} found that the time variation of the dark energy is very small or zero up to $z\sim6$ using the relation between the isotropic peak luminosity $L$ and $E_{\rm p}$. Ref.\refcite{Tsutsui et al.2009} extended the Hubble diagram up to $z=5.6$ based on 63 bursts using the $E_{\rm p}-L$ relation and shown that these GRBs were in agreement with the concordance model within $2\sigma$ level. In the meantime, a lot of works\cite{Firmani et al.2005,Xu et al.2005,Amati2006,Amati et al.2008,Liang et al.2006,Wang et al.2006,Liang et al.2008,Qi et al.2008a,Qi et al.2008b,Wei et al.2009,Yu et al.2009,Wei2010,Wang et al.2011a,Wei et al.2013,Ding et al.2015,Lin et al.2015,Izzo et al.2015,Wei et al.2015} have been done in this so-called GRB cosmology field. Please see Refs.~\refcite{Ghirlanda et al.2006,Amati et al.2013,Wang et al.2015} for recent reviews. With the improving observational techniques and a wider coverage in redshift, we now have a better understanding of the star formation history in the Universe. However, direct star formation rate (SFR) measurements are still quite difficult at high redshifts ($z\geq6$), particularly at the faint end of the galaxy luminosity function. Fortunately, the collapsar model suggests that long GRBs provide a complementary tool for measuring the high-$z$ SFR from a different perspective, i.e., by the means of investigating the death rate of massive stars rather than observing them directly during their lives. Due to the fact that long GRBs are a product of the collapses of massive stars, the cosmic GRB rate should in principle trace the cosmic SFR. However, the Swift observations reveal that the GRB rate does not strictly follow the SFR, but instead implying some kind of additional evolution. \cite{Daigne et al.2006a,Guetta et al.2007,Le et al.2007,Salvaterra et al.2007,Kistler et al.2008,Kistler et al.2009,Salvaterra et al.2009,Campisi et al.2010,Qin et al.2010,Wanderman et al.2010,Cao et al.2011,Virgili et al.2011,Lu et al.2012,Robertson et al.2012,Salvaterra et al.2012,Wang2013,Wei et al.2014} An enhanced evolution parametrized as $(1+z)^{\alpha}$ is usually adopted to describe the difference between the GRB rate and the SFR.\cite{Kistler et al.2008} In order to explain the observed discrepancy, several possible mechanisms have been proposed, including cosmic metallicity evolution,\cite{Langer et al.2006,Li2008} stellar initial mass function evolution,\cite{Xu et al.2009,Wang et al.2011} and luminosity function evolution.\cite{Virgili et al.2011,Salvaterra et al.2012,Tan et al.2013,Tan et al.2015} Of course, if we knew the mechanism responsible for the discrepancy between the GRB rate and the SFR, we could set a severe limit on the high-\emph{z} SFR using the GRB data alone. In the framework of hierarchical structure formation, a self-consistent SFR model can be calculated. Especially, the baryon accretion rate accounts for the structure formation process, which governs the size of the reservoir of baryons available for star formation in dark matter halos, can be obtained from the hierarchical scenario. First stars can form only in structures that are suitably dense, which can be parametrized by defining the minimum mass $M_{\rm min}$ of a dark matter halo of the collapsed structures where star formation occurs. In briefly, structures with masses smaller than $M_{\rm min}$ are served as part of the intergalactic medium and do not participate in the process of star formation. The minimum halo mass $M_{\rm min}$ must, therefore, play a crucial role in star formation. There are a few direct constraints on $M_{\rm min}$ as follows. In order to simultaneously reproduce the current baryon fraction and the early chemical enrichment of the intergalactic medium, Ref.~\refcite{Daigne et al.2006b} suggested that a minimum halo mass of $10^{7}-10^{8}M_{\odot}$ and a moderate outflow efficiency were required. With a minimum halo mass of $M_{\rm min}\simeq10^{11}M_{\odot}$, Ref.~\refcite{Bouche et al.2010} could explain both the observed slopes of the star formation rate--mass and Tully--Fisher relations. Ref.~\refcite{Munoz et al.2011} found that the minimum halo mass can be constrained by matching the observed galaxy luminosity distribution, in which $M_{\rm min}$ was limited to be $10^{8.5}\rm M_{\odot}<M_{\rm min}<10^{9.7}\rm M_{\odot}$ at the $95\%$ confidence level. The association of long GRBs with the death of massive stars provides a new interesting tool to investigate star formation in dark matter halos. In this work, we review the applications of GRBs in cosmology. The rest of this paper is arranged as follows. In Section~2, we construct the GRB Hubble diagram and describe its cosmological implications. The constraints on the high-$z$ SFR from GRBs are presented in Section~3, and the capability of GRBs to probe star formation in dark matter halos is presented in Section~4. Finally, the conclusions and discussion are drawn in Section~5. For the more details of full samples and analysis results we discussed here, please refer to Refs.~\refcite{Wei et al.2013,Wei et al.2014,Wei et al.2016}.

In this paper, we have briefly reviewed the status for the exploration of the early Universe with GRBs. A few solid conclusions and discussion can be summarized: (i) GRBs can be used as standard candles in constructing the Hubble diagram at high-$z$ beyond the current reach of SNe Ia observations. However, the dispersion of luminosity relations are still large, which restricted the precision of distance determination with GRBs. Since the classifications of GRBs are not fully understood, some contamination of the GRB sample in a correlation is unavoidable, which would make the correlation be dispersed. Note that among all SNe, only a small class SNe (SNe Ia) can be used as standard candles. On the other hand, the large dispersion may also due to that we have not yet identified the accurate spectral and lightcurve features to use for the luminosity correlations. In order to improve the precision of distance measurement, we should take efforts to investigate the classification problem of GRBs, and search for more precise luminosity relations, especially the relations with certain physical origins. (ii) GRBs provide an independent and powerful tool to measure the high-$z$ SFR. The central difficulty of constraining the high-$z$ SFR with GRBs is that one must know the mechanism responsible for the difference between the GRB rate and the SFR. The current \emph{Swift} GRB observations can in fact be used to explore the unknown mechanism. However, the results from Swift data have uncertainties because of the small GRB sample effect. Fortunately, the upcoming GRB missions such as the Sino-French spacebased multiband astronomical variable objects monitor (\emph{SVOM}) and the proposed \emph{Einstein Probe} (EP), with wide field of view and high sensitivity, will be able to discover a sufficiently large number of high-$z$ GRBs. With more abundant observational information in the future, we will have a better understanding of the mechanism for the SFR-GRB rate discrepancy, and we will measure the high-$z$ SFR very accurately using the GRB data alone. (iii) GRBs also constitute a new promising tool for probing star formation in dark matter halos. In order to constrain the minimum dark matter halo mass $M_{\rm min}$, we also have to know the relation between the GRB rate and the SFR. In addition to the obvious method of increasing the sample size of GRBs with the help of future dedicated missions, we suggest that a much more severe constraints on $M_{\rm min}$ will be achieved by combining several independent observations, such as the observational data of star formation history, the luminosity distribution of galaxies, and the redshift distribution of GRBs.

1607.01550

1607

1607.02487_arXiv.txt

{Emission lines from highly-excited states ($n \ge 5$) of $\mathrm{H}$- and $\mathrm{He}$-like ions have been detected in astrophysical sources and fusion plasmas. For such excited states, $R$-matrix or distorted wave calculations for electron-impact excitation are very limited, due to the large size of the atomic basis set needed to describe them. Calculations for $n \ge 6$ are also not generally available. We study the behaviour of the electron-impact excitation collision strengths and effective collision strengths for the most important transitions used to model electron collision dominated astrophysical plasmas, solar, for example. We investigate the dependence on the relevant parameters: the principal quantum number $n$ or the nuclear charge $Z$. We also estimate the importance of coupling to highly-excited states and the continuum by comparing the results of different sized calculations. We provide analytic formulae to calculate the electron-impact excitation collision strengths and effective collision strengths to highly-excited states ($n \ge 8$) of $\mathrm{H}$- and $\mathrm{He}$-like ions. These extrapolated effective collision strengths can be used to interpret astrophysical and fusion plasma via collisional-radiative modelling. } \authorrunning{L. F. Menchero et al.} \titlerunning{Scaling of collision strengths for highly-excited atomic states}

\label{sec:introduction} Spectral emission lines of H- and He-like ions have been used for diagnostics of both fusion and astrophysical plasmas for decades. Perhaps the most famous examples are the temperature and density diagnostics of the He-like ions in the X-rays, described by \cite{jordan1969}. However, the status of the atomic data for these ions still requires improvement. As described below, atomic data for ions in these sequences are generally only available up to the principal quantum number $n=5$. Atomic data for highly-excited levels are needed for a variety of reasons. First, lines from these levels (up to $n=10$) have been observed in laboratory plasma (see, e.g. the compilations in the NIST database \citealt{nist2013}) and recently also in X-ray spectra of solar flares (see, e.g. \citealt{kepa2006}). Second, transitions between highly-excited levels should be included for any appropriate collisional-radiative modelling of these ions. Third, even if in most astrophysical spectra lines from these levels are not readily visible, they do contribute to the X-ray pseudo-continuum, so they should be included in any spectral modelling. Calculating atomic data for highly-excited levels is not a trivial task and has various limitations, since it requires a significant increase in the size of the atomic basis set. In Section~\ref{sec:atomic} we review previous calculations for the electron-impact excitation of He- and H-like ions, and present the results of test calculations made with larger basis sets. These calculations are performed to see how well the extrapolated data agree with the calculated ones for higher $n$. In Section~\ref{sec:extrapolation} we study the behaviour of the electron-impact excitation collision strengths and effective collision strengths for several kinds of transitions. These are the most common transitions that decay producing the lines observed in astrophysics. We also estimate the importance of coupling to highly-excited states and the continuum by comparing the results of differently sized distorted wave and $R$-matrix calculations. We then provide analytic formulae to calculate electron-impact excitation collision strengths to highly-excited states ($n \ge 8$) of $\mathrm{H}$- and $\mathrm{He}$-like ions. This is done by extrapolating the results obtained with the $R$-matrix or distorted wave methods. Potentially, the method provides results up to $n=\infty$, although accuracy reduces as $n$ increases. In Section~\ref{sec:comparison} we compare the solar flare line intensities with those predicted by applying the extrapolation rules to the effective collision strengths. Finally, in Section~\ref{sec:conclusions} we summarise the main conclusions.

\label{sec:conclusions} We carried out several new calculations, both $R$-matix and distorted wave, for the electron-impact excitation of $\mathrm{H}$- and $\mathrm{He}$-like ions. We have shown the dependence of the effective collision strengths $\Upsilon$ with respect to the principal quantum number $n$ for several transition types of some benchmark ions of the $\mathrm{He}$-like isoelectronic sequence. We tested three models to reproduce the behaviour of the $\Upsilon(n)$ and extrapolate them to more highly-excited states. In general, the extrapolation rules do not give such accurate values for the effective collision strengths as explicit calculations do; instead, they give an approximation that can be used for estimation purposes in modelling. Clearly, it is first necessary to have a good starting calculation before applying the extrapolations to the data. We note that $R$-matrix results become increasingly uncertain for the highest energy states included in the CI and CC expansions of the target, due to their lack of convergence \citep{fernandez-menchero2015b}. The description of the atomic structure, energy levels and radiative data, and the corresponding effective collision strengths, increasingly lose accuracy as we approach the last states included in the basis expansions. The same happens with the distorted wave calculations. Even when the distorted wave results are not affected by the coupling or resonances, which are included in the $R$-matrix ones, the description of the atomic structure is its main limitation. In consequence, such calculations may not give more accurate results for high-$n$ than the extrapolation rules combined with accurate calculations for lower, but sufficiently excited, states. Among the three extrapolation models considered, we recommend the second one. The Model 2 two-point extrapolation provides results that are closer to the $R$-matrix data, and reproduces the behaviour at high-$n$. Nevertheless it is not a completely general rule, and each particular transition type should be analysed to check which is the most accurate extrapolation method for the ion. We do not recommend Model 1 the least-squares fitting for several reasons. First, having just four points is not enough for a good-quality fitting. Second, the first data points $n=2,3$ have not yet reached the required $n^{-3}$ behaviour. This occurs because they can not be considered Rydberg levels: they interact more strongly with the core and resonances play a larger role for excitation to these shells. Finally, the fitting method requires a considerably larger computational effort to obtain a result that may be worse than the simpler methods. The Model 3 one-point extrapolation gives, in general, a poorer estimate than the two-point one, and the computational work is not substantially reduced. On the other hand, sometimes the parameters estimated with the one-point rule are more stable with respect to the reference point $n_0$ than the two-point one, for example, the dipole transitions shown here. Due to the inherent uncertainty in data for the most highly-excited states, we do not recommend extrapolating the $\Upsilon$ from the last $n$-shell, but suggest dropping form the extrapolation the last one or two $n$-values. Also, the $n^{-3}$ behaviour applies only from a certain excited shell. The extrapolation has to start from a level when the atom can be considered as a Rydberg one. That is at least approximately two shells higher than the last occupied one ($n=3$ for $\mathrm{H}$- and $\mathrm{He}$-like sequences). We also do not recommend the use of models with an oversimplified atomic structure so as to reach high-$n$ shells. The extrapolation from an accurate calculation to lower-$n$ provides in general a better estimate than a larger explicit calculation with a poorer atomic structure. As an example application, we have compared lines ratios obtained with the presented extrapolation-rule Model 2 with observations of X-rays by RESIK. We obtain good agreement with the observed \ion{Si}{xiii} and \ion{S}{xv} ratios during the peak phase of solar flares, but the values during the impulsive phase are still outside the theoretical range. It is expected that during the impulsive phase non-equilibrium effects are present. We investigated whether a non-Maxwellian distribution such as a $\kappa$-distribution was able to increase the ratios, but did not find significant increases. We are currently investigating other possible causes. For photoionised plasmas the ions exist at temperatures lower than the peak abundance in an electron collision dominated plasma. At these lower temperatures, the $\Upsilon$ are affected more by resonance enhancement, and perhaps radiation damping thereof, and can depend greatly on the position of these resonances, which are determined by the atomic structure. Resonance effects can cause the $\Upsilon$ to deviate from the $n^{-3}$ rule. The extrapolation rules should therefore only be applied starting from a higher excited shell, so these effects are minimised. This occurs even if the calculations are perfectly accurate at low temperatures. In these case, we suggest to start the extrapolation at least from four shells above the last occupied ($n=5$ for $\mathrm{H}$- and $\mathrm{He}$-like sequences). We estimate that the accuracy of the extrapolation Model 2 is approximately $20\%$ for all transition types of moderately- and highly-charged ions, and approximately $50\%$ for low-charged ions. This estimate considers the core calculations to be perfect and the extrapolation carried out from a shell where the ion can be considered as Rydberg. The inaccuracies in the core calculations could lead to larger errors in some cases, especially for low-charged ions and/or low temperatures.

1607.02487

1607

1607.05328_arXiv.txt

{We investigate the effect of adiabatic regularization on both the tensor- and scalar-perturbation power spectra in \textit{nonminimally} coupled chaotic inflation. Similar to that of the \textit{minimally} coupled general single-field inflation, we find that the subtraction term is suppressed by an exponentially decaying factor involving the number of $ e $-folds. By following the subtraction term long enough beyond horizon crossing, the regularized power spectrum tends to the ``bare'' power spectrum. This study justifies the use of the unregularized (``bare'') power spectrum in standard calculations.}

\label{intro} Cosmic inflation \cite{Guth:1980zm, Starobinsky:1980te, Starobinsky:1979ty, Sato:1980yn, Linde:1983gd, Albrecht:1982wi} found its place in modern cosmology as a solution to the horizon and flatness problems and an explanation for the origin of primordial cosmological perturbations \cite{LiddlenLyth, Mukhanov} that served as seeds for the large-scale structures (e.g., galaxies and clusters of galaxies) nowadays constituting the Universe. Like any other system of ideas intended to explain nature or a part thereof, this theory has to satisfy at least the two-pronged requirement of science namely, self consistency and agreement with experiment. The latter requires a theory to explain existing measurement results before its inception and predict values of physical observables (e.g., in the case of inflation, non-Gaussianity, tensor-to-scalar ratio, spectral tilt, etc. \cite{Planck:2013jfk,Ade:2015lrj}). On the other hand, self-consistency ensures that pathways of ideas within the framework of the theory do not lead to illogical contradiction and problematic calculation results such as those involving non-removable divergences. On this side, inflation must pass checks involving loop corrections \cite{Senatore:2009cf,Seery:2007we}, divergences in predicted physical observables \cite{Parker:Toms}, and the more general question of renormalization (e.g., see Ref. \cite{Contillo:2011ag}). In this work, we focus on the side of self-consistency and investigate the effect of removing infinities from the two-point functions of scalar and tensor perturbations, on their power spectra. In particular, we consider the adiabatic regularization \cite{Parker:Toms, Parker:1974qw, Fulling:1974pu, Bunch:1980vc} of the power spectra for the perturbations in \textit{nonminimally coupled} single-field inflation \cite{Fakir:1990eg,Futamase:1987ua,Makino:1991sg,Komatsu:1999mt,Nozari:2010uu} (see \textbf{Sec. \ref{setUp}} for the details of the setup). It is an extension/generalization of the work of Urakawa and Starobinsky \cite{Urakawa:2009} involving the adiabatic regularization of the power spectra in the \textit{canonical} single-field inflation. In our previous work \cite{Alinea:2015pza}, we performed the extension/generalization of their study along a different pathway including cases where the speed of sound $ c_s^2, $ is not constant (that is, \textit{general} single-field inflation, or, as usually referred to in the literature, \textit{k}-inflation,) but maintained the \textit{minimal-coupling} aspect of the theory (see \textbf{Fig}. \ref{figFlowOfGen}). This time, adding a nonminimal coupling term on top of the canonical case ($ c_s^2 = 1 $), we wish to determine if adiabatic regularization will significantly modify the power spectrum contrary to what was found by Urakawa and Starobinsky and in our extension to the general single-field inflation of their study. \begin{figure}[hbt] \centering \makebox[\textwidth][c]{\includegraphics[scale=0.9]{generalization_02.eps}} \caption{Generalizations of the canonical inflation. The canonical case involves the Lagrangian $ \mathcal L = \frac{1}{2}M_\text{Pl}^2 R + X - V(\phi) $, where $ M_\text{Pl},\, R,\, X, $ and $ V $ are the Planck mass, Ricci scalar, kinetic term, and potential term respectively. On the one hand, nonminimally coupled chaotic inflation (right hand most box) involves the addition of nonminimal coupling (N.M.C.) term to the canonical case. On the other hand, $ k $-inflation involves a Lagrangian that is a general functional of $ X $ and $ \phi $; that is, $ P(X, \phi ) $.} \label{figFlowOfGen} \bigskip \end{figure} Before we proceed with the calculation to fulfill this aim, it is good to shed some light on adiabatic regularization in the context of the problem we are dealing in this work. The two-point function of the gauge-invariant scalar perturbation $ \mathcal R $ \cite{Bardeen:1980kt} can be written as \begin{align} \label{twoPtFn} \langle \mathcal R(\tau, \mathbf{x})\mathcal R(\tau, \mathbf{y})\rangle &= \int\frac{\dd k}{k} \frac{\sin(k|\mathbf x - \mathbf y|)}{k|\mathbf x - \mathbf y|} \Delta^2_{\mathcal R}(k,\tau), \end{align} where $ \tau $ is the conformal time, $ (\mathbf x, \mathbf y) $ are pairs of position vectors, $ k $ is the wavenumber, and $ \Delta^2_{\mathcal R} $ is the dimensionless power spectrum given by \begin{align} \Delta _{\mathcal R}^2(k,\tau ) = \frac{k^3}{2\pi ^2}\big|\mathcal R_k(\tau )\big|^2, \end{align} with the subscript $ k $ in $ \mathcal R_k $ indicating momentum space. (Similar expression holds for the tensor perturbation.) The adiabatic condition \cite{Parker:Toms,BirrellDavies} tells us that $ |\mathcal R_k(\tau )| $ scales as $ 1/k $ for large $ k $ leading to $ \Delta^2_{\mathcal R} \sim k^2 $. It follows that in the coincidence limit, $\mathbf x\,\rightarrow \,\mathbf y$, the integral diverges. As applied to this scenario, adiabatic regularization is a method of systematically subtracting out terms in the integrand to yield a finite two-point function, subject to the assumption that the adiabatic condition holds. Such a subtraction scheme is correspondingly reflected in the power spectrum. One may then see that the physical power spectrum is the difference between the ``bare'' and the corresponding subtraction term. There is a disagreement about this viewpoint. Parker and his collaborators \cite{Parker:2007,Agullo:2009,Agullo:2009:II} advocate the need for adiabatic regularization and claim that it could significantly modify the power spectrum (with both the ``bare'' and the subtraction term evaluated before or at the moment of horizon crossing), producing expressions not in accord with the ``bare'' power spectrum used in standard calculations. Some other people casted some doubt on the necessity of this subtraction scheme based on the claims that (a) the regularised power spectrum can become negative when one goes beyond the minimal subtraction scheme \cite{Finelli:2007fr} and (b) in the far infrared regime, the adiabatic expansion is no longer valid \cite{Durrer:2009ii}. Moreover, in the review paper of Bastero-Gil et al. \cite{Bastero-Gil:2013}, the authors argued that the cosmic microwave background (CMB) anisotropy variable $ C_l $ depends on the inflaton two-point function evaluated at $ \mathbf x \ne \mathbf y$; thus, practically avoiding the divergence in (\ref{twoPtFn}). We adopt a position that adiabatic regularization may be necessary and its application to calculation of physical observables produces physically consistent result. In so taking this position, we wish not to devote this paper as an argumentative defense of such position. What we do here is leave some words in favor of our chosen side and focus on the main subject of this work; that is, the \textit{application} of adiabatic regularization to the calculation of power spectra. First, the coincidence limit $ \mathbf x \,\rightarrow \,\mathbf y $ is a mathematical possibility that cannot be avoided by the mere existence of physical circumstances where such a limit may not hold. Furthermore, when one considers potentials involving self interaction or the energy-momentum tensor\footnote{We do not deal with loop corrections, higher-order corrections to the power spectrum, and the energy-momentum tensor. They deserve separate studies of their own.}, the divergences are certainly inevitable. Second, the adiabatic regularization \textit{includes} as its prescription that \textit{the subtraction scheme be minimal}; rather informally, those terms producing infinity should be the only subjects of regularization to remove them. Since the adiabatic expansion say, of $ \mathcal R $, is in general, not convergent but only asymptotic \cite{Parker:Toms}, such a prescription prevents a given quantity---one whose mathematical sickness we want to cure in the first place---from becoming unphysical or ill-behaved (e.g., power spectrum becoming negative). Third and last, determining the definite cut on the value of the spectrum of $ k $ as to when adiabatic regularization should be or should not be performed would seem to introduce an unnecessary complication to this method whenever quantities in momentum space are considered. If only to tame the doubts casted by Durrer et al. \cite{Durrer:2009ii} that the adiabatic expansion may not be valid in the far infrared limit, we note that there are good indications that the subtraction term becomes negligible for the quantities involving superhorizon modes (see for instance, Refs. \cite{Urakawa:2009,Alinea:2015pza}). In this work as in our previous work \cite{Alinea:2015pza}, we follow the track laid down by Urakawa and Starobinsky \cite{Urakawa:2009}. We perform adiabatic regularization using the minimal subtraction scheme and follow the subtraction term long after the first horizon crossing during inflation. As what they found out, the subtraction term exponentially decays with the number of $ e $-folds, resulting to the standard expression; i.e., the regularized power spectrum tends to the ``bare'' power spectrum. As already mentioned above, we have done the same for the extension to the minimally coupled general single-field inflation. This study showed that apart from cases that may significantly deviate from the condition of scale invariance, the subtraction ``leads to no difference in the power spectrum'' \cite{Alinea:2015pza}. It is then our aim in this investigation to see if the same holds for the case of nonminimally coupled inflation. This paper is organized as follows. In \textbf{Sec. \ref{setUp}}, we provide the setup for nonminimally coupled chaotic inflation. In so doing, we state the action and the background equations necessary for our analysis and calculations. In \textbf{Sec. \ref{``bare''Power}}, perturbations are introduced about the FLRW metric and the ``bare'' power spectra for these perturbations (scalar and tensor) are calculated up to first order in the slow-roll parameters. After this, we compute the subtraction terms (also up to first order in the slow-roll parameters) and perform adiabatic regularization of the power spectra in \textbf{Sec. \ref{adiabaticReg}}. Following this, we justify the use of both the Einstein and Jordan frames in our calculation in \textbf{Sec. \ref{remarksEinsteinFrame}}. In the last section, we state our conclusion and future prospects about further generalizing this work to cases involving non-constant speed of sound, higher-order corrections to the power spectrum, and gravity beyond general relativity. \bigskip

\label{conclude} Power spectrum is one of the most important quantities in inflationary cosmology. In this work, we investigated the effect of ultraviolet (UV) regularization on the power spectrum within the framework of nonminimally coupled chaotic inflation. Our calculation suggests that the adiabatic regularization leads to no difference between the ``bare'' and regularized power spectra for both scalar and tensor perturbations. Going beyond the time of horizon crossing, the subtraction term exponentially decays with respect to the number of $ e $-folds and becomes insignificant in the long run; the expansion of the Universe essentially erases the subtraction term. This is the same result we obtained for minimally coupled general single-field inflation in our previous study \cite{Alinea:2015pza}. Thus, for both the minimally coupled and nonminimally coupled cases, the current and the previous studies together with that of Urakawa and Starobinsky \cite{Urakawa:2009}, justify the use of the ``bare'' power spectrum in standard calculations. So far, we have only investigated the area of nonminimally coupled inflation where the speed of sound is constant. It would be interesting to see the effect of varying speed of sound $ c_s^2 $, in scenarios where the usual kinetic-plus-potential term is replaced by a general ``pressure'' function $ P(X, \phi ) $, where $ X \equiv -\frac{1}{2}g^{\mu \nu }\phi _{,\mu }\phi _{,\nu } $. The subtraction term would significantly be more complicated because of the involvement of the time derivatives of $ c_s^2 $ and the more complex relationship between the speed of sound in the Jordan and Einstein frames. From a bird's eye view, three main terms have to be considered namely, those due to (a) the speed of sound, (b) variation of the Hubble parameter, and (c) the effects of the nonminimal coupling term. Beyond this scenario, one may also consider nonminimal inflation beyond Einstein gravity; e.g., instead of $ h(\phi )R $, one may consider $ h(\phi )f(R) $\footnote{Such generalization may be relevant when quantum corrections are considered. For instance, in Ref. \cite{Salvio:2015kka}, it is shown that at the quantum level, ``the renormalization for large non-minimal coupling requires an additional degree of freedom that transforms Higgs inflation into Starobinsky $ R^2 $ inflation, unless a tuning of the initial values of the running parameters is made.''}, where $ f(R) $ is a function of the Ricci scalar. In addition to this, we have not addressed the more challenging problems involving interactions and loop corrections. Considering all these problems, we look ahead to establishing a general set of conditions within the framework of general single-field inflation, both minimal and nonminimal, that would guarantee the null effect of adiabatic regularisation on both scalar and tensor power spectra.

1607.05328

1607

1607.00027_arXiv.txt

Open clusters and associations are groups of young stars, respectively bound and unbound, that share the same origin and disperse over time into the galactic field. As such, their formation and evolution are the key to understand the origin and properties of galactic stellar populations. Moreover, since their members have about the same age, they are ideal laboratories to study the properties of young stars and constrain stellar evolution theories. In this contribution, I present our current knowledge on open clusters and associations. I focus on the methods used to derive the statistical properties (IMF, spatial distribution, IMF) of young stars and briefly discuss how they depend on the environment. I then describe how open clusters can be used as probes to investigate the structure, dynamics and chemical composition of the Milky Way. I conclude by presenting the Gaia mission and discuss how it will revolutionize this field of research.

It is commonly said that stars form in ``clusters''. However it would be more correct to say that stars form in groups with $N=10$ to $10^5$ and a disctinction should be made between young open clusters and associations. While both of them correspond to physically associated groups of population I stars, sharing the same origin and moving together through the Galaxy, a cluster is gravitationally bound and an association is not. Observationally, a cluster (like the Pleiades) shows a clear concentration of young stars above the surrounding stellar background while an association corresponds to a region where there is a higher than normal density of O-B stars (OB associations, such as Orion OB1) or T-Tauri stars (T associations, like TW Hydra). The average stellar density of open clusters can go from 0.1 to about 10 stars/pc$^3$ in a typical diameter of a few parsecs, yielding a few tens to 10$^5$ stars. An association contains typically 10 to 1000 stars which are spread over a much wider area (several tens to a few hundred parsecs in diameter), corresponding to a stellar density of $\sim0.01$ stars/pc$^3$ or less (see Figure~\ref{density}). \begin{figure}[htbp] \centering \includegraphics[width=0.9\textwidth]{density.png} \caption{Typical density of associations, open clusters and globular clusters.} \label{density} \end{figure} Another difference between a cluster and an association arising directly from their definition is the lifetime. An open cluster is loosely bound and its star members will stay together for about 100 Myr to a few Gyr for the densest ones, before it gets disrupted by the galactic tidal field. On the contrary, an association is unbound and disperses rapidly, over a timescale of 10 to 100 Myr. This means that it is much more difficult to detect associations and get a full census of their population. Very large areas of the sky need to be covered to identify a few members among thousands of unrelated field stars based on colors or proper motions. In particular, Hipparcos data have been very powerful in detecting OB association (e.g. \cite{deZ99}). However, caution must be taken when using kinematic informations for $\sim100$ Myr associations as lots of field objects may share the same proper motion due to resonant trap in the galactic field (e.g. ABDor; \cite{fam08}). To date, about 80 O-B associations and 3000 open clusters are known in the solar neighborhood (up to $\sim1.8$ kpc; \cite{kha13}), but they may be up to $10^5$ in the Galaxy. Despite the above differences, open clusters and associations show strong similarities. They are located in the arms of the Milky Way and other spiral galaxies, they are composed of young stars that recently formed in the disks of galaxies, and they continue to form (as opposed to globular clusters). All the stars belonging to the same cluster/association have the same age (with maybe a spread of a few Myr) and metallicity, having formed from the same cloud of gas and dust. Open clusters and associations can therefore be used to probe the galactic disk structure and formation rate, as well as to study the young star properties and their formation process. Where does the galactic disk population come from: open cluster, association, or both ? Is there any prefered mode of star formation ? Do clusters and associations form similarly but evolve differently or do they form from different physical conditions (density, turbulence, magnetic field...) ? How does the cluster/association environment affect the stellar properties ? In order to be able to answer the above questions, we first need to be able to distinguish a cluster from an association, or in other words to know whether a stellar group is gravitationally bound or not. The total energy of a system is $E=T+U$ with $U=-GM_c^2/R_g$ the potential energy, and $\displaystyle T=\frac{1}{2}M_c\sigma_v^2$ the kinetic energy. $M_c$ is the cluster mass, $R_g$ the {\it gravitational} radius and $\sigma_v$ the 3D velocity dispersion. A system is bound if $E<0$, i.e., if $\displaystyle \sigma_v<\sqrt{2\frac{GM_c}{R_g}}$ or $\sigma_v<\sqrt{2}\,\sigma_v^{vir}$ where $\sigma_v^{vir}$ corresponds to the {\it virialized} velocity dispersion given by $2T+U=0$. Therefore a subvirial or virialized system is bound while a supervirial system can be marginally bound or unbound. In practice, it is not such an easy task to know the dynamical state of a stellar group as one first needs to estimate its mass, gravitational radius, and 3D velocity dispersion. Masses are estimated either directly by summing up individual masses, or indirectly from the tidal radius. Indeed $\displaystyle M_c=\frac{4A(A-B)r_t^3}{G}$ (\cite{kin62}) where A and B are the Oort constants and $r_t$ is the tidal radius. $R_g$ is usually derived from the core radius $r_c$ ($\sim0.2-0.3 R_g$) that is obtained by fitting a King distribution to the radial profile -- but this can only be done if the group is centrally concentrated. Typical values for open cluster core radii are 1-2 pc and 10-25 pc for tidal radii. The velocity dispersions in open clusters are typically only a few km/s, indicating that individual stellar velocities have to be measured with an accuracy $<1$km/s which requires high resolution spectroscopic observations. Moreover $\sigma_v$ has to be corrected from bias introduced by binaries (e.g. \cite{chb14}, \cite{cot12}). Recent studies have shown that among the young massive star clusters Westerlund 1 is subvirial and strongly bound (\cite{cot12}) while NGC3603 is about virialized and marginally bound (\cite{pan13}). The Cyg OB2 association is indeed unbound (\cite{kim07,kim08}; \cite{wri14}) but does not show any coherent expansion pattern on large scale. And the nearby star forming region IC348 is supervirial but probably bound thanks to the gas mass (\cite{cot15}). For faraway regions that are barely resolved and for which $\sigma_v$ cannot be measured, another diagnosis has been proposed by \cite{gpz11}. Noticing that all the stellar systems that stayed together for more than 30Myr and are most probably bound have an age that is larger than their crossing time $t_{cr}$, they suggest to use the ratio $\pi=\textrm{age}/t_{cr}$ as a proxy to distinguish between clusters ($\pi>1$) and associations ($\pi<1$). However the $\pi$ value is only indicative for regions younger than 10Myr and the distinction is not so clear. For embedded regions, in which a lot of gas remains in particular, the differentiation between cluster and association is meaningless. Indeed a group that is bound thanks to the gas mass may become unbound after gas removal. In this lecture, I will first present the properties of young stars in clusters (section~\ref{stars}), before describing how the global properties of open clusters can be used as probes of the galactic thin disk in section~\ref{global}. I will then present the Gaia mission and discuss what we will learn from it in section~\ref{gaia}.

It is clear that the Gaia mission is going to revolutionize our understanding of the stellar cluster formation and evolution as well as our view of the Milky Way. To fully exploit the immense amount of data coming from both Gaia and complementary programmes one needs to be prepared however. The data analysis will require new statistical tools to efficiently study young stars and star clusters in 6-dimensional phase space (three position and three velocity dimensions), as well as folding in a wealth of multi-dimensional data. Moreover, hybrid hydrodynamic and N-body simulations of star forming regions are needed to link the collapse of molecular clouds with the dynamical evolution of young embedded stellar clusters. We will need many simulations that cover a wide variety of initial and final conditions as well as a number of new statistical diagnostics that will allow us to quantitatively compare observations and simulations. These diagnostics will have to be extensively tested against simulations to make sure that they reliably trace the cluster evolution and that they can be used to resolve the initial conditions.

1607.00027

1607

1607.01021_arXiv.txt

We describe a combined halo model to constrain the distribution of neutral hydrogen (\HI{}) in the post-reionization universe. We combine constraints from the various probes of \HI{} at different redshifts: the low-redshift 21-cm emission line surveys, intensity mapping experiments at intermediate redshifts, and the Damped Lyman-Alpha (DLA) observations at higher redshifts. We use a Markov Chain Monte Carlo (MCMC) approach to combine the observations and place constraints on the free parameters in the model. Our best-fit model involves a relation between neutral hydrogen mass $M_{\rm HI}$ and halo mass $M$ with a non-unit slope, and an upper and a lower cutoff. We find that the model fits all the observables but leads to an underprediction of the bias parameter of DLAs at $z \sim 2.3$. We also find indications of a possible tension between the \HI{} column density distribution and the mass function of \HI{}-selected galaxies at $z\sim 0$. We provide the central values of the parameters of the best-fit model so derived. We also provide a fitting form for the derived evolution of the concentration parameter of \HI{} in dark matter haloes, and discuss the implications for the redshift evolution of the \HI{}-halo mass relation.

Studying the evolution of neutral hydrogen (\HI{}) in the post-reionization universe offers key insights into cosmology and galaxy formation. Intensity mapping, in which the \HI{} intensity fluctuations are mapped out without the resolution of individual galaxies, is a novel technique which allows the study of \HI{} on large scales and promises unprecedented constraints on galaxy formation and evolution, cosmology \citep{chang10, masui13, switzer13, santos2015, bull2015} and stringent constraints on models of dark energy \citep[e.g.,][]{chang2008} and modified gravity \citep[e.g.,][]{hall2013}. The 21-cm intensity power spectrum is also an important tracer of the underlying large-scale structure in the post-reionization universe ($z \sim 0-6$), due to the absence of the complicated reionization astrophysics at these epochs. The key astrophysical ingredients in the estimation of the 21-cm intensity mapping power spectrum are the density parameter of \HI{}, $\Omega_{\rm HI}$, and its bias $b_{\rm HI}$, which measures how the \HI{} is clustered with respect to the underlying dark matter. Constraints on these parameters come from a variety of astrophysical probes (a detailed summary is available in \citet{hptrcar2015}, hereafter Paper 1). At low redshifts, \HI{} gas in galaxies and their environments is primarily studied through blind 21-cm emission line surveys of nearby galaxies like the \HI{} Parkes All Sky Survey \citep[HIPASS;][]{zwaan05} and the Arecibo Fast Legacy ALFA (ALFALFA) Survey \citep{martin10}. These provide measurements of the mass function of \HI{}-selected galaxies, from which the density parameter of \HI{} in galaxies, $\Omega_{\rm HI, gal}$ can be derived. The clustering of the \HI{}-selected galaxies, which constrains the galaxy bias $b_{\rm HI, gal}$, can also be measured from these surveys \citep{martin12}. At intermediate redshifts ($z \sim 0-1$), intensity mapping experiments typically constrain the combination $\Omega_{\rm HI} b_{\rm HI}$ \citep[e.g.,][]{switzer13}. At higher redshifts, the major reservoirs of \HI{} gas in the post-reionization universe are the Damped Lyman-Alpha (DLA) systems. These systems have very high column densities ($> 10^{20.3}$ cm$^{-2}$) and are found to contain more than 80 \% of the neutral hydrogen at redshifts 2-5 \citep[e.g.,][]{lanzetta1991}. DLAs have been studied through line-of-sight absorption against bright background quasars \citep[][]{noterdaeme12, prochaska09}, as well as direct imaging surveys \citep{fumagalli2014, fumagalli2015}. The study of DLA systems enables the measurements of the column density distribution, $f_{\rm HI}(N_{\rm HI})$ of the DLAs at various redshifts, the incidence $dN/dX$ of the DLAs per unit absorption distance interval, and the density parameter of \HI{} in DLAs, $\Omega_{\rm DLA}$. Recently, the large-scale bias parameter of DLA systems, $b_{\rm DLA}$ has also been constrained through a cross-correlation study with the Lyman-$\alpha$ forest \citep{fontribera2012}. Analytical and simulation techniques have typically been used to model the \HI{} radio observations at lower redshifts \citep[e.g.,][]{dave2013, bagla2010, guhasarkar2012, kim2016, villaescusa2015} and the DLA observations at higher redshifts \citep{fumagalli2011, bird2014, barnes2014} separately. Modelling the the 21-cm based observables requires knowledge of the \HI{} - halo mass relation, $M_{\rm HI} (M)$ which is combined with the radial distribution of \HI{} (the \HI{} density profile, $\rho_{\rm HI} (r)$) to derive the DLA parameters. In \citet[][hereafter Paper 2]{hptrcar2016}, we combined the analytical approaches in the literature towards a consistent picture of \HI{} evolution across redshifts. We fitted the available data both from the 21-cm based as well as from the DLA-based observations by modifying the 21-cm based approach in the literature to also include the DLA profile. We found that the free parameters can be fixed by fitting to the combined set of 21-cm and DLA observations, however, a model that is consistent with the low-redshift observations leads to a prediction for the bias parameter of the DLAs that is lower than the observed value. On the other hand, a model that is consistent with all observations requires very rapid evolution of the nature of \HI{} bearing host dark matter haloes between redshifts 0-2. Hence, it was found that bringing together the low-and high-redshift observations, while not having been attempted in the literature before, has interesting consequences for the evolution of the free parameters involved in the modelling of \HI{}. In this paper, we revisit our previous modelling of \HI{} to combine all the available data within the fully statistical framework of a Markov Chain Monte Carlo (MCMC) analysis, using a six free parameter halo model which is a generalization of the \HI{}-halo mass relation used in previous literature. The model also generalizes the radial distribution of \HI{} to take into account the effective redshift evolution of the concentration of \HI{} in the dark matter haloes. We use a larger dataset compared to the previous work, including also the \HI{} mass function at $z \sim 0$ \citep{zwaan05}. We find that the resulting combination of data, together with the generalized model, enables fairly strong constraints on the free parameters. This also enables us to explore the correlations between the parameters and their posterior distributions. We again find that a model that is reasonably consistent with all observations underpredicts the measured DLA bias at $z \sim 2.3$. We also find indications of a possible tension between the \HI{} mass function and the column density distribution at $z \sim 0$. The evolution of the free parameters in the model has important physical implications for the power spectrum of \HI{} 21-cm intensity fluctuations, as also for the connections between \HI{} and galaxies. The paper is organized as follows. In the next section, we describe the details of the halo model, focusing on the free parameters involved and their physical implications for the \HI{} distribution in the dark matter haloes. We also review the procedure for obtaining the DLA- and \HI{} based quantities from the analytical model. In Sec. \ref{sec:data}, we briefly describe the details of the data sets used in the analysis. In the next section, Sec. \ref{sec:mcmc}, we outline the MCMC analysis used for the combination of the data, and obtain the best-fitting model from the analysis. We compare the model predictions to the observations where available, and describe the model predictions at higher redshifts. We summarize the physical implications of our findings and the future outlook in a concluding section (Sec. \ref{sec:summary}).

\label{sec:summary} In this paper, we have carried out a detailed statistical study of the evolution of \HI{} across redshifts by utilizing all the available data, both from the 21-cm based as well as the DLA-based observables. We build upon the results of our previous fitting to the observations (Paper 2) by generalizing the form of $M_{\rm HI} (M)$ and the \HI{} profile from older analyses, thus introducing extra degrees of freedom. We thus explore how well we can use the data available at present to develop a halo-model based approach towards analyzing \HI{}. The model free parameters are constrained using a full MCMC analysis so that the posterior distributions and the correlations between parameters can be effectively explored. We find that a model which is reasonably consistent with all available data underpredicts the value of DLA bias at $z \sim 2.3$, as shown in Figs. \ref{fig:redshift23} and \ref{fig:bias}. This was also observed with the previous fitting analysis (Paper 2). The model predictions are nevertheless consistent with the results of DLA imaging studies, which favour the association of high-redshift DLAs with faint dwarf galaxies \citep{cooke2015, fumagalli2014, fumagalli2015}, and the results of hydrodynamical simulations \citep[e.g.,][]{rahmati2014, dave2013}. The introduction of the new parameters $\beta$ and $\gamma$ constrains the shape of the \HI{}-halo mass relation and also the form and evolution of the \HI{} density profile. Constraining $\beta < 0$ results in a smaller weight being given to the higher range of host halo masses, thus allowing the upper cutoff $v_{c1}$ to be higher than for the case with $\beta = 0$. It is also seen from the results of abundance matching \citep{papastergis2013} that a flatter $M_{\rm HI} - M$ relation is favoured by the observations of ALFALFA \HI{} galaxies. However, there are indications of a possible tension between the \HI{} column density distribution and the mass function of \HI{}-selected galaxies at $z \sim 0$. This tension may not be fully alleviated by the systematic effects in the measurements of the column density distribution and the mass function at $z \sim 0$. It is possible that the form of the profile at low redshifts needs to be modified in order to fit the observed column density distribution. We hope to consider other forms of the profile such as an exponential surface density profile, as well as a possible variation of the form of the profile across redshifts in future work. With the introduction of the parameter $\gamma$, the redshift evolution of the concentration $c_{\rm HI}$ can be effectively explored, and it can be seen that a mild increase in \HI{} concentration with redshift is favoured by the best-fitting model. In future work, it would be interesting to extend the results of this halo model framework to include small-scale clustering, measured from surveys of \HI{} selected galaxies and DLA systems.This serves to provide additional constraints on the profile parameters, as well as their evolution across redshifts. This extension would also enable comparison to the results of hydrodynamical simulations, and thus shed light on the connections between \HI{} and galaxy evolution over cosmic time. Finally, such a combined halo model would have important consequences for constraining the expected signal in the \HI{} power spectrum, to be measured with current and future intensity mapping experiments.

1607.01021

1607

1607.03448_arXiv.txt

We discuss the design considerations and initial measurements from arrays of dual-polarization, lumped-element kinetic inductance detectors (LEKIDs) nominally designed for cosmic microwave background (CMB) studies. The detectors are horn-coupled, and each array element contains two single-polarization LEKIDs, which are made from thin-film aluminum and optimized for a single spectral band centered on 150 GHz. We are developing two array architectures, one based on 160 micron thick silicon wafers and the other based on silicon-on-insulator (SOI) wafers with a 30 micron thick device layer. The 20-element test arrays (40 LEKIDs) are characterized with both a linearly-polarized electronic millimeter wave source and a thermal source. We present initial measurements including the noise spectra, noise-equivalent temperature, and responsivity. We discuss future testing and further design optimizations to be implemented.

\label{sec:intro} % Lumped element kinetic inductance detectors (LEKIDs) are superconducting resonators, which are also photon detectors. The resonance is set by a capacitor and inductor, the latter of which has both geometric and kinetic components. The LEKID inductor acts as the absorber and is impedance matched to the incoming radiation. The absorbed photons break Cooper pairs which changes the quasiparticle density and subsequently the surface impedance and kinetic inductance. This results in a shift in resonance frequency and quality factor both of which are read out through a transmission line. LEKIDs, and more broadly microwave kinetic inductance detectors (MKIDs), have been successfully demonstrated and deployed for a range of frequencies\cite{Mazin2013,nikatsz13}. Photon-noise limited performance has been shown in multiple frequency bands as well\cite{mauskopf14, Hubmayr2014}. For cosmic microwave background (CMB) polarization studies it is important that the detector noise to be sub-dominant to the photon noise. Current CMB experiments employ thousands of detectors. To further increase the sensitivity, it is necessary to increase the pixel count. LEKIDs are a natural candidate as hundreds of detectors\cite{McHugh2012, baselmans2015} can be read out on a single transmission line. We are developing dual-polarization LEKIDs that have two resonators within a single optical element. The two resonators are sensitive to orthogonal polarizations for observation at a frequency band centered at 150~GHz. Dual-polarization LEKIDs will effectively double the number of detectors for a given focal plane area compared to single polarization detectors. LEKIDs have been demonstrated as sensitive devices for absorbing single-polarization radiation at millimeter wavelengths\cite{flanigan_2016}, and dual-polarization radiation for far infrared\cite{Dober_2016}. In this proceedings we present (i) design considerations for dual-polarization LEKIDs at millimeter-wavelengths, (ii) initial test results, and (iii) steps for further optimization and testing. \begin{figure} [t!] \centering \begin{tabular}{c} % \includegraphics[height=6.6cm]{dual-pol_LEKID_160um1.pdf} \includegraphics[height=6.5cm]{20-element_dual_pol.pdf} \end{tabular} \caption[example] { \label{fig:detector_schematic} {\sl Left:} Schematic of a single detector. The resonators corresponding to the orthogonal polarizations are shown in red and blue. The dotted circle represents the waveguide exit aperture. The resonator inductor acts as the photon absorber. The resonators are capactively coupled to the transmission line. {\sl Center:} Cross-section view of a single-element. The horn aperture tapers to a cylindrical waveguide which also acts as a high-pass filter. A choke impedance matches between the waveguide and device while also controlling lateral radiation loss. The detectors are fabricated on silicon and directly illuminated. The package bottom acts as the backshort, the distance of which is set by the silicon wafer thickness. {\sl Right:} Layout of the test array. There are 20-elements, or 40 resonators, that are horn coupled and 2 dark elements. A single transmission line reads out all the devices. Each resonator has a unique resonance frequency set by the IDC value.} \end{figure}

We have successfully demonstrated dual-polarization devices for millimeter-wavelengths. The devices have flat noise within the device band and measured NETs of 36~and~52~$\mu K \sqrt{s}$ referenced to a 4~K load. The devices are responsive, with an optical responsivity of $\sim $20~ppm/K at a 4~K load. Initial tests show the two polarizations have similar responsivities. We are currently performing measurements to determine the polarization selectivity. Further optimizations to the detector design for ground-based observing have been implement in a 542-detector array that will be tested imminently. The large array will also allow the multiplexing capabilities to be further tested. The design presented here and initial test results show that LEKIDs work well and are a promising technology for future CMB polarimetry experiments. \begin{figure} [t] \centering \begin{tabular}{c} \includegraphics[height=6.2cm]{pol_A_1434_5_MHz.pdf} \includegraphics[height=6.2cm]{pol_B_1063_4_MHz.pdf} \end{tabular} \caption[example] {\label{fig:noise} {\sl Left:} Representative noise spectra for a resonator of polarization A design. The dotted line corresponds to a noise equivalent temperature (NET) of 36 $\mu$K$\sqrt{s}$ referenced to 4~K. {\sl Right:} Representative noise spectra for a resonator of the orthogonal polarization B to that in the left plot. The dotted line corresponds to a NET of 52 $\mu$K$\sqrt{s}$ referenced to 4~K. } \end{figure}

1607.03448

1607

1607.00094_arXiv.txt

{ The far-infrared - radio correlation connects star formation and magnetic fields in galaxies, and has been confirmed over a large range of far-infrared / radio luminosities, both in the local Universe and even at redshifts of $z\sim2$. Recent investigations indicate that it may even hold in the regime of local dwarf galaxies, and we therefore explore here the expected behavior in the regime of star formation {surface densities} below $0.1$~M$_\odot$~kpc$^{-2}$~yr$^{-1}$. We derive two conditions that can be particularly relevant for inducing a change in the expected correlation: a critical star formation surface density to maintain the correlation between star formation rate and the magnetic field, and a critical star formation surface density below which cosmic ray diffusion losses dominate over their injection via supernova explosions. For rotation periods shorter than $1.5\times10^7 (H/\mathrm{kpc})^2$~yrs, with $H$ the scale height of the disk, the first correlation will break down before diffusion losses are relevant, as higher star formation rates are required to maintain the correlation between star formation rate and magnetic field strength. For high star formation surface densities $\Sigma_{\rm SFR}$, we derive a characteristic scaling of the non-thermal radio to the far-infrared / infrared emission with $\Sigma_{\rm SFR}^{1/3}$, corresponding to a scaling of the non-thermal radio luminosity $L_s$ with the infrared luminosity $L_{th}$ as $L_{th}^{4/3}$. The latter is expected to change when the above processes are no longer steadily maintained. In the regime of long rotation periods, we expect a transition towards a steeper scaling with $\Sigma_{\rm SFR}^{2/3}$, implying $L_s\propto L_{th}^{5/3}$, while the regime of fast rotation is expected to show a considerably enhanced scatter, as a well-defined relation between star formation and magnetic field strength is not maintained. The scaling relations above explain the increasing thermal fraction of the radio emission observed within local dwarfs, and can be tested with future observations by LOFAR as well as the SKA and its precursor radio telescopes. }

Over the last years, magnetic fields have been detected in a significant number of local dwarf galaxies. This includes prominent examples such as the Large Magellanic Cloud \citep[LMC, ][]{Gaensler05}, the Small Magellanic Cloud \citep[SMC, ][]{Mao08}, and many additional examples such as NGC~4449 \citep{Chyzy00}, NGC~ 1569 \citep{Kepley10}, NGC~6822 \citep{Chyzy03}, IC~10 \citep{Chyzy03, Heesen11} and NGC~ 4214 \citep{Kepley11}. \citet{Chyzy11} pursued a dedicated investigation of radio emission and magnetic fields in an unbiased sample of $12$ Local Group (LG) irregular and dwarf irregular galaxies yielding both detections and upper limits, while \citet{Roychowdhury12} employed the stacking technique to improve the sensitivity in the radio and to probe average properties of the radio emission for the faintest end of dwarf galaxies. A central result of both studies is that the magnetic fields in dwarf galaxies are about three times weaker than in normal spirals, with a typical field strength of $<4.2\pm1.8$~$\mu$G as given by \citet{Chyzy11}. Both the detections and upper limits are consistent with the assumption that local dwarf galaxies lie on the far-infrared - radio correlation, with a typical scaling of the magnetic field strength $B$ with the star formation surface density $\Sigma_{\rm SFR}^{1/3}$. The correlation between the far-infrared and radio fluxes was originally observed by \citet{Kruit73b, Kruit73c, Kruit73a}. Subsequent investigations have been pursued by \citet{deJong85} and \citet{Helou85}, while the interpretation in terms of calorimeter models was pursued by \citet{Volk89}. \citet{Niklas97b} proposed a detailed scenario in which the far-infrared - radio correlation emerges due to a relation between the magnetic field strength, the gas surface density and the star formation rate. In particular, it is well-known that the gas surface density is strongly correlated to star formation activity, as reflected in the Kennicutt-Schmidt relation \citep{Schmidt59, Kennicutt98, Kennicutt08, Walter08, Bigiel11, Kennicutt12}. Massive stars emit UV radiation absorbed by dust grains, and re-emitted in the infrared and far-infrared. In addition, the supernova explosions of massive stars inject both cosmic rays and turbulence into the interstellar medium. Such turbulence efficiently amplifies magnetic field via the small-scale dynamo \citep{Kazantsev68, Subramanian99, Scheko02, Schober12b, Schleicher13, FederrathPRL, Grete15}, and as a result, the feedback from star formation provides the relevant ingredients to drive the radio emission \citep[see e.g.][]{Groves03, Schleicher13b}. The potential validity of the far-infrared - radio correlation even in dwarf galaxies was already suggested by \citet{Bell03}, as both the dust content in the dwarfs and the efficiency of non-thermal radio emission may decrease in a similar amount towards lower star formation rates. At that time, the correlation had been observationally established only for nearby spiral galaxies, including a sample of 1809 galaxies probed by \citet{Yun01}, for which the correlation has been confirmed over 5 orders of magnitude in luminosity. More recent work by \citet{Lacki10} has shown that in fact both the escape of UV photons and cosmic rays from the galaxy may be correlated to the characteristic surface densities, and contributes to our overall understanding of the observations. It is worth noting that these correlations do not only hold on a global scale, but have further been confirmed within the galaxies via dedicated investigations \citep{Dumas11, Taba13, Heesen14}. Recent studies by \citet{Murphy09}, \citet{Ivison10a}, \citet{Jarvis10}, \citet{Sargent10} and \citet{Casey12} in fact provide evidence that the far-infrared - radio correlation holds at least until redshifts of $z\sim2$, and it also holds in the context of galaxy mergers \citep{Drzazga11}. {In addition, work by \citet{Miettinen15} shows that the radio emitting region is more extended than the infrared emitting region at least in some cases. The latter is potentially consistent with Taffy-like systems \citep{Condon02}, mergers \citep{Murphy13} or systems undergoing tidal interactions \citep{Donevski15}, though we suggest that tidal tails as in the Antennae galaxies \citep{Chyzy04} may be the more frequent scenario. In the context of mergers, previous studies, e.g. \citet{Drzazga11} have preferentially considered the impact of magnetic fields, while \citet{Lisenfeld10} pointed out the potential importance of particle acceleration. The latter requires rather high Mach numbers, which may not be available in the interstellar medium \citep{Guo14a}, or the Firehose instability, if the acceleration only concerns the electrons \citep{Guo14b}. In that case, a high plasma beta would be required, while current estimates indicate that it may be rather small \citep{Beck15}.} The radio-infrared correlation holds for thermal and non-thermal (synchrotron) radio emission, though with different slopes. Thermal radio emission may dominate in dwarf galaxies \citep{Roychowdhury12} and hence mask the relation between non-thermal and infrared emission. The interpretation requires a careful separation of both emission components, e.g. with the help of spectral index data, which is non-trivial and often not possible. Alternatively, the thermal radio emission is assumed to be linearly proportional to the infrared emission, although not all infrared emission is directly related to UV radiation from young stars. Only H$\alpha$ emission can be safely assumed to be proportional to radio thermal emission, if extinction is properly corrected. On theoretical grounds, it is expected that magnetic fields can be efficiently ampified even in small systems and at high redshift, due to the efficient amplification by turbulence \citep{Arshakian09, Wang09, Schleicher10c, Souza10, Latif13, Schober13}, and it is thus conceivable that the far-infrared - radio correlation will be in place early on. As pointed out by \citet{Murphy09}, a potential difficulty is however the increasing strength of the cosmic microwave background at high redshift, enhancing the inverse Compton emission and providing an additional loss mechanism for the cosmic ray electrons. It is thus conceivable that the latter may lead to a modification or a breakdown of the correlation at very high redshift due to differences in the energy loss mechanisms of cosmic rays \citep{Lacki10b, Schleicher13b, Schober15}. In this paper, we explore whether a correlation between the far-infrared and radio emission can still be expected in dwarf galaxies, in particular in the limit of star formation {surface densities} below $0.1$~M$_\odot$~kpc$^{-2}$~yr$^{-1}$. In this respect, one needs to examine a couple of relevant issues. A first one concerns the behavior of star formation itself. In the sample explored by \citet{Chyzy11}, the Kennicutt-Schmidt relation appeared to hold even in the dwarf galaxy regime, while \citet{Roychowdhury09} reported potential deviations towards lower star formation rates. Through dedicated spatially resolved investigations of star formation and the HI-dominated gas both in nearby spirals and dwarf irregular galaxies using the THINGS \citep{Walter08} and FIGGS \citep{Begum08} survey, \citet{Roychowdhury15} have pursued a detailed comparison the the Kennicutt-Schmidt relation in both regimes, finding in particular that there is no dependence on the metallicity of the gas. It is therefore conceivable that this relation will hold in the dwarf galaxy regime, even though the scatter may potentially increase towards lower gas masses. An important difference compared to spiral galaxies is that the rotation in local dwarfs may be substantially reduced, implying slow or more chaotic rotation with a low differential rotation \citep{Chyzy03}. In the VLA-ANGST survey of 35 nearby dwarf galaxies, the lowest detected rotation velocities have been of the order $20$~km/s. Through the Westerbork HI Survey of Spiral and Irregular Galaxies (WHISP), rotation curves have been measured for a sample of 62 galaxies, finding typical rotational velocities of $20-80$~km/s on scales of a few disk scale lengths \citep{Swaters09}. In general, a strong variation is found from dwarf to dwarf, and in some cases like NGC4449 \citep{Theis01} or IC~10 \citep{Ashley14}, the rotation curves have been explained by tidal interactions of disks with other dwarfs. The implications of the large diversity of dwarf galaxy rotation curves has therefore been discussed also in recent studies \citep{Oman15}. In the radio regime, additional differences have been found compared to the typical conditions in local spirals, whose main properties were already described by \citet{Condon92}. In particular, while the fraction of thermal radio emission corresponds to about $8\%$ for local spirals \citep{Murphy06}, a non-thermal fraction of $\sim50\%$ has been reported recently by \citet{Roychowdhury12} in the dwarf galaxy regime. We will in fact show in this paper that such a behavior arises rather naturally due to the non-linear nature of the far-infrared - radio correlation. The structure of this paper is as follows. In section~\ref{model}, we outline our overall modelling framework, which is employed to derive characteristic timescales both for dynamical processes within the dwarfs as well as the timescales for thermal and radio emission. These are employed to derive critical star formation surface densities, which are required for these processes to be maintained in a steady fashion. In section~\ref{slope}, these results are employed to distinguish between four characteristic regimes of radio emission in the dwarf galaxies, and we discuss in particular the expected slope and the breakdown of the correlation at very low star formation rates. A discussion with our main conclusions is presented in section~\ref{discussions}.

\label{discussions} In this paper, we have discussed the far-infrared - radio correlation in dwarf galaxies and its potential evolution in the regime of low star formation rates. As a starting point, we have assessed under which conditions a correlation between the magnetic field strength and the star formation rate can be maintained through the continuous injection of turbulence, leading to the definition of a critical star formation rate required for the injection of a sufficient amount of turbulent energy (Eq.~\ref{SFRcrit}). The latter will ensure magnetic field amplification via the small-scale dynamo \citep[e.g.][]{Kazantsev68, Scheko02, Schober12b}, which happens on short timescales and effectively ensures that the magnetic field strength is coupled to the star formation rate. {Considering a typical size of the star-forming region of $1$~kpc, this relation will break down at critical star formation surface densities of $10^{-5}-10^{-6}$~M$_\odot$~kpc$^{-2}$~yr$^{-1}$ depending on the amount of rotation.} We also derived a critical star formation rate to ensure continuous thermal emission in the galaxy, which results from the requirement that the timescale of massive star formation should be smaller than the typical lifetime of massive stars. This criterion has been developed in Eq.~\ref{thermal}, even though our comparison has shown that it is likely not relevant in practice. {Assuming a star-forming region of $1$~kpc and a $5\%$ fraction of massive stars, the resulting star formation surface density corresponds to $\sim10^{-6}$~M$_\odot$~kpc$^{-2}$~yr$^{-1}$. } We further explored the dominant loss mechanisms for cosmic rays in the galaxy, which have a major impact on the non-thermal radio emission and its dependence on the star formation rate. In particular, we derived a critical star formation surface density for the importance of cosmic ray diffusion losses given by Eq.~(\ref{SFRdiff}), below which the injection time of cosmic rays becomes larger than the timescale for diffusion losses, thus strongly {reducing} the amount of cosmic rays in the galaxy. {Assuming a star-forming region of $1$~kpc, these effects become relevant for star formation surface densities of $10^{-4}-10^{-6}$~M$_\odot$~kpc$^{-2}$~yr$^{-1}$, depending on the scale height of the disk. Overall, the critical star formation surface density for this transition thus appears higher than the corresponding transition for the thermal emission. We have compared the cosmic-ray diffusion losses} with the losses due to cosmic winds, adopting both the thermally-driven wind model by \citet{Shu05} as well as a generalized model accounting for the effect of cosmic rays and magnetic fields as potential driving agents. Depending on the amount of rotation in the galaxy, we have found that the losses from winds can become relevant at somewhat higher or lower star formation rates compared to the cosmic ray diffusion processes, but do not change dramatically the details of the transition. From a comparison of these results, we have shown that the critical star formation surface densities for the maintainance of turbulence (Eq.~\ref{SFRcrit}) and for the relevance of cosmic-ray diffusion (Eq.~\ref{SFRdiff}) are most relevant for introducing a potential breakdown in the far-infrared - radio correlation. Which of the transitions occurs first depends on the rotation rate of the dwarf and the scale height of the disk, and can be expressed in terms of a critical rotation rate given in Eq.~(\ref{omegadiffdisk}), {corresponding to rotation period of $1.5\times10^7$~yrs for a galaxy with scale height of $1$~kpc. We note in particular that cosmic-ray diffusion losses become relevant before the breakdown of the magnetic field - star formation surface density relation in the case of low rotation. Whether the rotation period is above or below the critical value thus regulates the further evolution of the far-infrared - radio correlation in the regime of low star formation rates.} We have determined four different regimes employing the conditions derived above. In the limit of high star formation rates above the critical star formation surface densities {mentioned above}, we have shown that the ratio of non-thermal radio to far-infrared / infrared luminosity should scale as $\Sigma_{\rm SFR}^{1/3}$, as we expect that the cosmic ray abundance is limited by the strength of the magnetic field, effectively enforcing equipartition, implying a scaling of the cosmic-ray abundance with $B^2$ and of the overall non-thermal emission as $B^4$ or $\Sigma_{\rm SFR}^{4/3}$, {typically above a star formation surface density of $\sim10^{-6}$~M$_\odot$~kpc$^{-2}$~yr$^{-1}$ for a star-forming region of $1$~kpc.} For lower star formation rates in the limit of low rotation (below the critical value in Eq.~\ref{omegadiffdisk}), the correlation between magnetic field strength and star formation will still be in place, but cosmic ray diffusion losses will start becoming efficient. As a result, the cosmic ray abundance in galaxies will no longer reach equipartition with the magnetic field strength, but is instead expected to be proportional to the injection rate, i.e. the rate of star formation. In this regime, it can then be shown that the ratio of non-thermal radio to far-infrared / infrared luminosity will scale as $\Sigma_{\rm SFR}^{2/3}$, implying a steeper decrease with decreasing star formation rate. Such steepening may be expected below typical star formation surface densities of $3\times10^{-5}$~M$_\odot$~kpc$^{-2}$~yr$^{-1}$, with some dependence both on the size of the star-forming region and the scale height of the disk. On the other hand, in the regime of low star formation rates and {rotation rates above the before-mentioned critical value}, the cosmic ray diffusion losses are not yet significant, but the correlation between the magnetic field strength and the star formation rate is expected to break down, as a steady injection of turbulence cannot be maintained, leading to the dissipation both of turbulence and the magnetic fields. As a result, one may expect significant non-thermal emission only in those sources with recent injection events, and otherwise no significant emission in the vast majority of sources. Finally, for star formation rates both below the critical rates for turbulence driving and cosmic-ray diffusion, it is clear that the correlation is broken down due to both mechanism. Overall, we therefore expect modifications of the far-infrared - radio correlation in the regime of low star formation rates, which will be particularly relevant for future observations with the Square Kilometre Array (SKA)\footnote{Website SKA: https://www.skatelescope.org/} and its Key Science Project on ''The Origin and Evolution of Cosmic Magnetism''. Its precursors and pathfinders such as LOFAR\footnote{Website LOFAR: http://www.lofar.org/}, MeerKAT\footnote{Website MeerKAT: http://www.ska.ac.za/meerkat/} and ASKAP\footnote{Website ASKAP: http://www.atnf.csiro.au/projects/askap/} are in fact already in the process of further probing magnetic field structures in local dwarf galaxies, as in the LOFAR Key Science Project on ''Cosmic Magnetism of the Nearby Universe'' \citep{Beck13LOFAR}. Such dedicated investigations will be particularly valuable to determine the frontiers of cosmic magnetism and the origin of magnetic fields in galaxies.

1607.00094

1607

1607.05487_arXiv.txt

Molecular cloud structure is regulated by stellar feedback in various forms. Two of the most important feedback processes are UV photoionisation and supernovae from massive stars. However, the precise response of the cloud to these processes, and the interaction between them, remains an open question. In particular, we wish to know under which conditions the cloud can be dispersed by feedback, which in turn can give us hints as to how feedback regulates the star formation inside the cloud. We perform a suite of radiative magnetohydrodynamic simulations of a $10^5$ solar mass cloud with embedded sources of ionising radiation and supernovae, including multiple supernovae and a hypernova model. A UV source corresponding to 10\% of the mass of the cloud is required to disperse the cloud, suggesting that the star formation efficiency should be on the order of 10\%. A single supernova is unable to significantly affect the evolution of the cloud. However, energetic hypernovae and multiple supernovae are able to add significant quantities of momentum to the cloud, approximately $10^{43}$ g cm/s of momentum per $10^{51}$ ergs of supernova energy. This is on the lower range of estimates in other works, since dense gas clumps that remain embedded inside the HII region cause rapid cooling in the supernova blast. We argue that supernovae alone are unable to regulate star formation in molecular clouds, and that strong pre-supernova feedback is required to allow supernova blastwaves to propagate efficiently into the interstellar medium.

\label{introduction} Massive stars release large quantities of energy into their environment. They produce protostellar jets, winds, radiation across a wide spectrum and supernovae. The first phase of stellar feedback occurs in dense molecular cloud environments in which the stars are born. In order for the energy from stars to propagate into the wider \ISM, it must first escape this cloud environment, either by destroying the cloud or creating sufficient channels through which the propagating shocks can escape. In the previous paper, \cite{Geen2015b}, we determined a limit at which ionising radiation can escape molecular clouds using both numerical simulations and an analytic model. This model is based on arguments made in \cite{Tremblin2014a,Didelon2015}, which compare models of HII regions expanding into turbulent environments to observed HII regions. These models were constructed using previous analytic theory by \cite{KahnF.D.1954,SpitzerLyman1978,Whitworth1979,Franco1990,Williams1997,Hosokawa2006,Raga2012}. In \cite{Geen2015b} we proposed a limit at which ionising photons are able to destroy their host cloud. This extends the argument of \cite{Dale2012}, who consider the case where the ionisation front cannot expand beyond the initial \Stromgren radius. We argue that if a calculated ``stall'' radius \citep{Draine1991} is smaller than the radius of the cloud, the ionisation front cannot escape the cloud. This in turn sets the ability for ionising radiation to regulate the environments in which stars form, and determines whether ionising radiation can suppress the star formation rate of molecular clouds. Massive stars typically end their lives as supernovae \citep[for estimates of which stars become supernovae, see, e.g. ][]{Heger2003}. The evolution of the supernova remnant depends on the environment into which it expands. Understanding the momentum deposition from supernovae in star-forming environments is crucial to understanding processes in galaxies as a whole. Sub-grid models by, e.g., \cite{Hopkins2014,Kimm2015a} attempt to correct for a lack of numerical resolution by depositing a pre-calculated quantity of momentum around the supernova if the resolution is insufficient to resolve the blastwave properly \citep[see also ][ for a study of numerical limits on resolving supernova blastwaves]{Kim2015}. Analytic and 1D simulation work by \cite{Chevalier1974,Cioffi1988,Draine1991,Thornton1998,Haid2016} provides insights into this process, with simulations of supernova blastwaves by \cite{Iffrig2015,Kim2015,Martizzi2015,Kortgen2016} extending this to more complex environments using 3D numerical simulations. Supernovae shock against the surrounding medium, expanding adiabatically \citep{Sedov1946}. Eventually they reach a point where they begin to lose a significant fraction of their energy to radiative cooling \citep[see estimates by ][]{Cox1972}. After a longer period of time, the supernova remnant begins to merge with the surrounding medium \citep{Cioffi1988}. Pre-supernova feedback as either stellar winds or ionising radiation has been found in simulations to enhance the final energy and momentum of the supernova remnant by injecting additional momentum and reducing the density of the environment into which the supernova occurs \citep{Dwarkadas2007,Fierlinger2015,Geen2015a}. \cite{Rogers2013,Walch2015} have had some success in driving outflows in simulations of molecular clouds with both supernova and pre-supernova stellar feedback. However, \cite{Draine1991} suggests that if the medium is sufficiently turbulent, the HII region will re-collapse before the supernova occurs, depending on the mass of the progenitor and the density of the surrounding medium. \cite{Krause2016} find that stellar feedback in very massive extragalactic clouds is ineffective at reducing the star formation efficiency. In this paper we explore the competition between pre-supernova ionising feedback and turbulence in molecular clouds, and the resulting evolution of the supernova remnant as it expands into the environment resulting from this competition. We simulate ionising radiation and supernovae in a turbulent cloud using \textsc{RAMSES-RT} \citep{Teyssier:2002p533,Fromang2006,Rosdahl2013}. The cloud is $10^5$ \Msolar, ten times more massive than the one studied in the previous paper. We choose this cloud mass because the slope of the cloud mass function ($\mathrm{d}N/\mathrm{d}M_c = M_c^{-1.7}$) means that more mass is expected to be found in clouds above $10^5$ \Msolar \citep[see review by ][]{Hennebelle2012}. Therefore, most of the stars in our Galaxy are expected to form in these clouds. It is thus important to study these objects if we wish to understand how feedback from massive stars interacts with both the host cloud and the wider Galactic \ISM. We begin in Section \ref{methods} by presenting the simulations performed. We then extend the analysis of HII regions in the previous paper to a more massive cloud in Section \ref{beforesn}, and produce simple models that describe the time-dependence of the evolution of the ionisation front. In Section \ref{aftersn}, we analyse the results of simulations that introduce a supernova into the cloud and HII region after the source of UV photons is extinguished, and how this compares to previous analytic and numerical theory. We extend the single supernova scenario to more energetic events in Section \ref{moresne}.

\label{conclusions} We perform a series of radiative magnetohydroynamic simulations of a $10^5$ \Msolar turbulent molecular cloud with embedded sources of ionising radiation and a supernova. We compare the results of these simulations to analytic models given in the previous paper and find good agreement provided a good fit is found for the density field of the cloud. We study the time evolution of these analytic models, finding that the limit at which the ionisation front stalls (i.e. stops expanding) is reached over approximately one free-fall time in the cloud. This introduces a competition between the two processes of star formation and stellar feedback. The environment that the supernova blastwave expands into depends strongly on the emission rate of ionising photons from the cluster beforehand. An emission rate of $10^{51}$ photons per second, roughly equivalent to a cluster of $10^4$ \Msolar or 10\% of the total cloud mass, is capable of disrupting the cloud, though dense clumps remain. If the number of stars formed is much lower, the ionising photons will not be able to destroy the cloud and the supernova will transfer its momentum to the dense cloud gas rather than the diffuse interstellar medium. This suggests that for this cloud a star formation efficiency of approximately 10\% is expected if the main source of feedback is from ionising photons. The position of the source of ionising photons and the supernova is highly important due to the effect of gas density on photon recombination and radiative cooling in general. We inject supernovae as a single thermal pulse of $10^{51}$ ergs. We also perform simulations with ten supernovae of the same energy 0.1 Myr apart or one hypernova of $10^{52}$ ergs. The resulting total momentum from our supernovae is roughly $10^{43}$ g cm/s per $10^{51}$ ergs of injected energy. This is at the low end of the values given in other works, but it is not inconsistent provided the early evolution of the blastwave occurs in gas at around $10^4$ \atcc or higher. We argue that flows of dense, turbulent gas inside the cloud are capable of reducing the momentum added to the \ISM by supernovae as long as dense clumps remain embedded within the cloud at the time the supernova occurs. Most of the momentum from a single supernova is deposited into the dense gas rather than as fast, hot, diffuse flows, except in cases where the ionising photons have swept away most of the cloud. We speculate that supernovae will occur too late to prevent the bulk of star formation in the cloud, but sufficient supernovae will be capable of expelling the remaining gas and allowing future supernovae to drive shocks into the interstellar medium.

1607.05487

1607

1607.04024_arXiv.txt

{The existence of satellite galaxy planes poses a major challenge for the standard picture of structure formation with non-baryonic dark matter. Recently Tully et al.~(2015) reported the discovery of two almost parallel planes in the nearby Cen\,A group using mostly high-mass galaxies (M$_B$\,<\,-10\,mag) in their analysis.} {Our team detected a large number of new group member candidates in the Cen\,A group. This dwarf galaxy sample, combined with other recent results from the literature, enables us to test the galaxy distribution in the direction of the Cen\,A group and to determine the statistical significance of the geometric alignment.} {Taking advantage of the fact that the two galaxy planes lie almost edge-on along the line of sight, the newly found group members can be assigned relative to the two planes. We used various statistical methods to test whether the distribution of galaxies follows a single normal distribution or shows evidence of bimodality as has been reported earlier.} {We confirm that the data used for the Tully et al.~(2015) study support the picture of a bimodal structure. When the new galaxy samples are included, however, the gap between the two galaxy planes is closing and the significance level of the bimodality is reduced. Instead, the plane that contains Cen\,A becomes more prominent. % } {We found evidence that the galaxy system around Cen\,A is made up of only one plane of satellites. This plane is almost orthogonal to the dust plane of Cen\,A. Accurate distances to the new dwarf galaxies will be required to measure the precise 3D distribution of the galaxies around Cen\,A.}

The mere abundance and spatial distribution of faint dwarf galaxies provide a powerful testbed for dark matter and structure formation models on Mpc and galaxy scales. The standard picture of structure formation with dark matter is heavily challenged by the highly asymmetric features found in the distributions of dwarf galaxies in the Local Group, {which was first noted by \cite{2005A&A...431..517K}}. There is the vast polar structure \citep[VPOS;][]{2012MNRAS.423.1109P,2015MNRAS.453.1047P,2016MNRAS.456..448P}, a thin (rms height $\approx$ 30 kpc) highly inclined, co-rotating substructure of faint satellite galaxies, young globular clusters, and stellar streams, spreading in Galactocentric distance between 10 and 250\,kpc. A similar feature was found in the Andromeda galaxy surroundings, the so-called Great Plane of Andromeda \citep[GPoA;][]{2006AJ....131.1405K,2007MNRAS.374.1125M,2013Natur.493...62I}. On a {slightly} larger scale, two dwarf galaxy planes containing all but one of the 15 non-satellite galaxies have been identified in the Local Group \citep{2013MNRAS.435.1928P}. Such extreme satellite planes are found in only $<0.1\%$ of simulated systems in cosmological simulations \citep[e.g.,][]{2014ApJ...784L...6I,2014MNRAS.442.2362P}, making them difficult to accommodate in a standard $\Lambda$CDM scenario. An alternative analysis of cosmological simulations, based on including the look elsewhere effect but ignoring observational uncertainties and the non-satellite planes \citep{2015MNRAS.452.3838C}, finds that only about 1 per cent of Local Group-equivalent environments should host similarly extreme satellite structures. The fundamental question arises whether the relative sparseness and asymmetric distribution of low-mass dwarf galaxies encountered in the Local Group is a statistical outlier or a common phenomenon in the local universe. \citet{2014Natur.511..563I} have approached this question with a statistical study of velocity anticorrelations among pairs of satellite galaxies on opposite sides of their host, using data from the SDSS survey. They found a strong excess in anticorrelated velocities, which is consistent with co-orbiting planes, but their conclusions were based on a small sample of only 22 systems, and have been challenged since \citep[][{ but see \citealt{2015ApJ...805...67I}}]{2015MNRAS.453.3839P,2015MNRAS.449.2576C}. A different approach to address the question of satellite planes is to extend searches for such structures to satellite populations in other galaxy groups in the nearby universe. For example, \citet{2013AJ....146..126C} have found that dwarf spheroidal galaxies in the M81 group lie in a flattened distribution. Most recently, \cite{2015ApJ...802L..25T}, hereafter T15, reported evidence for a double-planar structure around Centaurus\,A (Cen\,A, NGC5128) in the nearby Centaurus group of galaxies, with properties reminicent of the two Local Group dwarf galaxy planes \citep{2013MNRAS.435.1928P}. Furthermore, \cite{2015MNRAS.452.1052L} found that the Local Group and the Cen\,A group reside in a filament stretched by the Virgo Cluster and compressed by the Local Void and smaller voids. Four out of five planes of satellites (including the two galaxy planes around Cen\,A) align with this filament with the normal vectors pointing in {the} direction of the Local Void and the planes almost parallel to the minor axis of the filament. These results demonstrate that systematic studies of the spatial distribution of low luminosity galaxies in nearby groups can provide important observational constraints for further testing of structure formation models outside of the Local Group. The aim of the present study is to test how recently discovered dwarf galaxies in the Centaurus group \citep{2014ApJ...793L...7S, 2014ApJ...795L..35C,2016ApJ...823...19C, 2015A&A...583A..79M,2016arXiv160504130M} are distributed in the double planar structure reported by Tully and collaborators. Comparing the footprint of the two planes, 20 of these dwarf candidates are located in that region of the sky. It is important to note that this analysis is feasible without distance information because of special geometry: the normal vectors of the two planes are almost parallel and perpendicular to the line of sight. Consequently, any distance uncertainties for the new galaxies does not move them in or out of the two planes. In section 2 we present the four different galaxy samples used in our analysis. In section 3 we show the transformation between the equatorial coordinate system to the Cen\,A reference frame and fit the two planes of satellites. Section 4 follows a discussion of the geometrical alignment of the planes and the distribution of galaxies. We test the statistical significance of the two planes in Section 5, followed by Section 6 where we summarize the results.

The discovery of a large number of new dwarf galaxies in the nearby Centaurus group \citep{2014ApJ...795L..35C,2016ApJ...823...19C,2015A&A...583A..79M,2016arXiv160504130M} opened the opportunity to conduct a significance test of the two planes of satellites reported by \cite{2015ApJ...802L..25T}. While this normally requires follow-up observations to measure galaxy distances, in the case of the Cen\,A subgroup, owing to the special geometric situation, one can take advantage of the fact that the line of sight from our vantage point runs along the postulated planes of satellites. Therefore, galaxies even without distance measurements can be used for the test, as distance uncertainties move galaxies along or parallel to the planes. In other words, distance uncertainties produce only little crosstalk between the planes and the space around them. This allows us to include all Cen\,A member candidates in the analysis, which doubles the sample size from 30 to 59. Sampling galaxy positions along an edge-on projection of the planes, we studied the distribution of Cen\,A subgroup members by two different techniques: a histogram and a more sophisticated adaptive kernel density estimation. A gap, or dip in the distribution marking the two planes, is visible in both representations of the data. However, it is also very evident that with the inclusion of the new galaxy data the gap between the two planes starts to be filled, raising the conjecture that the two plane scenario around Cen\,A might be an artifact of low number statistics. To put this under statistical scrutiny, we performed an Anderson-Darling test and a Hartigan dip test to see whether the distribution is in agreement with a unimodal normal distribution. We find that both tests fail with the sample used by \cite{2015ApJ...802L..25T}, rejecting the unimodal normal hypothesis for that orginal sample. This result is consistent with their finding that the satellites can be split into two planes. However, with the addition of the new dwarf members and candidates of the Cen\,A subgroup, the deviation from a normal distribution loses statistical significance in the sense of the two tests applied. Hence it is now conceivable that the satellites follow a normal distribution and the gap between the two planes is indeed an artefact from small number statistics. We performed Monte Carlos simulations to further strengthen these results by taking distance uncertainties into account and find that the results change only marginally. Given that distance measurements for the candidate galaxies are unavailable at the moment, it is theoretically possible to allocate 17 out of 25 galaxies to one of the two planes each by moving them along the line of sight. That means that technically the existence of the two planes cannot be completely ruled out at this point. Only distance measurements will tell whether the candidates lie on or near the planes. All we can say is that the case for two planes around Cen\,A is weakened by including the currently available data for 29 new dwarfs and dwarf candidates. At first glance, another secondary result of our testing is that in parallel with the weakening of the significance for bimodality, the galaxy population in plane 1 has now become dominant. This is not surprising: Cen\,A lies closer to plane 1 in projection (CaZ=0 in Fig.\,5), such that any newly discovered satellite galaxy in a distribution that is radially concentrated on Cen\,A will necessarily result in an additional member of plane 1. However, this amassing of Cen\,A satellites in the proposed plane 1 is important, as the planarity of Plane 1 has not been destroyed by adding the candidates. The alignment of the tidally disrupted CenA-MM-Dw3 with plane 1 (see Figs.\,1 and 2) is intriguing too. The Cen\,A plane 1 is certainly a good candidate analog of the local thin planes detected around the Milky Way (VPOS) and the Andromeda galaxy (GPOA). Future distance measurements for the many new Cen\,A subgroup member candidates will be able to tell how far the analogy will take us.

1607.04024

1607

1607.06810_arXiv.txt

We report nine years of optical spectroscopy of the metal-polluted white dwarf SDSS\,J104341.53+085558.2, which presents morphological variations of the line profiles of the 8600\,\AA\ Ca\,{\textsc{ii}} triplet emission from the gaseous component of its debris disc. Similar changes in the shape of the Ca\,{\textsc{ii}} triplet have also been observed in two other systems that host a gaseous disc, and are likely related to the same mechanism. We report the Mg, Si, and Ca abundances of the debris detected in the photosphere of SDSS\,J1043+0855, place upper limits on O and Fe, and derive an accretion rate of $(2.5$ - $12)\,\times\,10^{8}$\,gs$^{-1}$, consistent with those found in other systems with detected debris discs. The Mg/Si ratio and the upper limit on the Fe/Si ratio of the accreted material broadly agree with those found for the crust of the Earth. We also review the range of variability observed among white dwarfs with planetary debris discs.

The detection of metal pollution in white dwarf atmospheres provides strong evidence that 25-50\,per\,cent of white dwarfs host remnants of planetary systems \citep{zuckermanetal03-1, koesteretal14-1, koester+kepler15-1}. The survival of planets through the post main-sequence evolution of their host star is further supported by the detection of more than 35 dusty debris discs at white dwarfs with metal pollution \citep{kilicetal06-1, farihietal08-1, farihietal09-1, juraetal09-1, debesetal11-2, hoardetal13-1, bergforsetal14-1, rocchettoetal15-1}, and is also predicted by theoretical studies \citep{villaver+livio07-1, veras+gaensicke15-1}. Debris discs around white dwarfs are thought to be produced by the tidal disruption of asteroids (or comets) scattered onto a highly eccentric orbit by planets in the system \citep{grahametal90-1, jura03-1}. While recent simulations have begun to explore the formation and evolution of these discs; from the initial disruption of the planetesimal and the shrinking of its orbit \citep{debesetal12-1,verasetal14-1, verasetal15-1}, to the dynamics and interactions of the debris within the disc \citep{rafikov11-2, metzgeretal12-1}, many aspects remain poorly understood. Over the last decade, gaseous debris discs have been detected around eight white dwarfs, all of which also host circumstellar dust and have metal polluted photospheres \citep{gaensickeetal06-3, gaensickeetal07-1, gaensickeetal08-1, gaensicke11-1, farihietal12-1, melisetal12-1, wilsonetal14-1, guoetal15-1}. These systems were identified by the detection of Ca\,{\textsc{ii}} 8498.02\,\AA, 8542.09\,\AA, 8662.14\,\AA\ emission lines (henceforth the Ca\,{\textsc{ii}} triplet, air wavelengths are given). \begin{figure*} \centerline{\includegraphics[width=2\columnwidth]{fig1.eps}} \caption{\label{f-1043_plot} X-Shooter spectra of SDSS\,J1043+0855 (grey, obtained in January and May 2011) together with a model atmosphere fit (blue) using the atmospheric parameters listed in Table\,\ref{t-wds}. The strongest emission and absorption lines have been labelled (note that the Fe\,{\textsc{ii}} feature is in emission).} \end{figure*} \begin{table*} \centering \caption{Log of observations of SDSS\,J1043+0855. $^{a}$ Different exposure times for the individual X-Shooter arms (UVB / VIS). We did not use data collected by the NIR arm of X-Shooter as the signal-to-noise ratio was too poor. \label{t-dates}} \begin{tabular}{rrrrr} \hline Date & Telescope/Instrument & Wavelength Range [\AA] & Resolution [\AA] & Exposure Time [s]\,$^{a}$ \\ \hline 2003 April 05 & SDSS & 3800 -- 9200 & 2.9 & 2900\\ 2007 February 03 & WHT/ISIS & 7400 -- 9200 & 2.0 & 4800\\ 2009 February 16 & WHT/ISIS & 6000 -- 8900 & 3.7 & 3950\\ 2010 April 22 & WHT/ISIS & 8100 -- 8850 & 1.1 & 7200\\ 2011 January 29 & VLT/X-Shooter & 2990 -- 10400 & 1.12 & 2950 / 2840\\ 2011 May 30 & VLT/X-Shooter & 2990 -- 10400 & 1.15 & 2950 / 2840\\ 2012 January 03 & SDSS & 3602 -- 10353 & 3.2 & 2702\\ \hline \end{tabular} \end{table*} So far only four of these eight gas disc systems have multi-epoch spectroscopy over time scales of years, three of which show variability in the Ca\,{\textsc{ii}} emission of the gaseous disc; either as significant changes in the morphology of the Ca\,{\textsc{ii}} triplet \citep{wilsonetal15-1, manseretal16-1}, or as a decrease in strength of the lines \citep{wilsonetal14-1}. It is thought that the gaseous components of the debris discs in these systems are tracers of dynamic activity \citep{gaensickeetal08-1, wilsonetal14-1, wilsonetal15-1, manseretal16-1}. The gaseous disc around SDSS\,J104341.53+085558.2 (henceforth SDSS\,J1043+0855) was discovered by \cite{gaensickeetal07-1} via the detection of Ca\,{\textsc{ii}} triplet emission. An infrared excess was detected by \cite{melisetal10-1} and \cite{brinkworthetal12-1}, confirming the presence of a dusty disc in SDSS\,J1043+0855. Here, we report nine years of spectroscopic observations of SDSS\,J1043+0855 that reveal a change in the morphology of the Ca\,{\textsc{ii}} triplet, similar to those seen in other gaseous disc systems (SDSS\,J122859.93+104032.9 and SDSS\,J084539.17+225728.0, henceforth SDSS\,J1228+1040 and SDSS\,J0845+2257, \citealt{wilsonetal15-1, manseretal16-1}). We also present the accretion rates of the debris onto the white dwarf and the metal abundances of the debris.

\label{sec:conc} We report here the morphological variability of the Ca\,{\textsc{ii}} triplet in SDSS\,J1043+0855 on a time scale of nine years. The evolution of the Ca\,{\textsc{ii}} triplet reported here is similar to that of two other systems, SDSS\,J1228+1040 and SDSS\,J0845+2257. We have also analysed the optical spectra of SDSS\,J1043+0855 to determine its stellar parameters and the photospheric metal abundance. The Mg/Si and (upper limit to the) Fe/Si ratios of the planetary debris that has been accreted onto the white dwarf are broadly consistent with those of the crust of the Earth. The recent detection of the 'real time' disruption of a planetesimal at WD\,1145+017, along with the dynamical evolution seen at the gaseous discs SDSS\,J1043+0855, SDSS\,J1228+1040, SDSS\,J1617+1620, and SDSS\,J0845+2257 reveals that variability at planetary systems around white dwarfs is more common than was initially thought. Additional spectroscopic and photometric monitoring of all the gaseous discs known so far is key to developing a more detailed understanding of the dynamical processes present in planetary debris discs at white dwarfs.

1607.06810

1607

1607.03929_arXiv.txt

We present 3D magnetohydrodynamic numerical simulations of the adiabatic interaction of a shock with a dense, filamentary cloud. We investigate the effects of various filament lengths and orientations on the interaction using different orientations of the magnetic field, and vary the Mach number of the shock, the density contrast of the filament $\chi$, and the plasma beta, in order to determine their effect on the evolution and lifetime of the filament. We find that in a parallel magnetic field filaments have longer lifetimes if they are orientated more ``broadside'' to the shock front, and that an increase in $\chi$ hastens the destruction of the cloud, in terms of the modified cloud-crushing timescale, $t_{cs}$. The combination of a mild shock and a perpendicular or oblique field provides the best condition for extending the life of the filament, with some filaments able to survive almost indefinitely since they are cocooned by the magnetic field. A high value for $\chi$ does not initiate large turbulent instabilities in either the perpendicular or oblique field cases but rather draws the filament out into long tendrils which may eventually fragment. In addition, flux ropes are only formed in parallel magnetic fields. The length of the filament is, however, not as important for the evolution and destruction of a filament.

The interstellar medium (ISM) is known to be a highly dynamic and non-uniform entity containing regions of varying temperature and density (see the review paper by \citet{Ferriere01}). Studies of the interaction of hot, high-velocity gas with cooler, dense material (often referred to as ``clouds") are of great interest for a complete understanding of the gas dynamics of the ISM since it is evident that the evolution and morphology of large-scale flows can be determined by the far smaller clouds \citep{Elmegreen04, MacLow04, Scalo04, McKee07, Hennebelle12, Padoan14}. Clouds may either accrete material from, or lose material to, the ambient medium: clouds which are hit by shocks or winds are likely to be destroyed, with such destruction affecting the flow by ``mass-loading" it via processes such as hydrodynamic ablation, whereas clouds may also collapse after being struck by a shock and therefore trigger star formation, thus removing material from the ISM \citep{Elmegreen77, Federrath10, Federrath12}. Shock-cloud interactions have been previously inferred from observations (e.g. \citealt{Baade54}; \citealt{vandenBergh71}) while more recent observations have provided direct evidence, e.g. bow shocks, for shock waves interacting with clouds (e.g. \citealt{Levenson02}). Recently, \textit{Herschel} images have revealed the ubiquitous presence of filamentary structures throughout the ISM in both star-forming and non-star-forming regions (e.g. \citealt{Andre10, Andre14}). There is now a large amount of literature, beginning in the 1970s, concerning the idealised case of a planar adiabatic shock striking an isolated spherical cloud. Numerical studies where the shock Mach number $M$ and cloud density contrast $\chi$ were varied include \citet{Stone92} and \citet{Klein94}. Other studies have reported on the effects of additional processes on the interaction, such as magnetic fields (e.g. \citealt{MacLow94}; \citealt{Shin08}), radiative cooling (e.g. \citealt{Mellema02}; \citealt{Fragile04}; \citealt{Yirak10}) and thermal conduction (e.g. \citealt{Orlando05, Orlando08}). \citet{Pittard09, Pittard10} also explored the turbulent nature of cloud destruction, whilst \citet{Poludnenko02} and \citet{Aluzas12, Aluzas14} investigated the interaction of shocks with multiple clouds, and \citet{vanLoo10} explored the interaction of a weak, radiative shock with a magnetised cloud. The purely hydrodynamic shock-cloud interactions lead to the cloud becoming initially compressed, as the shock strikes it, and over-pressured before the cloud re-expands. The cloud is then destroyed via the growth of dynamical instabilities such as Kelvin-Helmholtz (KH) and Rayleigh-Taylor (RT) instabilities which deposit vorticity at the cloud surface, leading to the mixing of the cloud material with the ambient medium. The interaction is milder at lower shock Mach numbers (e.g. \citealt{Nakamura06}; \citealt{Pittard10}) and more marked differences are observed when the post-shock gas is subsonic with respect to the cloud. The presence of magnetic fields can strongly change the nature of the interaction. 2D axisymmetric simulations have shown that if there is a magnetic field present then the formation of the KH and Richtmyer-Meshkov (RM) instabilities are impeded and the mixing of the cloud with the flow is reduced \citep{MacLow94}. Thus the presence of a magnetic field can prevent the complete destruction of the cloud, allowing it to survive as a coherent structure, as opposed to mixing completely with the ambient flow (as in the field-free case). Furthermore, if the field is parallel to the shock normal a ``flux rope" is formed behind the cloud since the field is preferentially amplified at that point due to shock-focussing. 3D simulations show that when the magnetic field is strong and aligned either perpendicularly or obliquely to the shock normal the cloud takes on a sheet-like appearance at late times and becomes orientated parallel to the post-shock field \citep{Shin08}. A perpendicular field can better deflect the flow around the cloud and reduce mixing, whereas a parallel field allows the cloud to be permeated by the flow and this enhances mixing \citep{Li13}. This effect was also noted in the paper on wind-cloud interactions by \citet{Banda16}, who found that cloud models where the magnetic field component was transverse to the wind direction had higher mixing fractions and velocity component dispersions than models where the field component was aligned with the flow. More recent work has considered the optimum field strength needed to produce cloud fragments which can survive the destructive processes and has found that intermediate-strength fields are most effective, since strong fields prevent compression and weak fields do not insulate the cloud from cooling \citep{Johansson13}. There are very few numerical studies in the current literature which consider interactions involving non-spherical clouds, and (to our knowledge) none which describe the effects of a magnetic field on these interactions. One of the first such studies concerned a shock interacting with a cylindrical cloud of aspect ratio 3:1 \citep{Klein94}. The cloud was orientated along the axis of propagation. \citet{Klein94} used a modified equation for the cloud-crushing time and found their results comparable to those of a spherical cloud; thus they concluded that small changes to the initial shape of the cloud did not alter their main conclusions. Another study \citep{Xu95} focussed on 3D simulations of shock-cloud interactions for clouds with varying morphologies and orientations. Unlike \citet{Klein94}, who assumed a cylindrical cloud aligned in the direction of shock propagation, \citet{Xu95} were able to orientate their cloud of aspect ratio 2:1 in all directions. They found that by modifying the cross-section of the cloud its evolution could be significantly altered depending on the cloud geometry. They also found that, whilst the formation of a vortex ring is a feature of interactions with spherical clouds, a prolate cloud aligned perpendicularly to the shock normal does not form a vortex ring since the interaction of the shock is more complex. Additionally, an aligned cloud was also accelerated to the post-shock flow velocity at a much faster rate than a spherical cloud. In contrast, the evolution of an inclined prolate cloud was substantially different from the aligned cloud: in this case the cloud's inclination caused it to be spun around, drastically altering the development of instabilities. The most recent study, \citet{Pittard16}, investigated shock-filament interactions and studied the formation of turbulent vortices behind the filaments as a result of the shock-filament interaction. They found that varying the filament length and angle of orientation to the shock front significantly changed the nature of the interaction. Filaments orientated at $\theta \lesssim 60^{\circ}$ formed three parallel rolls, whilst filaments orientated sideways-on expanded preferentially along their minor axis and in the direction of shock propagation. Slightly oblique filaments tended to spill the high vorticity flow around the upstream end of the filament. These filaments had longer wakes and were less symmetrical. Highly oblique filaments, in contrast, had a dominant vortex ring at the upstream end of the filament which aided their subsequent fragmentation. The current study extends the purely hydrodynamic work conducted by \citet{Pittard16}. By nature, it represents an idealised scenario before more realistic simulations of filaments are conducted. We investigate the effects that magnetic fields have on shock-filament interactions by varying the Mach number, density contrast, and plasma beta, in addition to varying the orientation and length of the filament, for parallel, perpendicular, and oblique magnetic fields. The outline of this paper is as follows: in Section 2 we introduce our numerical method, initial conditions and the results of a convergence study. In Section 3 we present the results of our simulations. A discussion of the relevance of our work to shock-filament and wind-filament studies is given in Section 4. Section 5 summarises and concludes, and addresses the motivation for further work.

This is the second in a series of papers investigating the interaction between astrophysical shocks and filaments. In this paper, we employed a magnetohydrodynamic code to investigate the evolution and destruction by an adiabatic shock of a filament embedded within a magnetised medium. In comparison to the results from the previous hydrodynamical study of filaments by \citet{Pittard16} we found that the presence of magnetic fields and an increase in the density contrast of the filament had significant effects on the evolution of the filament. We summarise our main results for each orientation of the magnetic field as follows, noting that in all comparisons the time is normalised by $t_{cs}$: \begin{itemize} \item \textit{Parallel fields}: \begin{description} \item[(i)] Filaments which are orientated either broadside, or nearly-broadside, on to the shock front survive for far longer than those orientated end on. Unless the filament is very small, the length of the filament has no significant effect on its evolution; \item[(ii)] Well-defined linear structures situated on the axis behind the filament are formed only when the filament is end on with respect to the shock front (i.e. orientated at $\theta=90^{\circ}$); \item[(iii)] An increase in the cloud density contrast hastens the destruction of the cloud through the increased presence of turbulent instabilities located on the filament surface. As the density contrast increases, so does the amount of turbulence; \item[(iv)] Low shock Mach numbers restrict the filament from fragmenting, thus significantly prolonging its life. \end{description} \item \textit{Perpendicular/oblique fields}: \begin{description} \item[(vi)] Even if the filament is end on with respect to the shock front, filaments in a perpendicularly/obliquely-orientated magnetic field do not form flux ropes; \item[(vii)] Compared with parallel-orientated fields, perpendicular/oblique fields shield the filament to a degree from the surrounding flow, allowing the filament lifetime to be considerably extended. The filament is more greatly confined by the field and maintains a higher average density; \item[(viii)] Filaments are more rapidly accelerated to the velocity of the post-shock flow due to the effects of the magnetic pressure and field line tension; \item[(ix)] An increase in the filament density contrast does not initiate large turbulent instabilities, compared to the case of a parallel field; \item[(x)] A combination of a mild (e.g. $M=1.5$) shock and a perpendicular/oblique field allows the filament to survive almost intact for a considerable length of time. \end{description} \end{itemize} The work presented in this paper is difficult to apply observationally since the adiabatic simulations do not include realistic physical processes such as thermal conduction, radiative cooling, and self-gravity. In future work we will extend our investigation to include the effects of radiative cooling, and will compare synthetic observations of such simulations with actual observations in order to present a more complete picture of the evolution of filaments in the ISM. It should be noted that \citet{Banda16} explored the effects of using a quasi-isothermal equation of state to approximate the effect of radiative cooling in MHD wind-cloud simulations and found that this led to significantly longer cloud lifetimes compared to the adiabatic case; a comparison with future isothermal shock-filament interactions would, therefore, be of interest.

1607.03929

1607

1607.06217_arXiv.txt

We have carried out an extensive population synthesis study of the ensemble properties of the present-day population of cataclysmic variables (PDCVs) that takes into account the nuclear evolution of high-mass donors close to the bifurcation and dynamical instability limits. Assuming the interrupted magnetic braking paradigm, we confirm many of the general features associated with the observed CV population and find enormous diversity in their secular properties. We predict that nearly half of the non-magnetic CVs with $P_{orb} \ge 6$ hours are at least mildly evolved (i.e., $>$ 50\% of their MS turn-off age). Some of these systems contribute to the observed population of PDCVs in the period gap. We also see an enhancement by up to a factor of two in the probability of detecting CVs at the `minimum period'. This spike is quite narrow ($\approx 5$ minutes) and is attenuated because of the spectrum of WD masses and partly by the evolution of the donors. Our syntheses imply that there should be a very rapid decline in the number of ultracompact CVs (such as AM CVns). We find that between $\sim 0.05$ to 1\% of PDCVs could be UCs, and thus it is likely that the CV channel is probably not the primary contributor to the intrinsic population of UCs (especially for $P_{orb}$ < 30 minutes). Finally, a preliminary analysis of our results suggests that WDs in PDCVs experience a net gain in mass of $\lesssim 0.1 M_\odot$ as a result of high mass-transfer rates early in their evolution.

INTRODUCTION} Cataclysmic variables (CVs) are a very heterogeneous class of semi-detached, interacting binaries consisting of a white dwarf (WD) that is accreting matter from a companion (donor) star that is overflowing its Roche lobe (see Patterson 1984; Warner 1995; and references therein). The recent SDSS survey has produced a wealth of new observational information (see, e.g., Szkody et al. 2011). According to the conventional model for the formation of galactic-disk CVs, the progenitors are primordial binaries for which one star is sufficiently massive and close enough to its companion star that it can engulf the companion during the giant phase of its evolution. The time interval over which this happens is governed by the nuclear timescale of the more massive star (the primary). The binary then enters a short-lived (dynamical timescale) common envelope (CE) phase of evolution during which the companion spirals inside the envelope of the giant star until it approaches the degenerate core (the nascent WD). The transfer of energy leads to the envelope being completely ejected thereby producing a binary that is most likely detached (i.e., a post common-envelope binary [PCEB]). The donor (secondary) can begin to transfer mass if the binary separation is reduced on a sufficiently short timescale through orbital angular momentum loss (AML) or if the secondary expands sufficiently due to its own nuclear evolution. Once mass transfer from the donor to the WD commences, we refer to the binary as a zero-age cataclysmic variable (ZACV). The trajectory of the matter as it undergoes Roche-lobe overflow (RLOF) largely depends on the strength of the magnetic field of the accreting WD. If the WD is non-magnetic, an accretion disk will form around the WD; but if the magnetic field is sufficiently strong, the disk can either be truncated or may not even form if the magnetic field can entrain the matter and force it to flow directly onto the WD's surface. The subsequent evolution of the CV is either driven by orbital angular momentum losses or the nuclear evolution of the donor. According to the `canonical model' of CV evolution, donors whose masses exceed approximately 0.37 $M_\odot$ will experience some form of magnetic braking (MB). It is assumed that MB will become ineffective once the mass of the donor star has been decreased so much that its internal structure becomes completely convective. This has been referred to as the Interrupted Magnetic Braking (IMB) paradigm (see Rappaport et al. 1983 [RVJ]; Spruit \& Ritter 1983; Hameury et al. 1988; Davis et al. 2008). When MB is switched off, the donor shrinks beneath its Roche lobe leading to the cessation of mass transfer at orbital periods of $\approx 3$ hours. Gravitational radiation (GR) losses alone then drive the binary back into a semi-detached state ($\approx 2$ hours) and the system continues to evolve towards a minimum orbital period ($P_{min} \approx 80$ minutes) before returning back to higher periods. It is important to try to infer the evolutionary history of CVs by comparing the observed ensemble properties of CVs with theoretically computed population syntheses. Although there are many subclasses of non-magnetic CVs (e.g., dwarf novae, recurrent novae) whose behaviors are governed by different physical effects, their evolutionary histories share many common features. Cataclysmic variables are especially well-suited for the application of population synthesis (PS) techniques because they may be viewed as ``first order" systems relative to other classes of interacting binaries. Higher order systems would include, for example, low-mass X-ray binaries (LMXBs), and double neutron-star binaries. These systems have extra dimensions of uncertainty complicating the calculation of their formation probabilities. For example, they might have undergone two CE phases or experienced natal neutron star kicks, both of which are subject to large physical uncertainties. Thus it is important to first securely determine the relative probabilities of the various channels that lead to the formation of CVs. The overall population of CVs is sufficiently large ($\gtrsim 1000$) that statistically significant inferences can be made if unbiased samples are available. Unfortunately, selection effects are the bane of this type of study thereby making it difficult to reach unbiased conclusions. The use of PS techniques as a tool for understanding the formation and evolution of all types of interacting binaries has been extensively studied. Some of the key PS studies that have made significant contributions to the CV field include those of de Kool (1992), Kolb (1993), Politano (1996), Howell et al. (1997), Nelemans et al. (2001), Howell et al. (2001 [HNR]), Podsiadlowski et al. (2003), Politano (2004), Kolb and Willems (2005), and Willems et al. (2005, 2007). However, the relative complexity of the physical processes and the wide range of timescales associated with them still pose considerable difficulties even with currently available computing power. The large number of dimensions of parameter space has often required that many simplifying assumptions be implemented in order to make the computations tractable. Except for the pioneering study by Podsiadlowski et al. (2003), previous PS studies have discounted the effects of the internal chemical evolution of the donor star on the Present Day CV (PDCV) population\footnote{We also refer the reader to the work of Andronov and Pinsonneault (2004) who examined how the evolution of chemically evolved CVs is affected by various descriptions of magnetic stellar wind braking.}. In this study, we examine the effects that such evolved donor stars have on the properties of the PDCV population. Our approach is to first compute detailed evolutionary tracks using a Henyey-type code (with state-of-the-art input physics) in order to generate a grid of models that are representative of most types of non-magnetic CV evolution. We then interpolate this grid to obtain a track corresponding to a specific set of initial conditions describing a ZACV. Since the interpolation is computationally inexpensive (as opposed to the calculation of individual CV tracks), it is relatively easy to generate large populations of PDCVs from a previously synthesized group of ZACVs. This paper is organized as follows: In \S 2 we describe the stellar code that is used to create the CV evolution grid, the method of interpolation, and our assumptions concerning the physics of the population synthesis itself. In \S 3, we present the results of the population synthesis study and explore the observable properties of the ensemble of PDCVs; a discussion of the implications of these results can be found in \S 4. In particular, we analyze the importance of evolved donors with respect to the period gap, the minimum period, and the formation of ultracompact binaries, and we briefly comment on the possible increase in the mass of WDs in CVs. Our conclusions and plans for future work are summarized in \S 5.

Using a highly efficient approach to population synthesis, we pre-compute representative tracks for the evolution of CVs and then interpolate the grid for a specific set of initial conditions corresponding to the properties of a particular ZACV. A similar approach to population synthesis has been used by Podsiadlowski et al. (2003) and by Willems et al. (2005, 2007); however, both used stellar models whose input physics was not as intricate. Since the interpolations are relatively inexpensive (but must be treated very carefully near edges), we can use Monte Carlo methods to generate large datasets of ZACVs for any assumed set of parameters describing their formation. This allows us to explore the many dimensions of parameter space (e.g., CE efficiency, mass correlation) in an efficient manner. We have also demonstrated the fidelity of this approach by taking our results (where applicable) and comparing them with previous studies. This is the first time that a full CV population synthesis has been carried out for the present-day population that takes into account the effects of chemical evolution for all donor masses up to the bifurcation limit and for the complete spectrum of WD masses. The results show that the assumption that CV evolution can be approximated by ZAMS mass-losing donors is generally a valid one. However, the range of the values of the observables at a given epoch has been greatly underestimated relative to other studies. This is especially true for the relatively rare, long-period CVs ($P_{orb} > 6$ hr) where we find good agreement with the range of observed spectral types. We are also able to populate the entire gap with non-magnetic CVs and predict the relative numbers that should be observed on both sides of the gap and in the gap itself. We find that the relative distribution is not in contradiction with the observations, and we claim that a non-negligible fraction of systems in the gap may be derived from (partially) evolved donors. We show that there should be a significant enhancement in the number of CVs observed near $P_{min}$. If GR is the sole mechanism for AML below the period gap, then this increase in the detection probability should manifest itself in an 4 - 5 minute period range coincident with $P_{min}$. If we can obtain an unbiased and statistically significant sample of CVs whose orbital periods are close to the minimum value, the theoretical predictions of the population synthesis models could be used to determine if there is a supplemental AML mechanism acting contemporaneously and determine if its magnitude varies depending on the properties of the binary system (e.g., the chemical composition of the donor, the magnetic field strength of the WD, etc.). With respect to ultracompacts, including AM CVn binaries, we conclude that the `CV channel' is not likely to account for any substantial fraction of these systems. In fact, if this channel has led to the formation of any of the observed ultracompacts, then it is very likely that they will be the ones with orbital periods of not much less than 60 minutes. Two of the main difficulties encountered in using our PS approach concern the need to compute a suitable grid so that the interpolations are as accurate as possible and the numerical expense of computing a new grid if the physics governing the evolution of the donor and/or binary is changed. With regard to the first issue, the grid needs to be carefully constructed so that difficult regions in initial-condition space that cause any of the parameters to change rapidly can be treated carefully (e.g., dynamical instabilities and the bifurcation limit). We plan to add many additional tracks to our grid in an attempt to address this issue and to incorporate the contributions from brown dwarfs. We will also extend the grid above the bifurcation limit and carry out a PS on these long-period systems that lead to mergers or the formation of double degenerates (including those containing hybrid HeCO WDs). Finally, we plan to recalculate the grid by self-consistently incorporating the effects of mass gain (or erosion) of the WD accretor during the evolution. Based on the very preliminary results presented in this paper, it appears as if WDs in PDCVs can gain $\sim 0.1 M_\odot$, an amount that may be consistent with the observations.

1607.06217

1607

1607.01601_arXiv.txt

We consider limits on the local ($z=0$) density ($n_0$) of extragalactic neutrino sources set by the nondetection of steady high-energy neutrino sources producing $\gtrsim50$~TeV muon multiplets in the present IceCube data, taking into account the redshift evolution, luminosity function and neutrino spectrum of the sources. We show that the lower limit depends moderately on source spectra and strongly on redshift evolution. We find $n_0\gtrsim{10}^{-8}-{10}^{-7}~{\rm Mpc}^{-3}$ for standard candle sources evolving rapidly, $n_s\propto{(1+z)}^3$, and $n_0\gtrsim{10}^{-6}-{10}^{-5}~{\rm Mpc}^{-3}$ for nonevolving sources. The corresponding upper limits on their neutrino luminosity are $L_{{\nu_\mu}}^{\rm eff}\lesssim10^{42}-10^{43}~{\rm erg}~{\rm s}^{-1}$ and $L_{{\nu_\mu}}^{\rm eff}\lesssim10^{41}-10^{42}~{\rm erg}~{\rm s}^{-1}$, respectively. Applying these results to a wide range of classes of potential sources, we show that powerful ``blazar'' jets associated with active galactic nuclei are unlikely to be the dominant sources. For almost all other steady candidate source classes (including starbursts, radio galaxies, and galaxy clusters and groups), an order of magnitude increase in the detector sensitivity at $\sim0.1-1$~PeV will enable a detection (as point sources) of the few brightest objects. Such an increase, which may be provided by next-generation detectors like {\it IceCube-Gen2} and an upgraded KM3NET, can improve the limit on $n_0$ by more than two orders of magnitude. Future gamma-ray observations (by {\it Fermi}, HAWC and CTA) will play a key role in confirming the association of the neutrinos with their sources.

The detection of an extraterrestrial high-energy, $\sim30$~TeV to a few PeV, neutrino flux by the IceCube Collaboration~\citep{Aartsen:2013bka,Aartsen:2013jdh,Aartsen:2014gkd,Aartsen:2014muf,Aartsen:2015ita,Aartsen:2015rwa} marks the beginning of high-energy neutrino astrophysics. The observed signal is consistent with an isotropic arrival distribution of the neutrinos, and with equal contents of $\nu_e$, $\nu_\mu$ and $\nu_\tau$ and their anti-particles. Above $\sim100$~TeV, the flux and spectrum are consistent with the Waxman-Bahcall (WB) bound~\citep{Waxman:1998yy}, $E_{\nu}^2\Phi_{\nu_i}\simeq{10}^{-8}~{\rm GeV}~{\rm cm}^{-2}~{\rm s}^{-1}~{\rm sr}^{-1}$ per flavor, with a possible spectral break or cutoff at a few PeV. These properties together hint to a cosmological origin of the observed neutrino flux, most likely related to the accelerators of high-energy cosmic rays (CRs) (see Refs.~\cite{Waxman:2013zda,Halzen:2013dva,Meszaros:2014tta} for reviews). High-energy neutrinos are expected to be emitted in this case mainly by the decay of mesons and muons produced in interactions of CRs with ambient gas (nucleons) or radiation fields within or surrounding the CR sources, with neutrino to CR energy ratio typically given by $E_\nu/E_{\rm cr}\approx(0.03-0.05)/A$, where $A$ is the CR atomic number~\cite{Waxman:2013zda,Murase:2010gj}. IceCube's analysis of lower-energy neutrino events indicates an excess of events at $\sim30$~TeV above an extension to low-energy of the ``flat", $E_{\nu}^2\Phi_{\nu_i}=Const.$, higher-energy spectrum~\cite{Aartsen:2014muf,Aartsen:2015ita}. Assuming that the astrophysical neutrino spectrum is described by a single power-law, $E_{\nu}^2\Phi_{\nu_i}\propto E_{\nu}^{2-s}$, different analyses of IceCube's data lead to different constraints on the spectral index $s$. There is some ($\sim2\sigma$) tension between analyses with higher-energy thresholds, yielding values consistent with $s=2$, and those with lower-energy thresholds, yielding $s\sim2.5$ (see Refs.~\cite{Aartsen:2014muf,Aartsen:2015ita,Aartsen:2015zva}). This may indicate a new component contributing to the flux at $\lesssim100$~TeV energies (see Refs.~\cite{Murase:2015xka,Chen:2014gxa,Palladino:2016zoe,Neronov:2016bnp} for discussion). The existence of such a component does not affect the analysis presented in this work, which is focused on the higher-energy, $\gtrsim100$~TeV, neutrinos, the flux and spectrum of which are consistent with the WB bound. It should be noted in this context, that the observed neutrino flux is comparable to the sub-TeV diffuse gamma-ray background flux measured by {\it Fermi}~\cite{Ackermann:2014usa}. An extension with $s\gtrsim2.1-2.2$ of the $\gtrsim100$~TeV neutrino flux to low energies, $\sim0.1-1$~TeV, implies a diffuse gamma-ray flux that exceeds the {\it Fermi} gamma-ray background, while an extension with a smaller spectral index is consistent with the {\it Fermi} data~\cite{Murase:2013rfa} (see also Fig.~\ref{Unified}). The coincidence of the IceCube signal with the WB bound implies that the universal average of the energy production rate of ultrahigh-energy (UHE), $>10^{19}$~eV, CRs is similar to the rate of energy production of $\sim0.1-1$~PeV neutrinos. The observed neutrino signal may thus be explained by a model in which the sources of UHECRs produce protons with a ``flat" spectrum (equal energy per logarithmic particle energy interval), $E_{\rm cr}Q_{E_{\rm cr}}=E_{\rm cr}^2d\dot{n}_{{\rm cr}}/dE_{\rm cr}=Const.$, and reside in ``calorimetric" environments in which protons of energy $\lesssim50-100$~PeV lose all their energy to meson production (i.e. these CRs are confined for a time longer than their $pp$ energy-loss time). This is the simplest explanation in the sense that the CR sources are known to exist, the required CR spectrum is consistent with that observed at $\gtrsim10^{19}$~eV and with theoretical expectations, the model contains no free parameters (the production rate of CRs is determined by observations and the fraction of their energy converted to mesons is ${\rm min}[f_{pp},1]\simeq1$ below $50-100$~PeV), and there is a known class of objects, which are expected to act as ``calorimeters" for $\lesssim50-100$~PeV protons -- starburst galaxies (SBGs). In fact, the signal detected by IceCube has been predicted to be produced by sources residing in SBGs~\citep{Loeb:2006tw}. The only assumption that one needs to make is that CR production is related to star-formation activity (which would be the case for sources like gamma-ray bursts (GRBs), energetic supernovae, or perhaps stellar tidal disruptions by supermassive black holes). The main uncertainty in this model (see Ref.~\cite{Waxman:2013zda} for a detailed discussion) is related to the fact that galaxies rapidly forming stars are inferred to act as calorimeters for CR protons based on the observations of local ($z=0$) SBGs, while most of the neutrinos are produced by galaxies rapidly forming stars at redshifts $z\sim1-2$. The properties of these galaxies are less well-constrained, and hence the fraction of them which are ``calorimetric" is uncertain. While the above unified scenario for the production of UHECRs and of IceCube's neutrinos is simple and natural~\cite{Katz:2013ooa,Waxman:2013zda}, we have no direct confirmation for the emission of neutrinos from SBGs. A wide range of different models have been proposed for the origin of IceCube's neutrinos. Models predicting the production of high-energy neutrinos through the decay of mesons and muons produced by high-energy CRs may be divided into two types: ``CR accelerator models", where neutrinos are produced within the CR source, and ``CR reservoir models", where neutrinos are produced while they are confined within the environment surrounding the CR source. CR accelerator models for the emission of high-energy neutrinos have been proposed, for example, for gamma-ray bursts (GRBs, e.g., Refs.~\cite{Waxman:1997ti,Murase:2006mm,Bustamante:2014oka}) and blazars~\cite{Mannheim:1995mm,Atoyan:2001ey,Atoyan:2002gu,Dermer:2012rg}, while CR reservoir models for the emission of high-energy neutrinos have been proposed for SBGs~\citep{Loeb:2006tw}, galaxy clusters and groups (GCs/GGs)~\cite{Murase:2008yt,Kotera:2009ms}, and active galactic nuclei (AGN)~\cite{Kimura:2014jba,Hooper:2016jls}. In accelerator models, mesons are typically produced by interactions of CRs with radiation, while in reservoir models they are typically produced by inelastic hadronuclear collisions. Some AGN core models, where protons are accelerated and undergo hadronuclear collisions in the vicinity of the black hole (e.g., Refs.~\cite{Tjus:2014dna,Kimura:2014jba}), are an exception. In models where the mesons are produced by photohadronic ($p\gamma$) interactions with radiation, a low-energy cutoff may be expected in the neutrino spectrum. For a characteristic energy $E_\gamma$ of the ambient photons, the low-energy cutoff is expected at $\sim0.05 E_{\rm min}$, where $E_{\rm min}\sim m_p m_\pi c^4/E_\gamma$ is the minimum CR nucleon energy required to allow pion production (in case the source is moving relativistically with Lorentz factor $\Gamma$, $E_{\rm min}\sim\Gamma^2 m_p m_\pi c^4/E_\gamma$). In models where mesons are produced by inelastic hadronuclear ($pp$) collisions, we expect the neutrino spectrum to extend down to sub-GeV energies, since pion production is allowed for all relativistic CRs. In this case, the neutrino spectral index should satisfy $s\lesssim2.1-2.2$, since, as explained above, for steeper spectra the accompanying gamma-ray flux will be inconsistent with the {\it Fermi} gamma-ray background below $\sim0.1-1$~TeV~\cite{Murase:2013rfa}. If IceCube's neutrinos are produced by $pp$ interactions, the sources significantly contribute to the extragalactic gamma-ray background. This is not necessarily the case for $p\gamma$ scenarios~\cite{Murase:2015xka}, since the radiation field required to produce sub-PeV neutrinos via $p\gamma$ interactions naturally leads to a large two-photon annihilation optical depth for GeV-TeV gamma rays~\cite{Waxman:1998yy,Murase:2015xka}. The neutrino signal detected in IceCube is not consistent with the predictions of most CR accelerator models derived prior to the IceCube detection. Neutrino production within GRB sources is expected to produce a flux which is $\sim10(E_\nu/1~{\rm PeV})$\% of the WB flux at $E_\nu\lesssim1$~PeV~\cite{Waxman:1997ti}. The neutrino spectra predicted to be produced in AGN jet models (in particular blazar models) are typically inconsistent with (too hard compared to) IceCube's data~\cite{Murase:2014foa,Padovani:2015mba,Murase:2015ndr}. Nevertheless, the blazar models are not ruled out since their underlying assumptions (e.g., the maximum CR energy) may be modified. In fact, many of them have been revised after IceCube's discovery, with parameters appropriately chosen to reproduce IceCube's flux above $\sim100$~TeV; in particular see Refs.~\cite{Dermer:2014vaa,Tavecchio:2014xha,Tavecchio:2014eia,Petropoulou:2015upa} and Fig.~\ref{Blazar} for blazar models. In addition, models of CR accelerators obscured in gamma rays are considered (see Refs.~\cite{Murase:2015xka,Winter:2013cla,Kistler:2013my} and references therein), and AGN core models~\cite{Stecker:1991vm,Stecker:2013fxa,Kimura:2014jba,Kalashev:2015cma} have been modified such that their flux normalization is adjusted to IceCube's flux. Finally, we note that ``choked jet supernova" models~\cite{Murase:2013ffa,Nakar:2015tma,Senno:2015tsn,Tamborra:2015fzv,Meszaros:2001ms} may also account for the IceCube data. However, in our current analysis we derive constraints on steady sources and therefore do not address these transient models further. The main goal of this paper is to demonstrate that the limits that can be set by IceCube's measurements on the source density exclude some widely discussed candidate sources, and to show that an order of magnitude increase in the detector sensitivity at $\sim100$~TeV is likely to enable the detection (as point sources) of the few brightest objects for almost all other candidate source classes. The limit on the density of ``standard candle" sources is derived in Sec.~\ref{nuA}. Its implications to various classes of sources are described in Sec.~\ref{nuB}, taking into account the redshift evolution and the luminosity function (LF) of the sources. The increase in the detector sensitivity required to enable the detection of neutrino point sources (sources producing multiple neutrino events) is discussed in Sec.~\ref{multi1}. The nondetection of point sources has been used in earlier work \cite{Lipari:2008zf,Silvestri:2009xb,Murase:2012df} to set limits on the density of neutrino sources. The limits derived here are more stringent thanks to the completion of the full IceCube detector, as recently discussed~\cite{Ahlers:2014ioa,Kowalski:2014zda,Murase:2014JSI,Murase:2015ipa}. Moreover, our analysis goes beyond those of earlier work in taking into consideration the dependence on the redshift evolution and on the LF of the sources, and also on the neutrino spectrum of the sources. As explained below, the density limit is sensitive to the redshift evolution, and taking into account the LF of the sources implies that their ``effective" number density $n_0^{\rm eff}$ (the number density of sources dominating the flux), which is constrained by the derived density limit, may be significantly smaller than the total density, $n_0^{\rm tot}$. The non-blazar component of the sub-TeV extragalactic gamma-ray background flux measured by {\it Fermi}~\cite{TheFermi-LAT:2015ykq,Lisanti:2016jub} can be explained by the sum of hadronic gamma rays produced inside the sources and ``cosmogenic" gamma rays produced in CR interactions with the cosmic microwave background and extragalactic background light (see Fig.~\ref{Unified}). In particular, SBGs have been predicted to produce a significant contribution to the diffuse gamma-ray background~\cite{Thompson:2006np,Lacki:2010vs,Lacki:2012si}, consistent with the neutrino flux measured by IceCube. The source density and luminosity reached by gamma-ray observations are discussed in Sec.~\ref{multi2}, where we show that some neutrino source models like CR reservoir models should be testable with future gamma-ray observatories. Our conclusions are summarized and discussed in Sec.~\ref{sec:discussion}. We use $\Omega_m=0.3$, $\Omega_\Lambda=0.7$ and $H_0=70~{\rm km~s^{-1}~{Mpc}^{-1}}$ throughout. \begin{figure}[t] \includegraphics[width=3.00in]{f1.eps} \caption{ Diffuse CR (thin dotted line), gamma-ray (thick solid line, adapted from Ref.~\cite{Murase:2013rfa}), and all-flavor neutrino (thick dashed line, adapted from Ref.~\cite{Murase:2013rfa}) intensities predicted in our grand-unified cosmic particle model in which the UHECR flux is produced by an extragalactic distribution of proton sources, producing a ``flat" CR proton spectrum, $E_{\rm cr}Q_{E_{\rm cr}}=0.5\times10^{44}~{\rm erg}~{\rm Mpc}^{-3}~{\rm yr}^{-1}$, and residing in environments which are almost ``calorimetric" for $E_{\rm cr}\lesssim50-100$~PeV protons. The observed UHECR flux and spectrum (Auger data points from Ref.~\cite{ThePierreAuger:2015rha,Aab:2015bza}) and IceCube's neutrino flux and spectrum (IceCube data points from Ref.~\cite{Aartsen:2014gkd}) are both self-consistently explained (see Refs.~\cite{Katz:2013ooa,Waxman:2013zda} for detailed discussion). The non-blazar contribution of the diffuse gamma-ray background measured by {\it Fermi} (shaded region above 50~GeV), which amounts to $\sim30\%$~\cite{Lisanti:2016jub} (see also Ref.~\cite{TheFermi-LAT:2015ykq}) of the ``total'' extragalactic gamma-ray background, shown as {\it Fermi} data points~\cite{Ackermann:2014usa}, is simultaneously accounted for in this model (see Refs.~\cite{Murase:2013rfa,Murase:2015xka} for details). The model UHECR flux (thin dotted line) and corresponding cosmogenic neutrino (thin dashed line) and gamma-ray (thin solid line) fluxes are adapted from Ref.~\cite{Decerprit:2011qe}. \label{Unified} } \vspace{-1.\baselineskip} \end{figure} \begin{figure}[t] \includegraphics[width=3.00in]{f2.eps} \caption{All-flavor neutrino fluxes of ``post-IceCube" blazar models, with parameters chosen to explain the IceCube data. We consider in this paper three spectral templates, taken from Tavecchio \& Ghisellini (TG15)~\cite{Tavecchio:2014eia} and Petropoulou et al. (PDPMR15)~\cite{Petropoulou:2015upa} for BL Lac objects, and from Dermer et al. (DMI14)~\cite{Dermer:2014vaa} for flat spectrum radio sources (FSRQs). \label{Blazar} } \vspace{-1.\baselineskip} \end{figure}

\label{sec:discussion} We have derived in Sec.~\ref{nuA} constraints on the density and luminosity of steady ``standard candle" neutrino sources dominating the high-energy, $\gtrsim100$~TeV, neutrino flux detected in IceCube, based on the nondetection of ``point sources" producing high-energy multiple neutrino-induced muon tracks in the detector. The limits are given in Eqs.~(\ref{eq:Lnu}) and(\ref{eq:n0}), and illustrated in Figs.~\ref{IClim1} and \ref{IClim2} (an upper limit on the density of steady sources at a given luminosity, which is valid for sources that do not necessarily dominate the flux, is given in Eq.~\ref{eq:n0a}). These limits were applied in Sec.~\ref{nuB} to a wide range of potential source classes, taking into account their redshift evolution and LF. While the distribution of electromagnetic luminosities, i.e. the photon LF, of different classes of objects are known, the neutrino LFs of most source classes are not known and are model dependent. We therefore did not attempt a comprehensive analysis under different model assumptions regarding the neutrino LFs of various classes of objects. Rather, for each class of objects we defined an effective number density $n_s^{\rm eff}$ (see Eq.~\ref{eq:neff}), characterizing the density of sources dominating the flux. Our conclusions are not sensitive to the details of the relation between the photon and the neutrino LFs and to the exact definition of $n_s^{\rm eff}$. The classes of sources that are ruled out by the constraint of Eq.~(\ref{eq:n0}), and for which there is a large difference between $n_0^{\rm eff}$ and the total number density $n_0^{\rm tot}$ (see Table 1), are characterized by $n_0^{\rm eff}$ values which are orders of magnitude smaller than the limit of Eq.~(\ref{eq:n0}). The constraints imply that rare objects, such as powerful blazar jets, are unlikely to dominate IceCube's flux. For blazars, we showed that the conclusion does not change even if harder (possibly more realistic) neutrino spectra are used (see Fig.~\ref{IClim2}). This result is consistent with those obtained from stacking and cross-correlation analyses~\cite{Wang:2015woa,Glusenkamp:2015jca}. However, it should be noted that neutrino emission by transient AGN ``flares" \cite{Atoyan:2001ey,Dermer:2014vaa,Petropoulou:2016ujj,Kadler:2016ygj} is not constrained by the current analysis, as is the case for other types of transient sources. CR reservoir models and AGN core models, with source density of $n_0^{\rm eff}\gtrsim3\times10^{-6}~{\rm Mpc}^{-3}$, are not constrained by current IceCube data. An order of magnitude improvement in $F_{\rm lim}$, the minimum flux required for a source to be detectable as a point source, can improve the limit on $n_0$ by more than two orders of magnitude, and will likely enable the detection (as point sources) of the few brightest objects for almost all candidate source classes, including SBGs, RL AGN, and GCs/GGs (see Table 1 and Fig.~\ref{IClim1}). Such an improvement in $F_{\rm lim}$ requires an order of magnitude increase in the effective mass of the detector at $0.1-1$~PeV (where the background is negligible), which may be provided by {\it IceCube-Gen2} and an upgraded KM3NeT. Searches for the brightest neutrino sources, including stacking and cross-correlation analyses, are powerful especially for the SBG and RL AGN models. However, in general, they are model dependent. While the detection of a few point sources may confirm the validity of a suggested source model, nondetections may not necessarily rule out all the models for the suggested source class. This is due to the fact that large deviations from an ``average source luminosity" cannot be excluded when the source physics is not well understood, and model uncertainties often prevent accurate predictions. For example, testing the LL AGN core model is feasible in the canonical case since $n_0^{\rm eff}\sim{10}^{-3}~{\rm Mpc}^{-3}$ can be reached by {\it IceCube-Gen2} for nonevolving sources. However, accessing $n_0^{\rm tot}\gtrsim{10}^{-2}~{\rm Mpc}^{-3}$ may be difficult. At photon energies of 1~GeV to 1~TeV, which are well below the energy of the neutrinos observed by IceCube but accessible to gamma-ray telescopes, a gamma-ray luminosity of $E_\gamma L_{E_\gamma}\approx2(E_\gamma/2E_\nu)^{2-s}E_\nu L_{E_{\nu_\mu}}$ is expected for CR reservoirs (like SBGs and GCs/GGs) in which (a) the parent CRs are produced with a power-law spectrum, (b) the production of mesons is dominated by inelastic $pp$ collisions with nucleons, and (c) the internal absorption of gamma rays by two-photon annihilation interactions is negligible below $\sim1-10$~TeV. We showed that gamma-ray observations may be useful for testing models of this type. In particular, dedicated targeted observations by the CTA detector of the brightest objects of a complete catalogue of candidate neutrino sources will lead to the detection of individual bright sources for source classes with $n_0^{\rm eff}\lesssim10^{-4}~{\rm Mpc}^{-3}$. \medskip

1607.01601

1607

1607.08672_arXiv.txt

We used a consistent and robust solar model to obtain upper limits placed by neutrino telescopes, such as IceCube and Super-Kamiokande, on the Dark Matter-nucleon scattering cross-section, for a general model of Dark Matter with a velocity dependent ($p$-wave) thermally averaged cross-section. In this picture, the Boltzmann equation for the Dark Matter abundance is numerically solved satisfying the Dark Matter density measured from the Cosmic Microwave Background (CMB). We show that for lower cross-sections and higher masses, the Dark Matter annihilation rate drops sharply, resulting in upper bounds on the scattering cross-section one order of magnitude above those derived from a velocity independent ($s$-wave) annihilation cross-section. Our results show that upper limits on the scattering cross-section obtained from Dark Matter annihilating in the Sun are sensible to the uncertainty in current standard solar models, fluctuating a maximum of 20 \% depending on the annihilation channel.

The joint efforts of physicists in the last four decades has resulted in solid evidence, not only at astrophysical scales, but also cosmological, which leave no doubt that our Universe is mainly populated by a still undetected non-interactive type of matter, the so-called Dark Matter, which nature is still unknown. Amongst the numerous theories devised to solve this problem, the picture of a weakly interactive massive particle (WIMP) arises as the most favourable, since the dark matter abundance inferred today from the Cosmic Microwave Background (CMB) matches the abundance of a relic particle with an annihilation cross-section of the order of the weak scale. Furthermore, new particle physics theories motivated by different reasons, provide natural candidates for this type of matter, making this a highly interdisciplinary field of investigation. \\ In the picture of particle Dark Matter, WIMPs that populate the Milky Way can be gravitationally captured by the Sun~\citep{steigman78}. Since WIMPs are mandatorily stable, they will accumulate inside Sun and annihilate to become standard model particles. This annihilation will produce a distinctive neutrino signal that can be detected in current and projected neutrino detectors, providing an excellent indirect survey to Dark Matter properties, which has been extensively studied~\citep{silk85,gaisser86,griest87,wilkstrom09}.\\ However, most indirect Dark Matter searches focus on simple models where WIMPs are Majorana particles and annihilate through an $s$-wave, velocity independent thermally averaged cross-section. In these models, it is usual to fix $\langle \sigma v \rangle \sim 10^{-26} \ \text{cm}^3 \text{s}^{-1}$ , in order to respect the Dark Matter density~\citep{kolb89}, which is precisely determined from CMB measurements ~\citep{planck15}. Theoretically, the Dark Matter abundance is defined by the time of WIMP freeze out, which happens when the universe's temperature drops below the WIMPs mass, resulting in the decoupling of Dark Matter particles from the primordial universe thermal bath. This will result in a constant number of WIMPs per co-moving volume which corresponds to the one measured today. \\ Despite the fact that the usual approach in literature is to use a constant thermally averaged cross-section, there are a large number of well-motivated models which have a $p$-wave contribution, (i.e. a dependence in the relative velocity between WIMPs) to the annihilation cross-section which can be dominant. Models where Dark Matter is a Majorana particle, such as the \textit{neutralino}, a natural candidate which arises in the Minimal Super symmetric Standard Model (MSSM), $s$-wave annihilation to fermion anti-fermion pairs is helicity suppressed by a factor of $(m_f/m_{\chi})^2$, where $m_{\chi}$ is the WIMPs mass~\citep{sheldon10,goldberg83}. Furthermore, due to $CP$ conservation, final states with $CP$=+1, are only accessible through $p$-wave annihilations ($s$-wave states for two identical majorana fermions are $CP$=-1). Hence, neutralino annihilation to $HH$ or any combination of the vectorial bosons $W$ and $Z$, can only occur through $p$-wave annihilation ~\citep{drees92}. Another well known example is the case of parity conserving minimal extensions to the Standard Model (SM) with a fermionic Dark matter candidate - a gauge singlet Dirac fermion - in which annihilation to scalar states with even parity, such as $\chi \chi \rightarrow HH$, does not receive a contribution from the $s$-wave annihilations ~\citep{kim07}. \\ In this article we obtain new limits on the Dark Matter scattering cross-section from the upper limits on the neutrino fluxes measured by the Super-Kamiokande and IceCube neutrino telescopes, using a general model where $p$-wave annihilation is the leading contribution to the total annihilation cross-section, i.e., where we focus in the second term in the thermally averaged annihilation cross-section expansion, \\ \begin{equation} \left \langle \sigma v \right \rangle = a + b \left \langle v^2 \right \rangle + \mathcal{O}(\langle v^4 \rangle ) \simeq \frac{b'}{x}, \label{eq:annCS} \end{equation} with $x = m_{\chi}/T$, and where we also assumed that the $s$-wave contribution to the annihilation, as well as higher order terms with $ \mathcal{O}\left( \langle v^4 \rangle \right)$, are negligible. The coefficient $b$ in \ref{eq:annCS} is assumed constant and obtained taking into account the Dark Matter density at the time of freeze-out, which is roughly the same as today (see sec. \ref{sec:DMabundance}). It is important to note that, since we are mainly interested in the epochs from the moment of WIMP freeze-out, the expansion in \ref{eq:annCS} is only accurate if WIMPs freeze out at non-relativistic velocities, i.e. if $T_F < m_{\chi}$, where $T_F$ is the temperature of freeze-out. However, in most cases this happens at $T_F \simeq m_{\chi}/20 \ll m_{\chi}$ ~\citep{jungman96}, which means that the thermally averaged annihilation cross-section in \ref{eq:annCS} is a reliable approximation for our analysis.\\ In section \ref{sec:DMabundance} we obtain the coefficient $b$ by solving the Boltzmann equation for a relic particle with a velocity dependent annihilation cross-section. In sec. \ref{sec:DMsun} we present the formalism needed to compute the neutrino flux from dark matter annihilation in this picture, as well as the stellar evolution code used to compute the Dark matter capture and annihilation to SM particles. In sec. \ref{sec:DMresults} we present our results followed by some final remarks.

Dark Matter particles trapped in the Sun will annihilate and create a neutrino signal that can be used to survey its properties. In this article we studied the neutrino emission for a simple model in which the main contribution for the annihilation comes from the $p$-wave channel, i.e. a velocity dependent cross-section. To obtain the annihilation cross-section we numerically solved the Boltzmann equation for the density of a relic particle satisfying the Dark Matter abundance as measured today from the CMB. To convert the upper limits on the neutrino flux from the \textsc{IceCube} and \textsc{SuperKamiokande} detectors to upper limits on the WIMP scattering cross-section we used a robust stellar evolution code to model the Sun including Dark Matter capture and annihilation. Assuming that WIMPs distribute isothermally in the Sun's core, we derived an analytical expression for the annihilation coefficient which is directly proportional to the solar central temperature.\\ Differently from the usual case adopted in literature, where WIMPs annihilate through a velocity independent constant annihilation cross-section ($s$-wave channel), the neutrino signal will be directionally proportional to the annihilation coefficient, resulting in upper limits on the scattering cross-section of at least one order of magnitude above the $s$-wave case, which reduces the tension with results from other detection experiments.\\ We also studied the impact of the current uncertainty on solar models (mainly due to the imposition of different solar heavy elements mixtures) in our results. We found out that models with a higher metallicity to hydrogen ratio have an enhancement of $\sim 20 \%$ on the Capture rate for Spin Independent scattering, while models with lower $\left(Z/X \right)_{\astrosun}$ capture $\sim 5\%$ more Dark Matter for Spin Dependent scattering. This variations can reflect a maximum of $\sim 20 \%$ uncertainty on the upper limits for the scattering cross-section for both $s$ and $p$-wave annihilations.\\

1607.08672

1607

1607.01279_arXiv.txt

{} {The activity levels of the solar-twin candidates HD~101364 and HD~197027 are measured and compared with the Sun, the known solar twin 18~Sco, and the solar-like star 51~Peg. Furthermore, the absolute ages of these five objects are estimated from their positions in the HR diagram and the evolutionary (relative) age compared with their activity levels.} {To represent the activity level of these stars, the Mount Wilson S-indices were used. To obtain consistent ages and evolutionary advance on the main sequence, we used evolutionary tracks calculated with the Cambridge Stellar Evolution Code.} {From our spectroscopic observations of HD~101364 and HD~197027 and based on the established calibration procedures, the respective Mount Wilson S-indices are determined. We find that the chromospheric activity of both stars is comparable with the present activity level of the Sun and that of 18~Sco, at least for the period in consideration. Furthermore, the absolute age of HD~101364, HD~197027, 51~Peg, and 18~Sco are found to be 7.2, 7.1, 6.1, and 5.1 Gyr, respectively.} {With the exception of 51~Peg, which has a significantly higher metallicity and a mass higher by about 10\% than the Sun, the present Sun and its twins compare relatively well in their activity levels, even though the other twins are somewhat older. Even though 51~Peg has a similar age of 6.1 Gyr, this star is significantly less active. Only when we compare it on a relative age scale (which is about 20\% shorter for 51~Peg than for the Sun in absolute terms) and use the higher-than-present long-term S$_{\rm{MWO}}$ average of 0.18 for the Sun, does the S-index show a good correlation with evolutionary (relative) age. This shows that in the search for a suitably similar solar twin, the relative main-sequence age matters for obtaining a comparable activity level.}

Solar twins are very important since they allow astrophysicists to look into the past and future of the Sun. For a star to be considered as a solar twin, its respective stellar parameters must be very similar to those of the Sun. It is not sufficient that the candidate star be located in the same place of the empirical HR diagram, the candidate star must also have a metallicity similar to that of the Sun and possess a similar evolutionary age. For young stars, sensitive age indicators are rotation, magnetic activity \citep[e.g.][]{Mamajek2008ApJ687.1264M}, and lithium abundance, all of which considerably diminish during the main-sequence (MS) lifetime of a star. By contrast, evolutionary effects become more visible in the HR diagram only in the second half of this phase, just when the former indicators become less and less sensitive and more ambiguous. One of the best solar twins appears to be 18~Sco, a star quite similar to our Sun \citep{Porto_de_Mello1997ApJ...482L..89P}. 18~Sco was used as comparison star in a study of HD~101364 by \cite{Melendez2012A&A..543A.29M}, who described HD~101364 as a ``remarkable solar twin'' with an age estimate of (3.5$\pm$0.7) Gyr from the lithium abundance, not too far from the age of 18~Sco, which \cite{Melendez2012A&A..543A.29M} estimated as 2.7 Gyr. \cite{Monroe2013ApJ774L32M} presented the star HD~197027 as a solar twin with an age of 8.2 Gyr, while, by comparison, our Sun only has an age of $\approx$4.5 Gyr, which is a factor $\text{of about }$2 lower than the above age estimate of HD~197027; the same authors estimated an age of 2.9 Gyr for 18~Sco. The same star was studied by \cite{Ramirez2014A&A572A48R}, who presented log~R$^{'}_{\rm{HK}}$ values and ages for 18~Sco and HD~197027. While their log~R$^{'}_{\rm{HK}}$ values are similar, the ages are not, since they estimated the age of 18~Sco to be 3 Gyr and that of HD~197027 to be 6.7 Gyr. To study the later activity evolution of our Sun, a more evolved MS star is clearly desirable. We therefore selected 51~Peg, a star with a somewhat higher mass and metallicity than the Sun. This higher mass and metallicity mean that 51 Peg evolved somewhat faster than the Sun. In this way, 51~Peg serves as a good reference for the future solar activity evolution as a late MS star even though its absolute age is comparable. To study the activity--age relation, a reliable mean S$_{\rm{MWO}}$-index is needed to represent the mean level of chromospheric activity of the star. Therefore, measurements over a long period are required because chromospheric activity is a phenomenon exhibiting variability on both short and longer timescales. The Sun, the solar twin 18~Sco, and the evolved solar-like star 51~Peg are all included in our stellar activity long-term monitoring program to study chromospheric activity with our TIGRE telescope. We therefore have obtained quite a large number of S$_{\rm{MWO}}$ values for these three objects. To compare the chromospheric activity of the Sun, the solar twin 18~Sco, the evolved solar-like star 51~Peg, HD~101364, and HD~197027, we can directly use the S-index because all these stars and the Sun have very comparable $B-V$ colour indices. In this case, the Mount Wilson S-index is suited best because it is the more empirical and less $B-V$-dependent activity indicator compared to the log~R$^{'}_{\rm{HK}}$ value. In addition, we estimate the ages of all four object by matching their HR diagram positions with suitable evolutionary tracks to cross check these results with the determined activity levels. Interestingly, recent studies \citep{Reiners2012ApJ...746...43R,schroeder2013A&A554A50S} of solar-like MS stars with some range of mass showed that the relative age of a star, that is, its absolute age $\tau$ with respect to some reference age, compares best with the activity level and not with the absolute age itself; we return to this point below. \begin{figure*} \centering \includegraphics[scale=0.4]{s_tigre_plot.eps} \caption{TIGRE spectrum of the Sun (moonlight) with the band-passes used in the S$_{\rm{TIGRE}}$ measurement.} \label{Fig1} \end{figure*}

Chromospheric activity is a phenomenon showing variability on long timescales (years and to centuries) caused by activity cycles and short timescales (days to years) produced by the evolution and rotational modulation of individual active regions. Hence, snapshot observations such as we used in this study may under- or overestimate the long-term level of chromospheric activity even though we averaged over many different nights. For more certainty, long-term monitoring is therefore necessary. As an example, we note that the long-term S-average from 1966-1992 \citep{b95} for the Sun is close to 0.18, covering three solar activity maxima {\it and} minima. Around the current relatively weak maximum, we record by contrast an average solar S$_{\rm{MWO}}$ value of only about 0.17. We can rule out that this significant difference is caused by our calibration procedure because we used the same stars for our transformation to Mt. Wilson S$_{\rm{MWO}}$ values as were used by O.C. Wilson and his group. In particular, for modestly active stars with S$<$0.2, the standard deviation of the residuals of our calibrated S$_{\rm{MWO}}$ values compared with those from Baliunas et al. is as small as 0.004, and the mean variation amounts to only 2.3$\%$. Furthermore, during the period when we measured the solar S$_{\rm{MWO}}$ value of 0.17, sunspot counts were also on a lower level than in the solar cycles of 1966-1992. We therefore conclude that this difference is real and represents a real change of the solar activity on a longer (decades to century) timescale. This casts some doubt as to how well the present snapshot S$_{\rm{MWO}}$ value of 0.17 represents the long-term solar activity. However, when modern sunspot records are compared with historic observations, all uncertainties included, (SIDC 2015, communication of July 1 \footnote{http://sidc.oma.be/press/01/welcome.html}) the second half of the past century (when \citet{b95} obtained their S$_{\rm{MWO}}$ value of 0.18 ) appears to have been above-average and so may portray a slightly too young Sun. Interestingly, the snapshot observations of HD~101364 and HD~197027 we presented here show that their activity levels are more or less comparable with that of the current Sun and that of 18~Sco, but all of them clearly differ from 51~Peg, which exhibits a significantly and consistently lower activity. Exclusively from the point of view of stellar activity levels, we may therefore conclude that HD~101364 and HD~197027 are younger in the context of stellar evolution than the more evolved solar-like star 51~Peg, which has already come quite close to the basal flux level of its S-index. Compared to the Sun and 18~Sco, however, it might be assumed that HD~101364 and HD~197027 have a comparable age or are slightly older. Using stellar ages obtained from the comparison of matching evolutionary tracks with the observed stellar HRD positions, we find that 51~Peg is, in absolute terms, not the oldest object in this small sample. However, when the faster MS evolution of 51~Peg is considered, we can show the same trend in the evolutionary advance as the one shown by the observed activity levels (see Fig. \ref{Fig5}) because its mass is higher than those of the other four stars. These results confirm earlier work, that is, a better correspondence between activity level and {\it relative} MS age instead of absolute age, as shown by \citet{Reiners2012ApJ...746...43R} and \citet{schroeder2013A&A554A50S}. Another important characteristic of a star is its rotation period. The rotation period is correlated with the chromospheric activity and with stellar age. This can also be seen in our very small sample of only three periods, which is in sufficient to show this correlation clearly. Nevertheless, we may say that the Sun and 18~Sco are comparable in age, in chromospheric activity, and in their rotation periods. Furthermore, the rotation of the Sun and 18~Sco is clearly faster than that of the evolved star 51~Peg. This finding agrees with the relation between evolutionary (relative) age and activity level shown in Fig.~\ref{Fig5}. It is obvious from Fig. \ref{Fig5} that the same relation holds for the other two solar twins that we discussed here, HD~101364 and HD~197027. Even though this must still be confirmed by extending this study to a larger sample of solar-like stars, we may already conclude that first, the chromospheric activity level is a good age indicator, even for advanced evolutionary states, and that second, the two new solar-twin candidates (HD~101364 and HD~197027) and 18~Sco we discussed here are suitable matches to our Sun with respect to mass, metallicity, activity, and evolutionary age, even if they are perhaps slightly older, and third, in the search for a perfect twin, stellar age should be given more consideration because otherwise the activity level would not match that of the Sun.

1607.01279

1607

1607.06459_arXiv.txt

We explore the relationship between active galactic nuclei and star formation in a sample of 513 optically luminous type 1 quasars up to redshifts of $\sim$4 hosting extremely high star formation rates (SFRs). The quasars are selected to be individually detected by the \textit{Herschel} SPIRE instrument at $>$\,3$\sigma$ at 250\,$\mu$m, leading to typical SFRs of order of 1000\,M$_{\odot}$yr$^{-1}$. We find the average SFRs to increase by almost a factor 10 from $z\sim0.5$ to $z\sim3$, mirroring the rise in the comoving SFR density over the same epoch. However, we find that the SFRs remain approximately constant with increasing accretion luminosity for accretion luminosities above 10$^{12}$\,L$_{\odot}$. We also find that the SFRs do not correlate with black hole mass. Both of these results are most plausibly explained by the existence of a self-regulation process by the starburst at high SFRs, which controls SFRs on time-scales comparable to or shorter than the AGN or starburst duty cycles. We additionally find that SFRs do not depend on Eddington ratio at any redshift, consistent with no relation between SFR and black hole growth rate per unit black hole mass. Finally, we find that high-ionisation broad absorption line (HiBAL) quasars have indistinguishable far-infrared properties to those of classical quasars, consistent with HiBAL quasars being normal quasars observed along a particular line of sight, with the outflows in HiBAL quasars not having any measurable effect on the star formation in their hosts.

Nuclear activity and star formation in galaxies are observed to often coexist across all redshifts \citep{farrah03, alexander05, shi09, hernan09, hatzimi10, mainieri11, lamassa13,harris16,alag16}, up to extremely high active galactic nuclei (AGN) and starburst luminosities \citep[e.g.][]{gen98,carilli01,omont01,farrah02,efst14,magdis14,rosen15}. Moreover, there is a tight correlation between the stellar velocity dispersion in the bulge and the mass of the supermassive black hole (BH) residing in the centre for nearby quiescent galaxies (e.g. \citealt{magorrian98,ferrarese00,tremaine02,haering04}). These observations suggest that there exists a deep connection between stellar and black hole mass assembly events in galaxies. The nature of, and connection between, star formation and AGN activity in galaxies is unclear. For starbursts, the majority of star-forming systems lie on a `main sequence' whose mean star formation rate (SFR) rises with redshift, from about 10\,M$_\odot \mathrm{yr}^{-1}$ at $z=0.5$ to roughly 100\,M$_\odot \mathrm{yr}^{-1}$ at $z=2$ \citep[as in e.g.][]{elb11,rodig11,schreiber15}. The origin of this main sequence is however still debated, as is the trigger for star formation as a function of a galaxy's position relative to the main sequence. For AGN, it is difficult to complete a robust census of AGN from surveys, since they are obscured by dust and gas for a significant fraction of their lives, with this fraction possibly dependent upon both luminosity and redshift \citep[e.g.][]{mart05}. It is also unclear if star formation events and AGN activity can directly affect one another. A direct relation is motivated by models for galaxy assembly to improve consistency between predictions and observations, most often via `quenching' of star formation by an AGN \citep[e.g.][]{bower06, croton06, booth09, fabian12}. However, observational studies of quenching remain inconclusive. Indirect manifestations of feedback, such as large molecular gas outflows \citep[e.g.][]{feruglio10,spo13} and powerful AGN-driven winds \citep[e.g.][]{perna15}, are found routinely (see also e.g. \citealt{farrah09,bridge13,spo13}), but studies that claim a causal relation are rarer \citep{farrah12,page12}, and sometimes controversial \citep[e.g.][]{harrison14}. As a result, the scaling relations between star formation and AGN activity across their respective duty cycles remain uncertain. Some authors find a scaling between SFR and AGN luminosity \citep[e.g][]{ima11,Young2014,delv15}, while others do not \citep[e.g.][]{shao10,mullaney12,harrison14,ma15,stanley15}. Recently, \citet{harris16} have shown that a correlation between SFR and AGN luminosity may only exist over certain AGN luminosity, SFR and redshift ranges. An insightful way to study the connections between star formation and AGN activity in the context of galaxy assembly events is to examine their scaling relations in populations that signpost specific regions of the AGN luminosity and SFR parameter space. One such population are optically luminous type 1 quasars that host luminous, off-main sequence starbursts. This population is straightforward to find in optical surveys, corresponds to a specific phase in the AGN duty cycle, and may signpost the extremes of the starburst duty cycle. As such, they may illustrate how the processes that convert free baryons to stellar and BH mass may change at the most extreme luminosities. They are also an excellent population in which to search for the initial stages of AGN feedback. A related population within which it would appear intuitive to search for evidence of AGN feedback is that of the Broad Absorption Line (BAL) quasars \citep{lynds67,turnshek84}. The BAL quasars possess broad P-Cygni-like absorption features in their ultraviolet spectra that are blue-shifted with respect to the nominal wavelength of the emission line. These features may stem from outflows from the quasars and may further be associated with high mass-loss rates \citep{dekool02, chartas03}. These outflows could arise in two ways; the outflows could be a random process, present for only a random fraction of the quasar lifetime and/or over certain viewing angles \citep[e.g.][]{elvis00}. The observation of BALs in only 10-15 per cent of quasars \citep[e.g.][]{hewett03} then arises from a suitable combination of these two factors. Second, BALs may signpost outflows that occur only at a particular point in a quasar's lifetime, most likely `young' objects recently (re)fuelled by mergers that also trigger star formation in their hosts. It is this second scenario that would mean that BAL quasars may signpost quasar-mode feedback. This may manifest itself via BAL quasars having different far-infrared (FIR) properties, on average, to those of ordinary quasars. In this paper, we undertake such a study, exploring the relationship between star formation and AGN activity in optically luminous quasars over $0.5<z<4$ that host extremely high SFRs, from 40$-$4000\,M$_\odot$yr$^{-1}$, placing these hosts approximately one dex above the star formation main sequence. To do so, we start with quasars in the Sloan Digital Sky Survey (SDSS) that also lie within extragalactic survey fields covered by the Spectral and Photometric Imaging Receiver \citep[SPIRE;][]{griffin10} instrument onboard \textit{Herschel} \citep{pilbratt10}. We then restrict ourselves to those quasars that are individually detected by \textit{Herschel}. We infer SFRs from the \textit{Herschel} data, and compare these to the properties of the AGN, as measured from the SDSS catalogue data. We also examine the properties of BAL quasars to see if their properties show evidence for AGN feedback. The paper is structured as follows. The SDSS and {\it Herschel} data, as well as the matched quasar catalogue, are described in Sec. \ref{sec:data}. Further, Sec. \ref{sec:analysis} describes the spectral energy distribution (SED) fitting, the SFR derivation and uncertainties and discusses the way in which the current study complements previous works in terms of samples and methodology. Sec. \ref{sec:results_full} presents the SFRs in the full quasar sample and in the BAL quasar sub-sample separately, as a function of the quasars' intrinsic properties, namely accretion luminosity, BH mass and Eddington ratio. Lastly, Sec. \ref{sec:discussion} discusses our results and places them in a greater context. Throughout this work, we assume $H_0 = 72 \hspace{0.1cm} \textnormal{km} \hspace{0.1cm} \textnormal{s}^{-1} \hspace{0.1cm} \textnormal{Mpc}^{-1}$, $\Omega_\Lambda = 0.7$ and $\Omega_\textnormal{M}=0.3$. \begin{table*} \begin{center} \begin{tabular}{l C{2cm} crrrrrrr} \hline &&& \multicolumn{3}{c}{SDSS quasars} && \multicolumn{3}{c}{{\it Herschel}/SDSS quasars} \\ \cline{4-6} \cline{8-10} \noalign{\vskip0.025cm} Field & Overlapping Region (deg$^2$) && Total & In DR7 & In DR10 && Total & In DR7 & In DR10\\ \hline \vspace{-0.25cm} &&&&&&&&\\ HerMES & 60 && 1,293 & 586 & 739 && 83 & 43 & 44\\ HerS & 79 && 4,524 & 2,154 & 2,779 && 212 & 150 & 90\\ HeLMS & 210 && 8,827 & 2,488 & 6,796 && 218 & 92 & 144\\ \hline \end{tabular} \caption{Quasar detection statistics by survey. The overlapping region is the area of each \textit{Herschel} field that is also covered by SDSS. The SDSS detections are those quasars that lie in their respective \textit{Herschel} fields. The joint detections columns represent those SDSS quasars with SPIRE detections above 3$\sigma$ at 250$\micron$. We further describe these objects by the data release in which they are found: DR7 or DR10. The sum of the DR7 and DR10 columns is always higher than the columns showing the totals, as the latter only consider unique entries. The SDSS detections columns represent those SDSS quasars located in each respective \textit{Herschel} field, while the joint detections are the quasars with 250\,$\micron$ detections at $>$3$\sigma$.} \label{tab:jointDetections} \end{center} \end{table*} \begin{figure} \begin{center} \vspace{-.2cm} \includegraphics[width=8cm]{Mi_z_withHisto.pdf} \vspace{-2.5cm} \caption{Absolute \textit{i}-band magnitude versus redshift. The red circles represent those objects included in SDSS DR7, while the blue squares are those objects in DR10. The filled and open shapes differentiate between the BAL and non-BAL samples, respectively (see Section \ref{sec:BAL}). To characterise these BAL and non-BAL subsets we use the objects that lie within the black box, with its limits defined by 1.5\,$\le$\,\textit{z}\,$\le $\,3.3 and $-28.5$\,$\le$\,$M_i$\,$\le$\,$-24.3$, in which most of the BAL quasars lie. The redshift histograms are shown at the top, with the filled red, dashed blue and solid black showing the DR7, DR10 and full sample of objects, respectively. The green spikes represent the redshift distribution of the BAL quasars.} \label{fig:MI_z} \end{center} \end{figure}

\label{sec:discussion} The relationship between SFRs and AGN properties gives insights into the co-evolution of these activities. While we do not have stellar masses for our sample, their extremely high SFRs coupled with reasonable assumptions for quasar host stellar masses at this epoch mean they almost certainly lie above the `main sequence' for star formation at their respective epochs. In the following, we consider the relationship between SFRs and both redshift and AGN properties. Finally, we discuss the implications for BAL quasars. \subsection{Evolution with redshift}\label{sec:disczevo} We find that the SFRs of our sample increase by a factor of up to $\sim$\,8 from $z<1$ to $z \ge 2$. We do not see any evidence for a rise in SFRs with redshift beyond $z\sim2$. This is consistent with observations of the global comoving SFR density, which increases by a factor of at least ten over $0<z<1$ and peaks around $z=2$ \citep[e.g.][]{reddy08, magn09, karim11, wuyts11, bethermin13, burgarella13, wang13, chen16}. Thus, in terms of evolution over $0\lesssim z \lesssim 3$, extremely luminous star formation events in quasar hosts appear to evolve, on average, in a manner consistent with the global star formation history as derived from the general galaxy population. Our results are also consistent with those of MY15, who find that the fraction of IR-luminous quasars peaks at $z\sim2$, and with H16, who find that SFRs in quasars do not change appreciably as a function of redshift over $2<z<3$ (but see also e.g. \citealt{netzer16}). They are also consistent with several previous studies of the evolution with redshift of star formation in AGN hosts \citep[e.g.][]{serjeant09,serjeant10,shao10,bonfield11,mullaney12}. These previous studies mostly sample different parts of the $z$-SFR-$L_\textnormal{acc}$ parameter space for quasars to our study. For example, the sample in \citealt{mullaney12} spans a similar redshift range but has lower $L_\textnormal{acc}$ values and SFRs by $\gtrsim 2$ orders of magnitude, while the sample in H16 has comparable $L_\textnormal{acc}$ values, but lies at higher redshifts with a factor $\sim 3$ lower SFRs. The apparent consistency of our results with theirs on the evolution of SFRs with redshift in quasar hosts is thus remarkable. This consistency suggests that the physical processes that lead to star formation in AGN hosts, such as mergers or secular evolution, do so in a way that gives a similar evolution with redshift for star formation events spanning tens to thousands of $M_{\odot}$ per year and accretion luminosities $\gtrsim10^{10}$\,L$_\odot$. We caution however that this does not account for selection effects between studies, and only applies to time-scales significantly longer than the AGN and starburst duty cycles. \subsection{Star formation and accretion luminosity}\label{sec:discsfagn} We first consider the relation between SFR and $L_\textnormal{acc}$. The top panel of Fig. \ref{fig:SFR_trends} is consistent with an approximately flat $\textnormal{SFR} - L_\textnormal{acc}$ relation, in all redshift bins. This is in agreement with some previous studies \citep[e.g][]{priddey03, shao10, dicken12, mullaney12, harrison12, rosario13, ma15, baner15, stanley15}, but not others \citep[e.g.][]{netzer09, hatzimi10, ima11, rafferty11, mull12a, chen13, Young2014, delv15, xu15, bian16}. Several explanations for an observed lack of a relation between SFR and $L_\textnormal{acc}$ have been proposed. They fall, broadly, into four categories. First, the strength of the relation depends on redshift, with an observable relation only emerging at $z\gtrsim2$ \citep{hatzimi10, rovilos12, delv15, harris16}. This could arise due to, for example, a higher free gas fraction at higher redshifts. Second, AGN luminosities can vary substantially on time-scales of days to months. If AGN luminosities are measured using methods that are particularly sensitive to such variations, such as X-ray observations \citep{barr86,ulr97,grupe01,chitnis09}, then underlying trends could be masked \citep{gabor13, volon15}. Third, it is possible that at very high SFRs and/or $L_\textnormal{acc}$ values, we are sampling an intrinsically different quasar population, one in which e.g. a different trigger for star formation leads to a weaker or absent correlation between SFR and $L_\textnormal{acc}$. Fourth, at high SFRs or high $L_\textnormal{acc}$, the properties of the starburst become decoupled from those of the AGN, plausibly because internal self-regulation processes become the dominant effect in regulating luminosities on time-scales comparable to or shorter than the AGN or starburst duty cycles. The first possibility is, {\itshape on its own}, not plausible for our sample, as our sample spans a wide range in redshift and we do not see a significant correlation in any redshift bin. The second seems unlikely, since our sample is large and has AGN luminosities measured from optical, rather than X-ray data. We note however that there exists a conceptually similar possibility - that the duty cycles of starburst and AGN activity are sufficiently mismatched in duration that no correlation is observed between their luminosities. The available constraints on the duty cycle lengths of starburst and AGN phases are however not accurate enough for us to comment on this possibility. The third possibility also seems unlikely, since we found no differences between the {\it Herschel}/SDSS sample and the general SDSS quasar population in Section \ref{sec:sdss}, though we cannot formally exclude it. The fourth scenario, however, is plausible, since our sample harbours the highest SFRs seen in quasars, where self-regulation may be expected to be seen, if it occurs. Moreover, it is consistent with observations of systems with comparable SFRs at $z>1$, i.e. submillimetre galaxies (SMGs, \citealt{barger14}), which exhibit a flattening in their SFRs at extremely high SFR values. An interesting corollary comes from the work of MY15, who find that star-forming regions in IR-luminous quasars do not increase in size at very high SFRs, but instead may increase in SFR {\itshape density}. An increase in SFR density could mean that starburst self-regulation becomes more effective, via e.g. approaching the Eddington limit for star formation, where winds from supernovae can expel free gas and thus stall star formation \citep[e.g.][]{thom05,murray05,dia12}. We thus propose that the flat relation between SFR and $L_\textnormal{acc}$ in our sample arises because self-regulation by the starburst becomes the dominant factor controlling the SFRs. We cannot, however, rule out a contribution from redshift-driven effects, or from mismatches in the lengths of duty cycles of starburst and AGN episodes. The top panel of Fig. \ref{fig:SFR_trends} also shows no downturn in SFRs at the highest $L_\textnormal{acc}$ values. This is, in principle, inconsistent with the idea that luminous AGN can quench star formation in their host galaxies (see also e.g. \citealt{vil16}). An alternative interpretation is however possible; {\itshape assuming that AGN feedback occurs}, this instead suggests that $L_\textnormal{acc}$ (as derived from rest-frame UV observations) is not a good proxy for the strength of AGN feedback. This is consistent with AGN feedback being a brief phase in the AGN duty cycle, signposted by properties other than AGN luminosity \citep[e.g.][]{farrah12}. \subsection{Star formation, black hole masses, and Eddington ratios}\label{sec:discsfagn2} The middle panel of Fig. \ref{fig:SFR_trends} is consistent with there being no relation between SFR and $M_\textnormal{BH}$. This lack of a correlation is, in some senses, surprising. For example, it has been suggested that SFRs should scale with $M_\textnormal{BH}$ in quasars since SFRs should scale with the total available free gas mass or its large-scale distribution \citep{cen15}, and both the free gas mass and $M_\textnormal{BH}$ should scale with the total baryon density. Thus, on $\gtrsim100$\,Myr time-scales, SFRs and $M_\textnormal{BH}$ in quasars may be expected to correlate since a larger free baryon reservoir would favor both higher current SFRs and higher $L_\textnormal{acc}$ {\itshape in the past}. We, however, do not see such a correlation. A plausible explanation comes from considering our results in context with those of H16, who found that SFRs in quasars at $2<z<3$ do correlate with $M_\textnormal{BH}$ (and $L_\textnormal{acc}$), but only up to an SFR of $\sim600$\,M$_{\odot}$yr$^{-1}$. Above this SFR, they find no correlation. Our SFRs are mostly higher than those in H16 and are measured individually, rather than stacked as in H16, so it is interesting to compare their results to ours. In the top and middle panels of Fig. \ref{fig:SFR_trends}, we plot the relations that H16 derive between SFR and $L_\textnormal{acc}$ and $M_\textnormal{BH}$, respectively. For the SFR-$L_\textnormal{acc}$ plot our results are, given the sizes of the error bars on the average points in each redshift bin, consistent with the H16 relation. For the SFR-$M_\textnormal{BH}$ plot the consistency is weaker, but fig. 17 of H16 shows, plausibly, a flattening in the SFR-$M_\textnormal{BH}$ relation at about $600$\,M$_{\odot}$yr$^{-1}$ as well. Overall therefore, we argue that our results and those in H16 are both consistent with the idea that SFRs in quasars correlate with $L_\textnormal{acc}$ and $M_\textnormal{BH}$, but only up to an SFR when internal `self-regulation' by the starburst becomes important. The bottom panel of Fig. \ref{fig:SFR_trends} is consistent with there being no relation between SFR and $\lambda_\textnormal{Edd}$, at any redshift. This result is also in line with those in H16, who also find no evidence for an $\textnormal{SFR} - \lambda_\textnormal{Edd}$ relation, at any SFR (modulo possible enhanced SFRs at low $\lambda_\textnormal{Edd}$ values. Overall, this is consistent with there being no relation between SFRs in quasars and how efficiently the black hole is accreting, at {\itshape any} SFR. It is also consistent with the suggestion that $\lambda_\textnormal{Edd}$ values in quasars do not depend strongly on redshift, at least at $z\gtrsim0.5$ (e.g. \citealt{shan13,capl15}, see also \citealt{fine06,bluck11,luss12}). \subsection{Star formation in broad absorption-line quasars }\label{sec:dischibal} We find that there is no difference in the SFR, $L_\textnormal{acc}$, $M_\textnormal{BH}$, and $\lambda_\textnormal{Edd}$ distributions between the HiBAL and non-BAL quasars in our sample; all are consistent with being drawn from the same parent population. A minor caveat to this result is that our sample may be incomplete for HiBALs at very high redshifts (Sec. \ref{sec:sdss}), an effect that we cannot account for using available data; we think it unlikely that this incompleteness could be a significant factor leading to the lack of differences between BAL and non-BAL quasars that we observe. This result aligns with previous studies, which find no differences between HiBAL and classical quasars, at any redshift \citep{priddey03,priddey07,gallagher07,pu15,harris16}. In particular, it is consistent with the study of \citealt{cao12}, who also use {\it Herschel}-SPIRE data to compare the properties of BAL versus non-BAL quasars, from the H-ATLAS survey, and find no differences between the two populations. The two studies, however, sample different regimes in SFR; our study examines quasars with typical SFRs of $\sim1000$\,M$_{\odot}$yr$^{-1}$, whereas the typical SFRs of the quasars in \citealt{cao12} are $\sim240$\,M$_{\odot}$yr$^{-1}$. The combined results suggest that BALs are not seen preferentially in certain $L_\textnormal{acc}$ or $M_\textnormal{BH}$ regimes, and not over specific ranges in SFR, even for those quasars harbouring the most luminous star formation events seen in any quasar at any redshift. It has been argued that BAL quasars may be a promising quasar population within which to look for evidence of AGN feedback, since the BAL winds provide a natural mechanism to couple momentum from accretion-disk winds to ISM gas. The lack of any differences between the HiBALs and the non-BALs in our sample however argues that HiBAL quasars are, as a population, not sites for AGN feedback, unless the AGN feedback phase is much shorter than the lifetime of the BAL winds. Instead, our results are consistent with HiBAL quasars being normal quasars observed along a particular line of sight, with the outflows in HiBAL quasars not having any measurable effect on the star formation in their hosts (see also e.g. \citealt{vio16}).

1607.06459

1607

1607.06804_arXiv.txt

T~CrB is a symbiotic recurrent nova known to exhibit active phases, characterised by apparent increases in the hot component temperature and the appearance of flickering, i.e. changes in the observed flux on the time-scale of minutes. Historical UV observations have ruled out orbital variability as an explanation for flickering and instead suggest flickering is caused by variable mass transfer. We have analysed optical and X--ray observations to investigate the nature of the flickering as well as the active phases in T~CrB. The spectroscopic and photometric observations confirm that the active phases follow two periods of $\sim$1000d and $\sim$5000d. Flickering in the X--rays is detected and follows an amplitude--flux relationship similar to that observed in the optical. The flickering is most prominent at harder X--ray energies, suggesting that it originates in the boundary layer between the accretion disc and the white dwarf. The X--ray radiation from the boundary layer is then reprocessed by a thick accretion disc or a nebula into UV radiation. A more detailed understanding of flickering would benefit from long-term simultaneous X--ray and optical monitoring of the phenomena in symbiotic recurrent novae and related systems such as Z~And type symbiotic stars.

Cataclysmic variables (CVs) are interacting binaries in which a white dwarf (WD) is accreting matter from a late--type donor. In the classical picture of a CV a Roche lobe filling star transfer mass through an accretion disc to the WD. In some CVs the material accreted onto a WD reaches a pressure and temperature sufficient to trigger a thermonuclear reaction, which gives rise to a classical nova outburst. If more than one outburst is observed, then it is classified as a recurrent nova (RN; see e.g. \citealt{2003cvs..book.....W}). If the donor star is also a red giant (RG), it is classified as a symbiotic recurrent nova (SyRN; see e.g. \citealt{1999A&A...344..177A}; \citealt{2011A&A...527A..98S}), with the system also belonging to the family of symbiotic stars (SySt; see \citealt{2012BaltA..21....5M} for a recent review). T~CrB is a SyRN with two outbursts recorded to date (Nova CrB 1866, 1946). The RG in the system has a spectral type M4III \citep{1999A&AS..137..473M} and fills its Roche Lobe \citep{1998MNRAS.296...77B}. The masses of the components are M$_{\mathrm{WD}}=1.2\pm0.2$M$_\odot$ and M$_{\mathrm{RG}}=0.8\pm0.2$M$_\odot$ \citep{1998MNRAS.296...77B,2004A&A...415..609S}. The average quiescent luminosity of the WD is close to 40~L$_\odot$ \citep{1992ApJ...393..289S}. T CrB was discovered as an X--ray source by \citet{1981ApJ...245..609C} and the X--ray emission has a relatively hard spectrum for a symbiotic star \citep{2008ASPC..401..342L}. Occasionally between the nova eruptions, T~CrB and other SyRNe (V745~Sco, V3890~Sgr and RS~Oph) show enhanced activity of the hot component, manifested by an increase in the emission line fluxes and appearance of a blue continuum. In particular, the \mbox{He\,{\sc i}} and \mbox{He\,{\sc ii}} emission lines became clearly visible in the spectrum whereas there are very weak or absent during most of the quiescent phase \citep{1990JAVSO..19...28I}. The line variability is correlated with the changes in the UV continuum and there is no correlation with the orbital period \citep[e.g.][]{1985ESASP.236..213C}. Similar activity was reported in the episodic appearance of the \mbox{He\,{\sc ii}\,4686} emission line and a hot continuum in the optical spectrum before the 1946 outburst (e.g. \mbox{\citealt{1939ZA.....17..246H}}; \citealt{1943PASP...55..101M}). The nature of these active phases is yet fully understood. One of the most mysterious phenomena observed in T~CrB is flickering. Flickering is a stochastic variation in the light curve of a star with an amplitude of few tenths of a magnitude on time-scales of seconds or minutes. It is well known that flickering in CVs is associated with accretion onto the WD, but the exact physical process is unknown. One of the first models of flickering included unsteady accretion through a bright spot \citep{1971MNRAS.152..219W}. The first to conduct a systematic study was \citet{1992A&A...266..237B} who estimated theoretical energies and time-scales of flickering. The author considered unstable mass transfer and interaction of matter with the disc edge, dissipation of magnetic loops, turbulence in the accretion disc and unstable accretion in the boundary layer. As a result the most promising model contained unstable mass transfer through a boundary layer and turbulence in the accretion disc (see also \citealt{1993A&A...275..219B}). Different models of flickering include magnetohydrodynamic turbulence transporting the angular momentum outward \citep{1998RvMP...70....1B}, turbulent transport of angular momentum in the whole disc \citep{2010MNRAS.402.2567D}, occasional flare-like events and subsequent avalanche flow in the accretion disk atmospheres \citep{1997ApJ...486..388Y}, discrete flares in the accretion disc \citep{2006Ap&SS.304..291R}, or a more recent model by \citet{2014MNRAS.438.1233S} which is mainly based on a fluctuating accretion disc. The flickering in T~CrB was observed by a number of authors (see \citealt{2006AcA....56...97G} and references therein). The colours of the flickering source in T~CrB indicate a temperature of $\sim$9000~K, which is significantly lower than the average temperature of flickering sources in classical CVs ($\sim$20000~K; \citealt{2015AN....336..189Z}). Flickering in T~CrB seems to originate in the vicinity of the WD \citep{1998A&A...338..988Z}. The ratio of the amplitude of flickering in T~CrB to the average flux remains constant \citep{2004MNRAS.350.1477Z}, which according to a model by \citet{1993A&A...275..219B}, means that the size of the boundary layer remains roughly constant. This is not necessarily the case in all CVs \citep{1998A&A...332..586F}. Flickering was also observed in the X--ray observations of T~CrB \citep{2008ASPC..401..342L} and in the emission lines \citep{2005PASP..117..268Z}. In this paper we analyse optical and X--ray observations of T~CrB in order to shed light on the nature of active phases and flickering observed in this system. The collected observations are presented in section~\ref{obs_sec}. In section~\ref{active_phase_section} we present results concerning the active phases of T~CrB. Section~\ref{flickering_section} is dedicated to the optical and X--ray observations of flickering in the system. A brief discussion about the nature of the observed phenomena is presented in section~\ref{discussion_sec} and a summary of the results is given in section~\ref{sumarysec}.

\label{discussion_sec} \subsection{Active phases} Similar active phases to those observed in T~CrB have been discovered in the other SyRN RS~Oph \citep{2008ASPC..401..219G}. Namely, the authors observed brightenings of the system on time-scales of 1200--1800d with an amplitude of $\sim$1~mag. \citet{2008ASPC..401..219G} pointed out that the behaviour of RS~Oph outside of the nova outbursts resembles the classical Z~And variability, which is thought to be a result of an instability in the accretion disc (\citealt{2003ASPC..303....9M}; \citealt{2006ApJ...636.1002S}). The active phases in T~CrB and RS~Oph could have the same origin as in the Z~And type symbiotic systems with the difference that in the Z~And type systems there is steady hydrogen burning on the surface of the WD. If the active phases in T~CrB and RS~Oph are driven by the same mechanism as the multiple-outburst activity of Z~And type SySt, similar time-scales should be expected. In fact, outbursts in many Z~And type systems recur with quasi-periods of one to tens of years. Namely, the major outburst in AG~Dra occurs every 12--15 years, while minor outbursts are observed roughly once a year \citep{2016MNRAS.456.2558L}. Z~And showed a time-scale of $\sim$8400d of the outburst activity \citep{1994A&A...292..534F} and the interval between outbursts in BF~Cyg is $\sim$6376d \citep{2006MNRAS.366..675L}. Given the differences between the systems these time-scales are in rough agreement with the ones in T~CrB and RS~Oph. A useful comparison can also be made based on similar physical parameters of these systems. In the prototype star Z~And, after assuming M$_{\mathrm{WD}}\sim$0.65M$_\odot$ \citep{1997A&A...327..219S}, a hot component mass function of f(m)=0.0240M$_\odot$ \citep{2000AJ....120.3255F}, an inclination of i=$41^\circ$ \citep{2010AJ....140..235I}, and orbital period of 759d \citep{2000AJ....120.3255F}, the radius of the Roche Lobe around the WD is $\sim$140R$_\odot$. Similarly, for T~CrB, assuming M$_{\mathrm{WD}}=1.2$M$_\odot$ and M$_{\mathrm{RG}}=0.8$M$_\odot$ \citep{1998MNRAS.296...77B,2004A&A...415..609S}, and an orbital period of 227.67d \citep{1988AJ.....95.1505L}, the radius of the Roche Lobe around the WD is $\sim$80R$_\odot$. In the case of RS~Oph, taking M$_{\mathrm{WD}}=1.3$M$_\odot$, M$_{\mathrm{RG}}=0.74$M$_\odot$, and an orbital period of 453.6d \citep{2009A&A...497..815B}, the radius of the Roche Lobe around the WD is $\sim$135R$_\odot$. The estimated mass transfer rate in T~CrB is 2.5$\times 10^{-8}$~M$_\odot$/yr \citep{1992ApJ...393..289S}, whereas in Z~And the mass transfer rate varies between 4.5$\times 10^{-9}$~M$_\odot$/yr during quiescence \citep{1988ApJ...324.1016F} to 3.2$\times 10^{-7}$~M$_\odot$/yr during the active phase \citep{2004A&A...428..985T}. Therefore, the mean mass transfer rates are of the same order of magnitude in both systems. The ratio of Roche Lobe radii around the WDs in Z~And and T~CrB is $\sim$1.5. It is interesting to note that this value is similar to the ratio of $\sim$1.7 between time-scales of Z~And outburst activity and big active phases in T~CrB. On the other hand, the ratio of Roche Lobe radii around the WDs in RS~Oph and T~CrB is $\sim$1.7, which is similar to the ratio of $\sim$1.5 between time-scales of active phases in RS~Oph and small active phases in T~CrB. This is to be expected, because the time-scale of a disc instability scales with the size of the accretion disc, which in turn is proportional to the radius of the Roche Lobe around the WD. This is more consistent with the longer time-scale of outburst activity in Z~And than the time-scale of active phases in T~CrB. Recently, disc instability in RS~Oph has been modeled by \citet{2011MNRAS.418.2576A}, who found recurring outbursts in the system on 10--20 yr time-scales, which is comparable to the time-scales of active phases actually observed in RS~Oph and Z~And \citep{2008ASPC..401..219G}. In some models alternating small and big outbursts take place. This could be associated with big and small active phases in T~CrB. \citet{1997ppsb.conf..117A} proposed a solar--type cycle resulting in a variable mass transfer as an explanation of the active phases in T~CrB. In order to obtain solar-like cycles one must have a high magnetic field. In the case of T~CrB this is probable since it is argued that in SySt the RG is more magnetically active than a single RG \citep{2002MNRAS.337.1038S}. In general, the hypothesis that the RG is a source of the observed activity is particularly interesting since variability with a time-scale of $\sim$1000d was observed in a single semi--regular variable WZ~Cas \citep{2005A&A...440..295L}. Overall it seems to be clear that at least the big active phases are related to a change in mass transfer rate. Therefore it is important to determine the exact nature of the T~CrB variability, especially since SySt are thought to be promising SNIa progenitors \citep{2013IAUS..281..162M}. Recently, \citet{2016NewA...47....7M} have claimed that during 2015 T~CrB displayed ''super-active conditions never seen before''. They have also suggested that this super-active state is a new form of activity and it is different from the active phases observed e.g. by \citet{1990JAVSO..19...28I}. We argue that this is not true, and the activity observed by \citet{2016NewA...47....7M} is just a new maximum of the big active phases (see Sec.~\ref{active_phase_section}) and similar activity has been observed in the past. \citet{2016NewA...47....7M} based their claim on a large increase in emission line fluxes and optical magnitudes of the system. However, as the authors noted, it is difficult to study the history of emission line variability, since the observations are sparse and the authors rarely have given line fluxes. Here we expand the analysis of \citet{2016NewA...47....7M} using EWs. From Fig.~\ref{all_fluxes} one can see that variability of emission line fluxes is well represented by variability of EW(H$\alpha$). In particular, the $H\alpha$ emission line flux, estimated on the spectrum from MJD~57447, is F($H\alpha$)=4.5$\times$10$^{-11}$~erg~cm$^{-2}$~s$^{-1}$ , and the corresponding corresponding EW($H\alpha$)=48.8. This is much higher than the maximum flux F($H\alpha$)=3.9$\times$10$^{-11}$~erg~cm$^{-2}$~s$^{-1}$ observed by \citet{2016NewA...47....7M}. Therefore we adopt EW($H\alpha$)=48.8 as a maximum EW observed during the supposed super-active conditions. During the past big active phases, the EW($H\alpha$) varied from $\sim$5 to $\sim$40 \citep[see fig.~4 in][]{2004A&A...415..609S}. Moreover, the maximum EW observed during the 1989 big active phase, EW(H$\alpha$)$\sim$40, was bigger than that observed during the 1997 big active phase (EW(H$\alpha$)$\sim$30). This shows that the amplitude of emission line variability changes, which is consistent with the quasi-periodic nature of the known phenomenon rather than a new form of activity. This is further supported by the photometric variability observations reported by \citet{2016NewA...47....7M} with timescales and amplitudes consistent with previous big active phases \citep[see fig.~4 in][]{2004A&A...415..609S}. \subsection{Flickering} Since the flickering in T~CrB is mainly observed in the hard X--rays, which are thought to originate in the accretion disc boundary layer, our study seems to confirm the model of flickering originating from an unstable mass transfer through the boundary layer. However, it seems that more than one mechanism is responsible for the observed variability. The events in which there is brightening in the 5--10~keV range and fading in the 10--20~keV range seem to indicate Compton cooling. Moreover, it is obvious that the flickering observed in the optical range is not the same as that observed in the X--rays since the optical colours of the flickering source in T~CrB are consistent with a blackbody of $\sim$9000~K \citep{2015AN....336..189Z}. Nevertheless, it seems that flickering in the hard X--rays is related to the optical flickering since the maximum in hard X--rays was observed close to the maximum of the optical $\sim$1000d active phase, and for both wavelengths the amplitude of the flickering--mean flux relation holds. Our results seem to be best reproduced by a model proposed by \citet{2014MNRAS.438.1233S} in which the flickering observed in the visible range comes from X--ray reprocessing by a geometrically thick disc. This model provides a natural explanation for the difference between the flickering source colours in SySt and classical CVs \citep{2015AN....336..189Z} as the accretion disc in SySt is thought to be bigger than in CVs. Reprocessing of the X--ray radiation by the nebula is in principle not excluded, but reprocessing by a thick accretion disc seems likely since the UV to IR SED of the system was successfully modelled by \citet{2004A&A...415..609S} using a model containing an optically thick accretion disc. On the other hand, reprocessing of the radiation by a nebula would be consistent with flickering observed in emission lines of SySt \citep{2005PASP..117..268Z,2014MNRAS.442.2637W}.

1607.06804

1607

1607.02616_arXiv.txt

The possible short gamma-ray burst (GRB) observed by {\it Fermi}/GBM in coincidence with the first gravitational wave (GW) detection, offers new ways to test GRB prompt emission models. Gravitational wave observations provide previously unaccessible physical parameters for the black hole central engine such as its horizon radius and rotation parameter. Using a minimum jet launching radius from the Advanced LIGO measurement of GW~150914, we calculate photospheric and internal shock models and find that they are marginally inconsistent with the GBM data, but cannot be definitely ruled out. Dissipative photosphere models, however have no problem explaining the observations. Based on the peak energy and the observed flux, we find that the external shock model gives a natural explanation, suggesting a low interstellar density ($\sim 10^{-3}$ cm$^{-3}$) and a high Lorentz factor ($\sim 2000$). We only speculate on the exact nature of the system producing the gamma-rays, and study the parameter space of a generic Blandford Znajek model. If future joint observations confirm the GW-short GRB association we can provide similar but more detailed tests for prompt emission models.

\label{sec:intro} With the first detection of GWs we entered a new era in astrophysics \citep{Abbott+16gw1}. Electromagnetic counterparts are crucial for establishing the astrophysical context for the GWs and also for a more accurate localization to aid subsequent follow-up \citep{Connaughton+15loc}. GRB progenitors \citep[see ][for reviews]{Meszaros+14review, Kumar+15review} have been leading candidates for sources of GWs \citep{Kobayashi+03gwgrb, Corsi+09mag}. The most widely considered GW sources are compact binary mergers with components stemming from a combination of neutron stars (NS) or black holes (BH). Other than BH-BH mergers, substantial radiation is expected to accompany the GW signal, and indeed, the leading candidate for short-hard GRBs are merging neutron stars \citep{paczynski86,eichler89}. The GW~150914 event is best explained by the merger of two $\sim$30 M$_\sun$ black holes. {\it Fermi}/GBM detected a tantalizing counterpart, GW~150914-GBM \citep{Connaughton+16gbmgw}, consistent with a weak short GRB, broadly consistent with the GW location and temporally coincident with the GW signal (offset of $\Delta t_{\gamma-{\rm GW}} \approx t_{\rm GRB}-t_{\rm GW}$=0.4 s). {We note however that while Advanced LIGO and GBM locations are consistent, they both span a significant portion of the sky ($\sim$600 square degrees for Advanced LIGO at 90\% confidence level and $\sim$3000 square degrees for GBM at 68\% confidence level).} This observation potentially marks the beginning of the multi-messenger astrophysics. In this paper we assume the weak GBM burst is a GRB (we refer to it as GW~150914-GBM) associated with GW~150914 and investigate its implications for the physical parameters of the system and for its surroundings. This joint electromagnetic (EM), GW observation was already addressed in a significant number of early studies covering aspects of EM energy extraction from a binary BH system and its surroundings \citep{Li+16gw, Loeb16gw,Perna+16gw, Fraschetti16gw, Yamazaki+16gw, Zhang16gw}. {INTEGRAL ACS observations \citep{Savchenko+16gwgrb} set a constraining upper limit in terms of the source fluence in the ACS energy range (above $\sim$75 keV). The uncertainties on the GBM spectral parameters and on the direction of the possible source, however, weaken any tension between the two measurements \citep[for details, see Section 3.3 of ][]{Connaughton+16gbmgw}.} {Nonetheless,} we emphasize that the association between the GBM event and GW~150914 might have occurred by chance. However, because the false alarm probability of the two events being associated is P=$2.2\times10^{-3}$ \citep{Connaughton+16gbmgw}, we will {\it assume} a common origin and venture to discuss the implications for the GRB emission models. There has been considerable uncertainty in the GRB prompt emission model parameters, such as the compact object mass and rotation rate. For the first time however we can use realistic input parameters for modeling the BH central engine, because the gravitational wave observations yield these parameters to a precision that was previously unavailable. We calculate, to the extent that the gamma-ray observations allow, the constraints on the usual GRB models that can be placed. Jets and black hole central engines are thought to be ubiquitous in GRBs. Energy released from the central engine becomes collimated either by magnetic stresses or the ram pressure of a progenitor star. The initial dynamics of the jet are determined by the launching radius, the size of the base of the jet, where the Lorentz factor of the matter, which eventually produces the GRB, is around unity. In other words the launching radius ($R_0$) is the characteristic size of the volume in which energy is deposited. It is beyond this radius that the jet starts to accelerate. Current methods of determining the launching radius rely on the blackbody components in the GRB spectrum \citep{Peer+07lorentz}. \citet{Larsson+15jet} found the launching radius for GRB 101219B is approximately 10 times the horizon radius. This suggests the launching radius is defined by the scale of the BH central engine rather than larger scales ($\gtrsim 10^9 \cm$) such as the progenitor star (e.g. in the case of reconfinement shocks, \citet{Nalewajko11jetdiss}). Even considering substantial progress in jet modeling, the launching radius is one of the least well constrained physical parameters of the fireball model. The gravitational wave observations can determine the parameters of the resulting black hole and give a strict lower limit on the launching radius. In the next section we list the observational properties of the GW and $\gamma$-ray event. In Section \ref{sec:model}, we briefly speculate on the parameters of the gamma-ray emitting system. In Section \ref{sec:rad}, we mention GRB radiation models in the context of this source. Finally, we discuss our results in Section \ref{sec:disc}. For quantity Q, we use the $Q_x=Q/10^x$ scaling notation in cgs units and the physical constants have the usual meanings.

\label{sec:disc} Assuming the binary black hole merger is associated with the gamma-ray signal detected by GBM, we have taken the leading prompt emission models for GRBs and applied it to the observations of GW~150914-GBM, aided by the accurately determined central engine parameters through GW measurements. We find that the non-dissipative photosphere and the internal shock models have some difficulty in interpreting the observations, though at this point no model can be definitely ruled out, while a dissipative photosphere model is unconstrained. The external shock model is able to interpret both the high peak energy and the observed flux yielding constraints on the Lorentz factor of the explosion ($\gtrsim 1000$) and the interstellar density ($\sim 10^{-3} \cm^{-3}$). The lower than usual ISM density is in line with the expectation that the merger takes place far from the birthplaces of its components (e.g. in a galactic halo environment). Furthermore, the low density might be a more general property of the external shock model itself which, applied to model afterglow observations of GRBs with $\gtrsim$ GeV photons (e.g. GRB 090510A) yield similarly high $\Gamma$ and low $n$ \citep{DePasquale+09-090510ag}. If we assume spectral parameters characteristic of short GRBs, we still consistently find high $\Gamma$ and low $n$. Even though our results are not definitive, the strength of such an approach lies in constraining values of the launching radius through GW observations and address EM observations. Further GW observations with better coverage from GBM will settle if merging BH binaries indeed emit $\gamma$-rays. It is possible however, that due to observer angle effects, the GRB-GW association will only be settled once a sizeable sample of GW and gamma-ray observations has accumulated. Indeed, GW signals from compact mergers are not strongly dependent on the orientation of the binary while prompt gamma-rays are essentially not expected if we are not inside the jet opening angle. On the other hand, edge-on systems have on average $1/\sqrt{8}$ the signal of the face on cases. This results in an increased likelihood that the systems detected by Advanced LIGO are face-on than edge-on. By measuring the jet opening angle for a GRB, we can constrain the available parameter space for the inclination, measured by the Advanced LIGO. Furthermore, detailed multiwavelength afterglow modeling \citep[e.g.][]{Zhang+15offaxis} can also constrain the viewing angle. Once GW observations become routine, and their EM counterparts will be readily available, we will be able to address the question of the association of short GRBs with BH mergers on a more solid footing. As an example of investigating GW~150914-GBM as a member of the short GRB population, \citet{Li+16gw} argued that it is an outlier on the $E_{\rm pk}-L_{\rm iso}$ diagram for short GRBs. This may indicate a different progenitor for the GBM event. However, due to a small sample size, the correlation for previous short GRBs is not strong enough for a definite conclusion. With our current understanding of GRBs, variability timescale of the GRB lightcurve can provide a limit on the size of the jet launch. If the GW-GRB association is confirmed, the variability times can be compared to the marginally stable radii resulting from the GW detection. Thus it will be possible to rule out some classes of models more firmly with similar analysis to the one presented here. {\it Acknowledgements -} We thank Tyson Littenberg and Michael Briggs for discussions. This study was supported by Fermi grant NNM11AA01A. P.M. acknowledges support from NASA NNX13AH50G.

1607.02616

1607

1607.05200_arXiv.txt

{} {We aim to study the possibility of a hadron-quark phase transition in the interior of neutron stars, taking into account different schematic evolutionary stages at finite temperature. We also discuss the strange quark matter stability in the quark matter phase. Furthermore, we aim to analyze the astrophysical properties of hot and cold hybrid stars, considering the constraint on maximum mass given by the pulsars J1614-2230 and J0348+0432.} {We have developed a computational code to construct semi-analytical hybrid equations of state at fixed entropy per baryon and to obtain different families of hybrid stars. An analytical approximation of the Field Correlator Method is developed for the quark matter equation of state. For the hadronic equation of state we use a table based on the relativistic mean field theory, without hyperons. The phase transition was obtained imposing the Maxwell conditions, by assuming a high surface tension at the interface hadron-quark. We solved the relativistic structure equations of hydrostatic equilibrium and mass conservation for hybrid star configurations.} {For the different equations of state obtained, we calculated the stability window for the strange quark matter, lepton abundances, temperature profiles and contours profiles for the maximum mass star depending on the parameters of the Field Correlator Method. We also computed the mass-radius and gravitational mass-baryonic mass relationships for different hybrid star families.} {We have analyzed different stages of hot hybrid stars as a first approximation of the cooling evolution of neutron stars with quark matter cores. We obtain cold hybrid stars with maximum masses $\geq 2 M_\odot$ for different combinations of the Field Correlator Method parameters. In addition, our study based on the gravitational mass - baryonic mass plane shows a late phase transition between hadronic and quark matter during the proto-hybrid star evolution, in contrast with previous studies of proto-neutron stars.}

A proto-neutron star (proto-NS) is the remaining old degenerated core of an intermediate mass star (between 10 and 25 solar masses at the zero age main sequence) after a supernova type-II or type-Ib/Ic explosion \citep{Woosley:2002}. Its evolution could be described roughly by three characteristic isentropic stages: first, the new born neutron star (NS) has an entropy per baryon $s \sim 1$ (in units of Boltzmann constant), with an abundance of trapped neutrinos, $Y_{\nu} \neq 0$, whose fraction is set by requiring the total fraction of leptons $Y_{\it l} \simeq 0.4$. This stage is dominated by neutrino diffusion, but as the neutrino mean free path is much smaller than the star radius, they remain trapped. The first $\sim 15$ seconds, the proto-NS radiates its lepton excess (deleptonization) and the neutrino flux produces an increase in the temperature. The star reaches the maximum heating in this second stage, and its inner matter has entropy per baryon $s \simeq 2$. Next, the neutrino mean free path begins to increase, becoming much larger than the radius of the star. The fraction $Y_{\nu}\neq 0$ evolves to a state in which $Y_{\nu} = 0$, because the matter becomes transparent to neutrinos. Once the electron chemical potential reaches the muon mass value, muons start to appear. The arising of muons reduces the number of electrons and affects the proton fraction of the stellar matter. The deleptonization process continues, and after a few minutes the third stage takes place. The entropy per baryon decreases to $s = 0$ and the proto-star cools down remaining transparent to neutrinos with $Y_{\nu} = 0$. The resulting object after these three stages is a cold and stable NS \citep{Burrows:1986,Steiner:2000}. In the chaotic rearrangement period of deleptonization and cooling described above, a phase transition to an exotic state of matter, like hyperons, Bose condensates or quark matter, could appear \citep{Benvenuto:1999,Benvenuto:1989tf,Benvenuto:1989qr, Prakash:2001} inside the proto-NS. In the case of quark matter, studies of the Quantum Chromodynamics (QCD) phase diagram predicts that the crossover shown by Lattice calculations at vanishing chemical potential, will become a first-order phase transition at intermediate temperatures and high baryon chemical potentials (see \cite{Yin:2012} and references therein). It may be plausible that this scenario of matter under extreme conditions occurs in the interior of neutron stars, as we get closer to the innermost core of the star, from the outer to the inner core, where the matter is compressed to densities several times the nuclear matter saturation density ($n_0 \sim \unit[0.16] {fm^{-3}}$). Most of the mass of a NS is contained in the inner and outer cores and its equation of state (EoS) has not yet been well determined. The discovery of the massive PSR J1614-2230 and PSR J0348+0432 ($\sim 2 M_{\odot}$) neutron stars, whose masses have been very accurately determined ($ \pm 0.04 M_{\odot}$) \citep{Demorest:2010,Antoniadis:2013}, has placed limits on a wide variety of EoS. Many EoS had to be ruled out, and the internal composition of NS had to be reconsidered, because a viable EoS should, at least, reproduce the mass value of these observed NS. At the end of the process of the NS formation, the degenerate neutron pressure partially counteracts the gravity force. The repulsive nuclear force beats the neutron degeneracy pressure resulting in less compact and more rigid structures supported by a stiffer EoS. For this reason neutron star masses are larger than $\sim \unit[0.7] {M_{\odot}}$, the maximum star mass derived by Tolman-Oppenheimer and Volkoff equations \citep{tov1, tov2} if a non-interacting degenerated neutron gas is considered as forming neutron stars. Thus, strong interactions are crucial in obtaining neutron star configurations that are compatible with measured neutron star masses. Some alternative models for the EoS include strange matter in form of hyperons (see, for example \cite{Bednarek:2011gd,Yamamoto:2014jga,Katayama:2015}). Although hyperons soften the EoS, massive NS with hyperons would be possible, considering certain relationships between the coupling constants and interactions between the particles involved in the hadronic matter composing their interiors \citep{Oertel:2016}. Another possibility is that NS could contain pure quark matter. However, to obtain stars as heavy as two solar masses, it should be feasible to adjust some of the parameters which mediates the interaction between quarks \citep{Orsaria:2013hna, Weber:2014qoa} and/or that quark matter is a color superconductor \citep{Bonanno:2012,Zdunik:2013}. Whatever the matter composing the interior of a NS may be (quarks or hadrons or a mixing of both phases), the equation of state describing the matter should be stiff enough to obtain massive NS. Regarding quark matter inside compact objects, for many years the models most used to describe it have been the MIT bag model \citep{Chodos:1974je} and the Nambu Jona-Lasinio model \citep{Nambu:1961a,Nambu:1961b}. Some modifications to these models have been implemented to obtain a more suitable description of quark matter at high densities in the context of hybrid stars (HS), NS with a quark matter core surrounded by a shell of hadronic matter (see \cite{Chatterjee:2016} and references therein). Recently, for the description of quark matter phase inside cold NS, the Field Correlator Method (FCM) has been used \citep{Plumari:2013,Logoteta:2013,Burgio:2016}. The FCM is a non-perturbative approximation of QCD which includes, from the first principles, the confinement dynamics in terms of field correlators. The model is parametrized through the gluonic condensate, $G_2$, and the quark-antiquark static potential for long distances, $V_1$, taking into account the confinement. These two parameters control the EoS of the deconfined quark phase. In the framework of the isentropic stages described at the beginning, we have studied the occurrence of a first order phase transition from hadronic to quark matter in proto-NS, assuming the surface tension at the interface hadron-quark is such that a sharp phase transition is the favorable scenario for the deconfinement (\cite{Alford:2015, Ranea:2015ldr}). The hybrid EoS that models the proto-NS is obtained using the simplest method for constructing a first order phase transition by considering an isobaric transition with one conserved charge, for example, the baryonic chemical potential. This is known as Maxwell construction. To obtain the phase transition and the astrophysical quantities for the study of hybrid stars, a computational code called NeStOR (Neutron-Star-Object Research) has been developed. This code contains a set of routines and scripts to calculate several EoS at finite temperature and constant entropy. The combination of different programming languages provided in the code establishes the essential conditions for an unique determination of the matter equilibrium composition in the proto-hybrid star, that is, the requirements of beta equilibrium, charge neutrality and the conservation of baryonic and leptonic number. The program also computes particle abundances and the stellar structure of a spherically symmetric compact object by integrating the relativistic equations of hydrostatic equilibrium for different model parameters. Through this code it is also possible to compare the values of the maximum baryonic mass of each family of stars for the three stages considered. Supposing that the baryonic mass remains constant during the process of thermal evolution, the gravitational mass of the star, which corresponds to the total stellar energy, changes during the evolution of the proto-NS. Then, a study of the evolution in a gravitational mass - baryonic mass plane must be taken into account, giving the possibility of knowing in which stage takes place the phase transition from hadronic matter to quark matter. The description of proto-NS evolution along constant baryonic mass sequences have been already discussed in \cite{Bombaci:1996}. Thereby, using the NeStOR code, it is possible to analyze the thermal evolution of proto-hybrid stars through snapshot sequences at constant entropy, also considering their characteristic composition at each stage. A semi-analytical calculation of the FCM at finite temperature is used to describe the quark matter phase at the inner cores. For the hadronic phase, we employ a table containing an EoS created for astrophysical simulations of proto-NS \citep{Shen:2010a, Shen:2011, Shen:2010b}. The paper is organized as follows. In Section \ref{section eos quark}, we give the thermodynamic expressions for the quark matter as functions of temperature and chemical potential, and we discuss the choice of the parameters of the FCM as well as the stability of strange quark matter. The thermodynamic conditions for the construction of the hybrid EoS and the phase transition hadron-quark are presented in Section \ref{phase}, together with the description of the hadronic and leptonic EoS. In Section \ref{entop} we explain briefly the calculations at constant entropy, giving the snapshot sequences of thermal evolution. In Section \ref{hybrid} we show the results about the structure of hot and cold hybrid stars and we analyze the gravitational mass-baryonic mass plane for some representative combinations of the FCM parameters. A summary of the work and conclusions are provided in Section \ref{conclus}. Details of the calculations and treatment of quark matter are given in the Appendix.

\label{conclus} In this work we have studied, in a very simplified way, the thermal evolution of proto-NS modeled as HS, considering a successive sequence of constant entropy per baryon snapshots. The evolution started in a hot stage, with and entropy per baryon of $s \simeq 1$ and where trapped neutrinos have been subjected to the condition $Y_{L_e} = 0.4$. In the following stage, after $\sim 15$ seconds, neutrinos diffuse and warm the star being $s \simeq 2$ and $Y_{\nu_e} = 0$. Finally, after a few minutes, the star cools down and the last stage is reached, where $s = 0$. We used a tabulated EoS to describe the hadronic phase of the proto-hybrid stars and the FCM for the description of quark matter. To analyze how the results for proto-hybrid stars evolution depend on the FCM parameters, we have taken into account three different combinations of $V_1$ and $G_2$. The use of the NeStOR code to construct the hybrid EoS, allowed us to study the HS structure via the integration of the TOV equations as well as the gravitational mass - baryonic mass plane, $M_G-M_B$. The hybrid EoS was constructed at finite temperature, taking into account the characteristic composition for each evolutionary stage. For the hadronic phase, we used an EoS constructed for astrophysical simulation of proto-NS. We have obtained, for each constant entropy stage, a family of stable HS, with a characteristic mass-radius relationship. Considering the constraints from astrophysical observations of NS, we have studied the variation of the maximum mass star with the different combination of FCM parameters, for each family. Also in the framework of the FCM semi-analytical quark matter EoS, we analyzed the stability of strange quark matter within the strange quark matter hypothesis. The results obtained from our semi-analytical quark EoS are in agreement with previous works using the FCM, both the simplified phase diagram of the QCD as the stability of the strange quark matter. Under the simplified model of thermal evolution considered in this work, the temperature profiles obtained for the HHS are in agreement with the dynamical calculations with schematic EoS of the thermal evolution of NS. Regarding the concept of maximum mass for proto-NS, we have also analyzed the $M_G-M_B$ plane. The cooling evolution was studied considering the baryonic mass conservation during the whole process \citep{Bombaci:1996}. We have schematically analyzed the baryonic mass ranges, separating the newly born NS, which will evolve to stable CHS from those that will collapse afterwards to a black hole. The study of the gravitational mass-baryonic mass plane for proto-HS, has revealed in which of the three stages of constant entropy considered, a phase transition from hadronic matter to quark matter could occur in the context of the microscopic models used in this work. Our results suggest that an early ($s \simeq 1$) phase transition between quark matter and hadronic matter is not produced for the baryonic mass ranges in which the proto-HS end as stable CHS, but the formation of stable CHS is possible due a late phase transition, after the stage $s \simeq 2$. This contrasts with previous studies \citep{Benvenuto:1989tf, Benvenuto:1989qr, Benvenuto:1999} where pure hadronic stars are converted to quark stars within the first seconds after their birth. It is worth pointing out that our model is a simplified treatment of the thermal evolution. The calculation of the quark matter nucleation rate and the nucleation time, due to thermal nucleation mechanisms in the framework of the EoS considered, would provide a better understanding of the finite size effects of the phase transition between hadronic and quark matter. These further calculations are beyond the scope of this work. We have obtained stable CHS that easily reach the two solar masses for the maximum mass star for several combination of the FCM parameters. According our results, the pulsars J1614-2230 and J0348+0432 would contain quark matter in their inner cores. However, we note that we have neither considered hyperons in the hadronic phase nor the color superconductivity in the quark-gluon plasma phase. The inclusion of either of these two contributions will modify the EoS of each phase and therefore the critical point where the transition occurs. The effects of such contributions should be studied in a future work. Future observations of neutrinos in supernovae, as well as the observation of $x-ray$ and $\gamma-ray$ emissions of young pulsars will enable the astrophysicist community to carry out simulations on real thermal evolution that will contribute to the understanding of the underlying physics of NS.

1607.05200

1607

1607.05170_arXiv.txt

Sky-coverage in laser-assisted AO observations largely depends on the system's capability to guide on the faintest natural guide-stars possible. Here we give an up-to-date status of our natural guide-star processing tailored to the European-ELT's visible and near-infrared (0.47 to 2.45 $\mu m$) integral field spectrograph -- Harmoni. We tour the processing of both the isoplanatic and anisoplanatic tilt modes using the spatio-angular approach whereby the wavefront is estimated directly in the pupil plane avoiding a cumbersome explicit layered estimation on the 35-layer profiles we're currently using. Taking the case of Harmoni, we cover the choice of wave-front sensors, the number and field location of guide-stars, the optimised algorithms to beat down angular anisoplanatism and the performance obtained with different temporal controllers under split high-order/low-order tomography or joint tomography. We consider both atmospheric and far greater telescope wind buffeting disturbances. In addition we provide the sky-coverage estimates thus obtained.

In this work we layout the NGS modes processing for laser tomography AO systems. We start by addressing the common tilt mode (isoplanatic) considering different control strategies from the ESO-suggested double stage, double integrators with lead filters and Linear-Quadratic-Gaussian (LQG) controllers under two distinct operating scenarios: in stand-alone mode or in tandem with the telescope's image stabilisation controller averaging tip/tilt signals from three NGS off-axis. We then move on to handle the anisoplanatic tilt. Three models are presented: \\ \textit{i)} tilt-tomography using a combination of tilt and high-altitude quadratic modes that produce pure tilt through cone projected ray-tracing through the wave-front profiles \cite{correia13}; \\ \textit{ii)} a spatio-angular MMSE tilt estimation anywhere in the field that is more general. \\ \textit{iii)} the (straightforward) generalisation to dynamic controllers using near-Markovian time-progression models from \cite{correia15}. These controllers will be used for the HARMONI NGS modes \cite{thatte16, neichel16, sauvage16}.

1607.05170

1607

1607.08052_arXiv.txt

In this paper, inspired by the ultraviolet deformation of the Friedmann-Lema\^{\i}tre-Robertson-Walker geometry in loop quantum cosmology, we formulate an infrared-modified cosmological model. We obtain the associated deformed Friedmann and Raychaudhuri equations and we show that the late time cosmic acceleration can be addressed by the infrared corrections. As a particular example, we applied the setup to the case of matter dominated universe. This model has the same number of parameters as $\Lambda$CDM, but a dynamical dark energy generates in the matter dominated era at the late time. According to our model, as the universe expands, the energy density of the cold dark matter dilutes and when the Hubble parameter approaches to its minimum, the infrared effects dominate such that the effective equation of state parameter smoothly changes from $w_{_{\rm eff}}=0$ to $w_{_{\rm eff}}=-2$. Interestingly and nontrivially, the unstable de Sitter phase with $w_{_{\rm eff}}=-1$ is corresponding to $\Omega_m=\Omega_d =0.5$ and the universe crosses the phantom divide from the quintessence phase with $w_{_{\rm eff}}>-1$ and $\Omega_m> \Omega_d$ to the phantom phase with $w_{_{\rm eff}}<-1$ and $ \Omega_m<\Omega_d$ which shows that the model is observationally viable. The results show that the universe finally ends up in a big rip singularity for a finite time proportional to the inverse of the minimum of the Hubble parameter. Moreover, we consider the dynamical stability of the model and we show that the universe starts from the matter dominated era at the past attractor with $w_{_{\rm eff}}=0$ and ends up in a future attractor at the big rip with $w_{_{\rm eff}}=-2$.\vspace{5mm}\\ PACS numbers: 04.60.Bc; 95.36.+x\\ Keywords: Phenomenology of Quantum Gravity, Dark Energy

Cosmological observations indicate that the Universe accelerates positively at the small redshifts \cite{DE} which leads to the so-called dark energy problem \cite{DEP,DEP2}. In the standard $\Lambda$CDM model, cosmological constant dominates at the late time and derives cosmic speed-up. But, the models in favor of cosmological constant clue to the cosmological constant problem due to the possible identification of the cosmological constant with the vacuum energy of the quantum fields \cite{DEP,CCP}. Furthermore, increasing evidences from the cosmological data reveal that the energy density corresponds to the dark energy evolves very slowly in time and the associated equation of state parameter lies in a narrow strip around $w=-1$ \cite{DE}. Thus, cosmological constant with sharp value $w=-1$ for the equation of state parameter is an appropriate candidate in the first order of approximation \cite{DE-CC}. In order to explain the dynamical nature of the dark energy, the quintessence scenarios with $w>-1$ and phantom models with $w<-1$ are proposed. In this respect, one usually interested in models which support the transition from the quintessence era to the phantom phase. These scenarios are usually based on two postulates: i) assuming general relativity is applicable even on cosmological scales and then considering some sort of unusual matter component(s) costing violation of some energy conditions, ii) deformation of general relativity at the cosmological scales. For the first case the matter source is usually given by a scalar field \cite{DE-MG,chaplygin,cardassian} and for the latter case, there are many candidates such as the extra dimensions models , $f(R)$ theories \cite{DE-FR,DE-FR2} and recently proposed massive gravity models \cite{DE-NIR}. From the theoretical point of view, de Sitter spacetime is a maximally symmetric space and its constant curvature is completely determined by the cosmological constant. Apart from the very small variation of cosmological constant with time, it can be interpreted as a fundamental constant of nature much similar to the speed of light and Planck constant. It therefore provides a universal infrared (IR) cutoff (corresponding to the large length scale $\sim10^{-56}\mbox{cm}^{-2}$) for the universe. For instance, existence of cosmological constant as an IR cutoff is essential for the quantization of scalar field in de Sitter spacetime. More precisely, it provides a minimum scale for the momenta of modes through the uncertainty principle and removes the IR divergences in this setup \cite{maggiore}. In this respect the uncertainty principle will be modified in curved spacetimes in order to respect the existence of cosmological constant as a universal IR cutoff \cite{UR-dS}. On the other hand, existence of a minimal length scale is suggested by any quantum theory of gravity such as loop quantum gravity \cite{Loop} and string theory \cite{String}. It is also shown that the uncertainty principle is modified in the presence of a minimal length scale \cite{GUP}. Thus, the uncertainty principle gets modifications in IR and ultraviolet (UV) regimes in order to respect the existence of cosmological constant and minimal length scale respectively \cite{KempfIR}. Taking these universal IR and UV cutoffs into account, the quantum field theories turn out to be renormalizable \cite{IRCurve}. Therefore, natural IR and UV cutoffs would be emerged in the context of ultimate quantum gravity theory. While the existence of a universal IR cutoff is supported by the standard general relativity framework through the de Sitter spacetime\footnote{Note that the anti-de Sitter spacetime with negative cosmological constant is also an appropriate candidate from the theoretical point of view. But it rejects by cosmological observations.}, there is not any explanation for the UV cutoff (minimal length scale) in this setup. On the other hand, a minimal length scale as a UV cutoff emerges in loop quantum gravity framework \cite{LQG} but there is not a well-defined explanation for taking a cosmological constant into account in this setup (see however Refs. \cite{LQC-CC} where some attempts have done in this direction). In this paper, we follow the UV deformation of the Friedmann-Lema\^{\i}tre-Robertson-Walker (FLRW) universe in loop quantum cosmology and we construct the corresponding IR-deformed case. We show that the late time cosmic acceleration arises in this setup which is significantly different from the $\Lambda$CDM model such that the universe crosses the phantom divide from $w_{_{\rm eff}}>-1$ to $w_{_{\rm eff}}<-1$.

Cosmological observations show that the universe accelerates at small redshifts and the cosmological constant derives the desired acceleration in standard $\Lambda$CDM cosmology. In the context of general relativity, the cosmological constant can be interpreted as a universal IR cutoff. Quantum gravity candidates such as loop quantum gravity and string theory also suggest the existence of a minimum length scale of the order of the Planck length. The ultimate quantum theory of gravity then should contain universal IR and UV cutoffs. While the standard general relativity accommodates the existence of IR cutoff through the de Sitter spacetime with positive cosmological constant, it cannot support the existence of a UV cutoff. On the other hand, loop quantum cosmology scenario suggests the existence of a minimum length scale as UV cutoff for the system under consideration, but it does not support the existence of any IR cutoff. In this paper, following the UV deformation of the FLRW gravitational phase space in loop quantum cosmology, we have formulated deformed phase space which supports the existence of an IR cutoff. We obtained the associated IR-deformed Friedmann and Raychaudhuri equations and we showed that the IR corrections derives the late time cosmic acceleration. The model has the same number of parameters as $\Lambda$CDM such that the IR effects replace the effects of cosmological constant. But, the dynamics of the universe in our model is very different from the standard $\Lambda$CDM cosmology. For instance the Hubble parameter and energy density turned out to be bounded from the below in this model. As a particular example, we applied the setup to the simple case of CDM dominated universe. We divide the cosmic evolution in this model into the following three phases: \begin{itemize} \item {\it Quintessence phase} ($-1<w_{_{\rm eff}}\leq0$): The universe starts from the standard matter dominated era with equation of state parameter $w_{_{\rm eff}}\approx{w_m}=0$ for $\rho_{_{\rm eff}}\gg\rho_{_{\rm min}}$ when the IR effects are negligible. In this phase, as the universe expands and the total energy density (\ref{rho-total}) dilutes, the IR effects become more and more appreciable up to $\rho_{_{\rm eff}}\sim\rho_{_{\rm min}}$. The effective equation of state parameter, which is given by the plus sign of the relation (\ref{EoS}) in this regime, then smoothly decreases from $w_{_{\rm eff}}=0$ in matter dominated era to the negative values $w<0$ and the universe starts to accelerate when it crosses over the value $w_{_{\rm eff}}= -\frac{1}{3}$. From the relations (\ref{EoS}) and (\ref{DP-rho}), it is clear that the accelerating phase with $-1<w_{_{\rm eff}}<-\frac{1}{3}$ is corresponding to $\Omega_m>\Omega_d$. \item {\it Unstable de Sitter phase} ($w_{_{\rm eff}}=-1$); {\it Transition from} $w_{_{\rm eff}}>-1$ {\it to} $w_{_{\rm eff}}<-1$: The effective equation of state parameter (\ref{EoS}) then approaches to the value $w_{_{\rm eff}}=-1$ when the energy density of CDM approaches to the critical value $\rho_{_m}=\rho_{_{\rm min}}$. At this momentum the Hubble parameter approaches to its minimum $H=H_{\rm min}$ that is corresponding to an unstable de Sitter phase with effective cosmological constant $\Lambda_{_{\rm IR}}=3H_{\rm min}^2=2\kappa \alpha\rho_{0_m}a_0^3$. From (\ref{EoS}) and (\ref{DP-rho}), one can easily see that the model includes a transition from the quintessence era with $w>-1$ and $\Omega_m>\Omega_d$ to the phantom phase with $w_{_{\rm eff}}<-1$ and $\Omega_m< \Omega_d$ when crossing over the unstable de Sitter phase with $w_{_{\rm eff}}=-1$ and $\Omega_m=\Omega_d$. This result makes the model observationally viable. \item {\it Phantom phase} ($-2\leq{w_{_{\rm eff}}}<-1$): After the universe enters into a phantom era with $w_{_{\rm eff}}<-1$ and $\Omega_m <\Omega_d$, the effective equation of state parameter decreases until approaches to its asymptotic value $w_{_{\rm eff}}=-2$ where the universe ends up in a big rip singularity at the finite time (\ref{time-br}). \end{itemize} We have also considered the dynamical stability of the model which shows that, in this model, the universe starts at a past attractor in matter dominated era ($w_{_{\rm eff}}=0$) and after crossing an unstable point, corresponds to a de Sitter phase ($w_{_{\rm eff}}=-1$), it approaches to a future attractor ($w_{_{\rm eff}}=-2$) which is corresponding to the big rip singularity.\\ {\bf Acknowledgement}\\ The author would like to thank Babak Vakili for a critical reading of the manuscript, Kourosh Nozari for useful comments on the first draft of this work, and Achim Kempf for helpful discussions. He also thanks anonymous referee for very insightful comments which considerably improved the quality of the paper.

1607.08052

1607

1607.04560_arXiv.txt

We present a measurement of the cosmogenic activation in the germanium cryogenic detectors of the EDELWEISS III direct dark matter search experiment. The decay rates measured in detectors with different exposures to cosmic rays above ground are converted into production rates of different isotopes. The measured production rates in units of nuclei/kg/day are 82 $\pm$ 21 for $^3$H, 2.8 $\pm$ 0.6 for $^{49}$V, 4.6 $\pm$ 0.7 for $^{55}$Fe, and 106 $\pm$ 13 for $^{65}$Zn. These results are the most accurate for these isotopes. A lower limit on the production rate of $^{68}$Ge of 74~nuclei/kg/day is also presented. They are compared to model predictions present in literature and to estimates calculated with the ACTIVIA code.

Germanium is widely used as a detector material in experiments searching for a rare process like the interaction of weakly interacting massive particles (WIMPs)~\cite{dmreview}. It is possible to build detectors with very good energy resolution based on the measurement of the ionization produced in the particle interaction, or of the increase of temperature~\cite{gebolo}. In addition, the combination of the ionization and heat signals is a powerful tool to distinguish nuclear recoils from electron recoils. Moreover, the crystal-growing process used in the semiconductor industry purifies the material to a high level that matches well the stringent ultra-low radioactivity requirements of rare event searches. The potential of germanium detectors for achieving very low threshold below 1~keV is particularly attractive for searches of WIMPs with masses below 10~GeV/c$^{2}$. The background at energies below 20~keV in such a detector is thus of particular interest. Notably, the contribution from tritium beta decays may have a significant impact on the sensitivity of the next generation of these detectors. The crystallization process removes all cosmogenically-produced radioactive atoms, with the exception of unstable germanium isotopes like $^{68}$Ge (see below). Their populations grow back again when the crystal is kept above ground, and therefore exposed to cosmic rays and the associated hadronic showers. Short-lived isotopes decay rapidly as soon as the detectors are stored underground, deep enough to suppress the hadronic component of the cosmic rays~\cite{Farley2006451}. The isotopes that merit attention have lifetimes exceeding a few tens of days, since shorter-lived nuclei can be eliminated just by storing the detectors in the underground site for some reasonable time before starting data taking. The cosmogenic products that have the most noticeable effect on the low-energy spectrum recorded in germanium detectors are those that decay via electronic capture (EC). The capture is often followed by the emission of a $K$-shell X-ray with characteristic energy between 4 and 11~keV. $L$- and $M$-shell captures will produce weaker lines at approximately 1 and 0.1~keV, respectively. The sharp line shapes and known $K$:$L$:$M$ intensity ratios can be used to identify and subtract the associated events. However, it is preferable to reduce their initial intensities to the lowest possible level. Measurements of the production rates of EC-decaying isotopes is helpful in designing a detector-production procedure that limits these backgrounds to acceptable levels, and, more generally, to constrain models predicting the production rates of all isotopes, including those that may prove to elude direct measurements. Another type of background of particular interest is the beta decay of tritium ($^{3}$H) originated from nuclear reactions induced by the interaction of the hadronic component of cosmic rays with atoms in the material~\cite{avignone}. The electron emitted in the beta decay of tritium has an end point $Q_{\beta}$ of only 18.6~keV, and thus contributes to the background of low-energy events over the entire energy range relevant for low-mass WIMP searches. The lifetime of $^{3}$H is particularly long ($\tau$ = 17.79~y), so the tritium activity can essentially be expected to remain almost the same throughout the life of the detector. The only way to reduce this background is to limit the exposure of the crystal between the time it is grown and its installation underground. There are large uncertainties in model predictions for the production rate of $^{3}$H, and available measurements can only provide crude upper limits~\cite{mei}. The EDELWEISS collaboration has operated an array of 24 germanium heat-and-ionization detectors with the objective to perform searches for WIMPs with a total exposure close to 3000~kg$\cdot$d, and more specific searches for low-mass WIMPs with a subset of its detectors with the best experimental energy thresholds~\cite{lowmass}. The experiment is located in the Laboratoire Souterrain de Modane (LSM) and protected by a mean rock overburden of 1800~m (4800~m.w.e.) that reduces the cosmic ray flux to about 5~$\mu$/m$^2$/day~\cite{Schmidt:2013gdc}, i.e.\ 10$^6$ times less than at the surface. The detectors are covered by interleaved electrodes that provide an efficient tool to reject surface events (i.e.\ particle interactions taking place within $\sim$2 mm from the detector surface) down to energies of $\sim$1~keV~\cite{lowmass}. The resolutions achieved with these detectors, the reduction of the external gamma-ray background and the excellent surface-event rejection performance of the interleaved electrodes~\cite{broniatowski}, have made possible a precise measurement of decay rates of different nuclei in the bulk volume of germanium detectors, and in particular, for the first time, to measure unambiguously the intensity of the tritium spectrum. Efforts were made to keep to a minimum the exposure of each crystal to cosmic rays throughout the detector production. A history of the key steps in the detector production process is available. Despite this, there are non-negligible systematic uncertainties in the recorded history of exposure times. However, these uncertainties can be tested, because unforeseen production delays\footnote{ These delays occurred to solve problems related to surface current leakage, as described in \protect{\cite{leakage}}.} led to a relatively large spread (up to a factor of 4) in the exposures of the different detectors to cosmic rays. It was therefore possible to check on isotopes with the largest statistics that the observed activation rates scaled with the expectations from the recorded history of exposure times. In the following, we detail the EDELWEISS-III setup relevant to this measurement (Section~\ref{sec:setup}), as well as the data selection (Section~\ref{sec:data}). We present the expected properties of the activation of tritium and other isotopes and of the energy spectrum of the emitted electrons, and describe the analysis used to extract their intensities from the data~(Section \ref{sec:analysis}). These results are then converted into production rates during the exposure above ground (Section~\ref{sec-history}) and compared to a previous measurement and calculations (Section~\ref{sec:pr}).

The cosmogenic activation of various isotopes in the germanium detectors of the EDELWEISS-III experiment has been measured. The data for five isotopes and thirteen detectors with different exposure times lead to a consistent set of measurements. The first measurement of the $^{3}$H decay rate in germanium detectors is presented. It has been interpreted in terms of production rate of 82~$\pm$~21~nuclei/kg/day with statistical and systematic uncertainties included. The tritium production due to cosmic-ray neutrons is thus important and the present measurement provides valuable information needed to evaluate the reduction of the exposure to cosmic rays necessary for germanium detector arrays used for dark matter searches. The measured production rates on $^{49}$V, $^{55}$Fe and $^{65}$Zn of 2.8 $\pm$ 0.6~nuclei/kg/day, 4.6 $\pm$ 0.7~nuclei/kg/day and 106 $\pm$ 13~nuclei/kg/day, respectively, presented here are the most accurate to-date. A lower limit of 74~nuclei/kg/day at 90\% C.L.\ on production rate of $^{68}$Ge is discussed. The measured $^{65}$Zn production rate and the lower limit on that of $^{68}$Ge are a factor 2.7 $\pm$ 0.6 larger than the measurements reported in Ref.~\cite{avignone}. The origin of this discrepancy is unknown. The measurements agree within a factor of two with estimates performed with the ACTIVIA code within this work. The best agreement is found for $^{49}$V and $^{55}$Fe. The estimates for $^{3}$H, $^{65}$Zn and $^{68}$Ge tend to underestimate the measured rates, with significance ranging from 1.7$\sigma$ to 6$\sigma$. The difference between these predictions and those from other models can also differ by as much as a factor of two in most cases, with no single model giving a satisfying description for all measured isotopes. It can be foreseen that the precision of the present measurements will help constrain and further improve the models.

1607.04560

1607

1607.06615_arXiv.txt

{High-energy cosmic-ray electrons reveal some remarkable spectral features, the most noteworthy of which is the rise in the positron fraction above 10~GeV. Due to strong energy loss during propagation, these particles can reach Earth only from nearby sources. Yet, the exact nature of these sources, which most likely manifest themselves in the observed anomalies, remains elusive. The many explanations put forward to resolve this case range from standard astrophysics to exotic physics. In this paper, we discuss the possible astrophysical origin of high-energy cosmic-ray electrons through a fully three-dimensional time-dependent Monte Carlo simulation. This approach, which takes advantage of the intrinsic random nature of cosmic-ray diffusive propagation, provides valuable information on the electron-by-electron fluctuations, making it particularly suitable for analyzing in depth the single-source scenario.} \begin{document}

\label{sec:introduction} Cosmic rays are high-energy charged particles striking Earth from all directions since time immemorial. They originate mainly in outer space and comprise roughly 99\% of atomic nuclei and 1\% of electrons. They are commonly divided into primary and secondary cosmic rays. Primary cosmic rays, such as electrons, protons, helium, iron and other nuclei synthesized in stars, consist of particles directly accelerated at sources. Secondary cosmic rays, like lithium, beryllium and boron, which are not created by stellar nucleosynthesis, are composed of particles produced by the primary cosmic rays during their propagation in the interstellar medium. A very small fraction of cosmic rays are stable particles of antimatter: positrons and anti-protons. Whether or not cosmic-ray positrons and anti-protons are pure secondaries is still an open question. Even though electrons\footnote{Unless stated otherwise, we hereafter refer to both e$^-$ and e$^+$ as simply \textit{electrons}.} represent only a tiny part of cosmic rays, they are of great interest, especially at high energy. Indeed, above a few GeV these particles are subject to severe radiative energy losses, mainly by synchrotron radiation in magnetic fields and by inverse Compton scattering in radiation fields \cite{JON65, BLU70}. These processes are so drastic that high-energy cosmic-ray electrons (HECREs) cannot travel far distances from their sources. In stark contrast to hadrons, they can reach Earth only from sources in the local neighborhood \cite{COW79}. Therefore, they can be used as a powerful probe into local cosmic-ray accelerators, which is crucial to the long-standing problem of the origin of cosmic rays. HECREs are also very important to X-ray and $\gamma$-ray astronomies, to dark matter investigation and to many other issues in high-energy astrophysics. Researchers have taken a keen interest in cosmic-ray electrons for a very long time. The first direct observation of these particles was achieved in 1961 with balloon-borne experiments \cite{EAR61, MEY61}. During the following years many balloon flights using different detectors were carried out \cite{DAN65, SMI68, RUB68, SCH71}. These early experiments measured the flux of cosmic-ray electrons up to several hundred GeV. Experiments then explored higher energy region using more sophisticated devices with a larger geometrical factor, longer exposure and higher proton rejection power. These detectors were chiefly of two kinds: magnetic spectrometers and emulsion chambers. The first category observed negative electrons and positrons separately up to a few tens of GeV. Such was the case with the balloon-borne experiments CAPRICE \cite{BOE00, BOE01}, HEAT \cite{DUV01, BEA04} and MASS-91 \cite{GRI02}, in addition to AMS-01, which was flown on the space shuttle Discovery in 1998 \cite{AGU07}. The second category lacked the ability of discriminating between negative electrons and positrons. However, it had the merit of extending the energy spectrum measurements far beyond the range accessible to the first category. Examples include the balloon-borne experiments ECC \cite{KOB12}, BETS \cite{TOR01}, ATIC-2 \cite{CHA05} and PPB-BETS \cite{YOS08}. Moreover, the ground-based Cerenkov telescopes H.E.S.S. \cite{AHA08} and MAGIC \cite{BOR11} also observed cosmic-ray electrons at very high energy. The launch over the last decade of a new generation of high-precision instruments on-board satellites (PAMELA, the Fermi Gamma-ray Space Telescope, AMS-02 and CALET) opened a new era in the study of HECREs. The detection of HECREs implies the existence of local sources. These sources are expected to manifest themselves in the energy spectrum, which should display special features. Recent experimental results show that the energy spectrum of cosmic-ray electrons extends well beyond 1~TeV and does have features. Most notably, PAMELA uncovered in the energy range 10-300~GeV a significant increase in the positron fraction (ratio of the positron flux to the combined flux of positrons and negative electrons) \cite{ADR09, ADR10, ADR13}. This remarkable result, confirmed by Fermi-LAT \cite{ACK12} and by AMS-02 \cite{AGU13}, is not consistent with conventional models based on the assumption that positrons arise only from the secondary production of cosmic rays by collision with the interstellar medium. These models rather predict a positron fraction falling smoothly with energy \cite{PRO82, MOS98}. A new precision measurement by AMS-02 extending up to 500~GeV indicates that the positron fraction, on the one hand, levels off at about 200~GeV and, on the other hand, does not show any fine structure or sharp cutoff \cite{ACC14}. In parallel, ATIC reported an excess of cosmic-ray electrons over conventional model expectations between 300 and 800~GeV followed by a steepening above 1~TeV \cite{CHA08, PAN11}. Although PPB-BETS also showed a bump-like structure between 100 and 700~GeV \cite{YOS08}, Fermi-LAT and H.E.S.S. experiments found no evidence for a prominent peak \cite{ABD09, ACK10b, AHA09}. There are several possible scenarios, from standard astrophysics to exotic physics, to account for the special features of the electron energy spectrum. They may be ascribed to the contribution from some individual local sources such as supernova remnants (SNRs) (see, e.g., \cite{BLA09, AHL09, STA10}) or pulsars (see, e.g., \cite{HOO09, MAL09, YUK09, DEL10, BLA11, PRO11, DIM14}). However, additional sources such as the annihilation or decay of dark-matter particles cannot be excluded (see, e.g., \cite{KAN09, ARK09, CHO09, POR11, KOP13, BER13, IBA14, FEN14}). The main idea of these models is that equal amounts of negative electrons and positrons are produced by the source, be it an astrophysical object or dark-matter particles. This contribution emerges at high energy from a background formed by electrons coming from distant sources. The observed anomalies may also be due to just propagation effects \cite{SHA09,GAG13,BLU13,COW14}. It must be emphasized here that there are arguments for and against each scenario and none of these approaches is yet conclusive \cite{MOS13}. Recent reviews can be found in \cite{FAN10, SER12, CHO13,PAN13,PIC14,ISR14}. The energy spectrum of cosmic-ray electrons is typically interpreted within the context of propagation models and the traditional method consists in solving appropriate transport equations \cite{GIN76b, STR07}. Due to its inherent stochastic nature, the diffusive propagation of cosmic rays can also be treated using Monte Carlo simulation \cite{HUA07}. While it is widely believed that this technique is inefficient in this kind of application, the proximity of sources in the case of HECREs, coupled with the absence of hadronic interactions and the simplicity in energy loss processes, simplifies the modeling significantly. Monte Carlo approach comes in handy particularly at very high energy where only a few sources are expected to dominate the cosmic-ray electron flux. It provides very useful information about the electron-by-electron fluctuations, thus complementing the traditional method. In this respect, we implemented a fully 3-dimensional time-dependent Monte Carlo simulation of the propagation of HECREs in our galaxy. To speed up calculations we employed MPI parallel programming on an HPC cluster system. We restricted ourselves to the energy region above 10~GeV where cosmic rays are not affected by the solar wind modulation. We considered only pure diffusion since convection and reacceleration are negligible above a few GeV \cite{DEL09}. We focused on the most natural way to explain the spectral features of HECREs, assuming that some nearby pulsars and/or SNRs are the sources of such particles. The other possible scenarios were not examined. We used a two-component model, separating the local source contribution from the distant source contribution \cite{ATO95}. The latter, which constitutes the background, was estimated with the public code GALPROP\footnote{http://galprop.stanford.edu/} \cite{VLA11}. There are two different types of key quantities when interpreting the observed energy spectrum of cosmic-ray electrons. The first ones are the parameters associated with propagation; they are behind the distortion of the injection spectrum. The second ones are the parameters related to the spectral profile at injection; they control the overall form of the measured spectrum. In our study we first examined the propagation effects. Next, we inferred from these initial calculations the most likely astrophysical sources of HECREs. Then, we derived the expected flux from some typical sources. And finally, we addressed the anisotropy issue. The outline of this paper is as follows. After this introduction (\S\ref{sec:introduction}) we review in \S\ref{sec:propagation} the necessary mathematical background for describing cosmic-ray electron propagation and we detail our Monte Carlo procedure. In the following section (\S\ref{sec:results}) we report the main results achieved: electron energy and lifetime distributions in \S\ref{sec:prop_effects}, potential astrophysical sources in \S\ref{sec:sources}, cosmic-ray electron flux, positron fraction and anisotropy for some typical sources in \S\ref{sec:flux} and \S\ref{sec:anisotropy}. We end our study with a brief conclusion (\S\ref{sec:conclusion}).

\label{sec:conclusion} Regarding the spectral peculiarities observed at high energy for cosmic-ray electrons, which still resist a unified interpretation, there is now a manifest need to explore new avenues in the hope of solving this puzzle. Most often the way of tackling this problem revolves around the resolution of the transport equation describing the galactic propagation of these particles. This work demonstrates the feasibility and the relevance of fully three-dimensional time-dependent Monte Carlo simulation, which can supplement the reference method by providing additional information on the electron-by-electron fluctuations. This approach is proving to be particularly efficient at investigating more deeply the possible astrophysical origin of HECREs from nearby single sources and its benefit appears plainly in addressing the anisotropy issue. The proposed algorithm is quite simple, very flexible and highly scalable. When looking at the lifetime and energy distributions of observed cosmic-ray electrons, we came to a first list of candidate sources that includes Vela, B1737-30, Monogem, B1822-09, Geminga and B1742-30. When assuming the burst-like approximation, we found out that the signal from a young source like Vela would outweigh all other signals and would produce an exceedingly large anisotropy amplitude, which obviously disagrees with observations. Considering a time delay of 50~kyr with respect to the source birth, we obtained a new list of candidate sources with the middle-aged Monogem taking the lead. We showed then that Monogem is able to reproduce simultaneously all the experimental data, namely the e$^+$/(e$^-$+e$^+$), e$^-$+e$^+$, e$^-$ and e$^+$ energy spectra. However, Monogem is also in conflict with the upper bounds of dipole anisotropy set by Fermi. These calculations show, in fact, that the non-observation of anisotropy does not support the thesis of one single object dominating at high energy, even with the existence of an overwhelming background. But these calculations also indicate that there are other key players, specifically B1822-09 and Geminga. While Monogem and Geminga are located close together on the sky dome, B1822-09 is situated just in the opposite side. In this way, B1822-09 could very well negate the possible anisotropy caused by Monogem and Geminga. But to be in accordance with Fermi data, we had to consider in addition a significant isotropic background likely to arise from distant sources. Last but not least, we are aware of the shortcomings of this new software. The main line of thinking of this work is to build a workable program with a minimum version to save computing time, and make improvements afterwards. Future refinements will focus on the propagation of HECREs through more realistic magnetic field structures.

1607.06615

1607

1607.07049_arXiv.txt

We investigate the Milky Way Galaxy's radial and vertical metallicity gradients using a sample of 47,406 red clump stars from the RAdial Velocity Experiment (RAVE) Data Release (DR) 4. This sample is more than twice the size of the largest sample in the literature investigating radial and vertical metallicity gradients. The absolute magnitude of \citet{Groenewegen08} is used to determine distances to our sample stars. The resulting distances agree with the RAVE DR4 distances \citep{Binney14} of the same stars. Our photometric method also provides distances to 6185 stars that are not assigned a distance in RAVE DR4. The metallicity gradients are calculated with their current orbital positions ($R_{gc}$ and $Z$) and with their orbital properties (mean Galactocentric distance, $R_{m}$ and $z_{max}$), as a function of the distance to the Galactic plane: d[Fe/H]/d$R_{gc}=$-$0.047\pm0.003$ dex kpc$^{-1}$ for $0\leq |Z|\leq0.5$ kpc and d[Fe/H]/d$R_m=$-$0.025\pm0.002$ dex kpc$^{-1}$ for $0\leq z_{max}\leq0.5$ kpc. This reaffirms the radial metallicity gradient in the thin disc but highlights that gradients are sensitive to the selection effects caused by the difference between $R_{gc}$ and $R_{m}$. The radial gradient is flat in the distance interval 0.5-1 kpc from the plane and then becomes positive greater than 1 kpc from the plane. The radial metallicity gradients are also eccentricity dependent. We showed that d[Fe/H]/d$R_m=$-$0.089\pm0.010$, -$0.073\pm0.007$, -$0.053\pm0.004$ and -$0.044\pm0.002$ dex kpc$^{-1}$ for $e_p\leq0.05$, $e_p\leq0.07$, $e_p\leq0.10$ and $e_p\leq0.20$ sub-samples, respectively, in the distance interval $0\leq z_{max}\leq0.5$ kpc. Similar trend is found for vertical metallicity gradients. Both the radial and vertical metallicity gradients are found to become shallower as the eccentricity of the sample increases. These findings can be used to constrain different formation scenarios of the thick and thin discs.

How galaxies formed in general and how our Galaxy formed specifically remain as unsolved problem in astrophysics. In order to understand the formation mechanisms and put constraints on simulations, Galactic archaeology relies on chemo-dynamical information of various tracer objects with different ages and chemical compositions with large baselines in the Galaxy. While overlapping properties make it difficult to disentangle the different components in the Galaxy, there are clear chemical and kinematical signatures in stars of the Solar neighborhood which change with increasing distances both in radial and vertical directions. Sensitive metallicity gradients require tracer objects with well-measured distances and metallicities, which are obtained with either spectroscopic or photometric methods. Tracers include main-sequence stars, red clump (RC) stars, open or globular clusters, Cepheid variables, planetary nebulae, O-B type stars and HII regions. Various distance scales are used to determine radial and vertical metallicity gradients in the literature. In radial gradient calculations, if the tracer objects are distant, then the current distance from the Galactic centre is used ($R_{gc}$). If proper motions, radial velocities and distances of tracer objects are known, even though the objects are currently close to the Sun, then the mean Galactocentric distances of the tracer objects calculated from Galactic orbital parameters are used, under the assumption of that the metallicities are constant on the orbital integration timescale ($R_m$). Another distance scale presented to the literature is the guiding radius which is the distance of the tracer from the Galactic centre, if it is following its guiding centre, assumed to be a circular orbit ($R_{g}$). Also, current vertical distances and maximum vertical distances from the Galactic plane of the objects are generally considered in vertical metallicity gradient calculations. Radial and vertical metallicity gradient studies that appeared in the literature almost in the last 15 years are listed in Tables 1 and 2, respectively. The organization of the metallicity gradient information is based on the distance indicators mentioned in the previous paragraph. According to Table 1, radial metallicity gradients tend to have steeper values near the Galactic plane and they flatten and even have positive values as the distance of the tracers from the Galactic plane increases. This behavior holds for different tracer objects, as well. Radial metallicity gradients of young objects show steeper values than the ones calculated from older objects. Similar results are also found in Table 1 for $R_m$. Vertical metallicity gradients are found to be steeper than their radial counterparts. And besides, absolute vertical metallicity gradients found for the Galactic disc ($z<5$ kpc) give larger values than the ones found for the Galactic halo ($5<z<10$ kpc; see Table 2). This is also valid for objects with distances calculated from different methods. The {\it Hipparcos} \citep{ESA97} era calculated accurate distance estimations for stars in the Solar neighborhood. High-resolution spectroscopic follow-up observations with ground-based telescopes allowed astronomers to determine more precise metallicity gradients for the small {\it Hipparcos} volume in the Solar neighborhood \citep[i.e.][]{Nordstrom04}. To increase this limited volume, astronomers initiated several photometric and spectroscopic sky surveys, such as the RAdial Velocity Experiment \citep[RAVE;][]{Steinmetz06}, the Sloan Extension for Galactic Understanding and Exploration \citep[SEGUE;][]{Yanny09}, the APO-Galactic Evolution Experiment \citep[APOGEE;][]{Majewski10}, the Large sky Area Multi Object fiber Spectroscopic Telescope \citep[LAMOST;][]{Zhao12}, the GALactic Archaeology with HERMES \citep[GALAH;][]{Zucker12} and the Gaia-ESO Survey \citep[GES;][]{Gilmore12}. These surveys are designed to reveal the structure and test the formation mechanisms of the Milky Way's disc as well as to decode its evolution via metallicity, kinematics and dynamics of its stars. According to stellar evolution models, a star within a mass range of 0.8 to 2.2 $M_{\odot}$ could generate a helium (He) core mass of $\sim0.45M_{\odot}$ after the hydrogen burning reactions ceased \citep{Girardi99}. At this point, as stated by \citet{Carretta99}, when a star begins to evolve to the red giant phase, its surrounding envelope starts to interact with inner regions of the star, thus materials inside and in the outer surface of the star start mixing via convection processes. In a red giant, gravitational collapse of the star is supported via pressure of degenerate electrons, until the He core reaches 0.45 $M_{\odot}$. Then a series of He-flashes \citep{Thomas67} occurs and breaks the degeneracy. Right after the removal of degeneracy, luminosity of the star suddenly drops and stable He burning reactions begin. Star's luminosity remains almost constant. At this stage, if a star is metal rich then it becomes an RC star, or if it is not metal rich then it becomes a member of the blue horizontal branch stars \citep{Iben91}. Since the main-sequence mass of the RC stars are between 0.9 and 2.25 $M_{\odot}$ according to theory \citep{Eggleton68}, their main-sequence life times are between 13 to 1.3 Gyrs, respectively. Theory also predicts that the horizontal branch life times of stars is roughly 10\% of the main-sequence life times. This is the reason why these stars form a clump structure in Hertzsprung-Russell diagrams. There has been much debate on the nature of the RC stars and their use as a standard candle since their discovery. \citet {Cannon70} suggested that RC stars could be used for distance determination as standard candles. There are several studies which show that the absolute magnitude of the RC stars in different parts of the electromagnetic spectrum is always the same \citep {Paczynski98, Keenan99, Alves00, Udalski00, Groenewegen08, Laney12, Bilir13a, Bilir13b, Karaali13, Yaz13}: absolute magnitudes of the RC stars are in the near infrared $K_s$ band, $M_{Ks}=$-$1.54\pm 0.04$ mag \citep {Groenewegen08}, and in $I$ band $M_I=$-$0.23\pm 0.03$ mag \citep {Paczynski98}. RC stars are common objects in the Solar neighborhood and their intrinsic brightness probes large distances, which enables us to investigate chemical gradients both in radial and vertical directions. Rather than using spectroscopic parallaxes, we preferred photometric parallax method to calculate distances using the aforementioned well-known constant infrared absolute magnitude in the $K_s$ band. Metallicity gradients are calculated for the current orbital positions of the stars as well as for their calculated complete orbit parameters. In metallicity gradient calculations we used the mean Galactocentric distance of stellar orbits ($R_{m}$) instead of guiding radius ($R_{g}$). One other aspect of the study is that the dynamical properties of each star are calculated with {\em galpy} library of \citet {Bovy15} using test particle integration for {\it MWPotential2014} potential. In this paper, we present results on the metallicity gradients of the RC stars using data from RAVE DR4. In \S 2, we present the selection criteria of the RC stars and their distance, kinematic and dynamic calculations. In \S 3, we give the results on overall metallicity gradients obtained using various distance scales both in radial and vertical directions in the Galaxy. Also, a sub-sample separation with planar and vertical eccentricities are carried out and implications on radial migration of the RC stars are mentioned. We discussed the results on the RC stars in \S 4. \\ \begin{landscape} \textwidth = 650pt \begin{table*} \setlength{\tabcolsep}{2.5pt} \centering {\tiny \caption{Radial metallicity gradients appeared in the literature. Distances ($R_{gc}$, $Z$, $R_{m}$, $z_{max}$, $R_{g}$) in kpc, age ($\tau$) in Gyr.} \begin{tabular}{lcccccccr} \hline ~~~~~~~~~~~~~~Author & Tracer Object & d[Fe/H]/d$R_{gc}$ & Remark & d[Fe/H]/d$R_{m}$ & Remark & d[Fe/H]/d$R_{g}$ & Remark & Bibliographic Code \\ & &(dex~kpc$^{-1}$)& &(dex~kpc$^{-1}$)& &(dex~kpc$^{-1}$)& & \\ \hline \citet{Jacobson16} & OC & -$0.100\pm0.020$ & $6<R_{gc}<12$&$-$ &$-$ &$-$ &$-$ &2016A\&A...591A..37J\\ \citet{Cunha16} & OC/RG & -$0.035\pm0.007$ & $6<R_{gc}<25$&$-$ &$-$ &$-$ &$-$ &arXiv:1601.03099\\ \citet{Netopil16} & OC & -$0.086\pm0.009$ & $5<R_{gc}<12$&$-$ &$-$ &$-$ &$-$ &2016A\&A...585A.150N\\ & & -$0.082\pm0.013$ & $1<\tau<2.5$ &$-$ &$-$ &$-$ &$-$ &2016A\&A...585A.150N\\ \citet{Plevne15} & FG MSs & $-$ & $-$ & -$0.083\pm0.030$& $0<z_{max}\leq0.5$ ([Fe/H]) & -$0.083\pm0.030$ & $0<z_{max}\leq0.5$ ([Fe/H]) & 2015PASA...32...43P \\ & & $-$ & $-$ & -$0.048\pm0.037$& $0.5<z_{max}\leq0.8$ ([Fe/H]) & -$0.065\pm0.039$ & $0.5<z_{max}\leq0.8$ ([Fe/H]) & 2015PASA...32...43P\\ & & $-$ & $-$ & -$0.063\pm0.011$& $0<z_{max}\leq0.5$ (M/H]) & -$0.062\pm0.018$ & $0<z_{max}\leq0.5$ ([M/H]) & 2015PASA...32...43P \\ & & $-$ & $-$ & -$0.028\pm0.057$& $0.5<z_{max}\leq0.8$ ([M/H]) & -$0.055\pm0.045$ & $0.5<z_{max}\leq0.8$ ([M/H]) & 2015PASA...32...43P \\ \citet{Huang15} & RCs & -$0.082\pm0.003$ & $|Z|\leq0.1$ &$-$ &$-$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & -$0.072\pm0.004$ & $0.1<|Z|\leq0.3$ &$-$ &$-$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & -$0.052\pm0.005$ & $0.3<|Z|\leq0.5$ &$-$ &$-$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & -$0.041\pm0.005$ & $0.5<|Z|\leq0.7$ &$-$ &$-$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & -$0.028\pm0.005$ & $0.7<|Z|\leq0.9$ &$-$ &$-$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & -$0.020\pm0.007$ & $0.9<|Z|\leq1.1$ &$-$ &$-$ &$-$ &$-$ & 2015RAA....15.1240H\\ \citet{Xiang15} & Turnoff stars & -$0.100\pm0.003$ & $|Z|\leq0.1; 2<\tau<16$ &$-$ &$-$ &$-$ &$-$ & 2015RAA....15.1209X\\ & & -$0.050\pm0.002$ & $0.4<|Z|\leq0.6; 2<\tau<16$ &$-$ &$-$ &$-$ &$-$ & 2015RAA....15.1209X\\ & & -$0.010\pm0.002$ & $0.9<|Z|\leq1.1; 2<\tau<16$ &$-$ &$-$ &$-$ &$-$ & 2015RAA....15.1209X\\ \citet{Recio14} & FGK MSs & -$0.058\pm0.008$ & Thin disc ([M/H]) &$-$ &$-$ &$-$ &$-$ & 2014A\&A...567A...5R\\ & & +$0.006\pm0.008$ & Thick disc ([M/H]) &$-$ &$-$ &$-$ &$-$ & 2014A\&A...567A...5R\\ \citet{Mikolaitis14} & FGK MSs & -$0.044\pm0.009$ & Thin disc (main) &$-$ &$-$ &$-$ &$-$ & 2014A\&A...572A..33M\\ & & -$0.028\pm0.018$ & Thin disc (clean) &$-$ &$-$ &$-$ &$-$ & 2014A\&A...572A..33M\\ & & +$0.008\pm0.007$ & Thick disc (main) &$-$ &$-$ &$-$ &$-$ & 2014A\&A...572A..33M\\ & & +$0.008\pm0.007$ & Thick disc (clean) &$-$ &$-$ &$-$ &$-$ & 2014A\&A...572A..33M\\ \citet{Genovali14} & Cepheids & -$0.021\pm0.029$ & $5<R_{gc}<19$ & $-$ & $-$ & $-$ & $-$ & 2014A\&A...566A..37G \\ \citet{Andreuzzi14} & OC & -$0.04$ & $6<R_{gc}<22$ & $-$ & $-$ & $-$ & $-$ & 2011MNRAS.412.1265A \\ \citet{Boeche14} & RCs & $-$&$-$ &$-$ & $-$ &-$0.054\pm 0.004$ & $ 0<|Z|\leq0.4$ & 2014A\&A...568A..71B\\ & & $-$&$-$ &$-$ & $-$ &-$0.039\pm 0.004$ & $0.4<|Z|\leq0.8$ & 2014A\&A...568A..71B\\ & & $-$&$-$ &$-$ & $-$ &-$0.011\pm 0.008$ & $0.8<|Z|\leq1.2$ & 2014A\&A...568A..71B\\ & & $-$&$-$ &$-$ & $-$ &+$0.047\pm 0.018$ & $1.2<|Z|\leq2.0$ & 2014A\&A...568A..71B\\ \citet{Hayden14} & RG & -$0.090\pm 0.002$ & $0.00<|Z|<0.25$; $0<R_{gc}<15$ & $-$& $-$& $-$& $-$& 2014AJ....147..116H\\ & & -$0.076\pm 0.003$ & $0.25<|Z|<0.50$; $0<R_{gc}<15$ & $-$& $-$& $-$& $-$& 2014AJ....147..116H\\ & & -$0.057\pm 0.003$ & $0.50<|Z|<1.00$; $0<R_{gc}<15$ & $-$& $-$& $-$& $-$& 2014AJ....147..116H\\ & & -$0.030\pm 0.006$ & $1.00<|Z|<2.00$; $0<R_{gc}<15$ & $-$& $-$& $-$& $-$& 2014AJ....147..116H\\ \citet{Frinchaboy13} & OC & -$0.090\pm 0.030$ & $7.9\leq R_{gc} \leq 14.5$ & $-$ & $-$ & $-$ & $-$ & 2013ApJ...777L...1F\\ & & -$0.200\pm 0.080$ & $7.9\leq R_{gc} \leq 10.0$ & $-$ &$-$ & $-$& $-$& 2013ApJ...777L...1F\\ & & -$0.020\pm 0.090$ & $10<R_{gc} \leq 14.5$ & $-$ &$-$ & $-$& $-$& 2013ApJ...777L...1F\\ \citet{Boeche13} & FG MSs & -$0.059\pm 0.002$ & $4.5<R_{gc}<9.5$ & $-$& $-$& -$0.065\pm0.003$ & $0<z_{max}\leq0.4$ & 2013A\&A...559A..59B\\ & & $-$ & $-$ & $-$&$-$ & -$0.059\pm0.006$ & $0.4<z_{max}\leq0.8$ & 2013A\&A...559A..59B\\ & & $-$ & $-$ & $-$&$-$ & +$0.006\pm0.015$ & $z_{max}>0.8$ & 2013A\&A...559A..59B\\ \citet{Carrell12} & FGK MSs & +$0.015\pm 0.005$ & $0.5<|Z|<1.0$, Isoc. method &$-$ &$-$ &$-$ &$-$ & 2012AJ....144..185C\\ & & +$0.017\pm 0.005$ & $0.5<|Z|<1.0$, Photo. method &$-$ &$-$ &$-$ &$-$ & 2012AJ....144..185C\\ \citet{Bilir12} & RCs & $-$& $-$ &-$0.041\pm0.003$ & $TD/D\leq 0.1$ &$-$ &$-$ & 2012MNRAS.421.3362B\\ & & $-$& $-$ &-$0.041\pm0.007$ & $e_v\leq 0.07$ &$-$ &$-$ & 2012MNRAS.421.3362B\\ & & $-$& $-$ &-$0.025\pm0.040$ & $e_v\leq 0.12$ &$-$ &$-$ & 2012MNRAS.421.3362B\\ & & $-$& $-$ &+$0.022\pm0.006$ & $e_v>0.25$ &$-$ &$-$ & 2012MNRAS.421.3362B\\ \citet{Cokkun12} & FG MSs & $-$& $-$ &-$0.043\pm0.005$ & F stars, $TD/D\leq0.1 $&$-$ &$-$ & 2012MNRAS.419.2844C\\ & & $-$& $-$ &-$0.033\pm0.007$ & G stars, $TD/D\leq0.1 $&$-$ &$-$ & 2012MNRAS.419.2844C\\ & & $-$& $-$ &-$0.051\pm0.005$ & F stars, $e_v\leq 0.04 $&$-$ &$-$ & 2012MNRAS.419.2844C\\ & & $-$& $-$ &-$0.020\pm0.006$ & G stars, $e_v\leq 0.04 $&$-$ &$-$ & 2012MNRAS.419.2844C\\ \citet{Cheng12} & Turnoff stars & -0.066$^{+0.030}_{-0.044}$ & $0.15<|Z|<0.25$ &$-$ &$-$ &$-$ &$-$ & 2012ApJ...746..149C\\ & & -0.064$^{+0.015}_{-0.004}$ & $0.25<|Z|<0.50$ &$-$ &$-$ &$-$ &$-$ & 2012ApJ...746..149C\\ & & -0.013$^{+0.009}_{-0.002}$ & $0.50<|Z|<1.00$ &$-$ &$-$ &$-$ &$-$ & 2012ApJ...746..149C\\ & & +0.028$^{+0.007}_{-0.005}$ & $1.00<|Z|<1.50$ &$-$ &$-$ &$-$ &$-$ & 2012ApJ...746..149C\\ \citet{Yong12} & OC & -0.120$^{+0.010}_{-0.140}$ & $\tau \geq 2.5$; $R_{gc}<13$ &$-$ &$-$ &$-$ &$-$ & 2012AJ....144...95Y \\ \citet{Karatas12} & MSs & $-$&$-$ & -$0.070\pm0.010$ &$0<z_{max}<5$ &$-$ &$-$ & 2012NewA...17...22K\\ \citet{LucKnLambert11} & Cepheids & -$0.062\pm0.002$ & $4<R_{gc}<17$ &$-$ &$-$ &$-$ &$-$ & 2011AJ....142..136L \\ \citet{Ruchti11} & RGB,RCs,HB,MS,SG & +$0.010\pm0.040$ & $5<R_{gc}<10$, $|Z|<3$ &$-$ &$-$ &$-$ &$-$ & 2011ApJ...737....9R \\ \citet{Luck11} & Cepheids & -$0.055\pm0.003$ & $4<R_{gc}<16$& &$-$ &$-$ &$-$ & 2011AJ....142...51L \\ \citet{Friel10} & OC & -$0.076\pm0.018$ & all sample &$-$ &$-$ &$-$ &$-$ & 2010AJ....139.1942F \\ \citet{Wu09} & OC & -$0.070\pm0.011$ & $6\leq R_{gc}\leq13.5$ & $-0.082\pm0.014$ &$6\leq R_{m}\leq20$ &$-$ &$-$ & 2009MNRAS.399.2146W \\ \citet{Pedicelli09} & Cepheids & -$0.051\pm0.004$ & all sample &$-$ &$-$ &$-$ &$-$ & 2009A\&A...504...81P \\ \citet{Magrini09} & OC & -$0.053\pm0.029$ & $7<R_{gc}<12$, $\tau \leq0.8$ Gyr&$-$ &$-$ &$-$ &$-$ & 2009A\&A...494...95M \\ \citet{Lemasle08} & Cepheids & -$0.052\pm0.003$ & $5<R_{gc}<17$ &$-$ &$-$ &$-$ &$-$ & 2008A\&A...490..613L \\ \citet{Lemasle07} & Cepheids & -$0.061\pm0.019$ & $8<R_{gc}<12$ &$-$ &$-$ &$-$ &$-$ & 2007A\&A...467..283L \\ \citet{Allende06} & FG MSs & 0& $1<|Z|\leq3$ & $-$ & $-$ &$-$ &$-$ &2006ApJ...636..804A \\ \citet{Nordstrom04} & FG MSs & $-$&$-$ & -$0.076\pm0.014$ & $\tau\leq1.5$ &$-$ &$-$ & 2004A\&A...418..989N \\ & & $-$&$-$ & -$0.099\pm0.011$ & $4<\tau\leq6 $ &$-$ &$-$ &2004A\&A...418..989N\\ & & $-$&$-$ & +$0.028\pm0.036$ & $\tau>10 $ &$-$ &$-$ &2004A\&A...418..989N\\ \citet{Chen03} & OC & -$0.063\pm0.008$ & $6<R_{gc}<15$ & &$-$ &$-$ &$-$ & 2003AJ....125.1397C \\ \citet{Hou02} & OC & -$0.099\pm0.008$ & $6.5<R_{gc}<16$ &$-$ &$-$ &$-$ &$-$ & 2002ChJAA...2...17H \\ \citet{Friel02} & OC & -$0.060\pm0.010$ & $7<R_{gc}<16$ &$-$ &$-$ &$-$ &$-$ & 2002AJ....124.2693F \\ & & -$0.023\pm0.019$ & $\tau<2$ &$-$ &$-$ &$-$ &$-$ & 2002AJ....124.2693F \\ & & -$0.053\pm0.018$ & $2\leq \tau \leq 4$ &$-$ &$-$ &$-$ &$-$ & 2002AJ....124.2693F \\ & & -$0.075\pm0.019$ & $\tau>4$ &$-$ &$-$ &$-$ &$-$ & 2002AJ....124.2693F \\ \citet{Andrievsky02} & Cepheids & -$0.130\pm0.030$ & inner disc &$-$ &$-$ &$-$ &$-$ & 2002A\&A...392..491A \\ \hline \end{tabular} \\ Abbreviations: Main Sequence Stars: MSs; Red Clump Stars: RCs; Red Giants: RG; Open Cluster: OC; Red Giant Branch: RGB; Horizontal Branch: HB; Sub-giant: SG } \end{table*} \end{landscape} \begin{landscape} \textwidth = 650pt \begin{table*} \setlength{\tabcolsep}{3pt} \centering {\scriptsize \caption{Vertical metallicity gradients appeared in the literature. Distances ($d$, $R_{gc}$, $Z$, $z_{max}$) in kpc, age ($\tau$) in Gyr.} \begin{tabular}{lcccccr} \hline ~~~~~~~~~~~~~~Author &Tracer Object&d[Fe/H]/d$|Z|$ &Remark&d[Fe/H]/d$z_{max}$ &Remark&Bibliographic Code \\ & &(dex~kpc$^{-1}$)& &(dex~kpc$^{-1}$)& & \\ \hline \citet{Plevne15} & F-G MSs & $-$ & $-$ &-$0.176\pm0.039$ & $0<z_{max}\leq 0.825$ & 2015PASA...32...43P \\ & & $-$ & $-$ &-$0.119\pm0.036$ & $0<z_{max}\leq 1.5$ & 2015PASA...32...43P\\ \citet{Xiang15} & Turnoff stars & -$0.110\pm0.020$ & $\tau>11$ & $-$ & $-$ & 2015RAA....15.1209X\\ \citet{Huang15} & RC & -$0.206\pm0.006$ & $|Z|\leq1,7<R_{gc}\leq8$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & -$0.116\pm0.008$ & $|Z|\leq1,8<R_{gc}\leq9$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & -$0.052\pm0.010$ & $|Z|\leq1,9<R_{gc}\leq10$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & $~0.000\pm0.012$ & $|Z|\leq1,10<R_{gc}\leq11$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & +$0.008\pm0.008$ & $|Z|\leq1,11<R_{gc}\leq12$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & +$0.057\pm0.012$ & $|Z|\leq1,12<R_{gc}\leq13$ &$-$ &$-$ & 2015RAA....15.1240H\\ & & +$0.047\pm0.012$ & $|Z|\leq1,13<R_{gc}\leq14$ &$-$ &$-$ & 2015RAA....15.1240H\\ \citet{Mikolaitis14} & FGK MSs & -$0.107\pm0.009$ & Thin disc (main) &$-$ &$-$ & 2014A\&A...572A..33M\\ & & -$0.057\pm0.016$ & Thin disc (clean) &$-$ &$-$ & 2014A\&A...572A..33M\\ & & -$0.072\pm0.006$ & Thick disc (main) &$-$ &$-$ & 2014A\&A...572A..33M\\ & & -$0.037\pm0.016$ & Thick disc (clean)&$-$ &$-$ & 2014A\&A...572A..33M\\ \citet{Schlesinger14} & G MSs & -$0.243^{+0.039}_{-0.053}$ &all sample& $-$&$-$ & 2014ApJ...791..112S \\ \citet{Boeche14} & RCs & -$0.112\pm0.007$ & $0<|Z|\leq2$ &$-$ &$-$ &2014A\&A...568A..71B \\ \citet{Hayden14} & RG & -$0.305\pm0.011$ & $0<|Z|\leq2$ &$-$ &$-$ & 2014AJ....147..116H \\ \citet{Bergemann14} & FGK MSs & -$0.068\pm0.014$ & $|Z|\leq0.3$ &$-$ &$-$ & 2014A\&A...565A..89B \\ & & -$0.114\pm0.009$ & $0.3<|Z|\leq0.8$ &$-$ &$-$ & 2014A\&A...565A..89B \\ \citet{Carrell12} & FGK MSs & -$0.113\pm0.010$ & $7<R_{gc}<10.5$, Isoc. method &$-$ &$-$ &2012AJ....144..185C \\ & & -$0.125\pm0.008$ & $7<R_{gc}<10.5$, Photo. method &$-$ &$-$ &2012AJ....144..185C \\ \citet{Bilir12} & RCs & $-$ &$-$ &-$0.109\pm0.008$ & $TD/D\leq0.1$ & 2012MNRAS.421.3362B \\ & & $-$ &$-$ &-$0.260\pm0.031$ & $e_v\leq0.07$ & 2012MNRAS.421.3362B \\ & & $-$ &$-$ &-$0.167\pm0.011$ & $e_v\leq0.12$ & 2012MNRAS.421.3362B \\ & & $-$ &$-$ &-$0.022\pm0.005$ & $e_v>0.25$ & 2012MNRAS.421.3362B \\ \citet{Peng12} & MSs & -$0.210\pm0.050$ & $0<Z<2$ &$-$ &$-$ & 2012MNRAS.422.2756P \\ \citet{Katz11} & RGB,RCs,HB,MS,SG & -$0.068\pm0.009$ & Thick disc &$-$ &$-$ & 2011A\&A...525A..90K \\ \citet{Chen11} & RHB & -$0.120\pm0.010$ & $0.5<|Z|<3$ &$-$ &$-$ & 2011AJ....142..184C \\ & & -$0.220\pm0.070$ & $1<|Z|<3$ &$-$ &$-$ & 2011AJ....142..184C \\ \citet{Kordopatis11} & FGK MSs & -$0.140\pm0.050$ & $1\leq Z\leq 4$ &$-$ &$-$ & 2011A\&A...535A.107K \\ \citet{Ruchti11} & RGB,RCs,HB,MS,SG & -$0.090\pm0.050$ & $6<R_{gc}<10$, $|Z|<3$ &$-$ &$-$ & 2011ApJ...737....9R \\ \citet{Yaz10} & G MSs & -$0.320\pm0.010$ & $Z<2.5$ & $-$ &$-$ & 2010NewA...15..234Y \\ & & -$0.300\pm0.060$ & $3<Z<5.5$ & $-$ &$-$ & 2010NewA...15..234Y \\ & & -$0.010\pm0.010$ & $6<Z<10$ & $-$ &$-$ & 2010NewA...15..234Y \\ \citet{Siegel09} & FS & -$0.150$ & $Z<4$ &$-$ & $-$ & 2009MNRAS.395.1569S \\ \citet{Soubiran08} & RCs & -$0.300\pm0.030$ & $d<1$ &$-$ & $-$ & 2008A\&A...480...91S \\ & & $-$ & $-$ &-$0.310\pm0.060$ & $0\leq z_{max}<1.2$ & 2008A\&A...480...91S \\ \citet{Ak07a} & G MSs & -$0.380\pm0.060$ & $3 \leq Z<5$ &$-$ &$-$ & 2007NewA...12..605A \\ & & -$0.080\pm0.070$ & $5 \leq Z<10$ &$-$ &$-$ & 2007NewA...12..605A \\ \citet{Ak07b} & G MSs & -$0.160\pm0.020$ & $Z<3$ North &$-$ &$-$ & 2007AN....328..169A \\ & & -$0.220\pm0.020$ & $Z<3$ South & $-$&$-$ &2007AN....328..169A \\ \citet{Allende06} & FG MSs & 0.030&$1<|Z|\leq3$; [Fe/H]$>-1.2$ &$-$ &$-$&2006ApJ...636..804A\\ \citet{Marsakov06} & FS & -$0.290\pm0.060$&Thin disc&$-$ &$-$& 2006A\&AT...25..157M\\ \citet{Chen03} & OC & $-$ &$-$ &-$0.295\pm0.050$ &$6<R_{gc}<15$ &2003AJ....125.1397C \\ \citet{Karaali03} & G MSs & -$0.200$ & $Z\leq8$ & $-$ & $-$ & 2003MNRAS.343.1013K \\ \citet{Barta03} & FS & -$0.230\pm0.040$ & $Z<1.3$ & $-$ & $-$ & 2003BaltA..12..539B \\ \hline \end{tabular} \\ Abbreviations: Main Sequence Stars: MSs; Red Clump Stars: RCs; Red Giants: RG; Open Cluster: OC; Red Giant Branch: RGB; Horizontal Branch: HB; Red Horizontal Branch: RHB; Sub-giant: SG; Field Stars: FS. } \end{table*} \end{landscape}

We have investigated the Milky Way Galaxy's radial and vertical metallicity gradients using a sample of 47,406 red clump (RC) stars from the RAdial Velocity Experiment (RAVE) Data Release (DR) 4. The largest samples in the literature for investigating radial and vertical metallicity gradients are in \citet{Hayden14} and \citet{Boeche14}, who analyzed $\sim$20,000 stars, less than half of our sample. Our sample was selected using the following constraints: (i) T$_{\rm eff}$ and log $g$ from the RAVE DR4 were used to select the RC stars, $1-\sigma$ from the peak density of the RC; (ii) the proper motions used in their space velocity estimations are available in the UCAC4 catalogue \citep{Zacharias13}, ii) the iron abundance is available in the RAVE DR4 catalogue \citep{Kordopatis13}; (iii) $S/N\geq40$; and (iv) total space velocity errors would be less or equal to 21 km s$^{-1}$. The mean absolute magnitude of the RC stars by \citet{Groenewegen08} ($M_{K_s}$=-1.54$\pm$0.04 mag) is used to determine distances to our sample stars. The resulting distances agree with the RAVE DR4 distances \citep{Binney14} of the same stars (mean and standard deviation of difference is 87 and 220 pc respectively). Our photometric method also provides distances to 6185 star (13\% of our sample) for which distances are not assigned in RAVE DR4. The metallicity gradients are calculated with their current orbital positions ($R_{gc}$ and $Z$) and with their orbital properties (mean Galactocentric distance, $R_{m}$ and $z_{max}$), as a function of the distance to the Galactic plane: d[Fe/H]/d$R_{gc}$=-0.047$\pm$0.003 dex kpc$^{-1}$ for $0\leq |Z|\leq0.5$ kpc and d[Fe/H]/d$R_m$=-0.025$\pm$0.002 dex kpc$^{-1}$ for $0\leq z_{max}\leq0.5$ kpc. This reaffirms the radial metallicity gradient in the thin disc but highlights that gradients are sensitive to the selection effects caused by the difference between $R_{gc}$ and $R_{m}$. The radial gradient is flat in the interval 0.5-1 kpc from the plane and then becomes positive for the distances greater than 1 kpc from the plane. Both the radial and vertical metallicity gradients are found to become shallower as the eccentricity of the sample increases. The radial metallicity gradients in terms of Galactocentric distance, for the interval of $0\leq|Z|\leq0.5$ kpc, d[Fe/H]/d$R_{g}$=-0.047$\pm$0.003 dex kpc$^{-1}$ is compatible with the one of \citet{Boeche14}, i.e. d[Fe/H]/d$R_{gc}$=-0.054$\pm$0.004 dex kpc$^{-1}$ within the errors, estimated for the RC stars with $5.5<R_{gc}<11$ kpc. The radial metallicity gradient in terms of mean orbital distance, d[Fe/H]/d$R_m$=-0.025$\pm$0.002 dex kpc$^{-1}$ is flat relative to the one of \citet{Bilir12}, d[Fe/H]/d$R_m$=-0.041$\pm$ 0.007 dex kpc$^{-1}$ estimated for the RC stars. However, \citet{Bilir12} restricted the vertical eccentricities of their star sample with $e_v\leq 0.07$. $e_v$ is proportional to $z_{max}$ and steeper gradients are found near the plane. The vertical metallicity gradient estimated for the present position of all RC stars, d[Fe/H]/d$|Z|$=-0.219$\pm$0.003 dex kpc$^{-1}$, is steeper than \citet{Boeche14}'s, d[Fe/H]/d$|Z|$=-0.112$\pm$0.007 dex kpc$^{-1}$. As the vertical metallicity gradient gets steeper further from the plane, our steeper result suggests we are sampling a higher number of stars further from the plane than \citet{Boeche14}. The vertical metallicity gradient in terms of the $z_{max}$ distance estimated in this study, d[Fe/H]/d$z_{max}$= -0.157$\pm$0.002 dex kpc$^{-1}$, is comparable with the one in \citet{Plevne15}, i.e. d[Fe/H]/d$z_{max}$= -0.176$\pm$0.039 dex kpc$^{-1}$. The difference between them may originate in the different $z_{max}$ intervals that the two gradients were evaluated, $0<z_{max}<3$ and $0<z_{max}\leq0.8$ kpc in our study and in \citet{Plevne15}, respectively. The d[Fe/H]/d$z_{max}$ metallicity gradients estimated for the RC stars in \citet{Bilir12} which are also given in Table 2 are different than the corresponding one in our study, due to additional constraints in \citet{Bilir12}. Table 5 shows that the radial metallicity gradient depends on the eccentricity. Actually, the flat value of d[Fe/H]/d$R_m$= -0.025$\pm$0.002 dex kpc$^{-1}$ could be reduced to steeper values, -0.089$\pm$0.010, -0.072$\pm$0.007, -0.053$\pm$0.004 and -0.044$\pm$0.002 dex kpc$^{-1}$ by applying constraints $e_p\leq0.05$, 0.07, 0.10 and 0.20 in $0\leq z_{max}\leq0.5$ kpc distance interval, respectively. By this limitation, low $e_p$ steepen the gradients compared to including higher $e_p$. This could be due to many effects. $e_p$ is expected to increase with age and older stars tend to be more metal-poor so including higher $e_p$ stars includes more metal-poor stars, which will flatten the gradient. Similarly, thick disc stars tend to have higher $e_p$ than thin disc stars so including higher $e_p$ stars includes more thick disc stars, which also tend to be more metal-poor, again flattening the gradient. We plotted the RC stars in the sub-samples defined by their $e_p$ eccentricities in four Toomre diagrams and investigated their behaviour in terms of space velocities. The Toomre diagrams in Fig. \ref{Fig16} cover the RC stars with distances $0\leq z_{max}\leq0.5$, $0.5<z_{max}\leq1$, $1<z_{max}\leq2$, and $z_{max}>2$ kpc. There is a trend between the space velocity components $U$ and $W$, and the $e_p$ eccentricities of the stars in all panels, i.e. $(U^2+W^2)^{1/2}$ increases with increasing $e_p$. Surprisingly, the total space velocity remains constant, which it suggests that the rotational velocity of the RC sample decreases with increasing $z_{max}$ distances. On the contrary, $(U^2+W^2)^{1/2}$ velocities of the stars for a given sub-sample increase with increasing $z_{max}$ distances. For example, $(U^2+W^2)^{1/2}\leq50$ km s$^{-1}$ for the stars with $e_p\geq0.20$ in the $0\leq z_{max}\leq0.5$ interval, while it is $(U^2+W^2)^{1/2}<100$ km s$^{-1}$ in the $1<z_{max}\leq2$ interval. Thus, we can say that the space velocity components of the RC stars at the same $z_{max}$ distance, but with different $e_p$ eccentricities, are different. The same argument holds for the stars with the same eccentricities with different $z_{max}$ distances. If we assume that different velocities originate from the intergalactic gas clouds of different angular momentum as well as different chemical structure, then we can expect different metallicity gradients for different sub-samples, as in case of this study. \begin{figure} \centering \includegraphics[width=9 cm,height=15.11cm]{Fig16.eps} \caption{Toomre energy diagram of 47,406 RC stars in four $z_{max}$ intervals. Pink, red, green, blue and grey circles represent $e_p\leq0.05$, $e_p\leq0.07$, $e_p\leq0.10$, $e_p\leq0.20$ and $e_p\leq1$ samples. Black solid lines show total space velocity borders of 50, 100 and 150 km s$^{-1}$, respectively.} \label{Fig16} \end {figure} These findings can be used to constrain different formation scenarios of the thick and thin discs. The next step will be to repeat this analysis with improved astrometry, including trigonometric parallaxes, measured by {\it Gaia}. {\it Gaia}'s first data release will include these measurements for Tycho-2 stars \citep{Michalik15}, which are the brighter half of RAVE stars.

1607.07049

1607

1607.00343_arXiv.txt

We present the results of a {\it NuSTAR} study of the dynamically confirmed stellar-mass black hole GS 1354$-$645. The source was observed during its 2015 ``hard'' state outburst; we concentrate on spectra from two relatively bright phases. In the higher-flux observation, the broadband {\it NuSTAR} spectra reveal a clear, strong disk reflection spectrum, blurred by a degree that requires a black hole spin of $a = cJ/GM^2 \geq 0.98$ ($1\sigma$ statistical limits only). The fits also require a high inclination: $\theta \simeq 75(2)^{\circ}$. Strong ``dips'' are sometimes observed in the X-ray light curves of sources viewed at such an angle; these are absent, perhaps indicating that dips correspond to flared disk structures that only manifest at higher accretion rates. In the lower-flux observation, there is evidence of radial truncation of the thin accretion disk. We discuss these results in the context of spin in stellar-mass black holes, and inner accretion flow geometries at moderate accretion rates.

Low-mass X-ray binaries are binary systems comprised of a compact object accreting matter from a low-mass companion star. The accretion disk is luminous across the electromagnetic band pass, but it peaks in the X-ray band. The hot corona contributes in the hard X-ray band. Without a need for strong bolometric corrections, X-ray studies of these sources can accurately constrain the geometry of the inner accretion flow and its energetic properties (for a review of stellar-mass black holes in such settings, see Remillard \& McClintock 2006). Owing to their proximity, Galactic binaries are excellent laboratories for probing the extreme gravitational effects of black holes. Astrophysical black holes can be fully described by their mass and ``spin'' (dimensionless angular momentum; $a = cJ/GM^{2}$, where $-1 \leq a \leq 1$). The location of the innermost stable circular orbit (ISCO) around the black hole depends on the spin parameter. For a Schwarzschild black hole (zero spin), the radius of the the ISCO is at $R_{ISCO} = 6~GM/c^2$. In an extreme Kerr black hole (maximally spinning), $R_{ISCO} \simeq 1 GM/c^2$ (e.g., Bardeen, Press, \& Teukolsky 1972). Hard X-rays produced via Componization and/or sychrotron in the corona are ``reflected'' from the accretion disk, and this effect can be used to measure the spin of both stellar-mass and supermassive black holes (for recent reviews, see, e.g., Miller 2007, Reynolds 2014, Miller \& Miller 2015). The reflection spectrum is ``blurred'' by the strong Doppler shifts and gravitational red-shifts close to the black hole, effectively tracing the radius of the ISCO and the spin of the black hole. Resolution is important in such studies, but sensitivity and a broad spectral band pass are also very important. {\it NuSTAR} (Harrison et al.\ 2013) has unprecedented sensitivity in the 3-79 keV band and does not suffer from distorting effects such as photon pile-up; it is an ideal mission for studies of reflection and black hole spin (see, e.g., Risaliti et al.\ 2013; Miller et al.\ 2013a,b; King et al.\ 2014; Parker et al.\ 2014; Tomsick et al.\ 2014; Furst et al.\ 2015). Other important measurements can be obtained through reflection modeling. Among these are the inclination of the innermost accretion disk, though models can also measure the ionization of the inner disk, and can also constrain elemental abundances. In the best cases, the size of the hard X-ray corona can even be constrained (e.g., Miller et al.\ 2015). Progress has not only relied upon improved X-ray instrumentation, but also improved reflection models. In particular, \texttt{relxill} (e.g., Garcia et al.\ 2013, Dauser et al.\ 2014) offers many advantages over prior models, including internal relativistic blurring and angle-dependent scattering calculations. The object of this work, GS 1354-645, is a binary system comprising a dynamically confirmed black hole of mass $M_{BH} \geq ~7.6(7)~ M_\odot$ (Casares et al. 2009) and a low mass stellar companion. It was first detected using the All Sky Monitor (ASM) aboard the {\it Ginga} satellite during an outburst in 1987 (Kitamoto et al.\ 1990). The last outburst of this source was detected using {\it RXTE} in 1997 (e.g. Revnivtsev et al.\ 2000, Brocksopp et al 2001; also see Reynolds \& Miller 2013). The distance to the source is not well-constrained, with estimates lying between 25 and 61 kpc (Casares et al. 2009). Monitoring with the {\it Swift}/BAT detected a new outburst of GS~1354$-$645 in June 2015 (Miller, Reynolds, Kennea 2015). Better-known, recurrent sources like GX~339$-$4 often show multiple spectral states, but GS 1354$-$645 is interesting in that only the ``hard'' state was observed in its last outburst. GS~1354$-$645 may therefore belong to a small sub-class of black hole transients including the better-known GRO J0422$+$32 and XTE~J1118$+$480 (Brocksopp et al.\ 2001). Disk reflection was clearly required in {\it RXTE} spectra of the 1997 outburst of GS~1354$-$645 (Revnivtsev et al.\ 2000), potentially indicating a means to study the black hole and innermost accretion flow in a ``hard state'' transient. We therefore requested observations with {\it NuSTAR}.

We have analyzed two {\it NuSTAR} spectra of the dynamically confirmed black hole GS~1354$-$645. Both observations were obtained in the ``low/hard'' state. When fit with a relativistic reflection model, the spectrum obtained close to the outburst peak suggests a high spin parameter, and also implies that the inner disk may be viewed at a high inclination. The spectrum obtained at a lower flux level suggests that the inner disk may have been truncated, potentially consistent with radiatively inefficient accretion flow models. In this section, we discuss the strengths and weaknesses of our results, potential sources of systematic errors, and impacts on our understanding of GS 1354$-$645. Fits to Obs.\ 2 with \texttt{relxill} indicate that the accretion disk likely extends very close to the black hole. This is now common in the most luminous phases of the ``low/hard'' state, especially at the sensitivity afforded by {\it NuSTAR} (see, e.g., Miller et al.\ 2015). The data strongly require a very high spin parameter, $a = 0.998_{-0.009}$, consistent with the extreme upper bound of the model. The error is merely the statistical error, and systematic errors are likely to be much larger. All measures of black hole spin obtained through the accretion disk rely on the optically thick reflecting gas obeying the test particle ISCO. Simulations suggest that the disk is likely to be thin and to obey the ISCO at Eddington fractions below 0.3 (Shafee et al.\ 2008, Reynolds \& Fabian 2008); for plausible combinations of black hole mass and distance, this condition was met in our observations. It is quite possible, however, that no astrophysical disk respects the ISCO at the percent level. The best-fit model for Obs.\ 2 yields parameters similar to those expected in a ``lamp-post'' geometry (a compact, central source of hard X-rays located on the spin axis above the black hole). The \texttt{relxilllp} model (Garcia et al.\ 2013; Dauser et al.\ 2014) explictly assumes this geometry and calculates the reflection fraction self-consistently; it also gives a high spin parameter ($a > 0.85$). However, our best-fit model (see Table 1) is superior at the $6.6\sigma$ level of confidence, as determined by an F-test (for \texttt{relxilllp}, $\chi^{2}/\nu = 2806/2733$, even if the reflection fraction is not linked to the lamp-post value). This might imply that the corona is indeed compact but not quite an idealized lamp-post. Recent work has noted some physical difficulties with idealized lamp-post models (e.g., Niedzwiecki et al.\ 2016, Vincent et al.\ 2016). Stiele \& Kong (2016) have reported a nearly maximal retrograde spin based on a short {\it XMM-Newton} observation of GS 1354$-$645. A combination of factors including calibration uncertainties in the EPIC-pn camera and photon pile-up may have acted to falsely narrow the reflection features in the {\it XMM-Newton} spectrum (Miller et al.\ 2010). It is also possible that the disk was mildly truncated during the {\it XMM-Newton} observation. Fits to {\it NuSTAR} Obs.\ 2 with $a \leq 0.93$ are rejected at the $5\sigma$ level of confidence, via an F-test. Our reflection modeling also indicates that the inner disk is viewed at a relatively high inclination, $i = 75(2)$ degrees. This is within the eclipse limit derived by Casares et al.\ (2009). It is possible that the inner disk is not aligned with the inclination of the binary system itself (Maccarone 2002; also see Tomsick et al.\ 2014). Systems that narrowly avoid eclipses are often observed to exhibit ``dips'' in their X-ray light curve. These events may be due to vertical structures in the outer accertion disk that block some of the light from the central engine (see, e.g., Diaz-Trigo et al.\ 2006). We did not detect any dips in the light curve of GS 1354$-$645, possibly indicating that the inclination of the outer disk must really be lower than the value derived for the inner disk via reflection. However, dips may only be manifested at higher Eddington fractions (see, e.g., Kuulkers et al.\ 2000) than the luminosities inferred in our {\it NuSTAR} exposures. Esin et al.\ (1997) predict that the inner accretion flow will become advection-dominated and radiatively inefficient at Eddington fractions below 0.01. In the 0.01--0.08~$L_{Edd}$ range, the inner disk may still be truncated but the inner flow can be more luminous. For GS 1354-645, assuming a distance at the lower limit of $d=25$ kpc and mass at the lower limit of $M = 7.6~M_{\odot}$, the luminosities based on the flux values in Table 1 would be $0.1~L_{Edd}$ and $0.53L_{Edd}$ for Obs 1 and Obs 2 respectively. Smaller distances ($d<10$ kpc) would more easily accommodate the lower end of the luminosity range at which thin disks may truncate. Alternatively, some combinations including very high black hole masses ($M > 90~M_{\odot}$) can also meet the prediction, but these prescriptions greatly exceed the range of black hole masses inferred in X-ray binaries (e.g. Remillard \& McClintock 2006). We thank the anonymous referee for comments that improved this manuscript. This work was supported under NASA contract No.\ NNG08FD60C, and made use of data from the {\it NuSTAR} mission, a project led by the California Institute of Technology, managed by the Jet Propulsion Laboratory, and funded by NASA. \clearpage

1607.00343

1607

1607.05799_arXiv.txt

Integral field unit spectrographs allow the 2D exploration of the kinematics and stellar populations of galaxies, although they are generally restricted to small fields-of-view. Using the large field-of-view of the DEIMOS multislit spectrograph on Keck and our Stellar Kinematics using Multiple Slits (SKiMS) technique, we are able to extract sky-subtracted stellar light spectra to large galactocentric radii. Here we present a new DEIMOS mask design named \textit{SuperSKiMS} that explores large spatial scales without sacrificing high spatial sampling. We simulate a set of observations with such a mask design on the nearby galaxy NGC~1023, showing that the kinematic and metallicity measurements can reach radii where the galaxy surface brightness is several orders of magnitude fainter than the sky. Such a technique is also able to reproduce the kinematic and metallicity 2D distributions obtained from literature integral field spectroscopy in the innermost galaxy regions. In particular, we use the simulated NGC~1023 kinematics to model its total mass distribution to large radii, obtaining comparable results with those from published integral field unit observation. Finally, from new spectra of NGC~1023 we obtain stellar 2D kinematics and metallicity distributions that show good agreement with integral field spectroscopy results in the overlapping regions. In particular, we do not find a significant offset between our SKiMS and the \atlas\ stellar velocity dispersion at the same spatial locations.

\label{sec:introduction} In recent years the two-dimensional (2D) distribution of galaxy properties has provided a remarkably high number of useful constraints to understand galaxy formation and evolution processes. For example, the stellar and gas 2D line-of-sight kinematics are strongly linked with galaxy intrinsic shape, internal orbital structure and radial mass-to-light ratio ($M/L$) profile (e.g. \citealt{vanderMarel93, Gerhard93, Cretton00,Cappellari13b}). Furthermore, from absorption line index and/or full spectral fitting analyses it is possible to extract the luminosity- and mass-weighted parameters of the integrated stellar population (e.g. age, total metallicity $[Z/\rm{H}]$, $\alpha$-element abundance $[\alpha/\rm{Fe}]$, stellar initial mass function), as well as single chemical element abundances (e.g. \citealt{Kuntschner10, McDermid15}). Long-slit spectroscopy, along different position angles in a galaxy, has been used in the past to access this 2D information (e.g. \citealt{Davies88,Statler99, Saglia10}). Even though this approach provides information out to large radii, it is incapable of properly mapping the galaxy internal structure and requires a large amount of telescope time \citep{Statler94, Arnold94}. Furthermore, the spectra are obtained at different times, suffering from systematic effects in the case of imperfect sky subtraction. A way to overcome such issues is the use of Integral Field Unit (IFU) spectrographs. These instruments are able to obtain full spectral coverage of a 2D field-of-view (FoV) with a single exposure. Because of their efficiency, IFU spectrographs have been extensively used in surveys to explore the properties of large numbers of galaxies in the nearby Universe, although without being able to explore out to more than a few effective radii ($\rm{R_{e}}$). For instance, the \atlas\ survey used the \sauron\ IFU \citep{Bacon01} to explore the kinematics and the stellar populations of 260 local early-type galaxies (ETGs) in their inner regions (i.e. probing $R<1~\rm{R_{e}}$ in most cases, \citealt{Cappellari11a}). Similarly, the Sydney-AAO Multiobject Integral-field spectrograph (SAMI) is used in the SAMI survey to observe $3400$ galaxies, reaching not much beyond $2~\rm{R_{e}}$ \citep{Bryant15}. In some cases, IFU spectrographs have been used to explore the galaxy chemodynamics beyond $1~\rm{R_{e}}$. For example, using the SAURON spectrograph, \citet{Weijmans09} measured both kinematics and absorption line-strengths in NGC~821 and NGC~3379 out to almost $4~\rm{R_{e}}$. With a larger sample, the Visible Integral-field Replicable Unit Spectrograph prototype (VIRUS-P) has been used to study the kinematics \citep{Raskutti14} and stellar populations parameters \citep{Greene13} of ETGs out to $R\approx2.5~\rm{R_{e}}$, although with poor velocity resolution (i.e. $\sigma\approx150~\rm{km~s^{-1}}$) and low S/N ratio at large radii. The same instrument is also used for the MASSIVE survey, which targets the most massive ETGs (i.e. $M_\star \geq10^{11.5}~\rm{M_{\odot}}$, \citealt{Ma14}). Furthermore, the Calar Alto Legacy Integral Field Area Survey (CALIFA) takes advantage of the PMAS/PPAK instrument and aims to observe 600 galaxies in the local Universe out to generally $2~\rm{R_{e}}$ \citep{Sanchez12}. Finally, an even larger sample of galaxies is that explored by the Mapping Nearby Galaxies at APO (MaNGA, i.e. $\approx10000$ galaxies), but with a field-of-view still limited to the inner $1.5-2.5~\rm{R_{e}}$ \citep{Bundy15}. As shown by \citet{Brodie14}, in a typical ETG only a small fraction of the total galaxy mass and angular momentum is included in the inner few $\rm{R_{e}}$. Furthermore, while these inner regions are dominated in mass by the stellar component, at large radii the dominant component is the dark matter. In order to sample a sufficiently large fraction of total galaxy mass and angular momentum, good 2D spatial coverage is needed in both these regions. In particular, it is important to consistently measure the stellar kinematics and population parameters at both small and large radii. Applying the Stellar Kinematics from Multiple Slits (SKiMS) method \citep{Norris08, Proctor09, Foster09}, it is possible to measure both the kinematics and the metallicity of nearby galaxy stellar components out to large radii (e.g. \citealt{Foster13, Arnold14, Pastorello14}). In particular, \citet{Foster15} obtained SKiMS kinematic measurements out to $5~\rm{R_{e}}$ ($2.6~\rm{R_{e}}$ on average) and \citet{Pastorello14} metallicity measurements out to about $3~\rm{R_{e}}$. The original SKiMS method was developed to extract the background galaxy stellar light spectra from the same DEIMOS slits that were primarily targeting globular cluster (GC) candidates. In this way, the galaxy stellar component is probed at randomly scattered spatial positions. Furthermore, the bright inner galaxy regions are not targeted, since GCs are difficult to detect in the presence of a strong stellar background. In order to include the innermost regions in their dynamical models of 14 ETGs, \citet{Cappellari15} combined the SKiMS large-radii stellar kinematics with those from \atlas\ at small radii. The use of two different datasets required the accounting for a number of systematic issues given by the non-homogeneity of the data (e.g. different spatial sampling and kinematic uncertainties). A way to overcome such issues is to use an homogeneous, although optimal azimuthally and spatially distributed, kinematic dataset. Here we present a new mask design that takes advantage of the Keck/DEIMOS multislit spectrograph to reliably explore the stellar kinematics and stellar population parameters in ETGs out to large radii, with complete azimuthal coverage. We name this \textit{SuperSKiMS}, since it is an application of the SKiMS method on data obtained using specially designed multislit masks. In particular, the Keck/DEIMOS slit distribution in the \textit{SuperSKiMS} configuration allows for the optimal sampling of both the inner and the outer regions of nearby galaxies. Tho demonstrate this, we present the results from mock \textit{SuperSKiMS} simulations in order to show that such a technique can return data comparable to that from IFU spectroscopy. In particular, we use such simulated observations to extract the total mass profile slope of the nearby lenticular galaxy NGC~1023 from the modelling of the galaxy kinematics. These results are compared with those obtained from \atlas\ kinematic measurements in the centre and with the results by \citet{Cappellari15} (i.e. using a similar but smaller sample, with a shorter radial baseline). Moreover, we use two SuperSKiMS masks to obtain new stellar kinematics and metallicity measurements for NGC~1023. These values are used together with already published SKiMS measurements, thus extending the study of stellar kinematics and metallicity out to $3.6$ and $2.5~\rm{R_{e}}$, respectively. At such galactocentric radii the stellar light from the galaxy is just 1.6\% and 5.5\% of the sky flux at similar wavelengths. At the same time, our new data extend into the central regions of NGC~1023, overlapping with several published longslit and IFU observations. We then evaluate how the addition of these new datapoints to the dataset affects several results already presented in the SLUGGS survey. In particular, from the new stellar velocity and velocity dispersion 2D maps we measure the radial specific angular momentum profile and compare it with the results from \citet{Foster15}. From the stellar metallicity 2D distribution we extract the azimuthally averaged radial metallicity profiles, from which we obtain new measurements for the inner (i.e. $R<1~\rm{R_{e}}$) and outer (i.e. $R\geq1~\rm{R_{e}}$) metallicity gradients and compare them with the results presented in \citet{Pastorello14} and \citet{Forbes15}. In this work we assume the following NGC~1023 parameters: effective radius $R_{e}=48~\rm{arcsec}$, axial ratio $b/a=0.37$, position angle $PA=83.3~\rm{degrees}$ and distance $d=11.1~\rm{Mpc}$ \citep{Brodie14}. The NGC~1023 photometric decomposition by \citet{Savorgnan15} shows that a faint but spatially extended bar is present in the galaxy inner regions. Such a bar extends out to $R_{\rm{bar}}\approx40~\rm{arcsec}$ and has a width of $w_{\rm{bar}}\approx20~\rm{arcsec}$, oriented $\Delta PA = -22~\rm{degrees}$ from the galaxy major axis $PA$. The structure of the paper is as follows. In Section \ref{maskDesign} we describe the new \textit{SuperSKiMS} mask design and observation method. In Section \ref{sec:simulations} we then present how, under typical observing conditions, the data from the \textit{SuperSKiMS} technique would compare with \atlas\ data in the inner regions of NGC~1023. In Section \ref{sec:results} we model the galaxy mass density distribution with both the original SKiMS, the \atlas\ and the \textit{SuperSKiMS} mock datasets. Section \ref{sec:data} describes the observation and data reduction of two \textit{SuperSKiMS} masks on NGC~1023. In Section \ref{dataanalysis} we present the extraction of stellar kinematics and metallicity 2D maps and radial profiles, which we compare with available literature measurements in Section \ref{sec:analysis}. In the same Section we compare the results with those from \citet{Foster15} and \citet{Pastorello14}, in order to show the improvement given by the addition of the new data. Finally, in Section \ref{sec:conclusions} we present our conclusions.

\label{sec:conclusions} We present for the first time a new multislit mask design, called \textit{SuperSKiMS}, that maximizes the azimuthal and radial coverage of a galaxy's 2D stellar field. This slit mask design can be optimised to the best azimuthal coverage given the expected number of mask pointings. In addition, the slit lengths are designed in order to provide a high S/N ratio at most radii. The outermost regions of the mask are used to obtain pure-sky spectra necessary for sky subtraction and to target globular cluster candidates or background galaxies. We tested the technique by simulating the exposure of four \textit{SuperSKiMS} masks in the central regions of NGC~1023. For this galaxy, \atlas\ 2D kinematic and metallicity maps are available. We extracted the velocity, velocity dispersion and metallicity measurements from these \atlas\ maps, at the spatial locations where the \textit{SuperSKiMS} central slits would be. We obtained 2D maps which we compared with those from \atlas\ in the inner regions. We found that the observation of 4 \textit{SuperSKiMS} DEIMOS masks would allow us to recover the stellar kinematic and metallicity 2D distribution observed by the \atlas\ work with high accuracy. Moreover, we compared a partially simulated SLUGGS+\textit{SuperSKIMS} dataset with the composite SLUGGS+\atlas\ dataset used by \citet{Cappellari15} to model the total mass density profile of NGC~1023. In particular, from both datasets we retrieved a nearly isothermal total mass density distribution in the probed radial range, consistent with that published by \citet{Cappellari15} for this galaxy. As a following step, we tested the \textit{SuperSKiMS} technique using the large field-of-view DEIMOS multislit spectrograph on the nearby galaxy NGC~1023. Unfortunately, issues with the Keck alignment software prevented us from fully exploiting the capabilities of the \textit{SuperSKiMS} mask design. In any case, we observed two masks with the \textit{SuperSKiMS} design and added the new data to the pre-existing SLUGGS dataset. With such data we then produced 2D stellar velocity and velocity dispersion maps. Such maps allowed the extraction of virtual slit radial profiles, which we compared with longslit literature results. Despite random sampling of the field, we found a good consistency with the literature velocity profiles. In addition, comparing our kinematic radial profiles with those extracted from \atlas\ 2D maps, we found fair consistency in the overlapping regions. Since many of the new datapoints lie in the \atlas\ field-of-view, we also compared the differences in kinematics at the same spatial locations with \atlas\ data. We found a good agreement between the velocity and a fair agreement between the velocity dispersion values (in particular, we did not observe the velocity dispersion offset noted in \citealt{Arnold14} and \citealt{Foster15}). These results suggest that the previously observed offset in NGC~1023 might be driven by the small number of datapoints overlapping with the \atlas\ field-of-view, most of which also lie at the edges of this region, where the \atlas\ values have been obtained from the binning of low S/N data. However, other galaxies in \citet{Arnold14} and \citet{Foster15} show similar offsets even in the presence of a large number of datapoints overlapping with the \atlas\ field-of-view. From these kinematic measurements we also extracted the local specific angular momentum radial profile of NGC~1023, extending the results by \citet{Foster15} with our updated dataset. In addition to the kinematics, it is possible to reliably measure the equivalent width of the CaT absorption lines and hence produce 2D metallicity maps. From these maps we extracted the azimuthally averaged radial metallicity profile out to $R\approx~2\rm{R_{e}}$. This profile is consistent with the stellar metallicity measurements by \atlas\ in the overlapping radial regions, while extending to almost twice their probed galactocentric radii. In addition, we extracted the inner and outer stellar metallicity gradients for NGC~1023, updating the results already presented in \citet{Forbes16} with our new dataset. In future studies, the adoption of the \textit{SuperSKiMS} mask design to observe nearby galaxies with multislit spectrographs will allow one to probe the stellar component homogeneously from the innermost to the outermost regions. From this it will be possible to efficiently retrieve reliable kinematics and stellar population 2D maps comparable to those from integral field unit spectrograph studies, although reaching much larger radii.

1607.05799

1607

1607.03113_arXiv.txt

Cosmic acceleration may be due to modifications of cosmic gravity and to test this we need robust connections between theory and observations. However, in a model independent approach like effective field theory or a broad class like Horndeski gravity, several free functions of time enter the theory. We show that simple parametrizations of these functions are unlikely to be successful; in particular the approximation $\alpha_i(t)\propto\Omega_{\rm de}(t)$ drastically misestimates the observables. This holds even in simple modified gravity theories like $f(R)$. Indeed, oversimplified approximations to the property functions $\alpha_i(t)$ can even miss the signature of modified gravity. We also consider the question of consistency relations and the role of tensor (gravitational wave) perturbations.

The origin of cosmic acceleration is an extraordinary mystery in modern physics. The observation of cosmic acceleration \cite{perlmutter,riess,DEreview,DEreview2} must be connected to some fundamental theory beyond the current standard model of particle physics, but we do not know whether its origin lies in the structure of the quantum vacuum or an extension to Einstein's theory of gravitation. Considerable progress has occurred in the last decade in exploring aspects of modified gravity \cite{troddenreview,koyamareview,fabianreview} but the ability to connect theory and observations in a manner not highly dependent on a specific model is lacking in essential aspects. Here we examine the challenges for such a connection, and caution against oversimplification. Modified gravity is a much more complex arena than scalar field dark energy, with its one free function of time (e.g.\ the equation of state $w(a)$). In large part this is because of the role played by perturbations and the tensor sector. To begin, consider the case for cosmic acceleration not arising from modified gravity, e.g.\ quintessence dark energy. Here we also have challenges in connecting essential theory to observations, with perhaps the most information arising from the thawing vs.\ freezing classification of scalar fields \cite{caldlin}. This at least describes the steepness of the potential relative to the Hubble friction, and has distinct implications for whether the theory is approaching or departing from a cosmological constant-like (or sometimes de Sitter) state. Beyond that, the expansion history from dark energy, whether from quintessence or modified gravity, is extremely accurately characterized phenomenologically by two numbers \cite{depl}, $w_0$ and $w_a$, measures of the present and time variation of the dark energy equation of state. Indeed, this characterization has been shown valid to the 0.1\% level in the observables of distances and Hubble parameters for a wide range of quintessence, k-essence, modified gravity, etc.\ models. On the cosmic structure, i.e.\ perturbative, side of observations, dark energy not arising from modified gravity (or nonstandard couplings) has little to add: quintessence perturbations are small inside the Hubble scale and k-essence (noncanonical kinetic energy model) perturbations have little observational effect since they are suppressed by equations of state near $w=-1$, as observations indicate. For modified gravity effects, a successful, if limited, phenomenological parametrization is the gravitational growth index $\gamma$ \cite{lingam}, again accurately describing observables at the subpercent level for a variety of modified gravity models \cite{lincahn}. However, this has very restricted interpretable connection to fundamental theory. Better (pseudo)observables include effects not only on growth of structure, but on the deflection of light. These come from the nonrelativistic and relativistic modified Poisson equations \cite{bert11}, and can be written as effective gravitational coupling strengths $\gm(k,a)$ and $\gl(k,a)$, where we explicitly show their scale dependence (e.g.\ on the Fourier wavenumber $k$) and time dependence (e.g.\ on the cosmic scale factor $a$). The gravitational growth index $\gamma$ is directly related to $\gm$ in the scale independent limit. To connect with theory, however, we need to relate the scale and time dependences of these ``observables'' (we will henceforth refer to these quantities $\gm$, $\gl$, and their ratio, related to the gravitational slip function, as observables because, while not directly observable, they are so closely connected to observations, i.e.\ structure growth and gravitational lensing) to the theory -- or at least to phenomenological property functions $\alpha_i$ \cite{bellsaw}. The functional form of the scale dependence is a ratio of $k^0+k^2$ polynomials in many cases (see \cite{bert08,11084242,11091846} and the especially clear \cite{13021193}, but see \cite{12084163,14098284} for exceptions), and one simplification is that on scales below the Hubble scale (or more generally the sound horizon or braiding scale \cite{bellsaw}) the scale dependence ($k^2$ terms) is subdominant and one essentially has purely functions of time. While this seems to be considerable progress, the problem is that in order to know whether $\gm$ and $\gl$ have any simple parametrization of their time dependence one has to evaluate them from the underlying theory, ideally in as model independent a fashion as possible. An excellent framework for this is the effective field theory (EFT) of dark energy \cite{gubitosi,bloomfield,gleyzes,lsw}. Within EFT at quadratic (lowest) order there are seven free functions of time, and within the Horndeski class of gravity there are four free functions. The challenge of connecting such theory functions to realistically parametrized observables was highlighted in \cite{lsw}, who examined various limits. Here we go into greater depth and quantify the problems with oversimplification of the parametrization. In Sec.~\ref{sec:fR} we start with the workhorse of modified gravity, $f(R)$ theory. This corresponds to only one independent free function of time in the EFT formalism and so is a basic place to start in assessing parametrizations. We expand in Sec.~\ref{sec:early} to the four functions of Horndeski gravity and examine the motivation for a potential simplification in the early time (matter dominated) limit, deriving the asymptotic behavior of the property functions and observables. In Sec.~\ref{sec:propto} we identify how extending these limiting behaviors to the epoch of cosmological structure observations raises foundational issues, and we quantify the dramatic deviations that actually arise in generic circumstances. Section~\ref{sec:discuss} discusses the reasons why simplified parametrizations appear generally unviable, and solving a problem like modified gravity is so difficult. We conclude in Sec.~\ref{sec:concl} with some thoughts on further progress, while Appendix~\ref{sec:consist} explores the possibility of proving broad consistency relations that an entire class of modified gravity theories must obey.

\label{sec:concl} Modified gravity leading to cosmic acceleration is a much richer field than envisioned even a few years ago. The early models like DGP gravity with a single number (the crossover scale) or $f(R)$ gravity with a time dependent scalaron mass as described by a single power law index of scale factor have much less freedom compared to even the quite restricted covariant Galileon theories with constant coefficients, let alone the Horndeski class or EFT with their several free functions of time. This complexity, in both the theory and its connections to observables, means that accurate approximations to the observables -- being ``ratios of sums of products of ratios of sums of functions'' -- are rare. We derive analytic limits in the early time, matter dominated regime for general classes of Horndeski gravity, and show under what conditions they appear. These early time approximations, however, break down dramatically even at redshifts $z\approx10$, let alone in the heart of the observable epoch. Even percent level deviations in the property functions $\al_i(a)$ can lead to large misestimations in observable properties. In particular, we demonstrate that taking them proportional to the effective dark energy density, $\al_i(a)/\ode(a)\propto$ constant can lead to unphysical behavior and fine tuning and can miss significant signatures of departure from general relativity. This last property is perhaps most damaging: misestimation could just give a false alert, but lack of an alert will miss essential physics \cite{xkcd}. To meet the challenge of connecting theory and observations, we need some parametrization that can prove itself accurate on at least broad swathes of theories in the literature. The numerical solutions we have shown for $f(R)$ and covariant Galileon gravity, demonstrating the complexity of the evolution, indicate this may be a difficult task. In a real sense this is no surprise: the hallmark of modified gravity is that the physics of growth does not simply follow the expansion history, e.g.\ $\ode(a)$. If a nearer term goal is merely an alert that general relativity may not be matching observations, then bins in scale and time of $\gm$ and $\gl$, proposed in \cite{10021962,10080397}, work well. Moreover, they would give some indication of how the breakdown occurs, i.e.\ the trend in space and time variation. While the lack of an elegant parametrization such as exists for the background expansion (e.g.\ dark energy equation of state) or even simple linear growth (e.g.\ the gravitational growth index) is disappointing, it also points up the richness of the problem of modified gravity. In Appendix~\ref{sec:consist} we comment on a conjecture for a general consistency relation between observables that could apply to wide classes of modified gravity theories. At the same time, we should seek new gravitation theories that are neither overly simplified and so lacking model independence nor complicated but observationally unviable.

1607.03113

1607

1607.01785_arXiv.txt

Early quiescent galaxies at $z\sim2$ are known to be remarkably compact compared to their nearby counterparts. Possible progenitors of these systems include galaxies that are structurally similar, but are still rapidly forming stars. Here, we present Karl G. Jansky Very Large Array (VLA) observations of the CO(1--0) line towards three such compact, star-forming galaxies at $z\sim2.3$, significantly detecting one. The VLA observations indicate baryonic gas fractions $\gtrsim$5 times lower and gas depletion times $\gtrsim$10 times shorter than normal, extended massive star-forming galaxies at these redshifts. At their current star formation rates, all three objects will deplete their gas reservoirs within 100\,Myr. These objects are among the most gas-poor objects observed at $z>2$, and are outliers from standard gas scaling relations, a result which remains true regardless of assumptions about the CO-H$_2$ conversion factor. Our observations are consistent with the idea that compact, star-forming galaxies are in a rapid state of transition to quiescence in tandem with the build-up of the $z\sim2$ quenched population. In the detected compact galaxy, we see no evidence of rotation or that the CO-emitting gas is spatially extended relative to the stellar light. This casts doubt on recent suggestions that the gas in these compact galaxies is rotating and significantly extended compared to the stars. Instead, we suggest that, at least for this object, the gas is centrally concentrated, and only traces a small fraction of the total galaxy dynamical mass.

\label{intro} Galaxies with large stellar masses ($\Mstar \gtrsim 10^{11}$\,\msol) and little ongoing star formation have been observed at redshifts up to $z\sim4$ \citep{straatman14}, and begin to appear in large numbers by $z\sim2.5$ \citep[e.g.,][]{kriek06,whitaker10,cassata13}. Many studies have shown that these early quiescent galaxies were much smaller at $z\gtrsim1.5$ than equally massive star-forming galaxies at similar redshifts \citep[e.g.,][]{trujillo07,vanderwel14} or quiescent galaxies of similar mass in the local universe \citep[e.g.,][]{daddi05,trujillo06,buitrago08,cimatti08,vandokkum08,damjanov11}. With effective radii of just $\sim1-3$\,kpc, the stellar densities are of order 100$\times$ higher than present-day elliptical galaxies. Similarly massive and compact galaxies are extremely rare in the local universe \citep[e.g.,][]{trujillo09,taylor10}, implying significant size growth largely consistent with the effects of minor merging (e.g., \citealt{trujillo11,newman12}; see also \citealt{carollo13}). Although not all local massive galaxies had a compact progenitor \citep[e.g.,][]{franx08,vandokkum08}, most of the $z\sim2$ compact, massive galaxies likely now reside in the centers of present-day elliptical galaxies \citep[e.g.,][]{bezanson09,belli14a}. The formation mechanism(s) of the $z\sim2$ compact quiescent population is still unclear. Recently, however, a population of similarly compact yet highly star-forming galaxies at $z\sim2.5$ has been identified in deep \textit{Hubble Space Telescope} imaging \citep{barro13,barro14a,nelson14,williams14,vandokkum15}. Given their structural similarity, these compact star-forming galaxies (SFGs) are natural candidate progenitors of the early quiescent population, requiring only the cessation of star formation to superficially match the stellar distribution and structure of the $z\sim2$ quiescent galaxies. Additional evidence from dynamical studies \citep{barro14,vandokkum15} and number density evolution \citep{barro13} also indicates that compact SFGs will plausibly quench star formation on short timescales ($\lesssim$500\,Myr) to build up the growing quiescent population. The small sizes and non-exponential light profiles of compact SFGs present a clear contrast to the typical massive star-forming galaxies at $z\sim2$, which consist mostly of gas-rich, rapidly rotating disks \citep[e.g.,][]{wuyts11,tacconi13,wisnioski15}, evidence which suggests very different evolutionary histories. The aforementioned dynamical studies indicate that the stellar masses of compact SFGs are ubiquitously comparable to or somewhat in excess of simple estimates of the total dynamical masses, \Mdyn, determined through observations of \halpha. This implies that the dynamics of compact SFGs are almost completely dominated by the stars -- any gas present likely traces, but does not significantly contribute to, the gravitational potential. \citet{vandokkum15} argue that the \halpha-emitting gas is likely rotating and more extended than the stellar light, preventing the unphysical scenario of $\Mstar>\Mdyn$. This inference was motivated by observations of velocity gradients across the slit consistent with rotation for a few objects. Although selection effects are likely in play, these observations imply that the gas is extended by a factor $\sim$2.5 relative to the stars on average. Adaptive optics-assisted integral-field or interferometric synthesis imaging can be used to spatially and spectrally resolve the gas, providing a more robust tracer of the dynamics than can be inferred from long-slit spectroscopy. Recent cosmological simulations have matched the observed number counts of compact SFGs and quiescent galaxies \citep{wellons15}. This work suggests that the early quiescent population consists mostly of a combination of galaxies which formed their stellar mass early and remained compact since their formation time, and objects which have recently undergone gas-rich major mergers. High-resolution hydrodynamical simulations further indicate that many characteristics of compact massive galaxies can be reproduced through in-situ formation through gas accretion and cooling \citep[e.g.,][]{naab09,feldmann15}, strong central star formation during a gas-rich major merger \citep[e.g.,][]{wuyts10,ceverino15}, and/or dissipative contraction of gas-rich disks \citep[e.g.,][]{ceverino15,zolotov15}. In each case, the properties of the gas are key, as its collisional nature allows energy to dissipate and angular momentum to be transferred through the galaxy, permitting the formation of characteristically dense stellar structures. In this work, we present Karl G. Jansky Very Large Array (VLA) observations of the CO(1--0) line toward three compact SFGs at $z\sim2.3$. CO(1--0) has long been known as a tracer of molecular hydrogen in the interstellar medium (ISM), the direct fuel from which stars form. If, as we have outlined above, the compact SFGs are in the process of quenching star formation to become $z\sim2$ quiescent galaxies, we may expect molecular gas properties unlike those of normal SFGs at these redshifts. In particular, if compact SFGs are to quench in tandem with the rapid buildup of the quiescent population, we expect short gas depletion or quenching timescales, indicating that the ongoing high star formation rates (SFRs) cannot be sustained for more than a few hundred Myr. The outline of this paper is as follows. In Section~\ref{obs}, we describe our sample selection and VLA observations. Section~\ref{results} describes our results, including constraints on the molecular gas masses of our targeted objects (Section~\ref{mgas}), gas fractions and depletion timescales (Section~\ref{fgastdep}), and the physical extent of the molecular gas reservoirs in comparison to the stellar light (Section~\ref{dynamics}). We summarize our conclusions in Section~\ref{conclusions}. Throughout, we assume a flat $\Lambda$CDM cosmology, with $H_0=67.7$\,\kms\,Mpc$^{-1}$ and $\Omega_\mathrm{m}=0.307$ \citep{planck15}.

1607.01785

1607

1607.02053_arXiv.txt

Blazars can be divided into two sub-classes namely high energy and low energy peaked blazars. In spectral energy distribution, the first synchrotron hump of the former class peaks in UV/X-rays and in IR/optical bands for the latter class. The peak of the spectral energy distribution seems to be responsible for variability properties of these classes of blazars in X-ray and optical bands. Since, in low energy peaked blazars, the X-ray bands lies well below the synchrotron hump, one expects that the highest energy electrons available for the synchrotron emission would have slower effect of variability on X-ray intra-day timescale. In this paper, by taking the advantage of a sample of 12 low energy peaked blazars with total 50 observations from XMM$-$Newton since its launch, we confirm that this class is less variable in X-ray bands. We found that out of 50 observational light curves, genuine intra-day variability is present in only two of light curves i.e 4\%. Similar results we obtained from our earlier optical intra-day variability studies of high energy peaked blazars where out of 144 light curves, only genuine intra-day variability was detected in 6 light curves i.e $\sim$ 4\%. Since, X-ray bands lie below the peak of the spectral energy distribution of LSPs where inverse Compton mechanism is dominating rather than synchrotron radiation at the peak of the optical band, leads to slower variability in the X-ray bands. Hence, reducing their intra-day variability in X-ray bands as compared to the variability in optical bands.

A small sub-class of radio-loud active galactic nuclei (AGN) is known as blazars. BL Lac objects and flat spectrum radio quasars (FSRQs) collectively known as blazars. Optical spectrum of BL Lac objects are featureless i.e. absence of prominent emission or absorption lines, whereas FSRQs show prominent emission lines. The common properties of blazars include large amplitude violent flux variation in complete electromagnetic (EM) spectrum, high and variable polarization from radio to optical bands, core-dominated radio morphology, and emission being predominantly nonthermal. Blazars emit relativistic charged particle jets which pointed close (at angles $\leq$ 10$^{0}$) to our line of sight (e.g., Urry \& Padovani 1995) and it causes the observed emission to be relativistically beamed. Since blazars emit radiation in the complete EM spectrum, these are among ideal objects to study their multi-wavelength spectral energy distribution (SED). The SED of blazars show two well defined broad spectral components (Mukherjee et al. 1997). Based on the location of these SED peaks, blazars are further classified into low energy peaked blazars (LBLs) and high energy peaked blazars (HBLs). In LBLs the first SED component peaks in radio to optical while the second component peaks at GeV energies, and in HBLs the first component peaks in UV/X-rays while the second component peaks at TeV energies (Padovani \& Giommi 1995). Recently blazars classification is made on synchrotron peak frequency and divided into three sub-classes i.e low synchrotron peaked (LSPs), intermediate synchrotron peaked (ISPs) and high energy peaked blazars (HSPs) (Abdo et al. 2010a). Basically LSPs and ISPs collectively belong to LBLs class. \begin{table*} {\bf Table 1.} Observation log of XMM-Newton X-ray data for low energy peaked blazars$^{*}$ \small \begin{tabular}{lccclcclr} \hline Blazar Name & $\alpha_{2000.0}$& $\delta_{2000.0}$ & redshift &Blazar & Date of Obs. & Obs. ID & Window & GTI$^{b}$(s) \\ & & & $z$ &Class & yyyy.mm.dd & & Mode$^{a}$ & \\\hline TXS 0106$+$612 & 01h09m46.3s & $+$61$^{0}$33$^{'}$30$^{''}$ & 0.783 & ~LSP & 2010.02.09 & 0652410201 & ~~FF & 12034 \\ PKS 0235$+$164 & 02h38m38.9s & $+$16$^{0}$36$^{'}$59$^{''}$ & 0.94 & ~LSP & 2002.02.10 & 0110990101 & ~~FF & 16847 \\ & & & & & 2004.01.18 & 0206740101 & ~~SW & 29671 \\ & & & & & 2004.08.02 & 0206740501 & ~~SW & 11471 \\ & & & & & 2005.01.28 & 0206740701 & ~~SW & 16270 \\ PKS 0426$-$380 & 04h28m40.4s & $-$37$^{0}$56$^{'}$20$^{''}$ & 1.11 & ~LSP & 2012.02.11 & 0674330201 & ~~EFF & 20745 \\ PKS 0528$+$134 & 05h30m56.4s & $+$13$^{0}$31$^{'}$55$^{''}$ & 2.06 & ~LSP & 2009.09.11 & 0600121601 & ~~FF & 25672 \\ PKS 0537$-$286 & 05h39m54.3s & $-$28$^{0}$39$^{'}$56$^{''}$ & 3.104 & ~LSP & 2000.03.19 & 0114090101 & ~~FF & 18852 \\ & & & & & 2005.03.20 & 0206350101 & ~~FF & 80138 \\ S5 0716$+$714 & 07h21m53.4s & $+$71$^{0}$20$^{'}$36$^{''}$ & 0.31 & ~ISP & 2007.09.24 & 0502271401 & ~~SW & 71624 \\ 4C 71.07 & 08h41m24.3s & $+$70$^{0}$53$^{'}$42$^{''}$ & 2.172 & ~LSP & 2001.04.12 & 0112620101 & ~~FF & 33480 \\ OJ 287 & 08h54m48.9s & $+$20$^{0}$06$^{'}$31$^{''}$ & 0.3056 & ~ISP & 2005.04.12 & 0300480201 & ~~LW & 13192 \\ & & & & & 2005.11.03 & 0300480301 & ~~LW & 39074 \\ & & & & & 2006.11.17 & 0401060201 & ~~LW & 44972 \\ & & & & & 2008.04.22 & 0502630201 & ~~LW & 53568 \\ & & & & & 2011.10.15 & 0679380701 & ~~LW & 21669 \\ 3C 279 & 12h56m11.1s & $-$05$^{0}$47$^{'}$22$^{''}$ & 0.5362 & ~LSP & 2009.01.21 & 0556030101 & ~~FF & 25235 \\ & & & & & 2011.01.18 & 0651610101 & ~~SW & 125487 \\ BL Lac & 22h02m43.3s & $+$42$^{0}$16$^{'}$40$^{''}$ & 0.0686 & ~ISP & 2007.07.10 & 0501660201 & ~~SW & 18478 \\ & & & & & 2007.12.05 & 0501660301 & ~~SW & 19272 \\ & & & & & 2008.01.08 & 0501660401 & ~~SW & 22371 \\ 3C 454.3 & 22h53m57.7s & $+$16$^{0}$08$^{'}$54$^{''}$ & 0.859 & ~LSP & 2006.07.02 & 0401700201 & ~~SW & 15972 \\ & & & & & 2007.05.23 & 0401700401 & ~~SW & 2995 \\ & & & & & 2006.12.18 & 0401700501 & ~~SW & 15021 \\ & & & & & 2007.05.31 & 0401700601 & ~~SW & 28971 \\\hline \end{tabular} \\ $^{*}$ Observation log for the LSP 3C 273 is given in Table 1 of our paper Kalita et al. (2015). \\ $^{a}$ Extended Full Frame = EFF, Full Frame = FF, Large Window = LW, Small Window = SW \\ $^{b}$ GTI = Good Time Interval \\ \end{table*} \begin{figure*} \includegraphics[width=185mm]{f1.ps} \caption{X$-$ray light curves of the low energy peaked blazars. In each panel blazar name / observation ID is given. Light curves are generated with 100s binning. For light curves of 3C 273 see Fig. 1 of Kalita et al. (2015).} \label{f1} \end{figure*} Variation in the blazar flux of the order of a few hundredth to tenths over a time scale of few minutes to less than a day is called intraday variability (IDV) (Wagner \& Witzel 1995). Variability timescales of weeks to few months is commonly known as short term variability (STV) and timescales of months to years is called long term variability (LTV) (Gupta et al. 2004). The most puzzling flux variation in blazars are those which are happening on IDV timescales. The variability mechanisms in the blazars on IDV timescales are widely accepted that it is associated with the instabilities and irregularities in the jet flow. Hence, in order to understand IDV, it is very important to understand the fine details of jet formation. Astronomers seeks to image jet formation using very long baseline radio interferometry (VLBI) but it suffers severe lack of sufficient angular resolution. An alternative method to measure fine structure of jets i.e. to search for IDV with the fastest possible time sampled light curves (LCs). The first genuine optical IDV in blazar was reported by Miller et al. (1989). Since then extensive search for optical IDV in large number of blazars were done (e.g. Heidt \& Wagner 1996; Montagni et al. 2006; Gupta et al. 2008, 2012; Agarwal et al. 2015; Gaur et al. 2015; and references therein). In a detailed statistical analysis of optical IDV of blazars, (Gupta \& Joshi 2005) reported that if a blazar continuously observed for less than 6 hours and more than 6 hours, the chances of detecting IDV is $\sim$ 60 -- 65\% and 80--85\%, respectively. Till 2005, most of the known blazars belong to LSPs and ISPs class and so the report by Gupta \& Joshi (2005) was true only for LSPs and ISPs. Till 2005, there were only 6 HBLs or HSPs (Mrk 421, Mrk 501, 1ES 1426+428, 1ES 1959+658, PKS 2155$-$304, 1ES 2344+514) were known, and their studies were mainly focused in X-ray and $\gamma-$ray bands (e.g. Edelson et al. 2001; Zhang et al. 2002; Giebels et al. 2002; Krawczynski et al. 2002; Aharonian et al. 2002, 2005; Konopelko et al. 2003; Massaro et al. 2004; Daniel et al. 2005; and references therein). Thanks to the revolution in $\gamma-$ray astronomy due to HESS, MAGIC, Fermi, VERITAS, etc. which lead to a very rapid increase in detection of new HSPs (e.g. Abdo et al. 2010b, Nolan et al. 2012, Acero et al. 2015). The new sample of HSPs gave us an opportunity to see the optical IDVs of HSPs and compare its properties with optical IDVs of LSPs. We started a dedicated project to search for optical IDV in HSPs and after doing 62 nights of IDV observations of HSPs which gave us 144 LCs (41 in B band, 62 in R band, and 41 in B$-$R color) of five HSPs (Mrk 421, 1ES 1426+428, 1ES 1553+113, 1ES 1959+650, and 1ES 2344+514). Interestingly, we found that, 4 HSPs did not show any IDV (Gaur et al. 2012a, 2012b, 2012c), but only one HSP 1ES 1426+428 for which we have the least observations have shown IDV in 6 LCs out of 8 LCs (Gaur et al. 2012c). Our this pilot project gave us 6 IDV LCs out of 144 LCs searched for IDV i.e $\sim$ 4\% LCs have shown IDV. We explained it by density inhomogeneities and bends in the bases of the jets by Kelvin--Helmholtz instabilities (Romero et al. 1999). We gave an alternative explanation i.e. since in HSPs, the optical band lies below the SED peak, hence, we should see changes in the efficiency of acceleration of, and/or in the rates at which energy is radiated by, the highest energy electrons available for synchrotron emission would have a more retarded effect on optical variability in HSPs (Gaur et al. 2012b). In LSPs, the optical band is dominated by highest energy electrons emitting synchrotron radiation and probably the X-ray emission is dominated by the comparatively lower energy electrons emitting the inverse Compton radiation, hence their X-ray variability is less pronounced than optical variability. If SED peak is really responsible for IDV properties, then we suspected that X--ray IDV LCs in LSPs should not show any IDV at all or show on rare occasions. With this motivation, here we present the X-ray IDV study of almost complete sample of 10 LSPs and 2 ISPs observed by XMM-Newton since its launch and we found that the LSPs show very less IDV 2 out of 50 LC i.e. 4\% in X-ray bands. We have reported above the similar finding for HSPs in optical bands. The paper is structured as follows. In Section 2, we discuss XMM--Newton archival data and analysis. Section 3 reports our results and section 4 contains discussion.

We searched for X--ray IDV in the LCs of almost all the LSPs observed by XMM--Newton since its launch. The details of the observation log is given in Table 1 which gives the data of 9 LSPs and 2 ISPs with total 25 LCs. All the LCs are plotted in Fig. 1. There were 25 LCs of LSP 3C 273 which are already presented in Table 1 and Fig. 1 of Kalita et al. (2015). In Table 2, we report the results of variability parameters of all 50 LCs which are calculated using excess variance. It is clear from Table 2 that in only two pointings, \ one of the ISP S5 0716+714 and another one of LSP PKS 0235+164, we found significant IDV with variability amplitude of 23 and 11.6\%, respectively. Since, S5 0716+714 is observed in outburst state and is an ISP, it could be expected that the synchrotron peak of this blazar reaches up to soft X-ray regime. Similarly, PKS 0235+164 is also found to shift its peak up to soft X-ray region in its outburst state (Madejsi et al. 1996). Hence, in these two occasions, we expect to get the variability in X-ray band. In other pointings, we did not find any significant variability. Most of the sources listed in Table 1. are well known LSPs and are extensively monitored in optical bands and they have shown high duty cycles in optical bands up to $\sim$ 70--80\% (Gupta et al. 2008, 2012; Goyal et al. 2012; Agarwal et al. 2015; Gaur et al. 2015 and references therein). Similarly, IDV of HSPs in X-ray band is well studied and are highly variable in these bands (Lachowicz et al. 2009; Gaur et al. 2010; Kalita et al. 2015; and references therein) but not very extensively studied in optical bands. Till now, we put an extensive effort to observe the optical IDV of HSPs and found these sources to be less variable in these bands. The difference in the multi-frequency spectral properties of HSPs and LSPs requires a systematic change of intrinsic physical parameters such as magnetic field, jet size, maximum electron energy and it is investigated by Sambruna et al. (1996) that the change is in the sense that HSPs have higher magnetic fields/electron energies and smaller sizes as compared to LSPs. All the above factors lead to the difference in the cut off energies of LSPs and HSPs and hence have a more retarded effect on the X-ray variability of LSPs. ISPs lies in between these two classes and is difficult to mark the exact boundaries as it depends on the state of the source i.e whether it is in quiescent/outburst state. The variations in the acceleration efficiency of the relativistic electrons near the synchrotron hump could arise from the changes in the local number density of the most energetic electrons or the strength of the localized magnetic fields. Near the peak of the SED, acceleration processes dominates and produces the higher energy electrons while the lower energy electrons are available for the emission below the SED peaks. Below the SED peak, probably cooling processes dominate which involve mainly inverse Compton for LSPs in general (Joshi et al. 2014). Since, the most excepted model for the intra-day variability involves magneto-hydrodynamic instabilities; presence of turbulence behind or in the vicinity of the shock (Marscher, Gear \& Travis 1992); hence one can expect that the X-ray variability would be more pronounced for HSPs as compared to LSPs which is confirmed from our observations. These differences can lead to the apparent dichotomy between these two classes of blazars however we need statistically more observations to firmly conclude our findings.

1607.02053

1607

1607.05029_arXiv.txt

The advent of Wide Field Adaptive Optics (WFAO) systems marks the beginning of a new era in high spatial resolution imaging. The newly commissioned Gemini South Multi-Conjugate Adaptive Optics System (GeMS) combined with the infrared camera Gemini South Adaptive Optics Imager (GSAOI), delivers quasi diffraction-limited images over a field of $\sim$ 2 arc-minutes across. However, despite this excellent performance, some variable residues still limit the quality of the analyses. In particular, distortions severely affect GSAOI and become a critical issue for high-precision astrometry and photometry. In this paper, we investigate an optimal way to correct for the distortion following an inverse problem approach. Formalism as well as applications on GeMS data are presented.

By using multiple Laser Guide Stars (LGS), Wide Field Adaptive Optics (WFAO), improves the performance of high spatial resolution imaging: The AO-corrected images Field of View (FoV) is increased, as well as the fraction of the sky that can benefit from such correction. The Gemini Multi-Conjugated Adaptive Optics (MCAO) instrument GeMS is the first multi-Laser Guide Star operational system on sky. It has been implemented on the Gemini South telescope and commissionned in 2013. It works with two deformable mirrors conjugated at 0 and 9 km and a sodium-based LGS constellation composed of five spots: four are located at corners of a 60 arcsec square, with the fifth positioned in the center. GeMS, as a facility instrument, can direct its light output to different science instruments installed at the Cassegrain focus of the Gemini South telescope. Combined with the Gemini South Adaptive Optics Imager (GSAOI), it delivers near-diffraction limited images at Near-Infrared (NIR) wavelength (from 0.9 to 2.4 $\mu$m) over a FoV of 85$\arcs \times 85\arcs$. More details about the GeMS/GSAOI system and its commissioning results are described in detail in previous papers (see McGregor et al. (2004)\cite{McGregor2004}, Carrasco et al.(2011)\cite{Carrasco2011a}, D'Orgeville et al. (2012)\cite{DOrgeville2012}, Neichel et al. (2013 \cite{Neichel2013a} and 2014\cite{Neichel2014a}) and Rigaut et al. (2014)\cite{Rigaut2014}). However, despite the excellent performance of the GeMS/GSAOI system, the correction provided is not perfectly uniform and may generate variable Point Spread Function (PSF) over the field. For instance, we observe in the data an average Full Width Half Maximum (FWHM) over the field ranged between 80 mas and and 145 mas depending on the filter and the natural seeing. The standard deviation associated is typically around 15 mas, corresponding to a variation of 10-15\%. The average Strehl Ratio (SR) ranges between 3\% and 14\% with a standard deviation from 1\% to 3 \% (see Tab.\ref{tab1}). Those spatial variations of the PSF are present: \begin{enumerate} \item on single frames, mostly due to residuals from the AO correction ; \item on stacked images, critically amplified by optical distortions generated by the system ; \end{enumerate} Indeed, the optical components present in the instrument and the telescope as well as the GSAOI camera, introduce static and dynamical distortions. They depend on environmental parameters like the LGS spot size or the telescope pointing and may vary from one frame to another. These distortions degrade the resolution and strongly reduce the sensitivity when combining multiple frames. From there, the ability to deal with the spatial variation of the PSF is critical for high-precision astrometry and photometry studies (see Massari et al. (2015)\cite{Massari2015}, Turri et al. (2015)\cite{Turri2015}). Therefore, dedicated tools are needed to properly correct for distortion and preserve as much as possible the initial quality of the data when stacking them. In a first part of this paper we present a set of data obtained with the GeMS/GSAOI system: a recent observations of a very active and young star-forming region N159W located in the Large Magellanic Cloud (LMC) and that was previously studied by Deharveng et al. (1992)\cite{Deharveng1992}, Testor et al. (2007)\cite{Testor2006}, Chen et al. (2010)\cite{Chen2010a}. We obtained deep $J$, $H$, and \ks images in order to study the properties of the cluster stellar members and bring new elements to our understanding of the massive star formation process (for complete study see Bernard et al. (2016) \cite{Bernard2016}). The N159W field provides a large number of isolated stars and is therefore an ideal case to evidence limitations and experiment new methods of correction of distortions. In a second part of this paper we investigate an optimal way to correct for the distortion following an inverse problem approach. The method is based on the work presented in Gratadour et al. ($A\&A$ 2005)\cite{Gratadour2005} on image re-centering, but generalized to all distortion modes. We present here the formalism and simulation results as well as first application on the N159W GeMS/GSAOI data.

In this paper we have presented deep, high angular resolution, near-infrared images of the N159W star forming region located in the Large Magellanic Cloud. The data were obtained with the Near-Infrared Wide Field Adaptive Optic instrument GeMS/GSAOI recently implemented on the Gemini South telescope. These images aim at exploring the stellar content of the cluster and the massive star formation history of the region by doing a photometric study. Based on these real data, we evidence the limitations of the current reduction and analysis tools when it comes to Wide Field Adaptive Optics data. Indeed, despite the excellent performance of the adaptive optic correction, variable residuals are still limiting the quality of the image. Especially when combining multiple-frames, distortion effects consequently degrade the resolution and the sensitivity of the data. Thus, dedicated tools are needed to take into account these issues. In this article we presented a new method of correction of the distortion currently in development. This method is based on the least square minimization of a criterion, as previously done by D. Gratadour et al. (2005) \cite{Gratadour2005} for image re centering. We generalize here this method to any kind of distortion mode. The formalism as well as the implementation and performance results based on simulations as been presented. We showed a strong robustness to the noise that allows the use of the faintest stars in the field to compute a reliable correction of the distortion. As the quality of the correction is proportional to the number of stars used to apply the correction, the robustness to the noise appears to be an essential parameter for sparse field astrometry and photometry. Moreover, this method presents the advantage to estimate a "distortion-free" position of the stars in the sky, while existing tools base the correction on a non distortion-free reference. Finally, we have shown here a first application on real data. We estimated and corrected the distortion of ten random frames of the star forming region N159W. We found a typical standard deviation of position through the stack of 0.12 pixel $rms$ ($\sim$2.4 mas) which is consistent with simulations results. The perspectives for this work is now the optimization of several input parameters such as the number and the brightness of the stars used in the algorithm, the consideration of their flux or magnitude using weighting coefficients, the number of degrees of freedom considered, and the choice of the distortion base. We also might consider to correct independently each GSAOI detector and, in a second time, put them together. Finally, in order to evaluate the astrometric reliability of this method regarding real data, it is now essential to compare the distortion-free positions estimated with some distortion-free reference e.g. data obtained with the Hubble Space Telescope.

1607.05029

1607

1607.01393_arXiv.txt

We present the stellar mass ($M_{*}$) and {\it Wide-Field Infrared Survey Explorer} ({\it WISE}) absolute Band~1 magnitude ($M_{W1}$) Tully-Fisher relations (TFRs) of subsets of galaxies from the CO Legacy Database for the Galex Arecibo SDSS Survey (COLD GASS). We examine the benefits and drawbacks of several commonly used fitting functions in the context of measuring CO(1-0) line widths (and thus rotation velocities), favouring the Gaussian Double Peak function. We find the $M_{W1}$ and $M_{*}$ TFR, for a carefully selected sub-sample, to be $M_{W1} = (-7.1\pm0.6) \left[\log{\left(\frac{W_{50}/\sin{i}}{\text{km~s}^{-1}}\right)}-2.58\right] - 23.83\pm0.09$ and $\log{(M_{*}/M_{\odot})} = (3.3\pm0.3) \left[\log{\left(\frac{W_{50}/\sin{i}}{\text{km~s}^{-1}}\right)}-2.58\right] + 10.51\pm0.04$, respectively, where $W_{50}$ is the width of a galaxy's CO(1-0) integrated profile at $50\%$ of its maximum and the inclination $i$ is derived from the galaxy axial ratio measured on the SDSS $r$-band image. We find no evidence for any significant offset between the TFRs of COLD GASS galaxies and those of comparison samples of similar redshifts and morphologies. The slope of the COLD GASS $M_{*}$ TFR agrees with the relation of \citet{Pizagno:2005aa}. However, we measure a comparatively shallower slope for the COLD GASS $M_{W1}$ TFR as compared to the relation of \citet{TullyPierce2000}. We attribute this to the fact that the COLD GASS sample comprises galaxies of various (late-type) morphologies. Nevertheless, our work provides a robust reference point with which to compare future CO TFR studies.

\label{sec:intro} \subsection{Background} \label{subsec:background} The Tully-Fisher relation \citep[TFR;][]{Tully:1977aa}, one of the best studied galaxy scaling relations, can be derived easily by considering circular motion under gravity from a spherical mass distribution: $M(r) \propto v^2(r)r$, where $M$ is the mass of the body enclosed within a radial distance $r$ from its centre and $v$ is the rotational velocity at $r$. Defining the dynamical mass-to-light ratio $M/L$ and surface mass density $\Sigma \equiv M/\pi r^2$ of the body, one then obtains the TFR: \begin{equation} \label{eq:TFR} L \propto \frac{v^4}{\Sigma(M/L)}\,\,\,. \end{equation} Discovered observationally, the TFR was initially used as a galaxy distance indicator and, once coupled with the galaxy systemic velocity, as a tool to measure the Hubble constant \citep[e.g. ][]{Aaronson:1983, Bottinelli:1988, Giovanelli:1997, Sakai:2000,Tutui:2001}. Indeed, measuring a characteristic rotation velocity and an apparent luminosity, the TFR-predicted absolute luminosity then yields a distance estimate \citep[e.g.][]{Sandage:1976, Bottinelli:1980, Teerikorpi:1987, Fouque:1990, Bureau:1996, Ekholm:2000}. This of course assumes that one knows the mass-to-light ratio and surface density of the object studied, or more commonly that one compares populations of galaxies assumed to have identical $M/L$ and $\Sigma$ (e.g.\ galaxies of a given morphological type at a given epoch). The TFR implies a tight relationship between luminous mass (traced by $L$) and dynamical or total mass (traced by $v$), evidence that the growths over time of the luminous and dark matter present in galaxies are closely connected. Most importantly, given a reliable measure of distance, the TFR can be turned around and used to probe the mass-to-light ratio and surface density of galaxies. Indeed, many studies have used the zero-point, slope and scatter of the TFR to constrain cosmological models and models of galaxy formation \citep[e.g.][]{Cole:1989,Eisenstein:1996,Willick:1997,Steinmetz:1999,vandenBosch:2000,Yang:2003,Gnedin:2007,McGaugh:2012}. In addition, since a galaxy's redshift is itself a good distance indicator beyond the local Universe, the TFR is a powerful tool to directly probe the evolution of $M/L$ and $\Sigma$ over cosmological distances. Our goal here is thus to allow such studies by providing a local TFR benchmark, using molecular gas (CO) as the kinematic tracer. \subsection{Immediate Motivation} \label{subsec:motivation} The TFR was originally studied using blue optical restframe bands as proxies for the stellar mass of galaxies \citep[see e.g.][]{Tully:1977aa,McGaugh:1998aa,Ziegler:2002aa}. However, these tend to trace young stellar populations and are considerably affected by extinction, resulting in significant intrinsic scatter in the relation \citep[e.g.][]{Verheijen:2001aa}. In recent years and in the present work, near-infrared bands have been utilised instead \citep[see e.g.][]{Conselice:2005aa,Theureau:2007aa,Pizagno:2007aa,Puech:2008aa}, as these suffer little from extinction and better trace the bulk of the stellar mass \citep[e.g.][]{Dickinson:2003aa}, resulting in smaller intrinsic scatter. Whilst there is a large body of work on the TFR using atomic hydrogen (H{\small I}) 21~cm observations \citep[e.g.][]{Tully:1977aa,Sprayberry:1995aa,Bell:2001aa}, there are several advantages to using observations of carbon monoxide (CO), that traces the cold molecular gas in galaxies. First, atomic hydrogen in galaxies is currently only routinely detected in the local Universe. The Square Kilometre Array (SKA) precursors Australian Square Kilometre Array Pathfinder (ASKAP) and Karoo Array Telescope (MeerKat) can detect H{\small I} to moderate redshifts \citep[$z\lesssim0.4$;][]{Meyer:2009,Holwerda:2010,deBlok:2011,Duffy:2012}, but the SKA itself will be required to routinely detect galaxies at $z>1$ \citep{Abdalla:2015, Yahya:2015}. Multiple transitions of CO are however already routinely detected in the bulk of the star-forming galaxy population at $z\approx1$--$3$ \citep[e.g.][]{Daddi:2010, Magdis:2012, Magnelli:2012, Combes:2013, Freundlich:2013,Tacconi:2010,Tacconi:2013,Genzel:2015}, and in star-bursting galaxies up to $z\approx7$ \citep[e.g.][]{Walter:2004, Riechers:2008a, Riechers:2008b, Riechers:2009, Wang:2011, Wagg:2014}. CO thus allows to extend TFR studies probing the mass-to-light ratio and surface density of galaxies to the earliest precursors of today's galaxies. Second, previous work has shown that the H{\small I} discs of galaxies, that are typically more extended spatially than the molecular gas, can also be kinematically unrelaxed, with e.g.\ large-scale warps \citep[e.g.][]{Verheijen:2001aa}. Molecular gas is generally more dynamically relaxed and suffers less from such problems. More importantly, the atomic hydrogen in early-type galaxies is often significantly disturbed, with much of the gas lying in tidal features or nearby dwarf galaxies \citep[see e.g.][]{Morganti:2006aa, Serra:2012}, thus confusing low-resolution observations such as those obtained with single-dish telescopes. High-resolution interferometric observations are thus necessary to identify those early-type galaxies with a regular H{\small I} distribution appropriate to derive reliable TFRs \citep[see e.g.][]{denHeijer:2015}. On the other hand, \citet{Davis:2011aa} clearly showed that CO single-dish observations easily yield robust and unbiased TFRs for early-type galaxies. CO observations thus offer the attractive possibility to derive TFRs more accurate than those currently available, and this across the entire Hubble sequence. Our goal here is therefore to establish a benchmark CO TFR of local galaxies, as a pre-requisite to extend the relation to higher redshifts. There are clearly tracers other than CO that can be used at large redshifts, primarily optical ionised gas emission lines such as H$\alpha$, [O{\small II}] and [O{\small III}], and these should also be pursued to provide independent probes of the evolution of the TFR \citep[see e.g.][]{Cresci:2009aa,Gnerucci:2011aa,Miller:2011aa,Miller:2012aa}. However, it is known that ionised gas discs at $z\approx1$--$3$ are turbulent \citep[see e.g.][]{Schreiber:2006, Swinbank:2012}, and great care must be taken when measuring and interpreting their rotational motions \citep[e.g.][]{Wisnioski:2015,Stott:2016,Tiley:2016}. With the advent of the Atacama Large Millimeter/submillimeter Array (ALMA)\footnote{http://www.almascience.org/}, CO emission tracing dynamically cold gas is more easily detectable in high-redshift galaxies than ever before, and embarking on CO TFR studies across the Hubble sequence and redshift is particularly timely. In this paper, we thus take a step toward establishing a local benchmark for the CO TFR, using the CO(1-0) line as a kinematic tracer. In future work our TFRs will be compared to those of a local sample (Torii et al., in prep.) and a sample of $z\lesssim0.3$ galaxies (Topal et al., in prep.). The sample, photometric data and kinematic data used in this paper are described in \S~\ref{sec:data}. TFRs are derived in \S~\ref{sec:TFRs} (the rotation velocity measure adopted is defined in \S~\ref{subsec:measurew50} and extensively tested in Appendix~\ref{sec:velocity}). The results are discussed in \S~\ref{sec:discussion}. We summarise and conclude briefly in \S~\ref{sec:conclusions}.

\label{sec:conclusions} In an effort to firmly establish the CO Tully-Fisher relation (TFR) as a useful tool to probe the evolution of galaxies over cosmic time, we first tested the self-consistency and robustness of four functions appropriate to fit the integrated line profiles of galaxies, particularly in the low signal-to-noise ratio regime characteristic of molecular gas observations. The Gaussian Double Peak function was deemed to be the most self-consistent and to suffer the least from possible systematic biases as a function of the amplitude-over-noise ratio $A/N$, the galaxy inclination $i$ and the intrinsic flat circular velocity $V_{\text{c,flat}}$. We then constructed the {\it WISE} $W1$-band and stellar mass TFRs relations of galaxies drawn from the COLD GASS sample, both for an initial sample of all galaxies with available data, and for a restricted sub-sample of galaxies with high quality measurements (thus decreasing the scatter). The rotation of the galaxies was determined by fitting the Gaussian Double Peak function to the integrated CO(1-0) line profile of each galaxy, and then measuring the width at $50\%$ of the peak of the resultant fit ($W_{50}$). The $W1$ magnitudes were drawn directly from the {\it WISE} catalogue, and the stellar mass for each galaxy was determined via SED fitting of SDSS photometry. The TFRs obtained from unconstrained forward fits have shallower slopes than those expected from previous studies. Considering only the more robust reverse fits, however, the best-fit TFRs for the final COLD GASS sub-sample are \begin{equation} M_{W1} = (-7.1\pm0.6) \left[\log{\left(\frac{W_{50}/\sin{i}}{\text{km~s}^{-1}}\right)}-2.58\right] - 23.83\pm0.09 \end{equation} \noindent and \begin{equation} \log{(M_{*}/M_{\odot})} = (3.3\pm0.3) \left[\log{\left(\frac{W_{50}/\sin{i}}{\text{km~s}^{-1}}\right)}-2.58\right] + 10.51\pm0.04\,\,\,. \end{equation} The unconstrained reverse fit slope of the COLD GASS sub-sample $W1$-band TFR is still marginally shallower than that of \citet{TullyPierce2000}, but the slope of the stellar mass TFR agrees within the uncertainties with the relation of \citet{Pizagno:2005aa}. The intrinsic scatter (from forward fitting) is $0.59\pm0.05$~mag and $0.27\pm0.02$~dex for the $W1$-band and stellar mass sub-sample TFR, respectively. Fixing the slopes to those of the relations from the comparison samples, small offsets are found with respect to the comparison samples, that are however less than the intrinsic scatters. The COLD GASS samples therefore agree with the comparison samples, although they have slightly larger scatters than expected. Possible causes of the increased scatters were discussed and include the method adopted to measure inclinations, and fitting a single TFR to samples of galaxies with various primarily late-type morphologies. Importantly, we showed that for a subset of COLD GASS galaxies in the final sub-sample with both CO(1-0) and H{\small I} data, the intrinsic and total scatters of the CO(1-0) TFRs were less than those of the same TFRs constructed using H{\small I} integrated profiles. The COLD GASS initial sample and final sub-sample contain a number of galaxies comparable to those in previous CO TFR studies. Our work thus provides a robust local benchmark to be used for comparison with future CO work. In particular, Torii et al. (in prep.) will build a local reference sample based on observations of very nearby galaxies with the NANTEN2 telescope, and utilising identical fitting methods to ours. Topal et al. (in prep.) will compare the TFRs presented here with those measured using a sample of galaxies at $z\lesssim0.3$, including luminous infrared galaxies (LIRGs) and galaxies from the Evolution of Gas in Normal Galaxies (EGNoG) survey. With the dawn of ALMA (and NOEMA), it is now possible to relatively rapidly measure the CO emission of galaxies to large redshifts, when the first objects were forming and slowly settling into the discs we see today. In particular, significant samples of galaxies observed in CO are now being built to probe the epoch of peak global star formation activity ($1\lesssim z\lesssim3$), when turbulent gas-rich galaxies were building the bulk of their stellar mass. Our work therefore provide a robust reference point with which to compare future TFR studies of those objects.

1607.01393

1607

1607.08751_arXiv.txt

We describe a novel, very fast and robust, directed search incoherent method for periodic gravitational waves from neutron stars in binary systems. As directed search, we assume the source sky position to be known with enough accuracy, but all other parameters (including orbital ones) are supposed to be unknown. We exploit the frequency-modulation due to source orbital motion to unveil the signal signature by commencing from a collection of time and frequency peaks (the so-called $\textit{peakmap}$). We validate our pipeline adding 131 artificial CW signals from pulsars in binary systems to simulated detector Gaussian noise, characterised by a power spectral density $S_h = 4 \times 10^{-24}$~Hz$^{-1/2}$ in the frequency interval $[70,\, 200]$~Hz, which is overall commensurate with the advanced detector design sensitivities. The pipeline detected 128 signals, and the weakest signal injected and detected has a gravitational-wave strain amplitude of $\sim 10^{-24}$, assuming one month of gapeless data collected by a single advanced detector. We also provide sensitivity estimations, which show that, for a single-detector data covering one month of observation time, depending on the source orbital Doppler modulation, we can detect signals with an amplitude of $\sim 7\times 10^{-25}$. By using three detectors, and one year of data, we would easily gain more than a factor 3 in sensitivity, translating into being able to detect weaker signals. We also discuss the parameter estimate proficiency of our method, as well as computational budget, which is extremely cheap. In fact, sifting one month of single-detector data and 131~Hz-wide frequency range takes roughly 2.4 CPU hours. Due to the high computational speed, the current procedure can be readily applied in ally-sky schemes, sieving in parallel as many sky positions as permitted by the available computational power. The novel procedure has a sensitivity comparable and slightly higher than other competing pipelines present in literature, but is several orders of magnitude faster than those. We also introduce (ongoing and future) approaches to attain sensitivity improvements and better accuracy on parameter estimates in view of the use on real advanced detector data.

Introduction} Since the early 1960s, when the first gravitational-wave bar detector was developed~\cite{Weber1960}, the experiments that aimed at the detection of gravitational radiation, planned in laboratories throughout the world, have been in continuous progress~\cite{Auriga,Nautilus,Explorer,Allegro,Bonaldi,LeaciPRA,LeaciPRD,LeaciCQG,Virgo,Geo,TAMA,ligoref,IPWRA}. At present, LIGO~\cite{ADvLIGO} and Virgo~\cite{AdvVirgoRef:2009} are the most sensitive ground-based gravitational-wave detectors. Following a major upgrade lasted for 5 years, with consequent improvement in sensitivity~\cite{AdvDetO1}, the LIGO observatories resumed data taking with the first observing science run during which they have collected data between September 2015 and January 2016. On September 14, 2015, the advanced LIGO interferometers detected for the first time a coincident transient gravitational-wave signal produced by the coalescence of a pair of black holes~\cite{GWevent}, marking thus the official beginning of a new era: the era of gravitational-wave astronomy. There are however other classes of gravitational-wave signals, which have still to be detected, such as long-lived continuous waves (CWs), which are expected to be emitted by rapidly rotating neutron stars (NSs) with nonaxisymmetric deformations~\cite{Owen:2005fn}. We expect $\mathcal{O}(10^8)$ of these sources to exist in the Galaxy, but only $\sim2\,500$ NSs (mostly pulsars) have been electromagnetically observed~\cite{atnf}. Roughly 1\,300 of these observed radio pulsars are located in binary systems, and have rotation rates that can allegedly emit CWs in the advanced LIGO-Virgo sensitivity band. This promising class of sources is the target of the current paper. The detection of CW signals will enrich the understanding we have about the emitting objects (i.e., NSs), providing us insight about the equation of state of the matter at supranuclear densities inside the NSs, know exactly the NS degree of asymmetry, and have also demographic and evolutionary information about these sources. In the rest frame of the NS, CWs have a constant amplitude and are quasimonochromatic with a slowly decreasing intrinsic frequency. They are received at Earth-based detectors with a Doppler modulation due to the relative motion between the source and the detector. Consequently, the observed phase evolution depends on the intrinsic signal frequency, frequency time derivatives (also known as ``spindown'' terms), and source sky position. If the source is located in a binary system, there is a further frequency-modulation caused by the source orbital motion, which in general is described by five unknown Keplerian parameters~\cite{Dhurandhar:2000sd}, as detailed in Sec.~\ref{sec:binary-cw-phase}. The weakness of the expected signal requires long integration times, typically of the order of a few months or years, to accumulate a signal-to-noise ratio (SNR) sufficient for detection as, for a coherent (incoherent) search, the SNR scales as the square (fourth) root of the length of the observational time (i.e. the length of the data being analyzed)~\cite{Jaranowski:1998qm,Astone:2000jza}. All-sky, wide frequency searches over long observation times cannot be treated by using standard coherent methods (where the phase information is used), as is the case for targeted and narrow-band searches for known pulsars~\cite{VelaSpinDAbadie:2011md,NarrowBandVelaCrabAasi:2014jln}, because of the demanding computational burden. Hence, hierarchical approaches have been proposed~\cite{Brady:1998njStackSlide,Krishnan:2004sv,Cutler:2005pn}, where the entire data set is split into shorter segments. Each segment is analyzed coherently, and afterwards the information from the different segments is combined incoherently (which means that the phase information is lost). The hierarchical approaches allow us to dramatically reduce the analysis computational time at the cost of a relatively small sensitivity loss. The additional source orbital parameters make the sieved parameter space to blow up, resulting in a prohibitive computational cost. Hence, it becomes pressing to develop robust strategies to detect CWs emitted by NSs orbiting a companion object, and being able to reach a tradeoff between computational cost and sensitivity. Although CWs have not been detected so far by analysing data from initial LIGO and Virgo detectors, stringent upper limits have been set on the gravitational-wave signal strength for both isolated pulsars~\cite{Aasi:2012fw,Aasi:2015rar,GalcticCAasi:2013jya,NarrowBandVelaCrabAasi:2014jln} and pulsars in binary systems~\cite{TwoSpectAllSkyAasi:2014erp,ScoX1directAasi:2014qak}. A particularly interesting type of potential CW sources are NSs in low-mass x-ray binaries (LMXBs), with Scorpius X-1 being its most prominent representative~\cite{Watts:2008qw}. Several searches for CW signals from Scorpius X-1 have been performed (without any detections) on data from initial LIGO~\cite{Abbott:2006vg,ScoX1directAasi:2014qak}, and new pipelines have been developed~\cite{Leaci:2015bka} and recently tested in a Scorpius X-1 mock data challenge (MDC)~\cite{Messenger:2015tga}. We present here an incoherent strategy that allows us to perform directed searches for CWs emitted by NSs in binary systems (LMXB like sources), exploring a wide frequency range and source orbital parameters at a paltry computing cost. We also include investigations of pulsars in eccentric orbits, which we know to exist~\cite{HulseTay,EccOrbitWatts:2016uzu,atnf}. The method is based on selecting significant peaks from a short Fast-Fourier-Transform (FFT) database (SFDB), and exploiting the frequency modulation pattern produced by the source orbital motion to detect a potential CW signal. We show the performance of the current method to detect CW signals by analysing pure Gaussian noise data to which we add hundreds of fake signals. We consider one month of gapless data, i.e., data taken continuously during the assumed observation time from the Virgo (or equivalently LIGO) detector in its advanced configuration. We restrict our investigation to constant-frequency CW signals (i.e. we neglect frequency time derivatives). This is motivated by the assumed steady-state torque-balance situation in LMXB like sources, which are our main target of interest. However, the corresponding fluctuations in the accretion rate are expected to cause some stochastic frequency drift, and one will therefore need to be careful to restrict the maximal \textit{coherence time} (i.e. FFT duration) in order to limit the frequency resolution. In fact, in~\cite{Leaci:2015bka} it has been shown that the maximal FFT duration would be restricted by the astrophysical concern of \textit{spin wandering}, namely a stochastic variability of the spin frequency due to variations in the accretion rate. In the present work we neglect the spin wandering effect, as our new robust methodology is expected to be unaffected by these variations. Exhaustive future studies will be able to shed light on these considerations and will be presented in a subsequent paper. The novel procedure we present allows us to detect gravitational-wave signals with strain amplitude of $\sim 10^{-24}$, the weakest value we used for the simulations, assuming one month of gapeless data collected by a single advanced detector. Sensitivity estimation studies (see Sec.~\ref{Sec:SensEstim}) show however that, depending on the source orbital Doppler modulation, the current method can detect signals with an amplitude of $\sim 7\times 10^{-25}$. By using three advanced detectors, and one year of data, we would be able to further improve the sensitivity by a factor greater than 3. The paper is organized as follows. Section~\ref{Sec:signal} provides a general and brief introduction of the expected signal model. Section~\ref{Sec:AnM} describes details on data analysis approach necessary to produce input data set. In Secs.~\ref{Sec:PS} and~\ref{FFTfracOrbit} we discuss the choices of the investigated parameter space and FFT duration. In Sec.~\ref{Sec:method} we present a rigorous description of the novel strategy used to detect CW signals orbiting a companion object. Sections~\ref{RecoveryEstim} and~\ref{Sec:h0recovery} show the key results of the procedure. Sensitivity estimates and computing cost budget are detailed in Secs.~\ref{Sec:SensEstim} and~\ref{Sec:TeoCtot}, respectively. Finally, Sec.~\ref{Sec:Conc} contains concluding remarks, underway search improvements, and future prospects.

1607.08751

1607

1607.06861_arXiv.txt

BICEP3, the latest telescope in the BICEP/Keck program, started science observations in March 2016. It is a 550mm aperture refractive telescope observing the polarization of the cosmic microwave background at \SI{95}{GHz}. We show the focal plane design and detector performance, including spectral response, optical efficiency and preliminary sensitivity of the upgraded BICEP3. We demonstrate \SI{9.72}{\ukrts} noise performance of the BICEP3 receiver.

Measurements of the polarization of the Cosmic Microwave Background provide key information to further our understanding of the early universe. The $\Lambda$CDM model predicts a E-mode polarization pattern in the CMB at the level of a few $\mu$K and arc-minute B-mode polarization arises from gravitational lensing of E-mode power by the large scale structure of the universe. But inflationary gravitational waves may be a source of degree scale B-mode polarization and a detection of such signal can be use to constrain the tensor-scalar ratio $r$ and place limits on the energy scale and potential from inflation~\cite{Kamionkowski15}. However, several galactic mechanisms can generate B-mode foregrounds; to disentangle the cosmic signal from galactic ones, we need to probe the polarization of the CMB at multiple frequencies with high sensitivity. The BICEP/ Keck team has deployed multiple telescopes to the South Pole since 2006; we use small aperture, refracting telescopes with high sensitivity receivers to map the degree scale B-mode signal. The Keck Array is in its fifth season, currently observing at \SI{150}{GHz} and \SI{220}{GHz} and previously observed at \SI{95}{GHz}, \SI{150}{GHz} and \SI{220}{GHz}~\cite{Ogburn12} optical bands with 5 BICEP2 style cameras. BICEP3 is the latest addition to this program and was first deployed to the South Pole Station in 2015. It is a \SI{550}{mm} aperture, on-axis, refractive polarimeter designed to observe at \SI{95}{GHz}. During the first observing season, the focal plane was only partially filled with 1152 detectors~\cite{Wu15}, whereas in this year's observing season, the instrument is complete with 2560 detectors. Combining data with BICEP2/Keck, Planck and South Pole Telescope, BICEP3 is projected to set upper limits on $r < 0.03$ at $95 \%$ confidence. Ref.~\citenum{Grayson16} shows an overview of the BICEP3 telescope design and observing strategy.

In this proceedings, we present the the design of the BICEP3 focal plane module and its readout architecture. This compact design increases the packing density of the detectors and allows more efficient use of optically illuminated area on the focal plane. The modular design makes future replacements and upgrades easier. We also show great improvement in detector performance in the second season of BICEP3, increasing detector yield from 436 to 951 polarization-sensitive pixels, reducing the per-detector NET from \SI{395}{\ukrts} to \SI{333}{\ukrts}, and achieving a receiver NET of \SI{9.72}{\ukrts}.

1607.06861

1607

1607.06099_arXiv.txt

We use the scatter in the stellar-to-halo mass relation to constrain galaxy evolution models. If the efficiency of converting accreted baryons into stars varies with time, halos of the same present-day mass but different formation histories will have different $z=0$ galaxy stellar mass. This is one of the sources of scatter in stellar mass at fixed halo mass, $\slogm$. For massive halos that undergo rapid quenching of star formation at $z\sim 2$, different mechanisms that trigger this quenching yield different values of $\slogm$. We use this framework to test various models in which quenching begins after a galaxy crosses a threshold in one of the following physical quantities: redshift, halo mass, stellar mass, and stellar-to-halo mass ratio. Our model is highly idealized, with other sources of scatter likely to arise as more physics is included. Thus, our test is whether a model can produce scatter lower than observational bounds, leaving room for other sources. Recent measurements find $\slogm=0.16$ dex for $10^{11}$ \msol\ galaxies. Under the assumption that the threshold is constant with time, such a low value of $\slogm$ rules out all of these models with the exception of quenching by a stellar mass treshold. Most physical quantities, such as metallicity, will increase scatter if they are uncorrelated with halo formation history. Thus, to decrease the scatter of a given model, galaxy properties would correlate tightly with formation history, creating testable predictions for their clustering. Understanding why $\slogm$ is so small may be key to understanding the physics of galaxy formation.

In its simplest form, abundance matching connects galaxies with dark matter halos by the rank-order of both objects: the $N$th most massive galaxy resides in the $N$th most massive halo. The success of this paradigm rests on the assumption all halos of mass $\mhalo$ have galaxies with mass stellar $\mgal$ inside them, regardless of the formation history of each halo. Any scatter in the relation is put in by-hand, post-facto. In essence, abundance matching rests on the idea that galaxy formation is a `path-independent' process. Using the mean growth of halos, combined with measurements of the galaxy stellar mass function at various redshifts, one can use abundance matching to determine the average path of stellar mass growth in bins of halo mass (\citealt{conroy_wechsler:09, behroozi_etal:13_letter, behroozi_etal:13}, hereafter B13, \citealt{moster_etal:13}). From this, one can show the efficiency of converting accreted baryons into stars, $\fcon$. For massive halos---those whith $z=0$ masses of $10^{13}$ \msol, which will be the focus of this paper---this function monotonically increases with cosmic time up until it peaks at $z\sim 2$, whereupon it turns over and galaxy growth quickly stalls. These results are in agreement with analyses of the stellar populations of massive galaxies, which imply rapid galaxy growth at high redshift with limited growth after $z\sim 2$ (e.g., \citealt{thomas_etal:05}). These results reflect the {\it average} formation history of galaxies within halos. However, dark matter halos of fixed mass can have widely varying formation histories. Two halos with a present-day mass of $10^{13}$ \msol\ can differ by a factor of five at 1-$\sigma$ at $z=3$ (e.g., \citealt{wechsler_etal:02}). Any dependence of the baryonic conversion efficiency, $\fcon$, with redshift---either explicitly, or implicitly through a dependence on $\mhalo(z)$, $\mgal(z)$, or other quantity---with break the path-independence of galaxy formation. Two halos with the same $z=0$ dark matter mass will {\it not} have the same mass galaxy in them. The distribution of halo formation histories is then one of the prime sources of scatter in the stellar mass to halo mass relation. In this paper we will use the measurements of the scatter in stellar mass at fixed halo mass, $\slogm$, to put constraints on how $\fcon$ can vary with time for massive galaxies of present-day stellar mass $\mgal\approx 10^{11}$ \msol. These galaxies form in halos of $\mhalo\approx 10^{13}$ \msol (B13, \citealt{moster_etal:13}). We focus on massive galaxies for two reasons: first, these galaxies are nearly uniformly quiescent (\citealt{chen_etal:12, reid_etal:16}), thus the process that quenches star formation has already occurred in these halos. As we will show, because this process of quenching must occur over a short time span, it causes an extreme break of the path-independence of galaxy formation, and thus has a strong impact on $\slogm$. Second, although these galaxies are quite massive, abundance matching informs us that the buildup of stellar mass within these halos is due to in-situ star formation, and not by merging. B13 and \cite{moster_etal:13} both find that the fraction of stellar mass from in-situ growth in these halos if 90\% at $z=0$ and 95\% at $z=0.5$, which is the redshift of the Baryon Oscillation Spectroscopic Survey (BOSS; \citealt{dawson_etal:13}) galaxy sample from which we will take observations of $\slogm$. Thus, the dominant source of scatter in stellar masses within these halos is star formation and not merging, which can dominate $\slogm$ in higher mass halos (\citealt{gu_etal:16}). Using the clustering and abundance of BOSS galaxies, \cite{tinker_etal:16_boss} found $\slogm$ for $10^{13}$ \msol\ halos to be 0.16 dex. This value removes statistical errors from the stellar mass estimates, but does not remove systematic random errors incurred from the stellar mass estimation method itself. This value is in good agreement with other measurements of $\slogm$ from other galaxy samples at lower redshifts (\citealt{more_etal:11, reddick_etal:13, zu_mandelbaum:16}). Although this is an upper limit, 0.16 dex is a shockingly low value for a quantity---stellar mass---that is influenced by a series of disparate processes, all with their own intrinsic distributions, such as metallicity, AGN feedback, supernovae feedback and winds, gas-rich merging, and the different baryonic accretion rates that different halos experience. The models we present in this paper are highly simplified, incorporating none of the physical effects just listed. Thus, our test for whether a model for quenching is valid is whether it can yield a value of $\slogm$ below the observed value, leaving room for other sources of scatter from more physical effects. Throughout this paper, we assume a flat-\lcdm\ cosmology with $\om=0.3$, $\s8=0.8$, and $h=0.7$. We will use redshift as our time unit quite often, especially in our model parameterizations, but will show plots as expansion factor $a$, which is a more natural time unit for the growth of galaxies. \begin{figure*} \includegraphics[width=7in ]{galpath_hmass.ps} \vspace{-11.cm} \caption{ \label{galpath_hmass} The stellar mass growth of halos of various masses. In each panel, the solid blue curve is the median value from a series of halo merger trees. This curve is fit to the results of \citet{behroozi_etal:13}, which are shown with the red circles. The evolutionary tracks of a subsample of halos are shown in the thin gray curves. For each halo, the same $\fcon(z)$ function is applied, thus the differences in stellar mass are all driven by the differences in halo mass growth. In each panel, the amplitude of $\fcon(z)$ is varied but the redshift-dependence is the same. This figure shows three things: (1) a single function of redshift con describe $\fcon$ for halos $\le 10^{12}$ \msol, as well as the high-redshift growth in massive halos (2) that $10^{13}$ \msol\ halos undergo a rapid transition in their star formation efficiency at $z\sim 2$. This is a restatement of the previous results found in abundance matching studies And (3), different halo formation histories impart a scatter of $\la 0.1$ dex in $\log\mgal$ for a generic $\fcon(z)$ function. } \end{figure*} \begin{figure*} \includegraphics[width=7in ]{galpath6.ps} \vspace{-5.5cm} \caption{ \label{galpath6} {\it Top Row:} Best fit models for redshift quenching, halo quenching, and galaxy quenching. The solid blue curves are the median of the set of halo merger trees, which are fit to the data of \citet{behroozi_etal:13}, shown with the red circles. A sample of individual halos are shown with the gray curves. The value of the scatter, $\slogm$, is shown for each model in the panel. The blue dotted curves show the 68\% range of $\mgal$ around the median. {\it Bottom Row:} Same as the top row, but now for the three different models where the threshold for quenching is based on a critical $\mgal/\mhalo$ ratio, and after that point $\fq$ is parameterized by $z$, $\mhalo$, and $\mgal$, respectively. } \end{figure*}

We have implemented a simple model to explore how different models for quenching star formation in galaxies can impact the scatter of stellar mass at fixed halo mass, $\slogm$, for which we have excellent constraints from the clustering of massive galaxies. We test models in which quenching begins at some critical redshift, $\mhalo$, $\mgal$, or $\mgal/\mhalo$ ratio. We find: \begin{itemize} \item Under the assumption that the quenching threshold is constant with time, only galaxy quenching is consistent with the measurements of $\slogm$. \item The scatter imparted by halo and redshift quenching is somewhat larger than observations, while the scatter yielded by ratio quenching is nearly double that observed. \item To decrease the scatter induced by halo quenching, $\mhcrit$ must decrease with time. This goes in the opposite direction implied by the growth of metallicity in galaxies. \item There is little to no room for any stochasticity in $\mhcrit$ from halo-to-halo. The observed scatter in $\zz$, if uncorrelated with halo formation history, would raise $\slogm$ above the observed values. \item Decreasing the scatter in each model would require strong correlations between galaxy properties, such as metallicity or mean stellar age, and halo formation history. \end{itemize} Although galaxy quenching yields, by far, the lowest values of $\slogm$, there are some obvious questions that arise from this model. Observations of central galaxies within halos show that the quenched fraction varies smoothly with increasing galaxy mass, and not consistent with a threshold value (\citealt{weinmann_etal:06a}). Stochasticity in the $\mgcrit$ threshold may alleviate this tension, as well as some correlation of another galaxy property with the quenching threshold. Applying this model to the full galaxy population, rather than just massive galaxies, can resolve this question. This will be pursued in a future paper. There are many simplifications and assumptions that are used to construct the models described above, but many---if not most---create testable predictions that may be constrained by existing data. The current spectroscopic sample of massive galaxies presently contains upwards of $\sim 2$ million galaxies and reaches over 7 Gyr into the cosmic past through the combination of SDSS, BOSS, and now eBOSS data (for which the first clustering measurements have been published; \citealt{zhai_etal:16}). Although the quality of many of these spectra make detailed stellar population analysis untenable on a per-object basis, the clustering of these galaxies contains a wealth of information, well beyond the value of $\slogm$ used here. At fixed $\mgal$, the dependence of any galaxy property on halo formation history will show up in the clustering of those galaxies. This halo assembly bias has been shown conclusively in numerical simulations, and has recently been detected in cluster-sized dark matter halos observationally (\citealt{miyatake_etal:16, more_etal:16}). \cite{saito_etal:16} used two-point clustering to demonstrate that $z\sim 0.5$ BOSS galaxies are consistent with a model in which the colors of massive galaxies are correlated with halo age. For the fiducial implementations of galaxy and halo quenching, in which the threshold is constant in time, assembly bias is a natural consequence. Early-forming halos will have older stellar populations at fixed $\mgal$, which could impart an assembly bias signal based on galaxy color, luminosity, and metallicity at fixed mass. Redshift and ratio quenching, on the other hand, yield little correlation between mean stellar age and the halo formation time. Different implementations of the halo quenching model yield different assembly bias signals as well. Halo quenching models with a time-varying threshold tend to reduce the amount of assembly bias in the models because they delay quenching in early-forming halo and accelerate it in late-forming halos. Further investigation, both through the clustering of massive galaxies and by using stellar population synthesis models to calculate the observable properties of galaxies with various mean stellar ages, will be fruitful in further differentiating models or constraining a the parameter space of a specific model. We have assumed that other physical mechanisms that effect star formation within dark matter halos would add to the scatter induced by variations in halo formation history. We have specifically focused on metallicity as a probable source of such scatter within the halo quenching model, either in the form of stochastic variations of the quenching threshold or a redshift dependence that would widen the scatter over the fiducial model. It is always possible that these physical mechanisms correlate with halo formation history in a way to reduce the scatter in stellar mass. For example, if metallicity correlated with halo formation history such that early-forming halos have higher metallicity that later-forming halos, the dependence of $\mhcrit$ on $\zz$ would help reduce scatter by allowing early-forming halos to convert more of their baryons into stars than they would in the constant threshold model. Such a model would create an assembly bias signal on the metallicity of massive galaxies. Another assumption we have made in the construction of these models is that $\fcon$ is a universal function that only depends on $z$. Reducing the scatter in the pre-quenching phase of evolution would also reduce post-quenching scatter. Given the large variation in halo mass at $z=3$ for present day $10^{13}$ \msol\ halos---roughly a factor of five---it is difficult to construct a model that creates minimal scatter in $\mgal$ within these halos at high redshift that isn't highly ad hoc. For $z=0$ halos below the quenching threshold, there still exists a scatter in stellar mass that is larger than that shown in Figure \ref{galpath_hmass}. A model in which $\fcon$ depends on $z$ and some second parameter, such as $\mhalo(z)$ or $\mgal(z)$, may shrink the scatter at $z\sim 2$, but it is not clear that such a model would yield small scatter at $z=0$, as well as reproduce the measurements of SFR$(z)$ for galaxies of various masses, which is well fit by Equation (\ref{e.intmgal}) (B13, \citealt{moster_etal:13}). Any strong conclusions made here depend on strong assumptions. However, this work represents a proof-of-concept that the scatter in the stellar to halo mass relation contains significant information for constraining the physics of galaxy formation. The simplified models presented here are sufficient to test simple models of galaxy formation and evolution, and as we isolate the region of parameter space that is consistent with observations, more sophistication can be added to these models to properly explore this parameter space, and more data can be added by measuring clustering of massive galaxies to test for assembly bias in various physical quantities. Outside of empirical models of galaxy formation, semi-analytic and hydrodynamic explorations of galaxy formation physics should be utilizing $\slogm$ in the assessment of their models. The processes that regulate star formation will also determine the scatter in the total amount of star formation. Understanding why $\slogm$ is so small may be key to our understanding of how the present day galaxy population came to be.

1607.06099

1607

1607.01446_arXiv.txt

We present our multiwavelength analysis of a prototype \HI-excess galaxy, GASS 3505, selected based on having a large gas content ($M_{\rm HI} = 10^{9.9}$ \msun) compared to its little associated star formation activity ($\sim$0.1 \msun\ yr$^{-1}$) in the GALEX Arecibo SDSS Survey (GASS). Very Large Array (VLA) observations show that the \HI\ in GASS 3505 is distributed in a regularly rotating, extended ($\sim$50 kpc radius) gas ring. In the SDSS optical image GASS 3505 appears as a bulge-dominated galaxy, however deep optical imaging reveals low surface brightness ($\gtrsim25$ mag arcsec$^{-2}$) stellar emission around the central bulge. Direct evidence for accretion is detected in form of an extended ($\sim$60 kpc) stellar stream, showing that GASS 3505 has experienced a minor merger in the recent past. We investigate the possibility that the \HI\ ring in GASS 3505 was accreted in such a merger event using N-body and smoothed particle hydrodynamic (SPH) simulations. The best model that reproduces the general properties (i.e., gas distribution and kinematics, stellar morphology) of the galaxy involves a merger between the central bulge and a gas-rich ($M_{\star}$ = 10$^9$ \msun\ and $M_{\rm HI}$/$M_{\star}$ = 10) disk galaxy. However, small discrepancies in the observed and modeled properties could suggest that other sources of gas have to be involved in the build-up of the gas reservoir. This work is the first step toward a larger program to investigate the physical mechanisms that drive the large scatter in the gas scaling relations of nearby galaxies.

1607.01446