subfolder
stringclasses 367
values | filename
stringlengths 13
25
| abstract
stringlengths 1
39.9k
| introduction
stringlengths 0
316k
| conclusions
stringlengths 0
229k
| year
int64 0
99
| month
int64 1
12
| arxiv_id
stringlengths 8
25
|
|---|---|---|---|---|---|---|---|
1609 | 1609.06409_arXiv.txt | Pulsar timing observations have revealed planets around only a few pulsars. We suggest that the rarity of these planets is due mainly to two effects. First, we show that the most likely formation mechanism requires the destruction of a companion star. Only pulsars with a suitable companion (with an extreme mass ratio) are able to form planets. Second, while a dead zone (a region of low turbulence) in the disk is generally thought to be essential for planet formation, it is most probably rare in disks around pulsars because of the irradiation from the pulsar. The irradiation strongly heats the inner parts of the disk pushing the inner boundary of the dead zone out. We suggest that the rarity of pulsar planets can be explained by the low probability for these two requirements -- a very low--mass companion and a dead zone -- to be satisfied. | \label{intro} There are currently five exoplanets in three planetary systems that have been detected through pulsar timing to be orbiting pulsars, rapidly rotating highly magnetized neutron stars \citep{Lorimer2008}. The masses and semi--major axes of these planets are shown in Table~\ref{tab} and plotted in the blue points in Fig.~\ref{Mass_semi}. The first planets to be discovered (and immediately confirmed) outside of our solar system were in fact found around the pulsar PSR B1257+12 \citep{Wolszczan1992,Wolszczan1994, Wolszczan2012}. This pulsar has three very close--in planets. The outer two planets are coplanar within $6^\circ$ and all three have low eccentricity, implying a disk origin \citep{Konacki2003}. The outer two planets are close to a 3:2 mean--motion resonance. This suggests that the planets did not form at their current location but instead probably migrated in from farther out, possibly through a gas disk \citep[e.g.][]{Terquem2007}. A planet has also been observed around the pulsar PSR J1719--1438. This planet has a mass similar to Jupiter but a radius of less than about 40\% \citep{Bailes2011}. It is thought to be an ultra--low mass white dwarf companion that has narrowly avoided complete destruction \citep{VanHaaften2012}. Finally, a planet has been detected around the pulsar PSR B1620-26 \citep{Backer1993,Sigurdsson2003}. This pulsar is a part of a binary star system with a white dwarf and the planet is in a circumbinary orbit. The most likely formation mechanism for this pulsar is that a star and planet were captured by the pulsar, whose original companion was ejected into space and lost \citep{Rasio1994,Ford2000pp}. Thus, this planet probably did not form around the pulsar. All of the pulsar planets found so far are around old millisecond pulsars (MSPs) that are thought to have been spun up by accretion of matter from a companion star \citep[e.g.][]{Alpar1982,Bhattacharya1991}. \begin{table} \caption{Mass and semi--major axis of planets observed around pulsars} \centering \begin{tabular}{lllcc} \hline \hline Planet & $M_{\rm planet}$ & $a$/AU \\ \hline PSR B1257+12 A & $0.02\,\rm M_\oplus$ & 0.19 \\ PSR B1257+12 B & $4.3\,\rm M_\oplus$ & 0.36 \\ PSR B1257+12 C & $3.9\,\rm M_\oplus$ & 0.46 \\ \hline PSR J1719-1438 b & $1\,\rm M_{\rm J}$ & 0.004 \\ \hline PSR B1620--26 b & $2.5\,\rm M_{\rm J}$ & 23 \\ \hline \end{tabular} \label{tab} \end{table} \begin{figure} \epsscale{1.0} \begin{center} \includegraphics[width=8.4cm]{timing2.eps} \end{center} \caption{Pulsar companion mass and semi--major axis. The mass is the median mass assuming an orbital inclination of $i=60^\circ$. The assumed primary pulsar mass is $1.35\,\rm M_{\odot}$. The companions are planets (blue points), main--sequence stars (red points), ultra low-mass (black points), neutron stars (green points), CO or ONeMg white dwarfs (purple points) or helium white dwarfs (yellow points). The blue stars show the 8 planets in the solar system, the Moon, and the asteroid Ceres. The solid lines show the detection limit for timing residuals of $1 \, \rm \mu s$ (lower line, applicable for millisecond pulsars) and $1\,\rm ms$ (upper line, applicable for normal slow pulsars) found with equation~(\ref{mass}). \label{Mass_semi}} \end{figure} The precision of the pulsar timing allows detections of very low mass bodies outside of the solar system \citep{Wolszczan1994, Wolszczan1997}. The lower limit on the mass of an observable planet is \begin{equation} M_{\rm planet} \, \sin(i) \approx 0.90 \, \left(\frac{\tau_{\rm pl}}{1\,\rm ms} \right)\, \left(\frac{a}{1\,\rm AU}\right)^{-1}\,\rm M_\oplus \label{mass} \end{equation} \citep{Wolszczan1997}, where $\rm M_{\oplus}$ is the mass of the Earth, $\tau_{\rm pl}$ is the measured timing residual amplitude, $a$ is the orbital semi--major axis and $i$ is the orbital inclination \citep[e.g.][]{Thorsett1992,Blandford1993,Bailes1993,Wolszczan1997}. The pulsar is assumed to have a mass $M_{\rm p}=1.35\,\rm M_\odot$. We show the detectability limits in Figure \ref{Mass_semi} for timing residuals of $1\,\rm \mu s$ (applicable to MSPs) and $1\,\rm ms$ (applicable to normal pulsars) and compare to objects in the solar system. Jovian planets are easily detectable around a normal slow pulsar, whereas terrestrial planets, down to the size of large asteroids, can be found in the timing residual of a MSP. Despite the successes in finding planets around pulsars and the precision offered, pulsar planets are rare \citep[e.g.][]{Bailes1996,Bell1997pp}. In particular, the survey of \cite{Kerr2015} of 151 young pulsars revealed no planets. Similarly, the Australian Telescope National Facility (ATNF) pulsar catalogue\footnote{http://www.atnf.csiro.au/people/pulsar/psrcat/} currently contains $2536$ pulsars of which 436 are MSPs with spin period $<10\,\rm ms$ \citep{ATNF2005}. Less than 1\% of observed MSPs have a planetary mass companion. In Section~2 we discuss different potential pulsar planet formation models and find that the most likely place for a pulsar planet to form is in a disk formed by the destruction of a binary companion. In Section~3 we discuss the probability for the formation of an extreme mass ratio binary required for pulsar planet formation. In addition, most planet formation scenarios require the presence of a ``dead zone'' in the protoplanetary disk. This is a quiescent region where the magneto--rotational instability (MRI) is unable to drive turbulence and angular momentum transport \citep[e.g.][]{Gammie1996,Gammie1998,Currie2007}. The dead zone is thought to be a necessary component for planet formation for several reasons. First, it provides a quiescent region where solids can settle to the midplane \citep[e.g.][]{Youdin2002, Youdin2007,Zsom2011}. Second, it allows planetesimals to form \citep[e.g.][]{Chambers2010,Bai2010,Youdin2011}. Finally, it slows the rate of low--mass planet migration, preventing them from falling into the star \citep[e.g.][]{Ward1997,Thommes2005,Matsumura2006}. In Section~4 we consider whether a dead zone can form in a disk around a pulsar. We discuss and summarize our results in Sections~5 and~6. | We explain the small number of observed pulsar planets through a combination of two low--probability events. First, the most likely formation site is in a disk formed by the destruction of a companion star. This also explains why all the pulsar planets that we have found are around millisecond pulsars rather than young pulsars. A binary that allows for such a scenario must have an extreme mass ratio and has a very small chance of forming and surviving the supernova explosion. Second, planet formation is thought to require a dead zone, a region of low turbulence, within the protoplanetary disk. Because pulsars have a much stronger source of irradiation than a young star, the additional heating can lead to sufficient ionization that the dead zone does not form. The inner boundary of the dead zone is pushed farther out, due to the irradiation, into a region of the disk with lower surface density. A dead zone only forms if the surface density that is ionized by external sources is small, $\Sigma_{\rm crit}\lesssim 10\,\rm g\, cm^{-2}$. | 16 | 9 | 1609.06409 |
1609 | 1609.08622_arXiv.txt | We study the dust content of galaxies from $z=0$ to $z=9$ in semi-analytic models of galaxy formation that include new recipes to track the production and destruction of dust. We include condensation of dust in stellar ejecta, the growth of dust in the interstellar medium (ISM), the destruction of dust by supernovae and in the hot halo, and dusty winds and inflows. The rate of dust growth in the ISM depends on the metallicity and density of molecular clouds. Our fiducial model reproduces the relation between dust mass and stellar mass from $z=0$ to $z=7$, the number density of galaxies with dust masses less than $10^{8.3}\,\rm{M}_\odot$, and the cosmic density of dust at $z=0$. The model accounts for the double power-law trend between dust-to-gas (DTG) ratio and gas-phase metallicity of local galaxies and the relation between DTG ratio and stellar mass. The dominant mode of dust formation is dust growth in the ISM, except for galaxies with $M_*<10^7\,\rm{M}_\odot$, where condensation of dust in supernova ejecta dominates. The dust-to-metal ratio of galaxies depends on the gas-phase metallicity, unlike what is typically assumed in cosmological simulations. Model variants including higher condensation efficiencies, a fixed timescale for dust growth in the ISM, or no growth at all reproduce some of the observed constraints, but fail to simultaneously reproduce the shape of dust scaling relations and the dust mass of high-redshift galaxies. | Dust is a key ingredient in interstellar medium (ISM) and galaxy physics. For example, dust influences interstellar chemistry via surface reactions and acts as a catalyst for the formation of molecules \citep{Hollenbach1971,Mathis1990,Li2001,Draine2003}. Dust depletes metals from the gas phase ISM \citep{Calzetti1994,Calzetti2000,Netzer2007,Spoon2007,Melbourne2012}. Dust grains absorb stellar radiation in the ultraviolet (UV) and re-emit this radiation in the infrared \citep[IR,][]{Spitzer1978,Draine1984,Mathis1990,Tielens2005}. Dust contributes significantly to the metals in the circumgalactic medium (CGM) and can be an additional cooling channel for gas \citep[e.g.][]{Ostriker1973,Menard2010,Peeples2014,Peek2015}. Interstellar dust is produced in the ejecta of asymptotic giant branch (AGB) stars and supernovae \citep[SNe,][]{Gehrz1989,Todini2001,Nozawa2003,Ferrarotti2006,Nozawa2007,Zhukovska2008,Nanni2013}. After the initial formation, dust growth can occur in the dense ISM via accretion of metals onto dust particles \citep{Draine1990,Dominik1997,Dwek1998,Draine2009,Hirashita2011,Zhukovska2014}. The exact contribution to the dust mass of a galaxy by the different dust formation channels is still unknown, although several authors have suggested that dust growth via accretion in the ISM plays an important role \citep[e.g.,][]{Dwek2007,Zhukovska2014,Michalowski2015,Schneider2016}. Dust can be destroyed via thermal sputtering, collisions with other dust grains, and SN shocks \citep{Dwek1980,Draine1979,McKee1989,Jones1996}. Besides the aforementioned processes, AGN can act as an additional channel for the formation of dust \citep{Elvis2002}. This dust, however, is likely to dominate only in the very central regions of the galaxies, and should not have a major impact on the total dust content of galaxies. The dust content of galaxies at low and high redshifts has intensely been studied over the past decades. Such studies provide additional constraints for galaxy formation models and the baryonic physics that regulates the dust and gas content of galaxies. These observational constraints include for instance the relation between dust mass and stellar mass \citep{Corbelli2012,Santini2014}, the gas fraction of galaxies and dust mass \citep{Cortese2012}, dust mass and star-formation rate \citep[SFR;][]{daCunha2010,Casey2012,Santini2014}, and the dust mass function of galaxies \citep{Dunne2003,Vlahakis2005,Dunne2011,Eales2009,Clemens2013}. Two particularly interesting scaling relations are the ratio between dust mass and gas mass in the ISM (dust-to-gas ratio; DTG), or the ratio between dust mass and the total mass in metals (dust-to-metal ratio; DTM) as a function of metallicity or galaxy stellar mass \citep{Issa1990,Lisenfeld1998,Hirashita2002,James2002,Hunt2005,Draine2007,Engelbracht2008,Galametz2011,Magrini2011,Remy-ruyer2014}. \citet{Remy-ruyer2014} demonstrated that the DTG ratio in galaxies cannot be described by a single power-law as a function of metallicity, but is better represented by a double power-law with a break around a metallicity of 0.1 Z$_\odot$ \citep{Edmunds2001}. Absorption line studies using gamma ray burst and Damped Lyman-alpha absorbers have suggested that the DTM ratio in galaxies at redshifts $z=0.1$ to $z=6.3$ is surprisingly similar to the DTM ratio in the local group \citep{DeCia2013,Zafar2013,Sparre2014,DeCia2016,Wiseman2016}.The DTM ratios measured in these studies drop at metallicities lower than 0.05 Z$_\odot$. These results demonstrate that high-redshift absorbers can already be significantly enriched with dust but also that the dust production efficiency can vary significantly between different environments. Far-infrared (FIR) and submillimeter observations have shown that even at the highest redshifts ($z > 4$) galaxies can have significant reservoirs of dust \citep[$10^7\,\rm{M}_\odot$ or even greater,][]{Bertoldi2003,Hughes1997,Valiante2009,Venemans2012,Casey2014,Riechers2014}. \citet{Watson2015} found a galaxy at $z=7.5\pm0.2$ with a dust mass of $4\times10^7\,\rm{M}_\odot$ and a DTG ratio that is half of the Milky Way value. Although these dusty examples may not be representative of typical high-redshift galaxies, they set strong constraints on our understanding of dust formation and growth in galaxies in the early Universe. The Atacama Large sub/Millimeter Array (ALMA) and the James Webb Space Telescope (JWST) are expected to further revolutionise our understanding of dust physics in the low- and high-redshift Universe. It is therefore becoming important to develop cosmological galaxy formation models that include dust physics, in order to provide a theoretical context for the observations with these instruments. Despite the observational prospects and theoretical importance, cosmological models of galaxy formation typically do not include self-consistent tracking of the production and destruction of dust nor dust chemistry. Traditionally, a linear scaling between dust and metal abundance is assumed \citep[e.g.,][]{Silva1998,Granato2000,Baugh2005,Lacey2008,Lacey2010,Fontanot2011,Niemi2012,Somerville2012,Hayward2013,Cowley2016}. A few groups have started to include self-consistent tracking of dust in hydrodynamic simulations \citep{Bekki2013,Bekki2015,McKinnon2015,Aoyama2017}, but these studies used zoom-simulations of individual objects and didn't focus on trends between global galaxy properties and dust mass covering a large range of parameter space and cosmic time. Recently, \citet{McKinnon2016} used a hydrodynamic model to make predictions for the dust content of galaxies in cosmological volumes, focusing on the redshift regime $z<2.5$. \citet{Dayal2011} and \citet{Mancini2016} tracked the dust content of galaxies in cosmological simulations to look at the dust absorption properties of galaxies at redshifts $z>5$. Most implementations of dust chemistry in galaxy formation have been made using specialised models \citep[e.g.][]{Dwek1998,Hirashita2002,Inoue2003,Morgan2003,Calura2008,Zhukovska2008,Valiante2009,Asano2013,Calura2014,Zhukovska2014,Feldmann2015}. % These models have been essential for developing our understanding of the relevance of the individual channels of dust formation to the dust content of galaxies. However, these models are often idealised to reproduce specific objects and are not placed within a cosmological context. Furthermore, they generally do not include all physical processes thought to be relevant for galaxy formation. Semi-analytic models (SAMs) offer a good alternative approach for self-consistently tracking the production and destruction of dust in galaxies within the framework of a $\Lambda$ cold dark matter cosmology. Simplified but physically motivated recipes are used to track physical processes such as the cooling of hot gas into galaxies, star formation, the energy input from supernovae and active galactic nuclei into the ISM, the sizes of galaxy discs, and the enrichment of the ISM by supernovae ejecta and stellar winds \citep[see][for a recent review]{SomervilleDave2015}. The low computational cost of SAMs makes them a powerful tool to model a broad range of galaxy masses probing large volumes, provide predictions for future studies, and explore different recipes for physical processes in galaxies. In this paper, we include tracking of dust production and destruction in the most recent version of the Santa Cruz semi-analytic model \citep{Popping2014sam,Somerville2015}. We explore how the dust content of galaxies and our Universe evolves over time and how this is affected by different implementations of the processes that produce dust. We extend the \citet{Arrigoni2010} galactic chemical evolution (GCE) model to include the condensation of dust in stellar ejecta, the growth of dust in the dense ISM, the destruction of dust through thermal sputtering by supernovae (SNe), dusty winds from star-forming regions, dust destruction in the hot halos, and the infall of dust from the CGM. % In this work we only focus on the evolution of dust masses and the different dust formation channels, leaving the rest of the underlying galaxy properties unchanged from the models published in \citet{Popping2014sam} and \citet{Somerville2015}. In a future work we will extend this model by including a self-consistent treatment of the impact of dust on the galaxy formation physics (i.e., cooling through dust channels, \h2 formation recipes based on the dust abundance, and dust absorption based on the estimated dust abundance). This paper is structured as follows. In Section \ref{sec:model} we present the galaxy formation model and GCE used in this work. We present the newly implemented dust related processes in Section \ref{sec:dust_model}. We briefly summarise how observational estimates of dust masses in galaxies are typically obtained, and also discuss the uncertainties on these estimates in Section \ref{sec:observing_dust}. In Section \ref{sec:results} we present our predictions for the dust scaling relations in galaxies and how these evolve with cosmic time. We discuss our finding in Section \ref{sec:discussion} and summarise our work in Section \ref{sec:summary}. Throughout this paper we adopt a flat $\Lambda$CDM cosmology with $\Omega_0=0.28$, $\Omega_\Lambda = 0.72$, $h=H_0/(100\,\rm{km}\,\rm{s}^{-1}\,\rm{Mpc}^{-1}) = 0.7$, $\sigma_8=0.812$, and a cosmic baryon fraction of $f_b=0.1658$ \citep{Komatsu2009}. | \label{sec:summary} We have included the tracking of dust production and destruction in a semi-analytic model of galaxy formation and made predictions for the dust properties of galaxies from $z=9$ to $z=0$. We present results for different model variants for the dust production processes. The first is our fiducial model with dust condensation efficiencies in stellar ejecta of around 15 per cent and a density and metallicity dependent timescale for the accretion of metals onto dust grains. The second includes no accretion of metals onto dust grains, the third assumes much higher dust condensation efficiencies than in our fiducial model, whereas the fourth assumes a fixed accretion time scale of 100 Myr. We also explored two model variants that include the approach presented by \citet{Dwek1998} for the growth of dust in the ISM. We summarise our main findings below. \begin{itemize} \item Our fiducial model successfully reproduces the trends between stellar mass and dust mass in the local and high-redshift Universe, as well as the DTG ratio of local galaxies as a function of their stellar mass. It furthermore accounts for a double power law relation between DTG ratio and gas-phase metallicity, reproduces the dust mass function of galaxies with dust masses less than $10^{8.3}\,\rm{M}_\odot$, and the cosmic density of dust at $z=0$. \item The fiducial model has problems accounting for the slope and exact normalization of the DTG and DTM ratio of galaxies as a function of their gas-phase metallicity. \item The dust mass of galaxies at fixed stellar mass is almost constant from $z=2$ to $z=0$. It decreases by $\approx 0.2$ dex from $z=3$ to $z=0$. This is mainly driven by a decrease in galaxy gas fractions. At higher redshift the relation between stellar mass and dust mass remains constant with time. The dust mass function of galaxies on the other hand increases rapidly from $z=9$ to $z=3$, after which only the number density of the galaxies with largest dust masses ($10^8\,\rm{M}_\odot$) keeps increasing. \item The relation between the DTG ratio of galaxies and their gas-phase metallicity remains constant to within 50 \% up to $z=9$. There is no clear evolutionary trend in this relation. The DTG ratio of galaxies increases with cosmic time at fixed stellar mass, following the buildup of metals in a galaxy's ISM. \item Our model predicts a significant reservoir of dust in the CGM (hot halo) of galaxies. These reservoirs can be as large or even larger than the reservoir of dust in the ISM of the host galaxy. Our models predict that even more dust is ejected from galaxies. The amount of dust in the ejected reservoir is significantly larger than observational constraints on the cosmic density of dust, suggesting that additional processes are needed to destroy this dust. \item In our models, up to 25 \% of the gas-phase metals at redshift $z=0$ can be depleted onto dust. This lowers the gas-phase metallicity relation by $\sim0.1$ dex. This depletion should be taken into account when comparing model predictions to observations. Similarly, a significant fraction of the CGM metals may be locked up in dust. \item Within our fiducial model the accretion of metals onto dust grains is the dominant mode of dust formation in galaxies. The contribution from metal accretion becomes increasingly important with stellar mass. Only at the lowest stellar masses (less than $10^7\,\rm{M}_\odot$) does the condensation of dust in SN ejecta become the dominant mode of dust formation. \item The `high-cond' and `fix-tau' model variants cannot reproduce the DTG ratio of galaxies with metallicities less than 0.2 Z$_\odot$. Furthermore the `fix-tau' model cannot reproduce the high dust masses observed in galaxies at $z\sim6$. A model without accretion of metals onto dust grains can reproduce observations relatively well if an unrealistically high efficiency of 100\% is assumed for the condensation of dust in stellar ejecta, but predicts too high dust masses in low-metallicity galaxies. The model variants that include the \citet{Dwek1998} approach for the accretion of metals onto dust grains predict either too high dust masses in low-mass galaxies ('dwek98') or too large dust masses in high-redshift galaxies ('dwek-evol'). We conclude that a model in which the rate of accretion of metals onto dust grains is set by the metallicity \emph{and} the density of the cold gas is necessary to reproduce the shape of the observed scaling relations and dust mass budgets. \end{itemize} The results presented in this paper can serve as predictions for future surveys of the dust content of galaxies. We look forward to observations from current and upcoming facilities such as ALMA, the JWST, and single-dish sub-mm instruments, that will be able to confront our predictions. In this work we have ignored the effects dust has on the ISM physics and chemistry in galaxies, affecting the growth rate of molecular hydrogen, and the absorption of stellar light. In future work we will explore these effects, and their consequences for the stellar buildup and appearance of galaxies. | 16 | 9 | 1609.08622 |
1609 | 1609.06315_arXiv.txt | text{ The Gaia team has applied a renormalization to their internally-derived parallax errors $\sigma_{\rm int}(\pi)$ $$ \sigma_{\rm tgas}(\pi) = \sqrt{[A\sigma_{\rm int}(\pi)]^2 + \sigma_0^2} \qquad (A,\sigma_0) = (1.4,0.20\,\mas) $$ based on comparison to Hipparcos astrometry. We use a completely independent method based on the RR Lyrae $K$-band period-luminosity relation to derive a substantially different result, with smaller ultimate errors $$ (A,\sigma_0) = (1.1,0.12\,\mas) \qquad (\rm this\ paper). $$ We argue that our estimate is likely to be more accurate and therefore that the reported TGAS parallax errors should be reduced according to the prescription: $$ \sigma_{\rm true}(\pi) = \sqrt{(0.79\sigma_{\rm tgas}(\pi))^2 - (0.10\,\mas)^2}. $$ } \begin{document} \jkashead % | } The Tycho-Gaia Astrometric Survey (TGAS) has just been released with approximately 2 million parallaxes, having typical reported precisions of $\sigma(\pi)\sim 300\,\muas$. Thus, while constituting only a tiny fraction of the ultimate {\it Gaia} product, TGAS is by far the largest and (with the exception of a tiny handful of {\it Hubble Space Telescope} parallaxes, e.g., \citealt{benedict11}), the most accurate optical astrometric catalog now available \citep{gaia1,gaia2}. In order to validate the TGAS catalog, it is natural to compare with the best previously existing astrometric catalog, {\it Hipparcos}. This is not ideal as TGAS has, overall, significantly better parallax measurement compared to {\it Hipparcos}. However, it is feasible, in principle, because of the large number $(\sim 10^5)$ of overlapping entries and because such validation requires only the measurement of two error-renormalization parameters $(A,\sigma_0)$ \begin{equation} \sigma_{\rm tgas}(\pi) = \sqrt{(A\sigma_{\rm int}(\pi))^2 + \sigma_0^2}; \quad (A,\sigma_0) = (1.4,0.20\,\mas) \label{eqn:renorm} \end{equation} where $\sigma_{\rm int}$ and $\sigma_{\rm tgas}$ are respectively the internal and reported (renormalized) errors, $\sigma_0$ is the systematic error floor and $A$ is the renormalization factor. However, it is notoriously difficult to calibrate superior data from inferior data and, in particular, requires superb knowledge of the error-structure of the inferior data set. We therefore present an alternative method for calibrating the TGAS parallaxes, which does not require any external astrometric data. | \label{sec:discuss}} We provide an independent method to calibrate the TGAS catalog errors that does not rely in any way on previous generation astrometric data. This purely photometric method of the standard candles avoids the issues associated with relying on lower quality data and ultimately yields more precise values for the TGAS catalog. The method presented here is likely to be more accurate than one based on comparison to {\it Hipparcos} parallaxes. This is particularly true for measuring the zero-point floor $\sigma_0$, for which we find $\sigma_0=0.12\,\mas$ and the TGAS team found $\sigma_0=0.20\,\mas$. The {\it Hipparcos} catalog very likely has systematic errors at the $\sim 0.1\,\mas$ level. This is well below the typical statistical error for individual stars, so the only place that these systematics have surfaced is in the measurement of the distance to the Pleiades, which was based on combining measurements of many stars. The original {\it Hipparcos} parallax was larger than the ``traditional'' value by almost 1 mas, and it was suggested at that time that this could be due to correlated errors \citep{pin98,sod98}. \citet{ng99} demonstrated strong evidence for correlated {\it Hipparcos} errors in the Hyades field which, by chance, they showed did not lead to an error in the distance estimate. However, after reanalyzing the {\it Hipparcos} data, \citet{van09} could find no internal evidence of these correlations and published a similar Pleiades distance measurement as originally with yet smaller error bars. This conflict was resolved by TGAS in favor of the ``traditional'', longer Pleiades distance \citep{gaia1}. While the distance to the Pleiades is settled, the cause of the {\it Hipparcos} error in its estimate of this distance is not. It may be that the problem is entirely explained by correlations, but these may also be masking other problems. In particular, since the cause of these correlations (if they are in fact the root cause) have not been tracked down, it cannot be assumed that the {\it Hipparcos} error profile is understood at a level well below the precision of its measurement. By contrast, the external inputs into RRL PL relation ($P$, $K$, $A_K$) are quite well understood. Therefore, the approach of the current paper appears more secure and in any case independent from the purely astrometric approach. One important difference between the sample studied here (108 RRL) and the one studied by the {\it Gaia} team ($10^5$ {\it Hipparcos} stars) is that the RRL have intrinsically similar colors while the {\it Hipparcos} stars cover a full range of stellar colors. In principle, this difference could be important since the {\it Gaia} team did not attempt to correct TGAS for color-dependent astrometric deviations (as they will for subsequent releases). This issue clearly deserves further investigation, a path to which we outline below. However, if there are color-dependent systematic errors, these are likely to be larger for RRL than average stars because RRL have a larger color offset relative to the mean reference frame set by other stars. Nevertheless, this question can be further investigated by applying the same method that we have used here, but to Cepheids. Cepheids lie in the same instability strip and so, like RRL, they are systemically bluer than other stars. However, unlike RRL they virtually all lie in the Galactic plane. Moreover, they are more luminous than RRL and so are typically seen at greater distances and so through more dust. In our sample of 108 RRL, there are only 19 stars with $E(B-V)>0.1$ and only four of these have $E(B-V)>0.2$. This is not enough to probe a broad range of observed colors. By contrast, Cepheids will probe a broad range. Note in particular that while distant Cepheids provide relatively little information about the PL relation, they can provide excellent information on TGAS error characterization. This is because, once the PL relation is determined from nearby stars, the parallaxes of distant stars can be determined photometrically to much higher precision than the parallax errors. This is the same principle used in the \citet{gk2016} method of measuring $\pi_0$. | 16 | 9 | 1609.06315 |
1609 | 1609.04851_arXiv.txt | We present results from particle-in-cell simulations of driven turbulence in magnetized, collisionless, and relativistic pair plasma. We find that fluctuations are consistent with the classical $k_\perp^{-5/3}$ magnetic energy spectrum at fluid scales and a steeper $k_\perp^{-4}$ spectrum at sub-Larmor scales, where $k_\perp$ is the wavevector perpendicular to the mean field. We demonstrate the development of a non-thermal, power-law particle energy distribution, $f(E) \sim E^{-\alpha}$, with index $\alpha$ that decreases with increasing magnetization and increases with increasing system size (relative to the characteristic Larmor radius). Our simulations indicate that turbulence can be a viable source of energetic particles in high-energy astrophysical systems, such as pulsar wind nebulae, if scalings asymptotically become insensitive to the system size. | 16 | 9 | 1609.04851 |
||
1609 | 1609.04583_arXiv.txt | {Time-dependent injection can cause non-linear cooling effects, which lead to a faster energy loss of the electrons in jets. The most obvious result is the appearance of unique breaks in the SED, which would normally be attributed to a complicated electron distribution. The knowledge of the observation time and duration is important to interpret the observed spectra, because of the non-trivial evolution of the SED. Intrinsic gamma-gamma absorption processes in the emission region are only of minor importance.} \keyword{BL Lacertae objets; jets; $\gamma$-rays; relativistic processes} \conferencetitle{Blazars through sharp multiwavelength eyes} \begin{document} | Blazars are active galaxies, where the line-of-sight is closely aligned with the relativistic jet. Due to the strong Doppler boost of the produced radiation, the physics of the jet can be probed with great precision. The standard model invokes emission by highly relativistic electrons (in the frame of the jet) interacting with the magnetic field and ambient photon fields, such as the synchrotron photons (the synchrotron-self Compton (SSC) process), or external photon fields originating from the accretion disk, the broad-line region, the dusty torus, or the CMB. A review is given by, e.g., \cite{b07}. The observation of minute-short flares in blazars \cite{aea07,Aea07} challenges the standard one-zone model, which is usually invoked to explain blazar emission. The variability time scale implies either an extremely small emission region, which is in contradiction with the high luminosities recorded from these flares, or an extremely high Doppler factor, which is in contradiction to radio observations of moving knots in the jets of blazars \cite{hs06}. Several models have been developed to overcome these problems. These include the jet-in-a-jet model \cite{gea09}, the similar mini-jets-in-a-jet model \cite{bg12}, jet-star interactions \cite{bea12}, and others. These have in common that they invoke an emission region, which is smaller, denser, and faster than the surrounding jet material. Such features could develop from magnetic reconnection events within the jet. {Reconnection events might result in a time-dependent injection of particles into the radiation zone.} This changes the cooling of the electrons beyond the standard behavior. In fact, sources can enter a parameter space where the electron cooling becomes collective causing nonlinearities and time-dependencies. { Here, an analytical description of this effect is given, which} has profound implications for the resulting spectral energy distributions (SEDs) and lightcurves. Since internal absorption of photons is similar to the standard model, time-dependent injection should be considered as an important addition to the models explaining rapid variability in blazars. | We have discussed the effect of time-dependent injection on the cooling process and the influence on the observed spectra and lightcurves. Due to the time-dependent injection the electron distribution cannot reach the equilibrium state. Thus, the nonlinear nature of the SSC cooling becomes important resulting in a time-dependent cooling behavior. This has two major implications. Firstly, the cooling acts quicker than in purely linear scenarios, so that the ``half-energy'' cooling time is much reduced. Secondly, the cooling behavior changes, because the SSC cooling strength decreases with time and at some point the linear cooling takes over. { The resulting effects have been studied in the limiting case of a delta-function injection in both energy and time, which gives the opportunity to calculate analytically the entire SED and the lightcurves.} The SED is strongly influences by this change in cooling behavior, because it causes breaks in the SED, which are stronger than the usual cooling break. It is important to note that these breaks do not depend on the injected electron distribution, and therefore do not require any complicated electron distribution often invoked to explain such strong SED breaks. Furthermore, in a time-dependent scenario, the SED depends strongly on the time of observation. Therefore, only strictly simultaneous data should be compared with a single model curve. Otherwise the properties of the source might have changed dramatically. The internal photon-photon absorption is not very important, unless the source enters a rather extreme parameter space. Due to the time-dependency of the intensity distribution, the optical thickness of the source is also time-dependent, and mostly affects intermediate Compton energies located roughly around the MeV domain. The lightcurves are strongly influenced by the photon retardation at early times washing out the effects of the fast nonlinear cooling. Only at low energies, which evolve slower than the high energies, the effect of the nonlinear cooling becomes visible. The retardation effect implies that the nonlinear model faces the same difficulties as the usual one-zone model, namely that the variability timescale is given by the lightcrossing time. In conclusion, time-dependent injection and, thus, time-dependent cooling gives very interesting results for both the SED and the lightcurves. { Especially, the large-$\alpha$ case, where SSC cooling dominates initially, should be considered for modeling rapid flaring events in blazars. It could explain the change in the Compton dominance observed during flares of some sources \citep[e.g.,][]{aea09}.} \vspace{6pt} | 16 | 9 | 1609.04583 |
1609 | 1609.07255_arXiv.txt | {As part of the data processing for Gaia Data Release~1 (Gaia DR1) a special astrometric solution was computed, the so-called auxiliary quasar solution. This gives positions for selected extragalactic objects, including radio sources in the second realisation of the International Celestial Reference Frame (ICRF2) that have optical counterparts bright enough to be observed with Gaia. A subset of these positions was used to align the positional reference frame of Gaia DR1 with the ICRF2. Although the auxiliary quasar solution was important for internal validation and calibration purposes, the resulting positions are in general not published in Gaia DR1.} {We describe the properties of the Gaia auxiliary quasar solution for a subset of sources matched to ICRF2, and compare their optical and radio positions at the sub-mas level.} {Descriptive statistics are used to characterise the optical data for the ICRF sources and the optical--radio differences. The most discrepant cases are examined using on-line resources to find possible alternative explanations than a physical optical--radio offset of the quasars.} {2191 sources in the auxiliary quasar solution have good optical positions matched to ICRF2 sources with high probability. Their formal standard errors are better than 0.76~milliarcsec (mas) for 50\% of the sources and better than 3.35~mas for 90\%. Optical magnitudes are obtained in Gaia's unfiltered photometric $G$ band. The Gaia results for these sources are given as a separate table in Gaia DR1. The comparison with the radio positions of the defining sources shows no systematic differences larger than a few tenths of a mas. The fraction of questionable solutions, not readily accounted for by the statistics, is less than 6\%. Normalised differences have extended tails requiring case-by-case investigations for around 100 sources, but we have not seen any difference indisputably linked to an optical--radio offset in the sources. } {With less than a quarter of the data expected from the nominal mission it has been possible to obtain positions at the sub-mas level for most of the ICRF sources having an optical counterpart brighter than 20.5 mag.} | This paper presents and discusses the first Gaia astrometric solution for the optical counterparts of radio sources in the second realisation of the International Celestial Reference Frame (ICRF2). This is a complementary paper to the general presentation in \citet{2016GaiaL} of the astrometric solutions for Gaia Data Release~1 (Gaia DR1). The main goal of the present paper is to provide the detection statistics for the ICRF2 subset, including photometric data in Gaia's unfiltered $G$ magnitude band. We also discuss and validate the astrometric accuracy through a straight comparison first with the defining sources of the ICRF2, and then including the less accurate non-defining sources. We examine a few individual problem sources for possible systematic optical--radio offsets. When statistically significant differences are found between the optical and radio positions we attempt to trace the root cause in the observations or data processing, such as the presence of an extended host galaxy or the match to a nearby star brighter than the optical counterpart, rather than in the sources themselves. The paper starts with a presentation of the data used in this investigation and how the ICRF2 sources are matched to Gaia sources. We then discuss the overall properties of the Gaia solution for 2191 sources considered as optical counterparts of ICRF2 quasars. In the subsequent two sections we discuss the comparison of the Gaia positions to the radio ones respectively for the 94\% of the sources with good agreement and for the more troublesome cases (137 sources) where the distance between the two solutions exceeds 10~milliarcsec (mas). | As part of Gaia DR1 we present the optical positions of 2191 Gaia sources matched to ICRF2 sources, including 262 of the defining sources in ICRF2. These positions, which come from the special auxiliary quasar solution, are shown to be more accurate than the positions from the secondary solution given elsewhere in Gaia DR1. Magnitudes in Gaia's $G$ band are given for 2152 of the sources (260 defining). The properties of the optical data are discussed and detailed comparisons made between the optical and radio positions of the ICRF sources. The main conclusions are: \begin{itemize} \item The $G$ magnitudes span a range from 12.4 to 21.0~mag with the bulk of sources between 17 and 20~mag. \item The formal accuracy of the optical positions has a floor at $\sim\,$0.25~mas for $G<17$~mag, gradually increasing to a few mas at $G=20$. There is no systematic difference between defining and non-defining ICRF sources in terms of their optical accuracies versus magnitude. \item The overall agreement between the optical and radio positions is excellent: the angular separation is $<\,$1~mas for 44\% of the sources and $<\,$10~mas for 94\% of the sources. For the defining sources the corresponding numbers are 71\% and 98\%. \item Analysis of the large-scale systematic differences between the optical and radio positions in terms of VSH reveals no significant components except for a glide of amplitude $\sim\,$0.15~mas. \item The angular separations are in general consistent with the combined formal uncertainties in ICRF2 and the Gaia data, supporting the claimed accuracies. The uncertainties of the radio positions of VCS-only sources in ICRF2 may be overestimated. \item For most of the 6\% sources with angular separations above 10~mas the optical--radio offsets are consistent with the stated formal uncertainties of the data, but for a quarter of them the offsets are statistically significant. Individual examination of a number of these cases show that a likely explanation for the offset can often be found, for example in the form of a bright host galaxy or nearby star. \item Among the sources with good optical and radio astrometry we found no indication of physical optical--radio offsets exceeding a few tens of mas. For most sources the true offsets are likely to be less than 1~mas. \end{itemize} The last result is very encouraging for the future alignment of the very accurate optical reference frame to be built from Gaia observations and the corresponding radio frame, using common sources in two very different wavelength domains. | 16 | 9 | 1609.07255 |
1609 | 1609.00196_arXiv.txt | % {According to most stellar dynamo theories, differential rotation (DR) plays a crucial role for the generation of toroidal magnetic fields. Numerical models predict surface differential rotation to be anti-solar for rapidly-rotating giant stars, i.e., their surface angular velocity could increase with stellar latitude. However, surface differential rotation has been derived only for a handful of individual giant stars to date. } {The spotted surface of the K-giant KU\,Pegasi is investigated in order to detect its time evolution and quantify surface differential rotation.} {We present altogether 11 Doppler images from spectroscopic data collected with the robotic telescope STELLA between 2006--2011. All maps are obtained with the surface reconstruction code \emph{iMap}. Differential rotation is extracted from these images by detecting systematic (latitude-dependent) spot displacements. We apply a cross-correlation technique to find the best differential rotation law.} {The surface of \kupt shows cool spots at all latitudes and one persistent warm spot at high latitude. A small cool polar spot exists for most but not all of the epochs. Re-identification of spots in at least two consecutive maps is mostly possible only at mid and high latitudes and thus restricts the differential-rotation determination mainly to these latitudes. Our cross-correlation analysis reveals solar-like differential rotation with a surface shear of $\alpha=+0.040\pm0.006$, i.e., approximately five times weaker than on the Sun. We also derive a more accurate and consistent set of stellar parameters for \kupt including a small Li abundance of ten times less than solar.} {} | Quantifying differential surface rotation has proven difficult even for the Sun. Stellar observations are even more demanding and correspondingly ambiguous are the results. However, quantitative detections are now possible for those stars where we are able to spatially resolve the stellar disk by means of Doppler imaging. Such observations \citep[e.g.,][etc.]{2007A&A...476..881K,2007AN....328.1075W,2013A&A...551A...2K,2015A&A...573A..98K} as well as theoretical considerations \citep{2004AN....325..496K,2012AN....333.1028K} imply that stellar surface rotation could probably be more complex for evolved stars when compared to main sequence stars like the Sun. Differential rotation for main-sequence and pre-main sequence stars is found to decrease with effective temperature \citep{2005MNRAS.357L...1B, 2013A&A...560A...4R} just like predicted from mean-field dynamo models \citep{2015csss...18..535K}. However, the situation seems to be less well defined for post-main sequence stars with their much deeper convective envelopes. Because many of these giants are components in RS\,CVn-type binary systems, their surfaces are possibly distorted by and respond to the orbital dynamics. Moreover, as a star evolves up the red-giant branch, its core experiences a modification of nuclear reactions followed by core contraction and envelope expansion long before helium burning sets in. After the core hydrogen is exhausted, the hydrogen fusion keeps going in a surrounding shell, providing more helium onto the contracting inert core. The contraction heats up the core together with the interlocked shell, which expands inward. At a point the core becomes degenerate. The increasing density at the bottom of the H-rich shell yields a more efficient H-burning, which eventually blows up the envelope. The temperature of the envelope decreases and the outer layers become fully convective, transporting more flux outwards, which explains the rapidly increasing luminosity with decreasing surface temperature along the RGB. The shell material penetrates into the hotter regions below, decaying light elements and triggering a mixing process called the first dredge up, which is responsible for the dilution of the lithium. Indeed, according to \citet{1994A&A...282..811C,1995ApJ...453L..41C} in low mass ($\le2M_{\odot}$) stars, further rotationally induced mixing occurs after the completion of the first dredge up \citep[see also][]{1992A&A...265..115Z}. Such mixing episodes can be inferred from the lowering of the observed surface abundances of the most fragile elements ($^{7}$Li, $^{12}$C) and the $^{12}$C/$^{13}$C isotopic ratio \citep{2000A&A...354..169G}. Even the simple expansion appears to have an effect on the mixing of the convective envelope and eventually also alters the surface DR profile as well. In some cases, the DR profile can be even of anti-solar type, i.e., the equator rotating slower than the poles \citep[][etc.]{1991LNP...380..297V,2003A&A...408.1103S,2015A&A...573A..98K}. Whether such anti-solar DR was already present during the main sequence phase of such a star or explicitly developed during the expansion phase on the giant branch is not known. Besides, DR of either solar or anti-solar has been derived only for a handful of late-type evolved stars to date. A short list of the giant stars with known DR from Doppler imaging would include the following ones. Among the single (or effectively single) giants, solar type DR was reported for FK\,Com \citep{2007A&A...476..881K}, V390\,Aur \citep{2012A&A...541A..44K} and KU\,Peg \citep{2001A&A...373..974W}, i.e., the star to be revisited in this paper, while anti-solar DR was detected on HD\,31993 \citep{2003A&A...408.1103S}, DI\,Psc and DP\,CVn \citep{2013A&A...551A...2K,2014A&A...571A..74K}. In binary systems solar type DR was found e.g., on the evolved components of $\zeta$\,And \citep{2012A&A...539A..50K}, XX\,Tri \citep{2015A&A...578A.101K}, and IL\,Hya \citep{2014IAUS..302..379K}, while anti-solar DR was found on $\sigma$\,Gem \citep{2015A&A...573A..98K}, IM\,Peg, UZ\,Lib, etc. \citep[see][and the references therein]{2014SSRv..186..457K}. Our foremost aim is to enlarge the observational sample of reliable DR detections on giant stars. Time-series Doppler imaging has proven to be extremely useful for studying stellar DR \citep[e.g.][]{1996IAUS..176..245V,1997MNRAS.291....1D,1998A&A...330.1029W,2004AN....325..221P}. When having subsequent Doppler reconstructions of the spotted stellar surface, the rotation rates of individual spots can reveal the latitude-dependent stellar rotation profile. However, such Doppler reconstructions require high-resolution spectroscopic time-series data, covering at least two but better many consecutive rotation cycles. That this is indeed a challenge for stars with rotation periods of close to a month is obvious. A unique possibility for such long-term Doppler observations \citep[see, e.g.,][etc.]{2015A&A...573A..98K,2015A&A...574A..31S,2015A&A...578A.101K} is provided by the STELLA robotic observatory of the AIP in Tenerife \citep{2010AdAst2010E..19S}. In this paper we present and analyze such spectroscopic observations of the rapidly-rotating % ($P_{\rm rot}\approx24$\,days) K-giant \object{KU~Peg} (=HD\,218153). Chromospheric activity of \kupt was recognized by \citet{1983AJ.....88.1182B} who reported strong Ca\,{\sc ii}\,H\&K emission. The large chromospheric fluxes were later confirmed with IUE observations by \citet{1992A&A...254L..36D}. Just recently, \citet{2015A&A...574A..90A} detected magnetic fields on KU\,Peg and found that the star follows the magnetic field strength-rotation relationship established for active giants, indicating that probably a solar type magnetic dynamo was working inside. In addition, the authors reported an unusually strong X-ray luminosity of $L_X=11.8\times10^{30}$\,erg\,s$^{-1}$ confirming the existence of coronal activity as well. KU\,Peg was found to be a single-lined spectroscopic binary with an orbital period of $\approx$1400\,days \citep{1992A&A...254L..36D}, suggesting that it is \emph{effectively} a single star. Because differential rotation is supposed to be weakened (or totally quenched) by tidal forces in close binaries \citep{1981ApJ...246..292S,1982ApJ...253..298S}, \kupt is a good candidate for a comparison with theory. A projected rotational velocity of 29\,\kms\ was measured by \citet{1997PASP..109..514F}, which placed the star among the possible Doppler-imaging candidates \citep{2000A&AS..142..275S}. The first and so far only Doppler-imaging study of \kupt was carried out by \citet[][hereafter Paper~I]{2001A&A...373..974W} using high-resolution spectra taken with the McMath-Pierce solar telescope and the coud\'e feed telescope at Kitt Peak National Observatory over two months in 1996/97. The data allowed the reconstruction of two consecutive Doppler images that revealed an asymmetric polar spot and several other cool spots at lower latitudes. The time evolution of the spotted surface was followed by means of a cross-correlation analysis and revealed a complex DR profile that resembled the solar case only in the directional sense, i.e., lower latitudes rotating faster. Its lap time was twice as long as that of the Sun for the full pole-to-equator range but twice as short if only the latitudes where the Sun has spots was considered. Moreover, patterns of local meridional flows were detected, which likely play also an important role for stellar dynamos \citep{2004AN....325..496K,2011AN....332...83K}. The current paper is organized as follows. In Sect.~\ref{obs} we describe our photometric and spectroscopic observations. In Sect.~\ref{prot} photometric data from more than 18 years are employed to derive a precise average rotation period. These data are also used to search for photometric signals of surface DR. In Sect.~\ref{di}, we first redetermine the basic astrophysical properties of \kupt by including our new photometric and spectroscopic data. Then we give a brief description of our inversion code \emph{iMap} and its data assumptions for image reconstruction and, thirdly, we present the time-series Doppler images. In Sect.~\ref{ccf} the consecutive Doppler images are used to derive the surface DR of \kup. Lithium abundance determination is carried out in Sect.~\ref{lithium}. The results are summarized and discussed in Sect.~\ref{disc}. | \label{disc} Our new Doppler image reconstructions of \kupt indicate that its surface spot distribution is indeed very active and dynamic, even when compared to other overactive stars, and so must be the underlying dynamo. According to the theoretical estimation by \citet{2015A&A...574A..90A} the maximum convective turnover time of \kupt should be around $\tau\approx100$\,d. This would yield a moderate Rossby number of $Ro\approx0.24$, indicating that the star operates an $\alpha\Omega$ type dynamo. However, from the time evolution of the spotted surface, considering especially the polar spottedness between 2006-2011, we could only estimate a rough cycle length of a few years. On the other hand, from long-term photometry, such a cycle (of about 2--4 years) can be inferred at a very weak significance level, i.e., not conclusively. From our cross-correlation study, we derived a solar-like surface DR with a shear of $\alpha=+0.040\pm0.006$ and a lap time of $\approx$580~days. A similar solar-like DR was found in Paper~I with $\alpha=+0.09$ and a corresponding lap time of $\approx$260 days. Note that the higher value in Paper~I came from a cross-correlation of only two consecutive Doppler maps and the use of a less-robust correlation routine which both resulted in a less pronounced correlation pattern. A redetermination of $\alpha$ with a different cross-correlation program but the same data as in Paper~I by \citet{2005AN....326..287W}, revealed an $\alpha$ of +0.03, in agreement with our new value. In the present paper, we applied a more robust cross-correlation technique for altogether 8 cross-correlation maps and conclude that our new result is much more reliable and has now a reasonable error bar. The time-series Doppler reconstructions revealed evidence for systematic spot displacements that may be interpreted as evidence for local meridional flows. Examples are the poleward drift of the dominant feature in Fig.~\ref{2006di} at $\phi=0.25$ or the displacement of the low latitude spot in Fig.~\ref{2009di} at $\phi=0.25$ (as well as in Fig.~\ref{2011di} at $\phi=0.00$ and at $\phi=0.25$). Compared to the longitudinal displacements due to DR, the latitudinal displacements are more diffused and much weaker on average. At the same time the overall evolution of the spot distribution is generally more complex than thought and can have spots come and go from one rotation to the next (e.g. in 2011, Fig.~\ref{2011di}). Rapid rotation of an old, effectively single, evolved star like \kupt remains a challenge for theory. If tidal effects did not play a role in the past of the star, the most likely explanation for its rapid rotation would be that a deepening convective envelope eventually reaches the high angular momentum material around a fast rotating core, and thus transports high angular momentum material up to the surface on a comparably short convective time scale \citep{1979ApJ...232..531E}. This dredge up must take place before the star evolves up to the bump of the red giant branch. On the other hand, with the expanding envelope an increasing mass loss rate would be expected, which would again mean angular momentum loss. However, excessive mass loss can be excluded by the lack of any IR excess from a comparison of the measured (2MASS) $J$, $H$ and $K$ magnitudes and the color calibrations provided by \citet{2005ApJ...626..465R}. But out-flowing material can also be coupled to closed surface magnetic fields, generated by a dynamo, this way preventing the star from fast angular momentum loss \citep[cf.][]{2010ApJ...719..299C}. In any case, a rapidly rotating core on the main sequence is required to explain the spin-up by angular momentum transport from the deep. However, it is not likely that a star of $1.1\,M_{\odot}$ could provide such a fast rotating core. Moreover, according to \citet{2016A&A...591A..45P} the dredge up may not produce enough acceleration of the surface to be a reasonable explanation at all. As an alternative scenario, engulfment of one planet (or even more) may explain the rapid rotation \citep{1999MNRAS.308.1133S,2012ApJ...757..109C,2016arXiv160608027P}. Taken the expression from \citet{2008AJ....135..209M} we estimate the mass of the planet which would spin up the star to be $\approx$1.25\,$M_{\rm J}$. But this would also raise the Li abundance at the surface rather than lowering it, which is found for \kup. Note however, that according to \citet{2016arXiv160303038C} any close giant planet is likely to be engulfed well before the host star would evolve up the RGB. This would explain the low surface Li abundance, since by the end of the dredge up phase the extra Li coming from the planet will be destroyed together with the primeval Li of the stellar envelope. But anyhow, lithium can be affected by other less known processes too, thus the surface Li measurement itself can hardly account for or disprove any planetary interaction as pointed out by \citet{2016arXiv160608027P}. The position in the H-R diagram indicates that \kupt is past the RGB luminosity bump. Only low-mass stars that have a highly degenerate He core on the RGB, and later undergo the He flash, evolve through this phase \citep[see][]{2000A&A...359..563C}. At this time extra Li is produced and very high Li abundances are reached \citep[e.g., HD\,233517;][]{2015A&A...574A..31S}. However, this phase is extremely short lived because once the mixing extends deep enough the freshly synthesized Li is quickly destroyed. Immediately before the bump phase (and after the end of the first dredge-up), we expect relatively low Li abundances. The time-scale of the bump for $M=1.1 \mathrm{M}_\sun$ and $Z$=0.008 (almost equivalent with [Fe/H]=$-$0.37) is $\approx$10\,Myr, the time between the bump and the current position of \kupt is $\approx$40\,Myr according to the models of \citet{2008A&A...484..815B}. This must have been enough time to dilute KU\,Peg's surface Li to basically zero. | 16 | 9 | 1609.00196 |
1609 | 1609.07914_arXiv.txt | Some chemically peculiar stars in the upper main sequence show rotational period variations of unknown origin. We propose these variations are a consequence of the propagation of internal waves in magnetic rotating stars that lead to the torsional oscillations of the star. We simulate the magnetohydrodynamic waves and calculate resonant frequencies for two stars that show rotational variations: CU~Vir and HD~37776. We provide updated analyses of rotational period variations in these stars and compare our results with numerical models. For CU~Vir, the length of the observed rotational-period cycle, $\mathit\Pi=67.6(5)$\,yr, can be well reproduced by the models, which predict a cycle length of 51\,yr. However, for HD~37776, the observed lower limit of the cycle length, $\mathit\Pi\geq100$\,yr, is significantly longer than the numerical models predict. We conclude that torsional oscillations provide a reasonable explanation at least for the observed period variations in CU~Vir. | Chemically peculiar (CP) stars in the upper main sequence show light variability attributed to rotational modulation of magnetic spots of differing surface abundances. Radiative flux redistribution in the spots results from various bound-free \citep[ionization,][]{peter,lanko} and bound-bound atomic \citep[line,][]{vlci,ministr,molnar} transitions. Surface abundance maps derived from Doppler imaging \citep[e.g.,][]{leuma,silkow} can be used to predict the light variability in CP stars \citep[e.g.,][]{myteta,prvalis}. The brightness variability in CP stars allows measurement of their rotational periods. The strict periodicity observed in their light curves allows precise determination of their periods with typical relative uncertainties of the order $10^{-6}$\,--\,$10^{-5}$ \citep[e.g.,][]{adelm}. This facilitates the search for very minute changes in the rotational periods. For single CP stars, the usual mechanisms of period change are related to stellar evolution. Unfortunately, evolutionary changes in stellar rotation \citep{rotmod} are not detectable in main-sequence stars \citep{mikmos}. As a result, most CP stars have very constant rotational periods. However, there are exceptions. The hottest CP stars with surface magnetic fields and winds may show rotational braking as a result of angular momentum loss via the magnetized stellar wind \citep{brzdud}. This effect was discovered in the helium-rich star $\sigma$~Ori~E by \citet{town}. Period variations in HD~37776 discovered by \citet{mik901} were also attributed to angular momentum loss. However, subsequent analysis of HD~37776 by \citet{zmenper2} revealed a significant cubic term in the star's ephemeris, inconsistent with simple rotational braking. In the CP star CU~Vir, intervals of rotational braking were found to alternate with intervals of rotational acceleration by \citet{zmenper2}. These new findings in stars HD~37776 and CU~Vir need to be explained. Here we study the torsional oscillations that result from the interaction of rotation and magnetic field \citep[see][for a similar idea]{step}. Although originally introduced for other purposes (cf., \citealt{mestel}, pages 161--163, see also \citealt{mw}), we show that the torsional oscillations are able to explain the period variation in CU~Vir. | We simulated stellar torsional oscillations that result from the interaction of the internal magnetic field and differential rotation. The simulations were calculated for the chemically peculiar stars CU~Vir and HD~37776, which both have rotational period variations. We derived the internal structure of individual modes and calculated the wave resonance frequencies, for which the amplitudes of surface angular frequency variations are the largest. For each star we found a basic frequency and several high-order overtones. We provide a new analysis of period variations in the stars CU~Vir and HD~37776 assuming periodic rotational period variations. For CU~Vir, the length of the rotational period cycle $\mathit\Pi=67.6(5)$\,yr can be well reproduced by numerical models, which predict a cycle length of 51\,yr. The numerical model also predicts the variations on the scale of about 10~yr in agreement with observations. Consequently, the torsional oscillations provide a reasonable explanation of the observed period variations of CU~Vir. On the other hand, for HD~37776 the observed lower limit of the period cycle, $\mathit\Pi\geq100$\,yr, is significantly longer than the predicted cycle length of 5\,yr. It is immediately clear from the scaling of the wave equation with the magnetic field and from the observed strength of the field that the model cannot reproduce the observations of both stars. There may be other possible explanations for the observed period variations in HD~37776. | 16 | 9 | 1609.07914 |
1609 | 1609.00019_arXiv.txt | We quantitatively investigate the possibility of detecting baryonic acoustic oscillations (BAO) using single-dish 21cm intensity mapping observations in the post-reionization era. We show that the telescope beam smears out the isotropic BAO signature and, in the case of the Square Kilometer Array (SKA) instrument, makes it undetectable at redshifts $z\gtrsim1$. We however demonstrate that the BAO peak can still be detected in the radial 21cm power spectrum and describe a method to make this type of measurements. By means of numerical simulations, containing the 21cm cosmological signal as well as the most relevant Galactic and extra-Galactic foregrounds and basic instrumental effect, we quantify the precision with which the radial BAO scale can be measured in the 21cm power spectrum. We systematically investigate the signal-to-noise and the precision of the recovered BAO signal as a function of cosmic variance, instrumental noise, angular resolution and foreground contamination. We find that the expected noise levels of SKA would degrade the final BAO errors by $\sim5\%$ with respect to the cosmic-variance limited case at low redshifts, but that the effect grows up to $\sim65\%$ at $z\sim2-3$. Furthermore, we find that the radial BAO signature is robust against foreground systematics, and that the main effect is an increase of $\sim20\%$ in the final uncertainty on the standard ruler caused by the contribution of foreground residuals as well as the reduction in sky area needed to avoid high-foreground regions. We also find that it should be possible to detect the radial BAO signature with high significance in the full redshift range. We conclude that a 21cm experiment carried out by the SKA should be able to make direct measurements of the expansion rate $H(z)$ with measure the expansion with competitive per-cent level precision on redshifts $z\lesssim2.5$. | \label{sec:introduction} The spatial distribution of matter in the Universe is sensitive to the value of the cosmological parameters. Constraints on those can thus be placed by comparing the statistical properties of the density field against predictions from theoretical models. Unfortunately, the true matter density is not directly observable, and therefore one must resort to using proxies of it, such as the number density of galaxies or line emission intensity of cosmic neutral hydrogen (HI). A promising and new way of tracing the large-scale structure of the Universe is to carry out low angular resolution radio observations to detect the 21cm radiation from cosmic neutral hydrogen in the post-reionization epoch. The idea is not to detect individual galaxies through their 21cm emission, but rather to measure the combined flux in wide patches of the sky containing many galaxies. This technique is called intensity mapping \citep{Bharadwaj_2001A, Bharadwaj_2001B, Battye:2004re,McQuinn_2006, Chang_2008, Loeb_Wyithe_2008, Villaescusa-Navarro_2014a, Bull_2015}. Under the assumption that the measured 21cm flux traces the perturbations in the matter density on large-scales, we can use the clustering properties of the cosmic HI, as observed from 21cm intensity mapping surveys, to put constraints on the value of the cosmological parameters \citep{Bull_2015, Villaescusa-Navarro_2015a, Carucci_2015}. Baryonic acoustic oscillations (BAO), originated in the early Universe by the competition between the gravitational interaction and the radiation pressure of photons tightly coupled to baryons, leave an imprint in the late-time matter density in the form of a statistically preferred separation between density peaks of $r_s\sim110~h^{-1}{\rm Mpc}$, corresponding to the size of the sound horizon at the time of the baryon-photon decoupling. This translates into a distinct peak in the matter/galaxy two-point correlation function, or as a set of wiggles in the matter/galaxy power spectrum on scales $k\sim[0.05-0.3]~h{\rm Mpc}^{-1}$ with frequency $r_s$. The BAO signature thus constitutes a cosmological standard ruler, whose size depends on well understood physics of the early Universe. By measuring them in the temperature anisotropies of the cosmic microwave background (CMB) and in the clustering pattern of matter tracers, it is possible to measure the value of the Hubble rate and the angular diameter distance as a function of redshift. The main advantage of the BAO signal resides in its robustness against systematic effects: it is difficult for non-cosmological effects to mimic or shift the position of the BAO feature in the correlation function or power spectrum. Furthermore, given the large-scale nature of the BAO signal, the effects induced by the non-linear gravitational evolution are well captured by perturbation theory \cite[e.g.][]{Crocce_2008, Padmanabhan_2009, Baldauf_2015,peloso}. The BAO scale has been measured in the 2pt/3pt statistics of galaxy surveys \cite[see e.g.][]{Cole_2005, Eisenstein_2005, Anderson_2014, Gil-Marin_2015,Beutler_2016, Alam_2016,Slepian_2016}, in the Ly$\alpha$-forest \citep{Delubac_2015}, in the distribution of galaxy clusters \citep{alfonso} and in the spatial distribution of voids \citep{Kitaura_2016}. Upcoming and future radio experiments such as the Canadian Hydrogen Intensity Mapping Experiment (CHIME)\footnote{\url{http://chime.phas.ubc.ca/}}, the Ooty Radio Telescope (ORT)\footnote{\url{http://rac.ncra.tifr.res.in/}}, BINGO \cite{2012arXiv1209.1041B} and the Square Kilometre Array (SKA)\footnote{\url{https://www.skatelescope.org/}} will survey large areas of the sky using the intensity mapping technique in the post-reionization era. In this paper we investigate the prospects of detecting the BAO from 21cm intensity mapping observations, focusing on the SKA1-MID instrument. An ideal intensity mapping experiment would cover the largest possible field of view with as large angular resolution as possible. Since the angular scales probed by a radio interferometer are $\lambda/b_{\rm max}\lesssim\theta\lesssim\lambda/b_{\rm min}$, where $b_{\rm max/min}$ are the largest/smallest separation between two antenna elements, these two requirements can only be simultaneously met by building large interferometric arrays of tightly packed receivers. An alternative strategy would be to cover the desired sky footprint with single-dish observations, in which case the angular resolution has a lower bound $\theta\gtrsim\lambda/D_{\rm dish}$ determined by the dish diameter \cite[see][for a detailed discussion]{2015ApJ...803...21B}. In this paper we will focus on the latter case, the likely strategy of choice for the SKA1-MID instrument, described in \cite{braun15}. We demonstrate that the poor angular resolution inherent to single-dish 21cm observations smears out the BAO peak in the isotropic correlation function or power spectrum, and that in this case cosmological constraints would be driven by the overall shape of the 21cm power spectrum, which is more sensitive to systematic effects. We will however show that the BAO wiggles can be detected in the radial 21cm power spectrum, and thus can be used to make a direct measurement of the expansion rate $H(z)$. In our analysis we will focus on the impact of instrumental effects, such as the system noise, and the presence of Galactic and extra-Galactic foregrounds on our the results. This paper is organized as follows. In Section \ref{sec:BAO} we study the impact of the instrumental beam on the detectability of the isotropic BAO peak in single-dish experiments. In Section \ref{sec:methods} we describe the simulations and analysis methods used in this work. The results obtained from this analysis and their interpretation are presented in Section \ref{sec:results}, where we systematically investigate the impact of each complication (system noise and foregrounds) on the final uncertainties. Finally we discuss main conclusions of this paper in Section \ref{sec:conclusions}. | \label{sec:conclusions} The BAO scale is one of the most robust cosmological observables, due to the distinctive nature of their signature, a single peak on the 2pt correlation function or a set of wiggles on the power spectrum. This observable can be used to measure the Hubble function and the angular diameter distance as a function of redshift, and therefore represents a unique and robust probe to study the nature of dark energy. The purpose of this paper has been to investigate the accuracy with which the BAO scale can be determined through single-dish 21cm intensity mapping observations in the post-reionization epoch. We quote results for a possible intensity mapping experiment carried out with the SKA1-MID array covering more than half of the sky in the redshift range $z\sim[0.3-3]$. We however emphasize that our methodology is fully general and can be easily applied to other instruments. We have shown that the smearing caused by the beam size of the radio-telescopes will prevent a competitive measurement of the isotropic BAO scale in both the 21cm correlation function or power spectrum (see Fig. \ref{fig:beam_BAO}). However, we have shown that, given the good frequency resolution of radio telescopes, it should still be possible to measure the radial BAO signal down to high redshifts, thus placing competitive constraints on the expansion rate $H(z)$. In this paper we have proposed a method to recover the radial BAO scale in intensity mapping observations and implemented it in practice making use of a suite of 100 full-sky lightcone simulations in order to systematically study the effects of the instrumental noise and the robustness of the signal to foreground-related systematic effects. Our procedure follows three steps: \begin{itemize} \item A simulated sky is generated containing a realization of the full-sky HI cosmological signal as well as the most relevant Galactic and extra-Galactic foregrounds in the frequency range $\nu\in[350,1050]\,{\rm MHz}$. The simulated maps are smoothed to the angular resolution corresponding to the specifications of SKA1-MID, and white instrumental noise is added accordingly. \item We remove the foregrounds using a PCA algorithm, subtracting the first 8 principal components, which we have shown are dominated by foregrounds. \item We compute the radial power spectrum of the resulting maps by stacking the 1-dimensional Fourier transform of every pixel in the field of view along the frequency direction (further details about the method are given in Section \ref{subsec:pk}). \item For each estimated power spectrum, we determine the radial BAO scaling parameter $\alpha$ by fitting the template given in Eq. \ref{template} to the data. The mean value and uncertainty on $\alpha$ is then estimated by averaging over 100 simulations. \end{itemize} All our simulations take into account the sky area limitations of the SKA both in terms of accessible sky and Galactic foregrounds. Our results concerning the measurement of the BAO scale are summarized in Table \ref{tbl:best_fit}. We find that the BAO uncertainties become larger at both high and low redshifts, even in the absence of instrumental noise. We have shown that this is due to the low relative amplitude of the BAO signature in the radial power spectrum at low redshifts and to the lower signal-to-noise ratio of the total HI power spectrum at high redshifts caused by the larger size of the telescope beam, which overcomes the $\propto1/\sqrt{V}$ improvement factor due to the larger volume coverage. More importantly, we find that, while the effects of instrumental noise are irrelevant at low redshift, they come to dominate the error budget at redshifts $z\gtrsim2$, increasing the final BAO uncertainties by a factor of $\sim2$ with respect to the sample-variance limited result. Concerning the effect of radio foregrounds, we have shown that the large-scale bias induced by foreground removal on the radial power spectrum (as reported by e.g. \cite{2015MNRAS.447..400A}) does not cause a bias in the recovered BAO scale. This is thanks to the robustness of the BAO signal against broad-band variations in the shape of the power spectrum, as well as to the spectral separation between foregrounds and cosmological signal in the frequency direction. We have also shown that the contribution of foreground residuals to the final uncertainties is negligible, and that, therefore, the main effect of foregrounds is a reduction in the available sky area needed in order to avoid the regions of higher Galactic emission. Although we have not explicitly introduced this effect, it should be possible to mitigate the impact of correlated instrumental noise, of particular relevance to single-dish observations, using similar methods \citep{2015MNRAS.454.3240B}. Finally, we have studied the significance of the BAO signal. We have shown that, although the large variance of the instrumental noise at low frequencies reduces the signal-to-noise ratio of the BAO signature at high redshifts, we obtain significant detections ($>3\sigma$) of it in a large majority of our simulations in all redshift bins. Overall, we conclude that by a single-dish 21cm intensity mapping experiment carried out by SKA1-MID over $\sim50\%$ of the sky with an allocated observing time of 10000 hours should be able to place direct constrains the value of the Hubble function $H(z)$ with a relative uncertainty of $(2.4\%, 1.5 \%, 1.9\%, 3.1\%)$ at redshifts $\langle z\rangle=(0.6,1.0,1.6,2.5)$ by measuring the BAO scale in the radial 21cm power spectrum. This would correspond to a precision comparable with next-generation spectroscopic surveys (e.g. \citet{2014JCAP...05..023F}). | 16 | 9 | 1609.00019 |
1609 | 1609.05911_arXiv.txt | Star formation is primarily controlled by the interplay between gravity, turbulence, and magnetic fields. However, the turbulence and magnetic fields in molecular clouds near the Galactic Center may differ substantially from spiral-arm clouds. Here we determine the physical parameters of the central molecular zone (CMZ) cloud \brick, its turbulence, magnetic field and filamentary structure. Using column-density maps based on dust-continuum emission observations with ALMA+\emph{Herschel}, we identify filaments and show that at least one dense core is located along them. We measure the filament width $W_\mathrm{fil}=0.17\pm0.08\,\pc$ and the sonic scale $\lambda_\mathrm{sonic}=0.15\pm0.11\,\pc$ of the turbulence, and find $W_\mathrm{fil}\approx\lambda_\mathrm{sonic}$. A strong velocity gradient is seen in the HNCO intensity-weighted velocity maps obtained with ALMA+Mopra. The gradient is likely caused by large-scale shearing of {\brick}, producing a wide double-peaked velocity PDF. After subtracting the gradient to isolate the turbulent motions, we find a nearly Gaussian velocity PDF typical for turbulence. We measure the total and turbulent velocity dispersion, $8.8\pm0.2\,\km\,\s^{-1}$ and $3.9\pm0.1\,\km\,\s^{-1}$, respectively. Using magnetohydrodynamical turbulence simulations, we find that {\brick}'s turbulent magnetic field $B_\mathrm{turb}=130\pm50\,\mu\Gauss$ is only $\lesssim1/10$ of the ordered field component. Combining these measurements, we reconstruct the dominant turbulence driving mode in {\brick} and find a driving parameter $b=0.22\pm0.12$, indicating solenoidal (divergence-free) driving. We compare this to spiral-arm clouds, which typically have a significant compressive (curl-free) driving component ($b>0.4$). Motivated by previous reports of strong shearing motions in the CMZ, we speculate that shear causes the solenoidal driving in {\brick} and show that this reduces the star formation rate (SFR) by a factor of $6.9$ compared to typical nearby clouds. | Star formation powers the evolution of galaxies. However, the processes that control the conversion of gas into stars remain poorly understood. We now know that turbulence, magnetic fields and feedback are essential for regulating star formation in the Galactic disk, because gravity alone would produce stars at a $\sim\!100$ times higher rate than observed \citep{McKeeOstriker2007,PadoanEtAl2014,Federrath2015}. However, it is not so clear whether the same principles hold in the Central Molecular Zone---a much more extreme environment. For instance, despite the high gas densities and the large amount of available gas, there is about an order of magnitude less active star formation in the CMZ than expected \citep{LongmoreEtAl2013a,KruijssenEtAl2014,JohnstonEtAl2014}. In order to test theories of star formation, our main aim here is to measure the amount and structure of the turbulence and to determine the magnetic field. We do this for the CMZ cloud {\brick}, also known as the `Brick'. Besides constraining fundamental parameters of {\brick}, such as the density and mass of the cloud, we focus on determining the turbulent Mach number and driving, as well as the turbulent magnetic field component. We reconstruct the driving mode of the turbulence in {\brick} and find that it is primarily solenoidal. This is in stark contrast to spiral-arm clouds, where the turbulence seems to be significantly more compressive \citep{PadoanJonesNordlund1997,Brunt2010,PriceFederrathBrunt2011,GinsburgFederrathDarling2013}. The solenoidal driving of turbulence in {\brick} may provide a possible explanation for the unusually low efficiency of dense-core and star formation in this environment. Recent observations with the Atacama Large Millimeter/submillimeter Array (ALMA) have revealed that {\brick} is indeed a molecular cloud with a highly complex structure governed by turbulent motions \citep{RathborneEtAl2014,RathborneEtAl2015}. These high-resolution dust and molecular line observations indicate that {\brick} is filamentary, with networks of filaments having similar complexity as in nearby spiral-arm clouds \citep{AndreEtAl2014}. So far the filamentary structure inside {\brick} has not been quantified, because pre-ALMA observations did not have sufficient resolution. Here we measure the average filament column density and width in this CMZ cloud and compare our measurements to nearby spiral-arm clouds. \subsection{Turbulence driving?} The observations by \citet{RathborneEtAl2014,RathborneEtAl2015} demonstrate that {\brick} is highly turbulent, but it has been unclear what drives this turbulence \citep[for a discussion of potential drivers of turbulence in the CMZ, see \S5.2 in][]{KruijssenEtAl2014}. Numerical simulations have shown that turbulence decays quickly in about a crossing time \citep{ScaloPumphrey1982,MacLowEtAl1998,StoneOstrikerGammie1998,MacLow1999}. The fact that we see turbulence thus leads us to conclude that it must be driven by some physical stirring mechanism. In general, potential driving mechanisms include supernova explosions and expanding radiation fronts and shells induced by high-mass stellar feedback \citep{McKee1989,KrumholzMatznerMcKee2006,BalsaraEtAl2004,BreitschwerdtEtAl2009,PetersEtAl2011,GoldbaumEtAl2011,LeeMurrayRahman2012}, winds \citep{ArceEtAl2011}, gravitational collapse and accretion of material \citep{VazquezCantoLizano1998,KlessenHennebelle2010,ElmegreenBurkert2010,VazquezSemadeniEtAl2010,FederrathSurSchleicherBanerjeeKlessen2011,RobertsonGoldreich2012,LeeChangMurray2015}, and Galactic spiral-arm compressions of H\textsc{I} clouds turning them into molecular clouds \citep{DobbsBonnell2008,DobbsEtAl2008}, as well as magneto-rotational instability (MRI) and shear \citep{PiontekOstriker2007,TamburroEtAl2009}. Jets and outflows from young stars and their accretion disks have also been suggested to drive turbulence \citep{NormanSilk1980,MatznerMcKee2000,BanerjeeKlessenFendt2007,NakamuraLi2008,CunninghamEtAl2009,CarrollFrankBlackman2010,WangEtAl2010,CunninghamEtAl2011,PlunkettEtAl2013,PlunkettEtAl2015,OffnerArce2014,FederrathEtAl2014}. While different drivers may play a role in different environments (such as in spiral-arm clouds), \citet{KruijssenEtAl2014} found that most of these drivers are not sufficient to explain the turbulent velocity dispersions in the CMZ. Importantly, most of these turbulence drivers primarily compress the gas (e.g., supernova explosions, high-mass stellar feedback, winds, gravitational contraction, and spiral-arm shocks), but others can directly excite solenoidal motions (e.g., MRI, jets/outflows, and shear). Our goal here is to determine the fraction of solenoidal and compressive modes in the driving of the turbulence in {\brick}. This relative fraction of driving modes is determined by the \emph{turbulence driving parameter} $b$, which is proportional to the ratio of density to velocity fluctuations, $b\propto\sigma_\rho/\sigma_v$, in a supersonically turbulent cloud \citep{FederrathKlessenSchmidt2008,FederrathDuvalKlessenSchmidtMacLow2010}. \citet{FederrathKlessenSchmidt2008} showed that purely solenoidal (rotational or divergence-free) driving corresponds to $b=1/3$, while purely compressive (potential or curl-free) driving results in $b=1$. Increasing the fraction of compressive modes in the turbulence driving from zero to unity leads to a smoothly increasing driving parameter $b$ \citep[see Fig.~8 in][]{FederrathDuvalKlessenSchmidtMacLow2010}.\footnote{Note that even if the turbulence \emph{driving field} is fully compressive ($b=1$), there is still a substantial fraction of solenoidal modes that will be excited in the \emph{velocity field} via non-linear interactions \citep{Vishniac1994,SunTakayama2003,KritsukEtAl2007,FederrathDuvalKlessenSchmidtMacLow2010}, baroclinic instability \citep{DelSordoBrandenburg2011,PadoanEtAl2016,PanEtAl2016}, and by viscosity across density gradients \citep{MeeBrandenburg2006,FederrathEtAl2011PRL}.} Here we determine the turbulence driving parameter $b$ by measuring the standard deviation of the density fluctuations $\sigrho$ and the standard deviation of the probability distribution function (PDF) of the turbulent velocity field in {\brick}. We find that the turbulence driving in {\brick} is dominated by solenoidal shearing motions ($b<0.4$), while spiral-arm clouds have a substantial compressive driving component, $b>0.4$. Our results support the idea that shear is a typical driving mode of the turbulence in the CMZ and possibly in the centers of other galaxies, as proposed by \citet{KrumholzKruijssen2015} and Kruijssen et al., in preparation. This solenoidal driving mode can suppress star formation \citep{FederrathKlessen2012,PadoanEtAl2014} and may thus provide a possible explanation for the low SFR in the CMZ. \subsection{Universal filament properties?} Interstellar filaments are considered to be fundamental building blocks of molecular clouds, playing a crucial role in star formation \citep{SchneiderElmegreen1979,BalsaraEtAl2001,AndreEtAl2014}. Indeed, star-forming cores in nearby spiral-arm clouds are often located along dense filaments \citep{PolychroniEtAl2013,KonyvesEtAl2015} and young star clusters tend to form at their intersections \citep{Myers2011,SchneiderEtAl2012}. Recent observations and simulations of spiral-arm clouds show that filaments have coherent velocities \citep{HacarEtAl2013,MoeckelBurkert2015,HacarEtAl2016,SmithEtAl2016} and orientations preferentially (but not always) perpendicular to the magnetic field \citep{SugitaniEtAl2011,GaenslerEtAl2011,PalmeirimEtAl2013,Hennebelle2013,Tomisaka2014,ZhangEtAl2014,PlanckMagneticFilaments2014,PlanckMagneticFilaments2015a,PlanckMagneticFilaments2015b,PillaiEtAl2015,SeifriedWalch2015}. Most importantly, filaments seem to have a nearly universal width $W_\mathrm{fil}\sim0.1\,\pc$ \citep{ArzoumanianEtAl2011,JuvelaEtAl2012a,PalmeirimEtAl2013,MalinenEtAl2012,BenedettiniEtAl2015,KirkEtAl2015,WangEtAl2015,RoyEtAl2015,SaljiEtAl2015,KainulainenEtAl2016}.\footnote{Note that \citet{JuvelaEtAl2012a} and \citet{SaljiEtAl2015} found maximum variations of $W_\mathrm{fil}$ by a factor of $28$, while \citet{ArzoumanianEtAl2011} found maximum variations up to a factor of $10$. Thus, the term `universal' means in this context that $W_\mathrm{fil}$ definitely varies by less than two orders of magnitude, but more likely within factors of only a few around $0.1\,\pc$. Also note that \citet{SmithGloverKlessen2014} found somewhat larger values and variations of $W_\mathrm{fil}$ from simulations, in contrast to the observations in \citet{ArzoumanianEtAl2011}.} \citet{Federrath2016} provided a turbulence-regulated model for $W_\mathrm{fil}$, which is based on the sonic scale of the turbulence. Here we show that over-dense regions are located along filaments also in the CMZ cloud {\brick}, but the average filament column density is about 1--2 orders of magnitude higher compared to nearby clouds. Surprisingly though, the average filament width is similar in {\brick} to solar neighborhood clouds. Given the significant difference in gas temperature and magnetic fields in the CMZ, it seems surprising that $W_\mathrm{fil}$ is similar in {\brick} to nearby clouds. We explain the universal value for $W_\mathrm{fil}$ with the \emph{sonic scale}---the transition scale from supersonic to subsonic turbulence, following the theoretical model developed in \citet{Federrath2016}. We find excellent agreement between the measured filament width and the predicted sonic scale, both in {\brick} and in nearby clouds. The paper is organized as follows. Section~\ref{sec:observations} summarizes the observational data. In Section~\ref{sec:dens}, we identify filaments, measure their width and column density, and reconstruct the volume density dispersion of {\brick}. We measure the velocity PDFs of the total and turbulent (gradient-subtracted) velocity field in Section~\ref{sec:vels}. Numerical simulations to constrain the turbulent magnetic field are presented in Section~\ref{sec:mag}. We summarize all our measured and derived physical parameters of {\brick} in Table~\ref{tab:brick} of Section~\ref{sec:physics}. Sections~\ref{sec:sonicscale} and~\ref{sec:driving} provide a detailed discussion of derived sonic scale and turbulence driving parameter with comparisons to nearby clouds. A discussion of the limitations of this work are presented in Section~\ref{sec:caveats}. Our conclusions are summarized in Section~\ref{sec:conclusions}. | brick}. Here we list the most important results and conclusions: \begin{enumerate} \item Using the DisPerSE filament detection algorithm, we find 11 high-S/N filaments in the dense gas of {\brick} (see Fig.~\ref{fig:coldensimage}). Located along some of these filaments are three over-dense regions with a column density exceeding $2.5\times10^{23}\,\cm^{-2}$. As shown in previous studies, one of these cores has a water maser, which may indicate local active star formation. We find that the filling fraction of these cores is only 0.1\% of the total area of {\brick}, indicating inefficient dense-core and star formation. \item We construct the average radial profile of the filaments and find a typical filament column density of $\sim10^{23}\,\cm^{-2}$, which is an order of magnitude higher than the average filament column density observed in nearby spiral-arm clouds. We measure an average width of $W_\mathrm{fil}=0.17\pm0.08\,\pc$ (see Fig.~\ref{fig:filprof}). \item We find that the filament width does not significantly depend on the orientation of the filaments with respect to the ordered magnetic field component in \brick. \item Based on the column density PDF analyzed in \citet{RathborneEtAl2014} and combined with the column density power spectrum, we reconstruct the volume density dispersion, $\sigrho=1.3\pm0.5$, using the method developed in \citet{BruntFederrathPrice2010a}. \item Analyzing the spatial distribution of the HNCO intensity-weighted velocity, we see a strong large-scale velocity gradient across the whole cloud, which is likely associated with strong shearing motions (Kruijssen et al., in preparation). We subtract the large-scale gradient in order to obtain the distribution of turbulent velocities. From the Gaussian shape of the velocity PDF (Fig.~\ref{fig:velocities}), we find a turbulent velocity dispersion of $\sigma_{v,\mathrm{1D}}=3.9\pm0.1\,\km\,\s^{-1}$, which is significantly smaller than the total velocity dispersion ($8.8\pm0.2\,\km\,\s^{-1}$). \item Using magnetohydrodynamical turbulence simulations that take the measured turbulent velocity dispersion and the total (ordered+turbulent) magnetic field strength $B_\mathrm{tot}=2.2\,(0.9)\,\mathrm{m}\Gauss$ adapted from \citet{PillaiEtAl2015} as input, we determine the \emph{turbulent} magnetic field component $B_\mathrm{turb}=130\pm50\,\mu\Gauss$ (Fig.~\ref{fig:magnetic}). Given the velocity dispersion and strong ordered field in {\brick}, our simulations show that $B_\mathrm{turb}$ can only grow to $\lesssim\,B_\mathrm{tot}/10$. \item Using $B_\mathrm{turb}$ and adding the gas temperature $T=100\pm50\,\mathrm{K}$ constrained in the literature, we derive the sound speed, the Alfv{\'e}n speed and the ratio of thermal to magnetic pressure, plasma $\beta$ (Tab.~\ref{tab:brick}). Using these measurements, we derive a 3D turbulent sonic Mach number of $\mach=11\pm3$ and a turbulent Alfv{\'e}n Mach number of $\macha=4.6\pm2.1$ for {\brick}. \item We measure the effective cloud diameter $L=4.7\pm0.1\,\pc$ and combine it with the Mach number and plasma $\beta$ to derive the sonic scale $\ls$ of the turbulence in {\brick}. We find $\ls=L\mach^{-2}(1+\beta^{-1})=0.15\pm0.11\,\pc$, in agreement with our measurement of the filament width, $W_\mathrm{fil}=0.17\pm0.08\,\pc$. This supports the idea that the filament width is determined by the sonic scale, Equation~(\ref{eq:ls}), both in the CMZ and in spiral-arm clouds \citep{Federrath2016}. We caution that Equation~(\ref{eq:ls}) strictly only applies to the filament populations perpendicular to the ordered magnetic field; however, we find similar widths for parallel and perpendicular filaments (see Fig.~\ref{fig:filprof}). \item Our results imply that the filament width in {\brick} is similar to the filament width in nearby clouds, despite the orders-of-magnitude difference in some physical parameters of nearby clouds compared to the CMZ. The reason behind the similarity in $W_\mathrm{fil}$ is the sonic scale, Equation~(\ref{eq:ls}). It depends only on $L$, $\mach=\sigma_{v,\mathrm{3D}}/\cs$ and $\beta=p_\mathrm{thermal}/p_\mathrm{magnetic}$. While the thermal and magnetic pressure are both an order of magnitude higher in {\brick} compared to clouds in solar neighborhood, the ratio (plasma $\beta\sim0.3$) is similar in both environments. The same applies for the sonic Mach number---both $\sigma_{v,\mathrm{3D}}$ and $\cs$ are individually enhanced in {\brick} by factors of a few, but their ratio ($\mach\sim10$) is again similar to nearby clouds \citep{SchneiderEtAl2013}. \item Using the reconstructed volume density dispersion $\sigrho$ together with $\mach$ and $\beta$ allows us to derive the driving mode parameter $b$ of the turbulence, following Equations~(\ref{eq:b}) and~(\ref{eq:bvar}). We find $b=\sigrho\mach^{-1}(1+\beta^{-1})^{1/2}=0.22\pm0.12$, indicating solenoidal driving in {\brick}. \item We argue that the solenoidal driving in this Galactic-Center cloud is caused by strong shear, in agreement with the strong large-scale velocity gradient (c.f.~Fig.~\ref{fig:velocities}) and with detailed numerical simulations of CMZ clouds. We speculate that this solenoidal mode of turbulence driving might be the typical driving mode in the centers of galaxies, because of the enhanced shear in such environments. The solenoidal (shearing) mode of turbulence might explain the low SFRs observed in the CMZ compared to spiral-arm clouds, where the driving appears to have a significantly more compressive component, $b>0.4$ (see Fig.~\ref{fig:b}). Using SFR theory based on MHD turbulence, we find that $b=0.22$ yields a factor of $6.9$ lower SFR compared to $b=0.5$, emphasizing the role of the turbulence driving parameter. \end{enumerate} | 16 | 9 | 1609.05911 |
1609 | 1609.06387_arXiv.txt | We present the ${\it Kepler}$ photometry of KIC 6048106 exhibiting O'Connell effect and multiperiodic pulsations. Including a starspot on either of the components, light-curve synthesis indicates that this system is a semi-detached Algol with a mass ratio of 0.211, an orbital inclination of 73.9 deg, and a large temperature difference of 2,534 K. To examine in detail both spot variations and pulsations, we separately analyzed the {\it Kepler} time-series data at the interval of an orbital period by an iterative way. The results reveal that the variable asymmetries of the light maxima can be interpreted as the changes of a magnetic cool spot on the secondary component with time. Multiple frequency analyses were performed in the outside-eclipse light residuals after removal of the binarity effects from the observed {\it Kepler} data. We detected 30 frequencies with signal to noise amplitude ratios larger than 4.0, of which six ($f_2$--$f_6$ and $f_{10}$) can be identified as high-order (17 $\le n \le$ 25) low-degree ($\ell$ = 2) gravity-mode pulsations that were stable during the observing run of 200 d. In contrast, the other frequencies may be harmonic and combination terms. For the six frequencies, the pulsation periods and pulsation constants are in the ranges of 0.352$-$0.506 d and 0.232$-$0.333 d, respectively. These values and the position on the HR diagram demonstrate that the primary star is a $\gamma$ Dor variable. The evolutionary status and the pulsation nature of KIC 6048106 are discussed. | $\gamma$ Dor stars are A$-$F stars of luminosity class IV$-$V near the red edge of the $\delta$ Sct instability strip in the Hertzsprung-Russell (HR) diagram. Their observational properties are very similar to those of $\delta$ Sct stars, but they pulsate in high-order gravity ($g$) modes driven by convective blocking (Guzik et al. 2000; Dupret et al. 2004, 2005) with typical periods of 0.4$-$3 d and pulsation constants of $Q >$ 0.23 d (Kaye et al. 1999; Henry et al. 2005). These pulsating stars are of great interest to asteroseismic studies because the $g$ modes assist in probing the deep stellar interiors near the core region. The number of the $\gamma$ Dor-type stars has increased dramatically by space-based missions such as {\it CoRot} (Hareter et al. 2010) and {\it Kepler} (Balona et al. 2011; Bradley et al. 2015). Nonetheless, only thirteen eclipsing binaries (EBs) have been known to contain $\gamma$ Dor pulsating candidates (Maceroni et al. 2014; Kurtz et al. 2015; \c Cakirli \& Ibano\v{g}lu 2016), of which V551 Aur may be a $\delta$ Sct star as indicated by its pulsation frequencies (Liu et al. 2012). Because EBs are a primary source of the fundamental stellar properties such as mass and radius, the pulsating EBs are ideal targets for the study of the interior structure and evolution of stars from their binarity and pulsation features. Recently, \c Cakirli \& Ibano\v{g}lu (2016) suggested three possible relationships between the pulsation periods for $\gamma$ Dor components and other parameters (binary orbital periods for 11 EBs, surface gravities for 6 pulsating components, and the gravitational forces from companions for 5 EBs). As in case of the EBs with $\delta$ Sct components (Soydugan et al. 2006; Liakos et al. 2012), the longer the binary orbital period, the longer the pulsation period: $P_{\rm pul}$ = 0.425$P_{\rm orb}$ $-$ 0.355. In addition, as the surface gravity of pulsating components and the gravitational force from the companions decrease, their pulsation period increases. However, these relationships need further confirmation by new discoveries because the number of such stars remains small. KIC 6048106 (R.A.$_{2000}$=19$^{\rm h}$34$^{\rm m}$14$\fs028$; decl.$_{2000}$=+41$^{\circ}$23${\rm '}$43$\farcs$26; $K_{\rm p}$=$+$14.091; $g$=$+$14.303; $g-r$=$+$0.283) was announced to be an EB pulsating at frequencies of 0.43$-$3.28 d$^{-1}$ by Gaulme \& Guzik (2014). In this paper, we demonstrate that the binary system is a semi-detached Algol with a $\gamma$ Dor-type pulsating component, based on the precise and nearly-continuous {\it Kepler} data during a period of approximately 200 d. In Section 2, we carry out light-curve synthesis and present the binary parameters including the absolute dimensions. Section 3 describes the frequency analysis for the light residuals from each spot model. Finally, we summarize and discuss our conclusion in Section 4. | In this paper, we studied both the binarity and pulsation of KIC 6048106 from detailed analyses of the {\it Kepler} observations obtained during Quarters 14 and 15. The {\it Kepler} time-series data display mutiperiodic pulsations and the O'Connell effect with unequal light levels at the quadratures (Max I and Max II). The asymmetric light curve was modelled by applying a single spot to either of the components: a hot spot on the primary star and a cool spot on the secondary. Our light-curve synthesis indicates that KIC 6048106 is a classical Algol-type system with parameters of $q$=0.211, $i$=73$^\circ$.9, and ($T_{1}$--$T_{2}$)=2,534 K; the primary component fills about 52\% of its limiting lobe and is slightly smaller than the lobe-filling secondary. The locations of the components in the HR diagram are shown in Figure 7, together with those of other well-studied semi-detached Algols (\. Ibano\v{g}lu et al. 2006) and the $\gamma$ Dor stars in 9 EBs (Maceroni et al. 2014; \c Cakirli \& Ibano\v{g}lu 2016). Here, the dashed and dash-dotted lines are the instability strips of $\gamma$ Dor and $\delta$ Sct stars, respectively. The pulsating primary star of KIC 6048106 resides within the $\gamma$ Dor region on the zero-age main sequence (ZAMS), and the secondary lies in a location where the secondary components of other Algols exist. The {\it Kepler} data indicated that the light curve of KIC 6048106 has varied due to the combination of both spot and pulsation. To explore in detail the light variations, we individually analyzed the {\it Kepler} light curve at the interval of an orbital period by the iterative method described in the previous section. The variable asymmetries of the light maxima can be explained by the changes of the cool spot with time, which may be formed from magnetic dynamo-related activity because the system is rotating rapidly and the secondary component should be a deep convective envelope as surmised from its temperature. In order to understand the pulsational characteristics of the system, multiple frequency analyses were applied to the whole outside-eclipse light residuals, removing the binarity effects from the observed {\it Kepler} data. Thirty frequencies with S/N ratios larger than 4.0 were found in the range of 0.31$-$5.78 d$^{-1}$ with amplitudes between 0.14 and 3.29 mmag. Among these, six ($f_2, f_3, f_4, f_5, f_6$, and $f_{10}$) may be pulsation frequencies in $g$-mode region, which were stable during the observational interval of about 200 d. We computed the pulsation constants from the cool-spot model parameters in Table 1 and the well-known relation of $\log Q_i = -\log f_i + 0.5 \log g + 0.1M_{\rm bol} + \log T_{\rm eff} - 6.456$ (Petersen \& J\o rgensen 1972). The results are listed in the third column of Table 4. The $Q$ values and the position of KIC 6048106 on the HR diagram demonstrate that the detached primary component would be a $\gamma$ Dor-type pulsating star. On the other hand, all but the six frequencies appear to be harmonic and combination terms, some of which might arise from starspot activity or from alias effects caused by the orbital frequency. The orbital harmonics ($f_7, f_8, f_{11}, f_{12}, f_{19}, f_{20}$) can be stellar pulsations excited by the tidal forces of the secondary component (Welsh et al. 2011; Hambleton et al. 2013; Lee et al. 2016a). As binary stars are generally supposed to reach the synchronization before their semi-detached phases, the pulsating primary component of KIC 6048106 may have a synchronized rotation of approximately 50 km s$^{-1}$. Thus, we can make a possible identification of the radial order ($n$) and spherical degree ($\ell$) for the observed frequencies with the Frequency Ratio Method (Moya et al. 2005; Su\'arez et al. 2005), which is useful for $\gamma$ Dor stars with rotational velocities of $v \sin i \la$ 70 km s$^{-1}$ and at least three $g$-mode frequencies. Furthermore, it is possible to obtain the corresponding value for the integral of the Brunt-V\"ais\"al\"a frequency ($\cal J$). Following the procedure described by Lee et al. (2014), we determined the model frequency ratios ($f_i$/$f_5$)$_{\rm model}$ best-fitted to the observed ratios ($f_i$/$f_5$)$_{\rm obs}$, and identified the pulsation modes of the six frequencies. As listed in Table 4, the $f_2$, $f_3$, $f_4$, $f_5$, $f_6$ and $f_{10}$ frequencies are identified as degree $\ell$ = 2 for radial orders of $n$ = 25, 22, 24, 17, 21, and 18, respectively. The observed average value of $\cal J_{\rm obs}$ = 742.2$\pm$5.3 for the six frequencies is close to the theoretical integral of $\cal J_{\rm theo}$ $\approx$ 700 $\mu$Hz for a model of $\log$ $T_{\rm eff}$=3.845, 1.5 $M_\odot$, and [Fe/H]=0.0 in the $\cal J -$ $\log$ $T_{\rm eff}$ diagram given by Moya et al. (2005). Classical Algols are semi-detached interacting systems in which one type of interaction is mass transfer from the lobe-filling secondary to the detached primary component via the inner Lagrange $L_1$ point. The semi-detached configuration of KIC 6048106 permits some mass transfer between the component stars by means of a gas stream. Just as with the mass-accreting $\delta$ Sct components of semi-detached Algols (Mkrtichian et al. 2004; the so-called oEA stars), the secondary to primary mass transfer could at least be partly responsible for the $\gamma$ Dor-type oscillations detected in this paper. In addition, the pulsations may be influenced by the tidal and gravitation forces from the secondary component. As mentioned in the Introduction, \c Cakirli \& Ibano\v{g}lu (2016) presented three empirical relations for the $\gamma$ Dor stars in EBs, where the equation and figure between the pulsation period $\log P_{\rm pul}$ and the gravitational force $\log (F/M_1)$ are not consistent with each other. We think that their equation (9) is $\log (F/M_1) = -2.021 \log P_{\rm pul} +$ 2.093. The physical properties of KIC 6048106 match well these relationships of the orbital periods, the surface gravities, and the gravitational forces against the pulsation periods. However, because only thirteen stars, including KIC 6048106, have been identified as EBs containing $\gamma$ Dor-type components, additional discoveries and follow-up observations will help to reveal more accurate properties of the pulsating EBs. | 16 | 9 | 1609.06387 |
1609 | 1609.01272_arXiv.txt | We explore general scalar-tensor models in the presence of a kinetic mixing between matter and the scalar field, which we call Kinetic Matter Mixing. In the frame where gravity is de-mixed from the scalar this is due to disformal couplings of matter species to the gravitational sector, with disformal coefficients that depend on the gradient of the scalar field. In the frame where matter is minimally coupled, it originates from the so-called beyond Horndeski quadratic Lagrangian. We extend the Effective Theory of Interacting Dark Energy by allowing disformal coupling coefficients to depend on the gradient of the scalar field as well. In this very general approach, we derive the conditions to avoid ghost and gradient instabilities and we define Kinetic Matter Mixing independently of the frame metric used to described the action. We study its phenomenological consequences for a $\Lambda$CDM background evolution, first analytically on small scales. Then, we compute the matter power spectrum and the angular spectra of the CMB anisotropies and the CMB lensing potential, on all scales. We employ the public version of COOP, a numerical Einstein-Boltzmann solver that implements very general scalar-tensor modifications of gravity. Rather uniquely, Kinetic Matter Mixing weakens gravity on short scales, predicting a lower $\sigma_8$ with respect to the $\Lambda$CDM case. We propose this as a possible solution to the tension between the CMB best-fit model and low-redshift observables. | A key goal of current and future cosmic surveys is to constrain or possibly detect deviations from the standard $\Lambda$CDM model, which are expected if the origin of the present accelerated expansion is not a cosmological constant, but a dynamical field or a modification of General Relativity (see e.g. \cite{Amendola:2012ys,Amendola:2016saw}). To deal with the fact that there are many dark energy and modified gravity models (see for instance \cite{Clifton:2011jh,Joyce:2014kja}), effective approaches that describe these deviations for a large number of models in terms of a few time-dependent parameters have been proposed in the literature \cite{Creminelli:2008wc,Gubitosi:2012hu,Bloomfield:2012ff,Gleyzes:2013ooa,Bloomfield:2013efa,Gleyzes:2014rba,Battye:2012eu,Battye:2013ida,Baker:2011jy,Baker:2012zs,Skordis:2015yra,Lagos:2016wyv}. In most cases, these approaches are limited to a description of cosmological perturbations around a Friedmann-Lema\^itre-Robertson-Walker (FLRW) background in the linear regime (see however \cite{Bellini:2015wfa,Bellini:2015oua} for some nonlinear aspects), applicable to scales above $\sim 10$Mpc, where deviations from General Relativity are not yet well tested. This work focuses on the so-called Effective Theory of Dark Energy. Formulated for single scalar field models---i.e.~models where the time diffeomorphisms are broken while leaving the spatial ones preserved---in this approach the unitary (or uniform field) gauge action is given as the sum of all possible geometrical elements constructed from the metric and its derivatives that are invariant under the preserved diffs, i.e.~the spatial ones \cite{Creminelli:2006xe,Cheung:2007st}. It has been derived and studied for minimally and nonminimally coupled dark energy models, respectively, in \cite{Creminelli:2008wc} and \cite{Gubitosi:2012hu,Gleyzes:2013ooa} (see \cite{Piazza:2013coa,Tsujikawa:2014mba,Gleyzes:2014rba} for reviews). When restricting to the lowest order in derivatives, the final second-order action contains five free functions of time that parametrize any deviation from $\Lambda$CDM. As shown in \cite{Gleyzes:2013ooa}, four of these functions describe cosmological perturbations of effective theories of dark energy or modified gravity within the Horndeski class, i.e.~those with quadratic gravitational action with the same structure as Horndeski theories \cite{Horndeski:1974wa,Deffayet:2011gz,Kobayashi:2011nu}. This description has been reformulated in \cite{Bellini:2014fua} in terms of dimensionless functions that clearly parametrize deviations from General Relativity. The fifth function, denoted as $\alphaH$, describes scalar field models extending the Horndeski class, such as, e.g.,~the theories ``beyond Horndeski'' proposed in \cite{Gleyzes:2014dya,Gleyzes:2014qga} (see \cite{Zumalacarregui:2013pma} for an earlier proposal of theories beyond Horndeski). The Effective Theory formulation has been used to explore the observational consequences of deviations from $\Lambda$CDM (see for instance \cite{Piazza:2013pua,Kase:2014yya,Ade:2015rim,Lombriser:2015cla,Perenon:2015sla,Gleyzes:2015rua,Frusciante:2016xoj,Hu:2016zrh,Salvatelli:2016mgy,Renk:2016olm,Leung:2016xli,Pogosian:2016pwr}). In this direction, a few Einstein-Boltzmann solvers have been recently developed and employed \cite{Hu:2013twa,Raveri:2014cka,Bellini:2015xja,Zumalacarregui:2016pph,Huang:2015srv,zqhuang_2016_61166}. References~\cite{Gubitosi:2012hu,Gleyzes:2013ooa} assumed that all matter species are minimally coupled to the same metric, which we call Jordan frame metric for convenience. In general, however, there is no reason to impose this restriction. The universality of couplings is very well tested on Solar System scales \cite{Will:2014xja} but on cosmological scales constraints are much weaker and different species could have distinct couplings to the gravitational sector. If matter is universally but {\em nonminimally} coupled to the gravitational sector, in most cases it is convenient to perform a field redefinition of the metric that brings the system into the Jordan frame, where matter is minimally coupled. In general, this frame transformation depends on the scalar field and its derivatives and, as long as it is regular and invertible, it cannot change the physics (see e.g.~\cite{Domenech:2015tca}). The advantage of using the Jordan frame to derive predictions is that only the gravitational sector is non-standard; thus, one does not need to care about modifications of non-gravitational forces, which would otherwise greatly complicate the analysis. Along this line of thought, recently Ref.~\cite{Gleyzes:2015pma} extended the effective approach of \cite{Gubitosi:2012hu,Gleyzes:2013ooa} to allow for distinct conformal and disformal couplings of matter species to the gravitational sector. The treatment was restricted to effective theories within the Horndeski class and to conformal and disformal factors that depend only on the scalar field (not on its gradients). In this case, the full quadratic action depends on the four functions describing the gravitational sector and on two extra functions per species, describing the coupling to the scalar. However, two of these functions are redundant, because the structure of the action is preserved under transformations of the reference metric. This is expected, as it was shown that the structure of the Horndeski Lagrangians is preserved under disformal transformations with both conformal and disformal coefficients independent of the scalar field gradient \cite{Bettoni:2013diz}. The phenomenological aspects of general modifications of gravity described by Ref.~\cite{Gleyzes:2015pma} was studied in Ref.~\cite{Gleyzes:2015rua}, where constraints on the effective descriptions were derived from three observables: the galaxy and weak-lensing power spectra and the correlation between the Integrated Sachs-Wolfe (ISW) effect and the galaxy distribution. However, the study was restricted to the quasi-static limit, which is reliable on short enough scales and at late times, once the oscillations of the scalar fluctuations have been damped by the expansion of the universe. While this approximation is fairly good for current and future galaxy and weak lensing surveys,\footnote{The quasi-static approximation typically fails on scales $k \lesssim a H /c_s$, where $c_s$ is the sound speed of fluctuations of the scalar. As shown in \cite{Sawicki:2015zya}, this approximation should be reliable for surveys such as Euclid as long as the sound speed exceeds $10 \%$ of the speed of light, i.e.~$c_s \gtrsim 0.1$.} it fails on large scales or high redshifts. In this article we go one step forward, in two directions. First, in Sec.~\ref{section2} we extend the treatment of Ref.~\cite{Gleyzes:2015pma} and include in the gravitational action the fifth time-dependent function, $\alphaH$, describing models extending the Horndeski class. As shown in \cite{Gleyzes:2014dya,Gleyzes:2014qga}, the structure of the Lagrangian of theories beyond Horndeski is preserved under a disformal transformation of the metric with disformal coefficient that depends as well on the {\em gradient} of the scalar field, i.e.~of the form \be \label{disftr} \tilde g_{\mu \nu} = C(\phi) g_{\mu \nu} + D(\phi,X) \partial_\mu \phi \partial_\nu \phi \;, \qquad X \equiv g^{\mu \nu} \partial_\mu \phi \partial_\nu \phi \;. \ee Thus, in the following we consider the possibility that matter couples to a Jordan frame metric of this form.\footnote{Disformal transformations with $C=C(\phi,X)$ have been studied in the context of beyond Horndeski theories in \cite{Zumalacarregui:2013pma,Bettoni:2015wta} and in the context of degenerate higher-order theories in \cite{Crisostomi:2016tcp,Langlois:2015cwa,Achour:2016rkg,Crisostomi:2016czh}.} In particular, we denote the conformal and disformal coefficients of the nonminimal coupling of matter respectively as $C_{\rm m}(\phi)$ and $D_{\rm m} (\phi,X)$ (which can be distinct for different species). As shown in Sec.~\ref{section2}, the dependence of the disformal coupling on the derivative of the field introduces a kinetic mixing between the scalar and matter, which hereafter we call \emph{Kinetic Matter Mixing} (KMM), that has rather unique observational effects, as discussed below. To parametrize this direct kinetic coupling we introduce an additional function of time, \be \alpha_{\rm X, m} = \frac{X^2}{C_{\rm m}} \frac{\partial D_{\rm m}}{\partial X} \;, \ee where the right-hand side is evaluated on the background. Thus, the full quadratic action depends now on five functions describing the gravitational sector and three functions per species, describing the matter couplings. The structure of this action is preserved under transformations of the reference metric of the form \eqref{disftr}. Remarkably, $\alpha_{\rm X, m} $ is transformed into the beyond Horndeski parameter $\alphaH$ under a transformation which sets to zero the disformal coupling. Since KMM is a truely physical effect, it is possible to define a combination of these two parameters, proportional to $\left(\alphaH - \alpha_{\rm X, m} \right)^2 $ (c.f.~eq.~\eqref{gpar} below), that encodes in a frame-independent way the degree of kinetic mixing between matter and the scalar. While in Sec.~\ref{section2} we assume for simplicity that matter couples universally to the same Jordan frame metric, in App.~\ref{Lagmatter} we extend this treatment to multiple species with distinct couplings. Taking into account the invariance under the disformal transformation \eqref{disftr}, which reduces the number of independent functions of time by three, the whole system depends on a total of $2 + 3 N_S $ independent functions of time, where $N_S $ is the number of matter species. In the rest of the paper we assume that matter is universally coupled to the gravitational sector and work in the Jordan frame, where the coupling is minimal. In this frame, KMM is encoded in the beyond Horndeski parameter $\alphaH$. We then extend the treatment of Ref.~\cite{Gleyzes:2015rua} and explore the phenomenological consequences of general late-time modifications of gravity including beyond Horndeski theories (see also \cite{Kobayashi:2014ida,DeFelice:2015isa} and \cite{Saito:2015fza,Kase:2015zva,Sakstein:2016ggl,Babichev:2016jom} for an earlier study of the observational consequences of beyond Horndeski theories, respectively in cosmology and astrophysics). In Sec.~\ref{sec3}, we focus on short scales. In particular, we derive the eigenmodes of propagation of the scalar field and matter, which in the presence of a nonvanishing $\alphaH$ are mixed by their kinetic coupling. Moreover, we obtain the evolution equations in the quasi-static regime, which govern the dynamics once the oscillating modes have been damped by the expansion. Appendix \ref{Fullaction} contains the full action of perturbations in Newtonian gauge, derived for completeness, while the transition between the oscillating regime and the quasi-static limit is discussed in App.~\ref{app:QS}. In Sec.~\ref{sec:alphaH} we go beyond the quasi-static approximation and explore the full range of cosmological scales using the linear Einstein-Boltzmann solver of Cosmology Object Oriented Package (COOP)~\cite{zqhuang_2016_61166},\footnote{See \url{http://www.cita.utoronto.ca/~zqhuang/} for documentation.} which solves cosmological perturbations including very general deviations from $\Lambda$CDM in terms of the Effective Theory of Dark Energy description \cite{Gleyzes:2014rba}. In particular, assuming the background expansion history of $\Lambda$CDM, we compute the matter power spectrum, the Cosmic Microwave Background (CMB) anisotropies angular power spectrum, and the CMB lensing potential angular spectrum in the presence of KMM, for a non vanishing $\alphaH$ parameter. As we will see, on ``short'' scales, i.e.~for $k \gtrsim 10^{-3}h\,\text{Mpc}^{-1}$, the quasi-static approximation provides the correct amplitude for the linear growth factor, which is scale independent and suppressed with respect to the $\Lambda$CDM case. On larger scales, we compute the linear matter growth analytically using a perturbative expansion in $\alphaH$ that confirms the numerical results. To contrast with the effects of $\alphaH$, in App.~\ref{sec:alphaB} we compute the same observables in the case of a kinetic mixing between the scalar field and gravity, the so called kinetic braiding~\cite{Deffayet:2010qz,Pujolas:2011he} (see \cite{Creminelli:2006xe,Creminelli:2008wc} for an earlier study), and we find agreement with the results of Ref.~\cite{Zumalacarregui:2016pph}. We compare these results with the quasi-static approximation and a perturbative expansion in the braiding parameter. In contrast to kinetic braiding or other modifications of gravity within the Horndeski class, the exchange of fifth force in KMM suppresses the power of matter perturbations on redshift-survey scales. In Sec.~\ref{sec4.3}, we study the possibility that the lack of power measured in the large scale structures and in tension with that inferred from the CMB anisotropies observed by Planck~\cite{Ade:2013zuv,Ade:2015xua} can be explained by the KMM special signature. Finally, we conclude in Sec.~\ref{sec_last}. | \label{sec_last} Using the framework of the Effective Theory of Dark Energy, in this paper we studied the observational effects of Kinetic Matter Mixing, i.e.~a kinetic coupling between matter and the cosmological scalar field, which is present if matter is disformally coupled to the gravitational sector with a disformal coupling that depends on the first derivative of the scalar field or in theories beyond Horndeski. In Sec.~\ref{section2}, we started by discussing the most generic quadratic action for cosmological perturbations in the presence of conformal and disformal couplings of matter to the gravitational sector, under the assumption that the disformal factor depends as well on the first derivative of the scalar field, other than its value. Moreover, we showed that a change of frame does not change the structure of the action but redefines the coefficients of the various operators. In particular, the coefficient of the operator that characterizes theories beyond the Horndeski class is redefined only by the dependence of the disformal coupling on the field derivative. This is explicitly shown by the frame-independent parameter $\lambda^2$, defined in eq.~\eqref{gpar}, which measures the degree of Kinetic Matter Mixing. By diagonalizing the kinetic action, we derived the conditions that one must require for the perturbations to be free of ghosts and of gradient instabilities (the generalization to multiple matter species is given in App.~\ref{Lagmatter}). After this general frame-independent description, in Sec.~\ref{sec3} we assumed that matter is universally coupled and, without loss of generality, we considered the case where it is also minimally coupled, i.e.~the Jordan frame description, where observational predictions are more easily derived. We then discussed the short-scale regime and derived the eigenmodes of the acoustic oscillations, which are mixed states of matter and the scalar field waves. Focussing on the case where matter is made of nonrelativistic particles (such as cold dark matter or baryons) we derived the equations in the quasi-static approximation and discussed (see App.~\ref{app:QS}) how the quasi-static regime is reached during the cosmological evolution. These equations allow for a clear analytical understanding of the effects of modifications of gravity due to Kinetic Matter Mixing. In particular, while models in the Horndeski class only modify the Poisson equation with an effective Newton constant, Kinetic Matter Mixing also induces an additional friction term. Remarkably, requiring the stability conditions implies that gravity is weakened on short scales, an effect which is hard to reproduce in models within the Horndeski class. Finally, by comparing the quasi-static solution to the full numerical one, we showed that the quasi-static limit approximates very well the dynamics on scales shorter than the sound horizon. In Sec.~\ref{sec:alphaH} we focussed on the cosmological effects of the beyond Horndeski operator, obtaining the full numerical solutions using the publicly available Einstein-Boltzmann solver of COOP~\cite{zqhuang_2016_61166}. Using these solutions, we derived the matter power spectrum at two different redshifts, and the angular spectra of the CMB lensing potential and of the CMB anisotropies. On small scales, i.e.~for $k \gtrsim \text{few} \times 10^{-3} \textrm{Mpc}$, the solution matches the quasi-static regime and the matter power spectrum is suppressed independently of $k$. An analytical study of the large scales is complicated by the complexity of the full system of equations. However, we obtained analytical solutions on these scales by perturbing around the $\Lambda$CDM solutions for small Kinetic Matter Mixing. The agreement with the numerical solution is excellent. Moreover, its simplicity allows an immediate understanding of the behavior of the perturbations and their observables. Similarly to the matter power spectrum, also the angular spectrum of the CMB lensing potential is suppressed. The CMB anisotropy is affected at very low multipoles through the ISW effect, which is enhanced, and on very high multipoles because of the suppression of the lensing potential. In App.~\ref{sec:alphaB}, we compared this case with the one of kinetic braiding, which displays qualitatively opposite effects. Also in this case we studied analytically the large-scale behavior and derived the value of the crossing scale, i.e.~the scale at which the power spectrum displays the transition between the short-scale enhancement and large-scale suppression. As mentioned above, Kinetic Matter Mixing appears as the only modification of gravity in the context of single-field models that weakens the strength of gravity on small scales. Therefore, in Sec.~\ref{sec4.3} we entertained the possibility that the tension between the Planck data and small-scale observations can be explained by this effect. In particular, as shown in Figs.~\ref{fig:sigma8} and \ref{fig:fsigma8}, KMM predicts a lower value of $\sigma_8$ and $f\sigma_8$, which could be made compatible with those measured by weak lensing and redshift-space distortion observations. We postpone to future work a more consistent dedicated analysis that marginalizes over the other cosmological parameters. In summary, we presented a robust theoretical understanding of the effects of Kinetic Matter Mixing across different observables and scales. These effects may be a smoking gun of modified gravity for the next observational missions and a complete forecast, taking into account the characteristics of the next missions, is an obvious next step. \vspace{0.5cm} \noindent {\bf Acknowledgements:} It is a pleasure to thank J\'er\^ome Gleyzes, Martin Kunz, David Langlois, Jean-Baptiste Melin, Federico Piazza, Ignacy Sawicki and Jeremy Tinker for useful discussions. M.M. and F.V. also thank the Theoretical Physics Departments of CERN and Universit\'e de Gen\`eve for hospitality during large part of this work. G. D'A. thanks the CCPP at New York University for hospitality during the final stages of this work. \vspace{0.5cm} \appendix | 16 | 9 | 1609.01272 |
1609 | 1609.04162.txt | Using {\it HST} slitless grism data, we report the spectroscopic confirmation of two distant structures at $z \sim 2$ associated with powerful high-redshift radio-loud AGN. These rich structures, likely (forming) clusters, are among the most distant currently known and were identified on the basis of {\it Spitzer}/IRAC $[3.6] - [4.5]$ color. We spectroscopically confirm 9 members in the field of MRC~2036$-$254, comprising eight star-forming galaxies and the targeted radio galaxy. The median redshift is $z=2.000$. We spectroscopically confirm 10 members in the field of B3~0756+406, comprising eight star-forming galaxies and two AGN, including the targeted radio-loud quasar. The median redshift is $z = 1.986$. All confirmed members are within $500$ kpc ($1$ arcmin) of the targeted AGN. We derive median (mean) star-formation rates of $\sim 35~M_{\odot}\rm ~ yr^{-1}$ ($\sim 50~M_{\odot}\rm ~ yr^{-1}$) for the confirmed star-forming members of both structures based on their [\ion{O}{3}]$\lambda5007$ luminosities, and estimate average galaxy stellar masses $\la 1 \times 10^{11} ~M_{\odot}$ based on mid-infrared fluxes and SED modeling. Most of our confirmed members are located above the star-forming main-sequence towards starburst galaxies, consistent with clusters at these early epochs being the sites of significant levels of star formation. The structure around MRC~2036$-$254 shows an overdensity of IRAC-selected candidate galaxy cluster members consistent with being quiescent galaxies, while the structure around B3~0756+406 shows field values, albeit with many lower limits to colors that could allow an overdensity of faint red quiescent galaxies. The structure around MRC~2036$-$254 shows a red sequence of passive galaxy candidates.\\ | \label{sec:intro} At low to intermediate redshifts ($z \la 1.4$), massive early-type galaxies dominate galaxy cluster cores and form a tight red sequence (e.g., \citealp{Lidman08}, \citealp{Mei09}). The few studies at higher redshifts suggest that clusters at $z>1.5$ are still in the process of forming (\citealp{Snyder12}, \citealp{Mei15}), and that although clusters at these redshifts show a mixed population of both star-forming (SF) and quiescent galaxies, even the reddest (early-type) galaxies show on-going star-formation (\citealp{Mei15}). Star-formation activity has also been observed in the cores of massive galaxy clusters at $z > 1.4$ (e.g., \citealp{Tran10}, \citealp{Hayashi10}, \citealp{Zeimann13}, \citealp{Alberts14}, \citealp{Bayliss14}). For example, based on a sample of 16 spectroscopically confirmed clusters at $1 < z < 1.5$, \cite{Brodwin13} showed that at $z > 1.3$ the fraction of SF cluster members increases towards the cluster centers. These results suggest that the majority of star formation actually occurs in high-density environments at early epochs, implying that environmental-dependent quenching has not yet been established at $z > 1.3$. \cite{Brodwin13} predicted that this transition redshift should be a function of halo mass, with more massive halos transitioning earlier. This is consistent with the findings of \cite{Wylezalek14} of a $z\sim3$ transition period for clusters around radio-loud active galactic nuclei (RLAGN), which are extreme objects that tend to reside in the most massive dark matter halos (e.g., \citealp{Mandelbaum09}, \citealp{Hatch14}, \citealp{Orsi16}). To better understand the dependence of the formation mechanisms of massive galaxies on environment, we must focus on clusters at the relatively unexplored redshift range $z > 1.5$ where major assembly is in progress (e.g., \citealp{Mancone10}). Various selection methods are used to find (proto)cluster candidates, e.g., the red sequence (\citealp{GladdersYee00}, \citealp{Rykoff14}, \citealp{Licitra16}), a mid-infrared adaptation of the red sequence (\citealp{Muzzin13}, \citealp{Webb15}), photometric redshifts of infrared-selected samples (\citealp{Eisenhardt08}, \citealp{Stanford12}, \citealp{Zeimann12}), {\it Spitzer}/IRAC color selection (\citealp{Papovich08}, \citealp{Rettura14}), overdensities of sub-millimeter sources (\citealp{Smail14}, \citealp{Planck15}), X-ray emission (\citealp{Rosati98}, \citealp{Tozzi15}), and the Sunyaev Zel'dovich (SZ) effect (\citealp{Vanderlinde10}, \citealp{Bleem15}). These methods mostly rely on wide-field surveys. Discovering larger samples of galaxy clusters at high redshifts using these techniques therefore requires prohibitive amounts of telescope time over yet wider areas. Moreover, X-ray detections are limited by the surface brightness of the sources\footnote{With the caveat that based on the low-redshift \cite{Vikhlinin09} scaling relations, \cite{Churazov15} showed that at higher redshifts ($z \simeq 1-2$), clusters as massive as $z\sim0$ clusters should be as easily detectable.} dimming as $(1 + z)^4$. Both X-ray and SZ selections are also only able to detect very massive structures via their hot intra-cluster medium, which requires mature, collapsed clusters. Additionally, AGN activity increases for higher redshift clusters (e.g., \citealp{Galametz10}, \citealp{Martini13}), adding a complication for X-ray and SZ selections. These issues all conspire against finding galaxy clusters at high redshifts, although \cite{Mantz14} recently reported a massive cluster candidate ($M_{500}\sim(1 - 2)\times 10^{14}~M_{\odot}$) at $z\simeq1.9$ via a weak X-ray detection, the SZ decrement and photometric redshifts. Currently, only about ten galaxy clusters have been spectroscopically confirmed at $z > 1.5$ (e.g., \citealp{Papovich10} -- independently reported in \citealp{Tanaka10}, \citealp{Stanford12}, \citealp{Zeimann12}, \citealp{Gobat13}, \citealp{Muzzin13}, \citealp{Newman14}, \citealp{Mei15}). All these confirmed clusters are at $z \leq 2.0$. A few confirmed clusters and cluster candidates at $1.5 < z < 2$ have significant X-ray detections, implying they are likely virialized (e.g., \citealp{Santos11}, \citealp{Mantz14}, \citealp{Newman14}). This relatively small number of high-redshift confirmed clusters makes it challenging to draw a clear picture of cluster formation and evolution. Powerful high-redshift RLAGN are known to preferentially lie in overdense fields (with literature stretching back more than 50 years; e.g., \citealp{Matthews64}) and are efficient beacons for identifying large-scale structures and (proto)clusters. Indeed, targeted searches around RLAGN are a proven technique for identifying galaxy clusters at high redshifts (e.g., \citealp{Stern03}, \citealp{Venemans07}, \citealp{Galametz10}, \citealp{Hatch11}). Our team has made a major contribution to this effort with a targeted 400-hour {\it Warm} {\it Spitzer Space Telescope} program surveying 420 radio-loud AGN at $1.3 < z < 3.2$ across the full sky: Clusters Around Radio-Loud AGN (CARLA, \citealp{Wylezalek13, Wylezalek14}). Using a simple mid-infrared color selection technique, we successfully identified nearly 200 promising cluster candidates at $z > 1.3$. Ground-based observations are challenging for spectroscopically confirming high-redshift galaxy clusters because of atmospheric absorption and emission. In contrast, {\it Hubble Space Telescope} ({\it HST}) infrared spectra obtained with the Wide-Field Camera 3 (WFC3) slitless grism are free from atmospheric constraints and are thus ideal for obtaining spectra of high-redshift galaxies, albeit with a low dispersion ($46.5$ \AA/pix$^{-1}$) and a low resolving power ($ R = \lambda / \Delta \lambda = 130$; numbers for G141 grism, \citealp{Dressel14}). Following up on the {\it Spitzer}/CARLA survey, our team is using the WFC3 slitless G141 grism to study our $20$ densest cluster candidates at $1.4 \leq z \leq 2.8$. This paper presents early results confirming structures around MRC 2036$-$254 and B3~0756+406, two of the first fields to have their {\it HST} observations completed. Previous papers from the CARLA project include ground-based spectroscopic confirmation of two (proto)clusters, reported in \cite{Galametz13} and Rettura et al. (in prep.), ground-based imaging to study the formation histories of CARLA clusters, reported in \cite{Cooke15a, Cooke16}, and a comparison of mass-matched samples of radio-loud and radio-quiet galaxies at $z > 1.3$, showing that RLAGN indeed reside in significantly denser environments (\citealp{Hatch14}). We adopt here the \cite{Eisenhardt08} criteria of $z>1$ spectroscopic cluster confirmation: at least five galaxies within a physical radius of 2 Mpc whose spectroscopic redshifts are confined to within $\pm 2000(1+\left<z_{\rm spec}\right>) \rm~km~\sec^{-1}$. A significant concern is that this definition alone may also identify groups, protoclusters, sheets and filaments when applied to grism data. In the lower redshift universe, a more exacting definition for a confirmed galaxy cluster typically also requires: {\it (i)} detection of an extended X-ray emitting diffuse intracluster medium, {\it (ii)} a significant population of early-type (i.e., passive) galaxies, and {\it (iii)} a centrally concentrated distribution of galaxies. Current literature will often forego the first requirement, particularly for distant clusters, due to the challenges of acquiring such data --- and this is particularly problematic for our RLAGN targets due to Inverse Compton scattering of the cosmic microwave background by the hot plasma associated with AGN radio lobes into the X-ray regime. Nonetheless, we show that some CARLA systems have clear overdensities of passive galaxy candidates (see also \citealp{Cooke15a, Cooke16}, in prep.), and \citet{Wylezalek13} show that the CARLA cluster member candidates are, on average, centrally concentrated around the target RLAGN. This paper is organized as follows. Section \ref{sec:obs} briefly presents the CARLA sample and the overall strategy of our {\it HST}/CARLA program. Section \ref{sec:red} presents the analysis strategy carried out on the {\it HST} data, including detection limits. Section \ref{sec:res} presents the results, including cluster membership, star-formation rates (SFRs), and stellar masses. In Section \ref{s:discussion} we compare the {\it HST} results with the CARLA selection method and discuss the results. We summarize our work in Section \ref{sec:con}. We also present details on the data analysis in Appendix \ref{ap:notesdataprep}. Appendices \ref{ap:membprop} and \ref{ap:nonmembprop} list the line properties of structure members and non-members, respectively, and we present notes on individual sources in Appendix \ref{ap:clmemb}. Throughout, all magnitudes are expressed in the AB photometric system, and we use a flat $\Lambda$CDM cosmology with $H_{0} = 70 ~\rm km ~Mpc^{-1}~ s^{-1}$, $\Omega_{\Lambda} = 0.7$, and $\Omega_{m} = 0.3$.\\ %%OBSERVATIONS%% | \label{s:discussion} \subsection{Method Efficiency}\label{sec:dismethod} The two $z=2$ structures reported in this paper likely represent some of the highest redshift clusters currently spectroscopically confirmed. They illustrate the high efficiency of our approach to target RLAGN with IRAC mapping and {\it HST} grism follow-up. As the two fields are only a pilot study for our full sample of 20 cluster candidates, we next explore the efficiency of our grism spectroscopy and $[3.6] - [4.5]$ selection.\\ \subsubsection{Grism Efficiency} We assess in this section the outcome of our {\it HST} observations. In Figure \ref{fig:Globflowchart} we show the flowchart of the classification of {\it HST} sources from our master catalog. The classification is shown for both fields, with the numbers on the left side of each box corresponding to the MRC 2036$-$254 field and the numbers on the right side corresponding to the B3~0756+406 field. We also display in parentheses numbers corresponding to the classification for secure {\it Spitzer} color-selected candidates which have a single {\it HST} counterpart. Overall, we determine redshifts for $6\%$ of our exploitable {\it HST} sources (31/550 and 26/452, respectively, for the fields around MRC 2036$-$254 and B3~0756+406); where `exploitable' is defined to mean that $>75\%$ of the source continuum falls on the detector and contamination is less than $60\%$ of the cutout length. We find that $2\%$ of sources (9/550 and 10/452) are confirmed cluster members -- i.e., for every three redshifts that we measure, we find one cluster member. This is consistent with probing a biased environment, and is not an instrumental bias. For example, the redshift distribution of the 3D-HST field survey (\citealp{Momcheva15}) is roughly flat in the same redshift window ($0.7<z<2.3$), based on their sample of $46,256$ grism redshifts in this window obtained over $626~ \rm arcmin^{2}$ with a similar two-orbit per field $HST$ grism program. Of the exploitable {\it HST} sources, $71\%$ (370/550 and 344/452) are not detected in the dispersed grism data. Moreover, of the sources with spectral detections, $66\%$ (125/180 and 66/108) show continuum only, for which we could not measure a redshift because no emission lines were detected. These sources are likely a mixture of: {\it (i)} stars, {\it (ii)} old and passive (quiescent) galaxies with little star formation, {\it (iii)} galaxies at redshifts for which no strong features are covered by our grism observations, and {\it (iv)} SF galaxies for which we do not detect emission lines at the depth of our data (SFR $\la 20~M_{\odot} \rm~ yr^{-1}$, unless highly dust-obscured). Also, we have not been able to determine a redshift for $40\%$ (24/55 and 16/42) of sources with detected emission line(s) (with or without continuum) because of ambiguity on the nature of the line(s). This number improves by a factor of 3.5 when considering {\it Spitzer} color-selected sources since we have a color information which helps identify ambiguous emission lines. Overall, we also lose $26\%$ (195/745 and 170/622) of the source spectra because of full contamination of their spectral first orders ($15\%$, 112/745 and 88/622) or because their traces fell outside of the G141 detector ($12\%$, 83/745 and 82/622).\\ \subsubsection{IRAC Color Selection Efficiency} The primary aim of our {\it HST} program is to confirm galaxy clusters which were selected as overdense fields of mid-infrared color-selected galaxies. Hence, we also evaluate the success of the CARLA selection method given our {\it HST} observations. Focusing on the numbers in parentheses in the flowchart, we also find that a large fraction, $51\%$ (35/64 and 29/61, respectively for the fields around MRC 2036$-$254 and B3~0756+406) of the exploitable CARLA sources do not have any spectroscopic detection at the depth of our grism observations. This is better than the $71\%$ of all {\it HST} sources, implying that CARLA sources are brighter, on average, than our {\it HST} sources. This is unsurprising since rest-frame near-infrared luminosity strongly correlates with stellar mass (e.g., \citealp{Gavazzi96}), and our IRAC $\rm 4.5 \micron$ flux cut imposes a limiting stellar mass around $1 \times 10^{10} ~M_{\odot}$ on the CARLA sources, whereas the {\it HST} F140W imaging also detects sources below the IRAC flux limit. We determine the redshift of $18\%$ (10/64 and 12/61) of the exploitable {\it Spitzer} color-selected sources, and we find that $6\%$ (3/64 and 4/61) are cluster members. Again, we identify one cluster member for every three sources with redshift determinations. Among the CARLA sources with spectral detections, $57\%$ (17/29 and 18/32) show continuum without detectable emission lines. This suggests a large fraction of quiescent galaxies and/or galaxies with low or dust-extincted SFRs as potential cluster members (see Section \ref{sec:passivef}). Note the similarity with \cite{Brodwin13}, albeit at lower redshift, who found that $\sim 40\%$ of cluster members at $1.37<z<1.50$ are SF galaxies based on the {\it Spitzer} 24$\micron$ emission, using an infrared-selected sample with stellar masses $>10^{10.1}~M_{\odot}$, and a SFR lower limit of $50~M_{\odot}\rm ~yr^{-1}$ for the SF members. This is also consistent with the work of \cite{Cooke16}, who found that $76\%$ of CARLA galaxies ($M>10^{10}~M_{\odot}$) in the field of CARLA J1753+6311 ($z=1.58$) are quiescent. The CARLA selection method therefore finds both passive galaxies and SF galaxies, though the shallow grism data presented here are only capable of confirming the latter.\\ \subsection{Stellar Populations}\label{sec:passivef} \cite{Williams09} (see also \citealp{Labbe05}, \citealp{Wuyts07}, \citealp{Whitaker11}) have shown that it is possible to use rest-frame $UVJ$ colors to separate the SF, dusty SF and passive populations. To build our color-color diagrams (Fig. \ref{fig:ccd}), we use observed $z/i - \rm F140W$ vs. $\rm F140W - [3.6]$. Following \cite{Mei09} and using Mei's codes and the python version of {\ttfamily EZGAL} (\citealp{ManconeGonzalez12}), we transform \cite{Williams09} color limits into our observed apparent colors. We use a \cite{BruzualCharlot03} SSP model with galaxy formation redshifts averaged between $z_f=3$ and $8$, and metallicities equal to $40\%$ solar, solar and $2.5$ times solar; and we age the templates to $z=2$. We average the color conversions using a set of formation redshift ranges from $z_f=3-8$ to $z_f=5-8$ increasing by steps of $0.1$. This is a conservative simple model and assumes that passive galaxies are mostly located close to the passive boundary defined in \cite{Williams09}. A more detailed analysis of the CARLA cluster stellar population will be performed in future work. In Figure \ref{fig:ccd} we plot color-color diagrams for the two fields to investigate the presence of passive cluster members associated with CARLA J2039$-$2514 and CARLA J0800+4029, and in Figure \ref{fig:cmd} we plot their color-magnitude diagrams (CMDs). We use $i$-band data from \cite{Cooke15a} for CARLA J0800+4039 and a $4800~\sec$ $z$-band image obtained with VLT/ISAAC on UT 2002 July 17 for CARLA J2039$-$2514\footnote{Archival data from run ID 69.A-0234.}. For both CMDs, we use {\it HST}/F140W data to bracket the D4000 break at $z=2$. In the color-color diagrams, the passive candidates are located inside the upper left quadrant. The SF population is located below the horizontal line, whereas dusty SF galaxies lie on the right of the vertical boundary. We identify fourteen passive galaxies for CARLA J2039$-$2514 (two of them outside the plot) and two for CARLA J0800+4029.\\ \begin{figure*} \centering \includegraphics[width=8.97cm]{M2036_cmd_new.pdf} \includegraphics[width=8.97cm]{J0800_cmd_new.pdf} \caption[Color-magnitude diagrams of M2036 and J0800]{Color-magnitude diagrams of CARLA J2039$-$2514 (left) and CARLA J0800+4029 (right). Same markers as in Figure \ref{fig:ccd}, with the addition of sources detected in the optical and F140W but not detected in IRAC shown by small black dots. Sources below our optical detection limits ($2\sigma$) are set to these limits and shown with upward arrows. The gray and orange shaded areas represent estimates of a $z=2.0$ red sequence for delta-burst and exponentially decaying stellar populations respectively, described in Section \ref{subsec:redseq}. The thickness of these regions correspond to formation redshifts in the range $z_{f} = 3.0 - 8.0$. The solid blue lines represent the $z=1$ \cite{Mei09} color-magnitude relation passively evolved to $z=2.0$, as described in Section \ref{subsec:redseq} (the dashed blue lines represent the $3\sigma$ range).} \label{fig:cmd} \end{figure*} \subsubsection{Density of Passive Candidates}\label{subsec:passivedendity} We compare the density of passive, red ($([3.6]-[4.5])_{\rm AB} > -0.1$), sources in our RLAGN fields with densities of sources similarly selected in wide-field surveys. We make use of the 3D-HST multi-wavelength catalogs (\citealp{Skelton14}) in the five CANDELS fields (GOODS-North, GOODS-South, AEGIS, COSMOS and UDS; \citealp{Grogin11}). All CANDELS fields were at least partially covered by $i$-band, $z$-band, F140W, $3.6\micron$ and $4.5\micron$ observations. The F140W images have the most limited coverage; we therefore use the F140W image field of view to derive source densities (\citealp{Brammer12}). We isolate passive, red (i.e., CARLA), sources using the same color-selections as described above. At the depth of our RLAGN field data, we expect $\sim 9$ sources ($1.69 \pm 0.04$ arcmin$^{-2}$) selected by the $z$/F140W/[3.6] criterion in the $5.38$ arcmin$^{2}$ {\it HST} field of view of CARLA J2039$-$2514, and $\sim 2$ sources ($0.38 \pm 0.02 $ arcmin$^{-2}$) selected by the $i$/F140W/[3.6] criterion in the $5.54$ arcmin$^{2}$ {\it HST} field of view of CARLA J0800+4029. We identify $14$ ($\sim 2.6$ arcmin$^{-2}$) and $2$ ($\sim 0.4$ arcmin$^{-2}$) passive CARLA candidates for CARLA J2039$-$2514 and CARLA J0800+4029, respectively. CARLA J2039$-$2514 is therefore consistent with possessing an overdensity of $z \sim 2$ passive red galaxies relative to the field, while our comparison suggests that CARLA J0800+4029 has similar values than the field. However, these are likely underestimates of the true (over)densities given the limiting magnitudes in the CARLA imaging of those fields.\\ \subsubsection{Red Sequences and SF Populations}\label{subsec:redseq} In the CMDs (Fig. \ref{fig:cmd}), the orange and gray shaded areas show estimates of the expected color of $L^{*}$ early-type galaxies computed with a \cite{BruzualCharlot03} passive evolution model of solar metallicity and a \cite{Salpeter55} IMF\footnote{Using a \cite{Chabrier03} IMF instead does not significantly change the results.}. These were calculated using the python version of {\ttfamily EZGAL} (\citealp{ManconeGonzalez12}) with the appropriate filter transmission curves. We evolve a SSP (i.e., delta-burst), shown by the gray shaded areas, and an exponentially decaying stellar population of characteristic time-scale $\tau = 1 \rm ~Gyr$, shown by the orange shaded areas. The thickness of these regions correspond to formation redshifts over the range $3 < z_{f} < 8$. Both model evolutions are normalized to match the typical CARLA $L^{*}$ at $z=2.0$ (\citealp{Wylezalek14}), and assume slopes of $-0.1$, similar to what is observed in the Coma cluster at similar rest-frame wavelengths (\citealp{Eisenhardt07}). To compare with the color-magnitude relation (CMR) observed in confirmed X-ray and infrared detected clusters at redshift $z\sim1$, we plot the \cite{Mei09} early-type CMR (solid blue lines; the blue dashed lines show the $3\sigma$ dispersion around the mean) evolved passively to $z=2$. Following \cite{Mei09}, we use a \cite{BruzualCharlot03} SSP model with galaxy formation redshifts between $3$ and $8$, and metallicities equal to $40\%$ solar, solar and $2.5$ times solar, to convert rest-frame $(U-B)$ vs. $M_{B}$ derived at $0.8<z<1.3$ to our observed bandpasses at $z=2$. We again average the color conversions using a set of ranges from $z_f=3-8$ to $z_f=5-8$. The \cite{Mei09} early-type galaxy CMR parameters are derived using {\it HST}/ACS filters that correspond to rest-frame $(U-B)$ and $M_{B}$ in the range $0.8<z<1.3$. While the $z$-band and F140W-band correspond to the same rest-frame for CARLA J2039$-$2514, the $i$-band probes bluer stellar populations than the $U$-band rest-frame. Within the uncertainties, these relations are consistent with the predictions from \cite{BruzualCharlot03}. The majority of CARLA J2039$-$2514 passive candidates identified in Figure \ref{fig:ccd}, with $z-\rm F140W \ga 1$ mag and $\rm F140W > 20$ mag, agree well with the SSP models shown on the CMD (Fig. \ref{fig:cmd}), suggesting CARLA J2039$-$2514 hosts a population of quiescent galaxies consistent with a cluster red sequence at $z=2.00$. On the other hand, CARLA J0800+4029, unlike CARLA J2039$-$2514, does not exhibit $\rm F140W < 23$ mag passive candidates consistent with a cluster red sequence, and the depth of our $i$-band data prevents us from confirming $\rm F140W > 23$ mag sources for this field. Specifically, thirteen sources (marked with upward arrows in Figures \ref{fig:ccd} and \ref{fig:cmd}) do not have $i$-band detections and therefore require deeper data to determine whether they are passive candidates populating what could constitute a nascent cluster red sequence. In each color-color diagram, we only identify a handful of dusty SF CARLA candidates, none of which are spectroscopically confirmed as members or non-members. All spectroscopically confirmed members and non-members are identified as SF galaxies in the color-color diagrams, as expected, with the exception of the target RLAGN in CARLA J2039$-$2514 which falls in the passive region as expected for a type-2 RLAGN. Four of the confirmed members of CARLA J2039$-$2514 seem to be well described by the exponentially decaying model of star formation. The remaining two members that have both F140W and $z$-band detections, and most of the confirmed members of CARLA J0800+4029 do not seem to agree with this model, suggesting that the build-up of their stellar content followed a diversity of star-formation histories. This supports \cite{Cooke15a} who showed that the star formation histories (SFHs) of CARLA cluster galaxies are best described by multiple bursts of star formation normally distributed over a few Gyrs, or distributed following the cosmic SFH. Exploration of a range of SFHs is however beyond the scope of this paper and will be addressed in future work. Finally, we note that only one blue ($([3.6] - [4.5])_{\rm AB} < -0.1$ mag) source falls in the passive region of the color-color diagrams, and that the putative red sequence of CARLA J2039$-$2514 almost exclusively comprises a population of quiescent galaxies up to $\rm F140W<24$ mag. The spectroscopic confirmation of these passive CARLA candidates, however, will require significantly deeper data.\\ \subsection{Comparison to Other High-Redshift Clusters}\label{p:comparison} We compare our results to other high-redshift clusters spectroscopically confirmed using {\it HST}/WFC3 grism data. \cite{Gobat13}, hereafter G13, reported on a rich cluster at $z \simeq 2$. They first reported evidence of a fully established galaxy cluster at $z=2.07$ (\citealp{Gobat11}) from X-ray emission, ground-based spectroscopy, and ground- and space-based photometry, later revised by G13 to a slightly lower redshift using deep {\it HST} grism observations. The system consists of two unrelated aligned structures, with the background overdensity being ``sparse and sheet-like". The G13 median redshift of the {\it HST}/WFC3 confirmed members of the foreground overdensity is $z = 1.993$, with a standard deviation of $0.012$. With median redshifts of $z = 2.000$ (standard deviation 0.005) and $z = 1.986$ (standard deviation 0.014), CARLA J2039$-$2514 and CARLA J0800+4029, respectively, are among the most distant confirmed clusters currently known. G13 confirmed 22 members with 18 orbits and three orientations (12.5 hrs on source) including five quiescent sources, and found no evidence for an already formed red sequence within $20\arcsec$ of the cluster core (\citealp{Gobat11}). More recently, \cite{Newman14}, hereafter N14, confirmed another rich cluster at high redshift using {\it HST}/WFC3 slitless spectroscopy. With 14 orbits, they confirmed 19 members at $z = 1.80$, of which more than $75\%$ are quiescent. They found a clear red sequence of observed mean color $\left<z - J\right> = 1.98 \pm 0.02$, which includes 13 of their 15 quiescent cluster members. \cite{Zeimann12}, hereafter Z12, also confirmed both emission-line and quiescent sources in a cluster at $z = 1.89$ using six orbits of {\it HST}/WFC3 grism spectroscopy. Z12 showed that a significant fraction of early-type galaxies in the cluster field were consistent with forming a red sequence. These deep observations show that clusters can host a substantial fraction of quiescent galaxies even at early epochs. By design, our shallow two-orbit per field strategy only confirms SF members. In addition to the confirmed emission-line members, we find a large fraction ($79\%$, 52/64 and 47/61) of {\it Spitzer} color-selected cluster candidates below our detection limit or showing continuum only. CARLA J2039$-$2514 and CARLA J0800+4029 are therefore relatively robust confirmations, and have a high potential for being richer structures than what our shallow {\it HST} observations allowed us to unveil. Furthermore, we note that our two-orbit strategy spectroscopically confirms $\sim 5$ cluster members per orbit, which is approximately four times more efficient than the $>10$-orbit programs reported in G13 and N14. From X-ray emission and richness, G13 estimated a mass of $M_{200} \sim 5\times 10^{13} ~M_{\odot}$ for their cluster. N14 reported a massive cluster with $M_{200} = (2-3) \times 10^{14}~M_{\odot}$, including five very massive members whose stellar masses are in the range $(4-10) \times 10^{11}~M_{\odot}$. We do not find such high masses for our confirmed cluster members (with the exception of the RLAGN); our seven IRAC-detected sources have a median stellar mass of $1.1 \times 10^{11}~M_{\odot}$, and the remaining twelve non-IRAC detected members have masses $< 10^{10}~M_{\odot}$. Z12 derived SFRs from the [\ion{O}{2}] and H$\beta$ lines and find SFRs in the range $(20-40) ~M_{\odot} \rm~yr^{-1}$. Based on the admittedly crude [\ion{O}{3}] SFR indicator, we find similar SFRs for 9/16 of our SF cluster members. We also find higher SFRs, in the range $(40-140) ~M_{\odot} \rm~yr^{-1}$, for a significant number of SF cluster members (7/16). Note that our line detection limit imposes a SFR lower limit of $>20 ~M_{\odot}\rm ~yr^{-1}$ (similar to Z12 but based on a different line). According to the SFR/stellar-mass relation in \cite{Rodighiero11} for $1.5 < z < 2.5$ galaxies, our low-mass cluster members (typically all the non-IRAC detected members, with masses $\la 10^{10}~M_{\odot}$) are above the star-forming main-sequence towards starburst galaxies. Our detection limits prevent us from confirming main-sequence galaxies at these masses and redshifts. In Section \ref{p:sfr} we showed that confirmed emission line members of both structures imply total cluster SFRs of at least $\sim 400\, M_\odot\, {\rm yr}^{-1}$ within $\sim 500$~kpc of the cluster centers, which are assumed coincident with the target RLAGN. Based on {\it Spitzer} $24 \micron$ imaging of a sample of clusters from the {\it Spitzer}/IRAC Shallow Cluster Survey, \cite{Brodwin13} showed a steeply increasing SFR in cluster cores out to $z = 1.50$, with an average SFR of several hundred $M_\odot\, {\rm yr}^{-1}$ found in the cores of the highest redshift clusters in that study. \cite{Alberts14, Alberts16} find similar results based on longer wavelength {\it Herschel} data of a similar cluster sample. The results found here at $z \sim 2$ for CARLA~J2039$-$2514 and CARLA~J0800+4029 are consistent with those studies, though higher fidelity SFR indicators for these newly confirmed structures would be highly preferable to the current [\ion{O}{3}]-based values. Although G13 found no evidence for an already formed red sequence within $20\arcsec$ of the cluster core, they determined, from their best subsample of ($96$) cluster candidates comprising $14$ spectroscopic members and $82$ photo-$z$ candidates (expected to include $50$ interlopers), that $\sim(60-80)\%$ of candidates within $\sim 20\arcsec$ of the cluster core are passive regardless of mass, compared to $\sim 20\%$, $\sim 40\%$ and $\sim 60\%$ in the field for $\log M/M_{\odot} > 10, 10.5,$ and $11$, respectively (\citealp{Strazzullo13}). We additionally note that the structure reported in \cite{Spitler12} and \cite{Yuan14}, respectively discovered at $z=2.2$ from photometric redshifts and later spectroscopically confirmed at $\left<z\right>=2.095$ with $57$ members, comprises several $1\arcmin$ radius overdense groups covering a $12\arcmin \times 12\arcmin$ area in COSMOS (\citealp{Scoville07}). This structure has a slightly enhanced number of red galaxies for two groups compared to the field, with $N_{red}=0.5\pm0.2 \times N_{tot}$ compared to $0.2\pm0.03$ in the field. Our results are also indicative of the presence of passive CARLA candidates associated with CARLA J2039$-$2514, while CARLA J0800+4029 exhibits a number of passive candidates similar to the field in our comparison. A careful analysis accurately evaluating the passive fraction of galaxies in our structures will be addressed in Cooke et al. (in prep).\\ %%CONCLUSION%% | 16 | 9 | 1609.04162 |
1609 | 1609.04457_arXiv.txt | The formation of supermassive stars (SMSs) via rapid mass accretion and their direct collapse into black holes (BHs) is a promising pathway for sowing seeds of supermassive BHs in the early universe. We calculate the evolution of rapidly accreting SMSs by solving the stellar structure equations including nuclear burning as well as general relativistic (GR) effects up to the onset of the collapse. We find that such SMSs have less concentrated structure than fully-convective counterpart, which is often postulated for non-accreting ones. This effect stabilizes the stars against GR instability even above the classical upper mass limit $\gtrsim 10^5~\msun$ derived for the fully-convective stars. The accreting SMS begins to collapse at the higher mass with the higher accretion rate. The collapse occurs when the nuclear fuel is exhausted only for cases with $\dot M \lesssim 0.1~\msunyr$. With $\dot{M} \simeq 0.3 - 1~\msunyr$, the star becomes GR-unstable during the helium-burning stage at $M \simeq 2 - 3.5~\times 10^5~\msun$. In an extreme case with $10~\msunyr$, the star does not collapse until the mass reaches $\simeq 8.0\times 10^5~\msun$, where it is still in the hydrogen-burning stage. We expect that BHs with roughly the same mass will be left behind after the collapse in all the cases. | \label{sec:intro} Recent discovery of luminous quasars at $z > 6$ suggests the existence of black holes with mass exceeding $10^9~\msun$ when the age of the Universe was less than one billon years \citep[e.g.,][]{Mortlock11,Wu15}. The formation process of such supermassive black holes (SMBHs) is largely unknown, and poses a serious challenge to the theory of structure formation. The so-called direct collapse scenario \citep{BL03} provides an attractive pathway for the early BH formation via a peculiar mode of the primordial star formation. This scenario supposes the gravitational collapse of a primordial gas cloud forming a supermassive star with a mass of $\ga 10^5 \msun$. Specifically, the formation process proceeds in a two-step manner; a very small ($\ll M_{\sun}$) protostar first forms at the densest part of the cloud, and then grows in mass by rapid accretion from the surrounding envelope \citep[e.g.,][]{Latif13,Inayoshi14,Becerra15}. The accretion rate onto the protostar is expected to be as large as $\sim 0.1 - 1~\msunyr$, with which the stellar mass will exceed $\ga 10^5~\msun$ in less than the stellar lifetime. Such a supermassive star, if successfully formed, eventually collapses into a BH either by the core-collapse or by the general relativistic (GR) instability \citep[e.g.,][]{Iben63,Chandra64}. Given the massive remnant BH as a seed, a SMBH with more than $10^9~\msun$ can be assembled through further growth by accretion and/or mergers in less than one billion years \citep[e.g.,][]{DiMatteo12}. The size and the internal structure of an accreting protostar \citep[e.g.,][]{OP03,Ohkubo09} and its resulting radiative feedback effect, which potentially terminates the mass accretion \citep[e.g.,][]{McKee08, Hosokawa11, Hosokawa16, Stacy16}, are key ingredients to set the final stellar mass \citep[e.g.,][]{Hirano15}. Recent calculations show that the protostar structure is substantially modified if the accretion rate exceeds a few $\times 10^{-2}~\msunyr$; the star inflates to have an extended radius of $10-100$~AU \citep[e.g.,][]{Hosokawa12,Hosokawa13,Schleicher13}. Since such a fluffy protostar has low effective temperature and hence very low UV luminosity, the radiative feedback does not disturb the accretion flow at least until the stellar mass exceeds $\sim 10^5~\msun$ as long as continuous and efficient gas accretion ensues. In this {\it Letter}, we extend our stellar evolution calculations until the SMSs begin to collapse. Our calculations determine the final mass of the accreting SMSs by incorporating the realistic conditions within the context of the direct collapse scenario. Previous studies show that the GR instability causes the collapse of a non-rotating SMS with $\gtrsim$ a few $\times~10^5~\msun$ but without considering its evolution under rapid mass accretion \citep[e.g.,][]{Osaki66,Unno71,Fricke73,Fuller86}. We show that gas accretion significantly alters the stellar structure. With a very large accretion rate, a star continues to grow in mass without collapsing by the GR instability until the stellar mass reaches $\sim 10^6~\msun$. Our calculations show the final fate of such a SMS. | \label{sec:Discuss} We first compare our results with the previous work by \citet{Fuller86}, who study the GR stability of hydrogen-burning stars with fixed masses. They define a star $stable$ if the GR instability does not occur during the hydrogen burning stage. The critical stellar mass for the stability $M_{\rm crit}$ lies between $M_6$ = 0.1 and 0.25. However, our calculations show that all the models except for $\dot M=10~\msunyr$ are ``stable'' and begin collapsing after the exhaustion of the hydrogen at the center, when the stellar mass is well above $M_{\rm crit}$. Specifically, the star becomes GR-unstable during the helium-burning stage for $\dot M= 0.3 \sim 1.0~\msunyr$ with the final mass $M_6 = 0.19 \sim 0.35$. The main difference from \citet{Fuller86} is that accreting stars that we consider here have a high-entropy envelope. Hence the final masses (including the outer envelope) are larger than $ M_{\rm crit}$. The final stellar masses obtained in this {\it Letter} should be regarded as upper mass limits for the remnant BHs after the collapse. \citet{Chen14} argue that energetic supernova explosions occur at the collapse of SMSs in a very narrow mass strip around $M = 55000~\msun$. Since none of our models end up in this mass range, it is unlikely that such explosions are the final fate of our models. We will provide further results for this issue using a 1D-GR hydrodynamical code in a forthcoming paper. Preliminary results suggest that the BHs with approximately the same mass as the progenitors will be left behind after the collapse. If the progenitor star is a rapid rotator, some mass ejection or even explosion may occur at the final stellar collapse. An accretion disk can appear around a rotating BH and some outflows or jets can be launched. \citet{Shibata16}, for instance, show by way of a GR hydrodynamics simulation that some portion of the stellar envelope is ejected while most of the stellar matter is drawn into a newborn BH. We conclude that an accreting protostar at a very large rate of $\dot{M} = 0.1-10~\msunyr$ grows to become a SMS with mass $1.2-8.0 \times 10^5~\msun$. BHs with roughly the same mass will be left as remnants of such SMSs, which might have sown seeds for the formation of supermassive BHs in the early universe. { | 16 | 9 | 1609.04457 |
1609 | 1609.09355_arXiv.txt | A planet having protective ozone within the collimated beam of a Gamma Ray Burst (GRB) may suffer ozone depletion, potentially causing a mass extinction event to existing life on a planet's surface and oceans. We model the dangers of long GRBs to planets in the Milky Way and utilize a static statistical model of the Galaxy that matches major observable properties, such as the inside-out star formation history, metallicity evolution, and 3-dimensional stellar number density distribution. The GRB formation rate is a function of both the star formation history and metallicity; however, the extent to which chemical evolution reduces the GRB rate over time in the Milky Way is still an open question. Therefore, we compare the damaging effects of GRBs to biospheres in the Milky Way using two models. One model generates GRBs as a function of the inside-out star formation history. The other model follows the star formation history, but generates GRB progenitors as a function of metallicity, thereby favoring metal-poor host regions of the Galaxy over time. If the GRB rate only follows the star formation history, the majority of the GRBs occur in the inner Galaxy. However, if GRB progenitors are constrained to low metallicity environments, then GRBs only form in the metal-poor outskirts at recent epochs. Interestingly, over the past 1 Gyr, the surface density of stars (and their corresponding planets) that survive a GRB is still greatest in the inner galaxy in both models. The present day danger of long GRBs to life at the solar radius ($R_\odot=8$ kpc) is low. We find that at least $\sim$65\% of stars survive a GRB over the past 1 Gyr. Furthermore, when the GRB rate was expected to have been enhanced at higher redshifts, such as $z\gtrsim0.5$, our results suggest that a large fraction of planets would have survived these lethal GRB events. | With the success of extrasolar planet searches, and subsequent convergence on the fraction of stars that may host Earth-size planets, there is a growing interest in the effects of galactic environments on planetary biospheres~\citep{1995ApJ...444L..53T,2003ApJ...585.1169G,2004Sci...303...59L,2004IJA:240775,2005ApJ...634..509T,2005ApJ...622L.153T,2011AsBio..11..343M,2011AsBio..11..855G,2013RMxAA..49..253C,2013AsBio..13..491J,2013-637X-773-1-6,2014PhRvL.113w1102P,2015AsBio..15..207T,2015ApJ...810L...2D,2015ApJ...810...41L,IJA:10152160,Vukotic2016,2016ApJ...826L...3T,2016A&A...592A..96G}. One class of potential risks to the habitability of planets are transient radiation events, such as Gamma Ray Bursts (GRBs). These events are lethal to planets due to the subsequent depletion of ozone in planetary atmospheres. Furthermore, given that GRBs have collimated emissions that beam radiation on the order of a kpc, they are expected to pose a significant danger to life on planets in general, and Earth in particular. As such, these events may have been responsible for mass extinction events on the Earth~\citep{1995ApJ...444L..53T,2004IJA:240775,2005ApJ...622L.153T,2005ApJ...634..509T,Melott2009,2013-637X-773-1-6,2015AsBio..15..207T}. To estimate the effects that GRBs have on the habitability of the Milky Way, determination of the GRB rate from observational and statistical constraints is necessary. Applying the cosmologically local GRB rate to small volumes is a challenge, as many of the salient characteristics of GRBs, such as their progenitors and environments that give rise to the events are still not well understood. However, GRBs are found in metal-poor host galaxies~\citep{2006Natur.441..463F,2013ApJ...773..126J}, and thus correlate with low metallicity environments. One explanation for this correlation is the collapsar model of GRB formation, where low metallicity, massive helium stars are the progenitors of long GRBs~\citep{1999ApJ...524..262M}. Therefore, there is a metallicity bias, where environments that have undergone significant chemical evolution produce fewer GRBs than low metallicity environments. With the advanced chemical evolution of the Milky Way, it is expected that long GRBs are rare in the Galaxy. However, the degree to which metallicity quenches GRB formation is still uncertain when applying the locally observed rate to the Galaxy. Studies have suggested that the GRB rate is proportional to the star formation rate (SFR)~\citep{1997ApJ...486L..71T,2000MNRAS.312L..35B,2002ApJ...575..111B,2007ApJ...661..394L}; thus, disregarding metallicity bias, the GRB formation history follows the SFR history. However, this assumption has been challenged, partially due to the environments that GRBs are found~\citep{2006Natur.441..463F}. Despite varying explanations of the GRB formation history in the literature, using the SFR to trace the GRB rate is a reasonable approximation to understanding the GRB formation history in the Milky Way. Furthermore, this assumption yields a GRB formation history that follows the inside-out formation history of the Milky Way. This is an important effect to capture, as the majority of the stars in the early Galaxy were found within smaller galactocentric radii than at the present day. Thus, GRBs would have been more lethal to a greater fraction of the overall stars at high redshift than at the present day. Previous work on the Galactic Habitable Zone~\citep{2011AsBio..11..855G} considered the effects of supernovae on planetary biospheres. An interesting result was that the region with the greatest stellar density (and supernova rate) was found to host the greatest number of habitable planets, at a galactocentric radius of $R\approx2.5~\kpc$. While the fraction of stars that are nearby a supernova event is much higher at $R\approx2.5~\kpc$ than in the outskirts, or solar neighborhood, the comparatively higher stellar density of the inner Galaxy, and the average age of stars in the region outweigh the negative effects of supernovae. Another aim of the present study is to observe whether the same phenomenon holds true for GRBs. We utilize a model of the Milky Way that considers the inside-out star formation history, chemical evolution, 3-dimensional stellar number density distribution, and determination of the absolute GRB rate of the Galaxy at $z\sim0$. Using these properties, we model the collimated jet emission, and the influence of chemical evolution on GRB formation to examine the effects that long GRBs have on the habitability of the Galaxy. The GRB formation rate is a function of both metallicity and the star formation history (SFH)~\citep{2014ApJS..213...15W}; however, it is unclear to what extent metallicity evolution reduces the GRB rate over time and its influence on the location of GRB progenitors throughout the galactic disk. Therefore we compare two scenarios: (a) where GRB formation is a function of the SFH; and (b) GRB formation follows the SFH where there is a metallicity dependence on GRB progenitors that favor low metallicity host environments. The paper is outlined as follows. In Section~\ref{sec:methods}, we outline the construction of a model of the Milky Way Galaxy that includes a stellar population consistent with major observable properties. Additionally, we describe the properties of the GRBs, including the beamed emission and two formation history scenarios. Section~\ref{sec:results} illustrates the results, including two metrics of habitability: 1) the fraction, and 2) the surface density of stars that are within the beam of a GRB over a time period. Additionally, we compare these results to the relevant literature. Finally, we conclude the work in Section~\ref{sec:conclusions}. | \label{sec:conclusions} We examine the lethality of GRBs in the Milky Way by modeling the 3-dimensional stellar number density distribution, inside-out formation history, and chemical evolution of the Galaxy. We have proposed two models of the GRB formation history that differ as a function of the metallicity required of a GRB progenitor. When we compare the two models, we find that when the GRB rate is a function of the SFH, the rate is roughly consistent over time, favoring GRBs that are located in the inner galaxy over all epochs. Whereas if we include the constraint that assumes GRBs can only form in low metallicity environments, the outskirts primarily hosts GRBs at recent epochs. Due to declining stellar density with increasing galactocentric radius, low redshift GRBs located at the outskirts sterilize on average far fewer planets than GRBs at lower radii. Therefore the metallicity dependent progenitor model suggests that GRBs do not pose a significant danger to biospheres in the Milky Way at the present day. Even if the GRB formation rate only follows the SFH, we still find that over the past 1 Gyr, the inner Galaxy hosts the greatest density of stars (and associated planets) that survive GRB events. Thus, the region of the Milky Way with the greatest luminosity is the most favorable for life. These calculations suggest that GRBs may be less lethal than previous estimates~\citep{1995ApJ...444L..53T,2014PhRvL.113w1102P,2015ApJ...810...41L}. Additionally, studies of progenitor environments suggest that GRBs are unlikely in the Milky Way~\citep{2006Natur.441..463F}, and GRBs have only been found in galaxies having $<10^{10}$M$_{\odot}$~\citep{2013ApJ...773..126J}, which excludes the Milky Way. The results in this work, and others found in the literature imply that GRBs may be uncommon in the Milky Way and may not pose a significant danger to the propensity of planets to host life in the Galaxy. | 16 | 9 | 1609.09355 |
1609 | 1609.05667_arXiv.txt | RW Aur is a young binary star that experienced a deep dimming in 2010-11 in component A and a second even deeper dimming from summer 2014 to summer 2016. We present new unresolved multi-band photometry during the 2014-16 eclipse, new emission line spectroscopy before and during the dimming, archive infrared photometry between 2014-15, as well as an overview of literature data. Spectral observations were carried out with the Fibre-fed RObotic Dual-beam Optical Spectrograph on the Liverpool Telescope. Photometric monitoring was done with the Las Cumbres Observatory Global Telescope Network and James Gregory Telescope. Our photometry shows that RW Aur dropped in brightness to R = 12.5 in March 2016. In addition to the long-term dimming trend, RW Aur is variable on time scales as short as hours. The short-term variation is most likely due to an unstable accretion flow. This, combined with the presence of accretion-related emission lines in the spectra suggest that accretion flows in the binary system are at least partially visible during the eclipse. The equivalent width of [O\,I] increases by a factor of ten in 2014, coinciding with the dimming event, confirming previous reports. The blue-shifted part of the $H\alpha$ profile is suppressed during the eclipse. In combination with the increase in mid-infrared brightness during the eclipse reported in the literature and seen in WISE archival data, and constraints on the geometry of the disk around RW Aur A we arrive at the conclusion that the obscuring screen is part of a wind emanating from the inner disk. | Classical T Tauri stars (TTS) are pre-main-sequence objects with strong irregular variability, first noted in a study of highly variable stars by \cite{1945ApJ...102..168J}. In that same work, RW Aur is identified as a T Tau type star with a spectrum rich in emission lines. Today, it is known that RW Aur is a visual binary with a separation of 1.4\arcsec \,between the primary and secondary components. The spectrum of the primary shows strong signatures of accretion and winds (eg. \citealt{2001A&A...369..993P,2005A&A...440..595A}). The secondary, RW Aur B, is a late K type TTS \citep{1993AJ....106.2005G,1997ApJ...490..353G,2004ApJ...616..998W} with little or no evidence of ongoing accretion. There are indications for the existence of more stellar/substellar companions in the system. The multiplicity study of T Tau stars by \cite{1993AJ....106.2005G} reveals a possible tertiary component (RW Aur C) orbiting RW Aur B, with a separation of 0.12\arcsec\,between the B and C stars. Furthermore, spectroscopic studies by \cite{1999A&A...352L..95G} and \cite{2001A&A...369..993P} reveal periodic variations in the radial velocity of a number of lines in RW Aur A. The authors explore the possibility of a brown dwarf sized companion (RW Aur D) in close orbit around the primary as the cause for the observed spectroscopic variations. % \cite{2006A&A...452..897C} examined CO maps of RW Aur from interferometry. They estimated the size of the circumstellar disk around RW Aur A to be unusually small (40-57\,AU) and found evidence for a 600\,AU long structure of material trailing from the primary star. The disk size was calculated for a disk inclination between 45-60$^{\circ}$. Both the size of the disk and the $``$arm$"$ were suggested as evidence for the tidal disruption of the circumstellar disk from a recent interaction with the secondary RW Aur B. This tidal interaction hyphothesis is further backed up by recent hydrodynamical simulations of the RW Aur disk by \cite{2015MNRAS.449.1996D}. In 2010 RW Aur underwent a long lasting dimming event, that had never been observed before. The star dropped by $\sim$2 mag in the optical and remained in that state for several months. \cite{2013AJ....146..112R} presented a photometric investigation of this event, looking at KELT and American Association of Variable Star Observers (AAVSO) data. They conclude that the most likely explanation of this event is obscuration by material located $\sim$180\,AU from RW Aur A. It is noted that the estimated occulting body lies far beyond the outer edge of the disk, but is still within the size of the tidally disrupted arm and therefore likely a part of it. In 2014 RW Aur entered a second long lasting, even deeper ($\sim$3 mag, \cite{2015IBVS.6126....1A}) minimum. Figure \ref{aavso} shows an AAVSO visual lightcurve of RW Aur up until August 2016, including both dimming events. As can be seen in this lightcurve, there are clear signs that the star has recovered to its out-of-eclipse brightness in August 2016, indicating the end of the second eclipse. Our own photometry on Aug 25 and 26 confirms this, with R-band magnitudes of 10.7 and 10.6 \citep{2016ATel.9428....1S}. \cite{2015A&A...577A..73P} performed a spectroscopic analysis of emission lines associated with accretion, wind and jets in and out of the 2014 minimum. They found enhancement in the equivalent widths of [O\,I] and [S\,II] (forbidden lines associated with jets), but no increase in the fluxes. Additionally, an increase in the strength of the resonant lines Ca\,II and Na\,I was reported. Furthermore, they found no change in the emission of H$\alpha$ and He\,I, lines associated with accretion activity. The authors concluded that the obscuring body only covers the star and inner regions of the system, while the outer parts of the wind and jets remained unaffected. This effect was attributed to dust grains being lifted up into the line of sight through interactions with the winds. In addition, \cite{2015IBVS.6143....1S} report results of infrared photometry in JHKML in the period 2010-2015. They find a drop in JHK brightness corresponding to the 2014 dimming event, but a simultaneous increase in brightness in M and L. The excess IR emission is attributed to hot dust of about 1000\,K, located around the inner disk rim, 0.1-0.2\,AU away from the star. It is further speculated that the hot dust could then be lifted up into the wind, supporting the scenario presented in \cite{2015A&A...577A..73P}. The dusty wind idea is in contrast to the tidal arm dimming scenario \citep{2013AJ....146..112R,2016AJ....151...29R} where an occultation by the tidal arm is expected to at least partially obscure the inner disk and possibly cover the excess IR emission. \cite{2016AJ....151...29R} also report a smaller dimming in 2012-13 and argue that the kinematics of this event fit the values derived from their analysis of the 2010 eclipse. Their analysis of the 2014 dimming using the same method is hampered by the fact that the ingress for the eclipse happened during the seasonal observing gap and its duration is not constrained yet. The tidal arm dimming scenario is potentially supported by the hydrodynamical simulations by \cite{2015MNRAS.449.1996D}. Their simulations suggest the presence of material in a bridge between RW Aur A and B, with several particles also crossing the line-of-sight to the system. Finally, \cite{2016ApJ...820..139T} present a spectroscopic study of the system between Oct 2010 and Jan 2015 and argue that neither the tidal arm nor dusty wind scenario alone can account for the complex variability seen in the lines in and out of the eclipses. The authors further explore time-variable mass accretion as an alternative explanation to the observed dim and bright states of the system within the 2010-15 time frame. In this study we present photometric and spectroscopic observations of RW Aur taken just before and during the 2014 dimming. We explore the photometric variability in several optical and infrared bands along with equivalent width and profile analysis of the H$\alpha$, He\,I and [O\,I] lines to provide further evidence to the above scenarios. \begin{figure} \includegraphics[width=1.0\linewidth]{plots/rwaur_aavso.png} \caption{American Association of Variable Star Observers (AAVSO) visual lightcurve for RW Aur between 2007 and 2016. The 2010 and 2014 dimmings are clearly visible.} \label{aavso} \end{figure} | We present new spectroscopic and photometric observations of RW Aur during the dimming in 2014-2016. In addition, we review the previous and new observational constraints for the system. \subsection{Photometric variations} The dimming event that started in 2014 is continuing through our observing period. The system has become even fainter in the optical during the 2015-16 winter observing season (see Figures \ref{aavso} and \ref{LCOGTlc}). In addition to the long-term dimming trend, we clearly see irregular short timescale variations in the JGT (Figure \ref{JGTlc}) and LCOGT (Figure \ref{LCOGTlc}) data. These variations could be related to small scale changes in extinction along the line of sight. However, the irregular components in the morphology of the JGT lightcurves better resemble burst-like events. Such morphology is typically associated with changes in an unstable accretion flow for T Tau stars (see \citealt{1994AJ....108.1906H,2014AJ....147...82C} for an overview of the different types of photometric variability displayed by T Tau stars). Being able to see such short time-scale irregular variations in RW Aur implies that the accretion flow in the system remains at least partially visible during the dimming. In addition to the overall drop in photometric flux in the optical, colour-magnitude diagrams of our LCOGT data (Figure \ref{CMDplots}) reveal a significant long-term reddening trend since the start of the eclipse. We compare the observed reddening with models for interstellar extinction by \cite{1989ApJ...345..245C}. Extinction fails to fit the observed colours simultaneously. In particular, the model under-predicts the observed u\,\arcmin-band amplitude and is unable to reproduce the u\,\arcmin$-$\,g\,\arcmin reddening slope. We therefore conclude that the long-term reddening trend is not consistent with interstellar-like extinction. We also attempt to explain the reddening with a simple hot spot model, where the star and spots are approximated by black bodies and the spot temperature and filling factor are varied in order to fit the observed lightcurve amplitudes (eg. \citealt{2009MNRAS.398..873S}, \citealt{1995A&A...299...89B}). We find a number of parameter values that produce amplitudes and reddening slopes similar to the observed ones, especially in the red portion of the spectrum. However, hot spots have one major problem with explaining the observed long-term reddening. Hot spots form at the stellar surface as a result from accreting material near-free falling onto the star and creating shocks. As the star rotates the spots come in and out of view, resulting in a photometric variably on time-scales comparable with the stellar rotational period (typically several days for T Tau stars). This is in contrast to the long-term dimming of RW Aur which has lasted for months. While accretion may still play a role for the short-time scale variability of RW Aur, hot spots cannot be the cause for the observed long-term reddening, where the system has continuously been getting redder for months. An alternative explanation for the observed long-term reddening arises from the fact that our photometry does not resolve RW Aur A and B. \cite{2015IBVS.6126....1A} present resolved photometry of the system. Their measurements (after the start of the dimming) show an R-band magnitude of 13.08 and 11.97 for the A and B-components respectively. Our JGT photometry yields R$\sim$12 mag suggesting that the red part of the spectrum in our observations is most likely dominated by the B-component. We can therefore interpret the observed long-term reddening as the A-component becoming faint enough that the B-component becomes the dominant source of flux, making the entire system appear redder while simultaneously getting dimmer as we lose flux from RW Aur A. % \begin{figure} \includegraphics[scale=0.44]{plots/cmd_ug.png} \includegraphics[scale=0.44]{plots/cmd_gr.png} \includegraphics[scale=0.44]{plots/cmd_ri.png} \\ \caption{ Colour-magnitude diagrams constructed from all LCOGT observations. The epochs of observations are colour coded. The black solid line indicates the best linear fit to the data with thinner lines indicating the error in the slope estimation. The red arrow is the interstellar extinction slope for $R_V = 3.1$ and $A_V = 0.9$, where the $A_V$ value is selected such that the resulting extinction in the r-band best matches the observed amplitude in the r-band lightcurve.} % \label{CMDplots} \end{figure} \subsection{Spectral signatures} Spectroscopy provides more information about accretion and outflow in young stellar systems. In the classical magnetospheric accretion scenario the inner-most region of the circumstellar disk is threaded by stellar magnetic field lines. Material is channelled from the disk onto the star forming accretion columns at near free-fall velocities, with a shock forming at the surface of the star. The shock and accretion columns generate H$\alpha$ emission. The accretion-related H$\alpha$ emission lines are typically broad with equivalent widths of over 10\,$\AA$. Other emission lines related to CTTS accretion include permitted lines of He\,I, O\,I and Ca\,II. For a review of accretion signatures we refer the reader to \cite{2005ApJ...626..498M} and references therein. In addition to accretion indicators, CTTS often shows outflow-related forbidden line emission such as [O\,I] lines \citep{1995ApJ...452..736H,1984A&A...141..108A}. Our spectra show that the EW of the [O I] (6300\,$\AA$) line increases by a factor of 10 after the eclipse (Figure \ref{OISeries},\ref{HA_OI_EW}). The observed EW increase is in agreement with measurements by \cite{2015A&A...577A..73P} and \cite{2016ApJ...820..139T}. \cite{2015A&A...577A..73P} attribute the change in the forbidden line's EW to a drop in the photospheric continuum, as reported by the resolved photometry of RW Aur A by \cite{2015IBVS.6126....1A}. \cite{2015A&A...577A..73P} further imply that the obscuring body must be blocking the inner parts of the system while the outflow remains visible. The He I (6678\,$\AA$) and H$\alpha$ emission lines remain clearly visible in our spectra after RW Aur enters the eclipse. This again indicates that at least part of the accretion flow is still visible. The He I EW remains constant (Figure \ref{HeEW}), implying the He-emitting regions are obscured in the same way as the photosphere. The H$\alpha$ EW decreases after the eclipse (Figure \ref{HA_OI_EW}). Our EW measurements show lower values for the line during the dimming compared to the values reported by both \cite{2015A&A...577A..73P} and \cite{2016ApJ...820..139T}. We attribute this to the fact that our spectroscopy does not resolve the A and B-components, hence we get an additional flux in the continuum from RW Aur B, resulting in a lower EW value. The H$\alpha$ line profiles show a suppressed blue-shifted peak after the start of the dimming (Figure \ref{HaSeries}). This can also be seen in the H$\alpha$ profiles from \cite{2016ApJ...820..139T} from both the 2010 and 2014 dimmings. This blue-shifted absorption can be interpreted as the presence of an absorbing material moving towards the observer's line-of-sight, i.e. an outflow. \subsection{Origin of the eclipse} In Sect. 1 we introduce two possible explanations for the long dimming event observed in RW Aur. The obscuring screen could either be a tidal arm in the outer disk, as suggested by \cite{2013AJ....146..112R,2016AJ....151...29R}, or a wind emanating in the inner disk, suggested by \cite{2015A&A...577A..73P}. Our observations provide new evidence to support the occultation by a wind. In particular, we show that the blue-shifted portion of the H$\alpha$ line is suppressed since the beginning of the eclipse, indicating the presence of outflowing and thus blueshifted gas with velocities of 100-200\,kms$^{-1}$ in the occulting screen. This scenario is further supported by the increase in the mid-IR fluxes of the system in eclipse, reported by \cite{2015IBVS.6143....1S} and visible in archive NEOWISE data in the W2 band (Figure. \ref{neowise}). \cite{2015IBVS.6143....1S} fit a blackbody to the spectral energy distribution of RW Aur in its faint state and find that the excess IR emission is plausibly explained by the presence of hot dust with temperature around 1000\,K along the line of sight. Thus, the occulting body is outflowing and contains hot dust, which fits with a scenario where a wind emanating from the inner disk eclipses the star RW Aur A. These two facts together are also inconsistent with the tidal arm scenario, in which the occulting material would not be expected to move rapidly along the line of sight and the dust in the screen would be cold. If the eclipse indeed arises in a wind, the long duration of the current event would imply that the outflowing screen stretches azimuthally over a large portion of the inner disk. We note that the recent resolved X-ray detection of RW Aur shows a slight soft X-ray extension of RW Aur A along the blue-shifted jet axis close to the star \citep{2014ApJ...788..101S} - this could be related to the obscuring screen discussed here. Furthermore, \cite{2015A&A...584L...9S} present resolved X-ray and NIR observations of the system in eclipse and provide further evidence for the presence of an obscuring screen with hot material located close to the star. In addition, there is also a geometrical argument in favour of a blocking screen located above the inner disk. The CO map modelling from \cite{2006A&A...452..897C} found an inclination between 45-60$^{\circ}$ for the large-scale disk. The inclination of the jet to the line of sight is measured to be $46\pm 3^{\circ}$ \citep{2003A&A...405L...1L}. These values indicate that the large-scale disk is far from edge-on, which makes it difficult to lift material in the outer disk into the line of sight. In fact, the hydrodynamical simulations of the tidal arm scenario by \cite{2015MNRAS.449.1996D} give a best fitting value of 64$^{\circ}$ for the inclination, and thus does not match the constraints from observations. On the other hand, the modelling of K-band Keck interferometric observations tracing the continuum and gas in the inner region of the disk results in an inclination angle of 75$^{\circ}$ (\citealt{2014MNRAS.443.1916E}, Table 6) for an emitting region within a fraction of 1\,AU. This is a significantly higher inclination and closer to edge-on than for the outer disk. We note that some of the Keck interferometer data for RW Aur A was taken during the first eclipse in 2010/11. The discrepancy in the inclinations between inner and outer disk may be evidence for a large-scale warp in the disk, possibly caused by interaction with the companion RW Aur B. Thus, RW Aur B may have an effect on the disk, but it is unlikely to be the direct cause of the feature that eclipses RW Aur A. Adopting the value from the inner disk for the disk inclination, the material in the occulting screen needs to be vertically displaced from the disk midplane by at least 0.26 times the distance from the central object. This value is comparable to the maximum vertical extent of the inner disk found in the so-called $``$dippers$"$, T Tauri stars with periodic eclipses caused by disk warps caused by accretion channelled by a magnetic field tilted with respect to the rotational axis \citep{1999A&A...349..619B,2015A&A...577A..11M}. RW Aur is not known to be a dipper, perhaps because the tilt of its magnetic field is not sufficient. Based on the high value for the inner disk inclination, however, it is conceivable that a magnetically driven disk wind can lift material into the line of sight. The jet of RW Aur A has been extensively studied with high-resolution imaging. It is one of the densest among the T Tauri stars investigated in detail, but the mass outflow rate seems to be stable within an order of magnitude along the jet axis and in both lobes at $10^{-9}$ to $10^{-8}\,M_{\odot}$yr$^{-1}$ \citep{2009A&A...506..763M}. However, in the past there have been clear indications for strong variability in the forbidden lines caused by the wind as well as in the blue-shifted portion of H$\alpha$ \citep{2005A&A...440..595A}. In the standard scenario of magnetospheric accretion, the wind is thought to be coupled to the accretion, thus, we would expect that a change in the outflow rate or geometry, as indicated by our interpretation of the RW Aur dimming, is preceded by strong variability in the accretion flow. As far as we are aware, there is no evidence for a recent accretion burst in RW Aur. Its accretion rate is relatively high ($10^{-7}$ to $10^{-6}\,M_{\odot}$yr$^{-1}$) compared with other CTTS with similar mass, but seems to be relatively stable (e.g., \citealt{2014MNRAS.440.3444C}). Thus, changes in the wind characteristics might not always be coupled to changes in the accretion properties. One possible trigger for a wind outburst could be an interaction between the disk and the putative stellar or substellar spectroscopic companion to RW Aur A \citep{1999A&A...352L..95G,2001A&A...369..993P}. However, it is unclear how such an interaction can produce the sudden onset of wind activity multiple times over the course of the last $\sim$\,7 years, with no precedents observed in the previous $\sim$\,50 years. Another alternative may be that the fly-by of RW Aur B has caused an instability in the disk of RW Aur A, although modelling this scenario is beyond the scope of this study. We encourage further multi-band photometric and spectroscopic monitoring of accretion and outflow signatures of RW Aur to further test and solidify our understanding of the system. | 16 | 9 | 1609.05667 |
1609 | 1609.07512_arXiv.txt | Much of the inner Milky Way's (MW) global rotation and velocity dispersion patterns can be reproduced by models of secularly-evolved, bar-dominated bulges. More sophisticated constraints, including the higher moments of the line-of-sight velocity distributions (LOSVDs) and limits on the chemodynamical substructure, are critical for interpreting observations of the unresolved inner regions of extragalactic systems and for placing the MW in context with other galaxies. Here, we use SDSS-APOGEE data to develop these constraints, by presenting the first maps of the LOSVD skewness and kurtosis of metal-rich and metal-poor inner MW stars (divided at ${\rm [Fe/H]} = -0.4$), and comparing the observed patterns to those that are seen both in $N$-body models and in extragalactic bars. Despite closely matching the mean velocity and dispersion, the models do not reproduce the observed LOSVD skewness patterns in different ways, which demonstrates that our understanding of the detailed orbital structure of the inner MW remains an important regime for improvement. We find evidence in the MW of the skewness-velocity correlation that is used as a diagnostic of extragalactic bar/bulges. This correlation appears in metal-rich stars only, providing further evidence for different evolutionary histories of chemically differentiated populations. We connect these skewness measurements to previous work on high-velocity ``peaks'' in the inner Galaxy, confirming the presence of that phenomenon, and we quantify the cylindrical rotation of the inner Galaxy, finding that the latitude-independent rotation vanishes outside of $l \sim 7^\circ$. Finally, we evaluate the MW data in light of select extragalactic bar diagnostics and discuss progress and challenges of using the MW as a resolved analog of unresolved stellar populations. | \label{sec:intro} Boxy/peanut bulges and bars dominate the inner regions of a large fraction of galaxies --- at least $\sim$45\% of massive disk galaxies \citep[e.g.,][]{Lutticke_2004_BPbulgestatistics,Buta_2015_S4Gmorphologies}. These structures contain fossil records of the numerous events --- internal evolution and external interactions, both stochastic and secular events --- that have occurred throughout the lifetime of the galaxy \citep[see ][for a recent review]{Kormendy_2016_ellipticals-and-bulges}. However, both dust extinction and perspective pose difficulties to interpreting these mostly-planar structures in edge-on systems. The Milky Way (MW) presents a beautiful, close example of an edge-on boxy bulge that we can use not only to retrace the evolutionary history of our home galaxy, but also to characterize extragalactic edge-on bars and their impact on galactic properties. We can test the accuracy of extragalactic bar diagnostics using independent measurements of resolved stars, and understand the diversity of populations that can remain blended or hidden in those diagnostics of unresolved stars. Our understanding of the inner regions of the MW --- the boxy bulge or bar, and inner few kpc of the disk --- has been revolutionized in the past several years by multiple surveys probing the large number of stars needed to measure the mean kinematical and chemical properties of the highly extincted stellar populations \citep[e.g., BRAVA, ARGOS, GIBS, APOGEE;][respectively]{Rich_11_brava,Freeman_2013_argos,Zoccali_2014_GIBS1,Majewski_2015_apogeeoverview}. The stellar population is dominated by metal-rich stars entrained in a thick bar with half-length $R_b \sim 2$~kpc, beyond which a thin planar bar component (with $R_b \sim 4.5$~kpc) is visible \citep[the ``long bar''; e.g.,][]{Zasowski_2012_innerMW,Benjamin_05_glimpse,Wegg_2015_RClongbar}. The size, or even presence, of a ``classical'' bulge, resulting from the hierarchical merging of protogalaxies in the early days of the MW or from in situ star formation during this same phase, remains a significant unknown \citep[e.g.,][]{Shen_10_purediskbulge,Robin_2012_besanconbarmodel}. See Section~4 of \citet{BlandHawthorn_2016_theMWproperties} for a recent review of our understanding of the inner MW. \begin{figure*}[] \begin{center} \includegraphics[trim=1.4in 1.7in 1.3in 4.0in, clip, angle=180, width=\textwidth]{f1.pdf} % \end{center} \caption{ Distribution of stars used in this analysis. (a) Positions and approximate sizes of the APOGEE fields (blue circles). In grayscale is a smoothed [4.6$\mu$] WISE image of this region \citep{Lang_2014_unWISEcoadds}. Both the boxy shape of the bulge, and its asymmetry around the minor axis ($l=0^\circ$), are visible. (b) Distribution of APOGEE stars against the WISE contours from panel (a), colored by the weighted mean velocities of neighboring stars. } \label{fig:WISE_datamap} \end{figure*} Empirically, the boxy bulge/bar displays roughly cylindrical rotation \citep[e.g., as seen by the BRAVA and ARGOS surveys;][]{Kunder_2012_bravaDR,Ness_2013_argoskinematics}. \citet[][hereafter Paper I]{Ness_2016_apogeekinematics} used APOGEE data to map this mean rotation pattern and the velocity dispersion, which peaks in the center and falls off smoothly into the kinematically colder disk. However, these properties are more correctly stated as describing the dominant metal-rich stellar populations; the more metal poor stars show slightly different behavior \citep[slower rotation and higher velocity dispersion, known since, e.g.,][see also Paper I and \citet{Kunder_2016_bulgeRRLbrava}]{Harding_1993_bulgekinematics}. These stars comprise a relatively small, but potentially very interesting, fraction of the total population \citep[e.g., $<$5\% have ${\rm [Fe/H]} \lesssim -1$;][A.~E.~Garc\'ia~Perez, in preparation]{Ness_2016_bulgeMDF}. Numerous $N$-body models of the inner Galaxy have been developed and scaled to reproduce its mean rotation and velocity dispersion and begin to explain the evolutionary history that has resulted in today's stellar chemodynamical patterns \citep[e.g.,][]{Athanassoula_2007_barmodel,Shen_10_purediskbulge,MartinezValpuesta_2013_barmodel}. Among the most important outstanding issues to address, through the combination of observations and these models, are the mass fraction and dynamical impact of the long bar, the quantity and relevance of chemodynamical diversity in the inner Galaxy, the mass contribution of the inner halo and any classical bulge, and the relationships among kinematical, chemical, and morphological structures. To address these questions --- to constrain and discriminate among models --- we need more than the mean rotation and dispersion patterns that can be well-matched by models with a wide variety of initial conditions and histories. The {\it shape} of the kinematical distributions matters, along with the dependence of this shape on chemistry. In this paper, we present the skew and kurtosis (the shape parameters) of the line-of-sight velocity distributions for $\sim$19,000 stars observed predominantly in the inner $\sim$4~kpc of the MW. We discuss the implications for the claims of high velocity ``peaks'' in the inner Galaxy \citep{Nidever_2012_apogeebar} and compare these empirical results to those of MW-scaled $N$-body models and to common extragalactic bar diagnostics. We also examine the issue of cylindrical rotation in similar terms as it has been quantified in external galaxies \citep[e.g.,][]{Molaeinezhad_2016_cylindricalBProtation}. All of these are critical steps towards properly using the MW as a laboratory for understanding galaxy evolution at large. | \label{sec:finis} Understanding the finely detailed behavior of stars in the inner MW, including the relationships between their chemical and kinematical patterns, is critical not only to retracing the evolutionary history of the MW, but also to interpreting the unresolved inner regions of external galaxies and placing the MW in context with those systems. Along with the line-of-sight velocity mean and dispersion, we present here the first maps of the velocity skewness and kurtosis of inner Galaxy stars, independently for relatively metal rich (${\rm [Fe/H]} > -0.4$) and metal poor (${\rm [Fe/H]} \ge -0.4$) stars. We find patterns that, based on comparison to $N$-body simulations, are consistent overall with the metal rich stars being more representative of stars entrained in the bar --- e.g., an increase in the velocity skewness at the projected end of the bar --- and the metal poorer stars being kinematically hotter and less bar-like. We revisit the issue of the high velocity ``peaks'' in the inner Galaxy and verify the initial detections \citep{Nidever_2012_apogeebar}, but we conclude that the velocity skewness is a significantly more robust way to quantify this phenomenon than the fitting of discrete components to the velocity distribution. This conclusion is supported by the chemical similarity of the ``high velocity'' stars to the rest of the sample, which indicates these are not distinct and separable populations. We explore the extent of cylindrical rotation in our sample and find an apparent break in the pattern between $l=7^\circ$ and $l=10^\circ$, near the end of the boxy bulge, as expected. Throughout this work, we also discuss the advantages, disadvantages, and caveats of using MW data to interpret unresolved stellar behavior in external galaxies. Some qualitative diagnostics proposed to identify end-on bars in these galaxies are compared to the patterns we see in the MW, which we know from many lines of evidence to show a nearly end-on bar towards the Sun. These diagnostics include the midplane integrated flux profile and the correlation between mean velocity and velocity skewness. We find good agreement in these qualitative comparisons, and in future work, we will explore these questions quantitatively, framing the MW as an unparalleled Rosetta Stone for the signatures of chemodynamical evolution observed in galaxies throughout the Universe. | 16 | 9 | 1609.07512 |
1609 | 1609.08957.txt | Apparent exponential surface density profiles are nearly universal in galaxy discs across Hubble types, over a wide mass range, and a diversity of gravitational potential forms. Several processes have been found to produce exponential profiles, including the actions of bars and spirals, and clump scattering, with star scattering a common theme in these. Based on reasonable physical constraints, such as minimal entropy gradients, we propose steady state distribution functions for disc stars, applicable over a range of gravitational potentials. The resulting surface density profiles are generally a power-law term times a S\'{e}rsic-type exponential. Over a modest range of S\'{e}rsic index values, these profiles are often indistinguishable from Type I exponentials, except at the innermost radii. However, in certain parameter ranges these steady states can appear as broken, Type II or III profiles. The corresponding velocity dispersion profiles are low order power-laws. A chemical potential associated with scattering can help understand the effects of long range scattering. The steady profiles are found to persist through constant velocity expansions or contractions in evolving discs. The proposed distributions and profiles are simple and solve the stellar hydrodynamic equations. They may be especially relevant to thick discs, which have settled to a steady form via scattering. | Spiral galaxies have been known to have exponential radial profiles for a long time (\citealt{pa40}, \citealt{de59}, \citealt{fr70, fr07}, \citealt{va02}), with scale lengths that are independent of Hubble type for early and intermediate types \citep{de96}. The early observations did not extend to very faint surface brightnesses, nor over many scale lengths. Recent observations have gone much deeper (e.g., \citealt{bl05} \citealt{ga09}, \citealt{er05, er08}, \citealt{po06}, \citealt{he13}, and \citealt{zh15}), showing the continuation of exponential form over about 10 scale lengths in some cases. As originally noted by \citet{fr70} the radial profiles often have a break, either turning downward in the outer parts (Type II) or upward (Type III) (also see \citealt{er05, er08}, \citealt{po06}, \citealt{he13}). However, the slope changes are often modest, and both inner and outer profiles are well fit by exponential forms. In addition to large disc galaxies, the surface density profiles of dwarf Irregular galaxies, which have little shear and generally no spiral waves, also follow exponential profiles out to 6 or more scale lengths (e.g., to $31\ mag\ arcsec^{-2}$ in V-band; \citealt{hu11}). The dwarfs are easily harassed by encounters with other galaxies, subject to continuing gas accretion (\citealt{va98}, \citealt{wi98}), or cycles of gas expulsion and reassertion. Thus, the stars must continuously migrate to smooth out profile disturbances, and generally do so without the aid of spirals, bars or shear. These observations suggest that exponentials are the generic surface density forms for the full range of two-dimensional galaxy components, and that these profiles extend over a huge range of surface brightness. They must be able to reform promptly after major disturbances, especially in dwarfs, and initially form promptly as judged by their presence in high redshift galaxies (\citet{fa12}). These more recent results greatly stress some older theories for the origin of the exponentials. This includes the model of \citep{me63}, based on the collapse of a uniform density, uniformly rotating sphere, with no redistribution of angular momentum. The resulting configuration has a distribution of mass as a function of angular momentum that is nearly the same as that of an exponential profile out to a radius of order 6 scale lengths. The assumptions of this simple model are questionable in light of modern disc formation models, which include processes like cold accretion from large scale structures (e.g., \citealt{ro04}, \citealt{mi12}, \citealt{br12}, \citealt{ve14}), and the exponential extent is not great enough. The Mestel model was updated in subsequent decades, especially to incorporate the effects of dark halos, e.g., see \citet{fa80}, \citet{mo98}, and references therein. Like the original these works consider discs that form from generally self-similar halo collapses, preserving specific angular momentum and angular momentum distributions, though some viscous redistribution was also considered. They generally assume an exponential surface density profile in the discs, and do not advance the Mestel model significantly in this regard. However, the authors do note that the model discs are close to or exceed global gravitational stability thresholds, thus providing a basic theoretical understanding for the development of massive clumps or strong waves observed in young discs (\citealt{el07}, \citealt{guo15}), and seen in more recent models (\citealt{bo07}, \citealt{ok16}). Another popular model suggests that viscous accretion is proportional to star formation, and then the exponential profile results from secular evolution (see \citealt{li87}, \citealt{yo89}, \citealt{fe01}, \citealt{wa09}). This type of model is discomfited by the recent observations in a couple of ways. Firstly, it requires substantial shear, which is not found in the dwarf Irregulars, and some regions of spirals. Secondly, it is not clear that disturbed exponentials can be reformed sufficiently rapidly by these processes, especially following a disturbance that does not enhance star formation. By redistributing angular momentum and driving radial migration, bars and spiral waves can also change the surface density profile and produce exponentials. In particular, strong bars can generate double exponentials (\citealt{de06}, \citealt{fo08}). Of course, not all exponential discs have bars, including especially the dwarf Irregulars. While these several processes may play a role, they seem unable to form exponential profiles in all cases and promptly enough to account for the observations. On the other hand, simulations do show the resilience of the exponential profile in model discs. For example, the models of \citet{be15} show how steady accretion is promptly smoothed into the exponential form. \citet{el14} use analytic models to demonstrate the preservation of the exponential form when external accretion is balanced by star formation. In these models the disc can expand or shrink and the exponential scale length evolves (also see Sec. 4.1 below). These many observations and models urge the question of why the exponential form is ubiquitous, and promptly generated? It appears to result from a very basic and fundamental physical process. This, despite the fact that, as we will see below, it is not among the simplest possible equilibrium states. An important hint was provided by the numerical models of \citet{el13}, in which scattering clumps were introduced into simple test particle discs with an initially flat profile. In that work it was found that the scattering off the clumps generally resulted in exponential profiles. (We note that this scattering is rather different than that originating near the corotation radius of a spiral pattern discussed by \citet{se02}, \citet{ros12} and the radial evolution driven by accretion flows modeled in \citet{be15}.) This overall result did not depend on many details, such as the number or mass of the clumps, or the form of the gravitational potential, though other properties, such as the profile evolution time can depend on such parameters. Scattering is common to many, if not all, of the models above, especially those involving bars, waves or clumps. If scattering is indeed the underlying process that generates the exponential profiles, then the question remains of why that profile is the universal result? As does the question of the origin of the different kinds of profile (e.g., broken)? And finally, the question of how much can we learn from analytic models versus numerical simulations? Near equilibrium hydrodynamical processes can often be approximated analytically, but classical scattering problems are usually treated statistically, though with less information obtainable by analytic means. Nonetheless, biased scattering, like that proposed by \citet{el16}, can be viewed as a kind of flow to a steady state. In the following we will demonstrate that, at least in the cases where certain, reasonable assumptions are satisfied, families of near exponentials are steady states over the range of potentials relevant to galaxy discs. Although these steady solutions are not unique, they are the simplest forms that satisfy the physical constraints. In \citet{el16} it was shown with simple scattering models (henceforth `hopping models') that the nature of disk density profiles depends on the type of bias in the scattering. This result suggests that the details of the scattering kinetics determine which of the hydrostatic solutions apply to given galaxy discs. Ultimately, the detailed study of the specific scattering processes that generate near exponential profiles in different disc evolutionary histories requires self-consistent numerical simulations. In the next section we describe general approximations to the hydrodynamic (Jeans) equations used to reduce these to a radial hydrostatic equation, and how scattering can be incorporated into the stellar distribution function as a Gibbs chemical potential. This formalism provides background and context, but may be skipped by readers wanting to proceed directly to discussions of solutions of the two-dimensional continuity equation \eqref{eq1} in the following sections. Section 3 discusses additional physical constraints on the solutions to the hydrostatic equation, and the nature of the constrained solutions. This discussion is expanded with a presentation of variable scalings and profile examples in Sec. 4. The hydrostatic forms derived here are the result of that evolution in mature stellar discs, which are likely to be thick discs (due to vertical scattering) in most cases. In the models of \citet{el13, el16} it was shown that it takes some time for steady, near exponential profiles to develop via scattering. We do not study that evolution in any detail here, but simple self-similar profile evolution is discussed in Sec. 5. Section 4 described some steady broken, exponential profiles, and Sec. 6 reviews the more general causes of breaks, and their evolution in simple scattering models. A summary is provided by Sec. 7. | The goal of this paper has been to better understand the exponential surface brightness or surface density profiles in galaxy discs, whose phenomenology was briefly reviewed in the Introduction. Part of the mystery of exponential discs is that while they appear to be equilibrium states, they are not fit by simple, polytropic solutions of the stellar hydrodynamics equations. This paradox has been sharpened by recent observations showing that the exponential profiles can extend over many scale lengths, as shown by previous studies referenced in Sec. 1. Although the simplest solutions do not suffice, the stellar hydrodynamics or Jeans equations still provide strong constraints on the surface density and velocity dispersion profiles. In Sec. 3 we used physical constraints to narrow the range of possible solutions, and proposed specific forms. These resulted in an exponential or S\'{e}rsic-type radial dependence in the surface densities, but with an additional power-law term. This form also describes the disc hopping model of \citet{el16}. This power-law term has several effects. The first is slightly changing the exponential slope in some cases. These power-law modified profiles can be well fit by two distinct exponential segments in some parameter ranges. These fits usually resemble observed Type II and III disc profiles. The second effect is an upturn, or downturn, of the profile at the smallest radii, which is also seen in observed profiles, but may be difficult to separate in the observations from a bulge contribution. Additionally, in-plane velocity dispersions are predicted to follow moderate power-law functions with radius. The adoption of equation \eqref{eq8} (with power-law $\chi$) reduces the range of steady solutions. Consider the consequences of relaxing it. With these assumptions the effective pressure, $\Sigma \sigma^2$, is an exponential in all cases, since power-law terms cancel. The pressure gradient term in equation \eqref{eq1} (pressure gradient divided by $\Sigma$), is a pure power-law in radius with the same power as the other terms in that equation. However, if the pressure was not purely exponential, then generally there would be two or more pressure terms with different radial scalings. The scaling of the centrifugal term must be altered to balance them. Observations, e.g., comparisons between gas and stellar kinematics, where the former are assumed to represent near circular orbits in centrifugal balance, tend to suggest that the gravitational and centrifugal terms do scale similarly, so such extra terms are usually small. Measurements of the in-plane velocity dispersion scalings in discs are difficult, but would be very helpful for further testing the scalings predicted above. In sum, while stellar discs in galaxies are not simple, cylindrically symmetric, isothermal, exponential (or polytropic) atmospheres, they come rather close. Firstly the observations showing velocity dispersions do not vary by large factors across discs. Secondly, the distribution functions and surface density profiles of Sec. 2 - 4 above are locally, but not globally, isothermal. This despite the fact that galaxy discs are very cold, and nearly in centrifugal balance, so we might not expect even local thermal relaxation. Another difference is that, because stars can be scattered over large distances, and not confined to a local annulus, the distribution functions contain a chemical potential term. The surface density solutions generally contain a power-law dependence on radius as well as the exponential, in each case appropriate to the specific gravitational potential. This cored S\'{e}rsic-type profile, extending over a range of S\'{e}rsic index values from about $3/4$ to $2$ provides a unification with the equilibrium structure of bulges and ellipticals. Pseudo-bulges, in particular, are believed to have indices a bit lower than 2. The overlap with some disc profiles makes sense if the they are indeed secularly formed from discs. Another perspective is obtained by eliminating $r$ between equations \eqref{eq12} and \eqref{eq18}. This yields a density-pressure ($\Sigma - \sigma^2$) relation that differs from a polytropic one in several ways. Firstly, it is more complicated. Secondly, it contains the centrifugal imbalance term $\chi$. It appears that the solutions here are generalizations of the polytropes to cases with the additional effects of biased scattering over a broad range. Hopping models \citep{el16} and numerical scattering models \citep{el13, st16} also show that the stellar disc structure is not described by a constant entropy (polytropic) equilibrium state. A true equilibrium would be a globally isothermal structure, which could only be achieved on a long, two-body relaxation timescale. Even early-type galaxy discs retain a large kinetic energy in near circular rotation, which can be viewed as free energy that will ultimately be converted into thermal energy. This exponential structure is a flow, driven by biased scattering, with an entropy gradient. It is a slow and slowing flow, with a nearly hydrostatic structure, in which scattering minimizes the entropy gradient. Such nonequilibrium, hydrostatic structures may be useful models in a variety of other applications where scattering is important. | 16 | 9 | 1609.08957 |
1609 | 1609.01722_arXiv.txt | {Rapid and luminous flares of non-thermal radiation observed in blazars require an~efficient mechanism of energy dissipation and particle acceleration in relativistic active galactic nuclei (AGN) jets. \mbox{Particle acceleration} in relativistic magnetic reconnection {is being actively} studied by kinetic numerical simulations. Relativistic reconnection produces hard power-law {electron energy distributions $N(\gamma) \propto \gamma^{-p}\exp(-\gamma/\gamma_{\rm max})$ with index} $p \to 1$ and exponential cut-off {Lorentz factor} $\gamma_{\rm max} \sim \sigma$ in the limit of {magnetization} $\sigma = B^2/(4\pi w) \gg 1$ {(where $w$ is the relativistic \mbox{enthalpy density})}. Reconnection in electron-proton plasma can additionally boost $\gamma_{\rm max}$ by the mass ratio $m_{\rm p}/m_{\rm e}$. Hence, in order to accelerate particles to $\gamma_{\rm max} \sim 10^6$ in the case of BL Lacs, reconnection should proceed in plasma of very high magnetization $\sigma_{\rm max} \gtrsim 10^3$. On the other hand, moderate mean jet magnetization values are required for magnetic bulk acceleration of relativistic jets, $\sigma_{\rm mean} \sim \Gamma_{\rm j} \lesssim 20$ {(where $\Gamma_{\rm j}$ is the jet bulk Lorentz factor)}. I propose that the systematic dependence of $\gamma_{\rm max}$ on blazar luminosity class---the blazar sequence---may result from a systematic trend in $\sigma_{\rm max}$ due to homogeneous loading of leptons by pair creation regulated by the energy density of high-energy external radiation fields. At the same time, relativistic AGN jets should be highly inhomogeneous due to filamentary loading of protons, which should determine the value of $\sigma_{\rm mean}$ roughly independently of the blazar class.} \keyword{blazars; relativistic jets; magnetic reconnection} \begin{document} | Blazars are persistent extragalactic sources of non-thermal radiation extending from the radio to the gamma-ray band and characterized by stochastic variability over a wide range of time scales. Multiwavelength observations of blazars typically reveal two major broad spectral components, with the low-energy one (radio---UV/X-ray) interpreted universally as synchrotron radiation of electrons. High-resolution interferometric radio/mm observations reveal a core-jet structure with individual jet substructures propagating with apparently superluminal velocity \cite{Jor01,Lis16}. Blazars are associated with active galactic nuclei (AGN) equipped with jets (radio loud; \cite{Urr95}), with one of the jets pointing closely at the observer \cite{Bla79}, which leads to a dramatic relativistic boost of apparent luminosity \cite{Ree66}. \mbox{Blazars are} typically classified into more luminous flat-spectrum radio quasars (FSRQs) and less luminous BL Lac objects (BL Lacs). An anticorrelation between the radio luminosity and the synchrotron peak frequency is known as the blazar sequence \cite{Fos98}. High apparent luminosities of blazars, up to $L_\gamma \sim 10^{50}\; {\rm erg\cdot s^{-1}}$ in the $\gamma$-ray band \cite{Abd11}, require efficient {in situ} dissipation of the jet power, efficient particle acceleration and efficient radiative~{mechanisms}. Modeling of the spectral energy distributions of blazars can be performed in two basic scenarios \cite{Boe13}: in the leptonic scenario, the high-energy spectral component is interpreted as inverse Compton scattering of soft radiation fields by energetic electrons; in the hadronic scenario, it is interpreted as due to various mechanisms involving relativistic protons. In luminous blazars, the requirement of high radiative efficiency favors the leptonic scenario \cite{Sik09}. However, in any case, the inferred characteristic energies of electrons producing the low-energy synchrotron component are very different in FSRQs ($\gamma_{\rm max} \sim 10^3$) and in BL Lacs ($\gamma_{\rm max} \sim 10^6$). In this work, I attempt to address several fundamental questions about the physics of relativistic AGN jets. \begin{enumerate} \renewcommand\labelenumi{(\theenumi)} \item What is the origin of the blazar sequence? What determines $\gamma_{\rm max}$? \item How is matter introduced into relativistic magnetized jets? What determines $\Gamma_{\rm j}$? \item Can relativistic magnetic reconnection explain energy dissipation and particle acceleration in blazar jets? \end{enumerate} In Section \ref{recon}, I summarize the current understanding of particle acceleration in relativistic magnetic reconnection. In Section \ref{mass}, I sketch two independent mechanisms of mass loading of relativistic jets by leptons and protons. In Section \ref{seq}, I discuss why {the blazar sequence} is unlikely to be regulated by radiative cooling rates. In Section \ref{prop}, I propose a novel picture of the composition of relativistic jets and {an alternative explanation} of the blazar sequence. | \label{con} Relativistic reconnection is a promising dissipation mechanism in relativistic jets. Rapid progress in understanding particle acceleration during relativistic reconnection has been achieved recently by means of kinetic numerical simulations. We {have learned} that relativistic reconnection can accelerate particles very efficiently, producing hard power-law distributions with index $p = 1$ and exponential cut-off at $\gamma_{\rm max} \sim \sigma_{\rm e}$. Applying this mechanism to blazars requires very high jet magnetization values, $\sigma_{\rm max} \sim 10^3$ in the case of TeV BL Lacs. I suggest that such high magnetizations may be present locally in relativistic jets due to highly inhomogeneous mass loading of protons by plasma instabilities. Fast~reconnection rates $\beta_{\rm rec} \sim 0.1$ can easily compete with radiative cooling rates, allowing in principle {acceleration of} particles to much higher energies than required by SED modeling, especially in the case of FSRQs. Therefore, I argue that the blazar sequence can hardly be regulated by radiative \mbox{cooling rates}. Instead, I suggest that the blazar sequence arises due to homogeneous loading of leptons by pair creation regulated by external radiation fields. Several aspects of the proposed picture of blazar physics require detailed theoretical or numerical verification. \vspace{6pt} | 16 | 9 | 1609.01722 |
1609 | 1609.01514_arXiv.txt | {Arch filament systems occur in active sunspot groups, where a fibril structure connects areas of opposite magnetic polarity, in contrast to active region filaments that follow the polarity inversion line. We used the GREGOR Infrared Spectrograph (GRIS) to obtain the full Stokes vector in the spectral lines Si\,\textsc{i} $\lambda$1082.7\,nm, He\,\textsc{i} $\lambda$1083.0\,nm, and Ca\,\textsc{i} $\lambda$1083.9\,nm. We focus on the near-infrared calcium line to investigate the photospheric magnetic field and velocities, and use the line core intensities and velocities of the helium line to study the chromospheric plasma. The individual fibrils of the arch filament system connect the sunspot with patches of magnetic polarity opposite to that of the spot. These patches do not necessarily coincide with pores, where the magnetic field is strongest. Instead, areas are preferred not far from the polarity inversion line. These areas exhibit photospheric downflows of moderate velocity, but significantly higher downflows of up to 30\,km\,s$^{-1}$ in the chromospheric helium line. Our findings can be explained with new emerging flux where the matter flows downward along the fieldlines of rising flux tubes, in agreement with earlier results. } | The structure of solar filaments depends on the magnetic field at their photospheric footpoints. Thus, to understand the processes forming a filament, a good knowledge of the underlying magnetic vector field is mandatory. In active regions, one has to distinguish between active region filaments and arch filament systems (AFS). Active region filaments follow the neutral or polarity inversion line (PIL), while structures in AFS cross the PIL. Commonly, active region filaments are related to magnetic shearing at the PIL, and AFSs appear where the emergence of new flux is still ongoing. The physical processes that lead in both cases to dark structures in chromospheric lines need to be investigated in more detail. It is generally assumed that the dark structures of an AFS represent the emerging fluxtubes. In this work we investigate such an AFS. For active region filaments we refer to recent works of \cite{cit:kuck12mag}, \cite{cit:kuck12vel}, \cite{cit:sasso2014} and \cite{cit:schwartz2016} and references therein. \begin{figure}[t] \includegraphics[width=\columnwidth]{th12_balthasar_fig01} \caption{Intensity image taken at 08:16\,UT by HMI. The white box outlines the area investigated in this work. } \label{fig_hmi} \end{figure} \begin{figure*}[t] \includegraphics[width=\textwidth]{th12_balthasar_fig02} \caption{Temporal evolution of the slit-reconstructed continuum images. We apply an unsharp masking to enhance the displayed contrast. The black arrow points towards disk center and the white arrow to an outjutting part of the penumbra (OJP) of the nearby spot.} \label{fig_int} \end{figure*} \cite{cit:howardharvey} distinguished between filaments and fibrils, where filaments lie parallel to iso-contour lines of the magnetic field and fibrils traverse them perpendicularly. \cite{cit:martres66} observed fibrils (``traces filamenteuses'') that connect areas of opposite polarity crossing the PIL. Systems of fibrils were described by \cite{cit:bruzek67}, and he suggested to call them AFS. They appear in the inter spot area and connect spots of opposite polarities and cross the PIL. Strong downflows of about 50\,km\,s$^{-1}$ occur at the footpoints of arch filaments \citep{cit:bruzek69}. The dynamic evolution of an AFS was investigated by \cite{cit:spadaro}, who found upward motions in central parts of the fibrils and downward flows at their ends. Recently, \cite{cit:vargas} used Hinode data to study AFS related to granular scale flux emergence in an active region. \cite{cit:grigoreva} considered the formation of AFS in an active region as the onset of newly emerging flux before a strong flare occurred one day later. \cite{cit:ma} investigate an AFS in H$\alpha$, and they use magnetograms from the Helioseismic and Magnetic Imager (HMI) onboard the Solar Dynamics Observatory (SDO). In their case, the arch filaments are related to moving magnetic features of opposite polarity to that of the main sunspot and the arch filaments are controlled by the moving magnetic features. \cite{cit:saminat} and \cite{cit:lagg07} observed an AFS in the infrared helium line at 1083\,nm and determined its magnetic structure. The chromospheric magnetic field of another AFS is investigated by \cite{cit:xu10}. In the zone of flux emergence they find a magnetic field strength of 300\,G, and the magnetic field is almost horizontal. At the edge of the emerging zone, the field is more vertical and stronger with a field strength of up to 850\,G. Supersonic downflows occur at both footpoints of a fibril (or loop as the authors call the structures). In the present work, we investigate an AFS, and we will concentrate on the magnetic field and the Doppler velocities in the lower photosphere. According to \cite{cit:howardharvey} we will use the term ``fibril'' instead of ``arch filament'' for the single structures of the AFS in this work. Results of the upper photosphere and chromosphere will be topic of another article, except for a velocity determination from the infrared helium triplet line at $\lambda$1083\,nm. \begin{figure*}[t] \includegraphics[width=\textwidth]{th12_balthasar_fig03} \caption{Temporal evolution of the slit-reconstructed helium line core intensity. White contours mark the dark photospheric intensity structures as seen in Fig.\,\ref{fig_int}, and the black ones indicate the PILs. The small arrows in the mid panel indicate the splitting of the main fibril.} \label{fig_inthe} \end{figure*} | We observe blueshifts (upflows) in the helium velocities along the central part of the fibrils, which hints at rising flux tubes. At their summits, matter can move only along the field line, and thus the matter has to rise with the flux tube. At the footpoints of the flux tubes we see a downward motion. These results are in qualitative agreement with previous results obtained by \cite{cit:bruzek69}, \cite{cit:saminat}, \cite{cit:spadaro}, \cite{cit:lagg07} and \cite{cit:xu10}. Since we observe redshifts at both ends of the fibrils, we can exclude a siphon flow, in that case one would expect a blueshift on one side. In total, this picture is in good agreement with the sketch of \cite{cit:bruzek69}. The rising flux tube does not gain much matter by condensations out of the surroundings, and after a while, not enough matter remains to maintain the downward drain and to absorb light in the helium line. The fibrils are no longer visible, and the downflow ends. This process is repeated when a new flux tube emerges. The analysis of the helium and silicon data is ongoing work, and we have to refer to a future work answering the question: How is the magnetic field oriented in the layers of the fibrils? The analysis of the helium and silicon data will shed more light on the connection between the deep photosphere to the chromosphere. In addition, we plan new observations to elucidate the physical properties of AFSs. | 16 | 9 | 1609.01514 |
1609 | 1609.04041_arXiv.txt | We present the results of the Gould's Belt Distances Survey (GOBELINS) of young star forming regions towards the Orion Molecular Cloud Complex. We detected 36 YSOs with the Very Large Baseline Array (VLBA), 27 of which have been observed in at least 3 epochs over the course of 2 years. At least half of these YSOs belong to multiple systems. We obtained parallax and proper motions towards these stars to study the structure and kinematics of the Complex. We measured a distance of 388$\pm$5 pc towards the Orion Nebula Cluster, 428$\pm$10 pc towards the southern portion L1641, 388$\pm$10 pc towards NGC 2068, and roughly $\sim$420 pc towards NGC 2024. Finally, we observed a strong degree of plasma radio scattering towards $\lambda$ Ori. | Young star forming regions towards Orion have been the subject of much interest, as they are the closest regions of massive young stellar population. The star formation in the Orion Complex is concentrated in two molecular clouds, Orion A and B, with clusters such as the Orion Nebula Cluster (ONC) and L1641 in Orion A, and NGC 2023/2024, NGC 2068/2072 and L1622 in Orion B. These clusters represent the most recent episodes of star formation in the region, which belong to the Orion OB1c and 1d sub-association, containing stars spanning ages from $\sim$1 Myr up to 6 Myr \citep{2008bally}. In addition to the clusters in the main cloud, there are other stellar groups in Orion that host very young stars, like $\sigma$ Ori, in the OB1b sub-association, and the groups of the $\lambda$ Ori association at the northernmost end of the complex. Finally, a somewhat older (8--12 Myr) population is contained within the OB1a sub-association, where most of the parental gas has already been removed. Over the course of the last century, many attempts have been made to measure distances to the Complex, particularly towards the ONC. Some of the earliest measurements were as high as 2000 pc \citep{1917Pickering} and as low as 185 pc \citep{1918Kapteyn}. Eventually most measurements settled in the 350 to 500 pc range obtained through various means, most typically through zero-age main sequence fitting. Much of the scatter originated from inconsistent assumptions, models, and sample selection \citep[see review by][]{2008Muench}. For some time, the most widely used distance was 480$\pm$80 pc, obtained from proper motions of H$_2$O masers towards the Orion BN/KL region \citep{1981Genzel}. In the last decade, however, direct stellar parallax measurements of non-thermal emitting masers and stars were made possible through radio Very Long Baseline Interferometry (VLBI). \citet[hereafter MR]{2007menten} obtained a distance of 414$\pm$7 pc from observations of 4 stars - GMR A, F, G, and 12 - in the central (Trapezium) region of the ONC. \citet[hereafter S07]{2007sandstrom} also observed GMR A and obtained a somewhat closer distance of 389$^{+24}_{-21}$ pc. \citet{2007hirota} and \citet{2008kim} observed H$_2$O and SiO masers to obtain a distance of 437$\pm$19 pc and 418$\pm$6 pc respectively in the Orion BN/KL region. Other major efforts to measure a distance towards the ONC include \citeauthor{2007Jeffries}'s (\citeyear{2007Jeffries}). He used stellar rotation to estimate distances of 440$\pm$34 pc for his entire sample and 392$\pm$32 pc including only stars without active accretion. \citet{2004Stassun} obtained a distance of 419$\pm$21 pc through monitoring the kinematics of a double-line eclipsing binary system, assuming a value for the solar bolometric luminosity of $M_{bol,\odot}=$4.59, although their distance estimate decreased to 390$\pm$21 pc with $M_{bol,\odot}=$4.75. \citet{2009Kraus} obtained a dynamical distance of 410$\pm$20 pc based on modeling the orbit of the close binary $\theta^1$ Ori C. Some attempts have also been made to obtain distances from dust extinction maps towards not just the ONC, but towards several distinct regions in the Orion Complex. \citet{2011Lombardi} estimated 371$\pm$10 pc towards Orion A and 398$\pm$12 pc towards Orion B using extinction maps measured from 2MASS. \citet{2014Schlafly} provided distance estimates of 20 distinct regions through extinction from PanSTARRS photometry, although many of them are highly uncertain. While the distance measured by MR is currently considered as canonical, it is is based on a small sample of 4 stars. In addition, the MR stars all lie within the central regions of the ONC; however, the Complex spans 100 pc projected on the sky, so it would not be surprising if the different regions of the cloud have substantially different distances, and it would not be surprising if the regions have differing radial distances of the same order. Therefore, even if the distance towards the ONC is known with high accuracy, by applying this distance to other regions an inherent uncertainty of $\sim$20\%, for example, could be introduced, as the Complex is located at the distance of $\sim$400 pc. This propagates to an error of $\sim$40\% in luminosity, to $\sim$70\% in ages of young stars \citep{2001hartmann}. Currently an ongoing mission of the \textit{Gaia} space telescope is in process of obtaining astrometry towards optically visible sources across the entire sky in order to measure parallaxes accurate to 100 $\mu$as for G$<$17 mag stars \citep{2014deBruijne}, which should provide accuracy in distance measurements to within 5-10\% up to 1 kpc. VLBI observations can provide an important independent check on optical parallax measurements, as shown by the comparison of VLBI with \textit{Hipparcos} distances for the Pleiades \citep{2014Melis}. In addition, radio VLBI can be useful for measuring sources in regions of high extinction and/or significant nebulosity, as is the case in many regions of Orion. In this paper we present radio VLBI observations of stellar parallaxes of Young Stellar Objects (YSOs) identified towards the Orion Complex, hereby significantly expanding the number of stars in Orion with known distances and kinematics. This work is done as part of Gould's Belt Distances Survey \citep[GOBELINS,][]{gbds}, which is dedicated to measure stellar parallax towards the Ophiuchus \citep{oph2}, Serpens \citep{ser}, Taurus, Perseus, and Orion star forming regions. | We monitored 36 non-thermal radio emitting YSOs spread throughout the Orion Complex with VLBA over a period of 2 years, and we report measured stellar parallaxes towards 26 of them. Fifteen of them are located towards the ONC, and we find a distance of 388$\pm$5 pc to the cluster; somewhat closer than the canonical $414\pm7$ pc distance found by \citet{2007menten} that is typically used in the literature. This result has implications on the luminosity and ages of the cluster. If the cluster is 7\% closer than the previous estimation, this implies that it is 12\% fainter, and 20\% older (assuming a relation $t\propto L^{-3/2}$) than what was previously reported in surveys of the ONC such as the one by \citet{2010dario}. We also report distances towards other regions located in the Orion Complex, such as L1641, NGC 2024, and NGC 2068. While these values are somewhat more uncertain due to significantly smaller sample size, limited spatial coverage (particularly in case of L1641), and multiplicity, these are the first direct measurements of the stellar parallaxes towards these regions. This provides insight into the structure of the Complex. We identify a possible region of large degree of plasma scattering towards the $\lambda$ Ori star forming region. The degree of scattering is significant, with broadening of the observed size of the objects of up to 16.5 mas near the center of the cluster at 5 GHz. The scattering is spread all within the ring 2.5---3$^\circ$ in radius produced by supernova activity. Unfortunately, this effect made it impossible to measure astrometry accurately enough to obtain a parallax towards stars found in this region. A persistent problem in the analysis of both the parallax and proper motions of the stars is the multiplicity. We conclusively identify 5 of 27 stars that have been detected in at least three epochs belonging to a multiple system with orbital periods between 6 months to 10 years, with at least three more systems identified as likely binaries, although further monitoring would be necessary to confirm them. It is impossible to accurately determine parallaxes to these systems without solving for orbital motion of these systems, which we can presently do to only two of them. Six stars are known spectroscopic binaries with very short periods (one of them also has an aforementioned intermediate period companion); possibly a larger number of them could have very close companions that are yet to be identified, particularly since few surveys of spectroscopic binaries have been performed in the Orion Complex outside of the ONC. While understanding of their orbits is not detrimental for finding the parallax, it could still influence the solution somewhat. Finally, four stars in the sample have long period companions, although only one of them does not have a closer companion in a higher order multiple system. These wide companions should not affect the solution for the parallax, although they do affect proper motions. In all, at least 14 (possibly more) of 27 stars observed with VLBA belong to multiple systems. Whether this multiplicity fraction is consistent with that for the entire Complex is not yet known as it is presently difficult to identify companions with intermediate periods towards Orion due to its distance. Future generations of high resolution optical and IR telescopes would make it possible to identify the full extent of multiplicity towards this region. Further monitoring of the identified YSOs would be beneficial as due to variability in radio, currently only a limited number of detections are available to some stars. In the future it would be possible to more effectively measure their parallax and proper motions. It will also be necessary to confirm multiplicity and constrain orbital parameters towards some sources. The distance solutions produced by the GOBELINS survey will be used as independent constraint on the accuracy of \textit{Gaia}, as the systematic effects behind the sample selection and the individual observations are different between these two programs. Approximately half of the systems observed with VLBA towards the Orion Complex are optically visible, therefore it should be possible to compare the distance solutions towards them directly, at least in the ONC, although nebulosity could significantly degrade performance in the optical regime. On the other hand, star forming regions towards Orion B suffer from high extinction; therefore only a few members of NGC 2024 and NGC 2068 would be detectable with \textit{Gaia}. | 16 | 9 | 1609.04041 |
1609 | 1609.01452_arXiv.txt | { Average stellar radii in open clusters can be estimated from rotation periods and projected rotational velocities under the assumption that the spin axis has a random orientation. These estimates are independent of distance, interstellar absorption, and models, but their validity can be limited by lacking data (truncation) or data that only represent upper or lower limits (censoring). }{ We present a new statistical analysis method to estimate average stellar radii in the presence of censoring and truncation. }{ We used theoretical distribution functions of the projected stellar radius $R \sin i$ to define a likelihood function in the presence of censoring and truncation. Average stellar radii in magnitude bins were then obtained by a maximum likelihood parametric estimation procedure. }{ This method is capable of recovering the average stellar radius within a few percent with as few as $\text{about ten}$ measurements. Here we apply this for the first time to the dataset available for the Pleiades. We find an agreement better than $\approx$ 10 percent between the observed $R$ vs $M_K$ relationship and current standard stellar models for $1.2\ge M/M_{\odot}\ge 0.85$ with no evident bias. Evidence of a systematic deviation at $2\sigma$ level are found for stars with $0.8\ge M/M_{\odot}\ge 0.6$ that approach the slow-rotator sequence. Fast rotators ($P < 2$ d) agree with standard models within 15 percent with no systematic deviations in the whole $1.2 \apprge M/M_{\odot} \apprge 0.5$ range. } { The evidence of a possible radius inflation just below the lower mass limit of the slow-rotator sequence indicates a possible connection with the transition from the fast- to the slow-rotator sequence.\thanks{Table 1 is only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/} } | \label{sec:Introduction} The disagreement between theoretical and observed parameters of young magnetically active and of fully convective or almost fully convective low-mass stars remains one of the main long standing problems in stellar physics. Current investigations focus on the inhibition of the convective transport by interior dynamo-generated magnetic fields and/or by the blocking of flux at the surface by cool magnetic starspots \citep[e.g.][]{2001ApJ...559..353M, 2007A&A...472L..17C, 2013ApJ...779..183F, 2014ApJ...789...53F, 2014MNRAS.441.2111J}, which produce an increase in stellar radius and a decrease in $T_{\rm eff}$. The same effect is also thought to be linked to the observed correlation between Li abundance and rotation \cite[e.g.][]{ 2014ApJ...790...72S, 2015MNRAS.449.4131S, 2015ApJ...807..174S, 2014MNRAS.441.2111J}. The consequences of these discrepancies are manifold. These include, for example, the determination of the mass and the radius of exoplanets, whose accuracy depends on that of the hosting star \citep[e.g.][]{2004ASPC..318..159H,2015ApJ...804...64M}, the age estimate of young open clusters \citep[e.g.][]{2014prpl.conf..219S,2015ApJ...807..174S}, and the mass-luminosity relationship for magnetically active low-mass stars. Fundamental determinations of stellar masses and radii with a 3 percent accuracy or better are provided by the light-curve analysis of detached eclipsing binaries \citep[e.g.][]{2010A&ARv..18...67T,2012ApJ...757...42F}. Interferometric angular diameter measurements of single stars are available today for tens of stars \citep[e.g.][]{2012ApJ...757..112B} with diameters measured to better than 5 percent. Statistical methods based on the product of $P$ and \vsini, which produces the projected radius $R \sin i$ (Sect.\,\ref{sec:ProjectedRadius}), and the assumption of random orientation of the spin axis \citep[e.g.][]{2009MNRAS.399L..89J} have the advantage of providing mean radii estimates for a large number of (coeval) single stars independently of distance, interstellar absorption, and models. No evidence of preferred orientation of the spin axis in open clusters has been found so far \citep[e.g.][and references therein]{2010MNRAS.402.1380J}, and therefore the method seems to be sound in this respect. The main difficulty is that the data sample is always truncated at a combination of sufficiently low inclination angle $i$ and low equatorial velocities $v_{\rm eq}$. In these cases, depending also on the spectral resolution, \vsini\ cannot be derived and only an upper limit can be given. A low $i$ may also cause difficulties in measuring $P$ and therefore there may be cases in which either one or both \vsini\ and $P$ cannot be measured. At the other extreme, ultra-fast rotator spectra can be so smeared by the rotational broadening that in some cases only a lower \vsini\ limit can be given. To take a low $R \sin i$ truncation into account, \cite{2009MNRAS.399L..89J} considered a cut-off inclination such that stars with lower inclination yield no $R \sin i,$ and they corrected the average $\sin i$ accordingly. Mean radii are then derived by taking the average of the ratio $R \sin i / \left< \sin i \right> $ in suitable magnitude bins. Here we present a new method, based on the survival analysis concept \citep{Klein+2003}, that makes use of the whole information content of the dataset by also considering upper and lower limits and data truncation. Data may also come from inhomogeneous estimates, like those in which \vsini\ upper and lower limits are obtained from different analyses and instrumentation, as long as they are not affected by significant biases. Uncertainties due to surface differential rotation (SDR) are also estimated, with the most likely values derived from the recent work of \cite{Distefano_etal:2016}. The method is applied for the first time to the rich dataset available for the Pleiades. In Sect.\,\ref{sec:data} we present the data used in this work. The method is described in Sect.\,\ref{sec:Method}. The results obtained for the Pleiades dataset are discussed in Sect.\,\ref{sec:Pleiades}. We draw our conclusions in Sect.\,\ref{sec:Conclusions}. | \label{sec:Conclusions} We have set up a new method for deriving mean stellar radii from rotational periods, $P$, and projected rotational velocities, \vsini, based on the survival analysis concept \citep{Klein+2003}. This method exploits the whole information content of the dataset with an appropriate statistical treatment of censored and truncated data. Provided censoring and truncation do not significantly affect the peak of the $R \sin i$ distribution and that there is no significant bias in the data, the method can recover the mean stellar radius with an accuracy of a few percent with as few as $n \approx 10$ measurements. The total standard deviation, $\sigma$, which cumulatively takes the data noise and the intrinsic $R$ standard deviation into account, can also be estimated with an accuracy of a few percent except in extreme cases where the distribution is too broad ($\sigma / \mathcal{R} \sim 0.3$) or too narrow ($\sigma / \mathcal{R} \sim 0.01$). The method has been applied for the first time to the dataset available for the Pleiades. We found that deviations of the empirical $\mathcal{R}$ vs $M_K$ relationship from standard models \citep[e.g.][]{2013ApJ...776...87S,2015A&A...577A..42B} do not exceed 5 percent for $1.2 \ge M/M_{\odot} \ge 1.0$ and 10 percent for $1.0 > M/M_{\odot} \ge 0.85$, with no significant bias. Evidence of a systematic deviation at $1-2\sigma$ level of the empirical $\mathcal{R}$ vs $M_K$ relationship from standard models is found only for stars with $M \approx 0.7 \pm 0.1 M_{\odot}$ that are converging on the slow-rotator sequence. Deviations of the $\mathcal{R}$ vs $M_K$ relationship for fast rotators ($P<2$~d) do not exceed $\approx 15$ percent in the whole mass range with no evidence of a systematic deviation from standard models. No evidence of a radius inflation of fast rotators in the Pleiades is therefore found. | 16 | 9 | 1609.01452 |
1609 | 1609.07873_arXiv.txt | {High resolution stellar spectral atlases are valuable resources to astronomy. They are rare in the $1 - 5\,\mu$m region for historical reasons, but once available, high resolution atlases in this part of the spectrum will aid the study of a wide range of astrophysical phenomena.} {The aim of the CRIRES-POP project is to produce a high resolution near-infrared spectral library of stars across the H-R diagram. The aim of this paper is to present the fully reduced spectrum of the K giant 10 Leo that will form the basis of the first atlas within the CRIRES-POP library, to provide a full description of the data reduction processes involved, and to provide an update on the CRIRES-POP project.} {All CRIRES-POP targets were observed with almost 200 different observational settings of CRIRES on the ESO Very Large Telescope, resulting in a basically complete coverage of its spectral range as accessible from the ground. We reduced the spectra of 10 Leo with the CRIRES pipeline, corrected the wavelength solution and removed telluric absorption with \textit{Molecfit}, then resampled the spectra to a common wavelength scale, shifted them to rest wavelengths, flux normalised, and median combined them into one final data product.} {We present the fully reduced, high resolution, near-infrared spectrum of 10 Leo. This is also the first complete spectrum from the CRIRES instrument. The spectrum is available online.} {The first CRIRES-POP spectrum has exceeded our quality expectations and will form the centre of a state-of-the-art stellar atlas. This first CRIRES-POP atlas will soon be available, and further atlases will follow. All CRIRES-POP data products will be freely and publicly available online.} | Stellar atlases and reference spectra provide a useful observational counterpoint to laboratory- and computationally-produced atomic and molecular line lists, model stellar atmospheres, and synthetic spectra. These observational and theoretical resources are interdependent, and the research applications of both are numerous. A few uses include identification of spectral features, determination of relative abundances, estimation of stellar properties, assisting the planning of observing programs, and providing references for the search for faint companions. At very high resolution, an observational stellar atlas can provide useful feedback to line lists and model atmospheres. Although there are several high resolution stellar atlases available in the optical range \citep[for example the UVES-POP (Ultraviolet and Visual Echelle Spectrograph Paranal Observatory Project) library;][]{uves-pop}, there is a distinct lack of atlases in the near-infrared (NIR), especially at high resolutions (throughout this paper we use the term `near-infrared' to refer to the $1 - 5\,\mu$m region, following IPAC convention). Extant examples include the Arcturus atlas \citep{arcturus} and that of the Sun \citep{sun}, both of which were observed using Fourier Transform Spectrographs, the work of Wallace \& Hinkle and collaborators \citep{wallacehinkle96,wallacehinkle97,joyce98,wallacehinkle02}, the work of the Li\`ege group \citep[][and references therein]{zander} and, at lower resolution, the NASA Infrared Telescope Facility (IRTF) spectral library of cool stars \citep{irtflib}. The NIR range was understudied for several decades, due partly to limitations of instrumentation, and partly to telluric absorption. For several parts of the NIR, Earth's atmosphere is opaque, which defines the boundaries of atmospheric windows outside of which observing from the ground is impossible (the well-known \textit{Y, J, H, K, L}, and \textit{M} bandpasses), and telluric absorption lines are common even within these windows. In recent years, NIR instrumentation has improved significantly, and techniques for telluric line removal are much improved \citep{seifahrt,molecfit,molecfit2}. This opens up the NIR as a rich laboratory for studying cool phenomena such as stellar atmospheres, disks, exo\-planets, and circumstellar matter. Particularly for cool stars, the NIR allows us to see many spectral features whose counterparts are difficult to distinguish at optical wavelengths, including those of atomic transitions of heavy elements, and those of a variety of molecules. The CRIRES-POP project \citep[][hereafter Paper I]{crires-pop} is in the process of producing high resolution, high signal-to-noise ratio (S/N) stellar spectral atlases that cover the entire NIR, for 26 stars that span a range of spectral types and luminosity classes. The atlases will be based on observations from the ESO/VLT high resolution NIR spectrograph, CRIRES (Cryogenic High-Resolution Infrared Echelle Spectrograph). The 26 targets were selected so the library would cover as much of the Hertzsprung-Russell diagram (HRD) as possible while satisfying the observing quality criteria. Targets were also selected to be representative of their spectral type and luminosity class, with the aim of making the atlases widely useful. An important motivation for CRIRES-POP was the provision of template spectra for the Extremely Large Telescopes, which will be most efficient in the NIR because of adaptive optics (AO) performance. The raw data, pipeline-reduced data, and final atlases will all be publicly available on the CRIRES-POP web archive and at CDS. Raw and pipeline-reduced data are already available at www.univie.ac.at/crirespop/. Each final atlas will consist of the full high resolution NIR stellar spectrum, line identifications, isotope ratios, abundances of major species, and stellar properties. We chose the K giant 10 Leo (HD 83240) for the first stellar atlas, partly because of our interest in cool stars and partly because it is the same spectral type as Arcturus (although with a very different metallicity), allowing us to use the Arcturus atlas as a reference when refining our data reduction process. 10 Leo is a red clump giant in the thin-disk population \citep{soubiran}, and is a long-period binary with an undetected low-mass companion. Some stellar and orbital properties of 10 Leo collected from the literature are given in Table~\ref{10leotable}. We note that the values are not all consistent, possibly due to differing analysis techniques and assumptions \citep[see for example ][]{lebzelter12}. Using the $T_{\rm{eff}}$ and distance values in the table, and assuming $A_V=0.03$ \citep{fink} and the bolometric correction from \cite{castellikurucz}, we calculate $L = 67.61\,L_{\odot}$ and $R = 11.9\,R_{\odot}$, which are clearly at odds with the literature values quoted in Table~\ref{10leotable}. \begin{table} \caption{Properties of 10 Leo} \label{10leotable} \centering \begin{tabular}{lcc} \hline\hline Property & Value & Reference\\ \hline SpT & K1 III & 1 \\ RA & 09:37:13 & 2\\ Dec & +06:50:09 & 2\\ $T_{\rm{eff}}$ & $4801\pm 89$\,K & 3\\ $L$ & 59.35\,$L_{\odot}$ & 4\\ $R$ & 14\,$R_{\odot}$ & 5\\ log $g$ & $2.83\pm 0.23$ & 3\\ Age & $3.51\,\pm\,1.80$\,Gyr & 6\\ Distance & 75.3\,pc & 2\\ $\rm{[Fe/H]}$ & $-0.03\pm 0.08$ & 3\\ $P_{\rm{orb}}$ & $2834\,\pm\,4$\,d & 7\\ $\gamma$ & $+20.0\,\pm\,0.1\,\rm{km\,s^{-1}}$ & 7\\ $T_{0}$ & $3\,8888\,\pm\,31$ (MJD) & 7\\ \hline \end{tabular} \tablebib{(1)~\cite{roman52}; (2)~\cite{hipparcos}; (3)~\cite{dasilva}; (4)~\cite{mcdonald}; (5)~\cite{pasfrac}; (6)~\cite{soubiran}; (7)~\cite{griffin}} \end{table} The production of a CRIRES-POP stellar spectral atlas is a two-stage process. The first stage consists of reducing the observed spectra, correcting for telluric absorption, and combining them into a single fully reduced stellar spectrum spanning $1 - 5\,\mu$m. The second stage encompasses line identification, abundance determination, and the compilation of stellar properties into the final spectral atlas. For 10 Leo, as the first published star from the CRIRES-POP library, a separate paper will be devoted to each of these two production stages. This allows us ample space to fully describe the associated data reduction and analysis, and as the third full NIR stellar spectrum and the first full CRIRES spectrum ever to be published, we feel this in-depth description is warranted. Our idea is that these two papers (comprising this paper and a forthcoming one) will serve as references for future CRIRES-POP project publications. The purpose of this paper is to introduce the full contents of the CRIRES-POP library now that all observations are completed, to detail the data reduction and preparation process for CRIRES-POP spectra, and to present the first reduced spectrum, that of 10 Leo. Section~\ref{observations} briefly describes the observations and summarises the properties of all stars in the library. Section~\ref{methods} describes the data reduction, telluric correction, and subsequent extensive work required for the preparation of a CRIRES-POP final spectrum. Examples of the final 10 Leo spectrum are presented in Sect.~\ref{results}. Finally, Sect.~\ref{discussion} summarises the paper, describes the next steps for the project, and explains the data that will be available and how they can be accessed. | \label{discussion} We presented the first fully reduced CRIRES-POP spectrum, that of the K1\,III giant 10 Leo, which will form the basis of the first CRIRES-POP stellar spectral atlas. This is also the first spectrum to cover (essentially) the complete spectral range of the CRIRES spectrograph. The observed CRIRES spectra were reduced with the standard pipeline, telluric- and wavelength-corrected with \textit{Molecfit}, corrected for observed radial velocity and heliocentric motion, resampled to a common wavelength scale at a constant resolving power of 90\,000, continuum normalised, and median combined to produce one final spectrum. Some poor data had to be discarded or fixed, and areas of residual telluric contamination are masked in the final spectrum. Examples of the spectrum are shown in Figs.~\ref{finalspec1} to \ref{finalspec7}. The spectrum is indicative of the quality of the CRIRES-POP library spectra as a whole, and we expect it will make an exceptional atlas. Naturally, the spectrum has some limitations. This reference spectrum is produced from actual observations of a real star, which means that it is not perfectly smooth with absolute wavelength coverage like a synthetic spectrum would be. The S/N is not constant across the spectrum because observing conditions are variable and rarely ideal, and also because of the effect of the grating blaze function. Some unidentified telluric lines may remain in the spectrum, although we will be able to identify these as such once further spectra in the library are reduced. The continuum normalisation is not perfect everywhere, as the continua of the observed spectra are warped by broad telluric bands, influences of the spectrograph such as curvature of the spectrum across detector rows and partial polarisation of light as it passes through the slit, and the pipeline reduction process; and \textit{Molecfit} has a limited ability to fit a highly warped continuum. There are gaps due to masks, poor coverage of CRIRES settings, and wavelengths that simply cannot be observed from the ground. Some features that were visible in the Arcturus atlas are occluded by telluric lines in our spectrum, which is due to differences in radial velocity between the stars; but conversely, our spectrum includes features that were occluded for Arcturus. Some line depths are uncertain because of differing telluric subtractions on different chips. However, an observed reference spectrum or stellar atlas has some advantages over synthetic spectra. It demonstrates to those planning observations which parts of the spectrum are scientifically interesting or useful, what S/N can be achieved, which parts of the spectrum are unusable as a result of telluric absorption, and which telluric lines can be feasibly removed without undue loss of information from the underlying stellar features. It provides feedback to telluric line lists to help improve tools like \textit{Molecfit}, and feedback to stellar line lists and atmosphere models to help improve synthetic spectra. It can be used to estimate relative abundances for other stars, study nucleosynthetic processes through isotopic ratios, and investigate stellar atmospheres and circumstellar environs. There is a clear need for both high-quality synthetic spectra and high-quality observational stellar atlases in astronomy, and each rely on the other. While 10 Leo is a K giant like Arcturus, significant differences in their stellar parameters can be easily spotted in a direct comparison of our CRIRES-POP spectrum with the Arcturus atlas by \cite{arcturus}. We give two examples of this comparison in Figs.~\ref{finalspec3} and~\ref{finalspec6}. The effective temperature of 10 Leo of 4801\,K (Table~\ref{10leotable}) is about 500\,K higher than that of Arcturus \citep{jonsson}. The latter also has a lower log $g$ value (1.67 vs.~2.83) and is known to be a metal-poor star, with $\rm{[Fe/H]}$ about 0.6 dex below 10 Leo. To have a reference for the changes produced in the spectrum by differing temperature and log $g$, we computed synthetic spectra for a COMARCS model atmosphere at 4300 and 4800\,K, and at two different log $g$ values. Details on the models and the computation of synthetic spectra can be found in \cite{aringer}. At the higher temperature, all atomic and molecular lines in the part of the spectrum shown in Fig.~\ref{finalspec3} are weakened. In accordance with the temperature difference between the two stars, all molecular features show a lesser depth in 10 Leo than in Arcturus. This is very obvious both in the CO 3-0 lines longwards of $1.558\,\mu$m and in the OH 2-0 and 3-1 lines. For CO, the model spectra also reveal a significant sensitivity of the line depth with surface gravity, leading to a further weakening when moving from log $g$ of 1.5 to 2.5. However, the opposite trend is expected for the OH lines, which are thought to become stronger at higher log $g$, although this is not apparent in Fig.~\ref{finalspec3}. A detailed analysis of the resulting abundances of the various molecular species in 10 Leo will be presented in a forthcoming paper. The atomic lines, in particular iron, show much less difference in strength between 10 Leo and Arcturus, despite what would be expected from the temperature difference. This supports the results from previous investigations that 10 Leo is more metal rich than Arcturus, with the higher metal abundance compensating for the reduction in line strengths by temperature. The section of the spectrum in the $L$ band plotted in Fig.~\ref{finalspec6} is dominated by a series of fundamental OH lines. There is a very clear difference in strength between 10 Leo and Arcturus, and this agrees well with expectations from the model spectra. The Si lines in Fig.~\ref{finalspec6} show only minor differences between the two stars even though a slight temperature sensitivity is expected. As in the $H$ band, this probably reflects the higher metallicity of 10 Leo. We stress that these comparisons are based on a qualitative analysis using model spectra of similar temperature and surface gravity. This is done primarily for illustrative purposes. A detailed analysis of stellar parameters and abundances for 10 Leo based on the whole CRIRES-POP spectrum will be presented in a forthcoming paper. Figure~\ref{telluric+magnetic} shows two prominent Na lines at $2.2\,\mu$m in the 10 Leo spectrum. These lines are magnetically sensitive, and their Zeeman broadening can be exploited in cool stars to measure the magnetic field strength \citep[cf.][]{johns96,shulyak}, particularly in M dwarfs. These lines are among the sparse atomic transitions in the NIR with accurate available data, whereas magnetically sensitive molecular lines in the NIR (e.g. FeH) often lack the appropriate data. Atomic lines such as the Na lines in Fig.~\ref{telluric+magnetic} serve as a benchmark when comparing field measurements from molecular and atomic lines, and are of particular relevance for later-type stars, where optical magnetic proxies become inaccessible. The exploitation of these lines is challenged by the broadening effect of stellar rotation in later M stars and requires very accurate continuum normalisation. \cite{shulyak} has shown that without accurate telluric modelling, magnetic field measurements in early- to mid-type M stars does not yield reliable results. In the upper panel of Fig.~\ref{telluric+magnetic} we show the telluric model superimposed on the observed spectrum, which illustrates the presence of blends of telluric and stellar lines in the observed spectrum. Although the telluric blends affecting the Na lines are small, removing them is of great importance for correct stellar line fitting. Forthcoming spectral types from the CRIRES-POP library, for instance V2500 Oph (Barnard's star; M4 V), show many more weaker atomic (e.g. Na, Ti, Fe) and molecular lines (e.g. FeH). Magnetic field measurements in these lines, given the CRIRES-POP resolution, can be expected to greatly benefit from our telluric correction process. Many of the CRIRES-POP targets also appear in the UVES-POP library, meaning that these stars will have high resolution spectra with a wavelength coverage from 3000\,\AA \ to $5\,\mu$m. 10 Leo appears in both libraries, and as there is some overlap in the wavelength coverage of UVES and CRIRES, we can compare the spectra of this star taken with the two different instruments. The UVES-POP wavelength coverage ends at $1\,\mu$m, and that of CRIRES-POP starts at $0.962\,\mu$m. A small region of this overlap is shown in Fig.~\ref{uvescrires} as an example. The most striking difference is in the telluric lines, which are present in the UVES-POP spectrum and have been removed from the CRIRES-POP spectrum. The profiles of the stellar lines agree very satisfactorily, showing consistency between the two projects and the very real possibility of using the combined libraries for analysis. \begin{figure} \centering \includegraphics[width=0.5\textwidth]{UVES-CRIRES.pdf} \caption{Comparison of the UVES-POP (blue) and CRIRES-POP (red) spectra of 10 Leo in a small wavelength region of the overlap between the two spectra. The UVES-POP spectrum has not been corrected for telluric lines, and this causes the majority of the differences between the two spectra.} \label{uvescrires} \end{figure} Our next step is to produce the full CRIRES-POP stellar atlas of 10 Leo. In addition to the spectrum presented here, the atlas will include comprehensive line identifications, isotopic ratios, abundances of major species, and stellar properties. Some preliminary line identifications, made by overplotting the Arcturus line list on our spectrum, are shown in Figs.~\ref{finalspec1}~-~\ref{finalspec7}. The rest of the atlases for all stars in the CRIRES-POP library (see Table~\ref{librarytable}) will follow. The central philosophy of the CRIRES-POP project is universal access to high-quality data. The entire library of atlases, including the data used to produce them, will be freely available online. The raw and pipeline-reduced spectra are available on the CRIRES-POP webpage, and the spectrum of 10 Leo presented in this paper is also now available. The final atlas data will be provided in ASCII format, and possibly in a visual format similar to the printed Arcturus atlas. | 16 | 9 | 1609.07873 |
1609 | 1609.02498_arXiv.txt | { The muon content of extensive air showers is an observable sensitive to the primary composition and to the hadronic interaction properties. The Pierre Auger Observatory uses water-Cherenkov detectors to measure particle densities at the ground and therefore is sensitive to the muon content of air showers. We present here a method which allows us to estimate the muon production depths by exploiting the measurement of the muon arrival times at the ground recorded with the Surface Detector of the Pierre Auger Observatory. The analysis is performed in a large range of zenith angles, thanks to the capability of estimating and subtracting the electromagnetic component, and for energies between $10^{19.2}$ and $10^{20}$ eV. \PACS{ {96.50.sd}{98.70.Sa} {13.85.Tp} } % } % \authorrunning{L. Collica for the Pierre Auger Collaboration} \titlerunning{Measurement of the Muon Production Depths at Auger} | \label{intro} The spectrum and arrival directions of Ultra High Energy Cosmic Rays (UHECRs) above $10^{18}$ eV have been recently measured with unprecedented precision \cite{bib:valino,bib:arr}. The flux of cosmic rays at these energies is very low (less than 100 particles $\mathrm{km^{-2} yr^{-1}}$) and their origin is still not well understood. Establishing the cosmic-ray composition at the highest energies is of fundamental importance from the astrophysical point of view, since it could discriminate between different scenarios of origin and propagation of cosmic rays. Moreover, mass composition studies are of utmost importance for particle physics. As a matter of fact, knowing the composition helps in exploring the hadronic interactions at ultra-high energies, inaccessible to present accelerator experiments. \\ UHECRs properties cannot be determined from direct detection, due to their low flux, but must be inferred from the measurements of the secondary particles that the cosmic-ray primary produces in the atmosphere. These particles cascades are called Extensive Air Showers (EAS) and can be studied at the ground by deploying detectors covering large areas. \\ Composition studies on a shower to shower basis are challenging because of the intrinsic shower-to-shower fluctuations which characterise shower properties. These fluctuations come from the random nature of the interaction processes, in particular the height of the first interaction. However, showers originating from different primaries can be distinguished, at least statistically, given their different cross sections with air nuclei and distinct hadronic multiparticle production properties. Masses may be inferred from comparisons of the measured observables with predictions for these same observables from Monte Carlo simulations. These simulations rely on hadronic interaction models, which extrapolate interaction details from measurements in the accelerators domain to much higher energies and to different kinematic regions. Therefore, comparisons with simulations constitute the most prominent source of systematic uncertainties. \\ Information about the composition of the primary cosmic rays has been obtained using the Fluorescence Detector (FD) of the Pierre Auger Observatory \cite{bib:auger}. The FD allows the measurement of the depth at which the electromagnetic component of the air shower reaches its maximum number of particles, $X_\mathrm{max}$ \cite{bib:aab}. This observable is sensitive to the nature of the primary particles, as well as the standard deviation of its distribution, $\sigma(X_\mathrm{max})$. The interpretation of these measurements is hampered by uncertainties in hadronic interaction models. Besides, the number of events detected with the FD at high energy is low, due to the small FD duty cycle (about 15\%); the stringent cuts imposed to avoid a biased data sample in the analysis, such a field of view cut, further reduce the available statistics. \\ To gain additional information about mass composition and investigate the validity of the current hadronic interaction models, independent measurements with larger statistics are needed, together with a different set of systematic uncertainties and the possibility of reaching higher energies. \\ The Pierre Auger Collaboration has developed different methods to infer the composition of UHECRs through the measurements performed with the Surface Detector (SD), which has 100\% duty cycle. Among them, the study of the atmospheric depth at which the muon production rate reaches a maximum in air showers exploits the fact that the muon production depth is one of the most sensitive observables to the primary mass \cite{bib:MPD}. In addition, muons are sensitive to hadronic interactions since they come from the decay of pions and kaons, which form the hadronic core, and suffer small energy losses and angular deflections on their way to the ground.\\ In this paper a method for the reconstruction of the muon production depth for zenith angles between $45^{\circ}$ and $65^{\circ}$ and energies greater than $10^{19.2}$ eV is presented. The method is based on the model of muon time distributions discussed in \cite{bib:MPD} but it exploits a new kinematic delay parametrisation, tuned on post-LHC models, and a different technique to estimate the electromagnetic component. The latter allows one to reconstruct the muon arrival times at the ground closer to the shower core and for lower zenith angles, thus improving the muon sampling and the range of applicability of the analysis. \\ The paper is organised as follows. Sect.~\ref{sec:auger} briefly describes the Pierre Auger Observatory. In Sect.~\ref{sec:model} the muon time distributions model is described in details. Sect.~\ref{sec:featMPD} gives an overview of the properties of the muon production depth distribution while Sect.~\ref{sec:method} discusses all the steps through which the muon production depth is reconstructed. Finally, in Sect.~\ref{sec:sum} a brief summary is given together with some comments on the possible application of the method. | \label{sec:sum} The Pierre Auger Observatory employs water-Cherenkov detectors to measure particle densities at the ground and therefore has a good sensitivity to the muon content of air showers. \\ By means of a model which relates the arrival time of muons at the ground with their production depths, a new approach to study the longitudinal development of the hadronic component of EAS has been established. Studying the muon profiles helps to improve our understanding of hadronic interactions at the highest energies and sets additional constraints on model descriptions. \\ In this paper we presented a method to reconstruct the muon production depth distribution on an event-by-event basis in the angular range $\theta=45^{\circ}-65^{\circ}$ and for energies $\mathrm{E}=10^{19.2}-10^{20}$ eV.\\ The large range of applicability of this analysis has been obtained thanks to the capability of extracting the muon time distribution in the SD stations for a wide range of distances from the core. The maximum of the muon production depth distribution, $X_{\mathrm{max}}^{\mu}$, is estimated with a systematic uncertainty of at most 17 $\mathrm{g/cm^{2}}$ and its measurement could be exploited to constrain the most recent LHC-tuned hadronic interaction models, QGSJETII-04 and EPOS-LHC, and could give insights about the nature of UHECRs. In addition, thanks to the high statistics made available, deeper insights in the hadronic interaction models can be achieved by the study of the angular dependence of $X_{\mathrm{max}}^{\mu}$, and its correlation with the electromagnetic counterpart, $X_{\mathrm{max}}$. | 16 | 9 | 1609.02498 |
1609 | 1609.09575_arXiv.txt | We investigate the $\gamma$-ray and X-ray properties of the Flat Spectrum Radio Quasar (FSRQ) \4c50 at redshift $z= 1.517$. The {\it Fermi}-LAT data indicate that this source was in an active state since 2013 July. During this active period, the source's emission appeared harder in $\gamma$-rays, with the flux having increased by more than a factor of three. We analyze two distinct flares seen in the active state and find that the variability is as short as several hours. The {\it Swift}-XRT data show that the source was variable at X-ray energies, but no evidence was found for flux or spectral changes related to the $\gamma$-ray activity. The broad-band X-ray spectrum obtained with {\it Swift}-XRT and {\it NuSTAR} is well described by a broken PL model, with an extremely flat spectrum ($\Gamma_{1} \sim 0.1$) below the break energy, $E_{\rm break} \sim 2.1~{\rm keV}$, and $\Gamma_{2} \sim 1.5$ above the break energy. The spectral flattening below $\sim 3$ keV is likely due to the low energy cut-off in the energy distribution of the photon-emitting electron population. We fit the broad-band spectral energy distribution of the source during both the active and quiescent states. The X-ray and $\gamma$-ray emission from the jet is mainly due to the inverse-Compton scattering process, with the seed photons provided from the broad line region, and the jet is estimated to be larger than the accretion power if the jet is mainly composed of electron-proton pairs. | \label{sec:intro} Blazars are radio-loud active galactic nuclei (AGNs) with relativistic jets pointing towards the Earth \citep{1978bllo.conf..328B}. Because of the Doppler beaming effect, emission from a jet dominates the broad-band spectral energy distribution (SED) from radio to $\gamma$-rays energies \citep{1995PASP..107..803U}. The SEDs usually have two broad bumps in a $\log\nu-\log\nu f_{\nu}$ diagram. While the low-energy bump usually peaks from infrared to X-ray energies, which are believed to be the synchrotron emission of non-thermal electrons, the high-energy bump peaks from X-ray to $\gamma$-ray bands, which is considered to be the inverse Compton (IC) emission of the same electron population. For the IC emission, the seed photons can come from the low-energy synchrotron emission, broad line region (BLR), or dusty torus \citep[see e.g.,][]{1981ApJ...243..700K, 1985ApJ...298..128B, 1992ApJ...397L...5M, 2000ApJ...545..107B}. Because of the synchrotron self-absorption effect, blazars tend to have flat radio spectra with spectral index $\alpha < 0.5$. As a subclass of blazars, flat spectrum radio quasars (FSRQs) have strong optical emission lines (equivalent width $>5${\AA}), comparing to BL Lac objects that show no or very weak emission lines \citep{1997A&A...325..109S}. In the current third \fermi\ Large Area Telescope (LAT) source catalogue (3FGL), the dominant extragalactic $\gamma$-ray sources are blazars \citep{2015ApJS..218...23A}. Extreme variability is not common to all blazars detected in $\gamma$-rays. The minimal variable timescale detected with {\it Fermi}-LAT has reached less than half an hour (e.g. PKS 1510-089, \citealp{2013A&A...555A.138F}) and the variation amplitude can be two orders of magnitude (e.g., 3C 454.3; \citealp{2011ApJ...733L..26A}). Detailed studies of spectra and variabilities are essential for determining the location and mechanism of radiation from the jets of the blazars. The FSRQ \4c50 (also known as NRAO 150) is one of the strongest radio and millimeter AGN sources in the northern sky \citep{1966ApJS...13...65P,2008ASPC..386..249A, 2010ApJS..189....1A}. The VLBI monitoring observations showed that the inner jet (inner 0.5 mas from the core) exhibits superluminal motions with $\beta_{app} \sim (6.3\pm1.1)c$ and a large, $>100^{\circ}$ projected misalignment of the jet within the inner 0.5 mas to 1 mas from the core \citep{2007A&A...476L..17A, 2014A&A...566A..26M}. These properties imply that a relativistic jet points toward the Earth with a very small viewing angle \citep{2007A&A...476L..17A}. \citet{2010A&A...519A...5A} measured the redshift using near-IR spectroscopic data (exhibiting strong H$\alpha$ and H$\beta$ emission lines), and derived the cosmological redshift $z=1.517 \pm 0.002$, which corresponds to the luminosity distance $d_{L}=11.2 \times 10^{3}~{\rm Mpc}$. \citet{2010ATel.2517....1F} reported the detection of $\gamma$-ray emission from \4c50 with LAT on board the {\it Fermi} satellite. Using almost 20 months of data, he provided the $\gamma$-ray flux above 100 MeV, $F_{100 MeV} = 3.2 \pm 1.1 \times 10^{-8}~{\rm photons}~{\rm cm}^{-2}~{\rm s}^{-1}$, and photon index $\Gamma = 2.6 \pm 0.2$. After the $\gamma$-ray flaring activity around $\sim$MJD 56686 (2014 January 29), {\it Swift} target-of-opportunity observations were performed \citep{2014ATel.5838....1C, 2014ATel.5878....1K}. For the purpose of fully studying this high-energy source, we collected its \fermi-LAT and available X-ray data, which include 15 {\it Swift} observations and one {\it NuSTAR} observation, and performed detailed analysis of the data. In this paper, we present the results from our analysis. In the following, \S~2 describes the data analysis of the {\it Fermi}-LAT, {\it Swift} and {\it NuSTAR} observations. The obtained temporal and spectral results are presented in \S~3 and \S~4, respectively. We discuss the overall properties of the source in \S~5, including fitting to its multiwavelength SED, and we summarize our results in \S~6. | \subsection{Gamma-ray Properties} \label{gama_discu} We have studied the $\gamma$-ray properties of \4c50 by analysing the {\it Fermi}-LAT data, and confirmed the prediction in \citet{2010A&A...519A...5A} that \4c50 is a luminous $\gamma$-ray emitter. The observed $\gamma$-ray photon index of \4c50\ has a range of $\Gamma\approx 2.4-3.0$ (see Table~\ref{gamma}), which is roughly consistent with those of the {\it Fermi}-LAT $\gamma$-ray FSRQs ($\langle\Gamma\rangle\approx 2.4-2.5$; \citealt{2015ApJ...810...14A}). From the temporal analysis, we found that \4c50 has been in an active state since 2013 July. During the active period, the $\gamma$-ray flux increased by $>3$ times compared to the quiescence level and the emission was harder. Moreover, two distinct $\gamma$-ray flares were well seen in the $0.1-1$ GeV light curve during this period. Our temporal analysis has shown that \4c50\, exhibited variability on the time scale of as low as several hours, which is not commonly seen for high redshift blazars. A blazar jet is produced at the central region around the super-massive black hole (SMBH), and as the inner region can not be resolved with current telescopes (note that thus far, M87 is the only source resolved with the current observing facilities, which reaches several Schwarzschild radius, see \citealt{2011Natur.477..185H}), variability is a useful feature for probing this region. Given the variability timescale of \4c50, the causality implies that the size of the emission region is $R = t_{var} c\delta/(1+z) = 4.2 \times 10^{14} (\delta/10)(t_{var}/1\ {\rm h})\ {\rm cm} = 1.36 \times 10^{15}\ {\rm cm}$ (taking $t_{var}\simeq 4$ h and $\delta=7.9$; for the $\delta$ value, see below), which is comparable to the Schwarzschild radius. The central BH mass of this source is $\approx4.68\times10^{9}$ M$_{\odot}$ and corresponding Schwarzschild radius is $1.38\times 10^{15}$ cm \citep{2010A&A...519A...5A}. Assuming the low-energy X-ray emission is produced from the same region, the $\gamma$-ray photons could be absorbed by X-ray photons through the pair production effect. The strength of this absorption is mainly dependent on the X-ray energy density, which will decrease if emission is relativistic Doppler beamed. Therefore, the observed $\gamma$-ray and X-ray data can be used to constrain the jet Doppler factor. Because the $\gamma$-ray photons actually escape from the emission region, the several-hours timescale constrains the lower limit of the beaming factor, $\delta \ge 11.8 [(1\ {\rm h}/t_{\rm var}) (1\ {\rm keV}/\epsilon_{X}) (L_{\epsilon_{X}}/10^{46}\ {\rm erg~s}^{-1})]^{1/4}$, where $\epsilon_{X} E_{\gamma}= 20.61 (\delta /10)^{2}$ \citep{1995MNRAS.273..583D}. Considering $\gamma$-ray photons with energies of $\sim$1~GeV and the X-ray luminosity of the source obtained in this study, the Doppler beaming factor $\delta \ge 7.9$ ($t_{var}\sim 4~{\rm h}$). The VLBI observations show that the apparent superluminal motion reaches $\beta_{app}=6.3c$ and the central jet changes the direction about $\sim 100^{\circ}$ \citep{2007A&A...476L..17A, 2014A&A...566A..26M}. Combining these with the assumption of $\delta=7.9$, we have estimated the viewing angle $\theta=7.9^{\circ}$ and the bulk Lorentz factor $\Gamma_{bulk}=6.5$, which suggest that the jet is highly relativistic and has a small viewing angle with respect to our line of sight. \subsection{X-ray Properties} \label{subsec:xray_discu} We have investigated the X-ray properties of \4c50 using the {\it Swift}-XRT and {\it NuSTAR} observations. The source showed variability in the long-term {\it Swift}-XRT light curve. While its intensity was at the high end of the variation range during the $\gamma$-ray flaring period (Figure~\ref{lc}; only the 2014 February {\it Swift} observation was conducted in the time period), no significant correlated activity was seen. The X-ray spectral parameters obtained in the active period did not have drastical changes either. We considered that X-rays and $\gamma$-rays are produced from the IC scattering radiation by the same electron population. Because the cooling timescale of electrons in the lower energy part (in X-rays) is longer than the timescale of higher energy part (in $\gamma$-rays), one can expect that the X-ray variability timescale would be longer than that of the $\gamma$-rays. The spectral flattening of the soft X-ray spectrum has been widely found in high-redshift radio loud quasars \citep[e.g.][and references therein]{2006MNRAS.368..985Y}. The flattening may be due to either the intrinsic absorption with column densities of the order of $10^{22} -10^{23}~\rm cm^{-2}$ or the low energy cut-off in the energy distribution of electron population in the jet \citep{2001MNRAS.323..373F,2001MNRAS.324..628F,2004MNRAS.350L..67W,2004MNRAS.350..207W}. In the excess absorption scenario, high $\rm N_{\rm H}^{\it z}$ may be the dense plasma in form of a wind or outflow \citep{1999MNRAS.308L..39F}. However in the radio-loud quasars like \4c50, the relativistic jet along the line-of-sight can remove the gas column efficiently. Indeed, the VLBI observations \citep{2007A&A...476L..17A} have revealed a jet toward the Earth, suggesting that the excess absorption scenario is not likely the case. If there is a low energy cut-off in the energy distribution of the electron population, a spectrum is expected to flatten in the soft energy band \citep{2001MNRAS.323..373F,2007ApJ...665..980T,2007ApJ...669..884S}. This scenario requests a broken PL model, where the cut-off in the soft X-ray band can be naturally explained as the intrinsic curvature of the spectrum near the low-energy end of the IC component. The soft X-ray flattening is then an intrinsic feature of a source. Among the four examined observations of \4c50, the spectra were relatively well described by the broken PL model of $\Gamma_{1}\sim -0.3$ -- +0.9 below the break energy $E_{\rm break} =$ 2.1 -- 3.3 keV, and $\Gamma_{2} \sim$ 1.4 -- 1.8. The simultaneous observations of \4c50 by {\it Swift} and {\it NuSTAR} showed that the broad-band X-ray spectrum is better modelled by a broken PL than by a PL or a PL plus extra absorption model. We found $\Gamma_{1} = 0.05^{+0.88}_{-1.71}$ below the break energy, $E_{\rm break} = 2.13^{+0.62}_{-0.31}$ keV, and $\Gamma_{2} = 1.52^{+0.04}_{-0.03}$. In the 14--195 keV energy range of the {\it Swift}-BAT, the source was found to have photon index $\Gamma_{\rm X} = 1.51\pm0.35$ \citep{2013ApJS..207...19B}, which is well in agreement with $\Gamma_{2}$ obtained in our broad-band fit. Thus we suspect that the flattening is likely the intrinsic feature of the source. This possibility is supported by the broad-band SED modelling (see the following Section 5.3 and Table~\ref{tab:modfit}). From the modeling, it can be known that the low energy of non-thermal electrons is about $\gamma_{min}\sim1.1$ and the Doppler beaming factor $\delta\sim22.5$ for the active state. The electrons around the minimum energy will IC scatter external seed photons, and emit at $\nu_{IC}\approx(4/3)\delta\Gamma_{jet}\gamma_{min}^{2}\nu_{ext}/(1+z)\sim2.6$ keV (assuming $\Gamma_{jet}=\delta$), which is roughly consistent with observations (similarly, we have $\nu_{IC}\approx2.1$ keV for the quiescent state). \subsection{Spectral Energy Distribution Fitting} \label{sed_discu} We collected the archival radio and optical data for \4c50 from \citet{2010A&A...519A...5A} and NED\footnote{\url{http://ned.ipac.caltech.edu/}} respectively. These data were combined with X-ray and $\gamma$-ray data in this work and the broad-band SED of the source is shown in Figure~\ref{sed}. In this SED, emission from the relativistic jet dominated except at optical wavelengths. The optical emission reached a peak luminosity of $\sim10^{47}$ erg s$^{-1}$ and appeared as a significant bump, which should be thermal arising from the optically thick accretion disk. As \4c50 hosts a very massive BH, the thermal disk emission reaches $\sim30\%$ of the Eddington limit \citep{2010A&A...519A...5A}. The broad-band SED is not simultaneous, except the two sets of {\it Swift} X-ray and the corresponding \fermi\ $\gamma$-ray data in the active and quiescent states (the red and blue squares, respectively, in Figure~\ref{sed}). Nevertheless, we modelled the broad-band SED by using a standard blazar emission model: one zone synchrotron plus inverse Comptonization model. This model was widely used in blazar SED modelling \citep[e.g.,][]{2010MNRAS.402..497G, 2012ApJ...748..119C,2017ApJ...842..129C}. The emission region is assumed to be a homogeneous sphere with radius $R$ embedded in the magnetic field $B$. A broken power-law electron energy distribution, \begin{figure*} \begin{center} {\includegraphics[width=0.8\linewidth]{fig7.eps}} \end{center} \caption{Broad-band SED of \4c50, the black and grey data points are separately from \citet{2010A&A...519A...5A} and NED. The red points (lines) stand for active states, while the blue ones for quiescent states. Solid lines are for EC/BLR model. The central accretion disk/corona/dust torus emission component is represented using green dash line. The red/blue dash, dot and dash dot lines are synchrotron, SSC and EC emission in active/quiescent states, respectively.} \label{sed} \end{figure*} \begin{table*} \centering \setlength{\tabcolsep}{4.0pt} \caption{SED Model parameters for 4C 50.11} \begin{tabular}{@{}lcccccccccccccc@{}} \hline \hline State & $B$ & $R$ & $N_0$ & $\delta$ & $\gamma_0$ & $\gamma_{min}$ & $p_1$ & $p_2$ & $P_{jet}^{e-p}$ & $P_{jet}^{e-e^{+}}$ & $L_{diks}$\\ & Gs & $10^{15}$cm & & & & & & & $10^{48}$erg/s & $10^{45}$erg/s & $10^{47}$erg/s \\ \hline high & 0.8 & 3.80 & 0.628E+07 & 22.16 & 34.80 & 1.1 & 2.2 & 3.3 & 4.616 & 5.930 & 1.8 \\ low & 0.8 & 4.29 & 0.422E+07 & 20.00 & 30.62 & 1.1 & 2.2 & 3.6 & 3.203 & 3.973 & 1.8 \\ \hline \end{tabular} \label{tab:modfit} \end{table*} \begin{equation} N(\gamma )=\left\{ \begin{array}{ll} N_{0}\gamma ^{-p_1} & \mbox{ $\gamma_{\rm min}\leq \gamma \leq \gamma_{0}$} \\ N_{0}\gamma _{\rm 0}^{p_2-p_1} \gamma ^{-p_2} & \mbox{ $\gamma _{\rm 0}<\gamma\leq\gamma_{\rm max}$,} \end{array} \right. \label{Ngamma} \end{equation} was assumed in our calculation. The parameters of this model include the radius $R$ of the blob, the magnetic field strength $B$, electron break energy $\gamma_{0}$, the minimum and maximum energy, $\gamma_{\rm min}$ and $\gamma_{\rm max}$, of the electrons, the normalization of the particle number density $N_{0}$, and the indices $p_{1,2}$ of the broken power law particle distribution, the jet Doppler factor (assumed to be equal to the bulk Lorentz factor), and the spectrum of the external seed photons. The frequency and luminosity can be transformed from the jet frame to observational frame as: $\nu=\delta\nu'/(1+z)$ and $\nu L_{\nu}=\delta^{4}\nu'L_{\nu'}'$, where the Doppler factor $\delta=1/\left[\Gamma\left(1-\beta\cos\theta\right)\right]$, and the prime represents the value measured in the jet frame. The synchrotron self-absorption and the Klein-Nishina effect in the IC scattering were properly considered in our calculations. Both the self-synchrotron Compton (SSC) scattering and external Compton (EC) scattering (external seed photons from the BLR and dusty torus were taken into account) were included in the calculation of the Compton scattering in the blob. As mentioned above, the optical emission is multi-temperature annular blackbody radiation arising from the accretion disk, which was modelled with a standard optically thick, geometrically thin disk (Shakura \& Sunyaev 1973). Above the accretion disk, the corona reprocessed a fraction of disk luminosity (fixed at a level of 10\%) and had a power-law spectrum with cut-off energy 150 keV (we fixed the spectral index $\alpha=1.0$). Because the accretion disk's radiation is de-beamed in the jet comoving frame, seed photons from it were not important and thus not included in the EC scattering. In our SED modelling, the luminosities of the BLR and dust torus were assumed to be a fraction of the disk luminosity, 10\% and 50\%, respectively \citep{2008MNRAS.387.1669G}. The radii of the BLR and torus were $R_{BLR}=10^{17}L_{disk,45}^{1/2}=0.43$ pc and $R_{torus}=2.5\times10^{18}L_{disk,45}^{1/2}=10.9$ pc, respectively ($L_{disk,45}=179.0$ is the disk luminosity in units of $10^{45}$ erg s$^{-1}$; see \citealt{2010A&A...519A...5A}). In this case, the external photon energy densities are typical values $U_{BLR}=2.65\times10^{-2}$ erg cm$^{-3}$ and $U_{torus}=2.12\times10^{-4}$ erg cm$^{-3}$. The size of the emitting region was assumed to be equal to the radius of a circular conic section, $R=\psi R_{diss}$ ($R_{diss}$ is the distance of the emission region from the central black hole, where $\psi=0.1$; see \citealt{2008MNRAS.387.1669G}). The variability timescale can be used to set an upper limit on the emission size due to the causality, $R\lessapprox c\Delta t\delta/(1+z)$. During our SED modeling, the minimum variability timescale ($\Delta t\approx$4h) was used for estimating the size of the emission region for the active state. Note that the Doppler factor estimated in Section~\ref{gama_discu}, $\delta\gtrsim7.9$, was the lower limit to avoid the absorption of $\gamma$-ray photons through electron pair production effect. In Figure~\ref{sed}, we show the model fits to the SEDs in both the active and quiescent states, with seed photons dominantly coming from the BLR. The model parameters are given in Table~\ref{tab:modfit}. From the jet bolometric luminosity $L_{jet}$, we can obtain the jet non-thermal radiation power \citep{2014Natur.515..376G}, $P_{\rm rad}\approx2L_{\rm tot}/\delta^{2}=3.1\times10^{46}$ erg s$^{-1}$ for the active state, which is about $\sim 17\%$ of the disk luminosity of $1.8\times10^{47}$ erg s$^{-1}$. The jet radiative efficiency is believed to be on order of $P_{\rm rad}/P_{\rm jet}\sim$10\%, which holds for AGNs, gamma-ray bursts, and even for black hole X-ray binaries \citep{2012Sci...338.1445N, 2013ApJ...774L...5Z, 2014ApJ...780L..14M}, which gives a jet power, $P_{\rm jet}\thickapprox10P_{\rm rad}=3.1\times10^{47}$ erg s$^{-1}$, larger than the disk luminosity. This suggests that the jet launching processes and the way of transporting energy from vicinity of the black hole must be very efficient. Actually, having the model parameters, the jet power can be calculated as, $P_{\rm jet}\simeq\pi R^{2}\beta \Gamma^{2}cU_{\rm tot}'$, where the total energy density measured in the rest frame of the jet, $U_{\rm tot}^\prime=U_{\rm e}^\prime+U_{\rm B}^\prime+U_{\rm p}^\prime$. The energy density for electrons $U_{\rm e}^\prime=m_{\rm e}c^{2}\int N(\gamma)\gamma d\gamma$, while the proton energy density $U_{\rm p}^\prime=U_{\rm e}^\prime(m_{\rm p}/m_{\rm e})/\langle\gamma\rangle$ if charge neutrality for pure hydrogen plasma is assumed. The estimated values for the jet powers are given in Table~\ref{tab:modfit}. It can be seen that the jet power $P_{jet}$ is larger than the disk luminosity $L_{disk}$ by more than one magnitude and even larger than that of accretion power $P_{acc}=L_{disk}/\eta\approx(0.6-1.8)\times10^{48}$ erg s$^{-1}$, where the radiative efficient of the accretion disk is assumed to be $\eta\approx0.1-0.3$ \citep{2014Natur.515..376G}. However, we note that the minimum electron energy in our fitting is small, $\gamma_{min} = 1.1$ (Table~\ref{tab:modfit}). The small value of $\gamma_{min}$ may result in the overestimation of the jet power. In SED modeling, the reproduction of X-ray emission is important to constrain $\gamma_{min}$. In \citet{2014ApJ...788..104Z,2015ApJ...807...51Z}, X-ray is produced through SSC mechanism, in which $\gamma_{min}$ is much larger than the unit. In our SED modeling, we failed to model the X-ray with the SSC emission; instead, following \citet{1998MNRAS.301..451G,2010MNRAS.402..497G}, we modeled the X-ray through EC emission and therefore obtained a small $\gamma_{min}$ (similar to \citealt{1998MNRAS.301..451G,2010MNRAS.402..497G}). In addition, it should be noted that the estimated jet power is largely dependent on the assumed jet components. For exmaple, if the jet is mainly composed of electron-positron pairs instead of electron-proton plasma, the jet powers will be significantly decreased and smaller than the accretion disk luminosity for both the active and quiescent states (see Table \ref{tab:modfit}). | 16 | 9 | 1609.09575 |
1609 | 1609.00381_arXiv.txt | \noindent PyTransport constitutes a straightforward code written in \CC \S together with Python scripts which automatically edit, compile and run the \CC \S code as a Python module. It has been written for Unix-like systems (OS X and Linux). Primarily the module employs the transport approach to inflationary cosmology to calculate the tree-level power-spectrum and bispectrum of user specified models of multi-field inflation, accounting for all sub and super-horizon effects. The transport method we utilise means only coupled differential equations need to be solved, and the implementation presented here combines the speed of \CC \S with the functionality and convenience of Python. This document details the code and illustrates how to use it with a worked example. It has been updated to be a companion to the second version of the code, PyTransport\,2.0, which includes functionality to deal with models of inflation with a curved field space metric. | \begin{framed} {\bf \noindent PyTransport is distributed under a GNU GPL licence. The most recent version can be obtained by visiting \href{https://transportmethod.com}{transportmethod.com}. If you use PyTransport you are kindly asked to cite Ref.~\cite{Dias:2016rjq} as well as the archive version of this user guide in any resulting works. } \end{framed} The main purpose of this document is to teach those interested how to use, and if so desired adapt, the PyTransport package. It has now been updated to be a companion to the second version of the code, PyTransport\,2.0, which includes functionality to deal with models of inflation with a curved field space metric. We have had to make some minimal syntax changes in this second version in order to support new functionality, as discussed below. Users of the original package will unfortunately not be able to switch to the new one without amending their scripts. The original user guide can still be found as the arXiv version~1 of this document. The philosophy behind the implementation is simplicity and ease of use. Python was selected as the language though which to interact with the code because it enables rapid scripting and provides a flexible and powerful platform. In particular, it has many readily available tools and packages for analysis and visualisation, and for tasks such as parallelisation (using for example Mpi4Py). As an interpreted language, however, Python can be slow for some tasks. This is circumvented here by using \CC \S code, which is compiled into a Python module, to perform numerically intensive tasks with the result that the speed of the package is nearly indistinguishable from pure \CC. The \CC \S code itself is kept as simple and clean as possible and can therefore easily be edited if required. PyTransport has been developed on OS X using Python 2.7. We have also performed limited testing on Linux systems, and attempted to ensure compatibility with versions of Python 3. It can also be adapted to Windows systems, but this functionality has not yet been incorporated into the released package\footnote{We thank Sean Butchers for work related to installing PyTransport on a Windows machine.} The code is intended to be a reusable resource for inflationary cosmology. It enables users to quickly create a complied Python module(s) for any given model(s) of multi-field inflation. The primary function of the complied module is to calculate the power-spectrum and bi-spectrum of inflationary perturbations produced by multi-field inflation. To this end, the module contains a number functions that can be called from Python and that perform tasks such as calculating the background evolution of the cosmology, as well as the evolution of the two and three point functions. We also provide a number of further functions written in Python that perform common tasks such as calculating the power spectrum or bispectrum over a range of scales by utilising the compiled module. The true power of the approach, however, is that users can rapidly write their own scripts, or adapt ours, to suit their own needs. The transport approach to inflationary perturbation theory that the code employs can be seen as the differential version of the integral expressions of the In-In formalism. It is helpful numerically because it provides a set of ordinary differential equations for the correlation functions of inflationary perturbations. The code solves these equations from deep inside the horizon until some desired time after horizon crossing using a standard variable step size ordinary differential equation (ODE) routine with error control. Such off the shelf routines are extremely well tested, and provide an easy way to change the required accuracy. This is helpful in order to check convergence of the numerical solutions, or to respond to needs of models with very fine features. Details of the transport method itself that the code is based on can be found in the recent papers \cite{Dias:2016rjq} and \cite{xxx2}, the second of which updates the method to allow for the analysis of models with a curved field space metric. We highly recommend reading this guide in combination with those papers. In this guide, we first we give some brief background and motivation for the code, much more can be found in Refs.~\cite{Dias:2016rjq,xxx2}, before giving an overview of its structure and how it can be set up. In the appendices we give some more detail about the structure of the underlying \CC \S code, give full details of all the functions the complied module provides, and all the functions provided by Python scripts which accomplish common tasks. The best way to learn how to use the package, however, is by example. We present an extended example below spread between the ``Getting going" and ``Examples" sections, complete with screen shots of the code in use. Other examples that come with the distribution are discussed in the ``Examples" section. Throughout, familiarity with Python and to some extent \CC \S is assumed, though in reality users can just probably get a long way by looking at the examples and modifying to their needs. Finally, we would also like to refer readers to the complementary package developed in tandem with the work in Ref.~\cite{Dias:2016rjq} and with PyTransport: CppTransport \cite{Seery:2016lko}. This is a platform for inflationary cosmology developed fully in \CC and recently also updated to deal with curved field space metrics \cite{seeryNew}. In comparison with PyTransport it has more external dependancies (in the sense that the dependancies of PyTransport are mainly Python modules), but provides more sophisticated parallelisation and data management capabilities. In limited testing it is also found to be marginally faster. For users with modest aims in terms of CPU hours and data generation, however, it is likely to have a higher overhead in getting started, but may well be beneficial for intensive users. PyTransport is intended to be more lightweight with users encouraged to utilise the power of Python in combination with PyTransport to achieve their specific aims and data management needs. | We have presented the PyTransport package for the calculations of inflationary spectra. This package compliments and extends currently available tools, and is also complimentary to two related packages mTransport \cite{Dias:2015rca} and CppTransport \cite{Dias:2016rjq,Seery:2016lko} all described at this \href{https://transportmethod.com}{website}\footnote{https://transportmethod.com}. In its most recent version it has also now been updated to include functionality to deal with models that have a curved field space metric \cite{xxx2}. Through use of a detailed example we have shown how PyTransport can be used in practice. We have also summarised the structure of the code, with some more details provided in the appendices. | 16 | 9 | 1609.00381 |
1609 | 1609.00724_arXiv.txt | Quasi-Separatrix Layers (QSLs) are a useful proxy for the locations where current sheets can develop in the solar corona, and give valuable information about the connectivity in complicated magnetic field configurations. However, calculating QSL maps even for 2-dimensional slices through 3-dimensional models of coronal magnetic fields is a non-trivial task as it usually involves tracing out millions of magnetic field lines with immense precision. Thus, extending QSL calculations to three dimensions has rarely been done until now. In order to address this challenge, we present \qsl \ -- a public, open-source code, which is optimized for calculating QSL maps in both two and three dimensions on GPUs. The code achieves large processing speeds for three reasons, each of which results in an order-of-magnitude speed-up. 1) The code is parallelized using OpenCL. 2) The precision requirements for the QSL calculation are drastically reduced by using perturbation theory. 3) A new boundary detection criterion between quasi-connectivity domains is used, which quickly identifies possible QSL locations which need to be finely sampled by the code. That boundary detection criterion relies on finding the locations of abrupt field-line length changes, which we do by introducing a new Field-line Length Edge (FLEDGE) map. We find FLEDGE maps useful on their own as a quick-and-dirty substitute for QSL maps. \qsl allows constructing high-resolution 3D FLEDGE maps in a matter of minutes, which is two orders of magnitude faster than calculating the corresponding 3D QSL maps. We include a sample of calculations done using \qsl to demonstrate its capabilities as a QSL calculator, as well as to compare QSL and FLEDGE maps. | \label{sec:Intro} Many important questions in solar physics concern phenomena that take place in the low-$\beta$ environment of the corona, such as flares and coronal mass ejections (CMEs), active region (AR) evolution and dynamics, heating of the corona and sources of the solar wind. For studying these, it is often useful to have a model of the 3D magnetic field in the corona as it still cannot be observed and mapped directly. Such models can be potential \citep{Jiang12a}, linear force-free field \citep[LFFF; constant-$\alpha$; e.g.][]{Demoulin94, Abramenko96, Jiang12a}, or non-linear force-free field \citep[NLFFF; $\alpha(r)$; e.g.][]{vanBallegooijen04, Wiegelmann04, Valori05, Wheatland06, Schrijver06, Malanushenko12, Inoue12, Jiang12b}. Potential field source surface models have been in use for a long time and although well representative of the structure of the field at large heights in the corona, they by definition lack currents, and hence free energy, which is important for powering solar eruptions. An alternative are NLFFFs, which have gained significant popularity in recent years with the advent of numerous methods for their computation, which use either line-of-sight or vector photospheric magnetograms to produce a model of coronal magnetic fields or to extrapolate the observed photospheric magnetic field to the corona. However, these 3D magnetic fields are intrinsically complicated and although basic topological features, such as null points (NPs), fan-spine surfaces and flux ropes (FRs) can often be approximately identified just by inspecting field line plots, often, there is a need of quantitative topological analysis in order to make sense of the complicated 3D coronal magnetic field structure, its dynamics, and evolution. Topological features in 2D and 2.5D, such as NPs \citep[e.g.][]{GorbatchevSomov88, Parnell10}, separatrices \citep{GorbatchevSomov88}, separator field lines, and null lines have been explored in solar physics context since the 80s. They are known to separate the field in connectivity domains. However, in the mid-90s a new topological term arose, namely quasi-separatrix layers \citep[QSLs;][]{Priest95, Demoulin96b}, which are the 3D generalizations of the above-mentioned features, now separating the field into quasi-connectivity domains. While the linkage of magnetic field lines over separatrices and NPs is discontinuous, across QSLs it is continuous but drastically changes. In the early description of \cite{Demoulin96a}, the strength of QSLs, i.e. the amount of the change in field line linkage, is measured by the norm of the Jacobian of the mapping of neighboring field lines \citep{1994ApJ...437..851L} from one end of the photosphere to the other. However, this quantity is not invariant with respect to the direction of tracing of the field lines. Consequently, \cite{Titov07} came up with an alternative covariant quantity quantifying QSL strengths, called the squashing factor, $Q$. QSLs (as well as NPs and separatrices) are preferential sites for build-up of current sheets in the presence of footpoint motions, and hence are preferential sites where reconnection can take place \citep{Aulanier05b}. Although, there is still no quantitative relationship \citep[possibly because it depends on the exact field configuration and footpoint motions; see][]{Galsgaard03b}, it has been suggested that the higher the value of $Q$, the thinner the current layer at that particular QSL \citep{Aulanier05b}. This makes these topological features very important for studies of storage and release of magnetic free energy in the process of reconnection at all scales. Quantitative studies of topology by deriving QSL maps in 2D from potential, LFFFs, and NLFFFs have been used over the past decade to tackle many problems in solar physics. The existence of a QSL wrapping around the flux rope and crossing itself at a very high-$Q$ topological feature, a hyperbolic flux tube \citep[HFT;][]{Titov07, Savcheva12a, Savcheva12b, Zhao14, Liu14}, has become the basic feature in the standard flare model in 3D \citep{Aulanier12, Janvier13}, confirmed by observations \citep{Janvier14, Savcheva15, Savcheva16a, Janvier16, Zhao16}. In this picture, tether-cutting reconnection happens at the HFT under the FR between J-shaped oppositely directed field lines, which slip \citep{Aulanier06b} over the photospheric traces of the HFT \citep{Janvier13} and produce S-shape field lines that feed the FR and post-flare arcade. This scenario was put forward supported by data-constrained NLFFF models and MHD simulations by \cite{Savcheva12b}. In this picture, the photospheric traces of the HFT are 2J-shaped \citep{Titov07, Aulanier10} and they match the 2J-shaped flare ribbons of classical two-ribbon flares \citep{Chandra09,Schrijver11}. The match between the shapes of QSLs and flare ribbons has been achieved recently by \cite{Liu14}, \cite{Savcheva15}, and \cite{Zhao16}. These QSLs have been shown to move together with the flare ribbons in direction perpendicular to the polarity inversion line (PIL) \citep{Savcheva16a, Janvier16}. The QSLs derived in \cite{Savcheva16a} and \cite{Janvier16} have been derived based on NLFFFs constrained only by pre-flare observations (magnetograms, and EUV and X-ray images), but have managed to reproduce the flaring topology and its evolution to a large extent. That indicates that these kinds of studies have potential predictive power as the use of topology analysis can show us the likely sites of flare reconnection a few hours before the event, as shown in \cite{Savcheva12a}. One could imagine going further and using the flare ribbon information and QSLs to work backwards and improve the initial NLFFF, thus providing better initial conditions for global data-driven MHD simulations of CME initiation and propagation \citep{Savcheva16b}. Further studies show the evolution of QSL maps of solar ARs over several days, noting the effects of: flux cancellation on building sigmoidal flux ropes (\cite{Savcheva12a}, who showed the transition from bald-patch separatrix surfaces \citep[BPSS;][]{Titov93} to a HFT); quadrupolar topology on the possible breakout scenario \citep{Zhao14}; flux emergence on the development of a fan-spine NP topology \citep{Jiang16}. The global topology of active regions before eruption has been shown to be important for the characteristics of the dynamics, be it an eruption or just loop reconfiguration \citep{Janvier16, Jiang16, Pontin16, Chintzoglou16}. Knowing the locations, extent, shape, and connections between connectivity domains, and the features they contain or border, can prove vital for understanding links between seemingly unconnected faraway regions on the Sun that erupt sequentially or almost simultaneously, i.e. sympathetic eruptions. A detailed study of one such event (1-2 August, 2010) was conducted by \cite{Titov12}, who showed that filaments embedded in neighboring pseudostreamers are activated sequentially after the first filament erupts and destabilizes the system \citep{Torok11}. Even if it is a single CME, the potential of the CME to have a large longitudinal extent or to present with a significant energetic particle signature at any point in the heliosphere is most probably dependent on the specifics of the global 3D topology in the corona and heliosphere as the CME evolves and propagates \citep{Masson13}. As a related phenomenon, the propagation of EUV dimmings may also turn out to be dependent on the global solar topology, neighboring the directly related AR \citep{Downs16}. On a smaller scale, reconnection at QSLs have been potentially found important for the heating of the solar corona \citep{Schrijver10}. Reconnection in loop braiding has been theoretically and numerically explored for this purpose as well \citep{WilmotSmith09a, Pontin15}. QSLs in the outskirts of ARs have been shown to drive plasma outflows \citep{Baker09} as evidenced by blueshifts in Hinode/EIS velocity maps of ARs, which could be important for understanding the outflow of plasma from the corona that contributes to the slow solar wind. Potential solar wind sources can be further derived my means of the S-web model of \cite{Archontis09}, which utilizes QSLs at the source surface and below to look at the connectivity domains surrounding active regions and coronal holes, as well as the connections between them. Ultimately, with the speed-up and automation of NLFFF codes and QSL computation methods, we will be able to implement 3D QSL analysis in space weather predictive operations aimed at identifying the next likely region to erupt, studying the effect of the propagation of the CME ejecta and its particles, and predicting the direction and sign of the CME magnetic field when it reaches the Earth's magnetosphere. One step on this path is obtaining a fast, reliable 3D QSL code that can work on the whole Sun or in an AR in great detail. Such codes have been developed and used before for analyzing potential coronal magnetic field models \citep[][]{0004-637X-806-2-171}, as well as experimental flux rope configuration \citep[][]{PhysRevLett.103.105002}, yet they were never made public. In this paper, we introduce a fast, freely-available, open-source code, \qsl, aimed at calculating 3D QSL maps, whose development was motivated by several potential uses, such as: \begin{itemize} \item Studying large resolution QSL physics and its application to reconnection theory. \item Exploring large parameter spaces of possible topologies. \item 3D studies of active region evolution, CME initiation and propagation. \item Obtaining the evolution of topology over large periods of time with high cadence from data-driven or idealized MHD simulations at a wide range of scales. \end{itemize} The paper is organized as follows. In Section~\ref{over} we give an overview of the code. In Section~\ref{algo} we give details about the algorithm used in \qsl. We show illustrative results in Section~\ref{ex} and give our concluding remarks in Section~\ref{summary}. | \label{summary} In this paper we presented \qsl: a free, publicly available, open-source code for fast calculation of Quasi-Separatrix Layer maps in two or three dimensions. It requires an input magnetic field sampled on a rectilinear grid in Cartesian or spherical coordinates. We benchmarked the code by calculating 3D QSL maps for a model of the SOL2010-04-08 sigmoidal region on a consumer workstation GPU (AMD W8100). We found that the code achieves large processing speeds for three main reasons, each of which results in an order-of-magnitude speed-up: \begin{itemize} \item Running the code on the GPU as opposed to the workstation CPU results in about an order of magnitude speed-up. \item Compared to previous studies \citep[e.g. ][]{Pariat12}, we drastically relax the precision requirements for the QSL calculation. We do that by applying perturbation theory when calculating field-line deviations, which are necessary for calculating the squashing factor, quantifying the QSL strength. \item We use a new boundary detection criterion between quasi-connectivity domains, which quickly identifies possible QSL locations which need to be finely sampled by the code. That boundary detection criterion relies on finding the locations of abrupt field-line length changes. A map of these jumps in field-line length we dub a FLEDGE map. We find that using such FLL jumps as a refinement criterion, instead of a threshold in $Q$ (or its second derivative), results in an order of magnitude speed-up of the code. \end{itemize} For the realistic model discussed above, we clocked \qsl at several million $Q$ values per minute, which implies that a representative 3D QSL map can be obtained within a few hours. We also presented a quick-and-dirty alternative to QSL maps: FLEDGE maps, which can be optionally output by \qsl. We show that, for the most part, FLEDGE maps and QSL maps identify similar topological features. Constructing high-resolution 3D FLEDGE maps with \qsl can be completed in minutes -- two orders of magnitude faster than calculating the corresponding 3D QSL maps. The main reason for this difference is the fact that one does not need to perform adaptive refinements when computing FLEDGE maps as jumps in field-line length are readily identified even at low resolution, unlike local spikes in $Q$. Thus, we argue that FLEDGE maps offer a computationally cheap substitute of QSL maps that can be especially useful in the preliminary stages of any (quasi-)topological studies. The potential advantages to the solar physics community of having such freely-available, open-source codes are largely unexplored beyond published data-reduction pipelines. One of our goals in making \qsl public is stimulating others to get involved in a collaborative effort to produce codes open to inspection and verification. This has the benefit of avoiding the duplication of coding efforts and waste of public resources, as well as decoupling the scientific and coding efforts. \qsl can be found at \url{https://bitbucket.org/tassev/qsl_squasher/}. | 16 | 9 | 1609.00724 |
1609 | 1609.02735_arXiv.txt | {Massive binary systems are important laboratories in which to probe the properties of massive stars and stellar physics in general. In this context, we analysed optical spectroscopy and photometry of the eccentric short-period early-type binary HD~152218 in the young open cluster NGC~6231. We reconstructed the spectra of the individual stars using a separating code. The individual spectra were then compared with synthetic spectra obtained with the CMFGEN model atmosphere code. We furthermore analysed the light curve of the binary and used it to constrain the orbital inclination and to derive absolute masses of $(19.8 \pm 1.5)$ and $(15.0 \pm 1.1)$\,M$_{\odot}$. Combining radial velocity measurements from over 60 years, we show that the system displays apsidal motion at a rate of $(2.04^{+.23}_{-.24})^{\circ}$\,yr$^{-1}$. Solving the Clairaut-Radau equation, we used stellar evolution models, obtained with the CLES code, to compute the internal structure constants and to evaluate the theoretically predicted rate of apsidal motion as a function of stellar age and primary mass. In this way, we determine an age of $5.8 \pm 0.6$\,Myr for HD~152218, which is towards the higher end of, but compatible with, the range of ages of the massive star population of NGC~6231 as determined from isochrone fitting.} | In recent years, it has been found that the majority of the massive stars belong to binary or higher multiplicity systems \citep[][and references therein]{SF}. This situation has considerable implications for our understanding of the evolution of massive stars, but at the same time, it also offers enormous possibilities to observationally constrain the properties of these stars. Particularly interesting are unevolved double-line spectroscopic eclipsing binary systems (so-called SB2Es). The joint analysis of the photometric eclipses and radial velocities inferred from spectroscopy allows us to determine absolute masses and radii that are essentially model independent. Moreover, for unevolved systems, the properties of these stars should be good proxies of those of single massive stars. Of special interest are eccentric SB2Es that show significant apsidal motion \citep[e.g.][]{Bulut,Schmitt}. The latter arises from the fact that the gravitational field of stars in a close binary system can no longer be approximated as the gravitational field of a point-like mass. As a result, the orbits of the stars can only to first order be described as closed ellipses. A better description is achieved when it is considered that the argument of periastron $\omega$ undergoes a secular perturbation $\dot{\omega}$ \citep[e.g.][and references therein]{Schmitt}. The rate of apsidal motion is directly related to the internal structure of the components of the binary \citep[e.g.][]{Shakura,CG92,CG10}, thus allowing us to obtain information on the internal mass-distribution of the stars. Measuring the rate of apsidal motion can also provide an estimate of the masses of the components of non-eclipsing eccentric close binaries \citep{Benvenuto,Ferrero}, although this method is strongly model dependent. To date, there are only very few massive stars for which detailed studies of apsidal motion have been performed. Only four systems among the 128 eccentric eclipsing binaries listed by \citet{Bulut} host at least one star with a mass of more than 20\,M$_{\odot}$. In this paper, we discuss a fifth system: HD~152218. HD~152218 is a member of the young open cluster NGC~6231 in the Sco\,OB1 association. With a binary fraction of at least 0.63, the O-star population of this cluster appears to be very rich in binary systems, especially short-period systems \citep{NGC6231}. HD~152218 is a binary system consisting of an O9\,IV primary and an O9.7\,V secondary \citep{Sana}. The system has a rather high eccentricity ($e \simeq 0.28)$ given its orbital period of about 5.604\,days \citep{Hill,Stickland,Sana}. \citet{Otero} reported HD~152218 to display shallow eclipses. They classified the system as an Algol-type eclipsing binary, but did not perform a detailed analysis of its light curve. Comparing their radial-velocity curve obtained from data collected between 1997 and 2004 with previously published solutions from the literature \citep{Struve,Hill,Stickland}, \citet{Sana} noted the probable presence of a significant apsidal motion with $\dot{\omega}$ between 1.4 and 3.3$^{\circ}$\,yr$^{-1}$, depending on the data sets that were considered. {\it XMM-Newton} data furthermore revealed a modulation in the X-ray flux of HD~152218 with orbital phase that was interpreted as a result of a wind-wind interaction \citep{Sana,GRYN}. In the present paper, we re-address the issue of apsidal motion in the HD~152218 binary system. In Sect.\,\ref{observations} we introduce the observational data. The radial velocities and individual spectra of the binary components are analysed in Sect.\,\ref{spectroscopy}, notably to establish the rate of apsidal motion. Section\,\ref{photometry} presents our analysis of the photometric light curve, whilst Sect.\,\ref{theory} compares the observed rate of apsidal motion to predictions from theoretical models. Finally, in Sects.\,\ref{discussion} and \ref{conclusions} we discuss our results and present our conclusions. | } We have re-analysed the massive binary HD~152218 in the young open cluster NGC~6231. Spectral disentangling allowed us to reconstruct individual spectra and hence to derive stellar temperatures and gravities comparing the individual spectra with CMFGEN model atmospheres. Photometric data were used to constrain the orbital inclination and the Roche-lobe filling factors. Radial velocity data from the literature allowed us to establish a rate of apsidal motion of $(2.04^{+.23}_{-.24})^{\circ}$\,yr$^{-1}$ , corresponding to a period of 176 years. Comparison of this rate with the predictions of stellar structure models yields an age of $5.8 \pm 0.6$\,Myrs for HD~152218. This value is towards the higher end of the range of ages of the isochrones that encompass the location of the O-star population of NGC~6231 in a Hertzsprung-Russell diagram. | 16 | 9 | 1609.02735 |
1609 | 1609.00730_arXiv.txt | Constrained realisations of Gaussian random fields are used in cosmology to design special initial conditions for numerical simulations. We review this approach and its application to density peaks providing several worked-out examples. We then critically discuss the recent proposal to use constrained realisations to modify the linear density field within and around the Lagrangian patches that form dark-matter haloes. The ambitious concept is to forge `genetically modified' haloes with some desired properties after the non-linear evolution. We demonstrate that the original implementation of this method is not exact but approximate because it tacitly assumes that protohaloes sample a set of random points with a fixed mean overdensity. We show that carrying out a full genetic modification is a formidable and daunting task requiring a mathematical understanding of what determines the biased locations of protohaloes in the linear density field. We discuss approximate solutions based on educated guesses regarding the nature of protohaloes. We illustrate how the excursion-set method can be adapted to predict the non-linear evolution of the modified patches and thus fine tune the constraints that are necessary to obtain preselected halo properties. This technique allows us to explore the freedom around the original algorithm for genetic modification. We find that the quantity which is most sensitive to changes is the halo mass-accretion rate at the mass scale on which the constraints are set. Finally we discuss constraints based on the protohalo angular momenta. | \citet[][hereafter HR]{HR} presented a fast technique to build constrained realisations of Gaussian random fields. This method is exact and applies as long as the constraints can be expressed in terms of linear functionals of the random field. The algorithm has been widely used to generate `special' initial conditions for numerical simulations of structure formation, either by requiring the presence of uncommon features like high-density peaks \citep[e.g.][]{VB, RD06} or by imposing sets of observational constraints to reproduce the large-scale properties of the local universe \citep[][and references therein]{GH, Sorce16}. Recently, \citet[][hereafter RPP]{RPP} applied the HR algorithm to modify the initial conditions within the Lagrangian patches that form dark-matter haloes in numerical simulations (protohaloes). The basic idea is to alter the linear density field in a controlled way so that to produce `genetically modified' haloes (or, possibly, even galaxies) with some desired properties (e.g. the final mass or the merging history). Although the concept is intriguing, its practical implementation is problematic due to the complexity of characterising the statistical properties of protohaloes. This was already realised by \cite{MB} who considered (and then abandoned) the idea of pursuing a similar approach (see their Appendix A) in order to build analytical models aimed at explaining the origin of the seemingly universal halo mass-density profiles. This paper digs deeper into the matter. In Section \ref{const}, we review the theory of constrained random fields and provide several examples of increasing complexity. These are intended to guide the less experienced reader through the topic but also set the notation and provide the mathematical background to understand the rest of the paper. Some of the examples we give are unprecedented and form the basis for new applications. In Section \ref{gmh}, we demonstrate that the original execution of the genetic-modification idea by RPP is approximate because it suffers from the implicit assumption that protohaloes sample a set of random points with a fixed mean overdensity. We show that an exact implementation of genetic modification requires a mathematical understanding of the process of halo formation and in particular of the physics that sets the locations of protohaloes in the linear density field. Using toy models rooted on the idea that protohaloes might be associated with local maxima of the smoothed density field, we explore the degrees of freedom of genetic modification and clarify the meaning of probability of a constraint. Our results suggest new ways to enforce constraints within protohaloes. In Section \ref{secmah}, we illustrate how the excursion-set method \citep[e.g.][]{BCEK, Zentner} can be used to predict the accretion history and the final mass of the genetically modified haloes. This provides us with a tool to calibrate the constraints to set in order to produce a given growth history. We also use this method to estimate the size of the deviations in the assembly history of the haloes from the solution presented in RPP. We find that the quantity which is most affected is the mass-accretion rate at the mass scale of the constraints. Finally, in Section \ref{am}, we discuss how to set constraints based on the angular momentum of the haloes and, in Section \ref{con}, we conclude. | \label{con} The HR method provides an efficient tool to generate constrained realisations of Gaussian random fields in which certain linear functionals of the field variables assume pre-defined values. Although this technique has been around for 25 years, many researchers are not very familiar with it and still see it as arcane or esoteric. Motivated by the intent to improve this situation, in Section \ref{const}, we reviewed the basic principles of the HR method and made a number of examples for its application to cosmology, including peak-based constraints. We hope that our analytical results will provide a useful reference and help revealing the intrinsic simplicity of the algorithm. In Section \ref{gmh}, we discussed `genetically modified' haloes. RPP applied the HR algorithm to modify the initial conditions of $N$-body simulations within and around the regions that collapse to form dark-matter haloes. The gist of their initiative is to alter the linear density field at will so that to produce haloes with a set of desired properties after the non-linear evolution. At first sight, this project might appear a relatively straightforward application of the HR method. However, it contains a subtle complication: the points at which the constraints are applied are chosen after inspecting the unconstrained realisation. They are the Lagrangian locations at which haloes form and they must preserve this property after being genetically modified. From the mathematical point of view, this is equivalent to restricting the ensemble over which averages in the HR method should be taken in order to build the conditional mean field. RPP have disregarded this issue and used averages taken over the full ensemble. In other words, they treated protohaloes as randomly selected points with a given overdensity in Lagrangian space. This implicit assumption made the calculation possible but the results that follow from it are likely to suffer from a statistical bias. Our paper provides a first step towards understanding this issue. What makes the problem so challenging is that we do not know yet how to characterize protohaloes in mathematical terms. Although it is currently impossible to find an exact answer, reasonable lines of attack have been presented in the literature. Two common assumptions are that i) the Lagrangian sites for halo formation coincide with local density maxima of the smoothed density field \citep[e.g.][BBKS]{Doroshkevich70,K84, PH} and ii) the boundaries of protohaloes correspond to isodensity surfaces \citep[e.g.][]{HP, CT}. Detailed tests against $N$-body simulations give strong support to the validity of the first hypothesis, at least for haloes above the characteristic collapsed mass at each epoch \citep{LP}. On the other hand, protohaloes' principal directions and shapes have been found to strongly correlate with the local tidal field rather than with the density distribution \citep{LeeP, PDH2, LHP09, LP, Despali13, LBP}. All this suggests that it should be possible to characterize (at least to some extent) the properties of protohaloes in terms of the following variables: the density contrast, its first and second spatial derivatives, and the tidal field. Using the HR method we derived an analytical formula for setting simultaneous constraints on all these quantities. Our result is given in Eqs. (\ref{quasimain}) and (\ref{main}) while Eq. (\ref{chimain}) can be used to evaluate the relative probability of the constrained realisations with respect to the original one. If one wants to make sure that a protohalo in the unconstrained initial conditions, $\du$, remains a protohalo in the constrained linear density field, $\dc$, only some of the relevant field variable should be allowed to vary while some others should be kept fixed. There is some freedom here. For instance, one might want to require that a density peak in $\du$ stays a peak in $\dc$ (i.e. $\overline{\nabla \delta}=0$ and $\overline{\mH}$ is negative definite). With this in mind, we showed that the field transformation that sets a pure density constraint and marginalises over all the other field variables (Eq. (\ref{sol1dens}) which has been used by RPP) corresponds to setting correlated constraints in $\bar{\delta}$ and the mean curvature $\bar{\kappa}/3$ when the condition of being a local extremum and the traceless Hessian matrix are kept fixed. Although the expression of the HR correction is identical in these two cases, the likelihood of the constrained realisations is quite different. This demonstrates that $\Delta\chi^2$ values should be interpreted with care as they depend on the assumptions that are made on the nature of the constraints. We also provided several additional examples including the case in which a density constraint is imposed while keeping the density gradient, the Hessian matrix and the tidal field fixed, Eq. (\ref{eur}). In the second part of the paper (Section \ref{secmah}) we have developed a variant of the excursion-set formalism in order to predict the mass-accretion history of GM haloes. This is key to optimising the choice of the constraints that should be set in order to produce haloes with the desired properties after their non-linear collapse. Our method does not require any external input and can be used with all sorts of constraints. Basically, we first compute the change in the excursion-set trajectory induced by the HR method and then solve for the first-upcrossing of a threshold which has been calibrated using the mass-accretion history of the original unconstrained run. The entire algorithm is very simple to code and essentially takes no time to run. For constraints that require small changes we derived an analytical expression for the final halo mass which is given in Eqs. (\ref{masspred}) and (\ref{masspred2}). Our analysis indicates that, after all, the implementation by RPP generates halo mass accretion histories that are qualitatively similar to those obtained assuming a correspondence between protohaloes and local density maxima, at least on galaxy scales (see Figure \ref{fig4}). However, we found that the mass-accretion rate at the mass scale of the constraints is very sensitive to the detailed form of the imposed restrictions. This suggests that the method used by RPP might be suitable for investigating broad evolutionary scenarios but care should be taken when using it to make precise quantitative measurements. Future studies should test our semi-analytic results against $N$-body simulations. In particular, they should measure how big of an effect is obtained when additional conditions on the density gradient, the Hessian matrix and the tidal field are combined with the pure density constraints used by RPP. Finally, in Section \ref{am}, we discussed the possibility of using the HR method to constrain the angular momentum that a halo gains to leading order in perturbation theory. We concluded that this is impossible to achieve because the shape of protohaloes depends on the initial conditions in an unknown (and thus unpredictable) way. Nevertheless, the HR method can be used to set constraints based on the angular momentum gained by a fixed Lagrangian region. We derived the corresponding analytical solution for patches centered on random points with a fixed overdensity which is given in Eqs. (\ref{LL}), (\ref{Ldelta}) and (\ref{amrandom}). We also demonstrated that this solution does not hold true for density maxima or, more generally, when information on $\overline{\nabla \delta}$ is used to identify the location of the constraints (and thus, most likely, for protohaloes). On the other hand, using the tidal-torque theory to first order, we reduced the angular-momentum constraints to tidal-field constraints that can more easily be imposed at special locations identified using spatial derivatives of the density field. In conclusion, we would like to express the wish that future investigations will focus more and more onto the problem of characterising the locations and properties of protohaloes. | 16 | 9 | 1609.00730 |
1609 | 1609.07110_arXiv.txt | When a circumbinary disk surrounds a binary whose secondary's mass is at least $\sim 10^{-2}\times$ the primary's mass, a nearly empty cavity with radius a few times the binary separation is carved out of the disk. Narrow streams of material pass from the inner edge of the circumbinary disk into the domain of the binary itself, where they eventually join onto the small disks orbiting the members of the binary. Using data from 3-d MHD simulations of this process, we determine the luminosity of these streams; it is mostly due to weak laminar shocks, and is in general only a few percent of the luminosity of adjacent regions of either the circumbinary disk or the ``mini-disks". This luminosity therefore hardly affects the deficit in the thermal continuum predicted on the basis of a perfectly dark gap region. | \label{sec:intro} Circumbinary disks are ubiquitous astronomical objects. They are often observed around young binary stars \citep[e.g.,][]{Dutreyetal1994,AW2005} and may have already been detected in disk-planet systems \citep{Reggianietal2014,Billeretal2014,Sallumetal2015}. They may also be present around supermassive binary black holes (SMBBHs) as galaxies merge \citep{BBR1980,Ivanovetal1999,MeMi2005}. When the mass ratio is not too far from unity, the binary clears out a low density gap in the disk center, crossed by one or two narrow streams. Emanating from the inner edge of the circumbinary disk, some of the matter in these streams reaches ``mini-disks" attached to the members of the binary, while some of the matter suffers strong torques and swings back out to the circumbinary disk. \citep{AL1994,GK2002,MM2008,Cuadraetal2009,Hanawaetal,deValBorroetal2011, Roedigetal2012,Shietal2012,DOrazioetal2013,Farrisetal2014,ShiKrolik2015, DOrazio2016}. A natural question to ask is therefore: does this distinctive region produce a characteristic light signal? There is some controversy about the answer to this question. A number of papers have argued that the gap region of a circum-SMBBH disk, including the streams, should be relatively dim, and would therefore produce a dip in the thermal disk spectrum over a specific range of wavelengths dependent upon the parameters of the system \citep{RoedigSesana2012,Tanakaetal2012,GM2012,Kocsisetal2012,TH2013CQGra, Roedigetal2014}. Others have pointed to events on either the inside or the outside of the gap as creating distinct radiative features. It has been suggested, for example, that because the accretion rate across the gap is in general modulated strongly at a frequency comparable to the orbital frequency, so, too, should the bolometric accretion luminosity from the mini-disks \citep{MM2008,DOrazioetal2013}. However, this is unlikely to occur because most of the luminosity from the mini-disks is made in their inner radii, and the inflow time from the rim of a mini-disk is much longer than the binary orbital period \citep{Farrisetal2014}. On the other hand, \citet{Roedigetal2014} pointed out that the shock created where streams from the circumbinary disk hit the outer rim of the mini-disks should be quite bright, radiating primarily in the hard X-ray band, and, unlike the accretion luminosity, its output should reflect the modulation because the local Compton cooling time is short. \citet{Nobleetal2012} identified another possible signature in thermal emission from the inner edge of the circumbinary disk due to shocks driven by returning streams. The radially integrated luminosity of this component is consistent with the amount of work done by the binary torque on the disk, as argued in \citet{Shietal2012}. Recently, however, \citet{Farrisetal2014} have computed thermal emission from the gap proper as well as its edges based on their 2D hydrodynamic simulations, reaching the surprising conclusion that, far from being dim, the streams can actually be more luminous than the inner region of the circumbinary disk. Unfortunately, this last effort depended upon an assumption with potentially significant implications for the result. Their heating rate was calculated on the basis of a phenomenological ``$\alpha$ model" shear viscosity. This model was invented by \citet{ShakuraSunyaev73} to mock up internal disk accretion stresses now known to be due to MHD turbulence driven by the magneto-rotational instability \citep{BalbusHawley98}. Because the dynamical state of the streams is quite different from the nearly-circular orbits inside an accretion disk, it is not at all clear whether this model is applicable to them. In this paper we set out to test these assumptions by measuring the heating rate in the streams as determined by a pair of 3-d MHD simulations first reported in \citet{ShiKrolik2015}. These simulations make no phenomenological assumptions relevant to dissipative processes, and should therefore provide an accurate measure of local heating whether due to small-scale turbulence or coherent dissipative features like shocks. Despite their assumed isothermal equation of state, we can also evaluate $pdV$ temperature changes from their stored velocity data. | \label{sec:implications} Given that shock heating in the streams is at most a few percent of the accretion luminosity on the radial scale of the gap, and our best estimate of turbulent dissipation within the streams is that it should be $< 1\%$ of the shock heating, it seems unlikely that their luminosity can be more than a rather small fraction of the ordinary accretion luminosity produced either near the inner edge of the circumbinary disk or near the outer edges of the mini-disks around the members of the binary. For this reason, we disagree with the claim made by \citet{Farrisetal2015}, who applied an ``$\alpha$-viscosity" model to stream fluid dynamics, that the streams should, if anything, be brighter than the adjacent portions of the disks. Instead, our results support the prediction of \citet{Roedigetal2014} that the thermal continua of accreting binaries should be marked by a ``notch" centered on the frequencies corresponding to the temperature an ordinary disk would have on a radial scale of order the binary separation. For typical parameters of supermassive binary black holes, the characteristic disk temperature scales as \begin{eqnarray} T_0 \!\!&=& \!\!1.9 \times 10^4 \left(\frac{\eta}{0.1}\right)^{-1/4}\left(\frac{\dot{M}}{\dot{M}_{\rm E}}\right)^{1/4} \!\!\! \\\nonumber \!\!&\times &\!\! \left(\frac{M}{10^8{\rm M_\odot}}\right)^{-1/4} \!\!\! \left(\frac{a}{10^2\,R_{\rm g}}\right)^{-3/4} \!\! {\rm K} \,, \label{eq:T0_unit} \end{eqnarray} where $\dot{M}_{\rm E} \equiv L_{\rm E} /(\eta \, c^2)$, with $\eta$ denoting the accretion energy conversion rate, $R_{\rm g}$ the Schwarzschild radius, $\dot{M}$ the outer disk accretion rate, and $L_{\rm E}$ the Eddington luminosity. The typical energy scale shown in Figure~\ref{fig:spectra} is then $\sim 0.1$--$10\,{\rm eV}$ if the binary separation is $\sim 100 R_{\rm g}$, $\eta\sim 0.1$, and $\dot{M}$ and $M$ are not too far from $\dot{M}_{\rm E}$ and $10^8 {\rm M_{\odot}}$ respectively. We can also estimate the typical optical depth in the gas streams based on their surface density in our simulations. In code units, we find this varies within the range $\sim 0.1$--$0.5\Sigma_0$). If we assume that electron scattering opacity dominates, the implied optical depth is \beq \tau_{\rm T} \simeq (92{\rm -}460)\left(\frac{\eta}{0.1}\right)^{-1}\left(\frac{\dot{M}}{\dot{M}_{\rm E}}\right)\left(\frac{a}{10^2\,R_{\rm g}}\right)^{-1/2}\,. \label{eq:tau_sca} \enq Thus, for $\eta \sim 0.1$ and $\dot{M}\lesssim \dot{M}_{\rm E}$, we expect Thomson optical depths $\tau_{\rm T}\sim 100$ when the binary separation is about $100 R_{\rm g}$. Adopting Kramer's law for absorption opacity, we have \begin{eqnarray} \tau_{\rm a} \!\!&=& \!\!3.2\times 10^{22}\,\Sigma^2 H^{-1}\, T_{\rm eff}^{-7/2} \\ \nonumber \!\!&\simeq &\!\! (2.5 {\rm -} 65)\times 10^{-4} \left(\frac{\eta}{0.1}\right)^{-9/8} \left(\frac{\dot{M}}{\dot{M}_{\rm E}}\right)^{9/8}\!\! \\ \nonumber \!\!&\times&\!\! \left(\frac{M}{10^8 {\rm M_{\odot}}}\right)^{-15/8}\!\! \left(\frac{a}{10^2 R_{\rm g}}\right)^{13/8} \,. \label{eq:tau_abs} \end{eqnarray} Here we adopt $\Sigma=2\rho H$ and $H\equiv c_s/\Omega(r)=0.1 (r/a)^{3/2}\,a$ with $r=1.5a$ for the location of the stream. To obtain this number, we estimated the effective temperature using Equation~\ref{eq:teff_shock}; it is a factor of a few greater than $T_0$. Together with $\tau_{\rm T}$, we find the effective absorption optical depth is \begin{eqnarray} \tau_* \!\!&\simeq&\!\! \sqrt{\tau_{\rm T}\,\tau_{\rm a}} \\ \nonumber \!\!&\simeq&\!\! (0.15-1.73) \left(\frac{\eta}{0.1}\right)^{-17/16} \!\! \left(\frac{\dot{M}}{\dot{M}_{\rm E}}\right)^{17/16}\!\! \\ \nonumber \!\!&\times&\!\! \left(\frac{M}{10^8 {\rm M_{\odot}}}\right)^{-15/16}\!\! \left(\frac{a}{10^2 R_{\rm g}}\right)^{9/16} \,. \label{eq:tau_eff} \end{eqnarray} Again taking $\eta \sim 0.1$ and $\dot{M}$ close to $\dot{M}_{\rm E}$, we find the effective opacity is roughly order unity. However, the scale height of our simulation ($H/R \sim 0.1$) is likely an overestimate of the scale height in a realistic disk in these circumstances, so that the absorption optical depth, which is $\propto H^{-1/2}$, might be rather larger. If so, the spectrum radiated even by the streams would be reasonably well thermalized. The temperature at the stream surface should closely track the shock heating rate in the interior of the stream because the cooling time should be at worst comparable to the disk dynamical time, and likely shorter. Even using the simulation scale height, we find: \begin{eqnarray} t_{\rm cool}\Omega \!\!&\simeq&\!\! H\Omega\tau/c \\\nonumber \!\!&\simeq &\!\! (0.64{\rm -} 0.33) \left(\frac{\eta}{0.1}\right)^{-1}\!\! \left(\frac{\dot{M}}{\dot{M}_{\rm E}}\right)\!\! \left(\frac{a}{10^2 R_{\rm g}}\right)^{-1} \,. \label{eq:tcool} \end{eqnarray} In fact, the short cooling time gives some justification to our assumption of an isothermal equation of state in treating the shocks. We close this discussion with one last remark. In ordinary accretion disks, matter traversing a factor of several in radius must lose significant energy to do so; that raises the question of what happens to the energy the streams don't radiate. The answer is that when these shocks join the binary, they do so by striking the outer edge of one of the mini-disks surrounding each black hole. As \citet{Roedigetal2014} showed, there is a very strong shock where this happens, in which the Compton cooling time is very short. Thus, the energy not lost by slow thermal emission is instead lost rapidly by Compton cooling at the edge of a mini-disk. In a fashion attractive to observations, it would also be modulated with a frequency of order that of the binary. However, it plays no role in the thermal spectrum because the characteristic energy of the Compton-scattered photons is $\sim 100$~keV. | 16 | 9 | 1609.07110 |
1609 | 1609.09757_arXiv.txt | We explore the possible Proxima Centauri b's interiors assuming the planet belongs to the class of dense solid planets (rocky with possible addition of water) and derive the corresponding radii. To do so, we use an internal structure model that computes the radius of the planet along with the locations of the different layers of materials, assuming that its mass and bulk composition are known. Lacking detailed elementary abundances of the host star to constrain the planet's composition, we base our model on solar system values. We restrained the simulations to the case of solid planets without massive atmospheres. With these assumptions, the possible radius of Proxima Centauri b spans the 0.94--1.40 $\rearth$ range. The minimum value is obtained considering a 1.10 $\mearth$ Mercury-like planet with a 65\% core mass fraction, whereas the highest radius is reached for 1.46 $\mearth$ with 50\% water in mass, constituting an ocean planet. Although this range of radii still allows very different planet compositions, it helps characterizing many aspects of Proxima Centauri b, such as the formation conditions of the system or the current amount of water on the planet. This work can also help ruling out future measurements of the planet's radius that would be physically incompatible with a solid planetary body. | \label{sec:1} The recent discovery of a planet orbiting the Sun's nearest star, the red dwarf Proxima Centauri (Anglada-Escud\'e et al. 2016), has brought the search for extraterrestrial life on exoplanets to the gates of our solar system. The mass of the planet Proxima Centauri b (hereafter Proxima b), ranging close to that of the Earth, and the fact that it follows an orbit within a temperate zone (Anglada-Escud\'e et al. 2016), make real the possibility that has been found an Earth analog (or at least a habitable planet) within our nearest neighborhood. This planet was discovered from reanalysis of previous radial velocity measurements, which yield the minimum mass $m$~sin~$i$, where $i$ is the (unknown) orbital inclination angle. The reported minimum mass is 1.10--1.46 Earth masses ($\mearth$). We adopt this mass range here in our modeling (setting sin~$i = 1$) because it is astrobiologically interesting, and allows us to use equations of state well-tested for the Earth. Proxima b orbits its host star with a period of 11.2~days, corresponding to a semi-major axis distance of 0.05~AU. Unlike many other exoplanets discovered so far, this short distance does not imply a high surface temperature for the planet. Proxima Centauri being a red dwarf, its luminosity is only 0.15\% of that emitted by the Sun, and its effective temperature is 3,050~K. Therefore, a planet located at 0.05 AU from Proxima Centauri has an equilibrium temperature of only $\sim$234~K (Anglada-Escud\'e et al. 2016). This temperature is close to the melting point of water if one assumes that Proxima b is surrounded by an atmosphere with a surface pressure of one bar, implying that the planet lies within the habitable zone of its host star. Because no transit signal has been observed for Proxima b, its radius is unknown. Chances for having this quantity determined in the future remain low since Anglada-Escud\'e et al. (2016) predict a geometric probability of transit of only 1.5\% from their data. Here we explore the possible Proxima b's interiors assuming the planet belongs to the class of dense solid planets (rocky with possible addition of water) and derive the corresponding radii. To do so, we use an internal structure model that computes the radius of the planet along with the locations of the different layers of materials, assuming that its mass and bulk composition are known. We excluded from this study cases where Proxima b could harbor a thick atmosphere, unlike the Earth, to focus on planets that are likely to be habitable. | \label{sec:4} Assuming that Proxima b is solid, we find that its radius is in the $\sim$0.94--1.40 $\rearth$ range. The lower value corresponds to a planet made from 65\% iron located in the core and 35\% silicate present in the mantle, meaning that the planet surface is rocky, possibly surrounded by a thin atmosphere. The upper value corresponds to a planet made from 50\% silicate present in the mantle and 50\% water located in an outer layer, mainly constituted of high-pressure ices surrounded with a $\sim$200 km deep liquid ocean (6\% of the total water mass). Proxima b may also present a radius outside the aforementioned range, implying that it would be made from materials others than iron, silicates or water (e.g. if it harbors a thick atmosphere of H/He), or that particular formation and/or evolution conditions would have increased its CMF to a value higher than 65\%. Our model also helps placing constraints on the planet's formation conditions. Four different formation scenarios have been proposed for Proxima b, with only one case leading to a planet completely dry (planet formation on its current orbit via pebble accretion). The dry planet case corresponds here to WMF = 0 and this limitation reduces the allowed space of compositions to only one side of the ternary diagram shown by Figure~\ref{fig:figure1}. The range of radii for a completely dry Proxima b then becomes 0.94--1.19~$\rearth$. Note that any future measurement of the radius in this range does not necessarily mean that the planet is dry. The other formation scenarios (in situ accretion of planetesimals diffused from the outer parts of the system, or formation of a planetary body/planetary embryos beyond the snowline with following migration to the current orbit) predict a planet containing a significant amount of water and/or volatiles. If we assume this amount to be at least 10\% of the planet's mass (about the WMF of Europa; Sohl et al. 2002), then the smallest achievable planet radius becomes 1.02 $\rearth$ (with $M_P$ = 1.10 $\mearth$, and (CMF,WMF)=(0.59,0.1)). These cases almost cover the full ternary diagram, the band with WMF $<$ 10\% being excluded. Investigations of the irradiation conditions of Proxima b show that the loss of the planet's water inventory is an open question (Ribas et al. 2016, submitted). This implies that the large possible range of WMF we assumed is still relevant. The limitations we considered to disregard some regions of the ternary diagram are entirely based on solar system values. These assumptions could be irrelevant for the Proxima Centauri system, since the star is a red dwarf with metallicity [Fe/H] = 0.21. Complementary data on the Proxima Centauri system is of prime importance to better characterize Proxima b. In particular, high resolution spectroscopy of the host star and the determination of its Mg/Si and Fe/Si ratios would provide further constraints on the planet's interior. More generally, a higher precision on future measurements of Proxima b is essential to reduce the compositional variations allowed for this planet. | 16 | 9 | 1609.09757 |
1609 | 1609.02671.txt | We propose to use the flux variability of lensed quasar images induced by gravitational microlensing to measure the transverse peculiar velocity of lens galaxies over a wide range of redshift. Microlensing variability is caused by the { motions of the observer, the lens galaxy (including the motion of the stars within the galaxy), and the source}; %relative motion between the quasar source and the distribution of stars in the lens galaxy hence, its frequency is directly related to the galaxy's transverse peculiar velocity. The idea is to count time-event rates (e.g., peak or caustic crossing rates) in the observed microlensing light curves of lensed quasars that can be compared with model predictions for different values of the transverse peculiar velocity. To compensate for the large time-scale of microlensing variability % -- of the order of years -- we propose to count and model the number of events in an ensemble of gravitational lenses. We develop the methodology to achieve this goal and apply it to an ensemble of 17 lensed quasar systems %with published light curves . In spite of the shortcomings of the available data, we have obtained tentative estimates of the peculiar velocity dispersion of lens galaxies at $z\sim 0.5$, $\sigma_{\rm pec}(0.53\pm0.18)\simeq(638\pm213)\sqrt{\langle m \rangle/0.3 M_\odot} \, \rm km\, s^{-1}$. { Scaling at zero redshift we derive, $\sigma_{\rm pec}(0)\simeq(491\pm164) \sqrt{\langle m \rangle/0.3 M_\odot} \, \rm km\, s^{-1}$, consistent with peculiar motions of nearby galaxies} and with recent $N$-body nonlinear reconstructions of the Local Universe based on $\Lambda$CDM. We analyze the different sources of uncertainty of the method and find that for the present ensemble of 17 lensed systems the error is dominated by Poissonian noise, but that for larger ensembles the impact of the uncertainty on the average stellar mass may be significant. %In any case, in spite of the uncertainties, this is the first direct estimate of the peculiar velocity dispersion of galaxies at such high redshift and demonstrates the great potential of the method. Its application to thousands of gravitational lenses discovered and monitored with future facilities like the LSST will allow us to study the peculiar velocity dispersion over a wide range of redshifts and the evolution of the cosmological growth factor. | The motion of galaxies with respect to the smooth Hubble flow, i.e., the peculiar velocity field of galaxies, is a useful probe of cosmology and galaxy formation. Peculiar velocities allow us to trace the overall matter distribution (including dark matter) over a wide range of scales. On cosmological scales the coherent flows of galaxies toward overdense regions are determined by the overall amount of matter and its large scale distribution (Kaiser 1988). Consequently, galaxy peculiar velocities provide powerful tests of the cosmological model through measurements of the linear growth rate; these measurements are complementary to other cosmic probes\footnote{Specifically, while the expansion rate (constrained by geometrical probes such as CMB, BAO, and SNe Ia) can be consistently explained by different `dark energy' models, they predict measurable differences in the evolution of growth rate with cosmic time.} (Koda et al. 2014). On small scales, the random motions of galaxies are determined by the gravitational clustering of galaxies, allowing the study of the mass function of dark matter halos (Sheth 1996). Classical methods used to estimate peculiar velocities (e.g., Scrimgeour 2016 and references therein) compare the velocity derived from the source redshift with that obtained using the Hubble law combined with an independent distance indicator (such as Tully--Fisher, Faber--Jackson, the Fundamental Plane, or SNe Ia). The main drawback of these methods is the very large intrinsic scatter in distance estimate, which results in an error in the determination of peculiar velocities (from 5 to 20\% of the Hubble recession velocity, depending on the indicator) that grows linearly with distance. This limits the application of these methods to low redshift values ($z\lesssim 0.05$). Another technique to measure peculiar velocities, without these limitations in redshift, is the kinetic Sunyaev--Zeldovich effect that has recently started to be measurable (e.g., Lavaux et al.\ 2013). Alternatively, the effects of the peculiar velocity field can be studied in an indirect way from the anisotropic pattern of galaxy clustering (redshift-space distortion, Kaiser 1987) through the two-point galaxy correlation function, either focussed on small (e.g., Li et al.\ 2006) or large (e.g., de la Torre et al.\ 2013) scales of galaxy clustering. Finally, the peculiar velocity field can be also studied from numerical simulations, such as the recent high precision $N$-body reconstruction of the Local Universe by Hess \& Kitaura (2016), which has reconciled $\rm \Lambda$CDM with CMB-dipole measurements. Gravitationally lensed quasars provide a scenario in which the direct measurement of transverse\footnote{Note that the other direct methods measure line-of-sight velocities with subsequent large errors arising from the uncertainties in the subtraction of the Hubble flow.} peculiar velocities, even for $z$ greater than 1, is possible\footnote{Kochanek et al.\ (1996) also proposed to measure the peculiar velocity field but from a very different ground; the astrometric measurement of proper motions of gravitational lenses.} (Wyithe et al.\ 1999, Gil-Merino et al.\ 2005, Poindexter \& Kochanek 2010, Mediavilla et al.\ 2015). For an ideal galaxy with a smooth distribution of matter the flux magnification of one lensed image (mean magnification) is determined by the gravitational field and the lens system geometry. However, real galaxies include stars, which at parsec scales break the smoothness of the mass distribution and of the gravitational field, giving rise to large anomalies in the magnification. Thus, magnification can strongly change in the neighborhood of the lensed image, and, owing to the %relative motion between quasar and galaxy { effective motion of the quasar source}, the brightness of the image can experience fluctuations around the mean value (quasar microlensing, Chang \& Refsdal 1979, 1984; Wambsganss 2006). As this variability depends on the unknown spatial distribution of the stars, it is studied as a random process whose properties (in particular the spatial scales of variability) can be modeled. The basic idea of the present study is to measure the peculiar velocity of lens galaxies by comparing the fundamental frequency of the temporal variability of microlensed quasar images (inferred from observed light curves) with the fundamental frequency of the spatial variations induced by microlensing in modeled light curves (tracks on magnification maps that simulate microlensing variability). Both quantities should be related by the relative velocity between the lens galaxy stellar distribution, the source, and the observer. The simplest option to achieve this is to use time-event rate detection methods such as determining the zero-crossing or peak rates of the light curves. In favorable circumstances, this can be done by counting caustic crossings. In the neighborhood of a lensed source we find regions of more or less gentle magnification gradient and one-dimensional loci with very high magnification (caustic curves). The crossing of a caustic by the lensed source due to its %relative motion with respect to the galaxy { effective transverse motion} is the most conspicuous event of microlensing. This kind of event can appear, depending on the source size, as a very sharp feature in the lensed image light curve. Thus, for our purposes, caustics can be treated like randomly distributed milestones of known mean separation. The distance traveled by a source is proportional to the number of crossed caustics and, assuming Poissonian statistics, the typical deviation to its square root. Thus, we can compare the number of caustic crossings with the mean separation predicted by the models to infer the transverse velocity of the lens galaxy (Wyithe et al.\ 2000a, Gil-Merino et al.\ 2005, Mediavilla et al.\ 2015). One practical drawback of this method is that, given the typical transverse velocities, the mean temporal separation between caustics amounts to several years (see Einstein radius crossing times in Mosquera \& Kochanek 2011) and we need to count enough crossings to reduce the Poissonian noise. To surmount this problem, we propose to follow an ensemble of (properly selected) gravitational lens systems to add together a statistically significant number of crossings. Another problem is that, although caustics are intrinsically very sharp, they could be smeared by the source size. For large sources, the intensity of the caustic crossing can be drastically diminished and perhaps confused with other type of microlensing phenomenology (or two caustics can even be blended into one single event). Thus, the best option to identify and count caustics is to use as small a source as possible. The X-ray emitting region of an AGN would be the optimal choice but, unfortunately, the massive X-ray monitoring of lensed quasars seems unaffordable with present and planned facilities. At other wavelengths (mainly the optical) the microlensing events detected may not always be easily identified as a single caustic crossing. An obvious generalization that includes caustic crossing and other more complex events is to study the rate of Peaks Over a Threshold (POT), which have been considered in the literature mainly for the case of a relatively high threshold (High Magnification Events, HME; e.g., Kayser et al.\ 1986, Kundic \& Wambsganss 1993). Alternatively, we could also study the zero-crossing rate but, then, a no microlensing baseline is needed to set the zero. The main objetive of this paper is, thus, to propose and discuss the use of microlensing magnification event rates in an ensemble of gravitational lenses to estimate the galaxy peculiar velocities. In \S 2 the measurement of the transverse peculiar velocity from the statistics of microlensing events is introduced. In \S 3 we conduct an illustrative analysis based on POT counting in the microlensing light curves available in the literature. \S 4 is devoted to analyzing the sources of uncertainty, studying the selection of the ensemble of lens systems, and discussing future perspectives. Finally, the main results are summarized in Section 5. | We have proposed and developed a new method based on the count of microlensing events (peaks over a threshold) in an ensemble of lensed quasar images to infer the transverse velocity of lens galaxies and to study the peculiar velocity field. The main results are the following: 1 - To illustrate the method with a pilot study, we have considered an ensemble of 17 gravitationally lensed quasars with light curves available from the literature, to show explicitly that, in spite of the hetereogeneous quality of the data (in S/N ratio, time coverage, sampling, etc.), the application of the method is straightforward and consistent even when two different thresholds of 0.1 and 0.2 mag are alternatively considered to qualify the peaks. 2 - We have studied the impact of the uncertainties on the parameters involved in the modeling of the microlensing variability rate. We found that most of them have a weak impact and can be, in most cases, averaged out over the ensemble and/or controlled with a suitable selection of lens systems. The most important source of uncertainty is the mean mass of the stellar PDMF, which can nevertheless be expressed in an explicit way in the results (allowing itself to be easily marginalized with the best available information about the PDMF). 3 - Even with the obvious limitations of the data on which they are based, the results deserve attention. We found a tentative estimate of the peculiar velocity dispersion at $z\sim 0.5$, $\sigma_{\rm pec}(0.53)\simeq (638\pm 213) \sqrt{\langle m \rangle/0.3 M_\odot} \, \rm km\, s^{-1}$, independent of the cosmological model (except for the computation of angular distances) and, to our knowledge, the highest direct redshift determination of the peculiar velocity dispersion. 4 - Using $f \approx \Omega_m^{0.6}$ ({ Peebles 1980}, Lahav et al.\ 1991) to transform the velocity of each lens system to $z=0$, we have also obtained an estimate of the 1D peculiar velocity dispersion at zero redshift, $\sigma_{\rm pec}(0)\simeq (491\pm 164) \sqrt{\langle m \rangle/0.3 M_\odot} \, \rm km\, s^{-1}$. This result is compatible, to within the uncertainties, with { the results of local velocity surveys}, predictions of recent nonlinear $N$-body models of the Local Universe based on the standard $\Lambda$CDM cosmology and with the Local Group bulk flow. 5 - With the ensemble of 6000 monitored lensed images provided by the LSST, it will be possible to determine the growth factor rate dependence with redshift in the $z\sim 0$ to 2 range. With an adequate selection of lens systems, it would be possible to study the PDF of the peculiar velocity field. \appendix | 16 | 9 | 1609.02671 |
1609 | 1609.05053_arXiv.txt | Planet formation theories predict a large but still undetected population of short-period terrestrial planets orbiting brown dwarfs. Should specimens of this population be discovered transiting relatively bright and nearby brown dwarfs, the Jupiter-size and the low luminosity of their hosts would make them exquisite targets for detailed atmospheric characterisation with {\it JWST} and future ground-based facilities. The eventual discovery and detailed study of a significant sample of transiting terrestrial planets orbiting nearby brown dwarfs could prove to be useful not only for comparative exoplanetology but also for astrobiology, by bringing us key information on the physical requirements and timescale for the emergence of life. In this context, we present a search for transit-signals in archival time-series photometry acquired by the {\it Spitzer Space Telescope} for a sample of 44 nearby brown dwarfs. While these 44 targets were not particularly selected for their brightness, the high precision of their {\it Spitzer} light curves allows us to reach sensitivities below Earth-sized planets for 75\% of the sample and down to Europa-sized planets on the brighter targets. We could not identify any unambiguous planetary signal. Instead, we could compute the first limits on the presence of planets on close-in orbits. We find that within a 1.28 day orbit, the occurrence rate of planets with a radius between 0.75 and 3.25~R$_\oplus$ is $\eta < 67 \pm 1\%$. For planets with radii between 0.75 and 1.25~R$_\oplus$, we place a 95\% confident upper limit of $\eta < 87 \pm 3\%$. If we assume an occurrence rate of $\eta = 27\%$ for these planets with radii between 0.75 and 1.25~R$_\oplus$, as the discoveries of the Kepler-42b and TRAPPIST-1b systems would suggest, we estimate that 175 brown dwarfs need to be monitored in order to guarantee (95\%) at least one detection. | Radial velocity and transit surveys have revealed that close-in packed systems of low-mass planets are very frequent around solar-type stars \citep{Mayor:2011fj,Howard:2010zr,Batalha:2013lr} and red dwarfs \citep{Dressing:2013xy}. While still undetected, such systems could also be frequent around brown dwarfs. Observations confirm that brown dwarfs possess the same protoplanetary disc fraction as T\,Tauri stars do \citep{Scholz:2008kx}. In addition \citet{Ricci:2012fk,Ricci:2013lr} found evidence of dust growth, which has also been theoretically explored \citep{Meru:2013yu}. Planet formation is thus expected to occur within these discs \citep{Payne:2007lr}. Despite these encouraging signs, little evidence about short-period planets orbiting brown dwarfs exist, which is despite some of their very advantageous properties for planet detection, but can be explained by their intrinsic faintness. \begin{figure*} \centering \includegraphics[width=0.65\textwidth]{Paper_figures/BDs2S.pdf} \caption{For a given equilibrium temperature (here 255K, like Earth), the number of orbits (i.e. transits or occultations) per year (black), the transit depth (dashed red), and the probability of transit (dotted blue) as a function of the primary's mass. The green area is approximatively where a $5\sigma$ on spectral signatures can be reached with the mission lifetime of {\it JWST}. Stellar parameters were obtained from a 1 Gyr isochrone \citep{Baraffe:2003gf}. The slopes steepen with older stellar ages.} \label{fig:BD} \end{figure*} To this day, only a few long-period planetary mass objects have been associated to brown dwarf hosts using direct imaging and microlensing techniques \citep[e.g.][]{Chauvin:2004lr}. The mass ratios and orbital distances however make all these systems resemble substellar binaries more than planetary systems, at the exception of the $\sim 2$ $M_{\rm Jup}$ object OGLE-2012-BLG-0358Lb discovered by \citet{Han:2013qf}. The Carnegie Astrometric Planet Search Program attempts to detect gas giant planets around late M, L, and T dwarfs, as the smaller masses of these stars yield larger astrometric signals for a given planetary companion \citep{Boss:2009lr}. Another astrometric search for long-period planets around brown dwarfs is also currently under way \citep{Sahlmann:2014ys} and has demonstrated sensitivity reaching well into the planetary domain. \citet{Blake:2010yq} used the radial velocity technique and produced a null result. A first attempt to detect short-period terrestrial planets transiting nearby brown dwarfs was performed by \citet{Blake:2007qf}. They used the PAIRITEL infrared telescope to monitor a sample of 20 ultra-cool dwarfs, including some brown dwarfs. This survey did not detect any transiting object, and its precision was too low -- $\sim$1 \% -- with a sample too small to constrain the occurrence of short-period terrestrial planets around brown dwarfs. In 2013, the newly detected nearby binary brown dwarf Luhman-16AB \citep{Luhman:2013qf} was intensively monitored in photometry by the TRAPPIST telescope \citep{Jehin:2011dk,Gillon:2011qf} to search for transiting planets down to the radius of Earth. This project, that failed to detect any transit but revealed the fast-evolving weather of Luhman-16B \citep{Gillon:2013qv}, was done in the context of a transit survey targeting the $\sim$50 brightest Southern ultracool dwarfs ongoing since 2010 on TRAPPIST \citep{Gillon:2013qy}. This same survey identified recently a trio of Earth-sized planets transiting a nearby star with a mass only $\sim$10\% more massive than the Hydrogen-burning limit \citep{gillon:2016gh}. This recent discovery combined with the theoretical prediction that similar systems should be frequent around brown dwarfs is a strong motivation to intensify the search for transiting planets around the nearest brown dwarfs. In this context, we present here the results of the first space-based search for terrestrial planets transiting brown dwarfs, based on archive data gathered for a sample of 44 brown dwarfs by the {\it Spitzer Space Telescope}. In the next section we outline the importance of a search for planets transiting brown dwarfs. In Section~\ref{sec:sample} we describe the sample we use and in Sect.~\ref{sec:signal} perform early calculations on what type of planets can be detected. We then present a search algorithm in Section~\ref{sec:retrieval}, and test it using synthetically inserted transits. We apply this algorithm to first seek planets within the sample (Sect.~\ref{sec:search}) and then use it to compute upper limits on the occurence rate (Sect.~\ref{sec:occurence}). We then conclude. | \label{sec:discuss} We have searched for transiting planets signals within photometric timeseries obtained on 44 brown dwarfs by the Weather on other Worlds programme \citep{Metchev:2015rf}, that used the {\it Spitzer} space telescope. After studying our ability to recover transits with synthetic signals, we identify six transit event candidates. They have depths at the threshold of our detection criteria and are therefore likely to be false positives. We assumed the lightcurves contained no transits and generated a size/period map of upper limits on the occurence rates of planets orbiting brown dwarfs. \begin{figure} \centering \includegraphics[width=0.45\textwidth]{Paper_figures/irac.pdf} \caption{{\it Spitzer}'s photometric precision as a function of the magnitude of the source in IRAC's channel 1 and 2 (3.6 and 3.5 $\mu$m). We show a histogram of the magnitudes in the top panel, and indicate to which planet size, photometric precision correspond to (for a SNR of 3.5). } \label{fig:irac} \end{figure} Due to the limited number of targets, and short observation time, our results are not very constraining, but can inform us about how much more of an observational effort can be made to discover planets orbiting brown dwarfs, using the transit method. The discovery of an Earth-sized planet transiting a nearby and relatively bright brown dwarf would provide an exquisite target for the detailed characterisation of a terrestrial extrasolar world, one that will be similar to Earth in some aspects (size, temperature...), but different in others (irradiation, tidal locking...). With the current sample and assuming a 27\% occurence rate (as implied by the discoveries of TRAPPIST-1 \citep{gillon:2016gh} and Kepler-42 \citep{Muirhead:2012fk}), we estimate that 175 brown dwarfs need to be monitored for 21 hours, of which 44 have already been done. This sample of brown dwarfs, that composed the WoW survey, is centred on the L-T transition, the focus of most of brown dwarf variability. In addition many of these targets were not selected for their brightness (see Fig.~\ref{fig:irac}). A careful sample selection concentrating on brown dwarfs with magnitudes brighter than 12 (in both {\it Spitzer} channels) would likely lower this required number down. A doubling of the observing time only marginally helps transit recovery (also because we used a `single transit' detection approach). Finding one Earth-sized planet transiting a nearby brown dwarf, would provide an exquisite target for the detailed characterisation of a terrestrial extrasolar world using transmission \citep[like for TRAPPIST-1; ][]{gillon:2016gh,de-Wit:2016fk}, and emission spectroscopy. In addition, the discovery of multiple planet orbiting brown dwarfs may help answer fundamental questions in Astrobiology, yielding precious information on the physical requirement and timescale for the emergence of life. Stars evolve, which changes the irradiation received by a planet, and therefore its ability to retain liquid water at the surface. This implies that many planets remain within the habitable zone only for a while. Some enter it at a late stage in their history. This argument has been used by \citet{Ramirez:2016qy}. They argue that as solar-like stars leave the main sequence and evolve to the stage of giants, the habitable zone sweeps worlds that were previously too cold. We would counter-argue that as a star becomes a giant, it becomes harder to distinguish the planet from the star due to a worsening flux and size ratio. However the main idea, that of evolution remains valid, and very useful. Brown dwarfs evolve all the time. They shrink and cool down \citep[eg.][]{Burrows:1997fj,Baraffe:2003gf} constantly. Planets that were previously not within the habitable zone will eventually enter it, sometimes Gyrs after their formation. Identifying such planets, and finding traces of life inside their atmospheres would inform us that Archean conditions are not necessary for the emergence of life. In addition, the time a given planet will remain habitable will vary as a function of the mass of the brown dwarf host. Once a statistically significant sample of habitable worlds has been gathered orbiting a variety of brown dwarfs, it will become possible to measure a timescale for the emergence of biology from the amount of time spent being habitable and the relative number of planets. | 16 | 9 | 1609.05053 |
1609 | 1609.08567_arXiv.txt | We report a first measurement for ultra-high energy cosmic rays of the correlation between the depth of shower maximum and the signal in the water Cherenkov stations of air-showers registered simultaneously by the fluorescence and the surface detectors of the Pierre Auger Observatory. Such a correlation measurement is a unique feature of a hybrid air-shower observatory with sensitivity to both the electromagnetic and muonic components. It allows an accurate determination of the spread of primary masses in the cosmic-ray flux. Up till now, constraints on the spread of primary masses have been dominated by systematic uncertainties. The present correlation measurement is not affected by systematics in the measurement of the depth of shower maximum or the signal in the water Cherenkov stations. The analysis relies on general characteristics of air showers and is thus robust also with respect to uncertainties in hadronic event generators. The observed correlation in the energy range around the `ankle' at \logenr{18.5}{19.0} differs significantly from expectations for pure primary cosmic-ray compositions. A light composition made up of proton and helium only is equally inconsistent with observations. The data are explained well by a mixed composition including nuclei with mass $A > 4$. Scenarios such as the proton dip model, with almost pure compositions, are thus disfavoured as the sole explanation of the ultrahigh-energy cosmic-ray flux at Earth. | An important quantity to characterize the composition of cosmic rays is the spread in the range of masses in the primary beam. In theoretical source models regarding protons as the dominant particle type, the composition is expected to be (almost) pure, while in other scenarios also allowing heavier nuclei to be accelerated, a mixed composition is predicted. For instance, in the `dip' model~\cite{berezinsky_dip_prl2005,berezinsky_dip_ankle_2006}, two observed features of the energy spectrum could be naturally understood as a signature of proton interactions during propagation (ankle at $\lg(E/\mathrm{eV}) \simeq 18.7$ from pair-production and flux suppression at $\lg(E/\mathrm{eV}) \simeq 19.6$ from photopion production). Therefore, the dip model predicts an almost pure cosmic-ray composition with small spread in primary masses. In a recent publication, the distributions of depths of shower maximum \xmaxs\ (the atmospheric depth where the number of particles in the air shower reaches a maximum value) observed at the Pierre Auger Observatory were interpreted in terms of primary masses~\cite{longXmax_fits2014} based on current hadronic interaction models. The results suggest a mixed mass composition, but there are differences between the interaction models, and a clear rejection of the dip model is hindered due to the uncertainties in modeling hadronic interactions\footnote{For indirect tests of the dip model using cosmogenic neutrinos, see e.g.~\cite{dip_neutrinos2015} and references therein.}. Specifically, around the ankle, a very light composition consisting of proton and helium nuclei only is favoured using QGSJetII-04~\cite{qgsjetii04} and Sibyll~2.1~\cite{sibyll21}, while for \eposlhc{}~\cite{eposlhc}, intermediate nuclei (of mass number $A \simeq 14$) contribute. The spread of masses in the primary beam near the ankle, estimated from the moments of the \xmaxs\ distributions measured at the Pierre Auger Observatory~\cite{PAO_Xmax_JCAP2013,longXmax2014}, depends as well on the details of the hadronic interactions and the results include the possibility of a pure mass composition. Observations of \xmaxs\ by the Telescope Array in the northern hemisphere were found compatible within uncertainties to both a pure proton composition~\cite{TA_MD2014} and to the data from the Auger Observatory~\cite{WGmass_UHECR2014}. In this report, by exploiting the correlation between two observables registered simultaneously with different detector systems, we present results on the spread of primary masses in the energy range \logenr{18.5}{19.0}, i.e. around the ankle feature. These results are robust with respect to experimental systematic uncertainties and to the uncertainties in the description of hadronic interactions. | 16 | 9 | 1609.08567 |
|
1609 | 1609.04864_arXiv.txt | We derive the expression for the observed redshift in the weak field limit in the observer's past light cone, including all relativistic terms up to second order in velocity. We then apply it to compute the cluster-galaxy cross-correlation functions (CGCF) using N-body simulations. The CGCF is asymmetric along the line of sight (LOS) owing to the presence of the small second order terms such as the gravitational redshift (GRedshift). We identify two systematics in the modelling of the GRedshift signal in stacked clusters. First, it is affected by the morphology of dark matter haloes and the large-scale cosmic-web. The non-spherical distribution of galaxies around the central halo and the presence of neighbouring clusters systematically reduce the GRedshift signal. This bias is approximately 20\% for $M_{\rm min}\simeq 10^{14} {\rm M_{\odot}}/h$, and is more than $50\%$ for haloes with $M_{\rm min}\simeq 2\times 10^{13} {\rm M_{\odot}}/h$ at $r>$4 {\rm Mpc}/$h$. Second, the best-fit gravitational redshift profiles as well as the profiles of all other relativistic terms are found to be significantly different in velocity space compared to their real space versions. We discuss some subtleties relating to these effects in velocity space. We also find that the S/N of the GRedshift signal increases with decreasing halo mass. | In general relativity, photons receive a gravitational redshift when climbing out of potential wells. In the weak field limit, the magnitude of the redshift is proportional to the depth of the Newtonian potential $\Phi$. Photons from central galaxies sitting at the bottom of the potential well of galaxy clusters are expected to be gravitationally redshifted by a larger amount than satellites and other neighbouring galaxies. The difference of the gravitational redshift (GRedshift) signal with respect to the cluster centre is of the order of 10~km/s. It can in principle be detected by stacking a large sample of clusters. This has been predicted by \citep{Nottale1990, Cappi1995, Kim2004} and the first few tentative measurements from stacked clusters from SDSS data sets have been reported \citep{Wojtak2011, Sadeh2015, Jimeno2015}. In observations, the GRedshift signal extracted from stacked clusters is related to the distortion of the cluster-galaxy cross-correlation function (CGCF), or $\xi_{\rm cg}$, which originates from the distortions of the observed redshifts of galaxies with respect to the cluster centre (which may be the centroid of the galaxies or may be taken to be the brightest cluster galaxy (BCG)). In theory, ignoring the evolution of cosmic potentials and observational systematics, the observed redshift consists of five components: (1) the cosmological redshift (2) the 1st order Doppler redshift from the peculiar velocity of the galaxy (3) 2nd order special relativistic corrections from the peculiar velocity (4) the peculiar gravitational redshift (5) effects associated with the fact that we observe galaxies on our past light cone. The effects of (1) \& (2) result in an observed CGCF that should be front-back symmetric, while asymmetry of the CGCF along the line of sight will arise due to the presence of (3), (4) \& (5). The main goal of this study is to explore these effects on the CGCF and disentangle the GRedshift effect from them. There is also an additional effect, (6), the peculiar velocity of galaxies affects their surface brightness via beaming. Coupled to any surface brightness dependent selection (such as an apparent magnitude limit) this results in a bias of the redshift distribution of the selected galaxies at the same order of magnitude. This last effect, unlike the others, is highly dependent on details of the luminosity function of the galaxies and how they are selected in the surveys. Here we shall focus only on those effects that are independent of how galaxies are selected. On large scales, the relativistic corrections to the galaxy correlation function and the resulting asymmetry of the cross-correlation function between two different `tracer' populations, in our case clusters and galaxies, has been studied in \citet{Yoo2009, McDonald2009, Challinor2011, Bonvin2011,Yoo2012, Croft2013}, and in \citet{Bonvin2014} where some other effects such as density evolution and lensing are included. Our study will focus on the CGCF at around the scale of clusters and up to tens of Mpc/$h$. This is the (quasi-) non-linear regime where some of the theoretical predictions based on perturbative methods will break down. It is therefore necessary to employ N-body simulations for this study. A robust detection of the GRedshift signal may provide a constraint on theories of gravity. This requires an accurate prediction of the observed redshift. \citet{Wojtak2011}, for example, have modelled the effect by assuming a power-law mass function for clusters which are individually spherically symmetric and have a NFW \citep{NFW} profile, and that the observed redshifts are given as the sum of the first order Doppler shift and the gravitational redshift with respect to the cluster centre. It was subsequently realised that several additional physical processes, such as the transverse Doppler redshift, the past light cone effect and relativistic beaming, would cause additional contributions which are generally of the same order of magnitude as the GRedshift signal \citep{Zhao2013, Kaiser2013} and which complicate the analysis. These analyses, however, do not necessarily capture all of the relevant effects that need to be considered in order to make an accurate prediction. One shortcoming is that these analyses are not adequate to treat the `quasi-linear' regime -- outside the virial radius -- which is observationally relevant here. Another is that, of necessity, the quantity that is measured is a {\em galaxy weighted\/} measurement of the redshift; i.e.\ the mean of the gravitational redshift, plus other contributions, for galaxies at a given projected distance from the cluster galaxy centre. I.e.\ it is not the simple 2-point cluster density-potential cross correlation function, rather it is a third-order statistic $\langle n_c(0) n_g(\br) \Phi(\br) \rangle / \langle n_c(0) n_g(\br) \rangle$, where $n_c(0)$, $n_g(\br)$ and $\Phi(\br)$ are the number density of central galaxies at the origin, the number density of galaxies at $\br$ and the peculiar Newtonian potential at $\br$. Here we use N-body simulations to attempt to remedy these shortcomings. The outline of the paper is as follows: In the next section, we derive an expression for the observed redshift accurate to second order in the velocities (Hubble and peculiar) and to first order in the peculiar potential and allowing for the fact that we observe galaxies on the past light cone. This provides the redshift in terms of quantities defined on a hyper-surface of constant time, which is useful as the simulations provide snap-shots of the galaxy positions, velocities and the peculiar gravity on such hyper-surfaces. We analyse the simulations in \S\ref{sec:stacks}. This analysis reveals and quantifies two important new complicating factors. The first has to do with the fact that while, in a composite sense, clusters are spherically symmetric, individual clusters are aspherical and their surroundings are highly aspherical owing to the presence of neighbouring clusters. Coupled with the fact that the quantity one most naturally measures is the {\em galaxy weighted\/} redshift and clumps of galaxies are correlated with potential wells this results in a systematic bias which causes the weighted potential to increase more slowly with distance from the cluster centre than one would expect from simple models invoking an ensemble of spherical NFW profile clusters. The second effect has to do with the fact that the galaxies are observed in velocity space rather than in real space. | We have explored how the modelling of the gravitational redshift signal from stacked clusters is affected by a variety of systematics. $\bullet$ Since the GRedshift signal is a component on the observed redshift, we start by presenting the expression Eq.~(\ref{Eq:full2}) for the observed redshift on the past light cone of an observer including relativistic corrections. It is relative to the centre of a cluster, and is expressed in terms of properties on surfaces of constant proper time. The effect of the second order terms in this expression on the cluster-galaxy cross-correlation function are quantified using N-body simulations. We find that the the gravitational redshift term causes the strongest asymmetry of the CGCF. The recovered GRedshift signal is biased high by approximately 0.5-1~km/s depending on the minimum halo mass due to neglecting the other second order terms. This is relatively minor compared to the other two other systematics we have found. $\bullet$ The underlying gravitational potentials are usually deeper where there is a concentration of galaxies, which indicates a concentration of mass. The fact that observations of GRedshift are galaxy-number weighted causes the observed GRedshift signal to be biased low compared to models where volume weighting is assumed. This bias does not go away even if the stacked cluster is perfectly spherically symmetric. The non-spherical distribution of galaxies in individual clusters and the complex cosmic-web structures surrounding the cluster cause the bias to persist at nearly all scales of interest. This bias is stronger for lower mass clusters as the chance of having more massive neighbouring structures is higher. A pronounced bump at approximately 2~Mpc/$h$ from the cluster centre is expected for the observed GRedshift profile due to this bias. However, the bump tends to be flattened in velocity space. $\bullet$ Peculiar velocities of galaxies are the most dominant feature in the CGCF. The measurement of the GRedshift signal is in essence conducted in velocity space. It is strongly influenced by peculiar velocities since the observed galaxies are shifted from their original locations, e.g. galaxies at the bottom of the potential may appear far away from the cluster centre due to velocity-space distortions. This tends to flatten the bump of the GRedshift profile caused by the impact of neighbouring structures as mentioned in the previous bullet point. It also affects the predictions for all the other second order terms in Eq.~(\ref{Eq:full2}). $\bullet$ We find that the CGCF along the line of sight associated with the GRedshift signal is highly non-Gaussian. Therefore, extracting the signal by using a Gaussian function to fit for the peak positions of the CGCF as done in \citet{Wojtak2011,Sadeh2015} and \citet{Jimeno2015} may not be the optimal. There may be room for improvement in future analysis of this kind. The box-size of the simulation we use in this study is relatively small. The methods we have developed allow us to extract the relatively weak signal free from sampling variance. Simulations with larger box-size will be needed to study the noise properties. | 16 | 9 | 1609.04864 |
1609 | 1609.04263_arXiv.txt | {We present an intensive monitoring of high-resolution spectra of the Ca {\sc ii} K line in the A7IV shell star $\phi$ Leo at very short (minutes, hours), short (night to night), and medium (weeks, months) timescales. The spectra show remarkable variable absorptions on timescales of hours, days, and months. The characteristics of these sporadic events are very similar to most that are observed toward the debris disk host star $\beta$ Pic, which are commonly interpreted as signs of the evaporation of solid, comet-like bodies grazing or falling onto the star. Therefore, our results suggest the presence of solid bodies around $\phi$ Leo. To our knowledge, with the exception of $\beta$ Pic, our monitoring has the best time resolution at the mentioned timescales for a star with events attributed to exocomets. Assuming the cometary scenario and considering the timescales of our monitoring, our results indicate that $\phi$ Leo presents the richest environment with comet-like events known to date, second only to $\beta$ Pic. } | In spite of their low mass, Kuiper Belt objects, comets, and asteroids are key elements for understanding the early history of the solar system, its dynamics, and composition. While exoplanets are now routinely detected and hundreds of debris disks provide indirect evidence of planetesimals around main-sequence (MS) stars \citep{matthews14}, little is directly known about minor bodies around other stars than the Sun. The immense difficulty of a direct detection lies in their lack of a large surface area, which is required to detect their thermal or scattered emission. Dust features provide hints about the properties of $\mu$m-sized grains in debris disks that result from collisions of planetesimals \citep[e.g.][]{olofsson12}. Circumstellar (CS) CO emission around some AF-type MS stars \citep[e.g.,][]{moor15,marino16} has been interpreted to be the result of outgassing produced by comet collisions \citep{zuckerman12}. A complementary, somehow more direct, information on the exocomet nature and composition is provided by the detection of variable absorptions superposed on photospheric lines in the spectra of some stars. Variable absorption features in metallic lines have been known for about 30 years in the optical spectrum of $\beta$ Pic \citep{hobbs85}. These irregular features, mainly traced in the Ca~{\sc ii} K line, appear redshifted and to a much lower degree blueshifted with respect to the radial velocity of the star, and might vary on timescales as short as hours. These features have been interpreted to be the result of the gas released by the evaporation of exocomets grazing or falling onto the star \citep[and references therein]{ferlet87,kiefer15} that are driven into the vicinity of the star by the perturbing action of a larger body, that is, by a planet \citep{beust91}. Variable absorptions like this have also been observed toward several A-type stars \citep[e.g.,][]{redfield07,roberge08,welsh15}. We have initiated a short- and medium-term high-resolution spectroscopic project aiming at detecting and monitoring these sporadic events that are attributed to exocomets in a sample of MS stars, most of them A-type stars, but also some FG-type stars with a range of ages. The sample includes stars for which exocomet signatures have been detected previously, such as 49 Ceti or HR 10, and also stars that have not been scrutinized yet for such events. So far, we have obtained more than 1200 spectra of $\text{about }$100 stars (Rebollido et al., in preparation). The star $\phi$ Leo (HR 4368, HD 98058, HIP 55084) stands out as its spectrum exhibits conspicuous variations on timescales of hours, days, and months. This work presents our results concerning the Ca {\sc ii} K line obtained in five observing runs from December 2015 to May 2016. This line is particularly sensitive to these absorptions and is most frequently analyzed in the exocomet literature. Results for other relevant lines, including the Ca {\sc ii} IR-triplet, Ti {\sc ii}, Na {\sc i} D, and Balmer lines, together with the spectra of other stars, will be presented in a forthcoming paper. | Our intensive monitoring of $\phi$ Leo showed that its spectrum is very rich in redshifted absorption events, which might be accompanied in some cases by broad wings and even blueshifted absorptions. These sporadic events are similar to those in $\beta$ Pic and can be most plausibly explained as exocomets that graze the star or fall onto the stellar surface. Assuming this scenario, it is intriguing how a relatively old 500-900 Myr star such as $\phi$ Leo, which does not have any known associated debris disk, can possess such a rich environment that hosts minor bodies. Another interesting aspect is the origin of what might be a triangular-shaped stable CS absorption component in the Ca {\sc ii} lines. Additional monitoring is clearly needed to better characterize the sporadic events and the stable component by comparing them with similar stars. | 16 | 9 | 1609.04263 |
1609 | 1609.03505_arXiv.txt | Good explanations for the unusual light curve of Boyajian's Star have been hard to find. Recent results by Montet \& Simon lend strength and plausibility to the conclusion of Schaefer that in addition to short-term dimmings, the star also experiences large, secular decreases in brightness on decadal timescales. This, combined with a lack of long-wavelength excess in the star's spectral energy distribution, strongly constrains scenarios involving circumstellar material, including hypotheses invoking a spherical cloud of artifacts. We show that the timings of the deepest dimmings appear consistent with being randomly distributed, and that the star's reddening and narrow sodium absorption is consistent with the total, long-term dimming observed. Following Montet \& Simon's encouragement to generate alternative hypotheses, we attempt to circumscribe the space of possible explanations with a range of plausibilities, including: a cloud in the outer solar system, structure in the ISM, natural and artificial material orbiting Boyajian's Star, an intervening object with a large disk, and variations in Boyajian's Star itself. We find the ISM and intervening disk models more plausible than the other natural models. | \label{sec:intro} \subsection{Discovery} \citet{WTF} announced the discovery of an extraordinary star, \wtf, observed by {\it Kepler} \citep{Kepler} during its prime mission. First noticed by citizen scientists as part of the Planet \replaced{Hunter}{Hunters} project\footnote{\url{http://planethunters.org}} to examine {\it Kepler} light curves by eye \citep{PlanetHunters1}, this star exhibited a series of aperiodic dimming events, with variable timescales on the order of days, amplitudes up to 22\%, and a complex variety of shapes. \citeauthor{WTF} established that the data are good and that this behavior is unique to \wtf\ among {\it Kepler} stars. Extensive follow-up by \citeauthor{WTF} allowed them to determine the star appears to be in all other ways an ordinary, \added{main-sequence} early F star, showing no signs of IR excess (that would be indicative of a disk or other close-in material responsible for absorption) or accretion. Indeed, the {\it Kepler} field was above the Galactic Plane, and contains no known star-forming regions that might produce a star young enough to have significant circumstellar material. Using AO, \citet{WTF} did discover a 2\arcsec\ companion consistent with a bound M4V star at a projected distance of $\sim 900$ au, but this is not unusual, nor does it seem to provide any explanatory power for the star's {\it Kepler} light curve. \citeauthor{WTF} constructed ``scenario-independent constraints'' for the source of occulting material under the assumption that it is circumstellar, based on the duration and depths of the events, their gradients, the lack of IR excess, and other considerations. This allowed them to rule out many scenarios, and they offered a provisional explanation that {\it Kepler} had witnessed the passage of a swarm of giant comets. The hypothetical comets, which must be very large to block an appreciable amount of stellar flux, would only have produced a significant infrared excess during their periastron passage (presumably around the time of the {\it Kepler} mission), thus explaining the lack of IR excess at other times. \citet{Bodman16} modeled this scenario, and found it would required hundreds to thousands of comets, perhaps tidally disrupted from a Ceres-massed progenitor, to explain the final 60 days of the {\it Kepler} light curve. Despite this success, that also found that they could not reproduce the long, slow, deep event observed during {\it Kepler} Quarter 8, casting doubt on the comet hypothesis. Interest in \wtf\ (which we will refer to as ``Boyajian's Star''\footnote{The star has picked up other popular monickers, including ``WTF'' (ostensibly for ``Where's the Flux?,'' the subtitle to the \citet{WTF} paper), and ``Tabby's Star'' (with Dr.\ Tabetha Boyajian being its namesake). We agree that a more memorable name than \wtf\ seems warranted for this extraordinary object, and choose ``Boyajian's Star'' in keeping with the long astronomical tradition of similar eponyms, such as Barnard's Star, Kapteyn's Star, Teegarten's Star, etc.}) increased significantly in response to popular media accounts of the work of \citet{GHAT4}, who connected it to the speculation of \citet{Arnold05} that {\it Kepler} could discover large artificial structures orbiting other stars, if they exist. That is, rather than a swarm of comets, \citeauthor{GHAT4} noted that a swarm of planet- or star-sized structures would produce numerous transit anomalies, including arbitrary ingress and egress shapes, anomalous transit bottom shapes, variable depths, and aperiodicity---all of which characterize Boyajian's Star's dimming events. \citeauthor{GHAT4} recommended that, until Boyajian's Star's light curve had a more plausible natural explanation than had been offered to date, SETI\footnote{The search for extraterrestrial intelligence, e.g. \citet{Tarter01}} researchers prioritize it in their searches for communication from extraterrestrial civilizations. \subsection{Follow-up} \citet{Marengo15} and \citet{Lisse15} analyzed {\it Spitzer} and {\it IRTF} observations, respectively, taken after the {\it Kepler} observations, and showed that the lack of IR excess noted by \citet{WTF} (based on {\it WISE} data \citep{WISE} taken before the {\it Kepler} observations) continued to later epochs. This ruled out many scenarios involving a cataclysmic, dust-generating event in a planetary system that occurred between the {\it WISE} and {\it Kepler} epochs. \added{\citeauthor{WTF} showed that the gradients of the dips were consistent with material on a circular orbit at $\sim10$ au, where in equilibrium it would be quite cool and would escape detection at these wavelengths.} \citet{Thompson16} found no significant millimeter or submillimeter emission \replaced{putting}{and put} an upper limit of 7.7$M_\earth$ on the total circumstellar dust mass within 200 au. \added{This upper limit rules out a very massive debris disks orbiting \wtf. \citeauthor{Thompson16} note that only $10^{-9} M_\earth$ of dust is required to explain one of the deepest dimming events seen by {\it Kepler}, and give an upper limit of $\sim 5 \times 10^{-3} M_\earth$ for dust on elliptical orbits at the distances favored by the cometary hypothesis.} On the SETI front, \citet{Abeysekara16} found no evidence of optical flashes using the VERITAS gamma-ray observatory. \citet{Harp16} and \citet{Schuetz16} found no evidence of narrowband radio communication or pulsed laser emission during a simultaneous viewing campaign with the Allen Telescope Array and the Boquete Optical SETI Observatory\added{, respectively}. Only one part of the analysis of \citet{WTF} showing Boyajian's Star to be an otherwise ordinary F3V star has been called into question: \citet{Schaefer16} used archival DASCH photographic plate photometry \citep{Grindlay12,Tang13} to recover 100 years of brightness measurements for the field from 1890 to 1989. \citeauthor{Schaefer16}'s thorough analysis showed that Boyajian's Star ``faded at an average rate of $0.164 \pm0.013$ magnitudes per century,'' which he claimed ``is unprecedented for any F-type main sequence star'' and ``provides the first confirmation that \wtf\ has anything unusual'' beyond the {\it Kepler} dips. This claim, at least as extraordinary as the {\it Kepler} light curve itself, prompted multiple groups to attempt to confirm or refute it. \citet{Hippke16} and \citet{Lund16} found that the systematic errors in the DASCH photometry \replaced{to}{do} not permit measurements at the accuracy claimed by \citeauthor{Schaefer16}. In addition, \citet{Lund16} found several {\it other} F stars that they claimed {\it do} show such long-term variations in brightness. \citet{Montet16} performed an analysis of the full-frame images from the {\it Kepler} mission to determine if the secular dimming continued into the 21\textsuperscript{st} century. They found that Boyajian's Star indeed shows irregular, monotonic fading at an average rate of $\sim 0.7$ mag per century (four times the \citeauthor{Schaefer16} average) and was $\sim$ 4\% dimmer at the end of the mission than the beginning. They also show that many of the F stars that \citeauthor{Lund16} found to have secular photometric trends are revealed by {\it Kepler} to in fact be shorter-term variables and that the secular dimming of Boyajian's Star in the {\it Kepler} data is unique among the $> 200$ stars they studied. We agree with \citeauthor{Montet16}'s suggestion that the independent detection of an extraordinary, secular dimming of Boyajian's Star in the {\it Kepler} data makes \citeauthor{Schaefer16}'s result from the DASCH photometry more plausible, and with their assessment that such dimming finds little or no explanation from the comet hypothesis. We are thus left with no good explanation for the dimmings of Boyajian's Star --- neither the complex, short-period events during the {\it Kepler} epoch, nor the similar amplitude, secular trends seen in both the DASCH photometry and the {\it Kepler} full-frame photometry. \subsection{Plan and Purpose of This Letter} \citeauthor{Montet16} conclude their paper by stating that they ``strongly encourage further refinements, alternative hypotheses, and new data in order to explain the full suite of observations of this very mysterious object.'' This Letter's purpose is to be responsive to their encouragement. In Section \ref{Periodicity} we examine the periodicities of the deepest dips, and in Section \ref{Constraints} we examine the constraints on solutions imposed by the spectral energy distribution (SED) of Boyajian's Star. In Sections \ref{Possibilities}--\ref{Intrinsic} we discuss several families of solutions with various degrees of plausibilities in light of these constraints. \deleted{Because our purpose is to propose general directions for future research, we decline to develop these families into sets of detailed models.} In this work, we shall use the term ``dimming'' in a generic sense, to refer to the observed changes in the brightness of Boyajian's Star on all timescales. We adopt the term ``dips'' from \citeauthor{WTF} for the days-long events seen in the {\it Kepler} light curve, and ``secular'' or ``long-term'' dimming for the changes in brightness noted by \citeauthor{Schaefer16} and \citeauthor{Montet16}. In our discussion below, \replaced{will always}{we} assume that the long-term dimming identified by \citeauthor{Schaefer16} and \citeauthor{Montet16} is real. The combined $\sim V$-band dimming from the two works is 17\% $= 0.20$ mag, which, according to the reddening law of \citet{Fitzpatrick99} for interstellar dust, would imply the total extinction across the spectrum of Boyajian's Star of 15\%. | In response to \citeauthor{Montet16}'s strong encouragement to generate alternative hypotheses for the extraordinary light curve of Boyajian's Star, we have examined new and existing data and attempted to survey the landscape of potential solutions for plausibility. We have shown that the timings of the deepest dips exhibited by Boyajian's Star appear consistent with being randomly distributed in time, and so potential explanations should not be constrained by any perceived periodicities in the {\it Kepler} data. We argue that the star's secular dimming combined with a lack of long-wavelength excess in the star's SED strongly constrains scenarios involving circumstellar material. In particular, we find that no more than 0.2\% of Boyajian's Star's flux is being intercepted by the absorbing material, despite what appears to be at least a 15\% decrease in total flux toward Earth. We find that scenarios involving a spherical swarm of artificial structures absorbing the material are only just barely consistent with the data if they involve non-dissipative work done at the maximum (Carnot) efficiency at 65 K---other temperatures and lower efficiencies are ruled out, although other swarm geometries are not. We have briefly surveyed a range of explanations that do not invoke circumstellar material, and find two broad categories worthy of further consideration: an intervening Bok globule or other ISM overdensity, and an intervening stellar remnant with a large disk. We have shown that the star's color excess, absorption lines due to interstellar sodium, and predicted extinction due to interstellar dust are all consistent with $A_V\sim 0.34$, or about twice the total amount of secular dimming observed to date. Less compelling, but difficult to rule out, are intrinsic variations due to spots, a ``return to normal'' from a temporary brightening (due to, perhaps, a stellar merger) and a cloud of material in the outer solar system. We find instrumental effects, other intrinsic variation in Boyajian's Star, and obscuration by a disk around an orbital companion to Boyajian's Star very unlikely to be responsible. We have identified several additional lines of research that may help explain Boyajian's Star's light curve, including {\it JWST} MIR--FIR observations; optical broadband and spectroscopic observations during future dips; a study of the ISM toward Boyajian's Star; a hunt for similar variations in stars near on the sky to Boyajian's Star; and careful consideration of the Gaia parallax. | 16 | 9 | 1609.03505 |
1609 | 1609.06116_arXiv.txt | We present a physically consistent interpretation of the dc electrical properties of niobiumnitride (NbN)-based superconducting hot-electron bolometer (HEB-) mixers, using concepts of nonequilibrium superconductivity. Through this we clarify what physical information can be extracted from the resistive transition and the dc current-voltage characteristics, measured at suitably chosen temperatures, and relevant for device characterization and optimization. We point out that the intrinsic spatial variation of the electronic properties of disordered superconductors, such as NbN, leads to a variation from device to device. | 16 | 9 | 1609.06116 |
||
1609 | 1609.02924_arXiv.txt | The effects of near neighbors on the galaxy-galaxy lensing signal are investigated using a suite of Monte Carlo simulations. The redshifts, luminosities, and relative coordinates for the simulated lenses were obtained from a set of galaxies with known spectroscopic redshifts and known luminosities. As expected, when all lenses are assigned a single, fixed redshift, the mean tangential shear is identically equal to the excess surface mass density, scaled by the critical surface mass density: $\gamma_T = \Delta\Sigma \times \Sigma_c^{-1}$. When the lenses are assigned their observed redshifts and $\Sigma_c$ is taken to be the critical surface mass density of the central lens, the relationship $\gamma_T = \Delta\Sigma \times \Sigma_c^{-1}$ is violated because $\gtrsim 90$\% of the near neighbors are located at redshifts significantly different from the central lenses. For a given central lens, physically unrelated near neighbors give rise to a ratio of $\gamma_T$ to $\Delta\Sigma \times \Sigma_c^{-1}$ that spans a wide range of $\sim 0.5$ to $\sim 1.5$ at projected distances $r_p \sim 1$~Mpc. The magnitude and sense of the discrepancy between $\gamma_T$ and $\Delta\Sigma \times \Sigma_c^{-1}$ are functions of both $r_p$ and the velocity dispersions of the central lenses, $\sigma_v$. At large $r_p$, the difference between $\gamma_T$ and $\Delta\Sigma \times \Sigma_c^{-1}$ is, on average, much greater for low-$\sigma_v$ central lenses than it is for high-$\sigma_v$ central lenses. | Galaxy galaxy lensing, the systematic weak lensing of background galaxies by foreground galaxies, is powerful method by which the amount of dark matter and its relationship to luminous matter can be directly constrained. The first statistically-significant ($\gtrsim 4\sigma$) detections of galaxy-galaxy lensing to be published in the peer-reviewed literature demonstrated the viability of galaxy-galaxy lensing as a cosmological tool, but the sample sizes of foreground and background galaxies were too small to place particularly strong constraints on the properties of the dark matter halos of the foreground galaxies (e.g., Brainerd et al.\ 1996; Dell'Antonio \& Tyson 1996; Griffiths et al.\ 1996; Hudson et al.\ 1998). Later studies were able to take advantage of increasingly large samples, leading to improved constraints and demonstrating the importance of large surveys in the detection of the galaxy-galaxy lensing signal (e.g., Fischer et al.\ 2000; Guzik \& Seljak 2002; Hoekstra et al.\ 2004, 2005; Sheldon et al.\ 2004; Mandelbaum et al.\ 2005, 2006; Heymans et al\ 2006; Kleinheinrich et al.\ 2006; Tian et al.\ 2009). More recently, increasingly large data sets from the Sloan Digital Sky Survey (York et al.\ 2000), the Dark Energy Survey (Flauger 2005), the COSMOS survey (Koekemoer et al.\ 2007; Scoville et al.\ 2007ab), the Red Sequence Cluster Survey 2 (Gilbank et al.\ 2011), the Canada-France-Hawaii Telescope Lensing Survey (Heymans et al.\ 2012; Erben et al. 2013), and the combination of the Galaxy And Mass Assembly Survey (Driver et al.\ 2009, 2011; Liske et al.\ 2015) with the Kilo Degree Survey (Kuijken et al.\ 2015) have continued to dramatically improve the constraints that can be placed on the relationship between dark and luminous matter from galaxy-galaxy lensing. These most recent studies have placed strong constraints on the stellar-to-halo mass relation, the luminosity-to-halo mass relation, the dependence of the average halo mass on cosmic environment, and the masses of the halos surrounding passive versus star-forming galaxies (e.g., van Uitert et al.\ 2011, 2015, 2016; Leauthaud 2012ab; Tinker et al. 2013; Brimioulle et al.\ 2013; Velander et al.\ 2014; Coupon et al.\ 2015; Hudson et al.\ 2015; Zu \& Mandelbaum 2015; Brouwer et al.\ 2016; Clampitt et al.\ 2016; Mandelbaum et al.\ 2016). A key goal for galaxy-galaxy lensing studies is a measurement of the surface mass density surrounding central lens galaxies via observations of the mean tangential shear, $\gamma_T(r_p)$. In the case of isolated, axisymmetric lenses, it is well-known that the mean tangential shear is related to the excess surface mass density, $\Delta\Sigma(r_p)$, through \begin{equation} \Delta\Sigma(r_p) \equiv \left< \Sigma(< r_p) \right> - \Sigma(r_p) = \Sigma_c ~ \gamma_T(r_p) \label{eq:sigma_gamma} \end{equation} where $\left< \Sigma(< r_p) \right>$ is the mean interior surface mass density contained within a circle of radius $r_p$, centered on the lens, and $\Sigma(r_p)$ is the surface mass density at radius $r_p$ from the lens (e.g., Miralda-Escud\'e 1991). The quantity $\Sigma_c$ is known as the critical surface mass density and is given by \begin{equation} \Sigma_c = \frac{c^2}{4\pi G} \frac{D_s}{D_l D_{ls}} \label{eq:sigma_c} \end{equation} where $G$ is Newton's gravitational constant, $c$ is the velocity of light, $D_l$ is the angular diameter distance between the observer and the lens, $D_s$ is the angular diameter distance between the observer and the source, and $D_{ls}$ is the angular diameter distance between the lens and the source. Equation~\ref{eq:sigma_gamma} is also true in the case of an isolated non-axisymmetric lens so long as $\gamma_T(r_p)$ and $\Sigma(r_p)$ are interpreted as mean values, averaged over a ring of radius $r_p$, centered on the lens (see, e.g., Schneider 2006). In the case of multiple lens galaxies that all share the same redshift (as would be the case for, say, the halos of satellite galaxies that are contained within a larger system), Equation~\ref{eq:sigma_gamma} is also strictly true. It is this latter property of Equation~\ref{eq:sigma_gamma} that is commonly used to convert an observed value of the mean tangential shear into a constraint on the excess surface mass density in galaxy-galaxy lensing studies. Since the value of $\Sigma_c$ for a given lens-source pair depends upon the redshifts of both the lens and the source, the conversion of the shear into a constraint on the excess surface mass density requires a computation of the value of $\Sigma_c$ that is appropriate for the sample. This is sometimes done by computing a mean value of $\Sigma_c$ for the entire sample, and the mean is computed over the full lens-source distribution in redshift space. In that case, the right hand side of Equation~\ref{eq:sigma_gamma} is evaluated as the product of the mean tangential shear, computed over all lens-source pairs, and the mean value of $\Sigma_c$. More often, a value of $\Sigma_c$ is computed for each lens-source pair, then the products of the individual values of $\Sigma_c$ and $\gamma_T$ for each lens-source pair are averaged together to infer the mean excess surface mass density. Note that, in practice, various weighting schemes are typically used in observational studies in order to optimize the detection of the weak lensing shear and the reader is referred to recent papers on this subject for examples of these weighting schemes (see, e.g., Velander et al.\ 2014; Hudson et al.\ 2015; Brouwer et al.\ 2016). The published observational constraints on $\Delta\Sigma(r_p)$ have all acknowledged that, in addition to shear caused by the central lens galaxies around which $\gamma_T(r_p)$ is computed (the so-called ``one-halo'' term), there is an additional contribution to $\gamma_T(r_p)$ at large $r_p$ due to neighboring galaxies (the so-called ``two-halo'' term). For the most part, the two-halo term has been interpreted as being caused by neighboring galaxies that are physically associated with the central lens galaxies. That is, it is generally assumed that the two-halo term is dominated by neighboring galaxies that share the same redshift as the central lens galaxies (see, e.g., the review by Mandelbaum 2015). In the case that the two-halo term is, indeed, dominated by physically associated near neighbor galaxies, then Equation~\ref{eq:sigma_gamma} can be used to infer the excess surface mass density directly from the observed mean tangential shear, with $\Sigma_c$ being the value of the critical surface mass density that is appropriate for a given central lens, its physically associated near neighbors, and a given background source. It is, however, not necessarily the case that the two-halo term is dominated by physically associated near neighbors. The first simulations of galaxy-galaxy lensing showed that, for a given lens galaxy, the closest lens on the sky was not necessarily the most important lens and, typically, a given source was lensed at a comparable level by at least 3 or 4 physically unrelated foreground galaxies (Brainerd et al.\ 1996). Additionally, Brainerd (2010) showed that, for a realistic population of lens galaxies with median redshift $z_{\rm med} = 0.55 $, a population of source galaxies with median redshift $z_{\rm med} = 0.96$ would experience a significant amount of weak lensing due to multiple, physically unassociated galaxies. The goal of this paper is to examine the degree to which physically unassociated near neighbors affect the galaxy-galaxy lensing signal in the limit of a realistic, relatively deep sample of lens galaxies. Using a set of galaxies with known spectroscopic redshifts and known luminosities, a suite of Monte Carlo simulations of galaxy-galaxy lensing were constructed using a simple, analytic dark matter halo model. From the simulations, the relationship between the mean tangential shear and the actual surface mass density can be investigated directly, and the net effects of multiple, physically unassociated lens galaxies on the galaxy-galaxy lensing signal can be assessed straightforwardly. The paper is organized as follows. Details of the galaxy-galaxy lensing simulations are discussed in Section~2, and results for the mean tangential shear and the excess surface mass density, scaled by the critical surface mass density, are presented in Section~3. A summary and discussion of the results are presented in Section~4. Throughout, the present-day values of the cosmological parameters are taken to be $\Omega_0 = 0.25$, $\Omega_{\Lambda 0} = 0.75$, and $H_0 = 70$~km~sec$^{-1}$~Mpc$^{-1}$. | The observed celestial coordinates, spectroscopic redshifts, and rest-frame blue luminosities of galaxies in the region of the HDF-N were used as the basis of a suite of Monte Carlo simulations of weak galaxy-galaxy lensing. Since the simulations incorporate known galaxies with a uniform completeness in redshift space ($R \le 23$), the simulations naturally incorporate the intrinsic redshift distribution, mass distribution, and clustering of galaxies in our Universe. The simulations investigated the effects of near-neighbor galaxies on the mean tangential shear, computed around central lens galaxies, and the relationship between the mean tangential shear and the excess surface mass density, scaled by the critical surface mass density of the central lens. The main results from the simulations are as follows: \begin{enumerate} \item As expected, the relationship between the mean tangential shear and the scaled excess surface mass density is given by $\gamma(r_p) = \Delta\Sigma(r_p) \times \Sigma_c^{-1}$ when all lens galaxies are assigned a fixed, identical redshift. \item When the lens galaxies are assigned their observed spectroscopic redshifts, the value of $\gamma(r_p)$ may differ from that of $\Delta\Sigma(r_p) \times \Sigma_c^{-1}$. This is due to the fact that the majority of the near neighbors have redshifts that are significantly different from those of the central lenses. \item At large scales, the ratio of the mean tangential shear to the scaled excess surface mass density for a given central lens galaxy spans a wide range, from $\sim 0.5$ to $\sim 1.5$. \item The magnitude and sense (i.e., greater or less than unity) of the discrepancy between the mean tangential shear and the scaled excess surface mass density are functions of both the physical scale, $r_p$, and the velocity dispersions of the central lenses. In particular, at large scales the difference between the mean tangential shear and the scaled excess surface mass density is considerably greater for low-velocity dispersion central lenses than it is for high-velocity dispersion central lenses. \item The effects of physically unrelated near-neighbors on the weak lensing shear signal do not cancel out. Instead, physically unrelated near-neighbors increase the mean tangential shear on scales larger than the separation between the central lens and its neighbors. \end{enumerate} The significance of this work is a demonstration that, in the limit of a realistic redshift distribution for weak galaxy lenses, systematic errors should arise when the common method of converting an observed galaxy-galaxy lensing signal into a constraint on the excess surface mass density is used (i.e., multiplication of the value of $\gamma_T$ for a given source galaxy by the value of $\Sigma_c$ for the central lens and that particular source galaxy). This is because neighboring lens galaxies that are located at redshifts other than the redshift of the central lens give rise to a substantial component of the net weak lensing shear. An important goal for future, deep weak lensing surveys (for example, those that will be yielded by the Large Synoptic Survey Telescope, {\sl Euclid}, and {\sl WFIRST}) is a constraint on the mass distribution of the universe that is accurate to $\lesssim 1$\% on scales $r_p > 1$~Mpc. The results of this work suggest, therefore, that in order to achieve such an accurate measurement of the mass density from future galaxy-galaxy lensing studies, it may be important to move beyond the methods that are currently used to convert the observed galaxy-galaxy lensing signal into a measurement of the excess surface mass density. This is due to the fact that, to date, most theoretical predictions for the galaxy-galaxy lensing signal have adopted a model in which the two-halo term (i.e., the contribution to the weak lensing signal caused by galaxies other than the central lens) is attributed to neighboring galaxies that are physically related to the central galaxy (i.e., the neighboring lens galaxies that contribute to the net shear are assumed to be located at the same redshift as the central galaxy). Noteable exceptions to this include Schrabback et al.\ (2015) and Saghiha et al.\ (2016), both of which utilized direct ray-traying through N-body simulations to predict the galaxy-galaxy lensing signal using the methods described by Hilbert et al.\ (2009). The simulations presented here demonstrate that, in a sufficiently deep data set, the majority of near neighbors are not physically related to the central lens, and these physically unrelated near neighbors give rise to a significant component of the net galaxy-galaxy lensing signal on large scales. Based upon the results presented here, the need to explicitly include the effects of physically unrelated near neighbors may be especially important when attempting to constrain the dependence of $\Delta\Sigma(r_p)$ on the physical properties of the central lens galaxies (e.g., luminosity, color, stellar mass). Indeed, although the significance is somewhat low, there may already be an indication in existing galaxy-galaxy lensing observations that physically unassociated near neighbors are contributing to the observed signal. For example, Velander et al.\ (2014) found that their adopted model for the one- and two-halo terms provided good fits to the galaxy-galaxy lensing signal when the signal was computed using the entire population of central lens galaxies and also when it was computed separately for red central lenses. However, their model did not provide a good overall fit to the galaxy-galaxy lensing signal for blue central lenses with low luminosities and low stellar masses. The work presented here suggests that, if physically unassociated near neighbors are contributing to the galaxy-galaxy lensing signal, then the largest discrepancies between the model adopted by Velander et al.\ (2014) and the observed mean tangential shear should, in fact, manifest in the regime of the lowest-luminosity, lowest-mass central galaxies. The degree to which the observed mean tangential shear will differ from the actual scaled excess surface mass density in any large, future survey will, of course, depend upon the redshift distribution of the foreground lenses and the nature of the dark matter halos that surround them. For the sake of simplicity, here an unphysical isothermal sphere model was adopted for the halos of the lens galaxies. This was done in order to construct a demonstration of the effects of near neighbor galaxies on the galaxy-galaxy lensing signal that is straightforward to reproduce. Because of the simplicity of the adopted halo mass distribution, and because the effects of all possible foreground lens galaxies were not expressly included here (i.e., the simulations were limited to at most the contribution of the four nearest neighbors), the simulations presented here do not represent an accurate estimate of the degree of discrepancy between the mean tangential shear and the scaled excess surface mass density that would be expected to occur in our Universe. Rather, it is generally agreed that the mass density of our Universe is dominated by CDM, for which the dark matter halos of galaxies are not isothermal. The shear profile of a CDM halo differs from that of an isothermal sphere, being shallower than isothermal on scales less than the scale radius of the halo and steeper than isothermal on scales greater than the scale radius of the halo (see, e.g., Figure~1 of Wright \& Brainerd 2000). An accurate estimate of the degree of discrepancy between the mean tangential shear and the scaled excess surface mass density that could be expected in our Universe therefore awaits a more thorough analysis, such as could be obtained from high-resolution CDM simulations that include luminous galaxies via either semi-analytic galaxy formation or numerical hydrodynamics. | 16 | 9 | 1609.02924 |
1609 | 1609.02115_arXiv.txt | Blue stragglers (BSS) are stars whose position in the Color-Magnitude Diagram (CMD) places them above the main sequence turn-off (TO) point of a star cluster. Using data from the core of 47 Tuc in the ultraviolet (UV), we have identified various stellar populations in the CMD, and used their radial distributions to study the evolution and origin of BSS, and obtain a dynamical estimate of the mass of BSS systems. When we separate the BSS into two samples by their magnitude, we find that the bright BSS show a much more centrally concentrated radial distribution and thus higher mass estimate (over twice the TO mass for these BSS systems), suggesting an origin involving triple or multiple stellar systems. In contrast, the faint BSS are less concentrated, with a radial distribution similar to the main sequence (MS) binaries, pointing to the MS binaries as the likely progenitors of these BSS. Putting our data together with available photometric data in the visible and using MESA evolutionary models, we calculate the expected number of stars in each evolutionary stage for the normal evolution of stars and the number of stars coming from the evolution of BSS. The results indicate that BSS have a post-MS evolution comparable to that of a normal star of the same mass and a MS BSS lifetime of about 200-300 Myr. We also find that the excess population of asymptotic giant branch (AGB) stars in 47 Tuc is due to evolved BSS.\\ | \label{sec:intro} With the development of high resolution astronomical imaging, astronomers have been able to study globular star clusters (GC) in great detail, exposing the presence of different anomalous stellar populations. An important example of such stars are blue stragglers (BSS). First discovered by \cite{sandage} in the GC M3, BSS were described as an extension of the main sequence (MS) defying normal stellar evolution within a cluster. How these stars are formed in GC and where they go after they leave their MS stage has been a constant debate (\cite{eco_bss}, especially Chapter 9 and 11). One of the largest populations of BSS resides in the GC NGC 104 (47 Tucanae, 47 Tuc). Although 47 Tuc has been the target of many investigations, observations using ultraviolet (UV) filters, such as the one used to obtain the current data set, facilitate the selection of BSS and make them one of the brightest populations in the CMD. In the last two decades BSS have been found in many GCs \citep{ferraro2012} as well as open clusters \citep{deMarchi2006,ahumada}, in dwarf galaxies \citep{santana} and in the field of our galaxy \citep{santucci}. In older stellar clusters, there is generally no evidence of recent star formation episodes, thus, for these stars to look brighter and bluer than the turnoff they had to go through some rejuvenating process.The BSS formation mechanisms can be divided in many different ways but they all must comply with two main conditions: {\it i)} there must be at least one MS star involved, and {\it ii)} one of the stars involved must gain mass in order to become rejuvenated. In fact, the positions of BSS on the CMD suggest that these stars are in fact more massive than the TO stars. The first attempt to directly measure the mass of a BSS was done by \cite{shara}, studying one of the brightest BSS in the core of 47 Tuc. They found a mass of $1.7 \pm 0.4 M_{\odot}$, almost twice the cluster TO mass of $\sim 0.9M_{\odot}$ \citep{hesser,thompson}. Later, different studies, including some done on variable BSS, have yielded masses between $\sim 1$ and $2 M_{\odot}$ for BSS in different GCs \citep{gilliland,demarco}. Recent results for pulsating BSS have provided a lower upper limit of $\sim 1.5M_{\odot}$ \citep{fiorentino2014,fiorentino2015}. There are two possible ways for a star to gain mass: mass transfer or merger. We will separate the initial scenarios into three different categories following the divisions chosen by \cite{bookCh11}: {\it i)} direct collisions of stars \citep{hills}, {\it ii)} stellar evolution of primordial binaries \citep{McCrea}, and {\it iii)} dynamical evolution of hierarchical triple systems \citep{iben}. The last scenario became more important with the discovery of triple systems harbouring BSS (see \cite{s1082} for example), and the disagreement between the observed BSS populations and that obtained from combined N-body and stellar evolution simulations that considered only collisions and primordial binary evolution. \cite{perets2009} claimed that previous BSS formation studies demanded a fraction of the primordial binaries to be short period binaries. A previous publication by \cite{shortperiod} had shown that such systems actually come from longer period binaries that have been perturbed by a third star via the Lidov-Kozai mechanism \citep{lidov,kozai}. In fact, studies done on short period \citep{tokovinin} and contact \citep{pribulla} binaries showed that at least 40\% of these systems have distant companions. Recent studies following the formation channel proposed by \citeauthor{perets2009}, indicate that the Lidov-Kozai mechanism has a 21\% efficiency when it comes to forming tight binaries \citep{naoz2014}. And, when applied to GC systems, it can contribute up to 10\% of the total BSS population \citep{antonini2015}. This population should show some observational differences when comparing them to the BSS with different origins. For instance their mass could reach much higher values than binary mass transfer scenarios, where part of the mass of the system is left in the WD companion. The WD is also another difference as BSS from a triple system are more likely to be left with a MS companion \citep{bookCh11}. Which mechanism dominates in the different environments in which BSS live is still under debate. Although no definite answer has been reached most studies agree that the observed populations today are a result of a combination of all the formation channels, with one mechanism prevailing over the others depending on the system's properties. Attempts to find the dominating formation mechanism in different GCs, have been based on finding the strongest correlation between the number of BSS and parameters of the cluster, like total or core mass, binary fraction and collision rate. \cite{knigge2009} found a strong correlation between the number of BSS in the core and the core mass in GCs. With the results pointing towards a binary origin for BSS, researchers started to look for confirmation of the correlation between BSS frequency and binary fraction already found by \cite{sollima2008} in low density GCs. \cite{milone_binaries} reaffirmed this correlation for a sample of 59 GCs. \cite{leigh2013} also tried to find a relation between binaries and BSS but their results showed a much stronger correlation with the core mass as found by \cite{knigge2009}, despite the fact that binary fraction in GCs anticorrelates with core mass \citep{milone2008}. One of the latest studies that included dynamical effects and stellar and binary evolution yielded {\it ``a dependence of blue straggler number on cluster mass, a tighter correlation with core mass, a weak dependence on the collisional parameter, and a strong dependence on the number of binary stars"} \citep{sills2013}. Observational evidence supporting the fact that more than one mechanism for BSS formation can take place has also been found. \cite{ferraroM30}, \cite{dalessandro2013} and \cite{simunovic2014} found two sequences of BSS in M30, NGC 362 and NGC 1261 respectively, in both clusters a blue and a red sequence of BSS was visible on the CMD. Both authors suggest that this feature is a possible consequence of two very distinct formation mechanisms taking place in the same cluster. In the case of M30, the blue-BSS sequence matches collisional models \citep{sills2009}, while the red-BSS sequence agrees with binary evolution models \citep{xin2015}. \subsection{Blue Stragglers in 47 Tucanae} \label{subsec:intro_bss} In the particular case of 47 Tuc, the study of its population of BSS started with the discovery of 21 such stars in one of the first HST observations of the core of this cluster \citep{paresce1991}. This small sample of BSS already provided evidence that the density of BSS is higher in the central regions of the cluster. Many investigations on the topic have taken place since then, \cite{sills2000} modelled the formation rate of BSS using data outside the core. The results obtained by these authors suggested that 47 Tuc may have stopped making BSS several billion years ago. The cluster underwent an epoch of enhanced BSS formation around the same time, and this was possibly connected to the epoch of primordial binary burning. \cite{ferraro2004} discovered a bimodal radial distribution for the BSS in 47 Tuc, as seen in other GC like M3 \citep{ferraroM3}, and M55 \citep{lanzoni2007}. These distributions show a peak in the cluster center, decreasing at intermediate distances from the center, to rise again in the outskirts. \cite{mapelli2004} tried to reproduce the BSS radial distribution in 47 Tuc by choosing different formation mechanisms: collisional BSS in the innermost region and primordial binary evolution outside the core. The best representation of the observational data was obtained when 25\% of the BSS came from binaries and 75\% from collisions within $0.5r_{c}$ ($r_{c} = 1$ core radius). This result was later refined by \cite{mapelli2006} obtaining a best fit when 46\% of the BSS come from mass transfer and 54\% from collisions. The models were also able to predict the minimum in the radial distribution and its surrounding regions named by \cite{mapelli2004} as the {\it ``zone of avoidance"}, with the condition that external mass transfer BSS production began beyond $30r_{c}$. Later on, \cite{monkman2006}, tried to explain the bimodal distribution with a purely collisional model throughout the cluster. Their results agreed with those found by \cite{mapelli2004,mapelli2006} for the core of the cluster where the collisional model represents the observational data. For their middle region (between 23 and 130 arcseconds from the center) BSS formation would have needed to stop about half a billion years ago. But for the external regions the collisional models were not able to predict the BSS population, a result that they concluded is likely due to another formation mechanism dominating the outskirts of 47 Tuc. Around the same time the formation mechanisms debate was taking place, researchers found evidence that BSS in the core of 47 Tuc have masses larger than twice the MS TO mass. One result that suggested the presence of massive BSS was found by \cite{mclaughlin2006}, while studying the proper motion and dynamics of the cluster core. They determined that the velocity dispersion of BSS was smaller than that of the cluster giants by a factor of $\sqrt{2}$ (i.e. twice their mass). That same year, \cite{knigge2006} identified a detached binary system consisting of a $1.5M_{\odot}$ BSS primary with an active, upper MS companion. These massive BSS can only be the outcome of a process involving at least three progenitors. Another interesting area of research is the evolution of BSS. In 1994 \cite{bailyn}, studied the central regions of 47 Tuc and found an overabundance of stars in the AGB of the cluster. He concluded that these extra stars could come from the evolution of BSS. \cite{beccari2006} also found evidence of an overabundance of stars in the AGB, which, according to theoretical tracks used by the authors, correspond to massive stars currently undergoing their RGB phase. Additionally \cite{beccari2006} separated the HB into faint and bright HB stars finding that both the overabundace of stars in the AGB and the presence of a bright extension of the HB could be related to the evolution of binary systems. Recently \cite{ferraro2016} confirmed observationally the presence of a massive star in the region of the CMD slightly brighter than the bulk of the HB. \cite{ferraro2016} report a mass of $1.4M_{\odot}$, much higher than the turnoff mass of the cluster, strongly suggesting this star is the result of the evolution of a BSS. | \label{sec:conclusion} We have identified a large sample of over 200 BSS and evolved BSS in HST UV data of the core of 47 Tuc. Expanding our research using available data in the visible, we have studied the properties of this population including their masses, possible formation mechanisms, and their evolution. When we separate the bright and faint BSS we find that the bright BSS show a much more centrally concentrated radial distribution and higher mass estimates, properties that suggest an origin involving triple or multiple stellar systems. In contrast, the faint BSS are less concentrated, with a radial distribution similar to the MSBn pointing to this populations as their likely progenitors. Isolating a sample containing only evolved BSS had, until now, only been attempted on the HB. The evolved BSS selected on the UV CMD along with the MESA models and the agreement between the radial distributions of the BSS, evolved BSS, bright HB, and AGB bump, allowed us to construct the story of the evolution of BSS. The time scales and number of observed and expected stars agree nicely with the BSS having a post-MS evolution comparable to that of a normal star of the same mass. The disagreement between our estimated MS lifetime and those found by others indicate that a more detailed study of individual BSS properties is necessary to constrain these values. We have also been able to select clean samples in the different stellar evolutionary stages for the normal evolution of stars. Here we find that the cumulative radial distributions for the upper MS, RGB, faint HB and AGB, seem to all come from the same sample as expected for stars of the same mass. It is important to mention that in both the AGB and the AGB bump, we find stars from the evolution of normal stars as well as those coming from the evolution of BSS. But the number of stars and their radial distributions have allowed us to state the dominant population in each sample. Future studies using high quality spectra, will tell us more about the formation and evolution of BSS. Each formation mechanism leaves BSS with different chemical properties and possible companions, both of which could be identified and characterised through spectroscopy (\cite{ferraro2016} and references within). \clearpage | 16 | 9 | 1609.02115 |
1609 | 1609.05390_arXiv.txt | We test the parallaxes reported in the {\it Gaia\/} first data release using the sample of eclipsing binaries with accurate, empirical distances from \citet{Stassun:2016}. We find an average offset of $-$0.25$\pm$0.05~mas in the sense of the {\it Gaia\/} parallaxes being too small (i.e., the distances too long). The offset does not depend strongly on obvious parameters such as color or brightness. However, we find with high confidence that the offset may depend on ecliptic latitude: the mean offset is $-$0.38$\pm$0.06~mas in the ecliptic north and $-$0.05$\pm$0.09~mas in the ecliptic south. The ecliptic latitude dependence may also be represented by the linear relation, $\Delta\pi \approx -0.22(\pm0.05) -0.003(\pm0.001)\times\beta$~mas ($\beta$ in degrees). Finally, there is a possible dependence of the parallax offset on distance, with the offset becoming negligible for $\pi\lesssim 1$~mas; we discuss whether this could be caused by a systematic error in the eclipsing binary distance scale, and reject this interpretation as unlikely. | } The advent of trigonometric parallaxes for $\sim$10$^9$ stars from the {\it Gaia\/} mission promises to revolutionize many areas of stellar and Galactic astrophysics, including exoplanet science. For example, with eventual expected precision in the parallax of $\approx$20~$\mu$as for bright exoplanet host stars, it should be possible to determine the stellar and planet radii and masses directly and empirically with accuracies of 3--5\% \citep[see, e.g.,][]{Stassun:2016b}. Already, the {\it Gaia\/} first data release \citep[DR1;][]{Gaia:2016} provides parallaxes for $\sim$2 million {\it Tycho-2\/} stars (the TGAS stars) with a nominal precision of $\approx$0.3~mas and with a quoted systematic uncertainty at present of 0.3~mas. The results from DR1 are based on only 14 months of observations, and use external information in the form of earlier positions from the {\it Hipparcos\/} \citep{ESA:1997, vanLeeuwen:2007} and {\it Tycho-2\/} \citep{Hog:2000} catalogs to help remove degeneracies \citep[the {\it Tycho-Gaia\/} Astrometric Solution;][]{Michalik:2015}. Additionally, they rely on very provisional and as yet incomplete calibrations, and as a result the astrometric products including the parallaxes are still preliminary. Nonetheless, the new parallaxes represent such an improvement in both quality and quantity that they are certain to be used by the community for a wide range of astrophysical applications, at least until future {\it Gaia\/} releases supersede them. It is essential, therefore, to assess the on-sky delivered performance of these parallaxes from {\it Gaia\/} DR1, especially the presence of any unexpected biases. This is particularly important in light of the experience from {\it Hipparcos\/}, which suffered a significant bias in at least the case of the Pleiades cluster \citep[e.g.,][]{Pinsonneault:1998}. Such a check requires a set of benchmark stars whose parallaxes are determined independently and with an accuracy that is at least as good as that expected from {\it Gaia\/} DR1. \citet{Stassun:2016} assembled a sample of 158 eclipsing binary stars (EBs) whose radii and effective temperatures are known empirically and precisely, such that their bolometric luminosities are determined to high precision (via the Stefan-Boltzmann relation) and therefore independent of assumed distance. \citet{Stassun:2016} reported new, accurate measurements of the bolometric fluxes for these EBs which, together with the precisely known bolometric luminosities, yields a highly precise distance (or parallax). The precision of the parallaxes for this EB sample was predicted by \citet{Stassun:2016} to be $\approx$190~$\mu$as on average. This is a factor of $\sim$1.5 better than the median precision of 320~$\mu$as for {\it Gaia\/} DR1 \citep{Gaia:2016}. It is even somewhat superior to the expected {\it Gaia\/} DR1 precision floor of 240~$\mu$as. These EB parallaxes can therefore readily serve as distance benchmarks for the trigonometric parallaxes reported by {\it Gaia} DR1, and in particular can be used to assess the presence of any systematics. In this Letter, we report the results of testing the {\it Gaia\/} DR1 parallaxes against the \citet{Stassun:2016} EB benchmark sample. Section~\ref{sec:data} describes the EB and {\it Gaia\/} data used. Section~\ref{sec:results} presents the key result of a systematic offset in the {\it Gaia\/} parallaxes relative to the EB sample. Section~\ref{sec:disc} considers potential trends in the parallax offset with other parameters. Section~\ref{sec:summary} concludes with a summary of our conclusions. | } Here we present evidence of a small but systematic offset in the average zero-point of the parallax measurements recently released by the {\it Gaia\/} Mission of about $-0.25 \pm 0.05$~mas, in the sense that the {\it Gaia\/} values are too small. We also find evidence to suggest that the offset is a function of ecliptic latitude. The offset in the northern ecliptic hemisphere is $-$0.38$\pm$0.06~mas and $-$0.05$\pm$0.09~mas in the southern ecliptic hemisphere. Alternatively, the offset may also be represented as a linear function of the ecliptic latitude, $\beta$~($^\circ$), according to $\Delta\pi \approx -0.22(\pm0.05) -0.003(\pm0.001)\times\beta$~mas. To apply the correction, this (negative) offset must be {\it subtracted} from the reported {\it Gaia\/} DR1 parallaxes. At present we can only confirm that the offset is statistically valid for relatively large parallaxes, $\pi \gtrsim 1$~mas. The reference for this determination is a set of more than 100 independently inferred parallaxes from a benchmark sample of well-studied eclipsing binaries with a wide range of brightnesses and distributed over the entire sky. This paper presents evidence of a {\it difference} between the {\it Gaia\/} and EB parallaxes, which we have interpreted here as a systematic error in {\it Gaia\/} after discussing the alternative. In particular, we have considered the possibility of a systematic offset in the EB effective temperature scale as a possible, but unlikely alternative explanation. It is expected that future releases of the {\it Gaia\/} catalog will remove this small shift as the number of observations increases, calibrations are improved, and the astrometric solution transitions to a self-consistent global fit using only {\it Gaia\/} data, independent of external astrometric information. Indeed, these final trigonometric parallaxes may then be used to further refine the EB sample itself, such as improvements to the EB effective temperature scale. In the meantime, investigators using the parallax results from {\it Gaia\/} DR1 are encouraged to keep the systematic error reported here in mind. | 16 | 9 | 1609.05390 |
1609 | 1609.02579_arXiv.txt | We present fits to the broadband photometric spectral energy distributions (SEDs) of \nebs\ eclipsing binaries (EBs) in the {\it Tycho-2\/} catalog. These EBs were selected because they have highly precise stellar radii, effective temperatures, and in many cases metallicities previously determined in the literature, and thus have bolometric luminosities that are typically good to $\lesssim$\lbolprecper. In most cases the available broadband photometry spans a wavelength range 0.4--10~$\mu$m, and in many cases spans 0.15--22~$\mu$m. The resulting SED fits, which have only extinction as a free parameter, provide a virtually model-independent measure of the bolometric flux at Earth. The SED fits are satisfactory for \ngoodseds\ of the EBs, for which we achieve typical precisions in the bolometric flux of $\approx$\fbolprecper. Combined with the accurately known bolometric luminosity, the result for each EB is a predicted parallax that is typically precise to $\lesssim$\parprecper. These predicted parallaxes---with typical uncertainties of \parprecwunit---are 4--5 times more precise than those determined by {\it Hipparcos\/} for \nhip\ of the EBs in our sample, with which we find excellent agreement. There is no evidence among this sample for significant systematics in the {\it Hipparcos\/} parallaxes of the sort that notoriously afflicted the Pleiades measurement. The EBs are distributed over the entire sky, span more than 10 mag in brightness, reach distances of more than 5 kpc, and in many cases our predicted parallaxes should also be more precise than those expected from the {\it Gaia\/} first data release. The EBs studied here can thus serve as empirical, independent benchmarks for these upcoming fundamental parallax measurements. | } There is arguably no astronomical measurement more fundamental than distance from trigonometric parallax. Such parallax measurements are foundational to the cosmic distance scale generally and to stellar astrophysics specifically, including our basic understanding of stellar evolution, stellar populations, and Galactic structure. Thus, the parallaxes provided by the {\it Hipparcos\/} mission \citep{Perryman:1997, vanLeeuwen:2007} have had important, unrivaled impact across many areas of study for the past 20 years. Yet the {\it Hipparcos\/} parallaxes were not without problems, perhaps the most significant of which was the aberrant distance to the Pleiades which the {\it Hipparcos\/} parallax placed at $120.2\pm1.9$ pc \citep{vanLeeuwen:2009} as compared to the broadly accepted distance of $\approx$135 pc from a variety of methods \citep[e.g.,][and others cited therein]{Munari:2004, Zwahlen:2004, Soderblom:2005, Groenewegen:2007, Melis:2014, Madler:2016}. Several of these recent determinations are formally highly precise ($\sigma = 1.2$--1.7 pc). The discrepancy in the {\it Hipparcos\/} distance for the Pleiades was identified almost immediately \citep[e.g.,][]{Pinsonneault:1998, Soderblom:1998}, thanks to the unusual combination of proximity and cluster richness afforded by the Pleiades. In general, however, heretofore there have been very few {\it fundamental} benchmarks against which to test the {\it Hipparcos\/} parallaxes across the sky, across various stellar environments, and across stellar parameter space. One notable example is the use of spatially resolved double-lined spectroscopic binaries, which can yield highly accurate and precise orbital parallaxes \citep[see, e.g.,][]{Tomkin:2005}. Eclipsing binary (EB) stars have long served as fundamental benchmarks for stellar astrophysics. Through analysis of the light curve and radial velocities, EBs yield direct measurement of the component stellar masses and radii, and temperatures, with accuracies of $\sim$1\% for the best cases. As a result, the bolometric luminosity (\lbol) of an EB can be determined to very high precision and---importantly---without need of a distance measurement because it depends only on the measured radii and effective temperatures. EBs with such accurately determined \lbol\ can therefore serve as empirical, independent benchmarks for stellar distances obtained by other methods such as trigonometric parallax, if the bolometric flux at Earth (\fbol) can also be measured empirically and with sufficient precision. For example, in the recent analysis of newly discovered Pleiades EBs in the {\it K2\/} (successor to {\it Kepler\/}) mission data, \citet{David:2016} used the available broadband photometry from {\it GALEX\/} ultraviolet bands through {\it WISE\/} mid-infrared bands to determine an empirical distance of 132$\pm$5~pc to the Pleiades EB HCG~76, consistent with the consensus distance and again showing the {\it Hipparcos\/} distance to be biased. Impressively, this new empirical EB-based distance has a precision of better than 4\%, significantly more precise than the 12\% discrepancy in the {\it Hipparcos} distance, making it useful as a meaningful benchmark. Variants of this idea have been used to determine distances to EBs in the near field (relying on bolometric corrections) and even in external galaxies such as M31, the LMC or the SMC \citep[e.g.,][]{Ribas:2005, Pietrzynski:2009, Graczyk:2014}, sometimes based on application of surface brightness relations. These latter EB studies and others like them have been extremely important in establishing the lower rungs of the cosmological distance ladder, serving as calibrators for other methods reaching larger distances. The upcoming {\it Gaia\/} mission holds great promise for many areas of stellar and Galactic astrophysics through the provision of fundamental trigonometric parallax measurements for $\sim 10^9$ stars. At the same time, there now exists a large sample of benchmark-grade EBs with accurate radii and temperatures, and in many cases metallicities, which provide accurate and distance-independent \lbol. In addition, there now exist all-sky, broadband photometric measurements for stars spanning a very broad range of wavelengths, from the {\it GALEX\/} far-UV at $\sim$0.1~$\mu$m to the {\it WISE} mid-IR at $\sim$22~$\mu$m. These measurements permit construction of spectral energy distributions (SEDs) that effectively sample the majority of the flux for all but the hottest stars. Consequently the bolometric fluxes, and in turn the distances to the EBs, can in principle be determined in a largely empirical manner that preserves the accuracy of the fundamental EB parameters. In this paper, we use available broadband photometry to construct empirical SEDs and to calculate \fbol\ for \nebs\ EBs in the {\it Tycho-2\/} catalog whose fundamental physical properties have previously been established with accuracies of better than 3\% \citep[e.g.,][]{Torres:2010}. The wavelength coverage of the SEDs for most of the EBs in our sample is sufficiently large that the resulting \fbol\ are typically precise to \fbolprecper, leading to predicted parallaxes that are in most cases precise to better than \parprecper. For \nhip\ of the EBs in our study sample, previous {\it Hipparcos\/} parallaxes are available for direct comparison to the parallaxes predicted from the EBs. In Section~\ref{sec:data}, we present our study sample, the data that we use from the literature, and our SED fitting procedures. The main results of the work, including empirical \fbol\ and predicted parallaxes for the full sample, and a comparison to {\it Hipparcos\/} parallaxes where available, are presented in Section~\ref{sec:results}. In Section~\ref{sec:disc} we briefly discuss the upcoming applicability of this work to {\it Gaia}, which is expected to include all of these EBs as early as its first public data release. We summarize our conclusions in Section~\ref{sec:summary}. | } Eclipsing binaries with well-measured physical properties allow the determination of distances that are essentially model-independent, and can be both highly accurate and highly precise. They are therefore ideal as benchmarks for validating other methods of establishing distances, and can often be used out to many kiloparsecs without loss of precision. They have been employed to great advantage even in external galaxies such as the LMC, the SMC, and others. In this paper we have assembled an all-sky list of \nebs\ EBs contained in the {\it Tycho-2\/} catalog with high-quality determinations of the component radii and effective temperatures from the literature, and combined this information with constrained SED fits using existing photometric measurements over a wide range of wavelengths. The distances calculated from the accurate absolute stellar luminosities and bolometric fluxes lead to predicted parallaxes having typical precisions of \parprecwunit\ (\parprecper\ relative errors), which are 4--5 times better than the trigonometric parallaxes from {\it Hipparcos}. We find excellent overall agreement between our results and those from {\it Hipparcos\/}. To the extent that our EB parallaxes represent a test of the {\it Hipparcos\/} parallaxes, we find no obvious systematic deviations of the sort that appear to have affected the Pleiades, at least for this particular sample. In any case, the good agreement between our EB parallaxes and the {\it Hipparcos\/} parallaxes supports the accuracy of our measurements, and other tests suggest that our precision estimates are also realistic. The quality of our predicted parallaxes also compares very favorably with that expected for the {\it Gaia\/} parallaxes from the Tycho-Gaia Astrometric Solution \citep{Michalik:2015}, soon to be delivered as part of the mission's first Data Release. Our results will therefore serve as an important external check on the spacecraft's astrometric performance early on in the mission. While subsequent {\it Gaia\/} releases will feature steadily increasing parallax precision for larger numbers of stars as the time base lengthens, we anticipate that predicted parallaxes from EBs will continue to provide a valuable reference that is completely independent of astrometry and whose precision does not degrade with increasing distance or diminishing brightness. | 16 | 9 | 1609.02579 |
1609 | 1609.07489_arXiv.txt | We analyze 45 spectropolarimetric observations of the eclipsing, interacting binary star V356 Sgr, obtained over a period of $\sim$21 years, to characterize the geometry of the system's circumstellar material. After removing interstellar polarization from these data, we find the system exhibits a large intrinsic polarization signature arising from electron scattering. In addition, the lack of repeatable eclipses in the polarization phase curves indicates the presence of a substantial pool of scatterers not occulted by either star. We suggest that these scatterers form either a circumbinary disk coplanar with the gainer's accretion disk or an elongated structure perpendicular to the orbital plane of V356 Sgr, possibly formed by bipolar outflows. We also observe small-scale, cycle-to-cycle variations in the magnitude of intrinsic polarization at individual phases, which we interpret as evidence of variability in the amount of scattering material present within and around the system. This may indicate a mass transfer or mass loss rate that varies on the time-scale of the system's orbital period. Finally, we compare the basic polarimetric properties of V356 Sgr with those of the well studied $\beta$ Lyr system; the significant differences observed between the two systems suggests diversity in the basic circumstellar geometry of Roche-lobe overflow systems. | Single star evolution is characterized primarily by a star's initial mass as it reaches the zero-age main sequence. However, the evolutionary pathway of a star in a binary or multiple system is significantly harder to predict, because these stars can lose a large fraction of their mass through interactions with their companions. Companion-affected mass loss is particularly important for massive stars, which are predominantly found in binary systems; approximately 70\% of all O stars will have their evolution strongly affected by a companion \citep{sana2012}. One way that the stars in a massive binary system can interact is via Roche-lobe overflow \citep[RLOF;][]{sana2012}, during which a large fraction of the mass in a star's outer layers may be stripped and transferred to its companion \citep{Iben1991,Podsiadlowski1992,sana2012}. Despite the far-reaching effects of this process, it is not well understood. For example, the fractions of material transferred and lost from massive binaries undergoing RLOF are unknown \citep{deMink2007}. Because changes in mass loss rates as small as a factor of two can alter the evolutionary model tracks of massive stars enough to produce different types of supernovae than otherwise expected \citep[SNe;][]{Smartt2009}, quantifying mass loss in interacting binary systems is vital to understanding massive stellar evolution. Unfortunately, directly imaging RLOF systems to better understand their mass-loss geometries remains largely beyond our current technological abilities. While interferometric techniques have had some success at resolving the components of massive binary systems undergoing RLOF in recent years (e.g., \citealt{zhao2008}, who imaged the tidal distortion of the primary star in the $\beta$ Lyrae system), very few RLOF objects have been observed with this technique. Spectropolarimetry provides a powerful way to diagnose the distribution of circumstellar material without the need to resolve system components. This technique is particularly advantageous for eclipsing binary systems because their eclipses provide strong geometrical constraints that aid in the interpretation of polarimetric data. In the best-studied case, the eclipsing RLOF system $\beta$ Lyr, \cite{hoffman1998} used position angle rotations in emission lines and across the Balmer jump to infer the existence of bipolar outflows, independently confirming their interferometric detection by \cite{harmanec1996}. Similarly, \cite{lomax2012} used polarimetric light curves of the $\beta$ Lyr system to detect a hotspot on the edge of the disk, which reveals itself through a consistent offset between the timings of the secondary eclipse in total and polarized light. However, it remains unclear whether $\beta$ Lyr's circumstellar geometry is representative of most massive RLOF systems, or whether the system is unique. V356 Sagittarii (hereafter V356 Sgr) is a massive interacting binary system whose geometry is thought to be similar to $\beta$ Lyr's. The 3 $M_{\odot}$, A2 supergiant secondary star (hereafter ``donor'') has filled its Roche lobe and is currently transferring matter to the brighter, B3 V primary star \citep[hereafter ``gainer'';][]{wilson1978,peters2004,rensbergen2011}. The gainer, originally the less massive star, now has a mass of approximately 11 $M_{\odot}$ \citep{rensbergen2011}. A mass transfer stream between the two stars feeds an accretion disk around the gainer \citep{wilson1978}. Because the system's far ultraviolet \textbf{(UV)} emission lines exhibit no change in shape, velocity, or strength during eclipses, \citet{peters2004} suggested that V356 Sgr may also possess bipolar outflows or other material outside the orbital plane of the system. Evolutionary modeling by \cite{ziolkowski1985} suggested that the mass transfer in V356 Sgr is not conservative, and that some material has been lost from the system. In fact, \citet{hall1981} found that they needed a quadratic term in their ephemeris to account for the system's small change in period with time. However, many studies of the system have found a fairly stable 8.89-day period (e.g. Popper 1980). \citet{wilson1995} argued that mass transfer in the system must be intermittent to explain the lack of observed period change in some datasets. In fact, the quadratic term in \citet{hall1981}'s ephemeris produces only a 0.22 day (2.5\%) increase in period from its initial epoch to the date of the last measurement we present in this paper. \citet{polidan} first reported that V356 Sgr exhibited a polarization signal. These authors showed that the system is intrinsically polarized at UV wavelengths, although they were unable to establish if the system is polarimetrically variable because there was only one observation. Because the two stars in the system are hot, the primary polarization mechanism in V356 Sgr is most likely electron scattering in its circumstellar material. Electron scattering is a wavelength-independent process that preserves information about the geometry of the scattering medium within the polarization signal of the scattered light. Therefore, analyzing the polarimetric behavior of V356 Sgr allows us to investigate the geometric properties of its circumstellar material without the need for interferometric techniques to resolve the system. In addition, the eclipsing nature of V356 Sgr provides constraints on the interpretation of the spectropolarimetric data, offering a unique opportunity to investigate the three-dimensional structure of its circumstellar material. In this paper, we present the results of a broadband polarimetric study of V356 Sgr and compare its distribution of circumstellar material to that found in $\beta$ Lyr. In Section 2 we discuss the details of our observations. We present the results of our polarimetric data in Section 3, and interpret those results in the context of the geometry of the V356 Sgr system in Section 4. Finally, we summarize our findings in Section 5. | } Analyzing the phase variations of V356 Sgr's total brightness, intrinsic polarization magnitude, polarized flux, and position angle enables us to constrain the general geometry of the circumstellar material in the system. The lack of any significant change in the intrinsic polarization across primary eclipse indicates that when the donor passes in front of the gainer and its accretion disk, it does not substantially affect the amount of scattered light seen by the observer. This implies that the accretion disk around the gainer is not a significant scattering region, which suggests that there may be an additional region responsible for scattering photons within the system. The similarity between the morphology of V356 Sgr's polarized flux and total flux phase curves implies the same properties. The minima at primary and secondary eclipse in the polarized flux appear to be solely due to a decrease in the amount of total flux observed and not to any decrease in the percent polarization. Additionally, the position angle of V356 Sgr's polarization is roughly constant with phase. Taken together, these results suggest two possible scenarios for the system's scattering region(s); most of the visible light is scattered either by material out of the orbital plane of the binary or by circumbinary material in the orbital plane. Regardless of the scattering scenario present in V356 Sgr, the data appear consistent with electron scattering; the intrinsic polarization is constant with wavelength. Therefore, no matter where the scattering region(s) in V356 Sgr are, they must be close enough to the stars to be ionized and produce a sea of electrons with which photons in the system can scatter. There is no indication of the presence of scattering with dust, which would have revealed itself through a wavelength dependent polarization signature. In Be stars whose disk is denser than about $5\times 10^{ -12}$ g cm$^{ -3}$, the \ion{H}{1} opacity introduces a significant negative slope in the polarized spectrum (e.g., Fig. 1 of \citealt{Hau}). Therefore, the fact that the polarization of V356 Sgr is constant with wavelength implies that the amount of \ion{H}{1} opacity in the scattering region is not significant, and allows us to estimate an upper limit of roughly $5\times 10^{-12}$ g cm$^{-3}$ for the density of the scattering material. \subsection{V356 Sgr's Accretion Disk} We first consider the simplest possibility: that all the visible polarization we observe is produced within the known accretion disk surrounding the gainer. The fact that the intrinsic polarization does not vary substantially across primary eclipse suggests that this eclipse is an occultation of primarily unpolarized light. This in turn implies that the accretion disk surrounding the gainer is optically thin. In this scenario, light arising from the gainer scatters into our line of sight at the outer edges of the accretion disk, acquiring a small polarization magnitude and a position angle oriented perpendicular to the orbital plane; we can thus interpret the near-constant PA of the system as a measure of the orientation of the orbital plane on the sky. If the disk is symmetric, we can assume each of these scattering edges produces the same amount of polarization and the same PA. Thus, when the secondary star occults these edges in turn, the overall PA does not change. This is consistent with the PA phase curves in Figure 5. However, we would expect to see two sharp, repeatable dips in polarization, symmetric around the eclipse, as each scattering edge is occulted by the secondary star. This effect should also appear in the polarized flux curves, where it would be even more dramatic, because occultation of each edge would correspond to a loss of half the polarized flux out of eclipse. Neither effect appears in our data (Figure \ref{fig:pflux}). The fact that our phase curves are sparsely sampled in time may limit our ability to detect such a signal, and we intend to obtain more data at a higher sampling rate in the future to search for it. However, for purposes of this paper, we assume hereafter that the accretion disk in V356 Sgr is not the only source of polarization, and investigate two other possibilities to explain the polarization behavior we observe. Both may exist in combination with a small contribution from scattering in the accretion disk. In Section 3.3, we noted that V356 Sgr exhibits stochastic variability in its intrinsic polarization at individual phases in all bandpasses (Figures \ref{fig:phasepolsub} and \ref{fig:sigma}). Following \citet{wilson1995}, we posit that these variations may be due to small-scale changes in the mass transfer rate between the donor and gainer, which lead to small inhomogeneities in the density of scatterers in the gainer's accretion disk. This may change the amount of material in the accretion disk on small enough time scales to create the observed stochastic variability. For example, a large amount of material could be quickly injected into the accretion disk from the loser and then quickly accreted onto the gainer. This scenario likely holds even if the accretion disk is not the main source of polarization in the system, as discussed in the sections below. \subsection{Extraplanar Material} Light scattering in material outside of the orbital plane of the stars would explain the lack of observed eclipses in the polarization. If the extraplanar material is persistent and maintains a consistent geometry, it will imprint a PA on the observed polarization that is different from that produced by light scattering in the accretion disk. In this case, the observed near-constant PA does not reflect the orientation of the orbital plane on the sky. If scattering in the extraplanar material dominates the polarization, the observed PA traces the orientation of this material. A more likely scenario is that the observed PA represents an intermediate value produced by the combination of polarization from light scattering in the accretion disk and in the extraplanar material. If the extraplanar material is not persistent (i.e., exists sporadically or not in every orbit), then the average PA we observe may trace the orientation of the orbital plane, while the stochastic rotations away from this mean reflect transient concentrations of extraplanar material, perhaps related to variations in the accretion rate as discussed above. \cite{peters2004} found that the behavior of the system's UV emission lines across primary eclipse suggests that they form from material not near V356 Sgr's stars or their orbit, e.g., within bipolar outflows. This lends credence to the idea that there could be a region of scatterers outside of the plane of the orbit of the system. However, there is very little information about these bipolar outflows available in the literature; their extent above and below the plane of the disk is currently unknown and no other data set has confirmed their existence. If extraplanar material is the correct interpretation of V356 Sgr's polarimetric signature, it is worth noting that this material is quite different from the bipolar outflows in $\beta$ Lyr \citep{harmanec1996,hoffman1998}. In $\beta$ Lyr, bipolar outflows are responsible for scattering the optical and UV emission lines, as well as the UV continuum, while the disk scatters optical continuum light. The distinction between the two scattering regions manifests itself as a near $90\degree$ rotation in the position angle of the polarization across the lines and the Balmer jump. However, we do not observe this behavior in V356 Sgr. We detected no change in polarization across the optical lines in V356 Sgr, and we also found that the position angle of the UV polarimetric continuum is consistent (within $10\degree$) with that of its optical polarimetric continuum (Section 3.6). In this scenario, therefore, bipolar outflows or other extraplanar material would have to be responsible for the majority of the continuum polarization observed from V356 Sgr, without any significant contribution from V356 Sgr's disk. \subsection{Circumbinary \textbf{or Intrabinary} Material} Another possible interpretation of the observed polarimetric behavior of V356 Sgr is that it is due to scattering in circumbinary material in the system's orbital plane. Because this material is exterior to the orbit of the stars, very little of it is eclipsed at any given time. This explains the lack of observed eclipses and also the constancy of the observed position angle, which in this case would again trace the orientation of the orbital plane on the sky. Circumstellar material between the stars but aligned with the orbital plane may also produce the polarization signatures we observe, as long as parts of the scattering region remain visible outside of both eclipses. One possibility for intrabinary material is the mass stream connecting the loser to the accretion disk; this is likely an elongated structure that lies in the orbital plane and is visible at most phases \citep{wilson1978}. Variations in the mass loss rate could cause the stream to become more or less dense over time and explain the stochastic changes in visible polarization we observe (Figure \ref{fig:pflux}). Another possible coplanar scattering region is the distorted atmosphere of the Roche-lobe-filling loser, a scenario proposed for $\beta$ Lyr by \citet{hoffman1998}, but this should show "eclipses" as the loser rotates, so by itself this cannot account for all the polarization we observe. \subsection{Comparison to $\beta$ Lyrae} Morphologically, $\beta$ Lyr and V356 Sgr are similar systems. They both are undergoing RLOF, which has caused disks to form around their gainers. Naturally, their similar geometry also causes both systems to have similarly shaped light curves. Polarimetrically speaking, the position angle of both systems is constant with phase and the magnitude of their polarization is similar (on the order of tenths of a percent). In a broader sense, we note that the observed polarimetric behavior of V356 Sgr is substantially different than that of the canonical eclipsing RLOF system $\beta$ Lyr \citep{hoffman1998,lomax2012}. Whereas $\beta$ Lyr possesses clear intrinsic polarization changes across both primary and secondary eclipses, intra-eclipse polarimetric variations due to the orientations of the loser and disk, an optically thick accretion disk, and phase-dependent polarization in its emission lines, V356 Sgr exhibits none of these. We conclude that although both systems contain accretion disks and bipolar flows, the properties of their circumstellar material are quite different. First, as discussed above in this section, our results suggest that the accretion disk in V356 Sgr is not the dominant scattering region. Because of this, we are not able to detect a hotspot at the point where the mass stream impacts the accretion disk in V356 Sgr as was done in the case of $\beta$ Lyr \citep{lomax2012}. Thus, if V356 Sgr contains a hotspot, it must be detected through different means. In addition, while V356 Sgr may possess bipolar outflows \citep{peters2004}, they may not be evident in our spectropolarimetric data. In the case of $\beta$ Lyr, electron scattering in the outflows causes a PA rotation between the near-UV and visible continua and between the optical emission lines and continuum \citep{hoffman1998}. This is caused by UV photons being absorbed in the optically thick accretion disk in $\beta$ Lyr; the only way for UV photons to escape the system is to travel perpendicular to the accretion disk and scatter within $\beta$ Lyr's bipolar outflows. However, the accretion disk in the V356 Sgr system is optically thin (Section 4.1), while $\beta$ Lyr's disk is optically thick. Therefore, UV photons in V356 Sgr are able to travel to any scattering region in the system that optical photons also reach. Taken together, these results imply that the geometrical distribution of circumstellar material in RLOF systems exhibits more variety than previously understood. A comprehensive study of a larger number of such systems is needed to draw broad conclusions about the typical mode of mass transfer between massive stars in these systems. \subsection{Limitations of the Data} We searched our data set for trends, variations, and periodicities other than the orbital period of the system in an effort to better understand what may be causing the scatter seen in Figures \ref{fig:phasepolsub} and \ref{fig:pflux} and interpret the general behavior of our data. For example, we searched the data for long-term trends by binning the data on yearly time-scales, but found no obvious trends when comparing different years to each other. Similarly, we plotted the polarization and position angle versus the Julian date and found no trends spanning the time frame over which our data was taken; i.e. the measured polarization and position angle of V356 Sgr do not appear to be increasing or decreasing over time. For completeness we also used \texttt{PERIOD04} to search for important frequencies in addition to the orbital period of the system, but found nothing significant. Our efforts to characterize and interpret the orbital polarimetric variations of V356 Sgr and the scatter in our data set are hampered by a couple of factors. First, the data are not well sampled in time. V356 Sgr was observed in only seven of the 22 years that our data set spans. Most notably, there is a 13 year span, between 1999 and 2011, when the system was not monitored at all. Additionally, there is a wide range in the number of observations obtained per year during the seven years in which the system was monitored. The fewest observations, 2, occurred during 1994, while the most, 12, were obtained in 2012. Second, our data set does not adequately sample any single orbit of V356 Sgr. The combination of these two limitations make it difficult to assess the source of the scatter in our polarimetric observations. Using our existing data, we cannot rule out the possibility that the polarization of V356 Sgr is well-behaved over the course of one epoch, but exhibits large variations between cycles. Alternatively, it is possible that the system does not display repeatable periodic behavior on any timescale. Future polarimetric observing campaigns could better investigate the cause of observed variations and scatter by intensely focusing on V356 Sgr over the course of one or two epochs (e.g. with observations at least once or twice per night) and by continuing to monitor the system over several more cycles at a lower frequency (e.g. one observation every 2-3 nights for several weeks). Additionally, there is currently very little published information about the UV spectrum and polarization of V356 Sgr. UV spectra taken during the system's primary eclipse have been published \citep{peters2004}, but there are no data that cover the rest of V356 Sgr's orbital period. There also exists two UV polarimetric observations in the HST archive that have yet to appear in the literature outside of conference abstracts (e.g. \citealt{polidan}). While beyond the scope of this paper, we plan to analyze and compare these data to our WUPPE observation in the near future in order to better understand the UV polarimetric behavior of the system. Finally, our ISP estimate is not rigorous. As previously stated, changes in our estimate of the ISP's position angle by $10\degree$ did not significantly affect our results. Perhaps even more problematic is that the field stars we used are quite far apart from both each other and V356 Sgr. Because of this, we may be probing different ISM clouds with different field stars which may in turn lead to a poor understanding of how the ISM is affecting the observed V356 Sgr polarization. Future polarimetric work on this system might benefit from identifying an ISP probe star that is spatially close to V356 Sgr with a small angular separation. However this may prove to be difficult; V356 Sgr does not lie within a cluster of stars. | 16 | 9 | 1609.07489 |
1609 | 1609.07440_arXiv.txt | {We spectroscopically re-observed the gravitational lens system SDSS J1339+1310 using OSIRIS on the GTC. We also monitored the $r$-band variability of the two quasar images (A and B) with the LT over 143 epochs in the period 2009$-$2016. These new data in both the wavelength and time domains have confirmed that the system is an unusual microlensing factory. The C\,{\sc iv} emission line is remarkably microlensed, since the microlensing magnification of B relative to that for A, $\mu_{\rm{BA}}$, reaches a value of 1.4 ($\sim$ 0.4 mag) for its core. Moreover, the B image shows a red wing enhancement of C\,{\sc iv} flux (relative to A), and $\mu_{\rm{BA}}$ = 2 (0.75 mag) for the C\,{\sc iv} broad-line emission. Regarding the nuclear continuum, we find a chromatic behaviour of $\mu_{\rm{BA}}$, which roughly varies from $\sim$ 5 (1.75 mag) at 7000 \AA\ to $\sim$ 6 (1.95 mag) at 4000 \AA. We also detect significant microlensing variability in the $r$ band, and this includes a number of microlensing events on timescales of 50$-$100 d. Fortunately, the presence of an intrinsic 0.7 mag dip in the light curves of A and B, permitted us to measure the time delay between both quasar images. This delay is $\Delta t_{\rm{AB}}$ = 47$^{+5}_{-6}$ d (1$\sigma$ confidence interval; A is leading), in good agreement with predictions of lens models.} | \label{sec:intro} If there is a massive galaxy between a distant quasar and the Earth, the quasar is seen as a multiple system consisting of several images \citep[e.g.][]{schneider92}. This gravitationally lensed quasar may also suffer microlensing effects by stellar mass objects in the lensing galaxy. Quasar microlensing was firstly detected in 1989 \citep{irwin89,vander89}, and has become a powerful astrophysical tool in the current century \citep[e.g.][]{schneider06}. Microlensing particularly affects compact sources such as the X-ray emitting regions, the accretion disk or the innermost line emitting clouds, so microlensed quasars are being intensively used to probe the quasar structure from spectral studies \citep[e.g.][]{chartas04,richards04,pooley07,sluse07,bate08,floyd09, blackburne11,mediavilla11,mosquera11,munoz11,jimenez12,motta12,sluse12,guerras13,braibant14, jimenez14,rojas14,sluse15} and analyses of extrinsic variabilities \citep[e.g.][]{chartas02, shalyapin02,goico03,kochanek04,gil06,morgan06,paraficz06,eigenbrod08,morgan08a,morgan08b, poindexter08,chartas09,dai10,morgan10,poindexter10,chen11,sluse11,chartas12,chen12,morgan12, hainline13,mosquera13,blackburne14,blackburne15,macleod15,mediavilla15a,mediavilla15b}. Unfortunately, many lensed quasars have only experienced weak microlensing effects in their spectra and light curves. Thus, it is important to identify the systems showing substantial spectral distortions due to microlensing, sharp microlensing variations or both phenomena. These microlensing factories are excellent tools for detailed interpretations, as well as the best targets for subsequent follow-up through suitable facilities \citep[e.g.][]{moskoc11}. For example, some systems exhibited strong chromatic microlensing in the optical continuum, with a maximum signal of $\sim$ 0.8 mag at 5439 \AA\ \citep[HE 0047$-$1756 and SDSS J1155+6346;][]{rojas14}. A microlensing-induced distortion of the shape of several high-ionization emission lines was also unambiguously detected in SDSS J1004+4112 \citep[e.g.][]{richards04,motta12}. Additionally, the Einstein Cross (QSO 2237+0305) displayed prominent microlensing (extrinsic) events in its $V$-band light curves over the final years of the past century \citep{wozniak00,udalski06}, and these $\sim$ 0.7 mag variations on hundreds of days strongly stimulated interpretation tasks and further observations (see above). Lensed quasars are not only powerful laboratories of the quasar structure, but are also often used as a cosmological probe \citep[e.g.][]{refsdal64,schneider06}. For that purpose, the time delay between pairs of lensed images must be measured to high accuracy, requiring us to disentangle extrinsic and intrinsic variability. Intrinsic variations in a lensed quasar appear in two given images at different observing times, and a measurement of the time delay between these lensed images can be used to estimate the current expansion rate of the Universe (the so-called Hubble constant) and other cosmological parameters \citep[e.g.][]{oguri07,suyu10,suyu13,sereno14,rathna15}. The time delay is also sensitive to the distribution of the lensing mass \citep[e.g.][]{refsdal64, schneider06,goico10}, so time delay measurements are turning in critical data for cosmology and extragalactic astrophysics. However, lensed quasars with appreciable extrinsic variations in their optical light curves may be a challenge for time delay determinations, since, in general, the extrinsic variability should be modelled by microlensing simulations or appropriate functions \citep[e.g.][and references therein]{tewes13a}. Various techniques have been developed to deblend the intrinsic signal from the microlensing variability. However, as seems intuitive, if the intrinsic and extrinsic variations have similar timescales and amplitudes, it is not possible to fairly distinguish between both signals and accurately determine the time delay of a lens system \citep[e.g. Q J0158$-$4325;][]{morgan08a,morgan12}. Despite the presence of clear extrinsic variability in HE 1104$-$1805, the time delay between its two images was measured to 4\% precision \citep[1$\sigma$ confidence interval;][]{ofek03}. \citet{ofek03} modelled a long-term microlensing gradient and analised the influence of short-timescale microlensing events (having a mean amplitude of $\sim$ 0.07 mag and a duration of approximately 1 month) in the delay estimation. Using additional data, and Legendre polynomials for describing the intrinsic and extrinsic variabilities, the relative uncertainty in the time delay measurement was decreased to 2\% \citep{poindexter07}. For this system (HE 1104$-$1805), \citet{morgan08a} also re-estimate the delay by modelling the extrinsic fluctuations through microlensing simulations. They obtained an 1$\sigma$ confidence interval in very good agreement with the Ofek \& Maoz's measurement. More recently, \citet{tewes13b} and \citet{rathna13} modelled the long-timescale extrinsic variations of two lens systems (RX J1131$-$1231 and SDSS J1001+5027) by free-knot splines (among other techniques), incorporating short-timescale correlated noise in their analyses. They measured the longest delays with $\leq$ 2\% precision. \citet{hainline13} also modelled the microlensing fluctuations in SBS 0909+532 by intensive simulations. The light curves of this double quasar included an intrinsic deep dip and significant extrinsic variability on different timescales, allowing the authors to estimate the 50-d delay with 6\% precision. As a part of the road map to deeply analyse the system \object{SDSS J1339+1310}, this paper is mainly dedicated to characterising the microlensing signal in the wavelength and time domains, as well as to measure the time delay. The lens system \object{SDSS J1339+1310} consists of two quasar images (A and B) at the same redshift $z_{\rm{s}}$ = 2.231 and separated by 1\farcs70, as well as a lensing galaxy (G) at $z_{\rm{l}}$ = 0.609 and located 0\farcs63 from the B image \citep{inada09,shalyapin14}. While A and B are optically bright ($r \sim$ 18$-$19 mag), the galaxy G has an $r$-band magnitude of about 20.5. Spectra of the lens system were obtained in 2013 using the OSIRIS R500R grism on the 10.4 m Gran Telescopio Canarias (GTC). These GTC-OSIRIS spectra allowed us to measure the redshift of G and obtain constraints on the macrolens and extinction parameters \citep[][henceforth Paper I]{shalyapin14}. We also found evidence for strong chromatic microlensing in its optical continuum (reaching a maximum signal of $\sim$ 1.5$-$1.7 mag at $\sim$ 5000 \AA), and were able to predict a time delay between A and B of $\sim$ 40$-$50 d (A is leading). The presence of sharp extrinsic events in preliminary light curves with the 2.0 m Liverpool Telescope (LT) precluded direct measurement of the delay. In Sect.~\ref{sec:spec}, we present new high-quality GTC (optical) spectra of A, B and G in 2014. In Sect.~\ref{sec:microspec}, we discuss widely the several spectral distortions caused by microlensing. In Sect.~\ref{sec:lcur}, we also present $r$-band LT light curves of A and B in 2009, 2012$-$2015 and early 2016. In Sect.~\ref{sec:delmicvar}, these curves are used to estimate the time delay and the $r$-band microlensing variability. Our conclusions are summarised in Sect.~\ref{sec:end}. We also include two appendices to address some specific issues. | \label{sec:end} We presented GTC spectrophotometric data of the gravitational lens system \object{SDSS J1339+1310} in 2014. As regards previous GTC-OSIRIS observations with the R500R grism at a single epoch \citepalias{shalyapin14}, the exposures in 2014 were taken with the R500R and R500B grisms at two epochs separated by approximately the time delay of the system. The seeing and the $SNR$ were also better. The new spectra of the lensing galaxy G incorporate several absorption features (e.g. Ca\,{\sc ii} $HK$ doublet and G-band) that helped us improve our previous lens redshift determination. We found a practically irrelevant offset of $z_{\rm{l}}$ by a $-$0.3\%, so $z_{\rm{l}}$ = 0.607 $\pm$ 0.001 (1$\sigma$) from the 2014 data. The new spectra of the quasar images A and B include five prominent emission lines (Ly$\alpha$, Si\,{\sc iv}/O\,{\sc iv}], C\,{\sc iv}, C\,{\sc iii}] and Mg\,{\sc ii}), two of which were not formerly observed with the GTC. From a multi-component decomposition of the carbon line profiles \citep[e.g.][]{wills85,kurasz02, marziani10,sluse11}, we derived two narrow-line flux ratios $B/A$, which were then used to achieve a reliable macrolens-extinction solution for a standard (linear) extinction law in G. This solution allowed us to remove macrolens and extinction contributions in our quasar spectra. We note that the narrow components of the carbon lines are related to line emitting gas with typical velocities of $\sim$ 700 (C\,{\sc iii}]) and 1300 (C\,{\sc iv}) km s$^{-1}$, which belong to the inner NLER \citep[e.g.][]{sulentic99,denney12}. We did not find evidence for microlensing effects on the Ly$\alpha$ and Mg\,{\sc ii} line cores. However, the cores of the Si\,{\sc iv}/O\,{\sc iv}], C\,{\sc iv} and C\,{\sc iii}] emission lines were clearly affected by microlensing. We obtained a microlensing magnification ratio (B relative to A) $\mu_{\rm{BA}}$ = 1.15 (0.15 mag) for the C\,{\sc iii}] line core, as well as higher ratios of 1.3$-$1.4 ($\sim$ 0.3$-$0.4 mag) for the high-ionization line cores (Si\,{\sc iv}/O\,{\sc iv}] and C\,{\sc iv}). In addition, $\mu_{\rm{BA}}$ = 1.6 ($\sim$ 0.5 mag) and $\mu_{\rm{BA}}$ = 2 (0.75 mag) for the C\,{\sc iii}] and C\,{\sc iv} broad-line emissions, respectively. This last ratio is about 1/3 of $\mu_{\rm{BA}} \sim$ 5.7 ($\sim$ 1.9 mag) for the nuclear continuum at 5000 \AA. Therefore, if the size of the NCER is comparable to the Einstein radius of the microlensing objects \citep[stars; e.g.][]{schneider06}, the C\,{\sc iv} BLER should be more extended (but not much more) than the NCER. When comparing the C\,{\sc iv} line shape in A and B, we also detected a microlensing-induced distortion. Although the two blue wings coincide well with each other, there is a strong enhancement in the red wing of B relative to that of A. This type of distortion suggests that the C\,{\sc iv} BLER has an anisotropic structure \citep[e.g.][]{schneider90,abajas02,lewis04,sluse12}. For the nuclear continuum, $\mu_{\rm{BA}} \sim$ 5$-$6 at 4000$-$7000 \AA, with larger ratios at shorter wavelengths. All these spectral results (together with other time-domain results; see below) confirm that \object{SDSS J1339+1310} is an uncommon microlensing factory. We also conducted a monitoring campaign on \object{SDSS J1339+1310} in 2009, 2012$-$2015 and early 2016 with the LT in the SDSS $r$ passband. Hence, this campaign spans five observing seasons and one additional month in 2016. The $r$-band light curves of the lensed quasar are characterised by typical photometric accuracies of 1.7\% (A) and 1.4\% (B), and show parallel V-shaped variations of A and B. Besides these prominent dips in both light curves, we also detected significant microlensing variability on different timescales, including the presence of microlensing fluctuations lasting 50$-$100 d. This extrinsic variability should be taken into account to obtain an unbiased measurement of the time delay between images \citep[e.g.][]{goico98,hainline13,tewes13a}. First, considering only the seasonal (long-timescale) microlensing fluctuation over the period 2009$-$2015, we obtained a time delay of 47$^{+2}_{-1}$ d (1$\sigma$ confidence interval; A is leading). Second, the estimation of the delay error turned out to be difficult for a more detailed (realistic) microlensing model. However, using a reasonable prior on the delay distribution from simulated light curves (we exclusively focused on solutions in the interval 30$-$60 d), we found a 47-d value with $\sim$ 10\% precision. This broader range of delays practically coincides with the range of plausible solutions from the observed light curves and the realistic microlensing model, and we adopted 47$^{+5}_{-6}$ d as our final 1$\sigma$ estimation. We also note that the observed interval overlaps with the delay interval predicted by lens models \citepalias{shalyapin14}. Finally, the time delay and the microlensing effects we report here are useful tools for doing several types of astrophysical studies \citep[e.g.][]{schneider06}. | 16 | 9 | 1609.07440 |
1609 | 1609.08171_arXiv.txt | We study the large solar energetic particle (SEP) events that were detected by GOES in the $>$ 10 MeV energy channel during December 2006 to March 2014. We derive and compare solar particle release (SPR) times for the 0.25--10.4 MeV electrons and 10--100 MeV protons for the 28 SEP events. In the study, the electron SPR times are derived with the time-shifting analysis (TSA) and the proton SPR times are derived using both the TSA and the velocity dispersion analysis (VDA). Electron anisotropies are computed to evaluate the amount of scattering for the events under study. Our main results include: 1) near-relativistic electrons and high-energy protons are released at the same time within 8 min for most (16 of 23) SEP events. 2)There exists a good correlation between electron and proton acceleration, peak intensity and intensity time profiles. 3) The TSA SPR times for 90.5 MeV and 57.4 MeV protons have maximum errors of 6 min and 10 min compared to the proton VDA release times, respectively, while the maximum error for 15.4 MeV protons can reach to 32 min. 4) For 7 low-intensity events of the 23, large delays occurred between 6.5 MeV electrons and 90.5 MeV protons relative to 0.5 MeV electrons. Whether these delays are due to times needed for the evolving shock to be strengthened or due to particle transport effects remains unsolved. | The origin of energetic particles accelerated in solar events is still an open question. While flares and shocks driven by coronal mass ejections (CMEs) are believed to be two sources of solar energetic particle (SEP) acceleration in impulsive and gradual SEP events respectively \citep[e.g.][]{Reames1999}, it is not clear what the exact flare-related acceleration mechanism in the impulsive SEP events is or where the CME-driven shocks most efficiently accelerate particles and when the particles are released in gradual SEP events. Electron release has been observed to temporally coincide with type III radio bursts at the Sun and traveling along open field lines into interplanetary space \citep[see][]{Lin1985}. Using observations of the Three-dimensional Plasma and Energetic Particles instrument \citep[3DP;][]{Lin1995} on the Wind spacecraft, \citet{Wang2006} studied three electron events and found two distinct injections of electrons: that of low-energy electrons at energies $\sim$ 0.4 to 6--9 keV began 9.1 min before the type III radio burst and that of $\sim$ 13 to 300 keV electrons started 7.6 min after the type III burst. Delays of 10 min up to half an hour between electron release time at the Sun and solar electromagnetic emissions (EM) have been reported by other works \citep[e.g.][]{Cliver82,Kallenrode91,Krucker1999, Haggerty02}. \citet{Krucker1999} showed evidence that some electron events are not related to type III bursts. They found that the electron events appeared to be related to the passage of large-scale coronal transient waves, also called EIT waves or Extreme Ultraviolet (EUV) waves \citep{Thompson1998, Thompson2000}, over the footpoint of the field line connected to the spacecraft. Although the nature of EUV waves is still largely debated, past studies using low-cadence ultraviolet images ($>$12 minutes) showed that EUV waves are correlated with CMEs rather than flares \citep{Plunkett1998, Cliver1999}. Based on more recent three-dimensional stereoscopic analyses, the EUV waves are generally believed to be the imprint of the CME driven shock on solar surface \citep[e.g.][]{Veronig2008, Patsourakos2009}. Proton release is probably more complicated than electron release. \citet{Krucker2000} studied the timing of proton onsets in the energy range from 30 keV to 6 MeV. They found that the release of the protons appears to be energy-dependent. The most energetic protons are possibly released simultaneously with the electrons while lower-energy protons are released $\sim$ 0.5 to 2 hrs later than electrons. They also found that protons with energies between 0.03 and 6 MeV are released high in the corona, around 1--10 Rs above the electrons. Their results are consistent with studies by \citet{Kahler1994} and \citet{Gopalswamy2012} on the CME heights at the time of SEP release. \citet{Kahler1994} analyzed $>$ 10 MeV proton events and found that the peak of the intensity profile for $>$ 10 MeV protons occurs when the associated CME reaches heights of 5--15 Rs. \citet{Gopalswamy2012} examined the onset times and release heights of energetic particles using the ground-level enhancement (GLE) events. They found an earlier release time and a lower release height of CMEs for this highly energetic subset of events. Although both SEP solar particle release (SPR) times and EM onsets have been discussed at length in the past, comparison between electron release times and proton release times has been discussed only in a few papers \citep[e.g.][]{Cliver82,Haggerty09,Koulou15,Kahler2003,Posner07}. In \citet{ Posner07}'s study, the author adopted the prevailing assumption of simultaneous release of electrons and protons, but he also pointed out that ``release of protons before electrons (and vice versa) is possible [E. Roelof, D. Haggerty, personal communication, 2006]''. Using a group of 32 historic GLE events, \citet{Cliver82} found that delays of 10 minutes between 100 keV and 1 MeV electron SPRs and $\le$ 5 minute delays between 2 GeV proton and 1 MeV electron SPRs. Delays of 10 to 50 minutes in the proton SPRs relative to metric type II onsets for well connected events were found in the smaller GLE events. \citet{Kahler2003} compared the onset of relativistic electrons and protons of GLEs from solar cycle 23. They found that half of GLE events the relativistic proton injection preceded that of electrons, however, the low intensity GLEs tend to have a later time for the proton injection. Recently, \citet{Koulou15} compared the proton and electron release as inferred from VDA based on Wind/3DP and ERNE data, and found a 7-min average dalay of near-relativistic electrons with respect to deka-MeV protons. \citet{Haggerty09} studied 19 electron beam events using EPAM 38-315 keV data, and found that for 11 of the 19 events the arrival of 50-100 MeV protons followed by electrons within $\sim$ 3 min. On the other hand, the remaining 8 events show a broad 5-25 minute delays of the protons relactive to the electron injections. In this paper, we study large SEP events with a peak $>$ 10 MeV proton flux above 10 $cm^2 sr^{-1} s^{-1}$ as observed by GOES from October 2006 (the launch of the Solar and Terrestrial Relations Observatory (STEREO)) to March 2014. The proton SPR times at various energies from 10 MeV to 131 MeV are investigated and compared to the release times of 0.25 MeV--10.4 MeV electrons and solar EM onsets. % The exact time when energetic particles are first released at the Sun is crucial to understanding the particle acceleration and where it takes place. This is the first systematic study and comparison between electron and proton SPRs for large SEP events in the new STEREO era. The paper aims to address the following key issues: 1) Are protons and electrons accelerated by the same source and released simultaneously at the Sun? 2) What is the acceleration time needed for protons and electrons to reach high energies and are the acceleration times energy dependent? | \subsection{Summary} By choosing the smallest CA among the three spacecraft, we derive and compare the high energy electron and proton SPR times using SOHO/EPHIN electron fluxes in the 0.25--10.4 MeV channels, SOHO/ERNE proton fluxes in the 13.8--101 MeV channels, or in the similar energy channels of the SEPT and HET (LET) detectors on STEREO. Our main results are listed below. \begin{itemize} \item The e2 release times are found to be systematically larger than the e1 release times by an average of 6.8 min and 7.3 min, for the 12 SOHO SEPs and 10 STEREO SEPs, respectively. Among these 22 events, three events (6, S5, and S10) have a large 10--28 min delay. \item The p2 protons are shown to have similar SPR times with the p1 protons. The average delay between the p2--p1 SPRs are $\sim$ 2.5 min and -0.2 min, for the 12 SOHO SEPs and 9 STEREO SEPs, respectively. For the p0 protons, there are 12 SEP events showing small delays between the p2--p0 SPRs within 5 min and five events (9,S2, S3, S9 and S11) showing a large 10--32 min delay due to proton scattering effects. \item The proton VDA results show that protons are released simultaneously with the e1 electrons within 8 min for $\sim$ 70\%(16 of 23) SEP events, and the e2 electrons with 6 min for 13 of 19 events. There are $\sim$ 30\%(7 of 23) SEP events showing a delayed proton release time by $\sim$ 8--31 min. Among these 6 events, 3 events (6, S5, and S10) also have a large e2-e1 SPR delay. \item $\sim$ 65\% (15 of 23) protons events show a small scattered path length ($<$ 1.5 AU); 8 of 23 proton events have a large apparent path length ($>$ 1.5 AU), part of reason is due to higher background levels in the STEREO HET data. \item The delays between e1 SPRs and type III onsets range from 2 min to 42 min. The CME heights at the e1 release times range from 2.1 to 9.1 Rs. From the CME heights, it is likely that the e1 electrons are accelerated by the CME-driven and/or flare shock waves rather than flare reconnections. \end{itemize} \subsection{Discussion} \subsubsection{Association between Electrons and Protons} Our results are consistend with \citet{Haggerty09}'s study, where they suggested that near-relatic electrons and the energetic protons are accelerated and released by essentially the same mechansim(s). \citet{Haggerty09} studied the injection times of near-relativistic electrons and non-relativistic protons for 19 electron beam events using ERNE 50-100 MeV proton and EPAM 38-315 keV electron data, and found that 11 of the 19 events (60\%) are statistically consistent with zero delay between the proton and electron injection within the uncertainty of $\sim$ 3min. The remaining 8 events show a broad 5-25 minute delays of the protons relactive to the electron injections. They also compared the peak intensity of 175--315 keV elections with that of 1.8--4.7 MeV protons from ACE/EPAM and found a good correlation in the peak intensity of electrons and protons. \begin{figure} \noindent\includegraphics[width=0.7\txw]{int_corr.eps} \caption{ Logarithmic peak intensity correlation between the e1 electrons and the p0 protons.} \label{int_corr} \end{figure} \begin{figure}[t!] \mbox{ \includegraphics[width=0.98\txw, height = 0.45\txw ]{Fig13_merge.eps} }\par \caption{ (Left) over-plotted electron and proton intensity from ACE/EPAM, EPHIN, and ERNE on December 13 2006. (Right) over-plotted electron and proton intensity from WIND/3DP electrons, EPHIN, and WIND/EPACT protons on November 26 2011. The intensity has been normalized to the background flux level for easy comparison.} \label{figall} \end{figure} Among the 28 SEP events under study, we found similar correlations between the peak intensity of e1 electrons and p0 protons, as shown in Figure~\ref{int_corr}. Futhermore, the profiles between different spices are found to be very similar to each other although not identical, as shown in Figure~\ref{figall}. Our results support the conclusion that near-relativistic electron and high-energetic proton acceleration are closely related to each other. On the other hand, how the intensity profiles evolve with time, which result from the transport-modulated SEP particle accelerations at an evolving CME-driven shock, is not well understood. For example, at the SEP rise phase, it is not well understood why the e2 electrons are the last to reach their peak value for event 1 (left in Figure~\ref{figall}); while in the second example (event 8, right panel in Figure~\ref{figall}), the e2 electrons reach the plateau before the protons. \subsubsection{Direct Shock Accleration vs Tranverse Transport} Besides simultaneously released electron and proton events, there are seven events showing large delays of 8--31 min between proton release times $t_{SPR}(p_{vda})$ from VDA and e1 SPRs. These events are SEPs with small e2 and p2 intensities. Three possible reasons may account for these large delays: 1) the late formed shocks at high altitudes around DH type II onset times; 2) longer times needed for the evolving shocks to be intense enough to produce high-energy SEPs after DH type II onsets; 3) times needed for shocks in SEP events with large CAs to reach the magnetic connection footpoint to the observer. Among the above 7 events, events 8 and S6 have small CAs of 6$^\circ$ and 3$^\circ$, events 3 and S5 have large CAs of 30$^\circ$ and 32$^\circ$, and the other 3 events (3, 6, and S10) have intermediate CAs of 13-21 $^\circ$. 6 of these events have the similar e1 SPRs with the DH type II onsets within 5 min ( and a large 13-26 min delay between $t_{SPR}(p_{vda})$ and metric type II onsets) except event S10. The obtained timing comparising results are consistent with one (or two) of the above three hypotheses. Rouillard et al. (2012) investigated the 2011 March 21 SEP event using STEREO and SOHO observations. By tracking the CME shock lateral expansion they demonstrated that the delayed solar particle release times are consistent with the time required for the shock to propagate to the magnetic footpoint connecting to the observer. On the other hand, for large CA and/or high latitude SEPs, an alternative (or contributing) explanation is that the delay between the SEP release and electromagnetic emissions is caused by the propagation times needed for the SEP particles to transport across the field line to the connection footpoint of the observer \citep[e.g.][]{Dresing12, Qin13, Laiti15}. It is possible that both direct shock acceleration and cross-field propagation of SEPs play roles in the formation of SEP intensity time profile. At an evolving CME-driven shock near the Sun, many factors such as the shock obliquity, the compression ratio and transport parameters may affect the SEP intensity, further investigations are needed. \subsection{Conclusion} Our results suggest that near-relativistic electron and high-energy proton acceleration are closely related to each other. There exists a good association between high-energy electron and proton release time, intensity peak values and time profiles. % For small intensity SEP events, it takes longer times for the e2 and p2 to reach up to the detectable flux levels. However, whether this delay is due to the times that needed for the evolving shock to be strengthened or due to particle transport effects are not resolved. | 16 | 9 | 1609.08171 |
1609 | 1609.01050_arXiv.txt | The invariance of the Lagrangian under time translations and rotations in Kepler's problem yields the conservation laws related to the energy and angular momentum. Noether's theorem reveals that these same symmetries furnish generalized forms of the first integrals in a special nonconservative case, which approximates various physical models. The system is perturbed by a biparametric acceleration with components along the tangential and normal directions. A similarity transformation reduces the biparametric disturbance to a simpler uniparametric forcing along the velocity vector. The solvability conditions of this new problem are discussed, and closed-form solutions for the integrable cases are provided. Thanks to the conservation of a generalized energy, the orbits are classified as elliptic, parabolic, and hyperbolic. Keplerian orbits appear naturally as particular solutions to the problem. After characterizing the orbits independently, a unified form of the solution is built based on the Weierstrass elliptic functions. The new trajectories involve fundamental curves such as cardioids and logarithmic, sinusoidal, and Cotes' spirals. These orbits can represent the motion of particles perturbed by solar radiation pressure, of spacecraft with continuous thrust propulsion, and some instances of Schwarzschild geodesics. Finally, the problem is connected with other known integrable systems in celestial mechanics. | Finding first integrals is fundamental for characterizing a dynamical system. The motion is confined to submanifolds of lower dimensions on which the orbits evolve, providing an intuitive interpretation of the dynamics and reducing the complexity of the system. In addition, conserved quantities are good candidates when applying the second method of Lyapunov for stability analysis. Conservative systems related to central forces are typical examples of (Liouville) integrability, and provide useful analytic results. Hamiltonian systems have been widely analyzed in the classical and modern literature to determine adequate integrability conditions. The existence of first integrals under the action of small perturbations occupied \citet[][Chap.~V]{poincare1892methodes} back in the 19th century. Later, Emmy \citet{noether1918invariante} established in her celebrated theorem that conservation laws can be understood as the system exhibiting dynamical symmetries. In a more general framework, \citet{yoshida1983necessary,yoshida1983necessaryII} analyzed the conditions that yield algebraic first integrals of generic systems. He relied on the Kowalevski exponents for characterizing the singularities of the solutions and derived the necessary conditions for existence of first integrals exploiting similarity transformations. Conservation laws are sensitive to perturbations and their generalization is not straightforward. For example, the Jacobi integral no longer holds when transforming the circular restricted three-body problem to the elliptic case \citep{xia1993arnold}. Nevertheless, \citet{contopoulos1967integrals} was able to find approximate conservation laws for orbits of small eccentricities. \citet{szebehely1964elliptic} benefited from the similarities between the elliptic and the circular problems in order to define transformations connecting them. \citet{henon1964applicability} deepened in the nature of conservation laws and reviewed the concepts of isolating and nonisolating integrals. Their study introduced a similarity transformation that embeds one of the constants of motion and transforms the original problem into a simplified one, reducing the degrees of freedom \citep[][\S3.2]{arnold2007mathematical}. \citet{carpintero2008finding} proposed a numerical method for finding the dimension of the manifold in which orbits evolve, i.e. the number of isolating integrals that the system admits. The conditions for existence of integrals of motion under nonconservative perturbations received important attention in the past due to their profound implications. \citet{djukic1975noether} advanced on Noether's theorem and included nonconservative forces in the derivation. Relying on Hamilton's variational principle, they not only extended Noether's theorem, but also its inverse form and the Noether-Bessel-Hagen and Killing equations. Later studies by \citet{djukic1984integrating} sought integrating factors that yield conservation laws upon integration. Examples of application of Noether's theorem to constrained nonconservative systems can be found in the work of \citet{bahar1987extension}. \citet{honein1991conservation} arrived to a compact formulation using what was later called the neutral action method. Remarkable applications exploiting Noether's symmetries span from cosmology \citep{capozziello2009noether,basilakos2011using} to string theory \citep{beisert2008dual}, field theory \citep{halpern1977field}, and fluid models \citep{narayan1987physics}. In the book by \citet[Chaps.~4 and 5]{olver2000applications}, an exhaustive review of the connection between symmetries and conservation laws is provided within the framework of Lie algebras. We refer to \citet[][Chap.~3]{arnold2007mathematical} for a formal derivation of Noether's theorem, and a discussion on the connection between conservation laws and dynamical symmetries. Integrals of motion are often useful for finding analytic or semi-analytic solutions to a given problem. The acclaimed solution to the satellite main problem by \citet{brouwer1959solution} is a clear example of the decisive role of conserved quantities in deriving solutions in closed form. By perturbing the Delaunay elements, \citet{brouwer1961theoretical} solved the dynamics of a satellite subject to atmospheric drag and the oblateness of the primary. They proved the usefulness of canonical transformations even in the context of nonconservative problems. \citet[][pp.~81--82]{whittaker1917treatise} approached the problem of a central force depending on powers of the radial distance, $r^n$, and found that there are only fourteen values of $n$ for which the problem can be integrated in closed form using elementary functions or elliptic integrals. Later, he discussed the solvability conditions for equations involving square roots of polynomials \citep[p.~512]{whittaker1927course}. \citet{broucke1980notes} advanced on Whittaker's results and found six potentials that are a generalization of the integrable central forces discussed by the latter. These potentials include the referred fourteen values of $n$ as particular cases. Numerical techniques for shaping the potential given the orbit solution were published by \citet{carpintero1998orbit}. Classical studies on the integrability of systems governed by central forces are based strongly on Newton's theorem of revolving orbits.\footnote{Section IX, Book I, of Newton's Principia is devoted to the motion of bodies in moveable orbits (\emph{De Motu Corporum in Orbibus mobilibus, deq; motu Apsidum}, in the original latin version). In particular, Thm.~XIV states that ``The difference of the forces, by which two bodies may be made to move equally, one in a quiescent, the other in the same orbit revolving, is in a triplicate ratio of their common altitudes inversely''. Newton proved this theorem relying on elegant geometric constructions. The motivation behind this result was the development of a theory for explaining the precession of the orbit of the Moon. A detailed discussion about this theorem can be found in the book by \citet[pp.~184--201]{chandrasekhar1995newton}} The problem of the orbital precession caused by central forces was recently recovered by \citet{adkins2007orbital}, who considered potentials involving both powers and logarithms of the radial distance, and the special case of the Yukawa potential \citep{yukawa1935interaction}. \citet{chashchina2008remark} relied on Hamilton's vector to simplify the analytic solutions found by \citet{adkins2007orbital}. More elaborated potentials have been explored for modeling the perihelion precession \citep{schmidt2008perihelion}. The dynamics of a particle in Schwarzschild space-time can also be regarded as orbital motion perturbed by an effective potential depending on inverse powers of the radial distance \citep[][p.~102]{chandrasekhar1983mathematical}. Potentials depending linearly on the radial distance appear recursively in the literature because they render constant radial accelerations, relevant for the design of spacecraft trajectories propelled by continuous-thrust systems. The pioneering work by \citet{tsien1953take} provided the explicit solution to the problem in terms of elliptic integrals, as predicted by \citet[][p.~81]{whittaker1917treatise}. By means of a special change of variables, \citet{izzo2015explicit} arrived to an elegant solution in terms of the Weierstrass elliptic functions. These functions were also exploited by \citet{macmillan1908motion} when he solved the dynamics of a particle attracted by a central force decreasing with $r^{-5}$. \citet{Urrutxua2015} solved the Tsien problem using the Dromo formulation, which models orbital motion with a regular set of elements \citep{pelaez2007special,urrutxua2015dromo,roa2015singularities}. Advances on Dromo can be found in the works by \citet{bau2015non} and \citet{roa2015orbit}. The case of a constant radial force was approached by \citet{akella2002anatomy} from an energy-driven perspective. They studied in detail the roots of the polynomial appearing in the denominator of the equation to integrate, and connected their nature with the form of the solution. General considerations on the integrability of the problem can be found in the work of \citet{san2012bounded}. Another relevant example of an integrable system in celestial mechanics is the Stark problem, governed by a constant acceleration fixed in the inertial frame. \citet{lantoine2011complete} provided the complete solution to the motion relying extensively on elliptic integrals and Jacobi elliptic functions. A compact form of the solution involving the Weierstrass elliptic functions was later presented by \citet{biscani2014stark}, who also exploited this formalism for building a secular theory for the post-Newtonian model \citep{biscani2012first}. The Stark problem provides a simplified model of radiation pressure. In the more general case, the dynamics subject to this perturbation cannot be solved in closed form. An intuitive simplification that makes the problem integrable consists in assuming that the force due to the solar radiation pressure follows the direction of the Sun vector. The dynamics are equivalent to those governed by a Keplerian potential with a modified gravitational parameter. The present paper introduces a new class of integrable system, governed by a biparametric nonconservative perturbation. This acceleration unifies various force models, including special cases of solar radiation pressure, low-thrust propulsion, and some particular configurations in general relativity. The problem is formulated in Sec.~\ref{Sec:Dynamics}, where the biparametric acceleration is defined and then reduced to a uniparametric forcing thanks to a similarity transformation. The conservation laws for the energy and angular momentum are generalized to the nonconservative case by exploiting known symmetries of Kepler's problem. Before solving the dynamics explicitly, we will prove that there are four cases that can be solved in closed form using elementary or elliptic functions. Sections~\ref{Sec:conic}--\ref{Sec:sinusoidal} present the properties of each family of orbits and the corresponding trajectories are derived analytically. Section~\ref{Sec:summary} is a summary of the solutions, which are unified in Sec.~\ref{Sec:Weierstrass} introducing the Weierstrass elliptic functions. Finally, Sec.~\ref{Sec:connection} discusses the connection with known solutions to similar problems, and with Schwarzschild geodesics. | The dynamical symmetries in Kepler's problem hold under a special nonconservative perturbation: a disturbance that modifies the tangential and normal components of the gravitational acceleration in the intrinsic frame renders two integrals of motion, which are generalized forms of the equations of the energy and angular momentum. The existence of a similarity transformation that reduces the original problem to a system perturbed by a tangential uniparametric forcing simplifies the dynamics significantly, for the integrability of the system is evaluated in terms of one single parameter. The algebraic properties of the equations of motion dictate what values of the free parameter make the problem integrable in closed form. The extended integrals of motion include the Keplerian ones as particular cases. The new conservation laws can be seen as generalizations of the original integrals. The new families of solutions are defined by fundamental curves in the zero-energy case, and there are geometric transformations that relate different orbits. The orbits can be unified by introducing the Weierstrass elliptic functions. This approach simplifies the modeling of the system. The solutions derived in this paper are closely related to different physical problems. The fact that the magnitude of the acceleration decreases with $1/r^2$ makes it comparable with the perturbation due to the solar radiation pressure. Moreover, the inverse similarity transformation converts Keplerian orbits into the conic sections obtained when the solar radiation pressure is directed along the radial direction. The structure of the solutions, governed by the roots of a polynomial, is similar in nature to the Schwarzschild geodesics. This is because under the considered perturbation and in Schwarzschild metric the evolution of the radial distance takes the same form. The perturbation can also be seen as a control law for a continuous-thrust propulsion system. Some of the solutions are comparable with the orbits deriving from potentials depending on different powers of the radial distance. Although the trajectory may take the same form, the velocity will be different, in general. | 16 | 9 | 1609.01050 |
1609 | 1609.06446_arXiv.txt | {We present the stellar kinematic maps of a large sample of galaxies from the integral-field spectroscopic survey CALIFA. The sample comprises 300 galaxies displaying a wide range of morphologies across the Hubble sequence, from ellipticals to late-type spirals. This dataset allows us to homogeneously extract stellar kinematics up to several effective radii. In this paper, we describe the level of completeness of this subset of galaxies with respect to the full CALIFA sample, as well as the virtues and limitations of the kinematic extraction compared to other well-known integral-field surveys. In addition, we provide averaged integrated velocity dispersion radial profiles for different galaxy types, which are particularly useful to apply aperture corrections for single aperture measurements or poorly resolved stellar kinematics of high-redshift sources. The work presented in this paper sets the basis for the study of more general properties of galaxies that will be explored in subsequent papers of the survey.} | \label{S:intro} The motion of stars within galaxies is a fundamental property set very early on in their life. Ever since the detection of rotation of stars in the Milky Way and nearby systems \citep[e.g.,][]{lindblad1927,mayall1951,munch1960}, the study of stellar motions has been a fruitful avenue to pose important constraints on our knowledge about galaxy formation and evolution. The analysis of rotational over random motions in early-type galaxies, for instance, has led to the realization that bright early-type galaxies are likely triaxial objects supported by orbital anisotropy \citep[e.g.,][]{bertola75,illingworth77,binney78}, rather than rotation. The coupling of long-slit spectrographs with telescopes 2 to 4\,m in size has provided, over the last three decades, a wealth of spatially resolved observations that has greatly improved our understanding of the overall stellar motion and level of kinematic substructure in external galaxies \citep[e.g.,][]{defis83,bertola84,bender94,fisher97,simien97, rubin99, vegabeltran01,aguerri03,fb03,pizzella04,remco15}. While the first integral-field units (IFUs) were already in place in the mid-90's \citep[e.g.,][]{bacon95}, the first serious efforts to measure stellar kinematics on large samples of galaxies using these kinds of instruments did not occur until year 2001. One of the pioneer projects in this respect was the SAURON survey \citep{bacon01, dezeeuw02}. With a representative sample of 72 galaxies (24 ellipticals, 24 lenticulars, and 24 early-type spirals, later extended with observations of 18 late-type spirals), this survey has set the reference for stellar kinematic IFU studies \citep[e.g.,][]{emsellem04,fb06,ganda06}. The discovery of the slow and fast rotator families in early-type galaxies \citep{emsellem07} served as the trigger for a larger project: the ATLAS$^{\rm 3D}$ survey \citep{cappellari_etal_2011}, in which a volume complete sample of 260 early-type galaxies revisited many kinematic aspects, from the amount of global angular momentum \citep{emsellem_etal_2011} to a detailed account of kinemetric features \citep{krajnovic06,krajnovic_etal_2011}. In parallel, the DiskMass survey mapped, the stellar kinematic properties of nearby late-type spirals with the aid of the PPak IFU \citep{Roth_etal_2005, Kelz_etal_2006}. The CALIFA survey \citep{Sanchez_etal_2012} was born to fill in existing gaps in other IFU surveys and to provide a morphologically unbiased view of the stellar kinematics in galaxies based on a large ($\sim$\,600 galaxies) and homogeneous integral-field spectroscopic dataset. The main advantage of CALIFA over existing surveys resides in a sample selection that includes all morphological types, as well as a field-of-view (FoV) that extends up to several effective radii (\Reff). While CALIFA is no longer the IFU survey with the largest number of observed objects in the nearby Universe, it still provides the best compromise between spatial coverage (1.8--3.7\,\Reff) and sampling ($\sim$1\,kpc). Currently ongoing IFU surveys are hampered in one way or another by these factors, for example, SAMI covers areas within 1.1--2.9\,\Reff\ with a spatial sampling $\sim$\,1.7\,kpc \citep{sami,bryant15}, while MaNGA primary sample targets have a spatial sampling of $\sim$3\,kpc within 1.5\,\Reff\ \citep{manga}. The real revolution in this respect will take place when MUSE at the Very Large Telescope \citep{bacon10} is used in survey mode, as anticipated by the very spectacular stellar kinematic cases presented in the first few years of operations \citep[e.g.,][]{emsellem14,davor15,gadotti15,iodice15}.\looseness-1 The goal of this paper is to present the first stellar kinematic maps extracted from the CALIFA survey, describe all the technical details of the extraction, and provide basic stellar velocity dispersion aperture corrections for elliptical and spiral galaxies. The maps presented here have already been used within the survey to establish the effect of galaxy interactions on the stellar kinematics of galaxies \citep{barrera14,barrera15}, constrain the pattern speed of barred galaxies across the Hubble sequence \citep{aguerri15}, to present a volume-complete Tully-Fisher relation \citep{simona16}, and the velocity function of galaxies as a benchmark for numerical simulations \citep{simona16b}. Forthcoming papers of the survey will make use of this information, for example, to revisit the distribution of global angular momentum in nearby galaxies and determine their dark matter content. \citet{fb15} provides a preview of some highlights. For results on the kinematics of the ionized gas in CALIFA, see \citet{garcialorenzo15}. The paper is organized as follows. Section~\S\ref{S:sample} describes the sample of 300 galaxies used in our study and how this sample compares with the full CALIFA sample. Section~\ref{S:setup} summarizes the instrumental setup employed during the observations. In section~\ref{S:kinextraction} we provide details of our kinematic extraction and comparisons with other major IFU surveys. Section~\ref{S:sigma_limit} explains the limit set by our instrumental setup in the measurement of stellar velocity dispersions. In section~\ref{S:aperture_corr} we provide velocity dispersion aperture corrections for elliptical and spiral galaxies. Finally, we summarize our work and conclusions in section~\ref{S:conclusions}. \begin{figure} \centering \includegraphics[angle=0,width=\linewidth]{califa_sample_distribution-eps-converted-to_lowres.jpg} \caption{Distribution of galaxies in the sample of CALIFA galaxies presented in this paper (see \S\ref{S:sample}) as a function of Hubble type, stellar mass, and absolute magnitude in the $r$ band. For convenience, along with the color bar, we indicate the number of galaxies in each bin.} \label{fig:sample} \end{figure} | \label{S:conclusions} In this paper we present stellar kinematic maps for a sample of 300 galaxies that are part of the CALIFA survey. The sample covers a wide range of Hubble types, from ellipticals to late-spiral galaxies. This subset is a good representation of the CALIFA mother sample in terms of redshift, isophotal diameter, and absolute magnitude. The large footprint of the PPak IFU, together with the average distance of the survey, allow us to measure stellar kinematics well beyond 1.8\,\Reff\ for 50\% of the galaxies, reaching out to \mbox{4--5\,\Reff} in a few exceptional cases. The penalty, caused by the combination of spatial sampling and distance, is the inability to detect kinematically decoupled components at the centers of galaxies. Still our data is well suited for the study of large-scale kinematic twists or long-axis rotation, which occurs in a handful of objects. The measurements presented in this paper are in good agreement with those of other well-known IFU surveys (e.g., ATLAS$^{\rm 3D}$ and DiskMass). The detailed comparison with the DiskMass survey allowed us to establish that we can measure reliable velocity dispersion values down to $\sigma$\,$\sim$\,40\,\kms\ (i.e., $\sim$\,30\,\kms\ below the instrumental resolution). We also characterized the relative uncertainties of our measurements, which are around 5\% for $\sigma$\,$\ge$\,150\,\kms. Below that value, relative uncertainties increase up to 50\% for velocity dispersions all the way down to $\sigma$\,$\sim$\,20\,\kms. We also took advantage of our large sample to compute integrated stellar velocity dispersion aperture corrections for different sets of galaxies across the Hubble sequence. These corrections are particularly useful to homogenize dispersion values of galaxies at different distances. We find two main classes of integrated aperture radial profiles: steadily decreasing profiles representative of early-type galaxies, and a second class of systematically increasing profiles typical of late-type spiral galaxies. We provide aperture corrections for each class for different stellar masses and absolute magnitudes. The main properties of the sample and the stellar velocity and velocity dispersion maps introduced in this paper are available as part of the Online Material in Table~1 and Appendix~\ref{A:kin_maps}. The values of the maps themselves, together with many diagnostic parameters to assess the quality of the measurements, will be made available to the community at the CALIFA website (\url{http://califa.caha.es}). | 16 | 9 | 1609.06446 |
1609 | 1609.01785_arXiv.txt | It is generally accepted that silicate-metal (`rocky') planet formation relies on coagulation from a mixture of sub-Mars sized planetary embryos and (smaller) planetesimals that dynamically emerge from the evolving circum-solar disc in the first few million years of our Solar System. Once the planets have, for the most part, assembled after a giant impact phase, they continue to be bombarded by a multitude of planetesimals left over from accretion. Here we place limits on the mass and evolution of these planetesimals based on constraints from the highly siderophile element (HSE) budget of the Moon. Outcomes from a combination of N-body and Monte Carlo simulations of planet formation lead us to four key conclusions about the nature of this early epoch. First, matching the terrestrial to lunar HSE ratio requires either that the late veneer on Earth consisted of a single lunar-size impactor striking the Earth before 4.45~Ga, or that it originated from the impact that created the Moon. An added complication is that analysis of lunar samples indicates the Moon does not preserve convincing evidence for a late veneer like Earth. {Second, the expected chondritic veneer component on Mars is 0.06 weight percent.} Third, the flux of terrestrial impactors must have been low ($ \lesssim 10^{-6}~M_\oplus$~Myr$^{-1}$) to avoid wholesale melting of Earth's crust after 4.4~Ga, and to simultaneously match the number of observed lunar basins. This conclusion leads to an Hadean eon which is more clement than assumed previously. Last, after the terrestrial planets had fully formed, the mass in remnant planetesimals was $\sim10^{-3}~M_\oplus$, lower by at least an order of magnitude than most previous models suggest. Our dynamically and geochemically self-consistent scenario requires that future N-body simulations of rocky planet formation either directly incorporate collisional grinding or rely on pebble accretion. | \label{sec:int} The formation of the terrestrial planets is a long-standing problem that is gradually being resolved. In traditional ㄑdynamical models the terrestrial planets grow from a coagulation of planetesimals into protoplanets and subsequently evolve into a giant impact phase, during which the protoplanets collide with each other to lead to the terrestrial planets. Several variations of this scenario exist, of which the {\it Grand Tack} model is currently popular \citep{W11}. The Grand Tack relies on early gas-driven migration of Jupiter and Saturn to gravitationally sculpt the inner solid circum-solar disc down to $\sim$1~AU after which terrestrial planet formation proceeds from solids in an annulus ranging from roughly 0.7~AU to 1~AU. Grand Tack has booked some successes, such as its ability to reproduce the mass-orbit distribution of the terrestrial planets, the compositional gradient, and total mass of the asteroid belt \citep{W11}. Subsequent evolution of the solar system after terrestrial planet formation, all the way to the present, however, has mostly been studied in separate epochs with disconnected simulations. \\ \citet{B16} scrutinised the Grand Tack model in more detail { and built a database of simulations that is used here}. Since published simulations had rarely been run for much longer than 200~Myr into the evolution of the solar system, we sought to test the model predictions specific to the long-term evolution of the terrestrial system. In this work we calculate the { evolution of the terrestrial planets for up to 300~Myr. We aim to obtain the} amount of mass accreted by the terrestrial planets after the Moon-forming event and whether this accreted mass is compatible with { the highly-siderophile element (HSE) budgets of the inner planets}, the early lunar and terrestrial cratering records and the nature of the purported late veneer. We also consider the surface conditions on the Hadean Earth from geochemical data and conclude with the implications of our simulations for future models of terrestrial planet formation. | Here, we determined amount of mass accreted by the terrestrial planets after the Moon-forming event. We compared the outcomes of N-body simulations of the dynamical evolution of the terrestrial planets and remnant planetesimals with Monte Carlo simulations of impacts on the Earth, Mars and Moon. By combining the outcome of these simulations with the lunar HSE budget, cratering record, geochronology and geochemistry, it is shown that \begin{enumerate} \item from HSE analysis the total mass in planetesimals at the time of the Moon-forming event had to be $\sim10^{-3}$~$M_\oplus$ (Section 2); { \item the expected chondritic contribution to Mars during the late veneer epoch is 0.06 wt\% (Section 5); \item the high terrestrial HSE budget was most likely caused by a singular event unique to Earth (Section 6.1); \item the remnant planetesimals from terrestrial planet formation may account for most lunar basins, but not necessarily most of the small craters (Section 6.2);} \item the surface conditions of the Hadean Earth were much more clement than commonly thought because of a low-intensity bombardment (Section 6.3); \item new terrestrial planet formation models will need to take collisional evolution into account and explore alternatives to oligarchic growth such as pebble accretion (Section 6.4). \end{enumerate} | 16 | 9 | 1609.01785 |
1609 | 1609.08806_arXiv.txt | {The detection of bright, hard, and variable X-ray emission in \tr\ prompted spectropolarimetric observations of this star, which in turn led to the discovery of a surface magnetic field.} {We want to further constrain the properties of this star, in particular to verify whether X-ray variations are correlated to changes in optical emission lines and magnetic field strength, as expected from the oblique rotator model that is widely accepted for magnetic O stars.} {We have obtained new low-resolution spectropolarimetric and long-term high-resolution spectroscopic monitoring of \tr, and we also analyse new, serendipitous X-ray data. } {The new X-ray observations are consistent with previous data, but their addition does not help to solve the ambiguity in the variation timescale because of numerous aliases. No obvious periodicity or any large variations are detected in the spectropolarimetric data of \tr\ obtained over three months. The derived field values appear to be in line with previous measurements, suggesting constancy of the field (though the possibility of small, short-term field variations cannot be excluded). Variations in the equivalent widths of H$\alpha$ are very small, and they do not appear to be related to the X-ray timescale; the overall lack of large variations in optical emission lines is consistent with the magnetic field constancy. In addition, variations of the radial velocities indicate that \tr\ is probably a SB1 binary with a very long period. } {Our new measurements of optical emission lines and magnetic field strength do not show an obvious correlation with X-ray variations. Our current data thus cannot be interpreted in terms of the common model, which assumes the electromagnetic emission associated with a wind confined by a dipolar field tilted with respect to the rotation axis. However, the sampling is imperfect and new data are needed to further constrain the actual periodicity of the various observed phenomena. If inconsistencies are confirmed, then we will need to consider alternative scenarios.} | Strong magnetic fields have been detected in a dozen O stars during the last decade \citep[and references therein]{pet13,fos15}. In such objects, the stellar wind flows are channelled towards the equator \citep{bab97}, creating a dense region of confined winds. Part of this material is shock-heated to higher temperatures, leading to the emission of X-rays, while the cooler plasma is easily detected in the visible (e.g. through Balmer emission lines). When the rotation axis is not perfectly aligned with the magnetic axis, a periodic modulation of the electromagnetic emission associated with these confined winds is observed. Variations of the longitudinal field, of the broad-band photometry, of the visible emission lines, and of the X-ray flux thus occur simultaneously, as was found for example for $\theta^1$\,Ori\,C \citep{don02,gag05} or HD\,191612 \citep{don06,naz07,how07,naz10}. \begin{table*} \centering \caption{List of the new X-ray observations. Phases were calculated considering $1/P$=0.01838\,d$^{-1}$ and $T_0$ as the date of the oldest \xmm\ observation (No. 4 in Table 1 of \citealt{naz14}, $JD$=2\,451\,751.707). } \label{list} \begin{tabular}{lcccc} \hline ObsID & exp. & Start Date & Associated JD & $\phi$\\ & time & & & \\ \hline \hline 0691970101 & 87\,ks & 2012-Dec-20@19:21:54 & 2456282.307 & 0.27 \\ 0742850301 & 13\,ks & 2014-Jun-06@19:13:05 & 2456815.301 & 0.07 \\ 0742850401 & 33\,ks & 2014-Jul-28@15:32:43 & 2456867.148 & 0.02 \\ 0762910401 & 11\,ks & 2015-Jul-16@01:18:44 & 2457219.555 & 0.50 \\ \hline \end{tabular} \end{table*} Spectropolarimetric observations found \tr\ (O8.5V) to be strongly magnetic \citep{naz12carina,naz14}. In the X-ray range, \tr\ displays a bright, variable, and hard emission atypical of single, ``normal'' massive stars \citep{eva04,ant08,naz11,naz14}, but in line with theoretical expectations for confined winds \citep{nazsurvey}. The Fourier analysis of the X-ray data clearly indicated the presence of a periodicity; the favoured value was $\sim$54d but many aliases were present, leaving some ambiguity on the actual period \citep{naz14}. Therefore, to better constrain its physical properties, we needed additional data of \tr. In this paper, we analyse these new data. Section 2 presents the observations used in the study, Sect. 3 provides the results, and Sect. 4 reports our conclusions. | Spectropolarimetric, X-ray, and high-resolution spectroscopic data of \tr\ have been obtained. They show a good agreement with previous results: the X-ray properties are in line with those already reported and the magnetic field values are similar (within errors) to previous measurements. However, there are also surprises. The optical spectroscopy indicates the presence of RV changes in \tr, suggesting that it is a long-term SB1, but the sparse sampling of the optical data prohibits us from deriving a full orbital solution. The probable binary nature of \tr\ cannot explain its exceptional X-ray emission, however. First, colliding-wind emission in the system could be seen as one possible source of X-ray variability, but (1) the long timescale of the RV variations may be difficult to reconcile with the shorter timescale detected in X-ray data and (2) such late-type O stars do not exhibit bright colliding-wind X-ray emission. Second, a low-mass (PMS) companion could be envisaged as the source of the X-ray variations because such objects display X-ray flares, but (1) the observed variations are large (about a factor of 2 corresponding to an increase in flux of $>10^{32}$\,erg\,s$^{-1}$), an extreme value for typical PMS flares, (2) PMS flares occur on relatively short timescales while the 2003 data indicate that a month was needed to return to ``normal'' flux levels, and (3) the combination of a long period and a large RV amplitude (at least 20\,\kms) of the O-type star rule out a companion with very low mass. Finally, it should also be noted that the average luminosity of \tr\ ($\log[L_{\rm X}]\sim32.3$) is similar to expectations from confined wind models (in particular, see the right panel of Fig. 6 in \citealt{nazsurvey} where \tr\ is \#8). On the other hand, neither the magnetic field values nor the optical spectroscopy display an obvious modulation with the favoured (54\,d) timescale derived from X-rays. Moreover, the measured magnetic field is compatible with a constant value, or at most with a low-amplitude modulation, while the optical line strengths also remain remarkably constant in our data. This near-constancy is reminiscent of HD\,148937 \citep{naz08,wad12,nazsurvey}. It would imply a system with a long period, always seen close to pole-on, or with magnetic and rotational axes aligned, but in these latter cases the X-ray emission of the confined winds would also remain stable, as significant occultation of the X-ray emitting regions by the stellar body would not occur, and this clearly contradicts the X-ray observations of \tr. However, many aliases were present in the X-ray flux periodogram, rendering the choice of the best period difficult and the new data did not fully resolve this ambiguity. The new data have thus brought additional questions. The variation timescales of the X-ray emission, optical emission line strength, and longitudinal magnetic field remain poorly constrained and -- worse -- their correlated behaviour has not yet been established, while it is clearly seen in all other magnetic O stars. The near-constancy of magnetic field values and optical emission strengths may even be difficult, if not impossible, to reconcile with the much larger variations detected at high energies. Only additional optical and X-ray observations, carefully scheduled and with high signal-to-noise ratios, would be able to clarify the situation and firmly confirm whether there is a mismatch for \tr\ with the usual oblique rotator scenario for confined winds in massive stars. | 16 | 9 | 1609.08806 |
1609 | 1609.09781_arXiv.txt | {For the last 100 years, General Relativity (GR) has taken over the gravitational theory mantle held by Newtonian Gravity for the previous 200 years. This article reviews the status of GR in terms of its self-consistency, completeness, and the evidence provided by observations, which have allowed GR to remain the champion of gravitational theories against several other classes of competing theories. We pay particular attention to the role of GR and gravity in cosmology, one of the areas in which one gravity dominates and new phenomena and effects challenge the orthodoxy. We also review other areas where there are likely conflicts pointing to the need to replace or revise GR to represent correctly observations and consistent theoretical framework. Observations have long been key both to the theoretical liveliness and viability of GR. We conclude with a discussion of the likely developments over the next 100 years.} \keyword{General~Relativity; gravitation; cosmology; Concordance~Model; dark energy; dark matter; inflation; large-scale~structure} \begin{document} | 16 | 9 | 1609.09781 |
||
1609 | 1609.02493_arXiv.txt | Magnetic fields (B-fields) play a key role in the formation and evolution of protoplanetary disks, but their properties are poorly understood due to the lack of observational constraints. Using CanariCam at the 10.4-m Gran Telescopio Canarias, we have mapped out the mid-infrared polarization of the protoplanetary disk around the Herbig Ae star AB Aur. We detect $\sim$0.44\% polarization at 10.3 $\micron$ from AB Aur's inner disk ($r<80$ AU), rising to $\sim$1.4\% at larger radii. Our simulations imply that the mid-infrared polarization of the inner disk arises from dichroic emission of elongated particles aligned in a disk B-field. The field is well ordered on a spatial scale commensurate with our resolution ($\sim$50 AU), and we infer a poloidal shape tilted from the rotational axis of the disk. The disk of AB Aur is optically thick at 10.3 $\micron$, so polarimetry at this wavelength is probing the B-field near the disk surface. Our observations therefore confirm that this layer, favored by some theoretical studies for developing magneto-rotational instability and its resultant viscosity, is indeed very likely to be magnetized. At radii beyond $\sim$80 AU, the mid-infrared polarization results primarily from scattering by dust grains with sizes up to $\sim$1 $\micron$, a size indicating both grain growth and, probably, turbulent lofting of the particles from the disk mid-plane. | Magnetic fields (B-fields) play an important role in star formation. They regulate the gravitational collapse and fragmentation of molecular cores, thus having a strong influence on the global star formation efficiency \citep{dullemond2007,crutcher2012,li2014ppvi}. It can be expected that large-scale B-fields can be dragged inward during core collapse and disk formation, leaving a remnant field in the resultant protoplanetary disk. For a weakly magnetized protoplanetary disk, magneto-hydrodynamic (MHD) turbulence arising from magneto-rotational instability (MRI) is thought to be the primary source of disk viscosity, a crucial driving force for disk evolution and planet formation \citep{balbus1998,turner2014}. Despite this consensus, observations that constrain B-field properties (geometry and strength) in protoplanetary disks are virtually non-existent. Dichroic emission and absorption of aligned elongated grains produce linear polarization that can trace the B-field morphology. In particular, polarimetric observations of dust thermal emission at centimeter or millimeter wavelengths with single-dish telescopes (e.g., CSO and JCMT) or interferometric arrays (e.g., JVLA, SMA, BIMA, and CARMA) have been used to map B-field structure in young stellar objects (YSOs) at scales from $\sim$50 to thousands of AU (see \citealt{crutcher2012} for a review). However, due to the limited sensitivity and angular resolution offered by current facilities, most of those studies have been focused on B-fields in molecular clumps and cores or Class 0-I objects \citep[e.g.,][]{qiu2013,zhang2014,davidson2014,segura-cox2014,rao2014,liu2016}, rather than classical protoplanetary (i.e., Class II) disks. Using CARMA, \citet{stephens2014} spatially resolved the HL Tau protoplanetary disk in polarized light at 1.3 mm. Their best-fit model was consistent with a highly tilted (by $\sim$50\degr\, from the disk plane) toroidal B-field threading the disk. However, this conclusion is challenged by recent follow-up studies, which show that the 1.3 mm polarization of HL Tau could also arise solely from dust scattering \citep{kataoka2015,yang2016}. Mid-infrared (mid-IR) polarimetry provides an alternative or complementary approach to the study of B-fields in YSOs and disks \citep{smith2000,barnes2015}. With 8-10-m telescopes, mid-IR observations can achieve 0\farcs3-0\farcs4 angular resolution in the 10-$\micron$ band under most observing conditions, sufficient to map out B-field structure in nearby disks at sub-disk (40-50 AU) scales. Protoplanetary disks are generally thought to be optically thick in the mid-IR out to hundreds of AU from the star \citep{chiang1997}. Hence, mid-IR polarimetry usually probes the emitting particles and B-field near the disk surface (also called the disk atmosphere) rather than its interior. This thin and warm surface layer and layers immediately adjacent to it are a potentially important channel for accretion and angular momentum transfer, since the disk mid-plane at the same radius may be too cold and too well shielded from ionizing radiation to enable MRI (i.e., the ``dead zone'', \citealt{gammie1996}). To gain new insight into B-fields in protoplanetary disks, we observed AB Aur (HD 31293, MWC 93) with CanariCam \citep{telesco2005,packham2005}, the facility mid-IR camera of the 10.4-m Gran Telescopio Canarias (GTC). AB Aur is an archetypal Herbig Ae star (i.e., intermediate-mass pre-main-sequence stars of 2-4 $M_{\odot}$) at the distance of 144 pc \citep{dewarf2003}. At 4$\pm$1 Myr old, this source still shows evidence of significant accretion ($\sim$10$^{-7}$ $M_{\odot}$ yr$^{-1}$; \citealt{dewarf2003,tang2012}). AB Aur is surrounded by a prominent disk, with mid-IR and 1.3 mm dust emission detected out to $\sim$280 AU and CO line emission detected out to $\sim$500 AU from the star \citep{marinas2006,tang2012}. In both CO and near-IR scattered-light images, the disk is rich in morphological features such as spiral arms and gaps, suggesting a dynamical disk environment and, perhaps, on-going planet formation \citep{pietu2005,hashimoto2011,tang2012}. Previous observations at various wavelengths gave a fairly consistent disk inclination of 27\degr\, (where 0\degr\, corresponds to pole-on), with the major axis of the disk oriented at position angle (P.A.) of 70\degr\, (measured E from N) \citep{pietu2005,tang2012,rodriguez2014}. $H$-band (1.6 $\micron$) polarization of the AB Aur disk has been imaged by \citet{hashimoto2011}, showing a clear centrosymmetric pattern indicative of scattering, as expected at these short IR wavelengths. The paper is organized as follows. Section \ref{sec:observations} describes our data acquisition and reduction, with results presented in Section \ref{sec:results}. Disk models are presented in Section \ref{sec:analysis}. The implications of our study are discussed in Section \ref{sec:discussion}, with our findings summarized in Section \ref{sec:conclusions}. | \label{sec:conclusions} We present GTC/CanariCam mid-IR (10.3 $\micron$) polarimetric observations with $\sim$0\farcs5 (50 AU) angular resolution of the protoplanetary disk of AB Aur to explore its magnetic field. The key findings are summarized below: \begin{enumerate} \item Linear polarization is detected from the disk of AB Aur out to $r\approx$ 1\farcs2 (170 AU). The polarization map shows two distinct regions, which we call the inner disk and the outer disk. Polarization vectors in the inner disk ($r<70$ AU) are approximately parallel to each other, whereas those in the outer disk (70 $< r <$ 170 AU) form a clearly centrosymmetric pattern. The (azimuthally averaged) degree of polarization increases from 0.44$\pm$0.05\% near the star to 1.4$\pm$0.4\% at 170 AU. \item We modeled the observations using RADMC-3D with customized code to include polarization from emission and absorption by aligned elongated dust grains. Our results show that the observed polarization is well reproduced when both polarized emission and polarization from scattering are included in the model. In the best-fit model, the disk of AB Aur is threaded by a poloidal field tilted from the spin axis of the disk by $\sim$30\degr-40\degr. Polarization of the inner disk is dominated by dichroic emission from elongated grains aligned in the B-field. In contrast, polarization of the outer disk is largely due to scattering. \item The disk of AB Aur is almost certainly optically thick at 10.3 $\micron$, so mid-IR polarimetry probes emitting dust grains and the B-field in the warm disk surface. Our observations imply that this surface layer is indeed magnetized, a crucial condition for MRI to operate. Furthermore, we estimate that, for both RAT and MRI to operate, the field strength on the disk surface should be of order 1-10 mG. \item A poloidal B-field tilted relative to the disk spin axis supports theories requiring such a misalignment to mitigate the ``magnetic breaking catastrophe'' \citep[e.g.,][]{hennebelle2009,joos2012}. It also ensures a considerable vertical component of the field (i.e., perpendicular to the disk plane), which is needed to create a sufficiently high accretion rate through MRI-driven turbulence as suggested by observations \citep{simon2013,simon2015}. \item Significant polarization arising from scattering in the outer disk of AB Aur requires micron-sized grains near the disk surface, indicating grain growth and possible lofting of these particles to the disk surface by turbulence. \end{enumerate} Our study of AB Aur is the first to probe B-fields in a protoplanetary disk with mid-IR polarimetry, and it demonstrates the potential of this technique. While our observations provide critical boundary conditions that must be satisfied by the B-field interior to the disk of AB Aur, that interior B-field geometry remains otherwise undefined. Other observing facilities, such as the Atacama Large Millimeter Array (ALMA) and Very Large Array (VLA) observing at sub-millimeter and centimeter wavelengths, will permit probing of disk interiors all the way to the disk mid-plane, observations that will strongly complement those in the mid-IR. | 16 | 9 | 1609.02493 |
1609 | 1609.00175_arXiv.txt | In the centre of our galaxy lies a super-massive black hole, identified with the radio source Sagittarius $A^\star$. This black hole has an estimated mass of around 4 million solar masses. Although Sagittarius $A^\star$ is quite dim in terms of total radiated energy, having a luminosity that is a factor of $10^{10}$ lower than its Eddington luminosity, there is now compelling evidence that this source was far brighter in the past. Evidence derived from the detection of reflected X-ray emission from the giant molecular clouds in the galactic centre region. However, the interpretation of the reflected emission spectra cannot be done correctly without detailed modelling of the reflection process. Attempts to do so can lead to an incorrect interpretation of the data. In this paper we present the results of a Monte Carlo simulation code we developed in order to fully model the complex processes involved in the emerging reflection spectra. The simulated spectra can be compared to real data in order to derive model parameters and constrain the past activity of the black hole. In particular we apply our code to observations of Sgr B2, in order to constrain the position and density of the cloud and the incident luminosity of the central source. The results of the code have been adapted to be used in Xspec by a large community of astronomers. | \label{introduction} In the centre of our Galaxy lies a super-massive black hole, identified with the bright radio source Sagittarius $A^\star$ (Sgr $A^\star$). The black hole has an estimated mass of ~4 million solar masses \citep{gillessen09}, and currently it is very faint with an X-Ray luminosity of only 10$^{33}$ - 10$^{34}$ erg/s \citep{baganoff03}, orders of magnitude dimmer than other Active Galactic Nuclei (AGN). During times of flaring the luminosity of Sgr $A^\star$ is known to increase by a factor of 100, but even this luminosity is still 10$^8$ lower than the Eddington luminosity for a black hole of this mass \citep{baganoff01}. However it is possible that previously Sgr $A^\star$ was much brighter than it is now. The X-ray emission from the Central Molecular Zone (CMZ) was first interpreted as a reflection of a past flaring of Sgr $A^\star$ by \cite{sunyaev93, koyama96}. Since then several studies have confirmed this hypothesis. Indeed for the past decade XMM-Newton has been regularly monitoring the the CMZ, along with other X-ray satellites, such as \textit{INTEGRAL}, \textit{Suzaku}, \textit{Chandra} and \textit{NuSTAR}. These observations have shown that the X-ray emission from some of the molecular clouds varies significantly with time \citep{muno07, inui09, ponti10, terrier10, clavel13, clavel14, zhang15}. For example the most massive cloud in the region, Sgr B2 has shown a relatively constant X-ray flux up to 2003, and since then has been showing a large flux decay, while another cloud, known as "The Bridge", has shown an increase in intensity. Some clouds in the region show constant X-ray emission, and others display hardly any X-ray emission at all. This non-uniform behaviour disfavours the interpretation of the X-ray emission as a result of accelerated particle interactions. The hypothesis that the illuminating source be Sgr $A^\star$ is supported by the large luminosity ( $> 10^{39}$ erg/s) needed to explain the observed flux of the Giant Molecular Clouds (GMC) \citep{ponti13}. The morphological evidence for a reflection origin of the emission is also compelling. The overall symmetry around the source of the Fe $K_\alpha$ emission (observed from GMC with both positive and negative galactic longitudes) along with the fact that the emission from Sgr B2 is shifted toward the galactic centre by about ~2 arcmin when compared to the molecular mass distribution, suggesting the illuminating source is in the direction of the Galactic Centre (See \cite{murakami01, terrier10, ponti13} for an in-depth look at the work in this area to date). Sgr B2 is a good candidate for observing the reflected X-rays of Sgr $A^\star$. Sgr B2 is the largest, most massive and dense cloud of the CMZ \citep{protheroe08}. Sgr B2 has been studied as an X-Ray Reflection Nebula (XRN) by several authors using ASCA, \textit{Chandra}, \textit{INTEGRAL}, \textit{XMM-Newton}, \textit{Suzaku} and \textit{NuSTAR} observations \citep{koyama96, murakami01, revnivtsev04, koyama07, terrier10, ponti13, zhang15} and the general interpretation is a reflection nebula generated by a Sgr $A^\star$ outburst occurring between 100 and 300 years past. However to understand the timing, durations and intensities of any outbursts from Sgr $A^\star$ that are being reflected by Sgr B2, it is necessary to know the position of Sgr B2 relative to Sgr $A^\star$ \citep{ponti13}. The latter is not yet well constrained. \cite{reid09} have measured the parallax of Sgr B2 and Sgr $A^\star$ with VLBI observations and deduced that the former should lie $\approx 120 $ pc in front of the latter assuming circular rotation. Conversely, in the \cite{molinari11} model, it lies far behind. The improved model of \cite{kruijssen15}, which takes into account orbital motion in the GC, gives a line of sight position of \textbf{$\approx$38 pc} in front of Sgr $A^\star$. Many analyses of the X-ray reflection phenomenon in the GC have relied on the X-ray spectrum only to extract information on the cloud thickness and location and deduce constraints on the illuminating source luminosity and spectrum. The spectrum of an X-ray nebula is often assumed to follow a simple absorbed power-law shape. In some cases, specific models of reflection spectra are used such as {\tt pexrav} \citep{ponti10} or {\tt MyTorus} \citep{zhang15, mori15}. However, we argue that these models do not reproduce precisely the spectral shape of an appropriate XRN, because they do not model the geometry of an isolated illuminated cloud. Specifically, the {\tt pexrav} model is a cold disk viewed from a given angle, is only valid in compton thick cases, does not re-produce the iron line, and does not provide a determination of the cloud column density $N_H$. Both \cite{zhang15} and \cite{mori15} have argued that the {\tt MyTorus} model is valid and applicable. However, there are some notable shortcomings of using the {\tt MyTorus} model, which are outlined in the appendices of both \cite{zhang15} and \cite{mori15}. Primarily {\tt MyTorus} only deals with uniform density, has a fixed iron abundance and there is also the geometry of {\tt MyTorus}, with scattering from the far side of the torus being an issue. This issue can be critical when the spectral determination is used to constrain the geometry of the reflection, namely the position of the cloud with respect to the illuminating source. For example, some authors have used the fitted $\mathrm{N_H}$ of the absorbed power-law to estimate the column density of the cloud and hence its location along the line of sight using partial covering of the plasma emissions in the CMZ \citep{ryu09}. Others have used an estimate of the equivalent width (EW) of the 6.4 keV line to place the cloud along the line of sight \citep{capelli12}. Both approaches are sensitive on the determination of the underlying continuum level and require correct models of the scattered X-ray spectrum. To address this issue it is necessary to produce a reliable model of reflected X-rays from an isolated cloud that can be properly applied to X-ray observations (i.e. through forward folding methods). Reflection spectra for Compton thin constant density spheres can be calculated analytically (See Appendix A). Indeed in section \ref{code_output} we present the results from these semi-analytical calculations. However these calculations are only dealing with Klein-Nishina scattering and are not taking into account binding modified multiple scattering (as such they are invalid in the Compton thick case), complex geometries or non-uniform densities. To this end we create a Monte Carlo (MC) code to simulate X-ray reflection spectra from molecular clouds and use it to constrain such properties of Sgr B2 as the dense central core column density $N_H$ , the photon index $\Gamma$ of the incident X-ray emission, the incident Sgr $A^\star$ luminosity and the position of Sgr B2. We show that the position of a molecular cloud has a strong influence both on the shape of the continuum and the relative strength of the 6.4 keV neutral iron line. Using an Xspec table model, we can compare model spectra to existing data and place constraints on the location of the Sgr B2 molecular cloud. Other authors have developed Monte Carlo models of the reflection process in cold molecular material, for example the early work of \cite{sunyaev98, murakami01}. However these works were subject to large approximations such as a fixed photon index and isotropic scattering as well as fixed geometry (in particuler the viewing angle). The work of \cite{odaka11} is free from such approximations and is comparable to our own. Their work concentrated more on the reflection morphology, while we concentrate on the spectral features and produce the Xspec table model to allow for fitting to real observational data. There is also the work by \cite{molaro16} which concentrated on the effects of clumpiness in the cloud rather than on the spectral features. This paper is organised as follows. In Section \ref{the_mc_code} we describe the details of our MC code, then in Section \ref{code_output} we discuss the dependence of the reflected component on the parameters of the model. In Section \ref{observations} we present the data obtained by \textit{Chandra}, \textit{XMM-Newton} and \textit{Integral} of Sgr B2, while in Section \ref{sgrb2} we apply the model to this data. Finally we compare our findings with previous work and present our conclusions in Section \ref{discussion}. | \label{discussion} In this paper we presented a new Monte Carlo code for simulating X-ray reflection spectra from molecular clouds. The code is capable of modelling clouds of varying density distributions and of varying geometries. We show how different input parameters will result in highly divergent output spectra and provide an analysis of the processes behind these changes. Using \textit{Chandra}, \textit{XMM-Newton} and \textit{INTEGRAL} observations we constrain several parameters of the giant molecular cloud Sgr B2 using Xspec table models: the photon index $\Gamma$ of the incident spectrum, the luminosity of the illuminating source, the $N_H$ of the reflecting cloud and the angular position of the cloud relative to the line of sight. More recent observations (such as \textit{NuSTAR}) of Sgr B2 could not be used directly with this model since it assumes the cloud is at or close to full illumination, and the hard X-ray flux has decreased by 40\% from 2003 to 2010 \citep{terrier10} with a time constant consistent with the light crossing time of the cloud. Given the cloud optical thickness, this is mostly due to multiply scattered photons \citep{sunyaev98}, but the main illuminating front has already started to leave the cloud. The relative brightness of the densest cores observed by \textit{NuSTAR} in hard X-rays \citep{zhang15} also supports this. Time dependent effects have therefore to be properly taken into account to model the spectra observed at later times. However, for the time interval between our observations we do not expect there to be any significant change in spectral shape. Figure \ref{fig:lightcurve} shows the light curve of observed photons in the 1-5 keV energy range for a Compton thick ($N_H = 4\times10^{24} cm^{-2}$) cloud for various line of sight positions. There is expected to be an extremely rapid decay in soft X-Ray flux after the cessation of the incident flare. The decay rate is best described by a power law, particularly in the first 20 months. Given the low characteristic time scale in the low energy, the observed time behaviour is likely close to that of the illuminating source. As there is an order of magnitude decay in the first year after flare cessation and we observe a rather marginal flux variation in 2004 compared to 2000, we can infer we are seeing a flare that lasted at least 4 years or close to it. In this case the spectral shape will not change between observations, assuming the incident spectrum remains the same. \begin{figure} \includegraphics[angle =-90, scale=0.35]{lightcurve_1_5_lt.eps} \caption{Light curve of escaping photons in the energy range 1 - 5 keV, for a Compton thick ($N_H = 4\times10^{24} cm^{-2}$) cloud. With various line of sight angle positions shown.} \label{fig:lightcurve} \end{figure} Although the actual density profile of Sgr B2 is most certainly not uniform \citep{etxaluze13, jones11, protheroe08} it can be seen in Figure \ref{fig:profile_comparison} that differences between density profiles (excluding the exponential) are quite small, in fact it is unlikely that current X-ray data is capable of distinguishing the density profile. For this reason we presented both uniform Density and Gaussian density profile results. Both fit the data reasonably well. Importantly, they predict similar parameters. The Gaussian fit angle is slightly high $\approx90^{\circ}$ which is at odds with the literature but not completely unreasonable. However, the uniform density fit is technically superior with a $\chi^2_{red} = 1.18$ as opposed to $\chi^2_{red} = 1.32$ for the Gaussian. This poor Gaussian fit is perhaps due to the small centralized nature of the observations which precludes treating the observation as that of a cloud with a Gaussian distribution given the large size of the cloud. The column density predicted by both models ($N_{H_c} = 2.13\times10^{24} cm^{-2})$ and ($N_{H_G} = 2.5\times10^{24} cm^{-2}$) respectively, is higher than that obtained with X-ray observations previously, most likely due to the improper XRN models used. However, it is in good agreement with radio observations \citep{jones11, protheroe08}. The estimated incident luminosity on the order $10^{39}$ erg/s is in good agreement with the literature \citep{ponti13, zhang15}. There is a small decrease in Flux/Luminosity between the two periods, indicating a source luminosity decrease. The photon index found by both models $\Gamma_C = 2.16_{-0.16}^{+0.12}$ and $\Gamma_G = 1.83_{-0.11}^{+0.06}$ respectively, is also in good agreement with previous estimates \citep{revnivtsev04, terrier10, ponti13, zhang15}. This confirms that the illuminating source must have been a previous high activity period of Sgr $A^\star$. The fitted angle of the uniform density fit $\theta = 64_{-7}^{+8}$ degrees gives a Sgr $A^\star$ - Sgr B2 distance of $111^{+8}_{-6}$ pc (assuming a projected distance of $100$ pc) and places Sgr B2 $48^{+2}_{-1}$ pc closer to Earth than Sgr $A^\star$ in nominal agreement with \cite{kruijssen15} who give $\approx 38$ pc. Whereas the angle for the Gaussian fit $\theta = 89_{-11}^{+10}$ places the cloud at it's projected distance of 100 pc. The ultimate aim of this fitting is to facilitate the analysis of GMC's in the Galactic Centre Region as reflectors of previous outbursts or periods of higher activity from Sgr $A^\star$. We make available these Xspec table models on Zenodo\footnote{\url{http://dx.doi.org/10.5281/zenodo.60229}}, we note that Iron abundance is not a parameter within the model, this was a choice to keep the number of fitting parameters low, however table models with other iron abundances are also available. We anticipate the use of this code in further work, fitting to other broadband X-Ray data (NuSTAR) from various GMC's. Helping to determine their characteristics and primarily their line of sight positions, thus offering improved constraints on the time delay and duration of the illuminating events. | 16 | 9 | 1609.00175 |
1609 | 1609.07080_arXiv.txt | It is known that unconfined dust explosions typically starts off with a relatively weak primary flame followed by a severe secondary explosion. We show that clustering of dust particles in a temperature stratified turbulent flow ahead of the primary flame may give rise to a significant increase in the radiation penetration length. These particle clusters, even far ahead of the flame, are sufficiently exposed and heated by the radiation from the flame to become ignition kernels capable to ignite a large volume of fuel-air mixtures. This efficiently increases the total flame surface area and the effective combustion speed, defined as the rate of reactant consumption of a given volume. We show that this mechanism explains the high rate of combustion and overpressures required to account for the observed level of damage in unconfined dust explosions, e.g., at the 2005 Buncefield vapor-cloud explosion. The effect of the strong increase of radiation transparency due to turbulent clustering of particles goes beyond the state-of-the-art of the application to dust explosions and has many implications in atmospheric physics and astrophysics. | As is known, dust explosions can occur when an accidentally ignited flame propagates through a cloud of fine particles suspended in a combustible gas. The danger of dust explosions have been known for centuries in the mining industry and grain elevators, and it is currently a permanent threat in all those industries in which powders of fine particles are involved. A significant progress has been made in the development of prevention and safety measures (see, e.g., \cite{Eckhoff2003,AmyotteEckhoff2010,Proust2006,AbbasiAbbasi2007,Yuan2015}, and references therein). However, despite considerable efforts over more than 100 years, the mechanism of flame propagation in unconfined (large scale) dust explosions still remains the main unresolved issue. Based on the analysis of many catastrophic accidents, it is currently well established that unconfined dust explosions consist of a relatively weak primary explosion followed by a severe secondary explosion. While the hazardous effect of the primary explosion is relatively small, the secondary explosion may propagate with a speed of up to 1000 m/s, producing overpressures of over 8-10 atm, which is comparable to the pressures produced by a detonation. However, the analysis of damages indicates that a detonation is not involved, while normal deflagrations are not capable of producing such high velocities and overpressures. It is interesting to note that M. Faraday and C. Lyell, who in 1845 analyzed the Haswell coal mine explosion, where the first to point to the likely key role of dust particles \cite{Faraday1845}. Of special interest for understanding the nature of dust explosions is the Buncefield vapor-cloud explosion, for which a large amount of data and evidence were collected, providing unique and valuable information about the timing and damage caused by the event \cite{Buncefield2009}. Later, a comprehensive analysis was performed by two groups \cite{Atkinson-Cusco2011,Bradeley2012} in an attempt to understand possible mechanisms leading to the unusually high rate of combustion. Among different mechanisms, the hypothesis proposed by Moore and Weinberg \cite{Moore81,Moore83}, where they argue that the anomalously high rate of flame propagation in dust explosions can be due to the radiative ignition of millimeter sized fibrous particles ahead of the flame front, attracted much attention. In normal practice, emissivity of combustion products and radiation absorption in a fresh unburnt gaseous mixture are small and do not influence the flame propagation. The situation is drastically changed for flames propagating through a cloud of fine particles suspended in a gaseous mixture. In this case, the radiative flux emitted from the primary flame ($\approx$ 300 - 400 kW/m$^{2}$), is close to the black-body radiation \cite{Hadjipanayis2015} at stoichiometric flame temperatures (2200 - 2500~K). In dust explosions, the radiative flux into the reactants is significantly enhanced by the increased emissivity of the large volume of burned products. Microns-size dust particles, with diameter $d_{p}$ that is larger than the radiation wavelength $\lambda_{\rm rad}$, efficiently absorb radiation emitted by the flame. The absorbed heat is transferred to the surrounding gas by thermal conduction, so that the gas temperature lags behind that of the particle. If particles are evenly distributed, the maximum temperature increase ahead of the flame is \cite{Liberman2015}: \begin{eqnarray} \Delta T_p \approx \frac{0.63 \, \sigma T_b^4}{U_{\rm f}(\rho_p \, c_{\rm p} + \rho_g \, c_{\rm v})}, \label{AAA1} \end{eqnarray} where $U_{\rm f}$ is the normal flame velocity, $\sigma T_b^4$ is the black-body radiative flux, and $\rho_p$, $\rho_g$, $c_{\rm p}$, $c_{\rm v}$ are the mass density of particles, the gas mixture and their specific heats, respectively. As a flame propagates through a dust cloud of evenly dispersed particles it consumes the unburned fuel before the gas temperature will have risen up to ignition level, so that radiation can not become a dominant process of the heat transfer \cite{Liberman2015}. Instead of being uniformly distributed, the particles may be organized in the form of a distant optically thick layer or a clump ahead of the flame front. It was shown \cite{Liberman15} that if a transparent gap between the particle layer and the flame is not too small, the particle layer can be sufficiently heated by the flame-emanated radiation to ignite new combustion modes in the surrounding fuel-air mixture ahead of the main flame. In turbulent flows ahead of the primary flame (${\rm Re}\approx 10^{4}$ - $10^{5}$), typical for dust explosions, dust particles with $\rho_p \gg \rho_g$ assemble in small clusters of the order of the Kolmogorov turbulent scale, $\ell_\eta \approx $0.1 - 1~mm. The turbulent eddies, acting as small centrifuges, push the particles to the regions between the eddies where the pressure fluctuations are maximum and the vorticity intensity is minimum. The particles assembled in these regions make clusters with higher particle number densities than the mean number density. This effect, known as inertial clustering, has been investigated in a number of analytical, numerical and experimental studies \cite{TB09,W09,BE10,EKR96,EKLR02,EKLR07,BB07,SA08,XB08,SSA08,SBB14}. Recent analytical studies \cite{EKLR13,EKLR15} and laboratory experiments \cite{EKR10} have shown that the particle clustering can be much more effective in the presence of a mean temperature gradient. In this case, the turbulence is temperature stratified, and the turbulent heat flux is not zero. This causes correlations between fluctuations of fluid temperature and velocity, and, therefore, correlations between fluctuations of pressure and fluid velocity, thus, producing additional pressure fluctuations caused by the tangling of the mean temperature gradient by the velocity fluctuations. This enhance the particle clustering in the regions of maximum pressure fluctuations. As a result, the particle concentration in clusters rises by a few orders of magnitude compared to the mean concentration of evenly dispersed particles \cite{EKLR13}. In order for a secondary dust explosion followed by a shock wave to occur, the pressure from the mixture ahead of the flame ignited by the radiatively heated particle clusters must rise faster than it can be equalized by sound waves. This means that the penetration length of radiation, $L_{\rm rad}$, in the dust cloud must be sufficiently large to ensure that the clusters of particles even far ahead of the flame are sufficiently exposed and heated by the radiation, to become ignition kernels, i.e., $\tau_{\rm ign} \, c_{\rm s} \ll L_{\rm rad}$. Here $c_{\rm s}$ is the sound speed in the mixture ahead of the flame, and $\tau_{\rm ign}$ is the characteristic time of fuel-air ignition by the radiatively heated particle clusters (ignition kernels). The effect of spatial inhomogeneities with the scales larger than the wavelength of the radiation has been discussed in \cite{K93,R15b,K89} and studied using Monte Carlo modeling \cite{Farbar16,Frankel16} of the radiative heat transfer in particle-laden flow, taking into account the inertial particle clustering in non-stratified turbulence. In this paper we show that clustering of particles ahead of the primary flame gives rise to a strong increase of the radiation penetration (absorption) length. The effect ensures that clusters of particles even far ahead of the primary flame are sufficiently exposed and heated by the radiation from the flame to become ignition kernels, and the condition, $\tau_{\rm ign} \, c_{\rm s} \ll L_{\rm rad}$, is satisfied. The multiple radiation-induced ignitions in many ignition kernels ahead of the primary flame increase the total flame surface area, so that the distance, which each flame has to cover for a complete burn-out of the fuel, is substantially reduced. It results in a strong increase of the effective combustion speed, defined as the rate of reactant consumption of a given volume. The proposed mechanism of unconfined dust explosions explains the physics behind the secondary explosion. It also explains the anomalously high rate of combustion and overpressures required to account for the observed level of damages in unconfined dust explosions. | It is shown that clustering of particles in the temperature stratified turbulent flow ahead of the primary flame gives rise to a strong increase of the radiation penetration length (Fig.~\ref{Fig1}), but within a narrow interval of turbulent parameters (see Fig.~\ref{Fig2}). This explains the mechanism of the secondary explosion in unconfined dust explosions. According to analysis of the Buncefield explosion \cite{Atkinson-Cusco2011}: ``The high overpressures in the cloud and low average rate of flame advance can be reconciled if the rate of flame advance was episodic, with periods of very rapid combustion being punctuated by pauses when the flame advanced very slowly.'' Assuming $L_{a}\sim 10$ cm, we obtain $L_{\rm eff}\sim 10$ m as the result of the turbulent clustering of dust particles. Given that the radiative flux from the primary flame is about $S\sim$ 300 - 400 kW/m$^2$, the temperature of particles increases up to the ignition level $\Delta T_p\approx 1000$~K during ${\tau}_{\rm ign}= (\rho_p \, d_p \, c_{\rm p} \, \Delta T_p)/2S < 10$ ms. Time-scales of the fuel-air ignition by the radiatively heated cluster of particles (with $d_p$ in the range from 1 to $20 \, \mu$m) can be smaller than 10 ms \cite{Liberman2015,R7}. This ensures that the condition $L_{\rm eff} \gg {c_{\rm s}} \, \tau_{\rm ign}$ is satisfied, and the pressure produced by the ignition of the fuel-air mixture rises until the pressure wave steepen into a shock wave. The effective rate of the secondary explosion propagation is $\sim 1000$ m/s, which corresponds to a Mach number, ${\rm Ma}=$2.5 - 3, and overpressures of 8-10 atm. This explains the level of damage observed in the aftermath of unconfined dust explosions \cite{Buncefield2009,Atkinson-Cusco2011}. At the same time, the secondary explosion strongly change parameters of the turbulent flow. As a result, rapid combustion in secondary explosion is interrupted until the shock wave run away, and the slow combustion parameters of the turbulent flow ahead of a slow deflagration wave will be suitable for formation of a new transparent radiation window. This agrees fairly well with the Buncefield explosion scenario \cite{Atkinson-Cusco2011}. The effect of the strong increase of the radiation penetration length due to turbulent particle clustering goes beyond to the dust explosion applications and has many implications in astrophysical and atmospheric turbulence \cite{OS16,SOC12,EKLR15}. | 16 | 9 | 1609.07080 |
1609 | 1609.02941_arXiv.txt | Fluctuations of the surface brightness of cosmic X-ray background (CXB) carry unique information about faint and low luminosity source populations, which is inaccessible for conventional large-scale structure (LSS) studies based on resolved sources. We used \chandra{} data of the XBOOTES field ($\sim9$~deg$^2$) to conduct the most accurate measurement to date of the power spectrum of fluctuations of the unresolved CXB on the angular scales of $3\arcsec-17\arcmin$. We find that at sub-arcmin angular scales, the power spectrum is consistent with the AGN shot noise, without much need for any significant contribution from their one-halo term. This is consistent with the theoretical expectation that low-luminosity AGN reside alone in their dark matter halos. However, at larger angular scales we detect a significant LSS signal above the AGN shot noise. Its power spectrum, obtained after subtracting the AGN shot noise, follows a power law with the slope of $-0.8\pm0.1$ and its amplitude is much larger than what can be plausibly explained by the two-halo term of AGN. We demonstrate that the detected LSS signal is produced by unresolved clusters and groups of galaxies. For the flux limit of the XBOOTES survey, their flux-weighted mean redshift equals $\left<z\right>\sim0.3$, and the mean temperature of their intracluster medium (ICM), $\left<T\right>\approx 1.4$~keV, corresponds to the mass of $M_{500}\sim 10^{13.5}~M_\odot$. The power spectrum of CXB fluctuations carries information about the redshift distribution of these objects and the spatial structure of their ICM on the linear scales of up to $\sim$Mpc, i.e. of the order of the virial radius. | \label{sec:intro} Since the discovery of the cosmic X-ray background (CXB) about half a century ago \citep{Giacconi1962}, understanding of its origin has been one of the major drivers for the development of X-ray astronomy and most X-ray space telescopes, such as the currently active missions: \xmm{}, \chandra{}, and NuSTAR \citep[e.g.][]{Fabian1992,Giacconi2013,Tanaka2013}. Thanks to the many, in particular deep X-ray surveys of \chandra{} \citep[e.g.][]{Brandt2005,Alexander2013,Brandt2015}, we now know for certain that the CXB is dominated by extragalactic discrete sources, with Active Galactic Nuclei (AGN) leading the way \citep[e.g.][]{Comastri1995,Moretti2003,Hickox2006,Hickox2007,Gilli2007,Moretti2012,Lehmer2012}. This makes the CXB the prefect window to study the accretion history of the Universe up to high redshift ($z\sim5$) \citep[e.g.][]{Hasinger2005,Gilli2007,Aird2010,Ueda2014,Miyaji2015}, which is an essential base to understand galaxy evolution \citep[e.g.][]{Hopkins2006,Hickox2009,Alexander2012}. Since the first X-ray surveys, angular correlation studies of the CXB had two major applications. They are used to disentangle the components of the CXB and at the same time to perform large-scale structure (LSS) studies \citep[e.g.][]{Scheuer1974,Hamilton1987,Shafer1983,Barcons1988,Soltan1994,Vikhlinin1995b,Miyaji2002}. The advantage of such studies is that one can analyze the CXB beyond the survey sensitivity limit, since one does not require any source identification or/and redshift information. Thanks to these studies it has been long known that the CXB must be dominated by point-sources with a redshift distribution similar to optical QSOs but somewhat higher clustering strength. These results were confirmed in the last $\sim$two decades by very deep pencil beam surveys \citep[e.g.][]{Hickox2006,Hickox2007,Lehmer2012} and LSS studies with resolved samples of X-ray-selected AGN from wide but more shallow surveys \citep[see reviews of][]{Cappelluti2012,Krumpe2013}. This became possible thanks to the high-angular resolution of the current generation of X-ray telescopes, complemented with optical spectroscopic redshift surveys of sufficient size and depth. Due to this clustering measurements with resolved AGN developed in the last decade to an important branch of LSS studies in general. It led to major advances in understanding how AGN activity is triggered and how does it depend on its environment, such as the host galaxy and dark matter halo (DMH) properties, and how do supermassive black holes (SMBH) grow and co-evolve with their DMH over cosmic time, which are essential questions in the field of galaxy evolution \citep[e.g.][]{Cappelluti2012,Krumpe2013}. In the future, it will become possible to use AGN as a cosmological probe via baryon acoustic oscillation measurements \citep[for details see e.g.][]{Kolodzig2013,Huetsi2013} with the $\sim3$~million AGN to be detected in the upcoming SRG/eROSITA all-sky survey \citep[for details see][]{eROSITA,eROSITA.SB,Kolodzig2012}. Due the focus on resolved AGN, the current knowledge of AGN clustering properties and its implications for AGN and galaxy evolution are biased towards objects of $L_{\eSoft}>10^{42}\,\mathrm{erg\;s^{-1}}$, in particular for higher redshifts ($z>0.5$), due to the luminosity cut from the AGN identification process and the signal-to-noise ratio (S/N) cut for the spectroscopic redshift \citep[e.g.][]{Allevato2011,Allevato2012,Allevato2014,Krumpe2010,Krumpe2012,Miyaji2011,Krumpe2015}. An important question to ask is if we are able to extrapolate these clustering properties to less luminous AGN, which trace galaxies at an earlier evolutionary stage with a less massive SMBH and/or smaller accretion rate than luminous AGN? % A significant step towards answering this question is to study the surface brightness fluctuations of the unresolved CXB measured with the current generation of X-ray telescopes, which allows us to measure angular fluctuations on small scales down to the arc-second regime. This type of clustering measurement offers us a great window to the small-scale clustering regime ($<1\;\mathrm{Mpc\,h^{-1}}$). Clustering studies of spatially resolved AGN samples have difficulties to access this regime, because of the low spatial density of AGN in general, and because multiobject spectroscopy surveys are typically limited to an angular separation of $\sim1\arcmin$ \citep[e.g.][]{Blanton2003,Dawson2013}. Therefore, our best spatially resolved measurement of the small-scale clustering regime comes from the use of dedicated catalogs of close AGN pairs \citep[e.g. SDSS Quasar Lens Search,][]{Kayo2012} or the direct measurement of the halo occupation distribution (HOD) of AGN from galaxy groups \citep[e.g.][]{Allevato2012}. Both types of measurement require an extensive amount of multi-wavelength survey data. In terms of standard clustering studies, the best results come from studies of optically-selected AGN thanks to the sufficient size of the available survey data \citep[e.g.][]{Hennawi2006,Kayo2012,Richardson2012,Shen2012}. For X-ray-selected AGN, the situation is more difficult due to the so far rather limited survey data \citep[e.g.][]{Allevato2012,Richardson2013}. Here, non-spatially resolved studies, such as the brightness fluctuations of the unresolved CXB, may offer a true alternative for small-scale clustering measurements, which have not been fully utilized yet. Due to their scientific focus, the only two existing studies of the brightness fluctuations of the unresolved CXB at these angular scales used very deep surveys \citep[e.g.][]{Cappelluti2013,Helgason2014}. However, this also implies a very small sky coverage of these surveys ($\sim0.1\,\mathrm{deg^2}$). In our study, we aim to conduct the most accurate measurement to date of the brightness fluctuations of the unresolved CXB on angular scales below $\sim17\arcmin$. We are able to achieve this by using the XBOOTES survey \citep[hereafter \Ken{}]{Murray2005,Kenter2005}, the currently largest available continuous \chandra{} survey, with a surface area of $\sim9\,\mathrm{deg^2}$. The advantage in comparison to previous studies is that a higher S/N makes any comparison with current clustering models from known source populations much more meaningful and it enables us to do clustering measurements in an energy resolved manner in order to separate different source populations. In this first study we present our measurement of the brightness fluctuations of the unresolved CXB with angular scales up to $\sim17\arcmin$, and make novel tests for systematic uncertainties such as the brightness fluctuations of the instrumental background. % Covering the angular scales from the arc-second to arc-minute regime may allow us to study the clustering properties of AGN within the same DMH (one-halo-term) and AGN of different DMHs (two-halo-term) of low-luminosity AGN ($L_{\eSoft}<10^{42}\,\mathrm{erg\;s^{-1}}$) and redshifts of $z>0.5$. This parameter regime is inaccessible for conventional clustering studies of the resolved CXB with current X-ray surveys \citep[e.g.][]{Cappelluti2012}. Diffuse emission from the intracluster medium (ICM) of \GCG{} and the associated warm-hot intergalactic medium (WHIM) also contributes to the CXB \citep[e.g.][]{Rosati2002,Hickox2007,Kravtsov2012,Roncarelli2012}. Since \GCG{} are more difficult to detect and an order of magnitude more sparse than AGN, our knowledge about their population, in particular at low fluxes ($\lesssim10^{-16}\,\mathrm{erg\,cm^{-2}\,s^{-1}}$) is less certain \citep[e.g.][]{Finoguenov2007,Finoguenov2010,Finoguenov2015,Clerc2012,Boehringer2014}. Thanks to cosmological hydrodynamical simulations \citep[e.g.][]{Roncarelli2006,Roncarelli2007,Roncarelli2012,Ursino2011,Ursino2014} and analytical studies \citep[e.g.][]{Diego2003,Cheng2004} we have nevertheless some reasonable understanding of their clustering properties. \Mc{As we will demonstrate in this paper, angular correlation studies of CXB fluctuations can help to dramatically improve the situation from the observational side.} This paper is organized as following: In \Sref{sec:DataProc} we explain our data processing procedure, in \Sref{sec:Espec} we show the energy spectrum of the unresolved CXB of XBOOTES and estimate the contribution by different components of the CXB, in \Sref{sec:Ana} we present our measurement of the surface brightness fluctuations of the unresolved CXB, and in \Sref{sec:Excess} we \Ac{study the origin of the detected LSS signal at large angular scales.} In the Appendixes we present results of tests \Mc{for various systematic effects and investigate the impact} of the instrumental background on our measurements. For the work we assume a flat $\Lambda$CDM cosmology with the following parameters: $H_0 = 70\,\mathrm{km\,s^{-1}\,Mpc^{-1}}$ ($h=0.70$), $\Omega_\mathrm{m} = 0.30$ ($\Omega_\Lambda = 0.70$), $\Omega_\mathrm{b} = 0.05$, % $\sigma_8=0.8$. The values for $H_0$ and $\Omega_\mathrm{m}$ were chosen to match the values assumed in the X-ray luminosity function studies, which we use in our calculations (e.g.\ \Sref{ss:UnresoAGN} or \ref{ss:AGNmodel}), and $\Omega_\mathrm{b}$ and $\sigma_8$ are derived from the cosmic microwave background (CMB) study of WMAP\footnote{\url{http://map.gsfc.nasa.gov}} \citep{Komatsu2011}. We note that the results of this work are not very sensitive to the exact values of the cosmological parameters and if we used the recently published, more precise cosmological parameters of the CMB study by PLANCK\footnote{\url{http://www.cosmos.esa.int/web/planck}} \citep{Planck_CosPara} our results would not change. | Surface brightness fluctuations of the unresolved CXB present a great opportunity to study faint source populations which are yet beyond the reach of more conventional studies of resolved sources. The renaissance of this field was facilitated by the fact that wide angle X-rays surveys covering tens of deg$^2$ have been undertaken by the modern X-ray telescopes aboard \chandra{} and XMM-Newton, featuring superb angular resolution \citep[e.g.][]{Brandt2015}. In the work reported in this paper we used the data of XBOOTES, the presently largest continuous \chandra{} survey covering the area of $\sim9$~deg$^2$, to conduct the most accurate measurement to date of the brightness fluctuations of unresolved CXB in the angular scales ranging from $\sim3\arcsec$ to $\sim 17\arcmin$. The XBOOTES survey with its average exposure of $\sim5$~ksec per field has an average point-source sensitivity of $\sim2\times10^{-15}\,\mathrm{erg\,cm^{-2}\,s^{-1}}$ in the \eSoftB{} (\Ken) and an average extended-source sensitivity of $\sim3\times10^{-14}\,\mathrm{erg\,cm^{-2}\,s^{-1}}$ (\Sref{sss:Unr_Exgal}). Given these flux limits, unresolved AGN and normal galaxies make the major contribution to the surface brightness of the unresolved CXB, accounting each for $\sim30\,\%$, with the galaxy clusters making a more modest contribution at the level of $\sim6-8\,\%$ (\Sref{sss:Unr_Exgal}). However, estimates for normal galaxies and galaxy clusters are highly uncertain, thus explaining that about $\sim 1/3$ of the CXB flux remains unaccounted for in this calculation. At a point-source sensitivity level of $\sim10^{-15}\,\mathrm{erg\,cm^{-2}\,s^{-1}}$ normal galaxies are not expected to contribute significantly to the unresolved CXB fluctuations. The unresolved AGN have a median redshift of $z\sim1.0$ and median luminosity of $L_{\eSoft} \sim 10^{42.6}\,\mathrm{erg\;s^{-1}}$ (\Sref{ss:UnresoAGN}). The unresolved galaxy clusters are about twice as close, with a median redshift of $z\sim0.6$ \Ac{(flux-weighted mean of $\left<z\right>\sim0.3$)} and have comparable median luminosity of $L_{\eSoft} \sim 10^{42.7}\,\mathrm{erg\;s^{-1}}$. With the standard scaling relations, these numbers correspond to an ICM temperature of $T\sim1.4\,\mathrm{keV}$, and a DMH mass of $M_{500}\sim10^{13.5}$ (\Sref{sss:Unr_Exgal}). After masking out resolved (point and extended) sources, we obtained the \PoSp{} of surface brightness fluctuations of unresolved CXB in the $3\arcsec-17\arcmin$ range of angular scales (\Fref{fig:PS_eSoft}). At sub-arcminute angular scales, the obtained \PoSp{} is consistent with the predicted point source shot noise of unresolved AGN, corrected for the PSF-smearing. However, at the angular scales exceeding $\gtrsim1\arcmin$ we detect a clear and highly statistically significant LSS signal above the AGN shot noise. After subtracting the point-source shot noise, we obtain the \PoSp{} of the LSS signal, which follows an approximate power-law shape with the slope of $-0.8\pm0.1$ (\Fref{fig:excess}) and has normalization corresponding to the fractional RMS variation of $(36\pm2)\,\%$. The detected LSS signal is by \Mc{almost two orders} of magnitude stronger than that expected from the AGN \THT{} (\Fref{fig:agn_tht}). We present strong evidence that it is associated with the \Mc{ICM of unresolved galaxy clusters, namely with their one-halo term. } In particular, we show that the LSS signal is not present in the \PoSp{} of resolved AGN (\Fref{fig:Ex_PntScr_Limit}), and it is much enhanced when resolved galaxy clusters are retained on the images (\Fref{fig:Ex_ES_Limit}). The energy dependence of the mean power of the CXB fluctuations at a $\sim$few arcmin angular scales is consistent with the energy spectrum of an optically thin plasma with temperature of $T=1.4$~keV redshifted to $z=0.6$, corresponding to the \Ac{average} ICM temperature and \Ac{median} redshift of unresolved galaxy clusters in the XBOOTES survey (\Fref{fig:PS_spec_Excess}). The shape of the \PoSp{} can be remarkably well described by the model of \citet{Cheng2004}, although the normalization of the theoretical spectrum is by about a factor of $\approx 2$ smaller than observed (\Fref{fig:Ex_rES_Cluster}). \Mc{The power spectrum of fluctuations of unresolved CXB carries information about the ICM structure in the outskirts (out to $\sim R_{200}$) of nearby ($z\lesssim0.1$) \GCG{} \citep[e.g.][]{Cheng2004}.} These scales are difficult to reach by conventional studies based on the surface brightness distribution of individual and usually relatively nearby objects. The shape of the observed \PoSp{} is determined by the spatial structure of ICM, and the redshift distribution of clusters and groups of galaxies, while its normalization is proportional to the square of their volume density. This underlines the enormous diagnostic potential of the unresolved CXB fluctuation analysis. | 16 | 9 | 1609.02941 |
1609 | 1609.09112_arXiv.txt | We analyze the physical properties and energy balance of density enhancements in two SPH simulations of the formation, evolution, and collapse of giant molecular clouds. In the simulations, no feedback is included, so all motions are due either to the initial, decaying turbulence, or to gravitational contraction. We define clumps as connected regions above a series of density thresholds. The resulting full set of clumps follows the generalized energy-equipartition relation $\sigma_{v}/R^{1/2} \propto \Sigma^{1/2}$, where $\sigma_{v}$ is the velocity dispersion, $R$ is the `radius", and $\Sigma$ is the column density. We interpret this as a natural consequence of gravitational contraction at all scales, rather than virial equilibrium. Nevertheless, clumps with low $\Sigma$ tend to show a large scatter around equipartition. In more than half of the cases, this scatter is dominated by external turbulent compressions that {\it assemble} the clumps, rather than by small-scale random motions that would disperse them. The other half does actually disperse. Moreover, clump sub-samples selected by means of different criteria exhibit different scalings. Sub-samples with narrow $\Sigma$ ranges follow Larson-like relations, although characterized by their respective value of $\Sigma$. Finally, we find that: i) clumps lying in filaments tend to appear sub-virial; ii) high-density cores ($n \ge 10^5$ cm$^3$) that exhibit moderate kinetic energy excesses often contain sink (``stellar'') particles, and the excess disappears when the stellar mass is taken into account in the energy balance; iii) cores with kinetic energy excess but no stellar particles are truly in a state of dispersal. | \label{sec:intro} Ever since the pioneering work of \citet{Larson1981} it has been recognized that molecular clouds (MCs) obey scaling relationships, that have been interpreted as representative of approximate virial equilibrium in the clouds between their internal ``turbulent'' motions and their self-gravity. Subsequently, however, there have been suggestions that these relations may actually be the result of observational selection effects \citep[e.g.,] [] {Kegel1989, Scalo90}. In addition, there have been attempts at generalization of these relations \citep[e.g.,] [] {KM86, Heyer2009} and reinterpretations in terms of global cloud collapse rather than virialization \citep[][hereafter B11] {Ballesteros2011}. Moreover, there are structures that are observed to possess kinetic energies in excess of those that would be consistent with equilibrium (or more generally, energy equipartition). In this paper we aim to investigate whether clumps forming in numerical simulations of clouds undergoing global gravitational contraction exhibit similar properties as those in observational surveys such as Larson scaling relations, and search for a cause of the apparent kinetic energy excesses seen in some subsets of clumps in observational samples. \subsection{Larson's relations and their generalization} \label{sec:Larson_gen} For over three decades, it has been accepted that molecular clouds (MCs) satisfy the so-called \citet{Larson1981} scaling relations between velocity dispersion ($\sigmav$), mean number density ($\nav$) and size ($L$). In their presently accepted form, these relations are \citep[e.g.,][] {Solomon+87, HB04} \begin{equation} \nav \approx 3400 \left(\frac{L}{1 {\rm \, pc}}\right)^{-1} \pcc, \label{eq:larson_dens} \end{equation} and \begin{equation} \sigmav \approx 1 \left(\frac{L}{1 {\rm \, pc}}\right)^{1/2} \kms. \label{eq:larson_vel} \end{equation} \citet{Larson1981} additionally showed that these relations implied that the velocity dispersion is close to the value corresponding to virial equilibrium. In what follows, we will more generally refer to this as ``near equipartition'' between the nonthermal kinetic and the gravitational energies. Also, it should be remarked that eq.\ (\ref{eq:larson_dens}) implies that the column density of the clouds, $\Sigma = \int_{\rm LOS} \rho d\ell$ is approximately the same for MCs of all sizes. In this expression, $\ell$ is the length element, and the integration is performed along the line-of-sight (LOS) through the cloud. However, the validity of Larson's relations has been questioned by various authors. \citet{Kegel1989} and \citet{Scalo90} argued that the apparent constancy of the column density may arise from selection effects caused by the need to exceed a certain minimum column density in order to detect the clouds, and by a maximum apparent column density caused by line saturation (optical thickening). This possibility was in fact recognized by \citet{Larson1981} himself. Some time later, \citet{Ballesteros2002} showed that, in numerical simulations of turbulent clouds, clumps defined by means of a column density threshold exhibited a Larson-like density-size relation, but clumps defined by means of a {\it volume} density threshold did not. Several years later, using the Boston University-FCRAO Galactic Ring Survey \citep[]{Jackson2006}, \citet[][hereafter H09]{Heyer2009} re-analyzed the giant molecular cloud (GMC) sample of \citet{Solomon+87}. The higher angular sampling rate and resolution available to H09, as well as the use of the $^{13}$CO $J=1$--0 line allowed them to obtain a much larger dynamic range in column density than that available to \citet{Solomon+87}. Moreover, H09 considered two different definitions for the cloud boundaries, thus effectively obtaining two different MC samples.\footnote{Contrary to some claims in the literature, these two definitions of the cloud boundaries do amount to {\it two} different MC samples, as the masses and velocity dispersions were measured for each cloud within each of the two boundaries. Thus, the `A2' clouds in H09 constitute a sample of smaller, denser objects within the `A1' sample, just like dense clumps and cores are substructures of their parent MCs, with independent dynamical indicators.} With this procedure, the GMC sample of H09 spanned over two orders of magnitude in column density, making it clear that column density is {\it not} constant for GMCs \citep[see also][]{Heyer+01}. Nevertheless, H09 noted that, in spite of the non-constancy of the column density, the GMCs are still consistent with virial equilibrium. They showed this by noting that their GMC sample satisfied \begin{equation} \frac{\sigmav} {R^{1/2}} \approx \left(\frac{\pi G \Sigma} {5}\right)^{1/2}. \label{eq:gener_Larson} \end{equation} When the column density is not constant, this relationship corresponds to virial equilibrium; i.e., to $|\Eg| = 2 \Ek$, with $\Ek$ being the nonthermal kinetic energy and $\Eg$ being the gravitational energy for a spherical cloud of uniform density and radius $R$. Thus, eq.\ (\ref{eq:gener_Larson}) can be considered as the generalization of Larson's relations when $\Sigma$ is not constant. Shortly thereafter, \citet{Lombardi+10} claim\-ed that the column density of GMCs is constant after all. Using near-infrared excess techniques, these authors argued that the mean GMC column density in their sample remained constant in spite of being sensitive to very low extinctions, thus suggesting that the minimum-column density imposed by a sensitivity threshold was not an issue. However, it has subsequently been recognized that this effect is natural for clouds with a $\Sigma$ probability density function (PDF) that peaks at some value and drops fast enough at lower column densities \citep{Beaumont+12, Ballesteros2012}. The lack of pixels at low $\Sigma$ implies that the dominant apparent column density will be that of the peak, and it is now recognized that the presence of a peak may be an artifact of incomplete sampling at low column densities \citep{Lombardi+15}. Therefore, at present there is no compelling evidence for the validity of the density-size relation (eq.\ \ref{eq:larson_dens}) for GMCs nor their substructures in general. On the other hand, the velocity dispersion-size relation,\footnote{In what follows, we will refer to this relation as the linewidth-size relation as well.} expression (\ref{eq:larson_vel}) above, has often been interpreted as the signature of supersonic turbulence, with an energy spectrum $E(k) \propto k^{-2}$, where $k$ is the wavenumber. Indeed, the velocity variance, interpreted as the average turbulent kinetic energy per unit mass in scales of size $\ell \le 2\pi/k$, given by $\sigmav^2(\ell) = \int_{k>2 \pi/\ell} E(k) dk$, scales as $\ell^{1/2}$ \citep[e.g.] [] {VS+00a, ES04, MO07}. In this case, the velocity dispersion- size relation would have a completely independent origin from that of the density-size relation, and the reason for the observed near- equipartition between the gravitational and turbulent kinetic energies would require a separate explanation. However, massive star-forming clumps notoriously do not conform to the $\sigmav$-$L$ relation \citep[e.g.,] [] {CM95, Plume+97, Shirley+03, Gibson+09, Wu+10}, a situation that appears inconsistent with a universal turbulent energy cascade spanning the whole range from the scale of GMCs down to the scale of massive clumps. An alternative interpretation was suggested by \citet[] [hereafter B11] {Ballesteros2011}, who proposed that the origin of the $\sigmav$-$L$ relation was not turbulence, but rather gravitational contraction of the clouds, combined with the observational selection effect of a restricted column density range. This possibility was actually suggested over four decades ago by \citet{GK74}. Similarly, \citet{Liszt+74} suggested that their line profiles and LOS-velocity maps of the Orion MC were consistent with extended radial motions, although they could not discriminate between expansion and collapse. However, the extended-motion scenario was soon dismissed by \citet{ZP74}, who argued that, if that were the case, then the star formation rate in MCs should be much larger than observed, and that systematic shifts between emission lines produced by HII regions at the centers of the clouds and absorption lines produced in the radially- moving cloud envelopes should be observed, but they are not. \citet{ZE74} then proposed that the observed linewidths corresponded to supersonic, small-scale turbulence. The small-scale turbulence scenario, however, suffers from a number of problems \citep[see] [for a detailed discussion] {VS15}. Instead, B11 have suggested a return to the scenario of gravitational collapse at the scale of the whole GMCs, with the problem of an excessive SFR being solved by early destruction of the clouds by stellar feedback \citep[][]{VS+10, ZA+12, Dale+12, Colin+13, ZV14}. B11 noted that the generalized Larson relation, eq.\ (\ref{eq:gener_Larson}) is not only satisfied by GMCs, but also by massive clumps that do not satisfy Larson's velocity dispersion-size relation, eq.\ (\ref{eq:larson_vel}). Thus, B11 interpreted the near-equipartition as evidence for free-fall in the clouds \citep[see also][]{Traficante+15} rather than near-virial equilibrium, noting that the virial and free-fall velocities differ only by a factor of $\sqrt{2}$. Indeed, for a freely collapsing cloud, defining the total energy as zero, the nonthermal kinetic energy and the gravitational energy satisfy $\Ek =|\Eg|$, so that, instead of eq.\ (\ref{eq:gener_Larson}), we have \begin{equation} \frac{\sigmav} {R^{1/2}} \approx \left(\frac{2 \pi G \Sigma} {5}\right)^{1/2}. \label{eq:free-fall} \end{equation} Generally, the observational errors and uncertainties in cloud and clump surveys are larger than this slight $\sqrt{2}$ factor, so that, for all practical purposes, {\it any evidence in favor of virial equilibrium based on energetics of the clouds can just as well be interpreted as evidence in favor of free collapse.} Recent observational studies have shown signatures of infall motions in line profiles along filaments and massive clumps in the Cygnus X region \citep[][]{Schneider+10}, in massive star-forming cores of the infrared dark cloud SDC335.579-0.272 \citep[][]{Peretto+13}, and in massive starless cores \citep[][]{Traficante+15}, supporting the notion that this systems are consistent with a global gravitational collapse. Moreover, \citet{Traficante+15} introduce an equivalent analysis as in H09 to demonstrate that most of the non-thermal motions in their sample originate from self-gravity. \subsection{Deviations from energy equipartition} \label{sec:attempts} An additional important feature in the GMC sample studied by H09 is that the clouds tended to lie systematically {\it above} the virial-equilibrium line in a plot of $\sigmav/R^{1/2}$ {\it vs.} $\Sigma$, a plot that we will refer to as the Keto-Heyer or KH diagram \citep{KM86, Heyer2009}. This feature has received different interpretations by different authors. H09 themselves interpreted it simply as a systematic underestimation of the cloud masses, due to the various assumptions they used in determining the masses from $^{13}$CO emission. On the other hand, \citet{Dobbs2011} have interpreted this feature as evidence that most of the GMCs are gravitationally unbound, probably because they form by cloud-cloud collisions, which feed a large velocity dispersion that unbinds the GMCs. \citet{KM86} and \citet{Field2011}, instead, have assumed that the clouds are gravitationally unbound, but confined by an external pressure, while B11 suggested that the GMCs are actually collapsing, and that, at face value, the H09 data are slightly more consistent with free-fall than with virial equilibrium, since the free-fall velocity is slightly larger than the virial one. In addition to the slight systematic overvirial nature of cloud surveys, in several observational clump and core surveys, some objects appear to be {\it strongly} overvirial, exhibiting values of the {\it virial parameter}, $\alpha \equiv 5 \sigmav^2 R/GM \sim 10$--100 \citep[see,e.g., fig.\ 16 of] [] {Barnes+11}, especially in the case of low-mass objects. These objects are traditionally interpreted as having a kinetic energy significantly larger than their gravitational energy, perhaps due to driving by stellar feedback, and therefore requiring confinement by external pressure to prevent them from dispersing, as in the interpretation by \citet{Field2011} of the H09 sample. However, if we adopt the interpretation that star-forming GMCs are undergoing global and hierarchical collapse, pressure confinement is not satisfactory, since in this scenario the clumps should be gravitationally-dominated as well. Investigating the origin of these kinetic energy excesses within the scenario of collapsing clouds is one of the goals of this paper. \subsection{This work} \label{sec:this_work} In this work we create an ensemble of clumps in simulations of the formation and evolution of molecular clouds, in order to investigate their energy balance under a scenario of initial turbulence and subsequent gravitational collapse. Our simulations, of course, have a number of limitations, which are discussed in more detail in Sec.\ \ref{sec:caveats}, but here we note that they neglect magnetic fields, stellar feedback, and have relatively low masses that restrict the clumps and cores we obtain to values typical of low-mass star-forming regions, which form low-mass stellar groups or low-mass clusters. Nevertheless, we expect that the results we obtain can be extrapolated to regions of larger masses. The organization of the paper is as follows. In Sec. \ref{sec:sim} we briefly describe the simulations, and in Sec.\ \ref{subsec:tools}, we describe the clump-finding algorithm used to define clumps at various values of volume density, $\nth$, as well as the selection criteria we used in order to avoid considering unrealistic clumps (i.e. those that would be affected by stellar feedback in reality). Next, in Sec.\ \ref{sec:res}, we present our results on the energetics of the clumps and cores and their implications on Larson's relations. In Sec.\ \ref{sec:discussion} we discuss our work in context with recent related numerical studies, as well as the range of applicability, possible extrapolations, and limitations of our study. Finally, in Sec.\ \ref{sec:concls} we present a summary and some conclusions. | \label{sec:concls} In this paper we have investigated the intrinsic (rather than derived from synthetic observations) physical conditions of clumps and cores in two SPH simulations of the formation and evolution of molecular clouds formed by converging motions in the warm neutral medium (WNM). The two simulations attempt to span a range of likely motions in this medium. In both simulations, once the dense clouds form, they soon begin to contract gravitationally, and some time later (a few Myr) they begin to form stars, as in the general scenario described by \citet{VS+07} and \citet{HH08}. Neither of the simulations includes turbulence-driving stellar feedback nor magnetic fields, and so all of the kinetic energy is either driven by gravity or is a residual of the turbulent/compressive motions that initiated the formation of the clouds. Within this context of globally contracting molecular clouds (MCs), we have investigated whether the clumps within them follow the Larson scaling relations, or their generalization, as proposed by H09 and B11. We have also investigated the physical conditions in clumps that appear to have an excess of kinetic energy, in an attempt to understand the physical processes that cause this apparent over-virialization. We created an ensemble of clumps in each simulation by defining clumps as connected sets of SPH particles above a certain density-threshold $\nth$, so that a single clump at a lower threshold may contain several clumps at a higher value of $\nth$. The objects defined at the highest thresholds ($\nth \ge 10^5\, \pcc$) are refereed to as ``cores''. Our results and conclusions may be summarized as follows: \begin{itemize} \item The full ensemble of clouds, clumps and cores does not follow either of the Larson scaling relations, but mostly follows their generalization, as proposed by H09 and B11. Nevertheless, low column density clumps in particular exhibit a large scatter, with a significant fraction of the clumps having values of the $\calh \equiv \sigmav/R^{1/2}$ parameter of up to an order of magnitude larger than the virial value, similar to the situation in various observational studies. \item We noted that, as emphasized by B11, the kinetic energy implied by free-falling motions is only a factor of $\sqrt{2}$ larger than that for virial equilibrium. We therefore generically refer to this condition as ``energy equipartition''. \item In our simulations, the equipartition condition is due to gravitational contraction, by construction. \item The clumps defined at a single threshold $\nth$ do not exhibit density-size or velocity dispersion-size relations. Instead, the exhibit nearly constant volume density, in agreement with previous studies \citep{Ballesteros2002}. However, ensembles of clumps that exhibit near-equipartition and that are selected by column density ranges, do exhibit Larson-like relations, suggesting that these relations are special cases of the more general equipartition condition. \item We find examples of clouds, clumps and cores that exhibit excess kinetic energies over the equipartition level at both low- and high-column densities. Low-column density clumps that exhibit this excess are the least massive, while the more massive ones are closer to equipartition. Moreover, for more than 50$\%$ of the low-density clumps with an $\calh$ excess in both simulations, the velocity field in the clouds appears to be convergent (i.e. have negative net divergence). This suggests an evolutionary process in which a turbulent compression initially dominates the kinetic energy and exceeds the gravitational energy of the forming cloud. However, as the cloud becomes denser and more massive, the gravitationally-driven velocity becomes dominant. Also, this suggests that the observation of an excess kinetic energy does not necessarily imply that a clump will disperse or needs an external thermal confining pressure to avoid dispersal. The excess kinetic energy may simply reflect the initial compressive motions within the clump. In this case, instead of {\it confinement} of the cores by thermal pressure, we have {\it assembly} by ram pressure. \item Some of the high-column density cores that exhibit kinetic energy excesses contain stellar particles that increase the total gravitational potential in the volume of the clump. When this stellar mass is added to the gas mass in the energy budget of the core, the gas+stars system returns to near equipartition. \item Some high-column density clumps with kinetic energy excesses, however, do not contain stellar particles, so that the the above correction cannot be applied. Investigation of the velocity field in these cases does show a rotating and/or expanding motion, so that these objects are in the process of being disrupted, and will not form stars. Because this process is occurring at high densities, the driver of these disrupting motions is likely to be the turbulence generated by the large-scale collapse. \item We also investigated the possibility that excess kinetic energies in high-column density cores might be due to the cores being located in filamentary clumps, with net accretion from the filament onto the core, so that the velocity dispersion in the cores might represent the gravitational potential of the mass in the filament. However, this mechanism does not seem to be operational. We find that the filaments and their embedded cores begin their evolution roughly simultaneously, accreting material from the cloud mostly perpendicularly to the filament. Accretion from the filament onto the core begins later, when the core has become more massive and has already started to form stars. Thus, cores that are actively accreting from their parent filaments are already in advanced star-forming stages, and do not correspond to pre- or early protostellar objects. \end{itemize} | 16 | 9 | 1609.09112 |
1609 | 1609.06938_arXiv.txt | In the first paper of this series, we study the level population problem of recombining carbon ions. We focus our study on high quantum numbers anticipating observations of Carbon Radio Recombination Lines to be carried out by the LOw Frequency ARray (LOFAR). We solve the level population equation including angular momentum levels with updated collision rates up to high principal quantum numbers. We derive departure coefficients by solving the level population equation in the hydrogenic approximation and including low temperature dielectronic recombination effects. Our results in the hydrogenic approximation agree well with those of previous works. When comparing our results including dielectronic recombination we find differences which we ascribe to updates in the atomic physics (e.g., collision rates) and to the approximate solution method of the statistical equilibrium equations adopted in previous studies. A comparison with observations is discussed in an accompanying article, as radiative transfer effects need to be considered. | The interplay of stars and their surrounding gas leads to the presence of distinct phases in the interstellar medium (ISM) of galaxies (e.g. \citealt{field1969,mckee1977}). Diffuse atomic clouds (the Cold Neutral Medium, CNM) have densities of about $50~\mathrm{cm^{-3}}$ and temperatures of about $80~\mathrm{K}$, where atomic hydrogen is largely neutral but carbon is singly ionized by photons with energies between $11.2~\mathrm{eV}$ and $13.6~\mathrm{eV}$. The warmer ($\sim8000~\mathrm{K}$) and more tenuous ($\sim0.5~\mathrm{cm^{-3}}$) intercloud phase is heated and ionized by FUV and EUV photons escaping from HII regions \citep{wolfire2003}, usually referred to as the Warm Neutral medium (WNM) and Warm Ionized Medium (WIM). The phases of the ISM are often globally considered to be in thermal equilibrium and in pressure balance \citep{savage1996, cox2005}. However, the observed large turbulent width and presence of gas at thermally unstable, intermediate temperatures attests to the importance of heating by kinetic energy input. In addition, the ISM also hosts molecular clouds, where hydrogen is in the form of $\mathrm{H_2}$ and self-gravity plays an important role. All of these phases are directly tied to key questions on the origin and evolution of the ISM, including the energetics of the CNM, WNM and the WIM; the evolutionary relationship of atomic and molecular gas; the relationship of these ISM phases with newly formed stars; and the conversion of their radiative and kinetic power into thermal and turbulent energy of the ISM (e.g. \citealt{cox2005,elmegreen2004, scalo2004,mckee2007}). The neutral phases of the ISM have been studied using optical and UV observations of atomic lines. These observations can provide the physical conditions but are limited to pinpoint experiments towards bright background sources and are hampered by dust extinction \citep{snow2006}. At radio wavelengths, dust extinction is not important and observations of the 21 cm hyperfine transition of neutral atomic hydrogen have been used to study the neutral phases (e.g. \citealt{weaver1973,kalberla2005,heilesandtroland2003b}). On a global scale, these observations have revealed the prevalence of the two phase structure in the interstellar medium of cold clouds embedded in a warm intercloud medium but they have also pointed out challenges to this theoretical view \citep{kulkarni1987, kalberla2009}. It has been notoriously challenging to determine the physical characteristics (density, temperature) of the neutral structures in the ISM as separating the cold and warm components is challenging (e.g. \citealt{heilesandtroland2003a}). In this context, Carbon radio recombination lines (CRRLs) provide a promising tracer of the neutral phases of the ISM (e.g. \citealt{peters2011, oonk2015a}). Carbon has a lower ionization potential (11.2~eV) than hydrogen (13.6~eV) and can be ionized by radiation fields in regions where hydrogen is largely neutral. Recombination of carbon ions with electrons to high Rydberg states will lead to CRRLs in the sub-millimeter to decameter wavelength range. Carbon radio recombination lines have been observed in the interstellar medium of our Galaxy towards two types of clouds: diffuse clouds (e.g.: \citealt{konovalenko1981, erickson1995, roshi2002, stepkin2007,oonk2014}) and photodissociation regions (PDRs), the boundaries of HII regions and their parent molecular clouds (e.g.: \citealt{natta1994, wyrowski1997, quireza2006}). The first low frequency (26.1 MHz) carbon radio recombination line was detected in absorption towards the supernova remnant Cas A by \citet{konovalenko1980} (wrongly attributed to a hyperfine structure line of ${^{14}}\mathrm{N}$, \citealt{konovalenko1981}). This line corresponds to a transition occurring at high quantum levels ($n=631$). Recently, \citet{stepkin2007} detected CRRLs in the range 25.5--26.5 MHz towards Cas A, corresponding to transitions involving levels as large as $n=1009$. Observations of low frequency carbon recombination lines can be used to probe the physical properties of the diffuse interstellar medium. However, detailed modeling is required to interpret the observations. \citet{watson1980,walmsley1982} showed that, at low temperatures ($T_e \lesssim 100~\mathrm{K}$), electrons can recombine with carbon ions by simultaneously exciting the ${^2}P_{1/2}-{^2}P_{3/2}$~fine structure line, a process known as dielectronic recombination\footnote{ This process has been referred in the literature as dielectronic-like recombination or dielectronic capture by \citet{watson1980} to distinguish from the regular dielectronic recombination. Dielectronic capture refers to the capture of the electron in an excited $n$-state accompanied by simultaneous excitation of the ${^2}P_{1/2}$ core electron to the excited ${^2}P_{3/2}$ state. The captured electron can either auto ionize, collisional transferred to another state, or radiatively decay. Strictly speaking, dielectronic recombination refers to dieclectronic capture followed by stabilization. However, throughout this article we will use the term dielectronic recombination to refer to the same process as is common in the astronomical literature.}. Such recombination process occurs to high $n$ states, and can explain the behavior of the high $n$~CRRLs observed towards Cas A. \citet{walmsley1982} modified the code from \citet{brocklehurst1977} to include dielectronic recombination. \citet{payne1994} modified the code to consider transitions up to 10000 levels. All of these results assume a statistical distribution of the angular momentum levels, an assumption that is not valid at intermediate levels for low temperatures. Moreover, the lower the temperature, the higher the $n$-level for which that assumption is not valid. The increased sensitivity, spatial resolution, and bandwidth of the Low Frequency ARray (LOFAR, \citealt{vhaarlem2013}) is opening the low frequency sky to systematic studies of high quantum number radio recombination lines. The recent detection of high level carbon radio recombination lines using LOFAR towards the line of sight of Cas A \citep{asgekar2013}, Cyg A \citep{oonk2014}, and the first extragalactic detection in the starburst galaxy M82 \citep{morabito2014b} illustrate the potential of LOFAR for such studies. Moreover, pilot studies have demonstrated that surveys of low frequency radio recombination lines of the galactic plane are within reach, providing a new and powerful probe of the diffuse interstellar medium. These new observations have motivated us to reassess some of the approximations made by previous works and to expand the range of applicability of recombination line theory in terms of physical parameters. In addition, increased computer power allows us to solve the level population problem considering a much larger number of levels than ever before. Furthermore, updated collisional rates are now available \citep{vrinceanu2012}, allowing us to explicitly consider the level population of quantum angular momentum sub-levels to high principal quantum number levels. Finally, it can be expected that the Square Kilometer Array, SKA, will further revolutionize our understanding of the low frequency universe with even higher sensitivity and angular resolution \citep{oonk2015a}. In this work, we present the method to calculate the level population of recombining ions and provide some exemplary results applicable to low temperature diffuse clouds in the ISM. In an accompanying article (\citealt{salgado2016}, from here on Paper II), we will present results specifically geared towards radio recombination line studies of the diffuse interstellar medium. In Section~\ref{section_method}, we introduce the problem of level populations of atoms and the methods to solve this problem for hydrogen and hydrogenic carbon atoms. We also present the rates used in this work to solve the level population problem. In Section~\ref{section_results}, we discuss our results focusing on hydrogen and carbon atoms. We compare our results in terms of the departure coefficients with previous results from the literature. In Section~\ref{section_conclusions}, we summarize our results and provide the conclusions of the present work. | We have solved the level population equation for hydrogenic atoms using novel rates involved in the process. The level population equation is solved in two approximations: the $n$ and the $nl$ method. The departure coefficients obtained using the $n$ method are similar to values from the literature (e.g. \citealt{brocklehurst1970} and \citealt{shaver1975}). Our results using the $nl$ method reproduce those from \citet{hummer1987} well, once allowance is made for updates in the collisional rates. By including the dielectronic recombination process together with the $nl$ method we are able to model the level population of carbon in terms of the departure coefficients. Our results are qualitatively similar to those of \citet{watson1980,walmsley1982}. However, the values obtained here differ considerably from those from the literature. The differences can be understood in terms of the use of improved collision rates and the improved numerical approach using the $nl$ method. We confirm that dielectronic recombination can indeed produce an increase on the values of the departure coefficients at high $n$ levels compared to the hydrogenic values. In anticipation of low frequency radio recombination line surveys of the diffuse interstellar medium now being undertaken by LOFAR, we have expanded the range of applicability of the formulation to the conditions of the cold neutral medium. For this environment, external radiation fields also become important at intermediate principal quantum levels while at high levels the influence of radiation fields on the level population is less important In an accompanying paper \citep{salgado2016}, we discuss the expected line strength for low frequency carbon radio recombination lines and the influence of an external radiation field. Throughout this work we have used a zero radiation field. In this companion paper we compare our results to existing observations of CRRLs towards Cas~A and regions in the inner galaxy. We also describe the analysis techniques and diagnostic diagrams that can be used to analyze the forthcoming LOFAR CRRL survey. The departure coefficients obtained here will be used to analyze the LOFAR observations of Cas~A in a future article \citep{oonk2015b}. | 16 | 9 | 1609.06938 |
1609 | 1609.08052_arXiv.txt | % The G14.225-0.506 infrared dark cloud (IRDC G14.2) displays a remarkable complex of parallel dense molecular filaments projected on the plane of the sky. Previous dust emission and molecular-line studies have speculated whether magnetic fields could have played an important role in the formation of such long-shaped structures, which are hosts to numerous young stellar sources. In this work we have conducted a vast polarimetric survey at optical and near-infrared wavelengths in order to study the morphology of magnetic field lines in IRDC G14.2 through the observation of background stars. The orientation of interstellar polarization, which traces magnetic field lines, is perpendicular to most of the filamentary features within the cloud. Additionally, the larger-scale molecular cloud as a whole exhibits an elongated shape also perpendicular to magnetic fields. Estimates of magnetic field strengths indicate values in the range $320 - 550\,\mu$G, which allows sub-alfv\'enic conditions, but does not prevent the gravitational collapse of hub-filament structures, which in general are close to the critical state. These characteristics suggest that magnetic fields played the main role in regulating the collapse from large to small scales, leading to the formation of series of parallel elongated structures. The morphology is also consistent with numerical simulations that show how gravitational instabilities develop under strong magnetic fields. Finally, the results corroborate the hypothesis that a strong support from internal magnetic fields might explain why the cloud seems to be contracting on a time scale $2-3$ times larger than what is expected from a free-fall collapse. | \label{introduction} Filamentary structures in the interstellar medium (ISM) are commonly observed in many different types of environments, such as diffuse nearby clouds \citep{penprase1998,mcclure2006}, giant molecular clouds \citep{lis1998,hill2011}, H{\sc ii} regions \citep{anderson2012,minier2013}, and supernova remnants \citep{gomez2012}. Their presence in the Milky Way Galaxy has typically been revealed by numerous different observing techniques, including visual extinction, HI emission, molecular line surveys and dust thermal emission. In particular, dense molecular filaments are in general associated with star-forming regions. \citet{myers2009} pointed out, for instance, that all the nearest low-mass star formation sites (within $300\,$pc from the Sun) seem to present a hub-filament structure, with some of them showing evenly-spaced parallel filaments. Although filaments are known for many decades, during more recent years, observations of dust thermal emission from the {\it Herschel} space observatory \citep{pilbratt2010} provided a groundbreaking understanding of filaments in the ISM, showing that they are in fact ubiquitous especially within giant molecular clouds \citep{andre2010,molinari2010}, which includes both quiescent and star-forming regions. The recognition of filaments as active sites of star-formation was made clear by the fact that most of the observed pre-stellar cores seem to form in gravitationally unstable filaments \citep{andre2010,arzou2011}. That led to an increasing interest in explaining how these structures are formed and how they evolve. Although turbulent motions in the ISM might be responsible for the formation of some filaments \citep{arzou2011}, other plausible explanations are the convergence of flows, large-scale collisions between filaments, or gravitational instabilities \citep{schneider2010,jimenez2010,nakamura2012,nakajima1996,vanloo2014}, a scenario which is also supported by numerical models \citep{gomez2014}. In addition, it is well known that the interstellar medium is entirely threaded by a large-scale structure of magnetic field lines that pervades the whole Galaxy \citep{mathewson1970,reiz1998,heiles2000,santos2011}. This includes the filaments as well as the dense molecular cores where the star formation is taking place \citep[e.g.,][]{alves2008,girart2006,girart2009,zhang2014,li2015}. In general, a small level of ionization is sufficient to provide enough coupling between the magnetic fields and the interstellar gas \citep{heiles2005}. Indeed, magnetic fields might also play an important role in generating filamentary structures, as suggested by several authors \citep{nagai1998,nakamura2008,li2013,li2015}. G14.225-0.506 (hereafter IRDC G14.2) is an infrared dark cloud at a distance of $1.98^{+0.13}_{-0.12}\,$kpc \citep{xu2011, wu2014} that shows an intricate pattern of filaments. These filaments are clearly seen in absorption against the bright mid-infrared background Galactic emission, as identified by \citet{2009peretto} using {\it Spitzer Space Telescope} data\footnote{Based on the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire \citep[GLIMPSE,][]{2003benj} and the Multiband Imaging Photometer for Spitzer Galactic Plane Survey \citep[MIPSGAL,][]{2009carey}.}. This region is part of a larger complex of clouds including the well-known M17 star-forming area \citep{elm1976}. Later studies revealed star formation signs such as H$_{2}$O maser emission \citep{1981jaffe,palagi1993,wang2006} and emission from dense gas tracers toward IRAS 18153-1651, which is one of the bright infrared sources in the region \citep{plume1992,anglada1996}, with a luminosity of $1.1\times10^4~L_{\sun}$. Furthermore, many young stellar objects were later identified by \citet[][who labeled this region as M17 SWex]{povich2010}, including several Class 0 and I sources. Although IRDC G14.2 does not appear to host very massive stars, a few ultra-compact H{\sc ii} regions are located amongst its filamentary structures \citep{bronfman1996,jaffe1982}. \citet[][herafter Paper I]{busquet2013} presented ammonia observations in IRDC G14.2, inferring that the parallel arrangement of most filaments could be explained by the gravitational collapse of an unstable thin layer threaded by magnetic fields \citep{vanloo2014}. The sky-projected morphology of magnetic field lines may be mapped through studies of the interstellar polarization due to magnetically aligned dust particles, either through observations of background starlight, or direct thermal emission from dust. Although the detailed aspects of the alignment mechanism is one of the most long-standing issues in the physics of the ISM, it is now generally believed that radiative torques are a dominant effect \citep{dolginov1976,draine_wein_1996,lazarian2007}, as suggested by different studies \citep{whittet2008,andersson2011,alves2014,jones2015}. Even though different large-scale polarization emission surveys have been providing and unprecedented view of magnetic fields in the ISM (such as Planck -- \citealt{planck2014} -- and BLASTPol -- \citealt{2016fissel}), the spatial resolution needed to distinguish filamentary features at distant clouds is still a challenge, making optical and near-infrared (NIR) polarimetry of background starlight a viable option. IRDC G14.2 is an ideal target for investigating the role of magnetic fields in generating filamentary structures. In this work, we present a vast extension of a preliminary polarimetric dataset previously shown in Paper I. This includes optical and NIR observations encompassing all the filamentary network of IRDC G14.2, as well as the associated large-scale molecular cloud. In Section \ref{s:obsdata} we describe the polarimetric observations, as well as the data processing. Results and analysis are shown in Section \ref{s:results}, which includes studies of the relative orientation between magnetic fields with the cloud and internal filaments, as well as estimates of various important physical paramenters. A detailed discussion of the results is given in Section \ref{s:discussion} and the final conclusions in Section \ref{s:conclusions}. | \begin{enumerate} \item We compared the orientation of magnetic fields with filaments and hubs, and also with the molecular cloud in which these structures are embedded. It is clear that magnetic fields are perpendicular both to the small-scale filamentary features and to the large-scale cloud. For filaments, this condition holds true with few exceptions even when considering Monte Carlo simulations which account for sky-projection effects. These characteristics are consistent with a scenario in which magnetic fields regulated the gravitational collapse from large ($\approx 30\,$pc) to small scales ($\approx 2\,$pc); \item Combining the polarization data with dust emission and molecular line observations, we estimate total magnetic fields strengths, Alfv\'en Mach numbers and mass-to-magnetic-flux ratios. % The structures are predominantly in a sub-alf\'enic and in close-to-critical condition, suggesting that magnetic fields are strong enough to overcome turbulent motions, but not sufficient to prevent the gravitational collapse. The high magnetic field values corroborate previous numerical simulations that show that these conditions eventually lead to a gravitational instability developing along magnetic field lines, therefore generating filaments organized in a parallel arrangement; \item The range of magnetic field values obtained for the filaments and hubs ($\approx 320 - 550\,\mu$G) is consistent with estimates based on simple equipartition assumptions by \citet{elm1976}, who suggested that internal magnetic field strengths would be around $340\,\mu$G. According to their interpretation, the presence of such strong magnetic fields might be a necessary condition to explain why the large-scale cloud is possibly contracting in a time scale $2-3$ times larger than what expected from the free-fall time. \end{enumerate} As a precursor to a massive OB association presenting numerous filamentary interstellar features and young stellar sources, the IRDC G14.2 cloud proves to be an ideal star-forming site to study the underlying physical conditions regulating the gravitational collapse. This is an important target for additional analysis, particularly using high-resolution polarization emission surveys (in the far-infrared or submillimeter wavelengths) or even spectral data focused on Zeeman splitting. This would be a natural continuation of this work, given the significant role played by magnetic fields in shaping the filamentary morphology and regulating the collapse. More specifically, magnetic field strengths (along with $M_{\mathrm{A}}$ and $\lambda$ values) could be better constrained with this kind of observation, specially if comparisons with numerical simulations are made, assuming the specific physical conditions of this cloud and its sub-structures. | 16 | 9 | 1609.08052 |
1609 | 1609.04756_arXiv.txt | Emission in the Ca \II{} H \& K line cores has long been known to be a good proxy for magnetic activity in the Sun \citep{Hall:2008}. \cite{Wilson:1978} was the first to use this emission to demonstrate the magnetic variability for an ensemble of Sun-like stars, using a decade of synoptic Ca \II{} H \& K observations from the Mount Wilson Observatory (MWO). The MWO HK project began in 1966 and continued until 2003, with the largest compendium of stellar activity for 111 stars with up to 25 years of observations appearing in \cite{Baliunas:1995}. The MWO HK project was the basis of numerous investigations of activity, its relationship to stellar age and rotation, and implications for dynamo theory \citep[see][, and references therein]{Baliunas:1998}.A complimentary synoptic observation program began at Lowell Observatory in the mid-1990's using the Solar Stellar Spectrograph (SSS), designed to take low resolution spectra covering the Ca \II{} H \& K region for the Sun and stars with the same spectrograph \citep{Hall:1995,Hall:2007b}. The SSS program continues to this day, and 57 of its $\sim$100 targets overlap with the MWO HK project. We combine the data from these two surveys making time series of nearly 50 years in length. This was done for the first time in \cite{Egeland:2015} for the young solar analog HD 30495. In that case, the long time series allowed for the identification of three and a half stellar cycles, with a mean period of $\sim$12 years, for a star that previously appeared to be acyclically variable. Work is ongoing to calibrate, combine, and analyze MWO+SSS time series for a sample of 27 solar analog stars with $0.59 \leq (B-V) \leq 0.69$ \citep{CayreldeStrobel:1996}, in order to understand the solar dynamo in the stellar context. In particular, we seek to better understand (1) whether the pattern of solar variability is common among Sun-like stars (2) how the patterns of long-term variability in the ensemble depend on stellar properties such as mass, luminosity, radius, metalicity, and rotation. Preliminary results from this project were presented at this conference \citep{Egeland:2016:cs19talk,Egeland:2016:cs19poster} and are summarized in these proceedings. The full details and final results are to appear in a peer-reviewed journal in the near future. | Diligent long-term observation programs by the Mount Wilson and Lowell Observatory provide unique data for understanding the variability patterns of Sun-like stars, with composite time series now approaching 50 years in length. Questions on the uniqueness of the solar cycle, and the sensitivity in stellar dynamos to changes in fundamental properties can be approached using these data, improving our understanding of the dynamo and our Sun in context. Work is ongoing to carefully quantify what these data can tell us about these questions, but the initial results indicate that the clean, clear solar cycle may be an exceptional case in the limited parameter regime of solar analogs. | 16 | 9 | 1609.04756 |
|
1609 | 1609.05346_arXiv.txt | Small $^3$He-rich solar energetic particle (SEP) events have been commonly associated with extreme-ultraviolet (EUV) jets and narrow coronal mass ejections (CMEs) which are believed to be the signatures of magnetic reconnection involving\deleted{open}field \added{lines open to interplanetary space}. The elemental and isotopic fractionation in these events are thought to be caused by processes confined to the flare sites. In this study we identify 32 $^3$He-rich SEP events observed by the {\sl Advanced Composition Explorer} near the Earth during the solar minimum period 2007--2010 and examine their solar sources with the high resolution {\sl Solar Terrestrial Relations Observatory} (STEREO) EUV images. Leading the Earth, STEREO-A provided for the first time a direct view on $^3$He-rich flares, which are generally located on the Sun's western hemisphere. Surprisingly, we find that about half of the $^3$He-rich SEP events in this survey are associated with large-scale EUV coronal waves. An examination of the wave front propagation, the source-flare distribution and the coronal magnetic field \replaced{extrapolations}{connections} suggests that the EUV waves may affect the injection of $^3$He-rich SEPs into interplanetary space. | \label{sec:intro} Solar energetic particles (SEPs) can be accelerated either at coronal mass ejection (CME) driven shocks in gradual large SEP (LSEP) events or by processes related to magnetic reconnection at the flare sites in impulsive or $^3$He-rich SEP events \citep[see review by][]{rea13}. Anomalous elemental composition in impulsive SEP events, extremely different from coronal abundances, has not been so far sufficiently explained \citep[see][for a review]{koc84,mas07,rea15}. In addition to a preferential acceleration by resonant interactions with MHD turbulence \citep[e.g.,][]{mil98,liu06}, other non-wave mechanisms have been proposed \citep[e.g.,][]{dra12}. In contrast to LSEP events, $^3$He-rich SEP events have lower ion intensities and energies, shorter durations and are associated with minor X-ray flares. Though they are numerous these features make them a less explored phenomenon. Solar sources of $^3$He-rich SEPs have been often associated with coronal X-ray and EUV jets \citep[e.g.,][]{nit06,nit08,wan06,buc14,che15,inn16} which sometimes show a high altitude extension in jet-like CMEs \citep{kah01,wan06,wan12}. Jets have been considered to be a signature of magnetic reconnection involving open field \citep[e.g.,][]{shi92}. Predominantly narrow CMEs in these events have been also reported by \citet{nit06}. Recently, new observations from the {\sl Solar Dynamics Observatory} (SDO) and {\sl Solar Terrestrial Relations Observatory} (STEREO) have revealed a few $^3$He-rich SEP events associated with large-scale coronal EUV waves \citep{nit15,buc15}. This finding appears to be surprising in view of previous associations with narrow flare signatures. Four of the reported $^3$He-rich SEP events with a coronal wave were accompanied by slow ($\lesssim$300\,km\,s$^{-1}$) CMEs \citep{nit15} but two others were without a CME \citep{buc15}. Coronal waves have been so far reported in LSEP events in association with fast and wide CMEs \citep[e.g.,][]{tor99,lar14,mit14}. However, there is still some controversy on the nature of coronal waves \citep[see][and references therein]{pat12,liu14,war15}. Presently, it is believed they are true magnetosonic waves and not just coronal magnetic field reconfigurations induced by a CME. Indeed, a recent report on EUV waves associated with X-ray flares alone \citep{nit13} appears to be consistent with the true wave concept. The aim of this paper is to examine how common an association of $^3$He-rich SEPs with large-scale coronal waves is without setting any restriction criteria on the selection of the events. For this purpose we collect 32 $^3$He-rich SEP events within the solar activity minimum period 2007--2010 to avoid source confusion in the more active periods. We identify solar sources of $^3$He-rich SEPs using new viewing perspectives from angularly separated spacecraft. These observations are presented in Section \ref{sec:rad}. In Section \ref{sec:con} we discuss possible implications of associated coronal waves on energetic ion characteristics in these events. | \label{sec:con} We collected 32 $^3$He-rich SEP events observed by ACE at L1 during the solar minimum period 2007--2010. This extended period of low solar activity provided favorable conditions for the identification of $^3$He-rich solar sources characterized by faint flare signatures. We include in this study also weak events with $^3$He intensity close to the detection threshold. We examined the solar sources with STEREO EUV imaging observations. At the beginning of 2007 the STEREO-A was near the Earth while at the end of the investigated period the angular separation extended to $\sim$90$^{\circ}$ from the Earth-Sun line allowing for the first time a direct view of the near west solar limb SEP sources. Surprisingly, we find that more than half of the events of this survey (20 of 32) are clearly associated with large-scale coronal waves. The finding remains unchanged (16 of 28) after excluding the events with a high-energy solar proton component. Other events, not counted in this number, like the 2008 June 16 and 2010 November 14 events, might also have secondary ion injections related to the wave expulsion (see Appendix). The EUV wave fronts passed by or showed close approach the L1 magnetic foot-point in several ($\sim$13--17) investigated events. However, this occurred either too fast after the type III radio burst onset or the events were too weak to distinguish between the flare and the potential wave-front ion injections. At least 10 events of this study have been associated with EUV jets including the wave events starting as a jet. Some events classified as having a brightening might have shown a narrow ejection from a different observing angle. Several $^3$He-rich events in the investigated period were reported in previous surveys (see references in Table \ref{tab:tab1}). \citet{mas09} focused on solar activity minimum conditions associated with acceleration of $^3$He-rich SEPs during the period March 2007--December 2008, \citet{wie13} on $^3$He-rich SEPs with wide longitudinal spread observed on multiple spacecraft between January 2007 and January 2011, and \citet{nit15} on solar source characteristics derived from high resolution EUV images on SDO in the period May 2010--May 2014. The present study shows a quite broad source flare longitude distribution in the $^3$He-rich SEP events where the most remote locations were associated with EUV waves. It has been noted that the connection to a distant source can not be sufficiently explained by a divergence of the coronal field lines below the source surface \citep{wie13}. The magnetic connection controlled by large scale EUV waves in the corona has been discussed in several LSEP events \citep[e.g.,][]{tor99,kru99,mal09,rou12,nit12,par13,ric14} and recently suggested as a possible mechanism also for $^3$He-rich SEPs \citep{nit15}.\deleted{We found that all previously reported events with a wide longitudinal spread of $^3$He-rich SEPs are associated with large-scale coronal waves. This provides new insights on energetic ion transport from their sources on the Sun to 1 AU.}\added{We examined magnetic connection combining the PFSS extrapolations of coronal field with Parker spiral model of IMF.} The PFSS model has been considered to be adequate for the extrapolation of global magnetic fields in the corona. For our 12 $^3$He-rich SEP events with jet/brightening the\deleted{PFSS model (with the source surface at 1.9--2.5\,$R_{\sun}$) predicts} magnetic connection \replaced{via coronal field consistent with the in-situ polarity in 10 cases (events 1, 4, 5, 13, 23, 24, 25, 27, 29, 30) and fails only in 2 cases (events 2, 10).}{to the source AR is found in 9 cases (events 1, 2, 4, 5, 13, 23, 25, 27, 29) while} for 16\deleted{$^3$He-rich SEP} events with a coronal wave \replaced{, the magnetic connection is correctly predicted only for 8 cases (6, 11, 12, 14, 16, 18, 21, 26); in the remaining 8 events the PFSS model shows closed field (3, 8, 9, 19, 22) or open field that is inconsistent with the IMF polarity (17, 20, 28). The remaining 4 events with EUV waves (7, 15, 31, 32) have the AR not visible in the magnetograms}{only in 6 cases (6, 12, 20, 21, 26, 28)}. In addition none of the events with jets show closed model field but all cases with a closed configuration are found for the events with EUV waves. \added{Thus the connection to the flaring region on the Sun was generally not required for the events with a coronal EUV wave.} \replaced{Thus the model field extrapolation, with its known limitations, also}{This} suggests that in some $^3$He-rich SEP events the ion injection may be instigated by EUV wave expulsions. We found that \replaced{all previously reported}{9} events \added{in this study} with a wide longitudinal spread of $^3$He-rich SEPs are associated with large-scale coronal waves. \added{In six angle-separated events (6, 7, 11, 16, 21, 22) the spacecraft coronal foot-points without magnetic connection to the source AR were traversed by EUV waves, probably ensuring an injection of SEPs onto IMF lines. In the remaining three events (9, 17, 28), where EUV waves were not seen crossing the foot-points, the wide angular distributions could be due to other mechanizms. These have been thoroughly discussed elsewhere \citep[e.g.,][]{wie13,ric14} and may include particle cross-field diffusion (in corona or interplanetary space) and a distortion of magnetic fields by CMEs.} In summary, the multi-spacecraft EUV imaging and radio observations on angularly separated STEREOs provided the most reliable identification of $^3$He-rich SEP sources as we have ever had. This study reveals quite common ($\sim$60\%) association of $^3$He-rich SEPs with large-scale coronal EUV waves. While in LSEP events the presence of significant CMEs expanding in all directions can mimic the role of EUV waves, the relation with energetic ions can be better seen in $^3$He-rich SEP events which are often found without a CME or accompanied by weak coronal outflows. Note that half of the EUV wave events in this survey are accompanied by slow CMEs ($\lesssim$300\,km\,s$^{-1}$) or by no CME at all. | 16 | 9 | 1609.05346 |
1609 | 1609.05697.txt | We fit the spectral energy distributions (SEDs) of members of a large sample of Fermi 2LAC blazars to synchrotron and inverse Compton (IC) models. Our main results are as follows. (i) As suggested by previous works, the correlation between peak frequency and curvature can be explained by statistical or stochastic particle acceleration mechanisms. For BL Lacs, we find a linear correlation between synchrotron peak frequency and its curvature. The slope of the correlation is consistent with the stochastic acceleration mechanisms and confirm previous studies. For FSRQs, we also find a linear correlation, but its slope cannot be explained by previous theoretical models. (ii) We find a significant correlation between IC luminosity and synchrotron luminosity. The slope of the correlation of FSRQs is consistent with the EC process. And the slope of the correlation of BL Lac is consistent with the SSC process. (iii) We find several significant correlations between IC curvature and several basic parameters of blazars(black hole mass, broad line luminosity, the Lorentz factor of jet). We also find significant correlations between bolometric luminosity and these basic parameters of blazars which suggest that the origin of jet is a mixture of the mechanisms proposed by Blandford $\&$ Znajek and by Blandford $\&$ Payne. | Blazars are the most extreme form of active galactic nuclei (AGN), with their jets pointed in the direction of the observer(Urry $\&$ Padovani 1995). The typical spectral energy distribution (SED) of blazars is generally described by a double bump structure. Harvey, Wilking $\&$ Joy (1982) were the first to propose using a smooth spectrum instead of a segmented power-law spectrum when they researched the submillimeter band of 3C 345. Landau (1983) researched a sample in the centimeter, millimeter and optical bands that consisted of 9 quasi-stellar objects(QSOs) and BL Lacs. Landau et al. (1986) analyzed a sample of low-energy peaked BL Lac objects from the millimeter band to the ultraviolet band, and found that, although they could obtain spectrum successfully when combined with power-law spectrum, the data from the radio to the ultraviolet band showed good consistency with the smooth spectrum many cases. They used a quadratic function, namely: \begin{equation} log S_{\nu}=C+[(log \nu -B)^2]/2A \end{equation} to fit the spectrum. Massaro et al. (2004a) found that the log-parabolic law can well describe the X-ray peak value of SED of Mrk 421, and subsequent found that log-parabolic law can fit the X-ray spectrum(Massaro et al. 2008). There is, however, a problem in using the log-parabolic law to fit SEDs: some parameters in the equation do not have a physical explanation. It is remarkable that the SEDs of many sources can be fitted by a quadratic function that contains three parameters. The property indicates that these sources may have a similar structure, or similar relativistic particle energy distributions, or both(Landau et al. 1986). With a more general and simple situation about the particles increasing their energy, the log-parabolic law is used to statistical acceleration structure. One advantage of log-parabolic law is that it can be used to calculate various useful parameters, such as peak frequency, curvature and so on, which is simpler than the case for other models. When we obtain the value of peak frequency and peak luminosity of SEDs, the curvature of the SEDs will be another important parameter. The curvature is the property possessed by the curving of a line. If the peak frequency and peak luminosity are known, the curvature can be used to derive the bolometric luminosity. The relationship between the peak frequency and the curvature can be explained in terms of the acceleration process of emitting electrons(Massaro et al. 2004a, 2004b, 2006; Paggi et al. 2009a, 2009b; Rani et al. 2011; Tramacere et al. 2007, 2009, 2011; Chen 2014). Chen (2014) performed an analysis of the curvature-peak frequency connection, and re-derived a theoretical model that some authors had used before(Massaro et al. 2004a, 2006; Tramacere et al. 2011). He predicted two electron acceleration mechanisms according to the coefficient of the curvature-peak frequency relationship. For models of stochastic acceleration, the values of the energy-dependent acceleration probablility, the fluctuation of the fractional acceleration gain(the latter two are statistical acceleration) and the slope k(1/-c=k$log{\nu}_{p}+h$) are 2, 2.5, 3.333 respectively. The result of Chen (2014) was 2 which is consistent with stochastic acceleration. The SEDs generated by two emission components, namely the synchrotron component and inverse Compton(IC) component (Ghisellini et al. 1997; Massaro et al. 2004a, 2006). In the lepton model, it is generally accepted that low-energy peak is caused by the synchrotron emission of relativistic electrons in the jet, and that the high energy peak is caused by IC scattering (Massaro et al. 2004, 2006; Meyer et al. 2012). There is, however, disagreement concerning the origin of the soft photons of IC scattering. (1) They are derived from synchrotron emission, termed the synchrotron self-Compton(SSC) process(Rees 1967; Jones et al. 1974; Marscher $\&$ Gear 1985; Maraschi et al. 1992; Sikora et al. 1994; Bloom $\&$ Marscher 1996). (2) They are derived from the exterior of jets, termed the external Compton(EC) process. There are three possible sources of EC soft photons: accretion disk photons entering jets directly(Dermer et al. 1992; Dermer $\&$ Schlickeiser 1993); broad line region (BLR) photons entering jets(Sikora et al. 1994; Dermer et al. 1997); and dust torus infrared radiation photons entering jets (Blazejowski et al. 2000; Arbeiter et al. 2002). Ghisellini (1996) derived two relationships between the synchrotron luminosity and the IC luminosity so that it could be determined whether the IC component is dominated by the EC process or the SSC process($L_{EC}$$\sim$$L_{syn}^{1.5}$, $L_{SSC}$$\sim$$L_{syn}^{1.0}$). The bolometric luminosity is one of the most important parameters of blazars. It represents the amount of electromagnetic energy a body radiates per unit of time. Thus it represents the total radiant energy over a wide band (from the radio band to the $\gamma$ band). In research concerning the origin of jets, the current theoretical model is that the jet power is generated from accretion and the extraction of rotational energy or angular momentum from disc/black hole(Blandford $\&$ Znajek 1977; Blandford $\&$ Payne 1982). The black hole and accretion disc play an important role in the process, so it is important to research the relationships between the jet and black hole and between the jet and accretion. Ghisellini et al. (2010) found that Fermi blazars with higher luminosities may have a larger black hole. Xiong et al. (2014a) researched the relationship between the jet power and black hole mass and found that there is a significant correlation between them, which means that the the jet power is controlled by the spin of black hole. In both mechanisms, the magnetic field plays a major role in channelling power from the black hole or the disc into the jet. The process be sustained by matter accreting on to black hole, so it is reasonable that there is connection between the origin of jet and accretion(Maraschi $\&$ Tavecchio 2003). Many authors have studied the jet-accretion disk relationship, using a variety of methods(Rawlings $\&$ Saunders 1991; Falcke $\&$ Biermann 1995; Serjeant et al. 1998; Cao $\&$ Jiang 1999; Wang et al. 2004; Liu et al. 2006; Xie et al. 2007; Gu et al. 2009; Ghisellini et al. 2009a, 2009b, 2010, 2011; Sbarrato et al. 2012; Yu et al. 2015). The BLR luminosity can be taken as an indicator of the accretion power of blazars (Celotti, Padovani $\&$ Ghisellini 1997), and the bolometric luminosity can be taken as the index of jet power, so researching the relationship between the BLR luminosity and bolometric luminosity can provide information on the relationship between the jet and accretion disc, contributing to the knowledge about the origin of jets. Xie et al. (2007) found a significant correlation between $log L_{BLR}$ and $log L_{jet}$($log L_{BLR}$=$log L_{jet}$+$log {\eta}$+const). The Lorentz factor ($\Gamma$) can describe the speed of the jet flow (Hovatta et al. 2009). There is a relationship between the jet power and jet speed: more powerful jets will appear to be faster (Kharb et al. 2010). L$\ddot{u}$ et al. (2012) found that the $\Gamma$ and $\gamma$-luminosity are significant connection. We know that the $\gamma$-luminosity can represent the bolometric luminosity(Fan, Xie, $\&$ Bacon 1999, Xie et al. 2004), so we can predict that there will be a linear correlation between $\Gamma$ and bolometric luminosity. In this paper, the SEDs of both synchrotron and IC components of a sample are fitted by a log-parabolic law. The bolometric luminosity is the amount of electromagnetic energy a body radiates per unit of time,it is the total radiant energy over a wide band (from the radio band to the $\gamma$ band), we calculate the exact value of the bolometric luminosity and curvature by fitting the SEDs. Firstly, we analyse the curvature-peak frequency relation and try to verify the particle acceleration mechanism. Then, we analyse the correlation between the IC luminosity and the synchrotron luminosity and try to judge the IC component is dominated by the EC process or the SSC process. Finally, we analyse the correlations between curvature, bolometric luminosity and black hole mass, BLR luminosity and the Lorentz factor. The paper is structured as follows. In Section 2 we present the sample; In Section 3 we detail the fitting procedure; in Section 4, we present the results; In Section 5 we provide a discussion and conclusion. A $\Lambda$CDM cosmology with $H_{\rm{0}}={\rm{70Km s^{-1} Mpc^{-1}}}$, $\Omega_{\rm{m}}=0.27$, $\Omega_{\rm{\Lambda}}=0.73$ is adopted. | \subsection{The effect of the inaccurate redshift of BL Lacs in related linear correlation analysis} Redshifts are normally obtained by spectroscopic and photometric measurement. However, BL Lacs typically lack emission lines. Clear spectroscopic measurements can be obtained for only a small proportion of BL Lacs. Most redshifts of BL Lacs are obtained by photometric measurement or are based on dubious private communications. In Plotkin et al. (2008), photometric redshifts were estimated to be accurate to $\Delta$z$\approx$0.01 from comparison with spectroscopic redshifs. However, the two types of redshift measured for a few BL Lacs have a huge difference, $\Delta$z$\approx$0.5. The use of inaccurate redshifts may have a significant effect on the calculation of the luminosity and peak frequency in rest-frame. In this paper, the inaccurate redshifts may affect the correlation between $log {\nu}_{peak}^{syn}$ and -1/$c_{syn}$ and between $log L_{IC}$ and $log L_{syn}$, discussed in Sect. 4.1 and Sect. 4.2, respectively. In subsamples A and B, there are 17 BL Lacs without measured redshifts. Therefore, we discuss the effect of the inaccurate redshift for two scenarios: (i) all BL Lacs in subsamples A and B. (ii) BL Lacs with measured redshifts in subsamples A and B. On the basis of the redshifts in our sample, we add a normally distributed disturbance. The mean value of this normally distributed disturbance(${\mu}_{disturbance}$) are 10\%z, 20\%z, \textbf{30\%z, 40\%z and 50\%z}, respectively. The standard deviation($\sigma$) are 10\%$\times$0.3144, 20\%$\times$0.3144, 30\%$\times$0.3144, 40\%$\times$0.3144 and 50\%$\times$0.3144(0.3144 is the standard deviation of the redshift distribution of BL Lacs in 2LAC). We can then obtain random redshifts and calculate $log L_{syn}$, $log L_{IC}$ and $log {\nu}_{peak}^{syn}$ in rest-frame. Finally, we repeated the linear correlations 5000 times and analysed the distributions of the coefficient of correlation R, the significance level p and slope K. The results are given in Table 1 to 4. Table 1 and 2 show the distributions of R, p and K for the correlation between $log {\nu}_{peak}^{syn}$ and -1/$c_{syn}$ for all BL Lacs and for BL Lacs with a measured redshift in subsample A. From Tables 1 and 2, it can be seen that all p-values are less than 0.0001 and that all the distributions of R are normally distributed. The mean values of R(${\mu}_{R}$) in Table 1 are almost 0.701 and those of ${\mu}_{R}$ in Table 2 are almost 0.66. Furthermore, all the standard deviation of R(${\sigma}_{R}$) are small, ${\sigma}_{R}$$\leq$0.02. It can also be seen that all the distributions of K are normal. Their mean values(${\mu}_{K}$) are 1.87 and 1.89, respectively. Their standard deviations are all small, ${\sigma}_{K}$$\leq$0.01. Table 3 and 4 show the distributions of R, p and K for the correlation between $log L_{IC}$ and $log L_{syn}$ for all BL Lacs and for BL Lacs with measured redshift in subsample B. From these two table, it can be seen that almost all the p-values are less than 0.0001, but the distributions of R are not normal any more. In the correlation between $log L_{IC}$ and $log L_{syn}$, the coefficient of R is 0.901. Therefore, we want to know how many R-values range between 0.85 and 0.95 after repeating the same linear correlation analysis 5000 times. In these two tables it can be seen that fewer and fewer R-values are in the range between 0.85 and 0.95 with increasing disturbance. The smallest percentage is 61.0\%. If we expand the range to 0.8$\sim$0.95, the percentage of R for the case of the largest disturbance will rise to 83.1\%. All the distributions of slope K are normal, and their mean values(${\mu}_{K}$) are in the range between 1.05 and 1.16. Their standard deviation(${\sigma}_{K}$) are less than 0.08. In light of the above analysis, we suggest that the effect of the inaccurate redshift is small for the results obtained in Sect. 4.1 and Sect. 4.2. Even if a sample contains a small amount of BL Lacs without measured redshift, our main results will not change. The correlation between $log {\nu}_{peak}^{syn}$ and -1/$c_{syn}$ for BL Lacs can still be explained by stochastical particle acceleration mechanisms. In addition, the IC component of BL Lacs can also be considered as an SSC process. However, this conclusion only applies to our sub-sample. It is possible that, if the sample size or the disturbance of redshift increase, our results will be completely different. \begin{table*} \begin{minipage}{155mm} \centering \caption{The distribution of R, p and K in the correlation between $log {\nu}_{peak}^{syn}$ and -1/$c_{syn}$ for all BL Lacs in sub-sample A.} \begin{tabular}{@{}llrrr@{}} \hline\hline ${\mu}_{disturbance}$ & ${\mu}_{R}$ & The percentage of $p<0.0001$ & ${\mu}_{K}$ \\ & ${\sigma}_{R}$ & & ${\sigma}_{K}$ \\ \hline 50\%z & 0.700 & 100\% & 1.87 \\ & 0.01 & & 0.01 \\ 40\%z & 0.700 & 100\% & 1.87 \\ & 0.004 & & 0.01 \\ 30\%z & 0.701 & 100\% & 1.87 \\ & 0.003 & & 0.01 \\ 20\%z & 0.702 & 100\% & 1.87 \\ & 0.002 & & 0.005 \\ 10\%z & 0.702 & 100\% & 1.87 \\ & 0.001 & & 0.002 \\ \hline \end{tabular} \end{minipage} \end{table*} \begin{table*} \begin{minipage}{155mm} \centering \caption{The distribution of R, p and K in the correlation between $log {\nu}_{peak}^{syn}$ and -1/$c_{syn}$ for BL Lacs with measured redshift in sub-sample A.} \begin{tabular}{@{}llrrr@{}} \hline\hline ${\mu}_{disturbance}$ & ${\mu}_{R}$ & The percentage of $p<0.0001$ & ${\mu}_{K}$ \\ & ${\sigma}_{R}$ & & ${\sigma}_{K}$ \\ \hline 50\%z & 0.65 & 100\% & 1.89 \\ & 0.02 & & 0.04 \\ 40\%z & 0.66 & 100\% & 1.89 \\ & 0.01 & & 0.03 \\ 30\%z & 0.67 & 100\% & 1.89 \\ & 0.01 & & 0.02 \\ 20\%z & 0.67 & 100\% & 1.89 \\ & 0.01 & & 0.01 \\ 10\%z & 0.66 & 100\% & 1.89 \\ & 0.003 & & 0.01 \\ \hline \end{tabular} \end{minipage} \end{table*} \begin{table*} \begin{minipage}{155mm} \centering \caption{The distribution of R, p and K in the correlation between $log L_{IC}$ and $log L_{syn}$ for all BL Lacs in sub-sample B.} \begin{tabular}{@{}llrr@{}} \hline\hline ${\mu}_{disturbance}$ & the percentage of R $\in$(0.85, 0.95) & the percentage of $p<0.0001$ & ${\mu}_{K}$ \\ & & & ${\sigma}_{K}$ \\ \hline 50\%z & 70.0\% & 99.7\% & 1.05 \\ & & & 0.07 \\ 40\%z & 74.9\% & 99.9\% & 1.06 \\ & & & 0.06 \\ 30\%z & 82.2\% & 100\% & 1.07 \\ & & & 0.05 \\ 20\%z & 92.9\% & 100\% & 1.09 \\ & & & 0.04 \\ 10\%z & 98.5\% & 100\% & 1.12 \\ & & & 0.03 \\ \hline \end{tabular} \end{minipage} \end{table*} \begin{table*} \begin{minipage}{155mm} \centering \caption{The distribution of R, p and K in the correlation between $log L_{IC}$ and $log L_{syn}$ for BL Lacs with measured redshift in sub-sample B.} \begin{tabular}{@{}llrr@{}} \hline\hline ${\mu}_{disturbance}$ & the percentage of R $\in$(0.85, 0.95) & the percentage of $p<0.0001$ & ${\mu}_{K}$ \\ & & & ${\sigma}_{K}$ \\ \hline 50\%z & 61.0\% & 95.1\% & 1.09 \\ & & & 0.08 \\ 40\%z & 66.0\% & 97.3\% & 1.10 \\ & & & 0.08 \\ 30\%z & 77.7\% & 99.3\% & 1.11 \\ & & & 0.07 \\ 20\%z & 90.9\% & 100\% & 1.13 \\ & & & 0.06 \\ 10\%z & 99.5\% & 100\% & 1.16 \\ & & & 0.04 \\ \hline \end{tabular} \end{minipage} \end{table*} \subsection{Particle acceleration mechanisms for FSRQ and BL Lac} The correlation between $log {\nu}_{peak}^{syn}$ and -1/$c_{syn}$ can be explained in two different scenarios, namely the statistical(energy-dependent acceleration probability and fluctuation of fractional acceleration gain) and the stochastical acceleration mechanisms. FSRQ and BL Lac, the two different blazar subclasses have many differences. These include(Giommi et al. 2012b): (i) different optical spectral; (ii) different extended radio powers; (iii) very different redshift distributions; (iv) different cosmological evolutions; (v) widely different mix of FSRQs and BL Lac objects in radio and X-ray selected samples; (vi) widely different distributions of the synchrotron peak energy ${\nu}_{peak}^{S}$. Therefore, in this paper, we collect a large sample and separate them into FSRQs(N=200) and BL Lacs(N=79) in order to study the linear correlations between $log {\nu}_{peak}^{syn}$ and -1/$c_{syn}$, respectively. For BL Lacs, the slope of the correlation($k_{BL Lac}$=1.87$\pm$0.19) is consistent with the stochastical acceleration mechanisms. However, for FSRQs, the slope of the correlation($k_{FSRQ}$=3.69$\pm$0.24) has a big difference. It is not consistent with any theoretical values(k=5/2, 10/3 and 2) and cannot be explained by the two particle acceleration mechanisms(Chen 2014). Chen (2014) used a sample of 43 blazars in order to study the linear correlation between $log {\nu}_{peak}^{syn}$ and -1/$c_{syn}$. The slope of the correlation was 2.04$\pm$0.03, which is consistent with the stochastic acceleration mechanisms. The number of objects in the sample was too small to separate them into FSRQs and BL Lacs. Morever, the correlation was based on only eight HSP blazars. Perhaps this is why he did not find different slopes between FSRQs and BL Lacs. As noted above, the slope of BL Lacs is consistent with the results of Chen (2014) and can be explained by stochastic particle accelerations. For FSRQs, the slope of the correlation $k_{FSRQ}$=3.69$\pm$0.24 is close to 10/3, which can be explained by statistical particle acceleration for the case of fluctuation of fractional acceleration gain. Tramacere et al. (2011) pointed out that the statistic method does not give a complete physical description of the processes responsible for the systematic and stochastic energy gain, as it ignores various physical processes. Thus the slope of the correlation that we found for FSRQs implies that some physical processes may cause error in statistic method. Even so, from the perspective of particle acceleration mechanisms, FSRQs are different from BL Lacs. As far as we know, in terms of the particle acceleration mechanisms, this is the first time that a difference between FSRQs and BL Lacs has been found by using a large sample. Our results can provide an observational information that is relevant to particle acceleration models. \subsection{EC vs SSC and Curvature of FSRQ} A typical SED of blazars displays two peaks. The high-energy component is usually explained as arising from IC scattering of the same electrons as produce the synchrotron emission. Depending on the origin of the soft photons, the IC scattering can be divided into EC and SSC components. According to Ghisellini et al. (1996), there are two relations that can determine whether the high-energy component is dominated by EC or SSC($L_{EC}$$\sim$$L_{syn}^{1.5}$, $L_{SSC}$$\sim$$L_{syn}^{1.0}$). Here we study the correlation between $log L_{IC}$ and $log L_{syn}$ for FSRQs and BL Lacs in our subsamples. For FSRQs, the slope of the best linear regression is 1.45$\pm$0.11; this suggests that the high-energy component of FSRQs is dominated by the EC process(Ghisellini et al. 2002; Celotti $\&$ Ghisellini 2008; Ghisellini et al. 2010; Finke 2013). The result of BL Lacs($k_{BL Lac}$=1.12$\pm$0.10) suggest that the high-energy component is dominated by the SSC process(Zhang et al. 2013; Lister et al. 2011; Celotti $\&$ Ghisellini 2008; Ghisellini et al. 2002; Ghisellini et al. 2010; Ackermann et al. 2012). When characterizing the two components of the SEDs of blazars, the curvature is another important parameter. It can represent the value of bolomeric flux/luminosity if the peak frequency and peak flux/luminosity are known. By studying the correlation between $log L_{BLR}$ and -1/c, we found that there is no correlation for synchrotron component but that the IC component shows a significant correlation. The synchrotron component is generally explained by synchrotron emission from relativistic electrons in a jet(Maraschi, Ghisellini $\&$ Celotti 1992; Massaro et al. 2006; Hovatta et al. 2009), which means that there is no correlation between $log L_{BLR}$ and -1/$c_{syn}$. The origin of soft photons in IC scattering is complex. In one of the EC models, the soft photons could come from the BLR(Sikora et al. 1994; Dermer et al. 1997). The significant correlation between $log L_{BLR}$ and -1/$c_{IC}$ that we found might suggest that the soft photons of IC scattering are indeed mainly from the BLR. In addition, we studied the correlation between -1/c and the $log L_{bol}$, $log M_{BH}$ and $\Gamma$. Coincidentally, all these parameters have a significant correlation with -1/$c_{IC}$ and no correlation with -1/$c_{syn}$. Perhaps there is a deep-seated physical signification, or maybe it is just a statistical coincidence. \subsection{Jet of FSRQ} In current theoretical models of the formation of jet, power is generated through accretion and the extraction of rotational energy of disc/black hole(Blandford $\&$ Znajek 1977; Blandford $\&$ Payne 1982) and is then converted into the kinetic power of jet. Bolometric luminosity is one of the most important parameters of blazars and can function as an index of the jet power(Du et al. 2016). The broad-line luminosity can be taken as an indicator of accretion power(Celotti, Padovani $\&$ Ghisellini 1997). Black holes will be spun up through accretion, as these objects acquire mass and angular momentum simultaneously through accretion (Chai, Cao $\&$ Gu 2012). The presence of the jet implies that the gravitational potential energy of the falling matter not only can be transformed into heat and radiation, but can also amplify the magnetic field, allowing the field to access the large store of black hole rotational energy and transform part of it into the mechanical power of the jet, as discusses by Ghisellini et al. (2014). These authors predicted that jet power is depended on $(aMB)^2$, where a and M are respectively the spin and mass of the black hole and B is the magnetic field at its horizon. In our work, the correlations between $log L_{bol}$ and $M_{BH}$, $L_{BLR}$ are significant. Our results show that the jet is correlated both with black hole and the accretion disc, which suggest that the origin of jet is a mixture of the mechanisms proposed by Blandford $\&$ Znajek and by Blandford $\&$ Payne(Punsly $\&$ Coroniti 1990; Meier et al. 1999, 2001). Furthermore, the coefficient of the $L_{BLR}$$\sim$$L_{bol}$ correlation is 0.9032, very close to 1. This result is consistent with Xie et al. (2007)(($log L_{BLR}$=$log L_{jet}$+$log {\eta}$+const)). The correlation between $log L_{bol}$ and $\Gamma$ can reflect the $P_{power}^{jet}$$\sim$$\Gamma$ relationship. Our result shows that the correlation between $log L_{bol}$ and $\Gamma$ is significant, and thus the correlation between $P_{power}^{jet}$ and $\Gamma$ is significant, too. From the results concerning $L_{bol}$$\sim$$\Gamma$ and $L_{bol}$$\sim$$M_{BH}$, we found that more powerful jets, the plasma blob in jets will move faster(Kharb et al. 2010; Wu et al. 2011; L$\ddot{u}$ et al. 2012) and have larger black hole masses(Xiong et al. 2014a). In other words, FSRQs with a larger black hole mass have faster jets($M_{BH}$$\sim$$\Gamma$). According to Maraschi $\&$ Tavecchio 2003($\delta$$\sim$$\Gamma$), the Doppler factor($\delta$) is also the indicator of the jet speed, so $M_{BH}$$\sim$$\delta$ can be used to represent $M_{BH}$$\sim$$\Gamma$. In Arshakian et al. (2005), the significant correlation between $M_{BH}$ and $\delta$ was found using a sample of 12 objects. Torrealba et al. (2008) obtained the same result using the 15 objects. | 16 | 9 | 1609.05697 |
1609 | 1609.05941_arXiv.txt | { We study the correlation between far-infared/submm dust emission and atomic gas column density in order to derive the properties of the high Galactic latitude, low density, Milky Way cirrus in the foreground of the Virgo cluster of galaxies. Dust emission maps from 60 to 850~$\mu$m are obtained from SPIRE observations carried out within the {\it Herschel} Virgo Cluster Survey, complemented by IRAS-IRIS and {\em Planck}-HFI maps. Data from the Arecibo legacy Fast ALFA Survey is used to derive atomic gas column densities for two broad velocity components, low and intermediate velocity clouds. Dust emissivities are derived for each gas component and each far-infared/submm band. For the low velocity clouds, we measure an average emissivity $\epsilon^\mathrm{LVC}_\nu = (0.79\pm0.08) \times 10^{-20}$ MJy sr$^{-1}$ cm$^2$ at 250~$\mu$m. After fitting a modified blackbody to the available bands, we estimated a dust absorption cross-section $\tau^\mathrm{LVC}_\nu/{N_\ion{H}{i}} = (0.49\pm0.13) \times 10^{-25}$ cm$^2$ H$^{-1}$ at 250~$\mu$m (with dust temperature $T=20.4\pm1.5$K and spectral index $\beta=1.53\pm0.17$). The results are in excellent agreement with those obtained by {\em Planck} over a much larger coverage of the high Galactic latitude cirrus (50\% of the sky vs 0.2\% in our work). For dust associated with intermediate velocity gas, we confirm earlier {\em Planck} results and find a higher temperature and lower emissivity and cross-section. After subtracting the modelled components, we find { regions at scales smaller than 20$\arcmin$ where the residuals deviate significantly from the average, cosmic-infrared-background dominated, scatter. These large residuals} are most likely due to local variations in the cirrus dust properties (and/or the dust/atomic-gas correlation) or to high-latitude molecular clouds with average $N_{\mathrm{H}_2} \lesssim 10^{20}$ cm$^{-2}$. { We find no conclusive evidence for intracluster dust emission in Virgo. } } | Diffuse far-infrared (FIR) emission from high-Galactic latitude dust (the Milky Way {\em cirrus}) is an important test bed for models of interstellar grains. The thinness of the Galactic disk, our peripheric position in it, and a viewing direction away from the Galactic plane offer clear advantages to modelling: the main source of dust heating is provided by the ambient radiation \citep[the Local Interstellar Radiation Field - LISRF;][]{MathisA&A1983}; the diffuse radiation gradients perpendicular to the disk being small \citep[as shown by radiative transfer models of galactic disks; ][]{BianchiA&A2000b,BocchioA&A2013,PopescuMNRAS2013}; thus the mixing of dust at different temperatures (and/or with different properties) along a single line of sight is limited. Starting from the first detection of the cirrus emission by the IRAS satellite \citep{LowApJL1984}, a tight correlation of the FIR surface brightness with the atomic gas column density was observed \citep{BoulangerApJ1988}. The cirrus surface brightness per \ion{H}{i} column density (a quantity known as {\em emissivity}) measured in the 60 and 100~$\mu$m IRAS bands was then used to verify the predictions or constrain the properties of early interstellar grain models \citep{DraineApJ1984,DesertA&A1990}. With the advent of the instruments aboard the COBE satellite, and in particular the FIRAS spectrophotometer, it was possible to study the Spectral Energy Distribution (SED) of the cirrus up to the peak of thermal emission and beyond. \citet{BoulangerA&A1996} found a tight correlation with \ion{H}{i} column density and measured the emissivity up to 1~mm. After fitting a modified blackbody (MBB) to the data, the dust absorption cross-section per unit \ion{H}{i} column density was retrieved, which was found to be in agreement with the available dust models. Since then, the FIRAS emissivity has been one of the major constraints { to dust models, which were formulated} to reproduce its SED levels and shape, resulting in an average dust cross-section $\propto \nu^\beta$, with $\beta\approx 1.8-2$ \citep[see, e.g., ][]{DraineARA&A2003,ZubkoApJS2004,CompiegneA&A2011,JonesA&A2013}. The recent analysis of data from the {\em Planck} satellite, however, revealed an inconsistency between {\em Planck} and FIRAS, which prompted a recalibration on planet fluxes of the 350 and 550~$\mu$m data from the High Frequency Instrument \citep[HFI; ][]{Planck2013VIII}. As a result, the newly determined emissivities have been found to be lower than the FIRAS values, and with a reduced dust cross-section spectral index \citep[$\beta\approx 1.6$;][]{Planck2013XI,PlanckIntermediateXVII}. { These results suggest that a re-evaluation of the dust model parameters may be required,} and in particular of the grain size distributions and relative contribution of grains of different chemical composition to the SED at 100~$\mu$m$ \la \lambda \la$ 1 mm. Furthermore, the higher resolution of {\em Planck} data with respect to FIRAS highlighted local dust emissivity variations with the environment \citep{Planck2013XI,PlanckIntermediateXVII}. A tenuous cirrus is seen in the foreground of the Virgo cluster of galaxies ($l=283.8^\circ, b=74.4^\circ$) through scattered starlight in deep optical \citep{RudickApJ2010,MihosProc2015} and UV images \citep{CorteseMNRAS2010,BoissierA&A2015}; it corresponds to low extinction, with $0.02 < E(B-V) < 0.1$ \citep{BoissierA&A2015}. Its emission is clearly detected at 250, 350 and 500~$\mu$m in images obtained by the SPIRE instrument \citep{GriffinA&A2010} aboard the {\em Herschel} Space Observatory \citep{PilbrattA&A2010}, as part of the {\it Herschel} Virgo Cluster Survey \citep[HeViCS; ][]{DaviesA&A2010,DaviesMNRAS2012,AuldMNRAS2013}. In this work, we will derive the dust emissivity of the HeViCS cirrus in the SPIRE bands. This will allow to verify the latest {\em Planck} results, since the SPIRE calibration also is based on planet models. Besides, the 250~$\mu$m band will fill the gap between the IRAS and {\em Planck} data, and complement the lower sensitivity and resolution COBE-DIRBE emissivity at 240~$\mu$m \citep{PlanckIntermediateXVII}. When compared to the {\it Planck} large area estimates, the relatively small sky coverage of the HeViCS field will also highlight local emissivity variations. The emissivities will be derived using $\ion{H}{i}$ observations from the Arecibo Legacy Fast ALFA survey \citep[hereafter ALFALFA][]{GiovanelliAJ2005}. The high resolution of the ALFALFA data (FWHM=3$\farcm$5) is comparable to that of the longest wavelength data dominated by dust thermal emission (the 850~$\mu$m {\em Planck}-HFI band, with 4$\farcm$8); it is also a factor two better than other $\ion{H}{i}$ data used to derive dust emissivities: previous large-area studies \citep{BoulangerA&A1996,Planck2013XI} used observations at FWHM=0.6$^\circ$ from the Leiden/Argentine/Bonn (LAB) Survey of Galactic $\ion{H}{i}$ \citep{KalberlaA&A2005}. The FWHM=14$\farcm$5 GASS $\ion{H}{i}$ survey of the southern sky has been used in \citet{PlanckIntermediateXVII}, while \citet{PlanckEarlyXXIV} exploited Green Bank Telescope maps at FWHM=9$\farcm$1. Only in the recent work of \citet{ReachApJ2015} Arecibo telescope data from the Galactic ALFA \citep[GALFA; ][]{PeekApJS2011} survey have been used to study the correlation between $\ion{H}{i}$ column density of isolated high-latitude clouds and {\em Planck}-based dust column densities. { The resolution of the ALFALFA survey thus allows a characterization of the residuals to the dust-$\ion{H}{i}$ correlation at smaller scales. } In fact, one of the aims of HeViCS is the detection of dust emission from the intracluster medium (ICM), once the structure of the foreground cirrus is subtracted using Galactic $\ion{H}{i}$ as a template. { The presence of the cirrus is indeed the limiting factor in these studies: so-far the only claim of the detection of FIR emission from the ICM is that on the Coma cluster by \citet{StickelA&A1998}, using data from the ISOPHOT instrument aboard the ISO satellite. However, the analysis of five other clusters with analogous data and techniques yielded no detection \citep{StickelA&A2002}; a following analysis of Coma using data from the {\em Spitzer} satellite dismissed the putative emission as cirrus contamination \citep{KitayamaApJ2009}. } This paper is organised as follows: in Sect.~\ref{sec:data} we present the ALFALFA and HeViCS data, together with other ancillary maps used in this work. In Sect.~\ref{sec:ana} we describe the method used to derive the correlation between dust and $\ion{H}{i}$. The fitted emissivities are presented in Sect.~\ref{sec:res}, while Sect.~\ref{sec:sed} is dedicated to the derivation of the mean opacity cross section from the emissivity SEDs. The possible Galactic origin of the largest residuals is discussed in Sect.~\ref{sec:resi}, while in Sect.~\ref{sec:icd} we put constrains on the dust emission from the Virgo ICM. The work is summarised in Sect.~\ref{sec:sum}. | 16 | 9 | 1609.05941 |
|
1609 | 1609.02085_arXiv.txt | Magnetic fields significantly influence the evolution of protoplanetary disks and the formation of planets, following the predictions of numerous magnetohydrodynamic (MHD) simulations. However, these predictions are yet observationally unconstrained. To validate the predictions on the influence of magnetic fields on protoplanetary disks, we apply 3D~radiative transfer simulations of the polarized emission of aligned aspherical dust grains that directly link 3D global non-ideal MHD simulations to ALMA observations. Our simulations show that it is feasible to observe the predicted toroidal large-scale magnetic field structures, not only in the ideal observations but also with high-angular resolution ALMA observations. Our results show further that high angular resolution observations by ALMA are able to identify vortices embedded in outer magnetized disk regions. | Magnetic fields play a crucial role in the formation and evolution of protoplanetary disks \citep{2011ARA&A..49...67W, 2011ARA&A..49..195A}. Magnetic fields influence the transport of dust and gas \citep[e.g.,][]{2007ApJ...654L.159C, 2014prpl.conf..411T, 2014prpl.conf..339T}, the disk chemistry \citep[e.g.,][]{2011ApJS..196...25S, 2013ChRv..113.9016H}, and even the migration of plane\-tesimals and planets within the disk \citep[e.g.,][]{2010A&A...515A..70D, 2011MNRAS.415.3291G, 2011ApJ...736...85U} through magnetohydrodynamic (MHD) turbulence. Moreover, MHD turbulence can provide the source of viscosity that drives the accretion \citep{1974MNRAS.168..603L}, and thus, the evolution of the disk \citep{1973A&A....24..337S}. One of the most promising mechanisms for driving turbulence, respectively accretion is the magneto-rotational instability \citep[MRI;][]{1991ApJ...376..214B, 1996ApJ...467...76B, 1998RvMP...70....1B}. Turbulence generated by pure hydrodynamical instabilities \citep{2013MNRAS.435.2610N, 2014ApJ...789...77L, 2014ApJ...788...21K} is not able to explain the observed accretion rates and disk lifetimes \citep{2011ARA&A..49...67W}, despite such instabilities could become important in disk regions with low ionization \citep{2014prpl.conf..411T}. High enough ionization, and so magnetic activity, is expected in the inner disk regions due to thermal ionization \citep{ume88,des15} and also in the outer disk regions due to interstellar radiation \citep[e.g.,][]{2007ApJ...659..729T,2013ApJ...765..114D, 2015ApJ...799..204C, 2015A&A...574A..68F}. However, even the prediction of any magnetic field structure by observational constraints remains a difficult task. First polarimetric observations of protoplanetary disks performed with SMA, CARMA, and VLA indicate toroidal magnetic field structures \citep{2015ApJ...814L..28C, 2014ApJ...780L...6R, 2013ApJ...769L..15S}. Spatially resolved observations of polarized millimeter continuum emission of aligned aspherical dust grains are well suited to reveal the magnetic field structure in the protoplanetary disk \citep[e.g.,][]{2000prpl.conf..247W, 2007ApJ...669.1085C}. Though, polarimetry is strongly influenced by many factors \citep{Bertrang2016}. Taking the dust grain shape, dust grain alignment, magnetic field properties, the resolution of the observation, and the projection along the line of sight into account, we present a feasibility study on high-angular resolution polarimetric observations and their interpretation to validate MHD predictions on magnetic fields in protoplanetary disks. | We present a feasibility study on high-angular resolution polarimetric observations and their interpretation to validate MHD predictions on magnetic fields in protoplanetary disks. Based on the assumption of perfect grain alignment by the magnetic field, we compute an upper limit for the polarized dust emission for aspherical grains. This is consistent with previous work \citep{2007ApJ...669.1085C} which was performed with less sophisticated models and lower spatial resolution. Our main results are: \begin{itemize} \item In order to take radial and vertical magnetic field components correctly into account, a 3D~disk model is vital for the simulation and interpretation of the polarized emission. \item We find that the polarized signal which arises from aspherical grains aligned by the magnetic field does trace the magnetic field topology in both ideal observations and simulated ALMA observations. It has to be noted that the high-angular resolution of ALMA is critical for resolving the polarization pattern. Unresolved structure in the polarization pattern leads to annihilation of the signal. \item Observations of the polarized emission of protoplanetary disk with ALMA are able to reveal the toroidal magnetic field structure. We have shown that it is even possible to observe small-scale deviations from the toroidal structure which could appear inside vortices located in the outer protoplanetary disk regions. \item At an inclination of $45^{\circ}$ we find that small-scale structures in the magnetic field, such as a vortex, can only be observed with very high resolutions. \item For disks at lower inclination, we find polarization patterns corresponding to the observations of protoplanetary disks (class~I) that have been performed so far. Detailed comparisons with observations of class II disks await future detections. \item Gaps and jumps in the dust density distributions by magnetic activity as well as vertical magnetic field components in edge-on disks are still unresolved by the polarimetric observations. \end{itemize} | 16 | 9 | 1609.02085 |
1609 | 1609.02566_arXiv.txt | { We apply photo-dissociation region (PDR) molecular line emission models, that have varying degrees of enhanced mechanical heating rates, to the gaseous component of simulations of star-forming galaxies taken from the literature. Snapshots of these simulations are used to produce line emission maps for the rotational transitions of the CO molecule and its \thco~ isotope up to $J = 4-3$. We use these maps to investigate the occurrence and effect of mechanical feedback on the physical parameters obtained from molecular line intensity ratios. We consider two galaxy models: a small disk galaxy of solar metallicity and a lighter dwarf galaxy with 0.2 \zsun metallicity. Elevated excitation temperatures for CO($1 - 0$) correlate positively with mechanical feedback, that is enhanced towards the central region of both model galaxies. The emission maps of these model galaxies are used to compute line ratios of CO and \thco~ transitions. These line ratios are used as diagnostics where we attempt to match them These line ratios are used as diagnostics where we attempt to match them to mechanically heated single component (i.e. uniform density, Far-UV flux, visual extinction and velocity gradient) equilibrium PDR models. We find that PDRs ignoring mechanical feedback in the heating budget over-estimate the gas density by a factor of 100 and the far-UV flux by factors of $\sim 10 - 1000$. In contrast, PDRs that take mechanical feedback into account are able to fit all the line ratios for the central $< 2$ kpc of the fiducial disk galaxy quite well. The mean mechanical heating rate per H atom that we recover from the line ratio fits of this region varies between $10^{-27}$ -- $10^{-26}$~erg s$^{-1}$. Moreover, the mean gas density, mechanical heating rate, and the $A_V$ are recovered to less than half dex. On the other hand, our single component PDR model fit is not suitable for determining the actual gas parameters of the dwarf galaxy, although the quality of the fit line ratios are comparable to that of the disk galaxy. } | Most of the molecular gas in the universe is in the form of H$_2$. However, this simple molecule has no electric dipole moment. The rotational lines associated with its quadrupole moments are too weak to be observed at gas temperatures less than 100~K, where star formation is initiated inside clouds of gas and dust. This is also true for the vibrational and electronic emission of H$_2$; hence it is hard to detect directly in the infrared and the far-infrared spectrum. CO is the second most abundant molecule after H$_2$, and it has been detected ubiquitously. CO forms in shielded and cold regions where H$_2$ is present. Despite its relatively low abundance, it has been widely used as a tracer of molecular gas. \cite{solomon75} were the first to establish a relationship between CO$(1 - 0)$ integrated intensity ($W_{\rm CO(1-0)}$) and H$_2$ column density ($N({\rm H}_2$)). Since then, this relationship has been widely used and it is currently known as the so-called $X$-factor. The applicability and limitations of the $X$-factor are discussed in a recent review by \cite{bolatto2013-1}. Environments where cool H$_2$ is present allow the existence of CO, and many other molecular species. In such regions, collisions of these molecules with H$_2$ excite their various transitions, which emit at different frequencies. The emission line intensities can be used to understand the underlying physical phenomena in these regions. The line emission can be modeled by solving for the radiative transfer in the gas. One of the most direct ways to model the emission is the application of the large velocity gradient (LVG) approximation \citep{sobolev1960}. LVG models model the physical state of the gas such as the density and temperature but do not differentiate among excitation mechanisms of the gas, such as heating by shocks, far-ultraviolet (FUV), or X-rays and cosmic rays, hence do not provide information about the underlying physics. The next level of complexity involves modeling the gas as equilibrium photo-dissociation regions, PDRs \citep{tielenshb1985,Hollenbach1999,rolling07}. These have been successfully applied to star forming regions and star-bursts. However, modeling of Herschel and other observations for, e.g., NGC 253, NGC 6240 and M82, using these PDRs show that other heating source rather than FUV are required to reproduce observational data. In particular, such heating source can be identified in AGN or enhanced cosmic ray ionization \citep[][among many others]{maloney96, komossa03, martin06-1, papadopoulos10, meijerink13-1, Rosenberg14}, or mechanical heating due to turbulence \citep{loenen2008, pan2009-1, aalto13}. The latter is usually not included in ordinary PDR models and is the focus of this paper. Various attempts have been made in this direction in modeling star-forming galaxies and understanding the properties of the molecular gas. However, because of the complexity and resolution requirements of including the full chemistry in the models, self-consistent galaxy-scale simulations have been limited mainly to CO \citep{Kravtsov2002, wada05, cubick2008-1, Narayanan2008-1, inti2009-1, xu2010, perez11, narayanan2011b, shetty2011, Feldmann2012, Narayanan2013-1, olsen16-1}, but see also e.g. \citep{olsen15-1} for an effort to model $[\rm{C}II]$. The rotational transitions of CO up to $J=4 -3$ predominantly probe the properties of gas with densities in the range of 10$^2$ - 10$^5$~\cmt, and with temperatures from $\sim 10$K to $\sim 50$K. Higher $J$ transitions probe denser and warmer molecular gas around $\sim 200$K for the $J = 10-9$ transition. In addition to high-$J$ CO transitions, low-$J$ transitions of high density tracers such as CS, CN, HCN, HNC and HCO$^+$, are good probes of cold gas with $n \sim 10^6$~\cmt. Having a broad picture on the line emission of these species provides a full description of the thermal and dynamical state of the dense gas (where strong cooling and self-gravity dominate). Thus, potentially unique signatures of turbulent and cosmic ray/X-ray heating may lie in the line emission of these species in star-forming galaxies. In \cite{mvk12, mvk13-a} we studied the effect of mechanical feedback on diagnostic line ratios of CO, \thco~ and some high density tracers for grids of mechanically heated PDR models in a wide parameter space relevant to quiescent disks as well as turbulent galaxy centers. We found that molecular line ratios for CO lines with $J \le 4 -3$ are good diagnostics of mechanical heating. In this paper, we build on our findings in \cite{mvk13-a} to apply the chemistry models to the output of simulation models of star forming galaxies, using realistic assumptions on the structure of the ISM on unresolved, sub-grid, scales. mode to construct CO and \thco~ maps for transitions up to $J = 4 - 3$. Our approach is similar to that by \cite{perez11} where the sub-grid modeling is done using PDR modeling that includes a full chemical network based on \cite{umist1999}, which is not the case for the other references mentioned above. The main difference of our work from \cite{perez11} is that our sub-grid PDR modeling takes into account the mechanical feedback in the heating budget; on the other hand we do not consider X-ray heating effects due to AGN. The synthetic maps are processed in a fashion that simulates what observers would measure. These maps are used as a guide to determine how well diagnostics such as the line ratios of CO and \thco~ can be used to constrain the presence and magnitude of mechanical heating in actual galaxies. In the method section we start by describing the galaxy models used, although our method is generally applicable to other grid and SPH based simulations. We then proceed by explaining the procedure through which the synthetic molecular line emission maps were constructed. In the results section we study the relationship and the correlation between the luminosities of CO, \thco, and H$_2$. We also present maps of the line ratios of these two molecules and see how mechanical feedback affects them, how well the physical parameters of the molecular gas can be determined, when the gas is modeled as a single PDR with and without mechanical feedback. In particular, we try to constrain the local average mechanical heating rate, column density and radiation field and compare that to the input model. We finalize with a discussion and conclusions. | In this paper, we have presented a method to model molecular species emission such as CO and \thco~for any simulation which includes gas and provides a local mechanical heating rate and FUV-flux. The method uses this local information to model the spatial dependence of the chemical structure of PDR regions assumed to be present in the sub-structure of the gas in the galaxy models. These are then used to derive brightness maps assuming the LVG approximation. The method for the determination of the emission maps together with the ISM model of the hydrodynamic galaxy models constitute a complete, self-consistent model for the molecular emission from a star forming galaxies' ISM. We compute the emission of CO and \thco~ in rotational transitions up to $J=4-3$ for a model disk-like and a dwarf galaxy. From the emission maps line ratios were computed in order to constrain the physical parameters of the molecular gas using one component PDR models. We conclude the following: \begin{enumerate} \item Excitation temperature correlates positively with mechanical feedback in equilibrium galaxies. This in turn increases for gas which is closer to the center of the galaxy. The analysis presented in this paper allows estimates of mechanical feedback in galaxies which have high excitation temperatures in the center such as those observed by, e.g., \cite{muhle2007} and \cite{israel2009-1}. \item Fitting line ratios of CO and \thco~ using a single mechanically heated PDR component is sufficient to constrain the local \gm~throughout the galaxy within an order of magnitude, given the limitations of our modeling of the emission. The density of the gas emitting in CO and \thco~ is better constrained to within half dex. This approach is not suitable in constraining the gas parameters of the dwarf galaxy, although the statistical quality of the fits is on average better than that of the disk galaxy. \item Our approach fails in constraining the local FUV-flux. It is under-estimated by two orders of magnitude in mechanically heated PDR fits, and over-estimated by more than that in PDRs, which do not account for mechanical feedback. This discrepancy is due considering one PDR component, where $ J < 3 - 2$ CO lines are optically thick and trace very shielded gas. This discrepancy can also be explained by looking at the right column of Figure-\ref{fig:paper3_dist-app}, where most of the emission results from SPH particles, which constitute a small fraction, and the rest of the gas is not captured by the PDR modeling. \end{enumerate} The main features of our PDR modeling with mechanical feedback is the higher number of degrees of freedom it allows compared to LVG models and multi-component PDR modeling. We fit 10 line ratios while varying four parameters in the case where \gm~was considered, and 3 parameters when \gm~was not. We note that the actual number of {\it independent} measurements of line emission of CO and \thco~up to $J = 4-3$ is 8. In LVG modeling the number of free parameters is at least 5, since in fitting the line ratios we need to vary $T$, $n({\rm H}_2$), $N$(CO), N(\thco) and the velocity gradient. On the other hand, PDR models assume elemental abundances and reaction rates. Moreover, the uncertainties in the PDR modeling could be much larger and depend on the micro-physics used in the PDR modeling \citep{vasyunin04, rolling07}. Thus, LVG models have more free parameters compared to single PDR models used in our method. On the other hand, considering two PDR components increases the number of free parameters to 9. This renders the fitting problem over-determined with more free parameters than independent measurements. Although LVG models are simple to run compared to PDR models, they do not provide information about the underlying physical phenomena exciting the line emission, which renders them useless for constraining \gm. Although a two component PDR model produces better ``Xi-by-eye'' fits, these fits are statistically less significant than a single PDR model, but physically relevant. The main disadvantage of PDR modeling is the amount of bookkeeping required to run these models and they are computationally more demanding than LVG modeling\footnote{A typical PDR model requires 10 time more CPU time than an LVG model in our simulations}, which is why we resorted to using interpolation tables. \subsection{Improvements and prospects}\label{subsec:improvementsandprospects} Our models have two main limitations: 1) we make implicit assumptions on the small scale structure of the galactic ISM and 2) we assume chemical equilibrium. The first limitation can be improved by performing higher spatial resolution simulations, which would result in gas that is better resolved and is denser than the current maximum of $10^4$~\cmt in the disk galaxy. Gas densities $> 10^6$~\cmt allows us to consider transitions up to $J = 15 - 14$ which are more sensitive to \gm~compared to the $J=1-0$. Galaxy scale simulations, which reach densities higher than $10^4$~\cmt~have been performed by \cite{Narayanan2013-1} and required about $3 \times 10^7$ particles, and our method could be easily scaled to post-process such simulations. In a follow up paper \citep{mvk16-2} we re-sampled the gas density distribution of the disk simulation used in this paper in-order to probe the effect of $n \gtrsim 10^4$~\cmt gas on the emission of $ 4-3 < J \le 15-14$ transitions of CO and \thco, in addition to $J \le 7-6$ transitions of HCN, HNC, and \hcop and their associated diagnostic line ratios on galactic scales. Having more data to fit helps in finding better diagnostics of \gm~ and narrower constraints for it. Another advantage in having emission maps for these high $J$ transitions is having a higher number of degrees of freedom in fitting for the emission line ratios. This renders the best fit PDR models statistically more significant. Moreover, it would be natural to consider multi-components PDR models in these fits, a dense-component fitting the $J > 4 - 3$ transitions and low-density component fitting the lower ones. These components are not independent, low-$J$ emission is also produced by the high density PDR components. The high density component PDR would generally have a filling factor 10 to 100 times lower than the low density component, thus the high density model contributes mainly to the high-$J$ transitions, whereas the low density models contributes mainly to the low-$J$ emission and less to the high-$J$ emission. In producing the synthetic line emission maps, we have used the LVG emission from SPH particles, which assume PDR models as the governing sub-grid physics. The main assumption in computing the emission using LVG models was the fixed value of the micro-turbulence line-width $v_{\rm turb} = 1$~km/s for all the SPH particles. The peak of the distribution of $v_{\rm turb}$ of the SPH particles is located at $\sim 3$~km/s. The optical depth in the lines are thus in reality smaller than what we used in the calculations for the paper. When the lines become optically thick it effectively reduces the critical density of those transitions, and allows the excitation to higher energy states. The lines become more easily optically thick, for normal cloud sizes ($A_{\rm V} = 5 - 10$) at $v_{\rm turb} = 1$~km/s, which causes the peak of the CO ladder to be at a higher rotational transition. In order to quantify the shift of the peak of the CO ladder, we calculated a grid with different turbulent velocities, $v_{\rm turb} = 1.0, 2.7, 5.0$, and 10.0 km/s, representative of the turbulent velocities in the SPH simulation, and used that grid to produce the resulting CO ladder. We found that the peak of the CO ladder is located at CO(4-3) transition, when computed from the distribution of $v_{turb}$, while the peak is located at CO(6-5) for the calculations where only $v_{\rm turb} = 1.0$~km/s was used. Although this causes a significant quantitative change in the emission, it does not affect the general conclusions of the paper. A possible consequence of considering $v_{\rm turb}$ to be constant is the fact that the \thco/CO line rations increase towards the center of our model galaxies and are close to the elemental abundance ratio of $^{12}{\rm C}/^{13}{\rm C}$ $\sim 40$ in contrast to observations where this ratio ranges between 8 and 15 \citep[e.g.][and references therein]{Buchbender13-01}. The emission of CO are enhanced towards the center due to the increasing temperature of the gas. Accounting for $v_{\rm turb}$ would most likely enhance the relative emission of CO compared to \thco~where the computed line-ratios of the model galaxies would be in-line with the observed trends and could also reduce $^{12}{\rm C}(1-0)/^{13}{\rm C}(1-0)$ to the observed range of, which is currently a limitation to our proposed modelling approach. An accurate treatment of the radiative transfer entails constructing the synthetic maps by solving the 3D radiative transfer of the line emission using tools such as LIME \citep{lime}. This would eliminate any possible bias in the fits. Currently, we construct the emission maps by considering the mean flux of the emission of SPH particles modeled as PDR within a pixel. This results in a distribution of the emission as a function of gas density as seen in Figure-\ref{fig:paper3_dist} and Figure-\ref{fig:paper3_dist-app}. When fitting the line ratios, we recover the mean physical parameters of the molecular gas in that pixel associated with these emission. Using tools such as LIME, the micro-turbulence line width would be treated in a more realistic way by using the actual local velocity dispersions in the small scale turbulent structure of the gas, as opposed to adopting a fiducial 1 \kms as we have done throughout the paper. A major limitation of our sub-grid modelling is the assumption that the FUV flux at the surface of the PDR corresponds to that of the environment of the SPH particle. The FUV flux and the FUV-driven heating is strongly attenuated by the intervening dusty CNM, and the H$_2$ clouds embedded in it. This is in contrast to turbulent and cosmic ray heating that operate over large gas volumes. An accurate treatment of the FUV radiative transfer could influence the modeled abundances of the molecular species and consequently the computed line emission fluxes. that are used as diagnostics. Not all H$_2$ gas is traced by CO, since the H/H$_2$ transition zone occurs around A$_V = 0.1$~mag whereas the C/CO transition occurs around A$_V = 1$~mag. When looking at Figure-\ref{fig:paper3_dist} and Figure-\ref{fig:paper3_dist-app}, we see that most of the CO emission emanates from a fraction of the gas. The remaining ``major'' part of the gas can contain molecular H$_2$, which is not traced by CO, but by other species such as C or C$^+$, that do have high enough abundances to have bright emission \citep{madden97-1, papadopoulos02-1, wolfire2010-1, pineda2014-1, israel15}. Finally, we note that we have ignored any ``filling factor'' effects in constructing the synthetic maps and the fitting procedure. This is not relevant to our method since in considering line ratios, and filling factors cancel out in single PDR modeling of line ratios. However, it is important and should be taken into account, when considering multi-component PDR modeling and when fitting the absolute luminosities. | 16 | 9 | 1609.02566 |
1609 | 1609.07944_arXiv.txt | {The RV~Tauri stars constitute a small group of classical pulsating stars with some dozen known members in the Milky Way. The light variation is caused predominantly by pulsations, but these alone do not explain the full complexity of the light curves. High quality photometry of RV~Tau-type stars is very rare. DF~Cygni is the only member of this class of stars in the original \textit{ Kepler} field, hence allowing the most accurate photometric investigation of an RV~Tauri star to date.} {The main goal is to analyse the periodicities of the RV~Tauri-type star DF~Cygni by combining four years of high-quality \textit {Kepler} photometry with almost half a century of visual data collected by the American Association of Variable Star Observers.} {\textit{Kepler} quarters of data have been stitched together to minimize the systematic effects of the space data. The mean levels have been matched with the AAVSO visual data. Both datasets have been submitted to Fourier and wavelet analyses, while the stability of the main pulsations has been studied with the O--C method and the analysis of the time-dependent amplitudes.} {DF~Cygni shows a very rich behaviour on all time-scales. The slow variation has a period of 779.606 d and it has been remarkably coherent during the whole time-span of the combined data. On top of the long-term cycles the pulsations appear with a period of 24.925 d (or the double period of 49.85 d if we take the RV Tau-type alternation of the cycles into account). Both types of light variation significantly fluctuate in time, with a constantly changing interplay of amplitude and phase modulations. The long-period change (i.e. the RVb signature) somewhat resembles the Long Secondary Period (LSP) phenomenon of the pulsating red giants, whereas the short-period pulsations are very similar to those of the Cepheid variables. Comparing the pulsation patterns with the latest models of Type-II Cepheids, we found evidence of strong non-linear effects directly observable in the \textit{Kepler} light curve. } {} | The RV~Tauri stars are evolved low-mass F, G and K-type pulsating supergiants that are located above the Population II Cepheids in the instability strip. Traditionally, they are classified into two subgroups. The RVa type is characterised by alternating minima with typical periods longer than 20 days. The RVb stars show an additional long-term variation in the mean brightness, with typical periods of 700-1200 days. The short-period variation is interpreted with fundamental mode pulsation, while the long-term phenomena is commonly interpreted as being caused by periodic obscuration of a binary system by circumbinary dust disk (Lloyd Evans 1985, Pollard et al. 1996, Van Winckel et al. 1999, Fokin 2001, Maas et al. 2002, Gezer et al. 2015). The RV~Tauri stars have similar luminosities, but higher effective temperatures than Miras. The luminosity function of RV~Tauri stars mainly overlaps with the low-luminosity part of the Mira luminosity function. There are similarities between their observational characteristics. A few Miras show double maxima which is common in RV~Tauri stars (e.g. R~Cen, R~Nor). Some Miras and semiregular variables may exhibit the same (or physically similar) quasi-periodic, long-term mean brightness variations as do RVb stars (e.g. RU~Vir; Willson and Templeton 2009). Also, there is an extensive literature on the so-called Long Secondary Periods of the Asymptotic Giant Branch stars, a phenomenon that still puzzles the researchers (Wood et al. 1999; Takayama, Wood and Ita 2015) and which causes light curves that are surprisingly similar to those of the RVb stars. Originally, the RV~Tauri stars were classified as Cepheid variable due to the similarities and the poor quality of the light curves. The Type II Cepheids may be divided in groups by period, such that the stars with periods between 1 and 5 days (BL~Her class), 10–20 days (W~Vir class), and greater than 20 days (RV~Tauri class) have different evolutionary histories (Wallerstein, 2002). In each cases the shape of the light curve is nearly sinusoidal, but the RV~Tau-type stars show alternating minima (which means that every second minimum is shallower). The irregularity of the minima grows as the period becomes longer. The light curve of some RV~Tau-type stars randomly switches into a low-amplitude irregular variation then switches back into the previous state (e.g. AC~Her (Kolláth et al. 1998) and R~Sct (Kolláth 1990; Buchler et al. 1996)). The RV~Tauri variables have been placed among the post-AGB stars (Jura 1986), that are rapidly evolving descendants of stars with initial masses lower than 8 M$_{\odot}$. In the late stages of their evolution they are crossing the Hertzsprung–Russell diagram (HRD) from the cool asymptotic giant branch (AGB) to the ionizing temperature of the planetary nebula nuclei. During this process they cross the classical instability strip, in which large-amplitude radial oscillations are driven by the $\kappa$-mechanism, while the complexity of the post-AGB variability pattern depends on their location in the HRD relative to the classical instability strip (Kiss et al. 2007). Several models attempted to reproduce the light variations of RV~Tauri stars. The short period variation can be easily explained by summing two sine curves if their frequencys' ratio is 1:2 and the phase difference is $\pi$/2 (Pollard et al. 1996). Buchler et al. (1996) and Kolláth et al. (1998) established a more sophisticated model to describe the light curve of R~Sct and AC~Her by low-dimensional chaos. Due to the several phenomena that occur in the light curves and spectra of RVb stars, it is difficult to explain the typical long-term variation. A long-term photometric and a spectroscopic survey was done by Pollard et al. (1996, 1997). They revealed that in some RVb stars, the reddest colours occur slightly after long-term light minimum. Furthermore, the light and colour amplitude of the short-term period is smaller during the long-term minima. They found that the equivalent width of the H$\alpha$ emission lines in the spectrum are varying with the phase of the long-term period. They concluded that the damping of the pulsation amplitude is difficult to be explained by the popular model of light variation of RVb stars by a binary system which is periodically obscured by circumstellar or circumbinary dust disc (Percy 1993; Waelkens and Waters 1993; Willson and Templeton 2009). They also concluded that the H$\alpha$ emission is caused by passing shock waves throughout the stellar photosphere. Pollard et al. (2006) proposed a possible dust-eclipse model which arrangement can explain both the photometric and spectroscopic characteristics of the long-term phenomena. Unfortunately, the quality of the measurements did not allow to investigate the pulsation mechanism in details, except the case of R Sct and AC Her. Gezer et al. (2015) studied the Spectral Energy Distribution (SED) of Galactic RV~Tauri stars based on WISE infrared photometry. The objects with circumstellar disks have near-IR excess in the SED. The light curve of the members of this group shows variable mean magnitude, while there is a clear correlation between disk sources and binarity. Consequently, binarity is connected to the long-term changes of the mean brightness. On the other hand, both Gezer et al. (2015) and Giridhar et al. (2005) investigated the photospheric chemical anomaly called depletion, as a sign of dust-gas separation. Two scenarios were proposed to objects with anomalous abundances: (1) single stars with dust-gas separation in their stellar wind and (2) binary stars where dust-gas separation is present in a circumbinary disk (Giridhar et al. 2005). The latter is consistent with the RVb binary hypothesis. Gezer et al. (2015) found that the presence of a disk seems to be a necessary but not sufficient condition for the depletion process to become efficient. The subject of this paper, the bright RVb-type variable (V$_{\rm max}\approx 10.5$ mag, V$_{\rm min}\approx 13$ mag) DF~Cygni was discovered by Harwood (1927). The period was found to be 49.4 days between the principal minima. Some years later a long period of 790$\pm$10 days was found (Harwood 1936, 1937). The period of the radial pulsation is about 50 days, the long secondary period is about 775 days (Percy 2006). The spectral and the color variation was investigated by Preston et al. (1963) who found CN bands in the spectrum. Gezer et al. (2015), in their search for disk sources, labelled DF~Cygni as uncertain based on its SED, while the star only appears as marginally depleted, if at all (Giridhar et al. 2005, Gezer et al. 2015). Everything put together, DF~Cygni can be considered as one of the better known RV~Tauri-type variable stars, and the only member of this class in the original Kepler field. Here we present a detailed light curve analysis of DF~Cygni by combining about 48 years of visual observations with the ultra-precise space photometry from \textit{Kepler}, with a time-span of about 4 years. The two sources of the data allow investigations both with very high frequency resolution and with extreme photometric accuracy at least over two cycles of the RVb variability. In Sect. 2 we describe the data sources and the preparation of the sets before the analyses. The details of the light curve analysis are presented in Sect. 3, with the discussion of results in Sect. 4. A brief summary is given in Sect. 5. | How typical an RV~Tau-type star is DF~Cygni? To shed light on the answer, we plotted DF~Cygni's location in the V-band amplitude vs. effective temperature and the luminosity vs. effective temperature diagrams of selected post-AGB variables studied by Kiss et al. (2007). For the effective temperature we adopted $T_{\rm eff}=$4840 K (Brown et al. 2011; Giridhar et al. 2005), while the V-band amplitude was approximated by the mean peak-to-peak amplitude of the phase diagram in Fig. 5, which is about 1.0$\pm$0.2 mag. The luminosity has been estimated using the RV~Tauri period-luminosity relation in the Large Magellanic Cloud, following the same approach as Kiss et al. (2007). The estimated luminosity is $1200\pm300L_\odot$, where the uncertainty reflects the standard deviation of the LMC P-L relation. In the top panel of Fig. 9, DF~Cygni falls close to several high-amplitude stars, all being typical RV~Tau-type objects (such as U~Mon and AI~Sco, both in the RVb class). Similarly good agreement is found in the empirical Hertzsprung-Russell diagram in the bottom panel of Fig. \ref{hrd}. Although DF~Cygni lies somewhat beyond the red edge of the classical instability strip, its location is consistent with its classification. Hence its light curve properties are likely to be representative for the whole class. \begin{figure} \centering \includegraphics[width=8cm]{teff_amp_df.eps} \includegraphics[width=8cm]{hrd_df.eps} \caption{{\it Top panel:} V-band amplitude as a function of effective temperature for the Kiss et al. (2007) sample of post-AGB variables, supplemented with the location of DF~Cygni. Open circles: single periodic stars; triangles: multiperiodic/semiregular stars; squares: variability due to orbital motion. The bar in the upper left corner shows the typical cycle-to-cycle amplitude variation in the pulsating stars. {\it Bottom panel:} The empirical HRD of the pulsating sample of Kiss et al. (2007) and the location of DF~Cygni. The dashed lines show the edges of the classical instability strip, taken from Christensen-Dalsgaard (2003). The error bars in the lower right corner represent $\pm$3\% error in effective temperature (about 200 K in the range shown) and the $\pm$0.35 mag standard deviation of the LMC P–L relation.} \label{hrd} \end{figure} DF~Cygni shows a very rich behaviour on all time-scales. The slow variation has a period of 779.606 d and it has been remarkably coherent during the whole 48 years of visual observations. On top of the long-term cycles the pulsations appear with a period of 24.925 d (or the double period of 49.85 d if we take the RV Tau-type alternation of the cycles into account). Both types of light variation significantly fluctuate in time, with a constantly changing interplay of amplitude and phase modulations. In the Fourier-spectrum of the {\it Kepler} dataset, a characteristic series of subharmonics of the frequency {\it f} appears, which is often interpreted as sign of period-doubling associated with an underlying low-dimensional chaotic behaviour. Similar phenomena were detected in several long-period pulsating variable stars, such as the arch-type of chaotic stars, the RV Tau-type R~Sct (Buchler, Serre \& Koll\'ath 1995, Buchler et al. 1996), the less irregular RVa star AC~Her (Kolláth et al. 1998), several semi-regular variables (Buchler, Kolláth \& Cadmus 2004) and one Mira-type variable (Kiss \& Szatmáry 2002). More recently, \textit{Kepler} has opened a whole new avenue of RR~Lyrae studies based on the period-doubling phenomenon and related effects (e.g. Szabó et al. 2010, Plachy et al. 2013, 2014, Benk\H{o} et al. 2014, Moskalik et al. 2015). Considering RV~Tau-type variability, there is a strong need for more theoretical investigations. The only recent study that touches the domain in which DF~Cygni resides is that of Smolec (2016), who studied a grid of non-linear convective Type II Cepheid models. Although his study does not cover the full parameter range of RV~Tau-type stars, some of the models extend close to the temperature and luminosity of DF~Cygni. For example, he presented a detailed discussion of the resonances, which affect the shape of the model light curves and in his fig. 13 one can see the regions of the 2:1 resonances of the fundamental, first and second overtone modes. The corresponding Fourier-parameters are shown in figs. 14-15 of Smolec (2016), where DF~Cygni's location is just about outside at the top of the high-luminosity edge of the more massive model calculations. Here, near the red edge of the instability strip, the resonance between the fundamental mode and the first overtone is dominant (see the V-shaped dark features in figs. 13-15, in the rightmost panels with the M=0.8 M$_\odot$ models). The observed complexity of DF~Cygni's pulsations is consistent with the strongly non-linear behaviour of the models. As we have seen in Sect. 3, all the periods, amplitudes and phases vary in time, and none of these variations is strictly periodic. As has been revealed by the O--C diagram, there was a transient episode in the light curve around BJD 2456000, where the change of the period and the amplitude was remarkable. Right after the ascending branch of the RVb cycle, the amplitude of the pulsation increased by almost a factor of two, whereas the period decreased quite dramatically (approximately by 3-4 percent, that is about 1 day per cycle). After four to six cycles, the pulsation returned to the previous state. The fact that the sudden amplitude increase is associated with a period decrease is surprising, given that the typical period-amplitude correlation of classical pulsating stars is just the opposite. Interestingly, Smolec (2016) studied the differences between the linear and non-linear pulsation models and found that the relative difference between the non-linear and the linear periods can be up to 15\%, with a strong dependence in the sign across the instability strip. We find a striking feature in fig. 6 of Smolec (2016): models near the location of DF~Cygni show a strong and negative period difference, meaning that the non-linear models have periods that are shorter than those of the linear calculations. We may speculate that it is not a coincidence that DF~Cygni's location and the region where the non-linear pulsation period is significantly smaller agree so well. If the high-amplitude--short-period transient of the light curve is caused by the pulsations becoming more non-linear then the Smolec (2016) models offer a natural explanation to the observed phenomenon. So that the sudden change in the light curve characteristics perhaps can be explained by the emergence of non-linear effects. Both the variations of the phase diagram and the point-to-point scatter of the O--C diagram (Figs. \ref{Fig:phase_diagram}-\ref{Fig:O--C}) are much unlike the stability of the Cepheid-type variables. We note, however, that recent space observations of Type I Cepheids have found small cycle-to-cycle fluctuations (Derekas et al. 2012), which depend on the mode of pulsation, the first overtone being less stable than the fundamental mode (Evans et al. 2015). Although Type II Cepheids have not yet been covered extensively by space data, one can imagine the outlines of a progression of becoming more unstable on the time-scale of the pulsations as we move closer to the non-linear regime of RV~Tau-type stars. Besides the properties of the pulsations, the nature of the RVb-phenomenon, the long-term change of the mean brightness, is also worth some discussion. The most common explanation of the RVb-phenomenon is a binary system which is periodically obscured by a circumstellar or a circumbinary dust disc (Waelkens \& Waters 1993, Pollard et al. 1996, 1997, Van Winckel et al. 1999, Maas et al. 2002), although some of the observed characteristics were difficult to reconcile with the dusty disk model (Pollard et al. 1996, 1997). The most recent infrared results from the WISE satellite (Gezer et al. 2015) indicate that RVb stars are exclusively those objects that have well-detected circumstellar disks, while there is also a clear correlation between disk sources and binarity. The 48-years coherence of the long-period variability of DF~Cygni is indeed indicative of a stable mechanisms like binary motion. We find noteworthy the similarities between the RVb-type variability and the Long Secondary Periods (LSPs) of red giant stars, the latter still representing a mystery (Nicholls et al. 2009 and references therein, Soszynski \& Wood 2013, Saio et al. 2015, Takayama et al. 2015). Among the suggested solutions we find binarity, variable extinction in circumstellar disks or strange non-radial oscillations, {many of them} somewhat resembling the proposed explanations of the RVb phenomenon. LSP-like variations have also been found in pulsating red supergiants (Kiss et al. 2006, Yang \& Jiang 2012), suggesting that the phenomenon may not be restricted to the red giants. We find some correlations between the pulsations and the RVb-cycle; most notably the amplitude of the pulsations tends to be smaller in the faint states of DF~Cygni, although the case is not entirely clear. It is also interesting to note that van Aarle et al. (2011) cross-correlated their sample of candidate post-AGB stars with Long Period Variables (LPV) from MACHO. They found 245 variables falling on the distinct period-luminosity relation of the Long Secondary Periods. While some of those stars may not be genuine post-AGB objects, it is an interesting questions if there is some connection between the LSPs of red giants and post-AGB variability. As has been pointed out by the referee, with what we know today, it seems most logical to conjecture that these red giants with long cycles are also binaries, but with longer orbital periods than the RV Tauri stars, so that the same mass-transfer phenomena occurred later in the evolution of the primary, which could climb further on the AGB. | 16 | 9 | 1609.07944 |
1609 | 1609.09732_arXiv.txt | \citet{opher15}, \citet{drake15} have shown that the heliospheric magnetic field results in formation of two-jet structure of the solar wind flow in the inner heliosheath, i.e. in the subsonic region between the heliospheric termination shock and the heliopause. In this scenario the heliopause has a tube-like topology as compared with a sheet-like topology in the most models of the global heliosphere \citep[e.g.][]{izmod_alexash15}. In this paper we explore the two-jet scenario for a simplified astrosphere in which 1) the star is at rest with respect to the circumstellar medium, 2) radial magnetic field is neglected as compared with azimuthal component, \added{3) the stellar wind outflow is assumed to be hypersonic (both the Mach number and the Alfv\'enic Mach number are much greater than unity at the inflow boundary). We have shown that the problem can be formulated in dimensionless form, in which the solution depends only on one dimensionless parameter $\varepsilon$ that is reciprocal of the Alfv\'enic Mach number at the inflow boundary. This parameter is proportional to stellar magnetic field. We present the numerical solution of the problem for various values of $\varepsilon$. Three first integrals of the governing ideal MHD equations are presented, and we make use of them in order to get the plasma distribution in the jets. Simple relations between distances to the termination shock, astropause and the size of the jet are established. These relations allow us to determine the stellar magnetic field from the geometrical pattern of the jet-like astrosphere.} | First models of the stellar/solar wind (SW) interaction with the interstellar medium (ISM) were developed by \citet{parker61}. Parker has considered three problems 1) the solar wind outflow into the homogeneous interstellar gas at rest, 2) the solar wind outflow into the interstellar gas moving with subsonic speed, and 3) the solar wind outflow into the interstellar magnetic field. Later, \citet{baranov70} considered a model of the solar wind interaction with supersonic interstellar wind. The structure of the interaction in the latter model has three discontinuities (Fig.~\ref{astrosphere_sketches}a): 1) the termination shock (TS) that decelerates the stellar wind from supersonic to subsonic, 2) the heliopause/astropause that is a tangential discontinuity (TD) separating the stellar wind flow from the interstellar medium, 3) the bow shock (BS) that decelerates the supersonic interstellar flow from the supersonic regime to the subsonic one. If the interstellar flow is subsonic or subalfv\'enic in MHD case then the bow shock maybe absent e.g. \citep[e.g.][]{izmod2009,mccomas2012}. During last $\sim$45 years the models of SW/ISM have been significantly developed. Modern models are three dimensional and time dependent, they take into account the multi-component nature of both SW and ISM, the effects of magnetic fields, interstellar neutrals, energetic particles. For details, see, for example, reviews and papers by \citet{zank15}, \citet{opher15}, \citet{izmod_alexash15}. What is common in all modern models is the sheet-like topology of the heliopause (see~Fig.~\ref{astrosphere_sketches}a). In 2015 \citet{drake15} and \citet{opher15} have shown that the heliopause may in fact have a tube-like shape (Fig.~\ref{astrosphere_sketches}b). \citet{opher15} have obtained such a shape in their numerical 3D MHD code for the case when the interstellar gas flows with respect to the star. Later this result has been discussed by \citet{pogorelov15} and \citet{izmod_alexash15} and needs to be explored further. \begin{figure*} \includegraphics[width=0.45\textwidth]{sheet_like_heliosphere_sketch.pdf} \includegraphics[width=0.45\textwidth]{tube_like_heliosphere_sketch.pdf} \caption{Schematic picture of the heliospheric/astrospheric interface with a sheet-like topology (a) of tangential discontinuity (TD) and a tube-like topology (b).} \label{astrosphere_sketches} \end{figure*} \citet{drake15} has considered a simpler case, when the solar wind flows into the homogeneous interstellar gas at rest. This case is quite identical to one of \citet{parker61}. The only difference is in the heliospheric magnetic field that has not been taken into account by Parker. To understand the effects of the stellar (heliospheric) magnetic field qualitatively, let us start with \citet{parker61} model. There is a shock transition in his solution (i.e. the termination shock) at $R_{TS} \sim \sqrt{\frac{\dot{M}V_0}{4\pi p_{\infty}}}$, where $R_{TS}$ is the heliocentric distance to the termination shock, $\dot{M}$ is the stellar mass loss rate, $V_0$ is the terminal velocity of the supersonic stellar wind, $p_{\infty}$ is the interstellar gas pressure. In the supersonic SW (for $R<R_{TS}$) the solution is $V \sim V_0$, $\rho \sim 1/R^2$ and $p \sim 1/R^{2\gamma}$, where $R$ is the distance to the Sun or star. In the subsonic region ($R>R_{TS}$) the gas may be considered incompressible and the solution is $V \sim 1/R^2 $, $\rho \sim \rho_{\infty} $ and $p \sim p_{\infty}$. This solution can be used to calculate the frozen-in magnetic field in the kinematic approximation. Solving $\nabla \times [\mathbf{V} \times \mathbf{B}]=0$ and assuming that magnetic field is parallel to the velocity vector at the Sun: \begin{equation} R<R_{TS}: B_R \sim 1/R^2, \quad B_\phi \sim (1/R)\sin\theta, B_\theta = 0 \label{parker_field} \end{equation} \citep{parker58}; \begin{equation} R > R_{TS}: \quad B_R \sim 1/R^2, \; B_\phi \sim R\sin\theta, B_\theta = 0. \end{equation} Here $\theta$ is the angle counted from the stellar rotational axis, $\phi$ is the azimuthal angle. In the subsonic wind the magnetic field grows proportionally to $r = R\sin\theta $ that is the distance to the axis of stellar rotation. Alfv\'enic Mach number $M_A = \sqrt{ (4 \pi \rho V^2)/B^2} \sim 1/(R^3 sin \theta)$ in the subsonic wind, so it decreases with the distance rapidly. In the the supersonic wind $M_A$ remains constant. For the Sun, for example, the constant is about 15. At the strong shock $M_A$ decreases by factor of $((\gamma-1)/(\gamma+1))^{3/2}$ that is equal to $1/8$ for $\gamma=5/3$. Downstream the TS $M_A \approx 15/8\approx1.9$ and then it decreases in the equatorial plane as $1/R^3$. The Alfv\'enic Mach number becomes on the order of unity at the distances of $\sim1.23 R_{TS}$. Therefore at these distances and further one can expect a strong influence of the magnetic field on the plasma flow. Magnetic force $\mathbf{F}_\text{mag} = ([\nabla \times \mathbf{B}] \times \mathbf{B})/(4\pi)$ has a main component in $r$-direction (in cylindrical ($z$, $r$, $\phi$) coordinate system; $z$-axis is the axis of stellar rotation). As a result, the stellar wind flow deflects from the original radial direction and flows along the stellar rotation axis ($z$-axis). Therefore the two-jets structure of the flow is formed. In this paper we further explore the solution of the problem. Section~2 gives the mathematical formulation of the considered problem in dimensionless form. We have shown that the solution only depends on one dimensionless parameter. In Section~3 we have performed a theoretical study of the problem: this section presents three first integrals of the governing MHD equations. These integrals allow us to reduce the initial system of partial differential equations (PDE) to the system of algebraic equation \added{at the stagnation point} in the equatorial plane and to the system of ordinary differential equations (ODE) in jets far from the star. Section~4 presents the results of parametric numerical study of the problem. Section~5 gives summary and discusses problems remaining for future work. | In the present paper we consider a model of the stellar wind interaction with the ISM that is at rest with respect to the star. The two jets along the stellar rotation axis are formed due to the effects of the azimuthal stellar magnetic field in the model. The main results of the present paper can be summarized as following: \begin{itemize} \item It is shown that under the assumption of the hypersonic stellar wind outflow ($M_E \gg 1$,$M_{A,E} \gg 1$) the considered problem has only one dimensionless parameter, $\varepsilon$, which is inversely proportional to the Alfv\'en Mach number. This parameter increases linearly with the increasing stellar magnetic field. \item The three first integrals (\ref{Bernoulli})-(\ref{entrop}) of the MHD equations (\ref{continuity})-(\ref{rotbb}) allow us to establish analytical (or semi-analytical -- in the form of algebraic equations) relations between three parameters -- (1) the distance to the termination shock, $R_{TS,0}$, (2) the distance to the astropause, $R_{TD,0}$ (both are in the stellar equatorial plane i.e. the plane perpendicular to the axis of the stellar rotation), and (3) the parameter $\varepsilon$. For a given value of $R_{TS,0}$ one can obtain $R_{TD,0}$ as a function of $\varepsilon$. However this solution is not self-consistent. \item The distribution of the plasma parameters in the jet as well as the size of the jet have been obtained as a solution of an ODE under assumptions of the hypersonic stellar wind outflow and the spherically symmetric termination shock. One should \emph{a priori} give the distance to the TS in order to obtain this solution. Therefore it is not a self-consistent solution. \item Using the solutions described above, we propose a method that allows to estimate the magnitude of the stellar magnetic field from the geometrical picture of a two-jet astrosphere. In particular, the knowledge (for example, obtained from observations) of two ratios -- $R_{TD}/R_{TS}$ in the equatorial plane and $r_{jet}/R_{TS}$ -- allows us (see Figure~\ref{combined_contour_plot}) to determine the dimensionless quantities $\hat{R}_{TS}$ and $\varepsilon$. Then, knowing $\varepsilon$ and actual dimensional distance to the TS one can derive the magnitude of the stellar magnetic field \added{at any given distance from the star}. \item The numerical solution of the MHD equations (\ref{continuity})-(\ref{rotbb}) allowed us to establish the functional dependences \added{of $R_{TD,0}(\varepsilon)$ and $r_{jet}(\varepsilon)$} (see~ eqs.~(\ref{RTD_fit})~and~(\ref{rjet_fit})). \item We have performed the numerical parametric study by varying the parameter $\varepsilon$ in the range from $0.01$ to $0.5$. The details of the numerical solution are shown on Figures~\ref{2d_0.01}-\ref{1d_postshock_z}. It is interesting to note that there is a good agreement between analytical/semi-analytical and numerical results for plasma parameter distribution in the jet, as well as for \added{$R_{TD,0}$ and $r_{jet}$}. All discrepancies are due to approximations that lay on the base of ODE problem formulation (hypersonic pre-shock flow and spherically symmetric TS); these approximations were used for the sake of simplicity and could, in principle, be relaxed. \end{itemize} In conclusion we have to note that we do not present the results of our numerical calculations for $\varepsilon < 0.01$ and for $\varepsilon > 0.5$, because we are not completely sure that they are correct. Vortex flows beyond the termination shock are formed in the both cases. The extension of the presented parametric study will be elaborated in the future. We have noticed in Subsection 3.2 that the system of ODE in jet region has a remarkable property that may allow us in principle to determine $R_{TS}$ as a function of $\varepsilon$ even without solving the full problem. We plan to further elaborate this point in future work. Finally in this paper we restricted ourselves to a very limiting case when the interstellar medium is at rest with respect to the star. This two-jet solution can, in principle, be generalized by adding the interstellar flow. Let us consider an arbitrary plane perpendicular to $z$-axis. This plane cuts a circle from the astropause. In the case of subsonic ISM flow we can consider planar potential solutions around circles for each plane. According to the d'Alembert paradox the force acting on each circle is zero. Therefore the tube of the astropause should not be deflected into the tail, although the circle could be deformed to the ellipsoidal shape in the self-consistent solution. This scenario works\added{, if we consider the interstellar flow to be ideal and incompressible}. \added{However numerical results \citep[e.g.][]{opher15} show some bending of the jets toward the tail. This bending in numerical models (for slow incompressible ISM flows) is connected with the numerical dissipation inherent in the numerical schemes.} Numerical viscosity may cause the boundary layer breakage on the surface of the astropause, that consequently causes the pressure asymmetry that deflects the astropause. This may be the explanation for the fact, that the tube of the astropause is always deflected to the tail in the numerical models. \added{The described above numerical effects have nothing to do with physical dissipation phenomena responsible for the bending of real astrospheres. The physical dissipation mechanisms (e.g. magnetic reconnection, finite resistivity, Hall effects) should be explored as a possible cause of the astropause bending in the models with slow subsonic ISM flow. For the fast supersonic ISM flow, the bow shock is formed around the astropause. The post-shock ISM flow is vortical, and the d'Alembert paradox does not work in this case. Therefore the bending of the astrospheric jets into the tail direction is easier to justify for the supersonic relative ISM/SW motion.} Another important aspect that strongly influences the solar wind plasma flow is the interaction with the interstellar atoms due to the charge exchange. The charge exchange provides additional momentum to the plasma towards the tail as it was first shown by \citet{baranov93}. See also a recent paper by \citet{izmod_alexash15} for the self-consistent kinetic-MHD model where both heliospheric and interstellar magnetic fields are taken into account. Additional important aspect of the considered problem is the stability of the obtained two-jet solution and the tangential discontinuity (e.g. astropause). Theoretical considerations elaborated in the papers by \citet{baranov92} and \citet{ruderman93, ruderman95} could be applied here. | 16 | 9 | 1609.09732 |
1609 | 1609.08996_arXiv.txt | We have studied the distribution patterns of lateral density, arrival time and angular position of Cherenkov photons generated in Extensive Air Showers (EASs) initiated by $\gamma$-ray, proton and iron primaries incident with various energies and at various zenith angles. This study is the extension of our earlier work \cite{Hazarika} to cover a wide energy range of ground based $\gamma$-ray astronomy with a wide range of zenith angles ($\le 40^\circ$) of primary particles, as well as the extension to study the angular distribution patterns of Cherenkov photons in EASs. This type of study is important for distinguishing the $\gamma$-ray initiated showers from the hadronic showers in the ground based $\gamma$-ray astronomy, where Atmospheric Cherenkov Technique (ACT) is being used. Importantly, such study gives an insight on the nature of $\gamma$-ray and hadronic showers in general. In this work, the CORSIKA 6.990 simulation code is used for generation of EASs. Similarly to the case of Ref.\cite{Hazarika}, this study also revealed that, the lateral density and arrival time distributions of Cherenkov photons vary almost in accordance with the functions: $\rho_{ch}(r) = \rho_{0}\;e^{-\beta r}$ and $t_{ch}(r) = t_{0}e^{\Gamma/r^{\lambda}}$ respectively by taking different values of the parameters of functions for the type, energy and zenith angle of the primary particle. The distribution of Cherenkov photon's angular positions with respect to shower axis shows distinctive features depending on the primary type, its energy and the zenith angle. As a whole this distribution pattern for the iron primary is noticeably different from those for $\gamma$-ray and proton primaries. The value of the angular position at which the maximum number of Cherenkov photons are concentrated, increases with increase in energy of vertically incident primary, but for inclined primary it lies within a small value ($\le 1^\circ$) for almost all energies and primary types. No significant difference in the results obtained by using the high energy hadronic interaction models, viz., QGSJETII and EPOS has been observed. | The primary objective of the $\gamma$-ray astronomy is to detect $\gamma$-rays from celestial sources. For this purpose, in the ground based $\gamma$-ray astronomy, the Atmospheric Cherenkov Technique (ACT) \cite{Ong, Hoffman, Weekes, Lorenz, Holder0, Funk, Degrange, Goumard, Bhat, Acharya, Holder} is being used most widely within its operational energy range from around hundred GeV to tens of TeV. This technique is based on detection of Cherenkov photons emitted in the Extensive Air Showers (EASs) that are generated due to the interaction between the primary $\gamma$-rays and air nuclei. The $\gamma$-ray sources also emit Cosmic Rays (CRs), which are deflected by the intragalactic magnetic fields because of their charge and hence they loose their directional property. Whereas $\gamma$-rays, being neutral, they retain their direction of origin. Thus, by the detection of $\gamma$-rays one can make an estimate of the positions of those astrophysical objects. Because of the indirect nature of experiments in ACT as well as due to the presence of huge CR background, a complete Monte Carlo simulation study on atmospheric Cherenkov photons needs to be carried out for the detection of $\gamma$-rays with proper estimation of their energy from the observational data of such experiments. It is to be noted that, although both $\gamma$-ray and CR can generate EAS, the nature of two are different. EAS generated by $\gamma$-ray is purely electromagnetic (EM) in nature, whereas it is an admixture of EM and hadronic cascades in the case of CR. Although many studies have already been done, specially on the lateral density and arrival time distributions of Cherenkov photons in EASs using available simulation techniques \cite{Bardan, Chitnis, Hillas, Lafebre, Nerling}, still it would be a worthy task to have detailed studies on angular distributions as well as on lateral density and arrival time distributions of Cherenkov photons, initiated by $\gamma$-ray and hadronic particles, incident at various zenith angles with a wide range of energy, particularly at high altitude observation levels. Keeping this point in mind, in this work we have studied the lateral density, arrival time and angular distributions of Cherenkov photons in EASs at different energies and zenith angles over a high altitude observation level, using two different high energy hadronic interaction models, viz., QGSJETII and EPOS with FLUKA low energy hadronic interaction model available in the CORSIKA simulation package \cite{Heck}. This is the extension of our earlier work \cite{Hazarika} to cover a wide energy range of ground based $\gamma$-ray astronomy with a wide range of zenith angles ($\le 40^\circ$) of primary particles, and to study the angular distribution patterns of Cherenkov photons in EASs. CORSIKA is a four dimensional detailed Monte Carlo simulation code developed to study the evolution and various properties of EASs in the atmosphere. It can be used to simulate interactions and decays of nuclei, hadrons, muons, electrons and photons in the atmosphere up to energies of the order of 10$^{20}$eV. For the simulation of hadronic interactions, presently CORSIKA has the option of seven high energy hadronic interaction models and three low energy hadronic interaction models \cite{Heck}. It uses the EGS4 code \cite{Nelson} for the simulation of EM component of the air shower. The rest of the paper is organized as follows. In the next section, we discuss about the simulation process involved in this work. The analysis of the simulation work and consequent results are discussed in the Section III. We summarized our work in the Section IV. | Taking into consideration of the importance of effective techniques for the gamma hadron separation in ACT, we have made an effort here to study the Cherenkov photon's density, arrival time and angular distributions in EASs of vertically incident as well as inclined $\gamma$-ray, proton and iron primaries with different energies using the simulation package, the CORSIKA 6.990 \cite{Heck}. This is the sequel of our earlier work \cite{Hazarika} to generalize the study in extending energy and zenith angles of primary particles. Cherenkov photon's density increases with energy for all primaries and decreases almost exponentially with increase in distance from the shower core. This is an obvious experimental fact that at a particular observation level it is easier to detect a high energy shower than a low energy shower and for a proper estimation of energy of a shower, the shower has to be well contained within the detector array. Also with increase in angle of inclination, the density decreases gradually near the shower core, but remains almost constant far away from the core. This result also supports the well known observational situation that, it is harder to detect an inclined than a vertical shower of same energy. $\gamma$-rays have the highest Cherenkov photon yield followed by proton and iron for any combination of energy, angle and hadronic interaction model. Thus, the equivalent energy of the iron primary must be highest followed by proton and $\gamma$-ray for a given Cherenkov photon yield for any cited combination. The average arrival time of Cherenkov photon is found to increase according to an exponential function (see equation (3)) with increase in distance from the shower core for all combinations of energy and zenith angle. With increase in energy, the general trend shows an overall increase in the arrival time for all primaries. However, with the increasing zenith angle the arrival time profile becomes flatter and hence there is a decrease in arrival time. At a particular energy and an angle of incidence, the average arrival time is highest for the $\gamma$-ray primary and least for the iron primary. All these information along with the features of density distribution may be useful to disentangle the showers of $\gamma$-ray from the hadronic showers while analyzing the experimental EAS data, apart from usual determination of direction of a shower from the arrival time information. In general, the shower to shower fluctuation for density and arrival time of Cherenkov photons decreases with increasing energy of primary particle, and is highest for proton primary and least for the $\gamma$-ray primary at all zenith angle. While estimating the systematic uncertainties in the data of a $\gamma$-ray experiment, this information will provide an important input to be considered. As we have seen in our earlier work \cite{Hazarika}, the density of photons increases with increasing altitude of observation level for all primaries, but decreases with the increasing zenith angle of all these particles. At highest zenith angle (40$^\circ$) the characteristic hump is also seen for the iron primary at lower altitude of the observation level. Similarly, the Cherenkov light front is flatter for the lower observation altitude as well as for the larger zenith angle of all primary particles. Thus the recording time of an inclined shower for a $\gamma$-ray telescope array at a lower observation level is less than that for a vertical shower of same energy and observed at higher observation level. In the calculation of density and arrival time distributions of Cherenkov photons, on average four different atmospheric models (viz., U.S. standard atmosphere as parameterized by Linsley, AT 115 Central European atmosphere for Jan. 15, 1993, Malarg\"ue winter atmosphere I after Keilhauer and U.S. standard atmosphere as parameterized by Keilhauer) available in the CORSIKA give almost similar results. Thus this analysis shows that any one of them may be used within a reasonable limit of error. The angular distributions of Cherenkov photons have distinct features in connection with the type of primary particle, its energy and zenith angle. For vertical showers of all primaries, the Cherenkov photon's angular position with respect to shower axis at which maximum photons are concentrated shifts to higher value with increasing energy of the primary. This tendency is highest for the $\gamma$-ray primary and least for the iron primary. With increasing zenith angle, the maximum photons are found to be remained within very near to the shower axes for all energy primaries. Also distributions for all primaries and energies gradually become narrower with increasing zenith angle. At low energy the iron primary has largest angular distribution at all zenith angles, whereas at higher energy it is the proton which has the largest distribution. Moreover, the QGSJETII-FLUKA and EPOS-FLUKA model combinations produce almost similar results in the density, arrival time and angular distributions. So, any of these two high energy models can be used to analyze the experimental data of $\gamma$-ray astronomy. A clear understanding of Cherenkov photon's density, arrival time and angular distributions for different primary particle with different energy and at different zenith angle, and also their possible interdependence will be greatly helpful to develop a more efficient technique of gamma-hadron separation in future. In this context a full paramerizations study on these sensitive parameters is very essential. We hope to report such work in future as a part of complete simulation study on atmospheric Cherenkov photons. | 16 | 9 | 1609.08996 |
1609 | 1609.04389_arXiv.txt | We present empirical measurements of the radii of \nhosts\ stars that host transiting planets. These radii are determined using only direct observables---the bolometric flux at Earth, the effective temperature, and the parallax provided by the {\it Gaia\/} first data release---and thus are virtually model independent, extinction being the only free parameter. We also determine each star's mass using our newly determined radius and the stellar density, itself a virtually model independent quantity from previously published transit analyses. These stellar radii and masses are in turn used to redetermine the transiting planet radii and masses, again using only direct observables. The \kgsins{median} uncertainties on the stellar radii and masses are \rstarprecper\ and \mstarprecper, respectively, and the resulting uncertainties on the planet radii and masses are \rpprecper\ and \mpprecper, respectively. These accuracies are generally larger than previously published model-dependent precisions of \rplitprecper\ and \mplitprecper\ on the planet radii and masses, respectively, but the newly determined values are purely empirical. We additionally report radii for \nrvhosts\ stars hosting radial-velocity (non-transiting) planets, with \kgsins{median} achieved accuracy of \rvrstarprecper. \kgsins{Using our empirical stellar masses we verify that the majority of putative ``retired A stars" in the sample are indeed more massive than $\sim$1.2~\msun.} Most importantly, the bolometric fluxes and angular radii reported here \kgsins{for a total of 498 planet host stars}---with \kgsins{median} accuracies of \fbolprecperall\ and \thetaprecperall, respectively---serve as a fundamental dataset to permit the re-determination of transiting planet radii and masses with the {\it Gaia\/} second data release to \rpdrtwoprecper\ and \mpdrtwoprecper\ accuracy, better than currently published precisions, and determined in an entirely empirical fashion. | } Precise and accurate estimates of the radii and masses of extrasolar planets are essential for a wide variety of reasons. In the most basic sense, these parameters allow one to estimate the bulk density of an exoplanet, and thus broadly categorize its nature, in other words, determine if it is, e.g., a gas giant, ice giant, mini-Neptune, or rocky planet. Indeed, it was the discovery of the first transiting planet HD~209458~b \citep{Charbonneau:2000,Henry:2000} that ultimately cemented the interpretation of the Doppler signals of ``Hot Jupiters" \citep[first discovered with the detection of 51~Peg~b;][]{Mayor:1995} as due to roughly Jupiter-mass objects with roughly Jupiter-like densities, and thus that these objects must be primarily composed of hydrogen and helium. Given that the ``Hot Jupiters" have periods of only a few days, this discovery, along with the fairly robust theoretical conclusion that the majority of gas giants must form beyond the ``snow line" at several AU \citep{pollack1996,kk2008}, also cemented the paradigm-shifting idea that a significant subset of giant planets undergo large-scale migration, thus revolutionizing our ideas about the evolution of planetary systems. Similarly, estimates of the masses and radii of planets, when coupled with information about their demographics (e.g., their periods and host star properties), can provide important insight into both the physics of planetary atmospheres and interiors, and the physics of planet formation and evolution. For example, a significant fraction of planets in the range of $\sim$0.1--2~\mjup\ have much larger radii than are predicted from standard models of the evolution of hot Jupiters, given their probable irradiation history (e.g., \citealt{Burrows:2000}). Despite many suggested solutions to this `inflated Hot Jupiter' problem \citep{Burrows:2007,Guillot:2002,Jackson:2009,Batygin:2011,Chabrier:2007,Arras:2010}, no one explanation has emerged as the leading contender, although empirical trends with stellar insolation \citep{Demory:2011} and perhaps age \citep{Hartman:2016} may provide clues to the correct physical model. Regardless, whichever physical mechanism turns out to be dominant, measurements of their radii as a function of the other properties of the system will provide important constraints on the physics of, e.g., tides, magnetic fields, and/or winds in these planets. As another example, estimates of the density of `warm Jupiters', i.e., those which do not appear to be affected by the inflation effect discussed above, can provide constraints on their heavy element content, and therefore potentially on the existence of a solid core \citep{Sato:2005}. Such cores are a `smoking gun' of the core-accretion, bottom-up formation scenario for giant planets \citep{pollack1996}, but are generally not expected in the gravitational instability scenario \citep{Boss:1997}. Indeed, evidence for a correlation between the inferred core mass of warm Jupiters and the heavy element composition of the host star lends credence to the idea that most, if not all, close-in giant planets form via core accretion \citep{Miller:2011}. More recently, estimates of the masses and radii of less massive planets ($M_p\la 10 M_\oplus$) detected via Kepler \citep{Borucki:2010} have uncovered an apparent dichotomy in the properties of planets with radii $\la 1.5~R_\oplus$ compared to those larger than this \citep{rogers2015}. In particular, the larger planets appear to have significant hydrogen and helium envelopes, whereas the smaller planets appear to be much more similar to the terrestrial planets in our solar system, with little to no atmospheres. Indeed, the most precise estimates of the masses and radii of the smallest planets reveal densities that are consistent with a Mg-Si-O composition that is identical to that of the Earth \citep{Dressing:2015}. Thus, accurate and precise estimates of the masses and radii of exoplanets have played, and will continue to play, an essential role in understanding the physical processes at work in these planets, and their formation and evolutionary histories. \subsection{\kgsins{The challenge of direct and accurate measurements of host-star radii and masses}} Essentially all exoplanets with measured masses and radii are those found in transiting systems. Unfortunately, as is well known, the masses and radii of transiting planets are generally not measured directly; \kgsins{the planets' masses and radii depend, through direct transit observables, on the assumed masses and radii of their host stars}. The observables are the depth of the transit, which (in the absence of limb darkening) is simply $\depth = k^2$, where $k\equiv R_p/R_\star$ and \rplanet\ and \rstar\ are the radius of the planet and star, respectively, and the velocity semi-amplitude \rvamp, which is given by \begin{equation} \rvamp \equiv \left(\frac{2\pi G}{P}\right)^{1/3}\frac{\mplanet~\sini}{(\mstar +\mplanet)^{2/3}}(1-e^2)^{-1/2}, \end{equation} where $P$, $e$, and $i$ are period, eccentricity, and inclination of the planet's orbit, respectively, \mplanet\ is the mass of the planet and \mstar\ is the mass of the star. The eccentricity of the orbit can be determined from the precise shape of the Doppler reflex (radial velocity, or RV) motion of the star, and the inclination can be measured from the relative duration of the ingress/egress $\tau$ and full-width half-maximum $T$ of the transit \citep{Carter:2008}. Thus, in order to estimate \rplanet\ and \mplanet, one must be able to measure \rstar\ and \mstar. Unfortunately, it is not possible to estimate the mass or radius of the host purely from photometric follow-up of the primary transit and RV measurements of the host star, {\it regardless of how precise these measurements are}. This is due to a well-known degeneracy, first pointed out in the case of transiting planets by \citet{Seager:2003}. As they note, the only parameter about the star that can be directly measured from observables is the ratio of the semimajor axis of the orbit $a$ to the radius of the star \ar\ \citep{Winn:2010}, \begin{equation} \frac{a}{R_\star} = \frac{\depth^{1/4}}{\pi} \frac{P}{\sqrt{T\tau}} \left(\frac{\sqrt{1-e^2}}{1+e\sin\omega}\right), \end{equation} where $\omega$ is the argument of periastron, which is also an observable from the RV curve. However, this quantity is closely related to the density of the star \citep{Seager:2003,Winn:2010}, \begin{equation} \rho_\star = \frac{3\pi}{GP^2}\left(\frac{a}{R_\star}\right)^3 -k^3\rho_p \simeq \frac{3\pi}{GP^2}\left(\frac{a}{R_\star}\right)^3, \end{equation} where $\rho_p$ is the density of the planet (and is typically $\sim \rho_*$) and the last equality follows from the fact that typically $k^3 \ll 1$. Thus $\rho_*$ can essentially be inferred from direct observables. Nevertheless, there remains a one-parameter degeneracy that makes it impossible to estimate the mass and radius of the star independently. All transiting planet systems (with only photometry of the primary transit and RV observations) are subject to this degeneracy. To break the degeneracy, one must bring in additional external constraints, such as a measurement of the surface gravity of the star, \logg\ (which is a direct observable from high-resolution spectra), astroseismological inferences of the stellar mass and radius (e.g., \citealt{Huber:2013}), or a measurement of the radius of the star (which, as we will show is a direct observable from the bolometric flux and effective temperature of the star, and the distance to the system). Up until now, these observables \kgsins{of the host stars}, while preferred because they are direct, have been either poorly measured, subject to systematic errors, or not constrained at all. \subsection{\kgsins{The value of reducing reliance on, and testing, stellar models and empirical relations}} Instead, most authors typically use theoretical and/or empirically-calibrated relations between observable properties of the star. Stellar evolution is reasonably well understood, and it is known that a star of a given effective temperature, metallicity, and density cannot have an arbitrary mass and radius. Indeed, to first order, these three parameters essentially fix the luminosity and age of the star, and thus its radius and mass. Therefore, adopting these constraints, while not direct, typically lead to much more precise estimates of the parameters of the system. Nevertheless, they are subject to uncertainties in stellar evolution models and second-order parameters (i.e., stellar rotation), and/or inaccurate calibrations of the empirical relations. One might therefore be concerned that these estimates, while precise, are not {\it accurate}. One clear demonstration of this is the case of KELT-6b \citep{Collins:2014}, where the parameters inferred using the Yonsei-Yale isochrones \citep{Demarque:2004} to break the degeneracy disagreed significantly (by as much as $4\sigma$) from those inferred using the \citet{Torres:2010} empirical relations. Likely this was due to the low metallicity [Fe/H]$\approx$$-0.3$ of the KELT-6 host star, and the fact that neither the isochrones nor the empirical relations are well-calibrated at such low metallicities. While slightly erroneous inferences about the properties of individual systems (as in the case of \citealt{Collins:2014}) are troubling, the difficulty with estimating accurate parameters of host stars and thus their transiting planets can be, and indeed has proven to be, quite deleterious in some cases, sometimes leading to markedly incorrect or inconsistent inferences about individual systems or even entire populations of planets. An early example of this is the case of the supermassive ($\sim 12~\mjup$) planet XO-3~b \citep{JohnsKrull:2008}, in which initial estimates of \rplanet\ differed by nearly a factor of two, from $\sim 1.2~\rjup$ to $\sim 2.1~\rjup$. The latter value would have implied that the planet was highly inflated relative to standard models, which would have been particularly interesting given its relatively large inferred mass. With improved photometry and thus an improved estimate of $\rhostar$, \citet{Winn:2008} were able to demonstrate that the true planetary radius was likely at the lower end of this range. Indeed, as we show in this work, our revised determinations with the {\it Gaia\/} distance---which places the star at a significantly shorter distance than previously assumed---reveal the planet to be $\mplanet \approx 7$~\mjup\ and $\rplanet \approx 1.4$~\rjup. \kgsins{We defer additional case studies of problems arising from current poor constraints on models and empirical relations to the Discussion (Sec.~\ref{sec:disc}).} \kgsins{However,} the difficulties with interpreting the properties of planetary populations due to uncertainties about the properties of the host stars became quite prominent and acute with the discoveries of thousands of planets via {\it Kepler} \citep{Borucki:2010}. Here the difficulties were threefold. First, the {\it Kepler} transiting planet hosts tended to be fairly faint compared to those found via ground based transit surveys, making characterization of the host stars more difficult. Second, the shear number of hosts made systematic assays of their properties via high-resolution spectroscopy extremely resource-intensive; this was obviously exacerbated by the faintness of the hosts. Finally, the wide {\it Kepler} bandpass, poor cadence, and/or low signal-to-noise ratio of the majority of the transit signals made estimates of the ingress/egress time, and thus stellar densities, generally imprecise. This has led to herculean efforts to characterize the properties of the host stars, often resulting in quite different conclusions as to the radius distribution of the {\it Kepler} target stars and thus their transiting planetary companions (see, e.g., \citealt{Pins:2012,Mann:2012,Dressing:2013,Huber:2014,Gaidos:2016}). \subsection{\kgsins{Aim of this paper: A path to precision exoplanetology in the era of {\it Gaia\/}}} Three recent advances now permit the determination of accurate {\it and} empirical radii and masses for a large sample of transiting planets. First, there now exist all-sky, broadband photometric measurements for stars spanning a very broad range of wavelengths, from the {\it GALEX\/} far-UV at $\sim$0.15~$\mu$m to the {\it WISE} mid-IR at $\sim$22~$\mu$m. These measurements permit construction of spectral energy distributions (SEDs) that effectively sample the majority of the flux for all but the hottest stars. Consequently the bolometric fluxes (\fbol), and in turn the stars' angular radii ($\Theta$, via the stellar effective temperature, \teff), can in principle be determined in a largely empirical manner. Using a set of eclipsing binary stars as benchmarks, \citet{Stassun:2016} have shown that with such data \fbol\ can be measured with a precision that is typically $\lesssim$3\%. Second, the {\it Gaia\/} mission's first data release \citep{Gaia:2016} has delivered trigonometric parallaxes for $\sim 2\times 10^6$ stars in common with {\it Tycho-2\/} \citep{Hog:2000}, with a precision for the best $\sim$10\% of stars of \tycparprec. These parallaxes permit $\Theta$ to be converted to \rstar. Third, a large sample of transiting planets orbiting stars that are sufficiently bright to have been included in {\it Tycho-2\/}, and consequently in the {\it Gaia\/} first data release, have been published with quantities that follow from direct observables such as the stellar density, \rhostar, the ratio of planet-to-star radii, and the orbital radial-velocity semi-amplitude. With \rstar, these quantities yield $R_p$ as well as the stellar mass, which in turn yields the planet mass. In this paper, we perform this procedure \kgsins{to measure \fbol\ and $\Theta$ for 498 planet host stars, which will serve as fundamental stellar parameters for use with upcoming data releases from {\it Gaia\/}. We also report empirical stellar and planet radii and masses as described above} for \nhosts\ stars that host transiting planets, have the necessary direct observables published in the literature, and have parallaxes newly reported in the {\it Gaia\/} first data release \citep{Gaia:2016}. We additionally perform this procedure for \nrvhosts\ stars that host non-transiting (radial-velocity only) planets, for which the newly derived planet properties remain modulo factors of $\sin i$. Importantly, the stellar and planet properties that we determine are \kgsins{independent of stellar models and of empirically calibrated stellar relations; thus the properties that we determine for the stars and their planets do not require the assumption that individual systems behave according to theoretical expectation or within the limits of mean relations. We argue, moreover, that the stellar and planet properties that we determine are} {\it empirical and accurate}, even if the {\it Gaia\/} parallaxes do not yet yield {\it precisions} that rival those typically achieved via model-dependent analyses reported in the literature. However, the $\Theta$ measurements that we determine are sufficiently accurate and precise that upcoming, improved parallax measurements from {\it Gaia\/} should enable the stellar and planet properties to be re-determined with accuracies and precisions superior to those currently available in many cases---in an entirely empirical, model-free fashion. \kgsins{As we enter the {\it TESS\/} era of superb precision in the transit parameters of very bright stars, such empirical and accurate stellar properties will become even more important than they have been for {\it Kepler\/} targets for which precision followup often proved challenging.} In Section~\ref{sec:data} we describe our study sample, the data used, and our methodological approach. The primary results of this study are presented in Section~\ref{sec:results}, including \fbol, $\Theta$, \rstar, and \mstar, followed by the planet properties, $R_p$ and $M_p$. In Section~\ref{sec:disc} we \kgsins{present additional motivation for the importance of reducing reliance on stellar models, explore the degree to which the approach laid out in this paper is truly empirical,} discuss our results in the context of previously published results, and briefly discuss the prospects for improving on the stellar and planet properties reported here with the anticipated advent of improved parallaxes from the {\it Gaia\/} second data release. Finally, Section~\ref{sec:summary} provides a summary of our results and conclusions. | } We have demonstrated that several new observational advances have enabled direct measurements of the fundamental properties for a much larger sample of bright ($V\la 12$) stars than has heretofore been available. These advances include the availability of broadband photometric measurements spanning a very broad range of wavelengths, thanks to several all-sky panchromatic surveys (e.g., GALEX, APASS, 2MASS, WISE). These photometric measurements permit construction of empirical SEDs that encompass a very large fraction of the stellar SED, which in turn enable nearly direct measurements of the bolometric fluxes of all but the hottest of these bright stars, to a precision of typically $\lesssim$3\% \citep{Stassun:2016}. When combined with estimates of the stellar effective temperatures, ideally measured from high-resolution spectra, the angular diameter of the star (as well as the extinction) can in principle be determined in a largely empirical manner. Finally, the newly available {\it Gaia\/} parallaxes can then be used to estimate the radii of the stars, essentially directly. With the current data available, we find that stellar radii can be determined to a precision of $\sim$\rstarprecper. This precision is limited by the current {\it Gaia\/} parallaxes themselves. However, with the final {\it Gaia\/} data release, we can expect typical (systematics-limited) parallax uncertainties of $\sim$\drtwoparprec\ for stars with $V\la 12$ \citep{Gaia:2016}. In this regime, the uncertainty in \rstar\ will be dominated by the uncertainties in \fbol\ and \teff. Fortunately, there are excellent prospects for improving the uncertainties (and accuracy) in \fbol\ even beyond the few-percent precision already provided in this paper. {\it Gaia\/} DR3 will release low-resolution ($R\sim 10-20$) spectrophotometry covering wavelengths of 0.33--1.05~\micron, and the proposed Explorer mission Spectro-Photometer for the History of the Universe, Epoch of Reionization, and Ices Explorer \citep[SPHEREx;][]{Dore:2016}, should it be selected, will provide low-resolution ($R\sim 40-100$) spectrophotometry from 0.75--5.0~\micron. Together these would {\it directly} measure $\sim 98\%$ of the bolometric flux of the majority of the $V\la 12$ stars with $\teff\la 8500$K (as well as the extinction as a function of wavelength to these stars from $\sim 0.3-5$~\micron). Ultimately, the precision with which \rstar\ can be measured for these stars will likely be limited by the uncertainty in \teff, which can optimistically be reduced to $\sim$1\% with a combination of high-resolution spectroscopy and SED fitting. The precision with which the stellar mass, and planetary mass and radius, can then be inferred will then ultimately be limited by the precision of the follow-up photometry of the primary transit, and the radial velocity measurements of the stellar reflex motion due to the planet, which together allow one to estimate \rhostar\ and the velocity semiamplitude, \rvamp. These parameters, together with the measurement of \rstar, allow a complete solution of the system, and thus a measurement of \mstar, \mplanet, and \rplanet. In principle, with sufficient perseverance and strict control systematic errors, these can be reduced to arbitrarily low precisions. Ultimately, we can expect precisions on \rplanet\ and \mplanet\ of several percent in the best cases \citep{Stevens:2016}. We are therefore poised to enter the era of precision exoplanetology, whereby we will be able to accurately measure the host star and planetary masses and radii of bright transiting systems to a precision of, at best, a few percent, likely limited by the total amount of available follow-up resources. Most importantly, these measurements will be direct, and will not rely on stellar models (e.g., \citealt{Demarque:2004}) or (externally-calibrated) empirical relations (e.g., \citealt{Torres:2010}). Given that the Transiting Exoplanet Survey Satellite (TESS, \citealt{Ricker:2014}) is expected to find $\sim 1700$ transiting planets, including $\sim 600$ with $\rplanet \la 2\rearth$, this will clearly be transformative for our understanding of the physical properties of exoplanets, as well as their formation and evolution. For example, it will be possible to estimate the masses and radii of essentially all known transiting Hot Jupiters in a consistent, uniform, precise, and accurate manner, thereby potentially allowing for the identification of trends within this population that have been heretofore hidden by large uncertainties, systematics, and/or inhomogeneous analysis methodology. Perhaps even more exciting is the prospect of constraining the properties of low-mass terrestrial planets, given the availability of such precise estimates of their masses, radii, and densities. This may allow for the identification of terrestrial planets that have bulk heavy element compositions that differ significantly from that of the Earth (e.g., \citealt{Rogers:2010,Unterborn:2016}), and to look for trends of these parameters with the atmospheric composition of the host stars themselves. Importantly, these precise estimates of the planetary properties will be anchored to the equally precise estimates of their host star properties. For stars more massive than the sun, and older than a few Gyr, it will be possible to estimate a robust (albeit model-dependent) age of the star via its (directly-measured) luminosity. Therefore, it will be possible to better recognize and quantify trends of planet properties with host star mass, radius, luminosity, effective temperature, and age. Notably, these host stars will themselves provide stringent tests of stellar evolutionary models, increasing the sample of stars with precise (few percent) and directly-measured masses and radii by nearly an order of magnitude above the largest such samples current available \citep{Torres:2010}. Furthermore, these stars will sample a much broader range of effective temperatures (particularly at the low \teff\ range), and many will be effectively single, and therefore will not inherit any potential biases in their parameters arising from the formation and evolution of close binaries. \kgsins{Already, we have been able to use the empirical stellar masses newly determined here to verify that the majority of putative ``retired A stars" in the sample are indeed more massive than $\sim$1.2~\msun.} \kgsins{Most importantly,} the fundamental stellar bolometric fluxes and angular radii supplied in this work will help to enable these future improvements for the 498 planet-hosting stars studied here. And the much deeper reach of future {\it Gaia\/} data releases should enable application of the methods laid out here to most if not all of the known planet-hosting stars in the Galaxy. | 16 | 9 | 1609.04389 |
1609 | 1609.06351_arXiv.txt | We consider Big Bang nucleosynthesis and the cosmic microwave background when both the neutrino temperature and neutrino number are allowed to vary from their standard values. The neutrino temperature is assumed to differ from its standard model value by a fixed factor from Big Bang nucleosynthesis up to the present. In this scenario, the effective number of relativistic degrees of freedom, $N_{\rm eff}^{\rm CMB}$, derived from observations of the cosmic microwave background is not equal to the true number of neutrinos, $N_\nu$. We determine the element abundances predicted by Big Bang nucleosynthesis as a function of the neutrino number and temperature, converting the latter to the equivalent value of $N_{\rm eff}^{\rm CMB}$. We find that a value of $N_{\rm eff}^{\rm CMB} \approx 3$ can be made consistent with $N_\nu = 4$ with a decrease in the neutrino temperature of $\sim 5\%$, while $N_\nu = 5$ is excluded for any value of $N_{\rm eff}^{\rm CMB}$. No observationally-allowed values for $N_{\rm eff}^{\rm CMB}$ and $N_\nu$ can solve the lithium problem. | Cosmological observations currently provide some of the most informative probes of physics beyond the Standard Model of particle physics. An analysis of the temperature anisotropies from the Cosmic Microwave Background (CMB) and the predictions made by Big Bang Nucleosynthesis (BBN) are two of the most robust methods available to gain an understanding of the physics governing the early Universe. One of the first such cosmological constraints drawn from such an analysis was the derivation of an upper limit on the number of neutrinos inferred from the abundance of $^4$He using BBN predictions to be $N_\nu \leq 5$ \cite{Steigman:1977kc}. More recently, precise measurements of the CMB by the WMAP \cite{Hinshaw:2012aka} and PLANCK \cite{Ade:2015xua} collaborations placed stringent limits on the effective number of relativistic species $N_{\rm eff}^{\rm CMB}$. The parameter $N_{\rm eff}^{\rm CMB}$ is an estimate of the total energy density contained in relativistic particles at recombination, parametrized in terms of the number of effective two-component neutrinos $N_\nu$. The analysis done by the PLANCK collaboration reveals \cite{Ade:2015xua} \begin{equation} \label{CMBlimit} N_{\rm eff}^{\rm CMB} = 3.15 \pm 0.23, \end{equation} where this value and its best-fit estimate are inferred in a Bayesian treatment by combining the Planck measurements with other cosmological data sources. While $N_{\rm eff}^{\rm CMB}$ includes any particles that are relativistic at recombination, we focus on the case that such ``dark radiation" consists entirely of neutrinos and leave more exotic possibilities to future work. Though not necessarily the case, it is often assumed that CMB measurements of $N_{\rm eff}^{\rm CMB}$ provide a direct probe of the number of neutrinos $N_\nu$. CMB observations are generally only sensitive to the total neutrino energy density at recombination through its effect on the expansion rate. The neutrino energy density at recombination is therefore the direct CMB observable, which does not depend only on $N_\nu$, but also on the neutrino temperature $T_\nu$. The equivalence between $N_\nu$ and the value of $N_{\rm eff}^{\rm CMB}$ derived from CMB observations assumes a standard neutrino thermal history, which we refer to henceforth as the ``standard model" (SM) temperature of neutrinos $T_{\nu \rm SM}$. The equivalence between $N_\nu$ and $N_{\rm eff}^{\rm CMB}$ is broken if nonstandard processes take place after neutrino decoupling, resulting in a temperature deviation from the usual $T_{\nu \rm SM}$. The possibility that the neutrinos might have a nonstandard temperature has a rich history and has been explored in numerous papers \cite{Kolb:1986nf, Serpico:2004nm, Ho:2012ug, Ho:2012br, Boehm:2013jpa, Steigman:2013yua, Nollett:2013pwa, Nollett:2014lwa}. As an example of such a scenario, particles that decay after the neutrinos decouple at a temperature of a few MeV may raise the photon temperature relative to the neutrinos, giving an effective neutrino temperature that is lower than that of the standard cosmological scenario. A similar effect can also occur for MeV-mass dark matter, which can stay in thermal equilibrium long enough to heat the photons relative to the neutrinos. In this more general case, the statement that $N_{\rm eff}^{\rm CMB} = N_\nu$ no longer holds true and is now generalized to \begin{equation}\label{relation:cmb-bbn} N_{\rm eff}^{\rm CMB} T_{\nu \rm SM}^4 = N_\nu T_\nu^4, \end{equation} where the total neutrino energy density (as long as the neutrinos are relativistic) is proportional to $N_\nu T_\nu^4$. The usual case of $N_{\rm eff}^{\rm CMB} = N_\nu$ is recovered if a standard-model neutrino temperature $T_{\nu SM}$ is assumed, with the more general treatment given by Eq. (\ref{relation:cmb-bbn}). This degeneracy between $N_{\rm eff}^{\rm CMB}$ and $N_\nu$ can however be disentangled by combining CMB and BBN observables. For example, \citet{Nollett:2013pwa} recently used the combination of BBN abundance predictions and CMB measurements to constrain electromagnetically coupled dark matter particles that raise the photon temperature relative to that of the neutrinos. They also showed in \cite{Nollett:2014lwa} that the opposite effect can be accomplished if a coupling is introduced between the dark matter sector(s) and neutrinos. In this paper, we consider the most general case, in which $N_\nu$ and $T_\nu$ are treated as free parameters, and then determine the observational constraints obtained from a combination of BBN and the CMB. Before the neutrinos decouple at the temperature $T_d \approx 2-3$ MeV, the weak rates ensure that $T_\nu = T_\gamma$. We make the assumption that the change in $T_\nu$ is induced after $T_d$, but before BBN begins at $T_\gamma \sim 1$ MeV, and that the neutrino temperature subsequently evolves in the standard way as the inverse of the scale factor. This is admittedly a narrow window over which the change is assumed to occur, and it is also true that the onset of BBN is not a sudden process. We discuss these issues in more detail below. The paper is organized as follows: in section \ref{sec:cosmo_params} we discuss how $T_\nu$ and $N_\nu$ affect BBN and CMB observables, while in section \ref{sec:numerical} we explore these effects more precisely using numerical simulations of BBN and its effects on primordial abundance predictions of $^4$He and deuterium. Combining these results with observational limits on the primordial element abundances, we derive corresponding limits on $T_\nu$ and $N_\nu$ and then examine the effects on $N_{\rm eff}^{\rm CMB}$. We find that even current observational bounds on $^4$He and deuterium, when combined with CMB limits, can be consistent with one extra sterile neutrino for a change in $T_\nu$ from the standard model value of only $\sim -5\%$, while two additional sterile neutrinos are ruled out. We discuss certain implications of this analysis and conclude in section \ref{sec:conclusions}. | \label{sec:conclusions} While observations of the CMB yield very precise limits on cosmological parameters, our results show that Big Bang nucleosynthesis remains an indispensible tool. For models in which the neutrino number and temperature can both vary, the CMB alone cannot produce any limits on $N_\nu$, while a combination of the CMB and BBN yields a very useful bound. In the models examined here, a value of the neutrino number as determined from the CMB of $N_{\rm eff}^{\rm CMB} \approx 3$ can be consistent with a true neutrino number, $N_\nu$, as large as 4, thus allowing for an additional sterile neutrino. Such a model requires a reduction in the neutrino temperature of approximately 5\% relative to the standard model neutrino temperature. However, a value of $N_\nu =5$ is ruled out for any value of $N_{\rm eff}^{\rm CMB}$. The obvious direction for future investigation would involve more complex behavior for the evolution of the neutrino temperature, both during and following BBN. Some of these types of behavior have been discussed previously in Refs. \cite{Kolb:1986nf, Serpico:2004nm, Ho:2012ug, Ho:2012br, Boehm:2013jpa, Steigman:2013yua, Nollett:2013pwa, Nollett:2014lwa}, but these studies do not by any means exhaust all of the interesting possibilities. | 16 | 9 | 1609.06351 |
1609 | 1609.04024_arXiv.txt | We present the results of a broadband radio-to-GeV observing campaign organized to get a better understanding of the radiation processes responsible for the $\gamma$-ray flares observed in 3C 279. The total intensity and polarization observations of the source were carried out between December 28, 2013 and January 03, 2014 using the {\it Fermi}-LAT, {\it Swift}-XRT, {\it Swift}-UVOT, and KVN telescopes. A prominent flare observed in the optical/near-UV passbands was found to be correlated with a concurrent $\gamma$-ray flare at a confidence level $>$95$\%$, which suggests a co-spatial origin of the two. Moreover, the flaring activity in the two regimes was accompanied by no significant spectral variations. A peak in the X-ray light curve coincides with the peaks of the fractional polarization curves at 43 and 86~GHz radio bands. No prominent variation was noticed for the total intensity and the electric vector position angle (EVPA) observations at radio bands during this period. We noticed a possible hint of steepening of the radio spectrum with an increase in percentage polarization, which suggests that the radio polarization variations could be simply due to a spectral change. In a simple scenario, the correlated optical/$\gamma$-ray flares could be caused by the same population of emitting particles. The coincidence of the increase in radio polarization with the X-ray flux supports the picture that X-rays are produced via inverse-Compton scattering of radio photons. The observed fractional variability for the $\gamma$-ray flare $\sim$0.23 does not exceed that in the optical regime, which is inconsistent with what we usually observe for 3C~279; it could be due to different dependencies of the magnetic field and the external radiation field energy density profiles along the jet. | Powered by accretion onto super-massive black holes (with masses up to $\sim$10$^{10}$~M$_{\odot}$), active galactic nuclei are extremely bright objects in the extra-galactic sky. In a small sub-group of these objects, a substantial fraction of accretion energy is converted into kinetic energy forming highly collimated and relativistic outflows of energetic plasma and magnetic fields, called jets. Blazars with their jets pointing close to our line-of-sight are strong emitters of electromagnetic radiation over a range of more than 20 decades in energy. Because of relativistic beaming, these sources can be detected out to much larger distances than unbeamed objects. BL Lacertae objects (BL Lacs) and flat spectrum radio quasars (FSRQs) are clubbed together and called blazars, in spite of the dissimilarity of their optical spectra -- FSRQs show strong broad emission lines, while BL Lacs have only weak or no emission lines in their optical spectra. Blazars are notable for showing variability on a range of timescales that is often described as a superposition of multiple flares. Despite several efforts to understand the broadband flaring activity of blazars, the exact origin of variability is still debated. Of particular interest is the question why for the same source sometimes we do see correlated behavior in different energy bands and sometimes we do not. The flat spectrum radio quasar (FSRQ) 3C~279 \citep[at z = 0.538][]{burbidge1965} is one of the most intensively studied objects of its class. The source has an estimated black hole mass in the range of (3--8)$\times$10$^8$~M$_{\odot}$ derived independently from the luminosity of broad optical emission lines \citep{woo2002} and from the width of the H$_{\beta}$ line \citep{gu2001}. Owing to its high flux density and prominent variations in total intensity and polarization, 3C~279 is an excellent candidate to examine physics of extragalactic jets and to understand particle acceleration to high energies. It has been monitored intensively at radio, optical, and more recently also X-ray and $\gamma$-ray frequencies and has been the subject of intensive multi-wavelength campaigns \citep[e.g.,][]{maraschi1994, hartman1996, chatterjee2008, wehrle1998, larionov2008, collmar2010, bottcher2007, hayashida2012, hayashida2015, kang2015}. In 2006, emission at very high energies (E$>$100~GeV) was detected from 3C~279 with the Major Atmospheric Gamma-Ray Imaging Cherenkov (MAGIC) telescope \citep{albert2008}. The source has a bright radio jet extending up to kiloparsec scales. The very long baseline interferometry (VLBI) observations have measured apparent velocities from 4 to 20~$c$ in the parsec scale region of the jet, which is aligned close to the observer's line-of-sight \citep[$\leq$2$^{\circ}$,][]{lister2013, jorstad2004}. Polarimetric observations have detected both linearly and circularly polarized emission from the parsec-scale jet of 3C~279 \citep{wardle1998, taylor2000, homan2009}. Additionally, the observed optical radiation is also highly polarized \citep[linear polarization up to 45.5$\%$ in the U~band][]{mead1990}. \citet{wagner2001} detected variable optical circular polarization in 3C~279 exceeding 1$\%$. \citet{abdo2010c} reported a coincidence of an optical polarization angle swing with a bright $\gamma$-ray flare in the source, which suggests a highly ordered configuration of magnetic field during the emission of bright $\gamma$-ray flares. The source has been detected by the LAT (Large Area Telescope) on board the {\it Fermi Gamma-ray Space Telescope} since its launch in 2008 \citep{abdo2010a}. Being one of the brightest and most rapidly variable sources in the GeV regime, 3C~279 has been the subject of several recent multi-wavelength campaigns \citep{hayashida2012, hayashida2015}. Multiple $\gamma$-ray outbursts have been detected in the source \citep{hayashida2015, paliya2015}. In December 2013, a series of $\gamma$-ray flares were observed, reaching the highest flux level measured in this object since the beginning of the {\it Fermi} mission, with F (E $>$100~MeV) of 10$^{-5}$ photons cm$^{-2}$ s$^{-1}$ and a flux-doubling time scale as short as 2~hr \citep{hayashida2015}. In June 2015, 3C~279 was observed in an exceptionally bright state \citep{cutini2015} with the highest measured flux, F (E $>$100~MeV) of 3.9$\times$10$^{-5}$ photons cm$^{-2}$ s$^{-1}$ \citep{paliya2015}, breaking its own record. Extremely bright flashes of light in the high-energy (GeV/TeV) regime at minutes to hours timescales in blazars have attracted the attention of astronomers, as this suggests that particles can be promptly accelerated with very high efficiency in tiny magnetized emission regions. The rapid variations in the high-energy regime are not always accompanied by flaring activity at other frequencies \citep{rani2013,rani2015, abdo2010c, hayashida2012}, making it even more difficult to understand the radiation processes and acceleration mechanisms involved. In the following, we present the results of our study of 3C~279 from a multi-wavelength campaign organized in December 2013, when it went through a series of rapid $\gamma$-ray flares. We monitored the source using both ground- and space-based telescopes for a time period between December 28, 2013 and January 03, 2014. The aim of our study is to understand the origin of the observed flaring activity. More specifically, we investigate the correlation of the $\gamma$-ray activity with the emission at lower frequencies. The paper is structured as follows. Section 2 provides a brief description of observations and data reduction. In Section 3, we report our results. Finally, discussion and conclusions are given in Section 4. | Extremely bright GeV/TeV flares at minutes to hours timescales are a common characteristic of $\gamma$-ray-bright blazars. The rapid flares usually do not have a low-frequency counterpart \citep{rani2013,rani2015, abdo2010c, hayashida2012}, making it difficult to interpret the radiation processes and acceleration mechanisms involved. We organized a multi-wavelength campaign to understand the physical processes responsible for the origin of extreme $\gamma$-ray flaring activity observed in 3C~279 in December 2013. The source was observed from radio (including polarization) to $\gamma$-rays. A relatively mild $\gamma$-ray flare was observed in 3C~279 during our campaign period. Compared to the extremely bright and rapid flares ($t_{var} \sim$few hours) seen at other times, this flare had relatively longer variability timescales ($t_{rise}$ and $t_{decay} \sim2.5$~days). Nevertheless, the $\gamma$-ray flare was accompanied by a prominent flare at optical and near-UV passbands, and we found a significant correlation between the two. The optical/near-UV flare, however, had a relatively fast rise ($\sim$0.4~day) and slow decay ($\sim$1.5~days). The flux variations (both at optical and $\gamma$-ray frequencies) were accompanied by no significant spectral variations. Broadband spectral modeling of the flaring activity by \citet{hayashida2015} for a period that has 3 days overlap with our observations suggests a substantial contribution of external-Compton radiation from the broad-line region. Their analysis also indicates that the location of the $\gamma$-ray emission region is comparable with the broad-line region radius. The observed optical--$\gamma$-ray correlation suggests a co-spatiality of the two emission regions. We observed a slight hint of variability (significance $\sim$1.4~$\sigma$) in the 0.3--10~keV X-ray light curve. The formal cross-correlation analysis suggests no correlation between the X-ray and $\gamma$-ray variations. There is a hint of possible correlation between the X-ray and the optical variations with the former leading the latter by $\sim$1.8~days; however, the confidence level of the correlation is below 95$\%$. A micro-flare peaking in the fractional polarization curve at 43 and 86~GHz radio bands coincides with the first peak (X1) in the X-ray light curve. Observations at a better cadence would be needed to test this apparent correlation. Unlike the higher frequencies, the total intensity variations are not pronounced at radio bands. We noticed a micro-flare in polarization degree (PD), while the variations in EVPA curves are not statistically significant. An indication of spectral steepening at radio bands with an increase in PD suggests that the PD variations could be due to a spectral change without any change in the magnetic field (B-field). The presence of contemporaneous optical/$\gamma$-ray flaring activity is not always the case for blazars. However, studies over the past few years indicate a similarity in the variability properties of the two bands \citep{cohen2014, chatterjee2012}, suggesting the same population of electrons being responsible for the synchrotron (optical) and inverse-Compton ($\gamma$-rays) radiation. For 3C~279, \citet{hayashida2012} reported a good correlation between the optical and $\gamma$-ray flares in 2008--2012. However later in 2013--2014, this correlation seemed to be less obvious \citep{hayashida2015}. A comparison of long-term optical/$\gamma$-ray flaring activity between November 2013 and August 2014 including our campaign period suggests a significant correlation between the two (Rani et al.\ 2016, in preparation). The observed optical--X-ray correlations in 3C~279 also present a complicated picture. A detailed analysis of the optical--X-ray correlations between 1996 and 2007 by \citet{chatterjee2008} suggests that for some flares optical lead X-rays while for others it is vice-versa. Additionally, an orphan X-ray flare was observed in the source in 2008--2010 \citep{abdo2010c, hayashida2012}. Presence of multiple emission regions \citep{marscher2014, hayashida2015}, different dependencies of magnetic field density and energy density of the external radiation field \citep{janiak2012}, and/or particle acceleration in a stratified jet \citep{rani2013,rani2015} have all been suggested as plausible scenarios. During our campaign period, a short-duration ($\sim$5~day) flare at optical/near-UV frequencies is found to be significantly correlated with a contemporaneous flare at $\gamma$-rays. The $F_{var}$ at $\gamma$-ray does not exceed that at optical/near-UV bands, which is usually not what had been observed for 3C~279. The $F_{var}$ at $\gamma$-rays is usually much higher than or at least comparable to that at optical frequencies \citep[][Rani et al.\ 2016, in preparation] {hayashida2012, hayashida2015}. The optical flare has a marginally significant correlation with the X-ray variations. Nearly simultaneous spectral analysis \citep{hayashida2015} suggests superposition of two spectral components, as a single-zone leptonic model failed to explain the X-ray spectrum. A coincidence of a micro-flare in the fraction of polarization at radio bands with the X-ray flare certainly emphasizes its jet base origin. In a tentative scenario, a single population of electrons seems to be responsible for the optical/near-UV and $\gamma$-ray flares. X-rays could be produced via synchrotron self-Compton with seed photons coming from short-mm radio bands as is supported by the coincidence of the increase of radio polarization (at 43 and 86~GHz bands) with the X-ray flux. The X-ray flare leading the optical flare could be expected if the high-energy electrons take longer to accelerate compared to the lower-energy electrons \citep{boettcher2007}. As a consequence, the radio and X-ray flares lead the optical/$\gamma$-ray flares. An alternative explanation could be that the particles are accelerated in a different region. The presence of multiple sub-components in the emission could also be possible. In leptonic models, the relative amplitude variations of synchrotron and IC flares could be simply determined by the B-field, the number of emitting electrons ($N_e$), and the Doppler factor ($\delta$) \citep{chatterjee2008}. A spectral hardening at radio frequencies with an increase in polarization degree does not support the B-field variations. An increase in both $N_e$ and $\delta$ causes a larger enhancement in the IC-flux compared to the synchrotron flux \citep{chatterjee2008}. As a result, one expects to see larger variability amplitude for the IC flare compared to the synchrotron flares; however, for our observations, the reverse is true. This could be because of different dependencies of B-field energy density and external radiation field energy density on the distance along the jet from the central engine \citep{janiak2012}. It is important to note that we consider leptonic models for our interpretation; however, hadronic models are equally able to reproduce a similar kind of variability. | 16 | 9 | 1609.04024 |
1609 | 1609.02602_arXiv.txt | We model the mass accretion rate $\dot{M}$ to stellar mass $M_*$ correlation that has been inferred from observations of intermediate to upper mass T Tauri stars---that is $\dot{M} \propto M_*^{1.3 \pm 0.3}$. We explain this correlation within the framework of quiescent disk evolution, in which accretion is driven largely by gravitational torques acting in the bulk of the mass and volume of the disk. Stresses within the disk arise from the action of gravitationally driven torques parameterized in our 1D model in terms of Toomre's $Q$ criterion. We do not model the hot inner sub-AU scale region of the disk that is likely stable according to this criterion, and appeal to other mechanisms to remove or redistribute angular momentum and allow accretion onto the star. Our model has the advantage of agreeing with large-scale angle-averaged values from more complex nonaxisymmetric calculations. The model disk transitions from an early phase (dominated by initial conditions inherited from the burst mode of accretion) into a later self-similar mode characterized by a steeper temporal decline in $\dot{M}$. The models effectively reproduce the spread in mass accretion rates that have been observed for protostellar objects of $0.2\,\mbox{M}_\odot \le M_* \le 3.0\,\mbox{M}_\odot$, such as those found in the $\rho$ Ophiuchus and Taurus star forming regions. We then compare realistically sampled populations of young stellar objects produced by our model to their observational counterparts. We find these populations to be statistically coincident, which we argue is evidence for the role of gravitational torques in the late time evolution of quiescent protostellar disks. | \label{sec:introduction} Protostellar disks are a ubiquitous outcome of the rotating collapse of dense molecular cloud cores in the standard paradigm of low-mass star formation \citep[e.g.,][]{terebey1984,shu1987}. Their existence has been confirmed around young stellar objects across a broad range in mass---from objects in the brown dwarf regime, to those with masses of up to $2\mbox{--}3\,\mbox{M}_\odot$ \citep[e.g.,][]{beckwith1990,andrews2005}---as well as in a wide variety of star forming environments \citep[e.g.,][]{lada1984,o'dell1994,mccaughrean1996}. Numerical simulations of collapsing cloud cores reveal that disks can form within ${\sim}10^4\,\mbox{yr}$ from the onset of core collapse \citep{yorke1993,hueso2005}. These early so-called Class 0 systems are difficult to study observationally as they are still embedded within their progenitor cloud cores \citep{andre1993}. Numerical simulations \citep[e.g.,][]{vorobyov2005b,vorobyov2006,vorobyov2010,vorobyov2015} suggest that the earliest periods (${\sim}0.5\,\mbox{Myr}$) of disk formation are rather tumultuous, as infall from the parent cloud core induces gravitational instability--driven mass accretion. Depletion of the gas reservoir by this mechanism then gives way to a much more quiescent period of accretion in which gravitational torques act to transport mass inward while transporting angular momentum outward \citep{gammie2001,lodato2004,vorobyov2007}. Indeed, the subsequent Class I and II phases are respectively marked by a decline in the rate of accretion from the surrounding natal environment, and its eventual cessation \citep{vorobyov2005a}. Hence, it is during the Class II phase, once the central star is optically visible, that the disk properties are most easily amenable to observational investigation. One result to emerge from observational studies of young stellar objects and their disks is the correlation between protostellar mass $M_*$ and the inferred accretion rate $\dot{M}$ from the disk, for which the power law exponent is typically estimated to be $\beta \sim 1.5\mbox{--}2.0$ \citep[e.g.,][]{muzerolle2005,herczeg2008,rigliaco2011}. Although this correlation appears to hold across multiple orders of magnitude in both $M_*$ and $\dot{M}$, fitting the accretion rates of brown dwarfs and T Tauri stars together may be misleading. In the brown dwarf regime, as well as for low mass T Tauri stars (i.e., those objects with mass $M_* < 0.2\,\mbox{M}_\odot$), a least squares fit yields $\beta = 2.3 \pm 0.6$. For intermediate and upper mass T Tauri stars ($M_* > 0.2\,\mbox{M}_\odot$), the equivalent fit yields a value for $\beta$ of $1.3 \pm 0.3$; suggestive that different physical mechanisms may be responsible for accretion across the sequence of protostellar masses \citep{vorobyov2008}. Studies by \citet{alexander2006} and \citet{hartmann2006} have sought to explain the $\dot{M}\mbox{--}M_*$ scaling in the context of viscous models for the disk evolution, wherein the turbulent viscosity has ad hoc spatial dependence of the form $\nu \propto r^\xi$. \citet{dullemond2006} link the disk evolution to the properties of the parent cloud core, providing a self-consistent basis for the results of their study. However, their models require that the ratio of rotational to gravitational energy be uniform across all cloud core masses. \citet{rice2009} have even attempted to (weakly) incorporate the additional effects of magnetic fields (in high temperature regions of the disk) in quasi--steady state models, but were also unable to fully account for the observed correlation. In this paper we present a study of the quasi--steady state evolution of viscous circumstellar disks surrounding young stellar objects, following the cessation of mass accretion onto the protostar-disk system (definitively Class II objects). These disks inherit initial conditions roughly consistent with the results of numerical simulations of the earlier burst phase \citep[e.g.,][]{vorobyov2005b,vorobyov2006,vorobyov2010,vorobyov2015}, and undergo diffusive evolution wherein angular momentum redistribution is driven by self-gravity, which we parameterize in terms of an effective kinematic viscosity \citep[following][]{lin1987}. We add to this a simplified argument for angular momentum conservation that correlates disk size with protostellar mass at the start of our simulations. With these assumptions, we are able to reproduce many features of the observed correlation between $\dot{M}$ and $M_*$ for young protostellar systems. Recent observations of disks using near-infrared polarization imaging \citep{liu2016} have found that disks around four recently outbursting (FU Ori) sources have large-scale (hundreds of AU) spiral arms and arcs that are consistent with models of gravitational instability. Added to previous near-infrared detections of spiral structure in smaller disks \citep[e.g.,][]{hashimoto2011, muto2012,grady2013}, there is a growing realization that meaningful spiral structure, arcs, and gaps exist in Myr-old disks (see the review by \citet{tamura2016} on the SEEDS survey by the Subaru telescope). New efforts are being made to use gravitational instability driven disk evolution models to predict the near-infrared scattered light patterns as may be seen by the Subaru or Gemini telescopes, or the millimeter dust emission patterns that may be seen with the ALMA telescope \citep{dong2016}. Furthermore, numerical simulations are also being extended to include long-term residual infall from the molecular cloud to the disk (even after the parent cloud core may have dissipated), which may be needed to keep gravitational instability active after several Myr \citep{vorobyov2015b,lesur2015}. Our model in this paper studies gravitational torque driven evolution in a simplified manner. It does not however include residual mass infall from the cloud, which may be a subject of future work. In this paper we seek to characterize the bulk of transport within the disk through the action of gravitational torques, in the same spirit as the models of e.g., \citet{armitage2001} and \citet{zhu2009,zhu2010}. Our aim is to explain the global behavior of disks in which the mass accretion rate is predominantly set by the action of gravitational torques acting through most of the disk. Other accretion mechanisms may be necessary in the innermost sub-AU regions of the disk, possibly introducing short-term time variability. The above studies typicaly invoke the magnetorotational instability \citep{balbus1991} as the transport mechanism in the hot inner disk, however it is worthwhile to keep in mind that the region $0.1\mbox{--}1.0\,\mbox{AU}$ from the star is generally thought to be the outflow driving zone \citep[e.g.,][]{garcia2001,krasnopolsky2003} from which significant amounts of angular momentum and mass are carried away from the disk. | \label{sec:dnc} In this paper, we have shown that the observed power law correlation between mass accretion rate $\dot{M}$ and protostellar mass $M_*$ can be explained within the framework of gravitational--torque--driven transport. We parameterize the effects of the gravitational torques as an effective kinematic viscosity using Toomre's $Q$ criterion \citep{toomre1964}, noting that this prescription resembles but also differs from the classical $\alpha$ model of \citet{shakura1973}. We carry out more than 200 individual simulations of protostellar disks in order to examine the time evolution of their mass accretion rates in the $\dot{M}-M_*$ plane. The rates associated with a particular protostellar mass agree with those inferred from observational studies of T Tauri disks across a broad spectrum of protostellar masses. The observed scatter in $\dot{M}$ arises naturally as a result of the temporal evolution of the protostar-disk system through this plane. We are able to use a simple statistical argument, resampling our simulations onto the initial mass function of \citet{chabrier2005proc}, to show that even with limited sampling, our simulation results are sufficiently robust to be able to reproduce the observed correlation. The initial disk masses presented in this paper are somewhat greater than is often reported in the literature \citep[by a factor of ${\sim}10$, e.g.,][]{andrews2005}. However, current estimates for disk masses based on dust emission may have been systematically underestimated \citep{hartmann2006,dunham2014}. Nevertheless, as the efficacy of our transport mechanism is dependent on disk mass, it is possible that that additional physics may be required at late times to remove the remaining disk material within observed disk lifetimes \citep{hernandez2008}. Conversely, there is growing evidence that some disks may persist for several Myr and have noticeable spiral structure, arcs, and gaps that may be indicative of gravitational instability \citep{liu2016}. One way to understand this is by invoking gas accretion on to the disks for an extended period of several Myr \citep{vorobyov2015b,lesur2015}. Accretion on to the disk is not a part of our current model and is a subject of future work. In a recent paper, \citet{ercolano2014} offer a physically different model explanation. In their view, $\dot{M}$ is initially mass independent, and declines in a self-similar manner with a somewhat different power law index than in our model, due to the use of an $\alpha$-viscosity. Mass accretion is then quenched when a model wind mass loss rate (that depends on the X-ray luminosity of the protostar) equals the mass accretion rate. The physical view here is that the observed mass-dependent X-ray luminosity sets the $\dot{M}\mbox{--}M_*$ relation by reducing the accretion rate when it drops to the level of the wind mass loss rate. Mathematically however, both models depend partially upon a bimodal mass accretion rate history. It is only the physical explanation of the two phases, and the specific mathematical shape of their curves, that differ. The observed correlation in $\dot{M}\mbox{--}M_*$ may therefore be fit by a variety of models that have common mathematical elements, but differ substantially enough in their physics that they will hopefully lead to interesting observational comparator tests in the future. | 16 | 9 | 1609.02602 |
1609 | 1609.05214_arXiv.txt | We examine the spatial distribution of the oldest and most metal poor stellar populations of Milky Way-sized galaxies using the APOSTLE cosmological hydrodynamical simulations of the Local Group. In agreement with earlier work, we find strong radial gradients in the fraction of the oldest (t$_{\rm form} < 0.8$ Gyr) and most metal poor ([{\rm Fe}/{\rm H}]$ < -2.5$) stars, both of which increase outwards. The most metal poor stars form over an extended period of time; half of them form after $z=5.3$, and the last 10\% after $z=2.8$. The age of the metal poor stellar population also shows significant variation with environment; a high fraction of them are old in the galaxy's central regions and an even higher fraction in some individual dwarf galaxies, with substantial scatter from dwarf to dwarf. Overall, over half of the stars that belong to both the oldest and most metal poor population are found outside the solar circle. Somewhat counter-intuitively, we find that dwarf galaxies with a large fraction of metal poor stars that are very old are systems where metal-poor stars are relatively rare, but where a substantial old population is present. Our results provide guidance for interpreting the results of surveys designed to hunt for the earliest and most pristine stellar component of our Milky Way. | It is often assumed that the most metal poor stars today are the oldest objects in the Galaxy -- and as such they are the most extensively studied relics from the ancient Universe. While very useful as a first order approximation, it is nevertheless obvious that the adoption of metallicity as a clock is only valid in closed-box environments, where metals are well-mixed at all times. In reality, it is expected that chemical evolution starts and progresses at different rates in various environments. To complicate matters even more, the $\Lambda$CDM cosmological paradigm of hierarchical formation allows stars to find themselves in very different environments throughout their lives -- albeit still bearing the chemical imprint of their birth location. The very first stars in the Universe are expected to form in environments that are overdense at early times. In low-density environments on the other hand, pristine stars may have formed much later -- provided that the gas in them has remained unpolluted. Simulations suggest that pristine stars can still form at moderate redshifts well after the completion of global reionization \citep[e.g.,][]{Scannapieco03,Scannapieco06,Tornatore07,Trenti09, Muratov13a,Xu16}. A distinction is sometimes made between those stars that were formed out of unpolluted gas but have been affected by radiation or kinematical feedback from other nearby stars, and those that are truly the first stars in their environment \citep[e.g.,][]{McKee08}. From a modelling perspective, simulations of the formation of the very first stars have proved to be complex. Generally, it is thought that their initial mass function is more top heavy than that of present day star formation, due to the inability of the initial gas cloud to cool and collapse via metal line emission. However, it is unclear at what level later fragmentation of the protostellar cloud could facilitate the formation of subsolar mass stars which -- even if formed at the dawn of time -- would still be around today \citep[see][and references therein]{Karlsson13, Bromm13,Greif15}. Furthermore, a full treatment of metal-line, molecular, and dust cooling, stellar nucleosynthesis of zero-metallicity stars, their stellar feedback and the reionization epoch would be needed to model accurately the epoch of the first stars. None of these processes is well understood and even when certain assumptions are made, it is particularly challenging and computationally expensive to simulate them within a fully cosmological setting. Modelling studies have therefore predominantly either focussed on a detailed treatment of processes in the early universe in a limited volume and/or a limited time \citep[e.g.,][]{Abel02,Wise12,Muratov13a, Muratov13b,Smith15}, or resorted to a simpler treatment of physical processes - often through a so-called semi-analytical approach whereby physical processes are modelled by analytical relations on top of a cosmological merger tree or Press-Schechter formalism \citep[e.g.,][]{Scannapieco03,Yoshida04,Scannapieco06, Salvadori07,Salvadori10, Komiya10,Tumlinson10,Gao10}. Observational endeavours have been extremely helpful in constraining models of early epoch star formation in the Milky Way. Several survey efforts have been mining the Milky Way stars for the lowest-metallicity members \citep[e.g.,][]{Beers85, Christlieb03, Keller07, Schlaufman14, Casey15, Howes16} and in the coming decade it is expected that the search for metal poor stars will intensify. At present, nine stars are known to have $[{\rm Fe}/{\rm H}]<-4.5$ \citep{Christlieb02, Frebel05, Norris07, Caffau11, Hansen14, Keller14, Bonifacio15, Frebel15}. With such low numbers of stars at the lowest metallicities, any history of the Milky Way's earliest epochs that might be derived from them would remain anecdotal at best. Another limitation of observational work in this field is that age determinations of individual stars have large uncertainties, up to several gigayears for old ages. The best avenue available to link these stars to the very early Universe is by comparing their abundance patterns to the predicted yields for metal-free massive stars. For instance, the most iron-poor star, SMSS J031300.36-670839.3 at $[{\rm Fe}/{\rm H}]$ $<$ -7.0 could be linked to a $\sim$60 solar-mass, zero-metallicity progenitor \citep{Keller14}. It is particularly intriguing that, from analysis of the limited numbers of extremely metal poor stars ($[{\rm Fe}/{\rm H}]<-3$), it seems that their abundance patterns vary with environment in the Galaxy. Recently, \citet{Howes15} showed that in 23 stars in the bulge region none shows significant carbon enhancement, whereas this occurs in $\sim$32 per cent of halo stars of such low metallicities \citep{Yong13b}. Although always based on small samples, in earlier work it has already been noticed that the number of carbon-enhanced stars (or their abundance patterns) seems to change as a function of environment. Variations -- or hints thereof -- have been reported with increasing disk height \citep{Frebel06}, in the inner versus outer halo \citep{Carollo14} and in the Sculptor dwarf spheroidal galaxy compared to the halo \citep{Starkenburg13, Skuladottir15}. If confirmed by larger studies, this might indicate that the most metal poor stars in various Milky Way environments have been formed through different channels, or even at different times \citep[see also the discussion in][]{Gilmore13,Norris13,Starkenburg14,Salvadori15}. Similarly, the diversity in abundance patterns found among the metal poor stars in satellite galaxies can give us also clues on the early Universe and chemical enrichment processes \citep[one of the more striking examples are the r-process enhanced stars in Reticulum II, ][]{Ji16}. Observationally, dwarf galaxies -- and particularly those as close to us as the satellites of the Milky Way -- are of great value, because they allow us to study in detail the early formation of galaxies in a different mass scale. Some of the faintest system, belonging to the class of so-called ``ultra-faint'' galaxies, seem to qualify to be ``first galaxies'' or ``fossil galaxies'' -- meaning that the great majority of their total stellar population was formed before the epoch of reionization \citep{Bovill09,Bovill11a, Bovill11b}. However, which and how many systems exactly qualify fossil galaxies criteria is debated and it is unclear to what extent quenching through tidal disturbance by the Milky Way could influence these results \citep{Weisz14}. \citet{Brown14} derives ages from Hubble Space Telescope colour-magnitude diagrams for six ultra-faint galaxies and finds all of them are consistent with having formed 80 per cent of their stars by $z=6$. Additional evidence for only one short burst of star formation can come from chemical element analysis of their stars showing no enrichment by, for instance, supernovae Type\ Ia products or s-process elements \citep{Frebel12,Frebel14}. In this work, we study the early generations of stars in the APOSTLE simulations \citep{Sawala16,Fattahi16}. This is a suite of twelve volumes selected to resemble the Local Group simulated with sufficient mass resolution to study the Milky Way and Andromeda-type galaxies as well as their smaller companions and isolated dwarf systems. Moreover, the hydrodynamical nature of the simulations allows us to follow enrichment and feedback processes more self-consistently than in semi-analytic frameworks. The main purpose of this work is to present further predictions and guidelines for the interpretation of present and future surveys of metal-poor stars throughout the Galaxy. We will introduce our suite of simulations in more detail in Section \ref{sec:apostle}. We first review the general distribution of metallicity and age within the Milky Way and M31 analogues in Section \ref{sec:metageapostle}, before focussing our attention on the oldest and most metal poor stars in Sections \ref{sec:oldormetpoor} and \ref{sec:oldandmetpoor} where we study their distribution in the main galaxies. Additionally, we compare these findings with the old and most metal poor populations in the surviving satellites and isolated dwarf galaxies in Section \ref{sec:dwarfs}. To aid future surveys, we discuss how to find the systems in which a metal-poor star is most likely to be very old in Section \ref{sec:howto}. We finally discuss the assumptions and a comparison with literature results in Section \ref{sec:disc}. Throughout this paper we adopt the solar abundance for iron to be ${\rm log} \ \epsilon_{\rm Fe}$ = 7.50\footnote{We use here the customary astronomical scale for logarithmic abundances where log $\epsilon_{\rm X}$ = ${\rm log}_{10}$(N$_{\rm X}$/N$_{\rm H}$)+12, with N$_{\rm X}$ and N$_{\rm H}$ the number densities of element X and hydrogen, respectively.} \citep{Asplund09}. | \label{sec:disc} \subsection{Where are the first stars?} In Fig.~\ref{fig:XY4} we again show a view of our six high-resolution main galaxies. The right panels of each 2$\times$2 grid show the oldest and most metal poor populations separately. On the left we show additionally a view of the stars formed before reionization in our simulation and of the perfectly pristine stars (those that are unpolluted by any metals). As discussed in the introduction, we do not necessarily believe that these stars should still survive in the simulation (they might have a very different IMF), nor that we are appropriately set up to trace these populations. Nevertheless, we show them here to illustrate and discuss how the sites in our simulations where the first and pristine star formation has occurred are distributed in the present day, as they might also trace the sites where we can find second generation stars that show their chemical imprint. In all galaxies, the pristine population has a total mass of only 2-5 per cent of that of the most metal poor stars, a negligible fraction. However, their radial distribution shows great similarity to that of the most metal poor stars, as can also be seen from the percentage of each of the populations found at radii beyond 8 kpc. Similar to Fig.~\ref{fig:XY} we have labelled each panel in Fig.~\ref{fig:XY4} with a percentage indicating the mass fraction of these stars that are found outside the solar circle (at distances $>$ 8 kpc). The results clearly show that in our simulation the outskirts of the galaxies, beyond the solar radius, do not only contain metal poor and old stars, they also contain a significant number of pristine stars and stars formed before reionization (corresponding here to a redshift of 11.5). A similar conclusion can be reached for the smaller still-bound satellites, some of which contain pristine stars that were formed before reionization. It is interesting to note that there is significant variation from galaxy to galaxy. Additionally, we still see a trend showing an even more centrally concentrated population of pre-reionization stars compared to pristine stars. Of the latter we find that 75 per cent is found at larger radii than 8 kpc, whereas the former only has 50 per cent of its population at those radii. If we focus solely on even older stars, formed at $z >15$ (not shown), we find their population to be even more centrally concentrated. However, also here we find that typically at least a few percent of them reside at radii $>$ 8 kpc today, only in one galaxy all stars are contained within 8 kpc. Based on these results we conclude that surveys primarily interested in finding the \textit{oldest} stars are still best served by focussing on the central Galaxy, although this has of course to be balanced with the difficulties of observing those regions. Observationally, the satellites of the Milky Way are much more accessible: many are relatively close, often visible at high Galactic latitude, and reside in the low-density Galactic halo environment. We find that surveys interested in studying first star formation in various environments as well as very old stars are also well-off targeting the regions outside the solar radius and the Galactic halo with its streams and bound substructures. They are very likely to still uncover first stars, or, if these stars have not survived, their chemical imprints on the next generations. \subsection{Comparison with other work} As mentioned in Section \ref{sec:oldandmetpoor}, our finding that the stars that are both metal poor \textit{and} old are centrally concentrated agrees qualitatively very well with the conclusions reached by earlier work \citep{White00, Scannapieco06, Brook07,Tumlinson10, Salvadori10, Gao10}. Quantitatively, however, there are many differences. Because of the great advances in hydrodynamical simulation techniques and computational power, we can work here at much higher resolution than was possible before \citep[for instance when compared to][]{Brook07} and without having to resort to semi-analytical techniques used to tag stars in N-body dark matter simulations \citep[as in][]{White00,Scannapieco06,Tumlinson10,Salvadori10,Gao10,Cooper10}. \citet{Scannapieco06} and \citet{Brook07} explore the distributions of the first and second generations of stars with a semi-analytical technique grafted on an N-body simulation with a resolution similar to our L2 simulations, and subsequently with a resolution several order of magnitudes lower with a hydrodynamical technique. \citet{Scannapieco06} draw several conclusions that are robust to the particular metal dispersion prescription and other approximations in the semi-analytical code: in particular they find a population of metal free stars inside the Galactic halo environment that formed as late as redshift $z=5$ (or in some cases even $z=3$) and that extends to very large radii. \citet{Brook07} confirm these findings and show how also in their models the oldest stellar populations are more centrally clustered, but metal free stars can still be found at large radii. Comparing these and our results with the \citeauthor{Tumlinson10} model, one striking difference is that in \citet{Tumlinson10} \textit{all} stars with $[{\rm Fe}/{\rm H}]<-2.5$ are formed before $z=6$ (i.e., all their most metal poor stars are old according to the definitions used in this paper) and in their fiducial model with the epoch of reionization at $z=10.5$, all stars with $[{\rm Fe}/{\rm H}]<-3$ form before $z=10$. Their radial gradient result of more centrally concentrated older stars thus merely differentiates between $z=10$ and $z=15$ at most, corresponding to $\sim200$~Myr. Just like in our simulations, \citet{Tumlinson10} does not adopt any special IMF or formation mechanism for the first generations of stars. However, in our chemodynamical model, clearly many stars with low metallicities are formed later than $z=6$. Specifically, \citet{Tumlinson10} show that in their model, the fraction of extremely metal poor stars (selected as $[{\rm Fe}/{\rm H}] < -3.0$) that are born before $z = 15$ changes from $\approx$13 per cent in the inner regions to just a few per cent in the outer regions. When we make identical cuts in our simulations we find that a very low percentage of stars is formed at $z > 15$, even at such low metallicities. In the L2 galaxies, none are formed at all, indicating (as expected) a dependence of resolution on the very first star formation in the simulations. In the L1 galaxies, we always find the ratio of $[{\rm Fe}/{\rm H}] < -3.0$ stars that are formed at $z > 15$ to be below 3 per cent. No radial gradient is apparent for this fraction. \citet{Salvadori10} find that while the relative contribution of very metal poor stars increases with radius from the Galactic centre, the oldest stars populate the innermost region, in agreement with the conclusions reached in our work. In the fiducial model of \citet{Salvadori10} an instantaneous and homogeneous mixing prescription is used to mix metals within each halo's gas reservoir, and similarly also within the diffuse gas in the Milky Way for metals that are blown out by supernova explosions. This results in a pre-enrichment of the diffuse gas in the Milky Way up to $[{\rm Fe}/{\rm H}]=-3$ by a redshift of 7. They report that due to this prescription no stars with metallicities lower than $[{\rm Fe}/{\rm H}]=-3$ are found in 80 per cent of Milky Way satellite galaxies, a prediction that disagrees with our findings and that we now know to be contradicted by observational evidence \citep[e.g.,][]{Kirby08, Starkenburg10,Tafelmeyer10,Frebel10,Kirby10,Venn12,Starkenburg13,Frebel14,Jablonka15}. Their experiment with an inhomogeneous mixing prescription does not affect the simulation results significantly, but they stress that the true mixing at low metallicities or large radii remains very uncertain. In \citet{Gao10}, molecular hydrogen cooling in minihaloes is modelled as a recipe on top of the dark matter merger tree backbone of the Aquarius simulations \citep{Springel08}. The explicit cooling of molecular hydrogen requires very high resolution and strong assumptions on the photodissociating background radiation and is hence ignored in all other work mentioned here, including our own. They find the minihaloes to be strongly clustered in space, therefore limiting the number of final building blocks that contain them as most of them merge in time. At the present day, the environments where the first stars are found are again centrally concentrated, but they also report that a significant fraction (over 20 per cent, a very similar percentage as found in this work as discussed in Section \ref{sec:oldandmetpoor}) remains in the satellites that are presently surrounding the Milky Way. In \citet{Gao10}, none of the baryonic processes are modelled after the first generation of stars, so the present day properties of these satellites can only be compared on the basis of their dark matter properties, which, however, do seem to be comparable to the surviving satellites that we observe around the Milky Way today. \subsubsection{The role of metal mixing} \begin{figure} \includegraphics[width=\linewidth]{f_R_substr_recipes.pdf} \caption{The fraction of most metal poor stars that are old as a function of radius. Shown here is the median value for all L2 galaxies (corresponding to the blue thick line in the bottom panel of Fig.~\ref{fig:fR}; the blue line in this figure is identical to that) for two different metallicity calculations: an SPH-smoothed metallicity (our fiducial choice) and a particle metallicity (no metal mixing at all). We additionally show two different cutoff values for the most metal poor stars: $[{\rm Fe}/{\rm H}] < -2.5$ (as used throughout the paper) and $[{\rm Fe}/{\rm H}] < -3.0$.} \label{fig:recipes} \end{figure} The many differences between the models once more illustrate how the results at early epochs and low metallicities are sensitive to the assumptions that are made and the modelling prescriptions that are used. Here we investigate in more detail in particular the sensitivity of our results to the assumptions on metal mixing. As described in Section \ref{sec:metageapostle}, our fiducial model includes an SPH kernel-smoothed mixing prescription. However, in the APOSTLE simulations we also store the individual star particle metallicities, inherited from unmixed gas. In general, we find that all qualitative results are robust to either calculation of the metallicity and that most quantitative results show only very minor variations (for instance, the percentages of the different populations that are located outside an 8 kpc radius as labelled in Figures \ref{fig:XY} and \ref{fig:XY4} are robust to the metallicity prescription used). However, there are significant changes in the radial dependence of the fraction of the most metal poor stars that are old. This is illustrated in Fig.~\ref{fig:recipes} where the median fraction for all L2 galaxies is shown as a function of radius for both types of metallicity estimates (the blue line corresponds to the blue thick line in the bottom panel of Fig.~\ref{fig:fR}). Fig.~\ref{fig:recipes} additionally shows that there is a degeneracy between the amount of mixing and the cutoff $[{\rm Fe}/{\rm H}]$ used to determine whether a star is metal poor. For the particle metallicity scheme we still see a clear gradient when we adopt a lower metallicity threshold. Whereas SPH-based results will generally underpredict mixing and use particle metallicities, semi-analytic models often assume perfect mixing of all gas within a system. Our implementation of smoothed metallicities lies in between these two extremes, but while it decreases the severity of the sampling aspect of the mixing problem within SPH, it does not solve it either \citep{Wiersma09a}. Clearly, quantitative results from any prescription quoting the fraction of metal poor stars that are old in a certain Galactic region should be treated with caution. They are very dependent on the assumptions for metal mixing, a process that is uncertain. | 16 | 9 | 1609.05214 |
1609 | 1609.09383_arXiv.txt | Coalescing massive black hole binaries, formed during galaxy mergers, are expected to be a primary source of low frequency gravitational waves. Yet in isolated gas-free spherical stellar systems, the hardening of the binary stalls at parsec-scale separations owing to the inefficiency of relaxation-driven loss-cone refilling. Repopulation via collisionless orbit diffusion in triaxial systems is more efficient, but published simulation results are contradictory. While sustained hardening has been reported in simulations of galaxy mergers with $N\sim 10^6$ stars and in early simulations of rotating models, in isolated non-rotating triaxial models the hardening rate continues to fall with increasing $N$, a signature of spurious two-body relaxation. We present a novel approach for studying loss cone repopulation in galactic nuclei. Since loss cone repopulation in triaxial systems owes to orbit diffusion, it is a purely collisionless phenomenon and can be studied with an approximated force calculation technique, provided the force errors are well behaved and sufficiently small. We achieve this using an accurate fast multipole method and define a proxy for the hardening rate that depends only on stellar angular momenta. We find that the loss cone is efficiently replenished even in very mildly triaxial models (with axis ratios 1 : 0.9 : 0.8). Such triaxiality is unavoidable following galactic mergers and can drive binaries into the gravitational wave regime. We conclude that there is no `final parsec problem'. | Supermassive black hole binaries (BHBs) are expected to form efficiently through cosmic time as a result of galaxy mergers, if the original galaxies contain a central massive black hole (MBH) \citep{BBR1980}. In the early stages of the merger, the MBHs are dragged towards each other by dynamical friction, until they form a binary system. The binary continues to evolve and harden by gravitational interactions with stars that come within a few binary separations. As a result of such encounters, stars subtract energy and angular momentum from the binary, and they are ejected to larger distances. $N$-body simulations of galaxy mergers \citep[e.g.][]{QH1997, MM2001, GM2012} show a decrease in the central stellar density and the formation of a core. This process, often called {\it core scouring}, may be responsible for the observation of cores in massive elliptical galaxies \citep{Lauer1985,Lauer1995,faber1997,Graham2004,lauer2005,ferrarese2006,Graham2013}. If a sufficient supply of stars is available in the merger remnant to interact with the binary, hardening continues down to separations where emission of gravitational waves becomes important. The evolution then proceeds rapidly to inspiral and black hole coalescence. The recent detections of gravitational wave signals (GW150914 and GW151226) by Advanced Ligo \citep{GW150914, GW151226} from the merger of two stellar mass black holes proves very strongly that black holes exist, they form binaries, and merge within a Hubble time due to emission of gravitational waves. The same is expected for MBHs, and several missions and detectors are being planned in the hope to detect gravitational waves from BHBs. In the context of BHBs, The Pulsar Timing Array (PTA) \citep[e.g.][]{Babak2016}, which is already operational, is the most suitable to detect gravitational waves from the most massive BHBs, while the $e$LISA space-based interferometer \citep{Barausse2015} will be sensitive to the frequencies typical of lower mass BHBs, like the Milky Way's MBH. Detection of BHBs in the gravitational waves window would provide crucial information on the masses, spins, orientations and even distances of the two black holes. It would also provide an exciting new cosmological probe, given that BHBs can in principle be seen right back to the beginning of the Universe \citep{hogan2009}. All efforts to detect gravitational waves from BHBs are based on the assumption that the binaries do in fact coalesce efficiently, i.e. a significant fraction have a timescale for emission of gravitational radiation shorter than a Hubble time. This assumption seems to be corroborated by observations: despite strong efforts to detect BHBs with a variety of techniques, only a handful of candidates exist, and for most of these systems alternative explanations have been put forward \citep[for a review see e.g.][]{DSD2012}. It appears that BHBs find a way to coalescence, but the mechanism for this remains to be understood. While gas may play a role in bringing BHBs into the gravitational wave regime, there is evidence for little or no gas in massive elliptical galaxies, and works by \citet{cuadra2009} and \citet{lodato2009} show that gas discs are only efficient at driving mergers of BHBs for low mass MBHs. This suggests that if BHBs merge, in many cases they must do so in the absence of gas. In this paper, we consider only pure stellar systems devoid of gas. Theoretically, the ultimate fate of BHBs depends on the supply of stars to the binary's loss cone, i.e. the region in phase space characterised by low enough angular momentum to ensure interaction with the binary. If the reservoir of stars on intersecting orbits is depleted and the loss cone cannot be maintained sufficiently populated, the binary's evolution will dramatically slow down or even stall \citep{BBR1980}. There are two distinct classes of physical mechanisms for loss cone refilling in the absence of gas: collisional processes and collisionless processes. In idealised spherical galaxies, the evolution of BHBs slows down after all stars initially on loss cone orbits are ejected, which occurs on a dynamical timescale. Repopulation of the loss cone in such models is due solely to two body relaxation, i.e. a collisional process. This operates on the relaxation timescale, which for most galaxies is much longer than the Hubble time, and implies that the binary is in the empty loss cone regime at all times \citep{MM2001, MM2003}. In $N$-body simulations of isolated spherical models, this regime can be identified by the $N$-dependence of the binary hardening rate, since the relaxation time scales approximately as $N/{\rm log}N$ \citep{MF2004,BMS2005,merritt2007}. The earliest simulations, however, did not observe this scaling due to the low $N$ used which implied a very short relaxation time and therefore a full loss cone \citep[e.g.][]{MM2001}. Collisional processes are rather inefficient at repopulating the loss cone in real galaxies, and are not sufficient to lead BHBs to coalescence in a Hubble time. However, they need to be treated very carefully since they introduce a spurious population of loss cone orbits in any $N$-body simulation where $N$ is smaller than the true number of stars. In non-spherical galaxies, refilling of the loss cone can proceed due to collisionless ``diffusion'' in angular momentum. Torques from a non-spherical potential cause stellar angular momenta to change on a timescale which is much shorter than the relaxation time, though typically longer than radial orbital periods \citep{Merrittbook,pontzen2015}. In axisymmetric nuclei, stellar orbits conserve only energy and the component $J_z$ of angular momentum parallel to the symmetry axis. While the total angular momentum $J$ is not conserved, conservation of $J_z$ implies that there is a lower limit to how small $J$ can become due to diffusion. In this sense, orbits in axisymmetric potentials are not truly centrophilic, but they reach very small radii and contribute to loss cone refilling. There are two families of orbits in flattened potentials: the tube orbits, regular orbits for which $J$ is approximately conserved, and saucer orbits, for which $J$ can become very small. The fraction of stars on saucer orbits depends on the degree of flattening in the system. In triaxial nuclei, $J$ is not conserved and truly centrophilic orbits are possible. In addition to two families of tube orbits and saucer orbits, there exists a new family of orbits, called pyramids, which can achieve arbitrarily low values of $J$ \citep{MV2011}. All stars on pyramid orbits eventually reach the centre, though the timescale can vary. While the fraction of saucer orbits in axisymmetric models is expected to be small, the fraction of pyramid orbits in triaxial models can be very high \citep{PM2004}. This suggests that collisionless loss cone refilling may be efficient at driving BHBs to coalescence \citep[e.g.][]{NS1983, PM2004}. Supporting evidence for collisionless loss cone refilling appeared with the first simulations of merging galaxies hosting central MBHs, in which the merger was followed from early times \citep{preto2011,khan2011,GM2012,khan2012} as well as in cosmologically motivated mergers \citep{khan2016}. In these simulations, the hardening rate of the binary is found to be to largely independent of $N$, and the merger remnant shows significant triaxiality, at least in the central regions. The supply of stars in these cases is sufficiently high to ensure MBH coalescence in much less than a Hubble time. However, the complexity of galaxy merger simulations leaves open the possibility that a different process (that could be physical or a numerical error) is responsible for the sustained binary. If torques from a non-spherical background are responsible for efficient loss cone refilling, the same behaviour should be observed for binaries placed in isolated non-spherical models. The first $N$-body simulations of BHBs in triaxial galaxy models were performed by \citet{BMSB2006}, who considered flattened models with net rotation and $N \leq 10^6$. In this case, an equal mass BHB is found not to stall but to sustain an hardening sufficient to lead to coalescence. On the other hand, \citet{VAM2014} follow the evolution of BHBs in spherical, axisymmetric and triaxial non-rotating isolated models, using direct $N$-body simulations with particle numbers up to one million. They find an $N$ dependence of the hardening rate in all models, with only a mild flattening at the largest $N$ for the non-spherical models. Hardening rates are higher in non-spherical models than in spherical ones, but always lower than the full loss cone rate. They conclude that collisional effects still contribute significantly to the loss cone repopulation at these $N$, and prevent a reliable extrapolation to real galaxies. $N$-body simulations of collisional stellar systems are very expensive, especially when regularisation and/or very small softening are employed to model the evolution of the BHB accurately. A possible solution to limit the effects of collisional repopulation is to adopt a collisionless numerical method. \citet{VAM2015} follow the evolution of BHBs in isolated models with a Monte Carlo code able to suppress the effects of two-body relaxation. In this case, they find that hardening rates tend to become $N$-independent in triaxial models in the limit of large $N$ ($N \gtrsim 5\times10^6$). Rates are always lower than the full loss cone rate, but in triaxial models they are sufficient to drive BHBs to coalescence in less than a Hubble time. Axisymmetric models, however, have hardening times that are too long to be of interest. This is in contradiction with the results of \citet{khan2013} for axisymmetric models, and the reason for the discrepancy remains unknown, though likely of numerical origin. In this study, we present an alternative approach to model the loss cone refilling of BHBs. By means of direct summation simulations of galaxy mergers, we show that a good proxy for the binary hardening rate is given by the integrated number of stars with angular momentum smaller than the loss cone angular momentum of the binary. We estimate this at the hard-binary separation, counting each star only once in the time interval of interest. This is because we expect stars to be scattered out of the loss cone following an interaction. We then consider different isolated galaxy models, spherical and flattened, and follow their evolution with both a direct summation code and a tailored fast multiple method. We show that the hardening rate is a strong function of particle number $N$ in spherical models, with a scaling consistent with expectations for collisional processes. On the other hand, the hardening rate becomes independent of N for $N \gtrsim 5\times 10^6$, in triaxial models. The behaviour of axisymmetric models is intermediate between that of spherical and triaxial models, with a slow but significant dependence on $N$. We reach values of $N$ equal of 2 million particles with direct summation and 64 million particles with the fast multiple method. Efficient loss cone refilling is seen even in mildly triaxial models (with axis ratios $1:0.9:0.8$). Such triaxiality is unavoidable following galaxy mergers and drives binaries into the gravitational waves regime. We conclude therefore that there is no `final parsec problem' in the evolution of BHBs even in gas-free systems due to the efficiency of collisionless loss cone refilling in triaxial potentials. | \label{sec:disc} Black hole binary hardening in simulations of galaxy mergers is believed to be sustained by a collisionless mode of loss cone refilling which owes to global torques in non-spherical potentials. This is necessary to ensure hardening down to separations where emission of gravitational waves becomes dominant and leads the black holes to coalescence. Merger remnants typically show a significant degree of flattening and a modest departure from axisymmetry \citep{preto2011,khan2011,GM2012}, which argues for a significant population of stars on centrophilic orbits. However, if this is indeed the case then binary hardening should also be efficient in isolated triaxial models. Simulations by \citet{VAM2014} show that collisionless losscone refilling is masked by collisional refilling due to stellar scatterings in simulations with modest particle numbers and small softening. Here we take a different approach to the problem and study losscone refilling in isolated galaxy models. In order to reduce the effects of collisionality we perform simulations with the fast multiple method code {\Griffin} \citep{Dehnen2014}, which allows us to increase particle number to 64 million. We find a proxy for binary hardening, which is given by the refilling parameter, i.e. the fraction of stars with angular momentum smaller than the angular momentum of an hypothetical BHB of given mass, where each star is counted only once per simulation. Our key findings are: \begin{itemize} \item The refilling parameter, i.e. the fraction of stars which can be counted at least once to have angular momentum smaller than the loss cone angular momentum, is a good proxy for the hardening rate in merger simulations. \item Loss cone refilling in spherical models depends critically on particle number, a clear signature that it is driven by two body scatterings. In real galaxies, this process becomes extremely inefficient and BHBs are not expected to reach the gravitational wave phase. \item There is no $N$-dependence of the refilling parameter in triaxial models above $N \sim 10^7$. Refilling is more efficient in triaxial models than in spherical models, and a higher degree of triaxiality also leads to more efficient refilling. \item Axisymmetric models have properties in between spherical and triaxial models. While refilling is consistently more efficient than in spherical cases, we observe a marked $N$-dependence with no obvious flattening even at the largest $N$ values. \item Hardening rates computed directly from the $N$-body data in isolated triaxial models match those computed in merger simulations. On the other hand, spherical isolated models and axisymmetric models have significantly lower hardening rates. \item The hardening rates measured for the triaxial models are large enough to ensure coalescence of the binaries within a Hubble time for Milky Way type galaxies as well as more massive ones. Hardening rates in axisymmetric models are only marginally sufficient to bridge the gap to the gravitational waves regime in the case of low mass binaries ($\mbin \lesssim 10^6 \msun$) but imply coalescence times that are longer than a Hubble time for typical binaries ($\mbin \sim 10^8 \msun$). Spherical models are generally characterised by hardening rates too low to lead to coalescence. \end{itemize} | 16 | 9 | 1609.09383 |
1609 | 1609.06747_arXiv.txt | We summarize the progress in neutrino astrophysics and emphasize open issues in our understanding of neutrino flavor conversion in media. We discuss solar neutrinos, core-collapse supernova neutrinos and conclude with ultra-high energy neutrinos. | \noindent Nature has provided us with a variety of neutrino sources, from the not yet observed 1.9 K cosmological background to the IceCube PeV neutrinos \cite{Aartsen:2014gkd}, whose origin is still mysterious. Neutrinos are intriguing weakly interacting particles. After 1998 many unknown properties have been determined thanks to the discovery of neutrino oscillations, first proposed in \cite{Pontecorvo:1957cp} and observed by the Super-Kamiokande experiment using atmospheric neutrinos \cite{Fukuda:1998mi}. This discovery is fundamental for particle physics, for astrophysics and for cosmology. Neutrino oscillations is an interference phenomenon among the $\nu$ mass eigenstates, that occurs if neutrinos are massive and if the mass (propagation basis) and the flavor (interaction basis) do not coincide. The Maki-Nakagawa-Sakata-Pontecorvo matrix relates these two basis \cite{Maki:1962mu}. Within three active flavors, such a matrix depends on three mixing angles, one Dirac and two Majorana CP violating phases. In the last two decades solar, reactor and accelerator experiments have precisely determined most of the oscillation parameters, including the so-called atmospheric $\Delta m_{23}^2 = m_{3}^2 - m_{2}^2 = 7.6 \times 10^{-3} $eV$^2$, and solar $\Delta m_{12}^2 = m_{2}^2 - m_{1}^2 = 2.4 \times 10^{-5} $eV$^2$ mass-squared differences \cite{Agashe:2014kda}. Moreover the sign of $\Delta m^2_{12}$ has been measured since $^{8}$B neutrinos undergo the Mikheev-Smirnov-Wolfenstein (MSW) effect \cite{Wolfenstein:1977ue,Mikheev:1986gs} in the Sun \cite{Ahmad:2002jz,Eguchi:2002dm,Robertson:2012ib}. The sign of $\Delta m_{23}^2$ is still unknown, either $\Delta m_{31}^2 > 0$ and the lightest mass eigenstate is $m_1$ (normal ordering or "hierarchy"), or $\Delta m_{31}^2 < 0$ it is $m_3$ (inverted ordering). Most of neutrino oscillation experiments can be interpreted within the framework of three active neutrinos. However a few measurements present anomalies that require further clarification. Sterile neutrinos that do not couple to the gauge bosons but mix with the other active species could be the origin of the anomalies. Upcoming experiments such as STEREO or CeSox will cover most of the mixing parameters identified in particular by the "reactor anomaly" \cite{Mention:2011rk}. Among the fundamental properties yet to be determined are the mechanism for the neutrino mass, the absolute mass value and ordering, the neutrino nature (Dirac versus Majorana), the existence of CP violation in the lepton sector and of sterile neutrinos. The combined analysis of available experimental results shows a preference for normal ordering and for a non-zero CP violating phase, currently favouring $\delta = 3 \pi/2$, although statistical significance is still low \cite{Marrone:Nu2016}. In the coming decade(s) experiments will aim at determining the mass ordering, the Dirac CP violating phase, the neutrino absolute mass and hopefully nature as well. Moreover Super-Kamiokande with Gadolinium should have the sensitivity to discover the relic supernova neutrino background \cite{Xu:2016cfv}. | 16 | 9 | 1609.06747 |
|
1609 | 1609.01167_arXiv.txt | {Relativistic shocks are one of the most plausible sites of the emission of strongly variable, polarized multi-wavelength emission from relativistic jet sources such as blazars, via diffusive shock acceleration (DSA) of relativistic particles. This paper summarizes recent results on a self-consistent coupling of diffusive shock acceleration and radiation transfer in blazar jets. We demonstrate that the observed spectral energy distributions (SEDs) of blazars strongly constrain the nature of hydromagnetic turbulence responsible for pitch-angle scattering by requiring a strongly energy-dependent pitch-angle mean free path. The prominent soft X-ray excess (``Big Blue Bump'') in the SED of the BL Lac object AO 0235+164 can be modelled as the signature of bulk Compton scattering of external radiation fields by the thermal electron population, which places additional constraints on the level of hydromagnetic turbulence. It has further been demonstrated that internal shocks propagating in a jet pervaded by a helical magnetic field naturally produce polarization-angle swings by 180$^o$, in tandem with multi-wavelength flaring activity, without requiring any helical motion paths or other asymmetric jet structures. The specific application of this model to 3C279 presents the first consistent, simultaneous modeling of snap-shot SEDs, multi-wavelength light curves and time-dependent polarization signatures of a blazar during a polarization-angle (PA) rotation. This model has recently been generalized to a lepto-hadronic model, in which the high-energy emission is dominated by proton synchrotron radiation. It is shown that in this case, the high-energy (X-ray and $\gamma$-ray) polarization signatures are expected to be significantly more stable (not showing PA rotations) than the low-energy (electron-synchrotron) signatures. } \keyword{Active galaxies: BL Lac objects; Jets and bursts; Radiation mechsnisms: polarization; Magnetohydrodynamics and plasmas} \begin{document} | 16 | 9 | 1609.01167 |
||
1609 | 1609.02508_arXiv.txt | A 3.56-hour white dwarf (WD) - M dwarf (MD) close binary system, AR Scorpii, was recently reported to show pulsating emission in radio, IR, optical, and UV, with a 1.97-minute period, which suggests the existence of a WD with a rotation period of 1.95 minutes. We propose a model to explain the temporal and spectral characteristics of the system. The WD is a nearly perpendicular rotator, with both open field line beams sweeping the MD stellar wind periodically. A bow shock propagating into the stellar wind accelerates electrons in the wind. Synchrotron radiation of these shocked electrons can naturally account for the broad-band (from radio to X-rays) spectral energy distribution of the system. | A white dwarf (WD) - M dwarf (MD) binary, AR Scorpii (henceforth AR Sco), was recently reported to emit pulsed broad-band (radio, IR, optical and UV) emission \citep{Marsh16}. The brightness of the system varies in long time scales with the orbital period of 3.56 hours, and shows pulsation in short time scales with a period of 1.97 minutes. Interpreting the high pulsating frequency as the ``beat'' frequency of the system, the inferred rotation period of the WD is 1.95 minutes. The spectral energy distribution of the pulsed emission supports a synchrotron origin for the radiation \citep{Marsh16}. These peculiar observational properties make AR Sco a unique system. Although AR Sco could be in the evolutionary stage of the so-called intermediate polar \citep{Bookbinder87,Patterson94,Oruru12}, of which accretion is the main power source, the absence of accretion features in AR Sco demands another mechanism to explain the observations. It has been suggested that if the dipole magnetic field of a WD is strong enough, it would behave like a WD radio pulsar \citep{Zhang05}.\footnote{The so-called anomalous X-ray pulsars were suggested as magnetized WDs \citep{Paczynski90,Usov93,Malheiro12,Lobato16,Mukhopadhyay16}, but they are now widely accepted to be extremely magnetized neutron stars known as magnetars \citep{Duncan92,Thompson96}. } The unique pulsating properties of AR Sco suggest that such WD pulsars indeed exist. Here we propose a WD-MD interaction model for the system. We show that the interaction between the WD pulsar open field line beams with the stellar wind of MD naturally accounts for all the observational properties of the system. | We have shown that the peculiar observations of the pulsating AR Sco system can be understood with the framework of interaction between the WD pulsar's open field line beams and the wind of the MD. The observational data demand a nearly perpendicular rotator for the WD pulsar, and a near edge-on orbital configuration for the observer on Earth. In order to interpret the observed SED, the required electron number density is too high for a WD wind. Rather electrons accelerated by a bow shock into the MD wind can produce the right amount of electrons to interpret both the shape and the normalization of the SED. In our model, although the magnetic field lines of the WD are likely ordered, an observer sees a hemisphere where magnetic field lines have different directions so that on average the directional information cancels out (e.g. in Fig. 2, the field lines above and below point B have opposite orientations). One would therefore do not expect significant circular polarization \citep{Matsumiya03}, which is consistent with the observations \citep{Marsh16}. Our model suggests that rapidly rotating, highly magnetized WDs can indeed behave like radio pulsars, as has been speculated in the past \citep{Zhang05}. The rarity of these WD pulsars \citep{Kepler13} may be due to the conditions to produce an active magnetosphere via pair production are much more stringent for WDs than NSs. The peculiarity of AR Sco lies in its extremely short period and its close proximity with its MD companion. According to our modeling, the observed emission is from the shocked MD wind rather than from the WD pulsar itself. However, if some WDs indeed behave as pulsars, one would expect to directly detect emission from WD pulsars in the future. GCRT J1745-3009 might be another, less energetic, transient WD pulsar \citep{Zhang05} at a distance beyond 1 kpc from the earth \citep{Kaplan08}. | 16 | 9 | 1609.02508 |
1609 | 1609.03280_arXiv.txt | We conducted VLA C-configuration observations to measure positions and luminosities of Galactic Class II 6.7~GHz methanol masers and their associated ultra-compact \ion{H}{2} regions. The spectral resolution was 3.90625 kHz and the continuum sensitivity reached 45 \uJypb. We mapped 372 methanol masers with peak flux densities of more than 2 Jy selected from the literature, 367 of them were detected. Absolute positions have nominal uncertainties of 0.3\arcsec. In this first paper on the data analysis, we present three catalogs, the first gives information on the strongest feature of 367 methanol maser sources, and the second on all detected maser spots. The third catalog present derived data of the 279 radio continuum sources found in the vicinity of maser sources. Among them, 140 show evidence of physical association with maser sources. Our catalogs list properties including distance, flux density, radial velocity and the distribution of masers on the Galactic plane is then provided as well. We found no significant relationship between luminosities of masers and their associated radio continuum counterparts. | \label{sec:introduction} Although massive stars make up only a few percent of the stellar population in the Milky Way, they play a central role in many astrophysical processes, such as shaping the interstellar medium, regulating star formation and ultimately governing the evolution of their host galaxy~\citep{2005IAUS..227....3K}. However, to date, understanding massive star formation remains a challenge, because of numerous observational problem. Nevertheless, investigations of the natal environment of massive stars provide rich clues for a better understanding of its physical and chemical properties and the star formation process itself. Massive star forming regions (MSFRs) offer a variety of tracers that can be detected in different wavebands, including massive dense cores~\citep{2007ARA&A..45..481Z}, infrared dark clouds~\citep{2010ApJ...723L...7K}, mid-infrared sources~\citep{2013ApJS..205....1P}, Extended Green Objects~\citep[EGOs,][]{2008AJ....136.2391C}, molecular outflows~\citep{2002A&A...383..892B}, masers~\citep[most common species are hydroxyl, methanol and water;][]{2014IJAA....4..571F}, compact and ultra-compact~(UC) \ion{H}{2} regions~\citep{2002ARA&A..40...27C} and others. Understanding what these signposts tell us may help to construct an evolutionary sequence for young massive stars. Methanol masers at 6.7~GHz were discovered by \citet{1991ApJ...380L..75M}; they, along with 22.2 GHz H$_2$O masers, are the brightest and most widespread maser species in massive star forming regions~\citep[e.g.][]{1995MNRAS.272...96C,2003A&A...403.1095M, 2008A&A...485..729X}. These masers may be one of the earliest and ubiquitous tracers of the massive star formation process~\citep{2015MNRAS.446.3461U}. Most 6.7~GHz methanol masers have been observed with a variety of telescopes and with differing resolutions. For example, \citet{2011ApJ...730...55P} observed 57 methanol masers at 6.7~GHz with the Multi-Element Radio-Linked Interferometer Network~(MERLIN) with an angular resolution of 60~mas. At this resolution they found a close correspondence between the masers and mid-infrared sources and claimed their results to support the theoretical model that these masers are pumped by infrared dust emission in the vicinity of massive proto-stars. Compact and ultra-compact~(UC) \ion{H}{2} regions are clear indicators of sites of massive stars~\citep{2005dmgp.book..203H}, whose ionizing photons are provided by embedded OB stars in the later stages of massive star formation~\citep{2002ARA&A..40...27C}. The thermal bremstahlung radiation from the ionized gas makes \ion{H}{2} regions directly observable in radio continuum bands. In recent decades, a number of radio survey have identified many hundreds of UC \ion{H}{2} regions. For example, \citet{2009A&A...501..539U} observed 659 MSFR candidates with the VLA at 4.9~GHz and an angular resolution of 1.5", achieving RMS noise levels of $\approx0.2$~\mJypb\ and detected 391 compact or UC~\ion{H}{2} regions. Many observations~\citep[e.g.][]{2009PASA...26..454C,2010MNRAS.407.2599C, 2010A&A...517A..78S} have established an association of these \ion{H}{2} regions with 6.7~GHz methanol masers. Since compact and UC~\ion{H}{2} regions present a reliable snapshot of MSFRs withing the last $\sim10^5$ years~\citep{2011MNRAS.416..972D,2011ApJ...730L..33M}, a census on this later phase of MSFRs can help us to better understand the evolutionary sequence of massive star formation and their association with different tracers. In order to study the relationship of 6.7~GHz methanol masers to compact and UC \ion{H}{2} regions, we conducted a large survey to simultaneously observe the methanol maser line and C-band continuum emission with the Karl G. Jansky Very Large Array~(VLA) in its C-configuration. We mapped 372 class II methanol maser sources at 6.7~GHz, selected from literature with peak fluxes $>2$~Jy. Absolute positions were established based on phase-referenced observations to astrometric calibrators. Our observations also served as precursor observations for the Bar and Spiral Structure Legacy~(BeSSeL) Survey~\citep{2011AN....332..461B} to provide high accuracy position of maser sources necessary for VLBA observations. Our uniform and high angular resolution survey includes the vast majority of known 6.7~GHz methanol masers and their nearby compact or UC~\ion{H}{2} regions. The survey produced three catalogs: (1) a methanol maser catalog containing accurate positions, radial velocities and flux densities of the brightest spots; (2) a catalog of maser spectra and spot maps; (3) a catalog of \ion{H}{2} regions toward the methanol masers. Compact \ion{H}{2} (referring to both compact and UC \ion{H}{2}) regions and methanol masers offer a unique opportunity to peer into the deep interior of massive star forming regions at high angular resolution since at centimeter wavelength their natal clouds are optically thin, and the large database should allow statistically meaningful estimates of many characteristics of 6.7~GHz methanol masers and identified \ion{H}{2} region counterparts. In combination with other signposts of massive star formations, the catalogs may provide important clues to solve several key issues related to massive star formation, such as whether 6.7~GHz methanol masers are located in accretion disks or in outflow walls (or both), when in the process of star formation these two signposts occur, how MSFRs evolve with time. Some prominent class II methanol masers are found in the dense molecular material associated UC \ion{H}{2} regions, W3OH and NGC~7538 being prominent examples~\citep{1988ApJ...333L..83M}. However, most have no radio continuum emission at the, say, few mJy level, that could be easily detected with the original VLA or the ATCA~\citep[e.g.][]{1998MNRAS.301..640W, 2001A&A...369..278M}. This leaves open the question of whether weaker radio emission could be associated with masers at the level that could be expected from younger, hyper compact \ion{H}{2} regions. In Section~\ref{sec:Observ} we describe the sample and the design of the observations. The data analysis is described in Section~\ref{sec:Reduc}. The three catalogs mentioned above are presented in Section~\ref{sec:Produc}. In Section~\ref{sec:Prop}, statistical properties of the 6.7~GHz methanol masers and their associated compact \ion{H}{2} regions are discussed. A brief summary is given in Section~\ref{sec:Summary}. This paper is the first of a series investigating massive star formation and its evolutionary sequence. | \label{sec:Summary} We report a survey toward a large sample of Galactic class II 6.7~GHz methanol masers and their C-band radio continuum counterparts. With the VLA in its C-configuration, we observed 372 targets selected from a variety of publications and produced a catalog of 367 methanol masers with sub arc~second level positional accuracy, and high resolution radial velocity. We believe our catalog is a uniform and comprehensive collection of class II methanol maser in the first and second Galactic quadrants. High accuracy J2000 coordinates, radial velocities and fluxes are presented. Maser spots detected for each source are reported in 367 individual on-line available files, and we plot spot maps and spectra for each maser source. We also search for compact radio continuum emission toward all masers. This resulted in detections of 279 radio continuum sources, of which 157 are likely to be physically associated with a maser source. Statistical characteristics of various source properties are presented. These three catalogs may aid future investigations of massive star forming and the maser process. | 16 | 9 | 1609.03280 |
1609 | 1609.03249_arXiv.txt | We present an application of machine-learning (ML) techniques to source selection in the optical transient survey data with Hyper Suprime-Cam (HSC) on the Subaru telescope. Our goal is to select real transient events accurately and in a timely manner out of a large number of false candidates, obtained with the standard difference-imaging method. We have developed the transient selector which is based on majority voting of three ML machines of AUC Boosting, Random Forest, and Deep Neural Network. We applied it to our observing runs of Subaru-HSC in 2015 May and August, and proved it to be efficient in selecting optical transients. The false positive rate was 1.0\% at the true positive rate of 90\% in the magnitude range of 22.0--25.0 mag for the former data. For the latter run, we successfully detected and reported ten candidates of supernovae within the same day as the observation. From these runs, we learned the following lessons: (1) the training using artificial objects is effective in filtering out false candidates, especially for faint objects, and (2) combination of ML by majority voting is advantageous. | The 8.2-m Subaru telescope has been running a 300-night Strategic Survey Program (SSP) over 5 years since 2014 March\footnote{http://www.naoj.org/Projects/HSC/surveyplan.html}, in order to elucidate the mystery of dark matter and dark energy as well as the evolution of galaxies. The survey utilizes Hyper Suprime-Cam (HSC; \cite{Miyazaki+2012}) with a wide field of view of 1.77 square degrees. In 2016--18, we observe two ultra deep fields, COSMOS and SXDS, for six-month each. The observations will be performed around the new moon in each month, with a typical cadence of 3--4 days. Note that the HSC is not installed on Subaru for two weeks around the full moon. Among other transient surveys, the Subaru HSC/SSP survey is the deepest for this survey area ($1.77 \times 2$ square degrees), and thus, will provide the unique dataset for transients. For example, the HSC/SSP survey will triple the number of type-Ia supernovae (SNe) beyond redshift $z>1$ and will also discover a few tens of superluminous SNe at $z>1$. Difference imaging is the standard method to search for optical transient objects, and so we use it in this study. We define a transient object as the one that appears only in the image at the later epoch (newer image), but not in the one at the earlier epoch (reference image or template image), out of the two images taken at different epochs. After the standard data reduction is made, two images are astrometrically aligned, and the reference image is subtracted from the newer image by matching point spread functions (PSFs) of the two images. The source-finding algorithm is applied to the difference image, and the detected sources are the candidates of transients (section \ref{sec: Data Analysis and Feature Extraction}). \begin{figure} \begin{center} \includegraphics[width=7cm, angle=0]{bogus_real_sample.eps} \end{center} \caption{Examples of real and bogus objects obtained with Subaru-HSC. The left, middle, and right columns show the reference, new, and difference images, respectively. The first, second, and third rows show the cosmic ray (a--c), ghost near a bright star (d--f), and inaccurate image convolution or astrometric alignment (g--i), respectively. The bottom row shows a real transient located in a galaxy (j--l). }\label{fig:bogus_sample} \end{figure} In an ideal situation, all the sources detected in difference images would be transient/variable astronomical sources, such as supernovae, variable stars, moving objects, and so on. In reality, however, they also include artifacts (see panels a--i in figure \ref{fig:bogus_sample}), such as cosmic-ray events, spikes around bright stars, and residuals related to inaccurate image convolution or astrometric alignment. These artifacts are present in every optical survey project \citep{Bailey+2007, Bloom+2012, Brink+2013}. Hereafter, we call them ``bogus'' \citep{Bloom+2012}. In the HSC/SSP survey, not only a few hundred transients, including SNe, but also $\sim 10^5 - 10^6$ bogus objects are expected to be detected each night. After the scheduled 300 nights, the number of candidates of transients, real and bogus combined, will reach $\sim 10^8$, which is well qualified as Big Data. We need to filter out bogus objects to select SNe and other real transients. Processing of filtering must be performed swiftly in order to increase the chances of new findings in an early phase of transient phenomena. The primary method to distinguish transient/bogus objects is, traditionally, visual inspection by human checkers, as many surveys have been adopting. However, the expected size of our data is so big that human checkers would not be able to go through all the data in a reasonable time. We have decided to introduce machine-learning (ML) techniques to select real transients. In the filtering process, we should not miss real objects, while a vast number of bogus objects are filtered out. Throughout the development, we try to minimize the false positive rate (FPR), while we maintain the true positive rate (TPR) of 90 \% or larger (namely the false negative rate, FNR $< 10\%$). We performed two HSC observations in 2015 May and August \citep{tominaga15atel1, tominaga15atel2}, aiming to detect short transients with a time scale of a few hours to a few days (e.g. \citet{Tanaka+2016}; \citet{Morokuma+2016}). An example of such short transients is an optical flash at the time of shock breakout of a supernova, of which the time-variance would be detectable during an observation for a single night. We use three kinds of ML methods, AUC Boosting, Random Forests (RF), and Deep Neural Network (DNN), both individually and in a combined way. We verify the performance of these machines by making the receiver operation characteristic (ROC) curves. Conditions of observations (the noise and seeing) vary every night. Hence, the ML classifiers must be robust against the change of environment. We use normalized features to reduce the influence of the variation. To validate the performance of our method, we show the ROC curves of our machine trained with the data on one-night observation, applied to the data on the other night observation. The results show the proposed classifier is robust. In optical surveys, ML techniques have been introduced by \citet{Bailey+2007} for the data from the Nearby Supernova Factory. They applied Boosted Decision Trees, RF, and Support Vector Machines (SVM) and succeeded in reducing the number of bogus candidates by a factor of ten. The Palomar Transient Factory team \citep{Bloom+2012} used RF, and achieved the TPR of 92.3\% at the FPR of 1\% \citep{Brink+2013}. The Pan-STARRS1 Medium Deep Survey used Artificial Neural Network, SVM and RF, and achieved TPR of 90\% at FPR of 1\% \citep{Wright+2015}. \citet{Goldstein+2015} applied the RF for Dark Energy Survey Supernova program (DES-SN), and reduced the number of transient candidates by a factor of 13.4, which were then fed to human scanning. \citet{du Buisson+2015} applied the RF, $k$-nearest neighbor, and the SkyNet artificial neural net algorithm, using features trained from eigen-image analysis for the Sloan Digital Sky Survey supernova survey. The rest of this paper is structured as follows. In section~\ref{sec: Data Analysis and Feature Extraction}, we explain the HSC data reduction and feature extractions. In section~\ref{sec: ML}, we introduce three machine learning methods we used. In section~\ref{sec: experiment}, the applications to the actual Subaru data are presented, and then we discuss the result of real vs bogus segregation in section~\ref{sec: Discussion and Conclusion}. Prior to the forthcoming HSC/SSP Transient survey, this paper will provide a `path finder' to identify real astronomical objects and to demonstrate the power of machine learning. | \label{sec: Discussion and Conclusion} \begin{figure*} \begin{center} \includegraphics[width=16cm, angle=0]{score_scat3_sel10.eps} \end{center} \caption{Cross-relations of scores obtained with the three machines for the entire sample (black) and the selected sample of ten candidates of supernovae reported to Astronomers' Telegram \citep{tominaga15atel2} (red) for the observation data on 2015 August\ 19. The left, middle, and right panels plot those of Random Forest v.s. AUC Boosting, AUC Boosting v.s. Deep Neural Network, and Deep Neural Network v.s. Random Forest, respectively. Horizontal and vertical dashed lines show the thresholds corresponding to the TPR of 90\%.} \label{fig:score_scat3_sel10} \end{figure*} We performed Subaru/HSC observations in 2015 May and August, made transient searches, and applied machine-learning techniques to the result to reduce the bogus transient objects. We have developed real-bogus classifiers as the core function for it, using the three machine-learning methods of AUC Boosting, Random Forests, and Deep Neural Network, and then made the combined classifier as their majority voting. We have installed our machines in the analysis pipeline of the HSC, and successfully found real supernovae within the same day as the observation, demonstrating the power of our method. Now, we have completed the preparation for the forthcoming HSC/SSP transient survey observation. In training our machines, we used artificial objects, because the data are highly imbalanced between real/bogus objects, and it was found to be crucial to make good machines efficiently. Although the HSC survey data are technically more difficult to deal with than other survey data, given that the HSC survey is deeper than other surveys, we have achieved the results comparable to the similar studies in other surveys. The cross-relations between three machines in figure \ref{fig:score_scat3_sel10} show that none of the three is significantly better than the other two. Therefore, combining the machines is beneficial. We used a moderate selection with all the combinations of multiple machines, ``majority voting'', to avoid missing some real objects. For the machine, robustness against variation of environment was confirmed. In this paper we have focused on real-bogus separation with machine-learning methods, and have demonstrated that they were indeed useful. Their use in extracting scientific results out of Big-data in astronomy is promising. As the next step, we will use machine-learning for classification of types of transients, by combining timing and color information, as well as the shape of objects. \begin{ack} We thank the Subaru Hyper Suprime-Cam team. This work is supported by Core Research for Evolutionary Science and Technology (CREST), Japan Science and Technology Agency (JST). It is also supported by the research grant program of Toyota foundation (D11-R-0830) and was in part supported by Grants-in-Aid for Scientific Research of JSPS (15H02075), MEXT (15H00788), and the World Premier International Research Center Initiative, MEXT, Japan. This paper makes use of software developed for the LSST. We thank the LSST Project for allowing their code available as free software at http://dm.lsstcorp.org. \end{ack} \appendix | 16 | 9 | 1609.03249 |
1609 | 1609.03908_arXiv.txt | {} {We aim to present {\coco a} generalized Bayesian inference method for constraining interiors of super Earths and sub-Neptunes. Our methodology succeeds in quantifying the degeneracy and correlation of structural parameters for high dimensional parameter spaces. Specifically, we identify what constraints can be placed on {\coco composition and thickness of core, mantle, ice, ocean, and atmospheric layers} given observations of mass, radius, and bulk refractory abundance constraints (Fe, Mg, Si) from observations of the host star's photospheric composition. } {We employed a full probabilistic Bayesian inference analysis that formally accounts for observational and model uncertainties. Using a Markov chain Monte Carlo technique, we computed joint and marginal posterior probability distributions for all structural parameters of interest. We included state-of-the-art structural models {\coco based on} self-consistent thermodynamics of core, mantle, high-pressure ice, and liquid water. Furthermore, we tested and compared two different atmospheric models that are tailored for modeling thick and thin atmospheres, respectively.} {First, we validate our method against Neptune. Second, we apply it to synthetic exoplanets of fixed mass and determine the effect on interior structure and composition when (1) radius, (2) atmospheric model, (3) data uncertainties, (4) semi-major axes, (5) atmospheric composition (i.e., a priori assumption of enriched envelopes versus pure H/He envelopes), and {\coco (6) prior distributions} are varied.} {Our main conclusions are: (1) Given available data, the range of possible interior structures is large; quantification of the degeneracy of possible interiors is therefore indispensable for meaningful planet characterization. (2) Our method predicts models that agree with independent estimates of Neptune's interior. (3) Increasing the precision in mass and radius leads to much improved constraints on ice mass fraction, size of rocky interior, but little improvement in the composition of the gas layer, whereas an increase in the precision of stellar abundances enables to better constrain mantle composition and relative core size. (4) For thick atmospheres, the choice of atmospheric model can have significant influence on interior predictions, {\coco including} the rocky and icy interior. The preferred atmospheric model is determined by envelope mass. This study provides a methodology for rigorously analyzing general interior structures of exoplanets which may help to understand how exoplanet interior types are distributed among star systems. This study is relevant in the interpretation of future data from missions such as TESS, CHEOPS, and PLATO.} | \sloppy The characterization of planet interiors is one of the main foci of current exoplanetary science. For the characterization of super Earths and sub-Neptunes, we mostly rely on mass and radius measurements. Direct measurements of atmospheres are, thus far, mostly limited to transiting hot Jupiters and {\coco a} few Sub-Neptunes \citep{iyer}, with the exception of super Earth 55 Cnc E \citep{Tsiaras, demory}. For interior characterization, common practice is the use of mass-radius-plots where mass and radius of exoplanets are compared to synthetically computed interior models \citep[e.g.,][]{sotin07, seager2007, fortney, dressing, howe}. {\coco However}, it is difficult to know (1) how well one interior model compares with the generally large number of other possible interior scenarios that also fit data and (2) which structural parameters can actually be constrained by the observations. Thus, this approach fails to address {\coco the degeneracy problem} that is, that different interior models can have identical mass and radius. In order to draw meaningful conclusions about an exoplanet's interior it is therefore necessary to account for this inherent degeneracy {\citep[e.g.,][]{rogers2010, schmitt2014, carter, weiss, dorn}}. The Bayesian analysis of \citet{rogers2010} to exoplanets of three to four parameters was generalized for purely rocky exoplanets by \citet{dorn}. Here, we extend the full probabilistic analysis of \cite{dorn} to more general interior structures by including volatile elements in form of icy layers, oceans, and atmospheres. {\coco The previous work of \citet[][]{rogers2010} uses a grid search method which calls for strong a priori assumptions on structure and composition of exoplanets to significantly reduce the parameter space. However, the number of parameters that affect mass and radius is large (e.g., it comprises composition and size of core, mantle, ice layers, and gas, as well as internal energy). Here, we present a generalized Bayesian inference scheme that incorporates the following aspects}: \begin{itemize} \itemsep0pt \item Our method is applicable to a wide range of planet-types, including rocky super Earths and sub-Neptunes. \item We employ a full probabilistic Bayesian inference analysis using a Markov chain Monte Carlo (McMC) technique to constrain core size, mantle thickness and composition, mass of water-ice, and key characteristics of the atmosphere (e.g., mass, intrinsic luminosity, composition). \item We test two different atmospheric models, tailored to thick and thin atmospheres, that account for enrichments in elements heavier than H and He. \item We employ state-of-the-art modeling to compute interior structure based on self-consistent thermodynamics for a pure iron core, a silicate mantle, high-pressure ice, water ocean, and atmosphere (to some extent). \item {\coco Compared to previous work of \citet[][]{rogers2010}, our scheme can also be used for high dimensional parameter spaces.} \end{itemize} Besides mass and radius estimates, additional constraints are crucial to reduce model degeneracy \citep[e.g.,][]{dorn, grasset09}. \citet{dorn} demonstrate that the use of relative bulk abundance constraints of Fe/Si and Mg/Si taken from the host star (henceforth referred to as abundance constraints) leads to much improved constraints on core size and mantle composition in the case of purely rocky exoplanets. The validity of a direct correlation between stellar and planetary relative bulk abundances is suggested by observational solar system studies and planet formation models \citep{carter, lodders03,drake,mcdono, bond,elser,johnson,thiabaud}. Here, we also assume solar bulk abundance constraints based on spectroscopic measurements \citep{lodders03}. Our generalized interior structure model is based on previous studies of mass-radius relations. Generally, H$_2$O in liquid and high-pressure ice form \citep[e.g.,][]{valencia07a, seager2007}, and H$_{2}$-He atmospheres \citep[e.g.,][]{rogers2011, fortney} are considered. Although it would not be surprising if the compositional diversity of ices and atmospheres exceeds the one found in the solar system \citep[e.g.,][]{newsom}, the few observational data on exoplanets limit us to relatively simple planetary interior models. The structural parameters that we {\coco investigate} include: (1) internal energy, mass, and composition of the gas layer, (2) mass and temperature of the ice layer, (3) mantle size and composition, and (4) core size. For present purposes, we assume a general planetary structure consisting of a pure iron core, a silicate mantle, a water ice layer and an atmosphere. To compute the resultant density profile for the purpose of estimating mass and radius, we follow \citet[][]{dorn} and assume hydrostatic equilibrium coupled with a thermodynamic approach based on Gibbs free-energy minimization and Equation-of-State (EoS) modeling. In this study, we wish to quantify the influence of the following parameters on predicted interior structure and composition: (1) planet radius, (2) data uncertainty (e.g., mass, radius, bulk abundances), (3) semi-major axis, (4) atmospheric model, (5) atmospheric composition (i.e., a priori assumption of enriched envelopes versus pure H/He envelopes), and {\coco (6) prior distributions.} {\coco In a companion paper \citep{dornA}, we present results on the application of our proposed method to six exoplanets (\mbox{HD 219134b}, Kepler-10b, Kepler-93b, CoRoT-7b, 55 Cnc e, and \mbox{HD 97658b}) for which spectroscopic measurement of their host star's photospheres are available \citep{hinkel}.} The outine of this study is as follows: we describe the iterative inference scheme (Section \ref{inversion}), model parameters (Section \ref{parametrization}), data (Section \ref{data}), and the forward model (Section \ref{model}). In Section 3, we validate our method against Neptune and present results for different synthetic planet cases. In Sections \ref{Discussion} and \ref{Conclusions}, we discuss results and conclude. | \label{Conclusions} {\coco We present a} generalized inference method that enables us to make meaningful statements about the interior structure of observed exoplanets. Our full probabilistic Bayesian inference analysis formally accounts for data and model uncertainties, as well as model degeneracy. By employing a Markov chain Monte Carlo technique, we quantify the state of knowledge that can be obtained on composition and thickness of core, mantle, water ice, and gaseous layers for given data of mass, radius, and bulk abundance proxies for $\fesi$ and $\mgsi$ obtained from spectroscopic measurements. We have built upon the work of \citet{dorn} and extended the dimensionality of the interior characterization problem to include volatile elements in the form of gas, water ice and ocean. {\coco Our method succeeds at constraining planet interior structure even for high dimensional parameter spaces and thereby overcomes limitations of previous works on mass-radius relationship of exoplanets.} We have validated our method against Neptune. Using synthetic planets, we have determined how predictions on interior structure depend on various parameters: bulk density, data uncertainties, semi-major axes, atmospheric composition (i.e., a priori assumption of enriched envelopes versus pure H/He envelopes), {\coco and prior distributions}. Furthermore, we have investigated two different atmosphere models and quantify how parameter estimates depend on the choice of the atmosphere model. We summarize our findings as follows: \begin{itemize}\itemsep0pt \item It is possible to constrain core size, mantle size and composition, mass of water ice, and key characteristics of the gas layer (e.g., internal energy, mass, composition), given observations of mass, radius, and bulk abundance proxies $\fesi$ and $\mgsi$ taken from the host star. \item A Bayesian analysis is key in order to rigorously analyse planetary interiors, as it formally accounts for data and model uncertainty, as well as the inherent degeneracy of the problem addressed here. The range of possible interior structures is large even for small data uncertainties. Our method is able to quantify the probability that a planet is rocky and/or volatile-rich. \item Our method has been successfully validated against Neptune for which independent structure estimates based on geophysical data (e.g., gravitational and magnetic moments) are available. \item Model parameters {\coco have different sensitivity to the various data}. Constraints on bulk abundances $\fesi$ and $\mgsi$ determine relative core size and mantle composition. Mass mostly determines the size of the rocky and icy interior, whereas radius mainly determines structure and composition of the gas and the water ice layers. \item Increasing precision in mass and radius leads to a much better constrained ice mass fraction, size of rocky interior {\coco (confidence regions of \mice and \rsolid in case B are three times smaller compared to case E)}, and some improvement on the composition of the gas layer, whereas an increase in precision of stellar refractory abundances enables improved constraints on mantle composition and relative core size. \item We have proposed two different atmospheric models: model I solves for radiative transfer; whereas model II uses a simplified scale-height pressure model. Both models yield different insights about possible gas layer characteristics that are subject to prescribed assumptions. In particular, for thick atmospheres, we see a clear discrepancy between model I and II which result in different estimates of rock and ice layers. The validity of model II is strictly limited to thin atmospheres (\menv $\lesssim 10^{-3}$~\ME). \item {\coco We have investigated the effect of prior distribution on estimated parameters and observed that the assumed prior distribution significantly affects the posterior distribution of those parameters, that are weakly constrained.} \end{itemize} In a companion paper \citep{dornA}, we present the application of our method to six observed exoplanets, for which mass, radius, and stellar abundance constraints are available. The method presented here is valuable for the interpretation of future data from space missions (TESS, CHEOPS, and PLATO) that aim at characterizing exoplanets through precise measurements of $R$ and $M$. Improving measurement precision, however, is costly as it depends on observation time. Our method helps to quantify the scientific return that could be gained as data precision is increased. Moreover, our study is relevant for the understanding on how interior types are distributed among stars and the implications of these for planet formation. | 16 | 9 | 1609.03908 |
1609 | 1609.05659_arXiv.txt | In conjunction with \textit{ab initio} potential energy and dipole moment surfaces for the electronic ground state, we have made a theoretical study of the radiative lifetimes for the hydronium ion H$_3$O$^{+}$ and its deuterated isotopologues. We compute the ro-vibrational energy levels and their associated wavefunctions together with Einstein coefficients for the electric dipole transitions. A detailed analysis of the stability of the ro-vibrational states have been carried out and the longest-living states of the hydronium ions have been identified. We report estimated radiative lifetimes and cooling functions for temperatures $<$~200~K. A number of long-living meta-stable states are identified, capable of population trapping. | In the Universe, molecules are found in a wide variety of environments: From diffuse interstellar clouds at very low temperatures to the atmospheres of planets, brown dwarfs and cool stars which are significantly hotter. In order to describe the evolution of the diverse, complex environments, it is essential to have realistic predictions of the radiative and cooling properties of the constituent molecules. Such predictions, in turn, require reasonable models for the energetics of each molecular species present. Although interstellar molecular clouds are usually characterised as cold, they are mostly not fully thermalized. Whether a species attains thermal equilibrium with the environment depends on the radiative lifetimes of its states and the rate of collisional excitations to the states: This is normally characterised by the critical density. In non-thermalized regions, radiative lifetimes are also important for modelling the maser activity observed for many species. The long lifetimes associated with certain excited states can lead to population trapping and non-thermal, inverted distributions. Such unexpected state distributions have been observed for the H$^+_3$ molecule both in space\cite{02GoMcGe.H3+, 05OkGeGo.H3+} and in the laboratory.\cite{02KrKrLa.H3+,04KrScTe.H3+} Dissociative recombination of hydronium H$_3$O$^{+}$ has been extensively studied in ion storage rings.% \cite{96AnHeKe.H3O+,00NeKhRo.H3O+,00JeBiSa.H3O+,10BuStMe.H3O+,10NoBuSt.H3O+} The lifetimes calculated in the present work suggest that H$_3$O$^{+}$ and it isotopologues will exhibit population trapping in a manner similar to that observed for H$_3^+$ in storage rings. Dissociative recombination of hydronium has been postulated as a possible cause of emissions from super-excited water in cometary comae\cite{09BaMiDe} and as the mechanism for a spontaneous infrared water laser.\cite{01SaKeWa.H3O+} Hydronium and its isotopologues play an important role in planetary and interstellar chemistry.\cite{00JeBiSa.H3O+,01GoCexx.H3O+} These molecular ions are found to exist abundantly in both diffuse and dense molecular clouds as well as in comae. Moreover, H$_3$O$^{+}$ is a water indicator and can be used to estimate water abundances when the direct detection is unfeasible.\cite{92PhVaKe.H3O+} Consequently, the ions have been the subject of numerous theoretical and experimental studies (see, for example, Refs.~% \onlinecite{00JeBiSa.H3O+,01GoCexx.H3O+, 92PhVaKe.H3O+, 73LiDyxx.H3O+, 80FeHaxx.H3O+, 82SpBuxx.H3O+, 83BeGuPf.H3O+, 83BoRoRe.H3O+, 85BoDeDe.H3O+, 84DaHaJo.H3O+, 84BuAmSp.H3O+, 85BeSaxx.H3O+, 85LiOkxx.H3O+, 85PlHeEr.H3Op, 86DaJoHa.H3O+, 86LiOkSe.H3O+, 87GrPoSa.H3O+, 88HaLiOk.H3O+, 88VeTeMe.H3O+, 90OkYeMy.H3O+, 90PeNeOw.H3O+, 91HoPuOk.H3O+, 97UyWhOk.H3O+, 99ArOzSa.H3O+, 00ChJuGe.H3O+, 01AiOhIn.H3O+, 02ErSoDo.H3O+, 02CaWaZu.H3O+, 03RaMiHa.NH3, 03HuCaBo.H3O+, 05YuBuJe.H3O+, 06DoNexx.H3O+, 08FuFaxx.H3O+, 09YuDrPe.H3O+, 10MuDoNe.H3O+, 12PeWeBe.H3O+,10BuStMe.H3O+} and references therein) mainly devoted to the spectroscopy and chemistry of the species. Whereas the cooling function of the H$_3^{+}$ ion has been extensively studied by Miller et al,\cite{96NeMiTe.H3+,10MiStMe.H3+,13MiStTe.H3+} no information about the radiative and cooling properties of H$_3$O$^{+}$ and its deuterated isotopologues has been available thus far. In the present work, we remedy this situation by determining theoretically the ro-vibrational states of the ions H$_3$O$^{+}$, H$_2$DO$^{+}$, HD$_2$O$^{+}$, and D$_3$O$^{+}$. We use \ai\ potential energy (PES) and dipole moment surfaces (DMS) for the ground electronic states of H$_3$O$^{+}$ from Ref.~\onlinecite{15OwYuPo.H3O+} to compute for each of the four ions considered here, ro-vibrational energy levels, the accompanying wavefunctions, and Einstein coefficients for the relevant ro-vibrational (electric dipole) transitions by means of the nuclear-motion program TROVE.\cite{TROVE} Lifetimes of individual ro-vibrational states are calculated and analyzed together with the overall cooling rates. Recently, the same methodology was used to estimate the sensitivities of hydronium-ion transition frequencies to a possible time variation of the proton-to-electron mass ratio.\cite{15OwYuPo.H3O+} We present a detailed analysis of the stability of the ro-vibrational states of the hydronium ions and identify the states with the longest lifetimes. This study is based on the methodology\cite{16TeHuNa.method} developed very recently as part of the ExoMol project. \cite{12TeYuxx.db} The ExoMol project aims at a comprehensive description of spectroscopic properties of molecules important for atmospheres of exoplanets and cool stars. The molecular lifetimes and cooling functions determined for H$_3$O$^{+}$ and its deuterated isotopologues in the present work are available in the new ExoMol data format.\cite{16TeYuAl.db} | We have carried out a theoretical study of the ro-vibrational states of the hydronium ion \ohhh\ and its deuterated isotopologues. \textit{Ab initio} potential energy and electric dipole moment surfaces were used to calculate ro-vibrational energy levels, corresponding wavefunctions and Einstein coefficients for the low-lying ro-vibrational transitions of these ions. We have analyzed stability of the ro-vibrational states and computed the radiative lifetimes and cooling functions for temperatures below 200~K. Taking into account only spontaneous emission as cause of decay of ro-vibrational states (and neglecting collisions and stimulated emission) we find the longest-lived hydronium state for \oddd: the population in the rotational state with $(J,K,\Gamma)$ $=$ $(5,5,E'')$ is trapped for 3816.0 years, which is relatively `hot' (152~K), at least in the context of molecular cooling, for example in storage rings. In this work we have identified a number of relatively hot ($E/k > 100$~K) meta-stable states with a lifetime longer than 10~s (typical timescales of ion storage experiments). Such meta-stable states which will be populated and hamper the cooling of hydronium ions to a temperature of a few Kelvin. The molecule with the shortest-lived meta-stable states is \ohdd\ with lifetimes of a few days. The timescale of interstellar collisions in diffuse clouds is longer (about a month), and thus some of these states undergo spontaneous emission. Our calculations show that deuteration influences significantly the hydronium lifetimes. This effect is mostly caused by the symmetry lowering from \Dh{3}(M) to \Cv{2}(M) and the ensuing perpendicular dipole moment component. A number of long-living meta-stable states are identified, capable of population trapping. Compared to the deuterated species, the cooling of the lightest isotopologue \ohhh\ is most efficient at higher temperatures ($T>30$~K). However, this changes at very low temperatures where the \ohhh\ ions are trapped at relatively high energy. The results obtained can be used to assess the cooling properties of the hydronium ion in ion storage rings and elsewhere. | 16 | 9 | 1609.05659 |
1609 | 1609.02881_arXiv.txt | The VERITAS imaging atmospheric Cherenkov telescope array has been observing the northern TeV sky with four telescopes since summer 2007. Over 50 gamma-ray sources have been studied, including active and starburst galaxies, pulsars and their nebulae, supernova remnants and Galactic binary systems. We review here some of the most recent VERITAS results, and discuss the status and prospects for collaborative work with other gamma-ray instruments, and with multimessenger observatories. | VERITAS (Figure~\ref{VERITAS}) is an imaging atmospheric Cherenkov telescope array, now entering its tenth year of operations. The array consists of four identical telescopes, located at the Fred Lawrence Whipple Observatory (FLWO) in Arizona, each with a $12\U{m}$ diameter tessellated reflector. Cherenkov light from gamma-ray and cosmic-ray initiated particle cascades is focused by the telescope reflectors onto 499-pixel photomultiplier tube (PMT) cameras, which cover a $3.5^{\circ}$ field of view. \begin{figure}[h] \centerline{\includegraphics[width=1.0\textwidth]{VERITAS.png}} \caption{The VERITAS array in its current (2016) configuration.} \label{VERITAS} \end{figure} The array has undergone two major upgrades over the past decade. The first, in 2009, involved the relocation of the original prototype telescope to a more favorable location, resulting in an approximate diamond-shaped layout with sides of $\sim 100\U{m}$. This improved the angular reconstruction capabilities, and enhanced the sensitivity of the array. In 2012, a major overhaul of the telescope electronics saw the installation of new trigger systems, and the replacement of all of the photosensors with higher quantum efficiency (\textit{super bialkali} photocathode) PMTs. This again improved sensitivity, and led to a dramatic enhancement of the low energy response. Thanks to these upgrades, and to improvements in analysis and calibration tools, VERITAS now detects a source with 1\% of the steady Crab Nebula flux in under 25\U{hours}; less than half of the exposure required in the original array configuration. Of similar importance is the increase in duty cycle provided by moonlight observations. VERITAS now commonly conducts observations with the lunar disk up to 50\% illuminated, including with reduced PMT voltage for higher illuminations. Moonlight observations account for as much as $\sim40\%$ of the total annual observing yield \cite{2015ApJ...808..110A}, which averages approximately 1300 hours per year. Figure~\ref{sensitivity} shows the differential sensitivity of VERITAS for three different epochs. Further details are available in \cite{2015arXiv150807070P}. \begin{figure}[h] \centerline{\includegraphics[width=0.6\textwidth]{sensitivity-crop.pdf}} \caption{VERITAS sensistivity for the three different configurations of the instrument.} \label{sensitivity} \end{figure} The VERITAS source catalog (Figure~\ref{catalog}) now stands at 56 sources, and is composed of eight different source classes. To make the most efficient use of limited observing time, the balance of the observing plan has shifted from attempts to detect new sources, to precision measurements requiring deep exposures, or observations of variable sources during exceptional high flux states. We summarize here some of the most interesting results from the past two years of VERITAS operations. \begin{figure}[h] \centerline{\includegraphics[width=1.0\textwidth]{catalog.png}} \caption{The VERITAS source catalog, in Galactic coordinates, as of July 2016. Shaded regions indicate visibility to VERITAS above $55^{\circ}$ elevation. Figure modified from TeVCat ({http://tevcat.uchicago.edu}). } \label{catalog} \end{figure} | VERITAS continues to operate smoothly, and the instrument is currently in its most sensitive configuration to-date. Ongoing analysis developments promise further incremental improvements in the coming years, while new scientific opportunities are presented by the introduction of HAWC, IceCube, and other facilities. Operational funding for VERITAS is secure, and the collaboration plans to continue observing with the array until at least 2019, with a 10-year anniversary workshop planned for 2017. The construction of the prototype mid-size Schwarzchild-Couder Telescope \cite{pSCThere} at the FLWO provides a firm link to CTA, and to the future of ground-based gamma-ray astronomy. | 16 | 9 | 1609.02881 |
1609 | 1609.02553_arXiv.txt | We present the measurement of the projected and redshift space 2-point correlation function (2pcf) of the new catalog of \textit{Chandra} COSMOS-Legacy AGN at 2.9$\leq$z$\leq$5.5 ($\langle L_{bol} \rangle \sim$10$^{46}$ erg/s) using the generalized clustering estimator based on phot-z probability distribution functions (Pdfs) in addition to any available spec-z. We model the projected 2pcf estimated using $\pi_{max}$ = 200 h$^{-1}$ Mpc with the 2-halo term and we derive a bias at z$\sim$3.4 equal to b = 6.6$^{+0.60}_{-0.55}$, which corresponds to a typical mass of the hosting halos of log M$_h$ = 12.83$^{+0.12}_{-0.11}$ h$^{-1}$ M$_{\odot}$. A similar bias is derived using the redshift-space 2pcf, modelled including the typical phot-z error $\sigma_z$ = 0.052 of our sample at z$\geq$2.9. Once we integrate the projected 2pcf up to $\pi_{max}$ = 200 h$^{-1}$ Mpc, the bias of XMM and \textit{Chandra} COSMOS at z=2.8 used in Allevato et al. (2014) is consistent with our results at higher redshift. The results suggest only a slight increase of the bias factor of COSMOS AGN at z$\gtrsim$3 with the typical hosting halo mass of moderate luminosity AGN almost constant with redshift and equal to logM$_h$ = 12.92$^{+0.13}_{-0.18}$ at z=2.8 and log M$_h$ = 12.83$^{+0.12}_{-0.11}$ at z$\sim$3.4, respectively. The observed redshift evolution of the bias of COSMOS AGN implies that moderate luminosity AGN still inhabit group-sized halos at z$\gtrsim$3, but slightly less massive than observed in different independent studies using X-ray AGN at z$\leq2$. | \label{sec:intro} \begin{figure*} \plottwo{zmeanpdf1.eps}{lid_766_3.eps} \caption{\footnotesize Left Panel: Mean normalized phot-z Pdf of all CCL AGN with best-fit phot-z $>$ 2.9. Right Panel: Normalized phot-z Pdf for the source lid766. This source has a best-fit photo-z value $<$ 2.9, but a phot-z Pdf(z$_i >$2.9)$>$0.001 (red thick line). The redshifts above this threshold, weighted by their Pdf, have been taken in account in the catalog used to estimate the 2pcf.} \label{fig2} \end{figure*} The presence of a nuclear supermassive black hole (BH) in almost all galaxies in the present day Universe is an accepted paradigm in astronomy (e.g. Kormendy \& Richstone 1995; Kormendy \& Bender 2011). Despite major observational and theoretical efforts over the last two decades, a clear explanation for the origin and evolution of BHs and their actual role in galaxy evolution remains elusive. Diverse scenarios have been proposed. One possible picture includes major galaxy merger as the main triggering mechanism (e.g. Hopkins et al. 2006; Volonteri et al. 2003, Menci et al. 2003,2004). On the other hand, there is mounting observational evidence suggesting that moderate levels of AGN activity might not be always causally connected to galaxy interactions (Lutz et al. 2010, Mullaney et al. 2012, Rosario et al. 2013, Villforth et al. 2014). Several works on the morphology of the AGN host galaxies suggest that, even at moderate luminosities, a large fraction of AGN is not associated with morphologically disturbed galaxies. This trend has been observed both at low (z $\sim$ 1, e.g., Georgakakis et al. 2009; Cisternas et al. 2011) and high (z $\sim$ 2, e.g., Schawinski et al. 2011, 2012; Kocevski et al. 2012, Treister et al. 2012) redshift. Theoretically, in-situ processes, such as disk instabilities or stochastic accretion of gas clouds, have also been invoked as triggers of AGN activity (e.g. Genzel et al. 2008, Dekel et al. 2009, Bournaud et al. 2011). AGN clustering analysis provides a unique way to unravel the knots of this complex situation, providing important, independent constraints on the BH/galaxy formation and co-evolution. In the cold dark matter-dominated Universe galaxies and their BHs are believed to populate the collapsed dark matter halos, thus reflecting the spatial distribution of dark matter in the Universe. The most common statistical estimator for large-scale clustering is the two-point correlation function (2pcf, Davis \& Peebles 1983). This quantity measures the excess probability above random to find pairs of galaxies/AGN separated by a given scale $r$. By matching the observed 2pcf to detailed outputs of dark matter numerical simulations, one can infer the typical mass of the hosting dark matter halos. This is derived through the so called AGN bias b, enabling then to pin down the typical environment where AGN live. This in turn can provide new insights into the physical mechanisms responsible for triggering AGN activity. The 2pcf of AGN has been measured in optical large area surveys, such as the 2dF (2QZ, Croom et al. 2005; Porciani \& Norberg 2006) and the Sloan Digital Sky Survey (SDSS, Li et al. 2006; Shen et al. 2009; Ross et al. 2009). These optical surveys are thousands of square degree fields, mainly sampling rare and high luminosity quasars. The amplitude of the 2PCF of quasars suggests that these luminous AGN are hosted by halos of roughly constant mass, a few times 10$^{12}$ M$_{\odot}$, out to z =3-4 (Shanks et al. 2011). Models of major mergers between gas-rich galaxies appear to naturally reproduce the clustering properties of optically selected quasars as a function of luminosity and redshift (Hopkins et al. 2007a, 2008; Shen 2009; Shankar et al. 2010; Bonoli et al. 2009). This supports the scenario in which major mergers dominate the luminous quasar population (Scannapieco et al. 2004; Shankar et al. 2010; Neistein \& Netzer 2014; Treister et al. 2012). \textit{Chandra} surveys have contributed significantly to the study of the AGN clustering (e.g. CDFS-N, Gilli et al. 2009; Chandra/Bootes, Starikova et al. 2011, Allevato et al. 2014). Deep X-ray data can be used to draw conclusions on the faint portions of the AGN luminosity function, where a significant fraction of obscured sources is present. In particular, the \textit{Chandra} survey in the two square degree COSMOS field (C-COSMOS, Elvis et al. 2009, Civano et al. 2012; \textit{Chandra} COSMOS Legacy Survey, Civano et al. 2016) has allowed the investigation of the redshift evolution of the clustering properties of X-ray AGN, for the first time up to z$\sim$3. Interestingly, over a broad redshift range (z $\sim$ 0 - 2) moderate luminosity AGN occupy DM halo masses of log M$_h$ $\sim$ 12.5-13.5 M$_{\odot}$ h$^{-1}$. The clustering strength of X-ray selected AGN has been measured by independent studies to be higher than that of optical quasars. Merger models usually fail in reproducing the data from X-ray surveys, opening the possibility of additional AGN triggering mechanisms (e.g., Allevato et al. 2011; Mountrichas \& Georgakakis 2012) and/or multiple modes of BH accretion (e.g., Fanidakis et al. 2013). Recently, Mendez et al. (2015) and Gatti et al. (2016) have suggested that selection cuts in terms of AGN luminosity, host galaxy properties and redshift interval, might have a more relevant role in driving the differences often observed in the bias factor inferred from different surveys. \begin{figure*} \plottwo{zsmoodistr.eps}{fig1.eps} \caption{\footnotesize Redshift (Left Panel) and 2-10 keV X-ray luminosity (Right Panel) distribution for 107 CCL AGN with known spec-z (red dashed line), 212 AGN with known spec or best phot-z (solid grey line) and 221.6 AGN with known spec-z or phot-z weighted by the Pdf (black dotted line), at 2.9$\leq$z$\leq$5.5} \label{fig1} \end{figure*} The measurement of the AGN bias is crucial at high redshifts, especially at z$>$2-3, i.e. at the peak in the accretion history of the Universe. At z$>$3, Shen et al. (2007,2009) measured for the first time the 2pcf of luminous SDSS-DR5 quasars (log L$_{bol} \sim 10^{47}$ erg/s) at $\langle z \rangle$ = 3.2 and 3.8. Even if with very large uncertainty, they found that these objects live in massive halos of the order of 10$^{13}$ M$_{\odot}$ h$^{-1}$. This result is consistent with models invoking galaxy major mergers as the main triggering mechanism for very luminous AGN. Recently, Eftekharzadeh et al. (2015) studying a sample of spectroscopically confirmed SDSS-III/BOSS quasars at 2.2$\leq$z$\leq$3.4, performed a more precise estimation of the quasar bias at high redshift. They found no evolution of the bias in three redshift bins, with halo masses equal to 3$\times$ and $\sim$0.6$\times$ 10$^{12}$ M$_{\odot}$ h$^{-1}$ at z$\sim$2.3 and $\sim 3$, respectively. There are only a few attempts of measuring the clustering properties of X-ray AGN at z $\sim$ 3. Francke et al. (2008) estimated the bias of a small sample of X-ray AGNs (L$_{bol}$ $\sim$ 10$^{44.8}$ erg s$^{-1}$) in the Extended Chandra Deep Field South (ECDFS), with very large uncertainty. They found indications that X-ray ECDFS AGNs reside in dark matter halos with minimum mass of log M$_{min}$ = 12.6$^{+0.5}_{-0.8}$ h$^{-1}$ M$_{\odot}$. On the other hand, Allevato et al. (2014) used a sample of \textit{Chandra} and XMM-Newton AGN in COSMOS with moderate luminosity (log L$_{bol} \sim 10^{45.3}$ erg/s) at $\langle z \rangle$=2.86. For the first time they estimated the bias of X-ray selected AGN at high redshift, suggesting that they inhabit halos of logM$_h$ = 12.37$\pm$0.10 M$_{\odot}$ h$^{-1}$. They also extended to z$\sim$3 the result that Type 1 AGN reside in more massive halos than Type 2 AGN. Recently, Ikeda et al. (2015) estimated the clustering properties of low-luminosity quasars in COSMOS at 3.1$\leq$z$\leq$4.5, using the cross-correlation between Lyman-Break Galaxies (LBGs) and 25 quasars with spectroscopic and photometric redshifts. They derived a 86\% upper limit of 5.63 for the bias at z$\sim$ 4. In this paper we want to extend the study of the clustering properties of X-ray selected AGN to z$>$3 using the new \textit{Chandra} COSMOS-Legacy data. To this goal, we perform clustering measurements using techniques based on photometric redshift in the form of probability distribution functions (Pdfs), in addition to any available spectroscopy. This is motivated by the development in the last years of clustering measurement techniques based on photometric redshift Pdfs by Myers, White \& Ball (2009), Hickox et al. (2011, 2012) and Mountrichas et al. (2013) and Georgakakis et al. (2014). One of the advantages of this new clustering estimator is that one can use in the analysis all sources not just the optically brighter ones for which spectroscopy is available. For this reason it is well suited to clustering investigations using future large X-ray AGN surveys, where the fraction of spectroscopic redshifts might be small. Throughout the paper, all distances are measured in comoving coordinates and are given in units of Mpc $h^{-1}$, where $h=H_0/100$km/s. We use a $\Lambda$CDM cosmology with $\Omega_M=0.3$, $\Omega_\Lambda=0.7$, $\Omega_b=0.045$, $\sigma_8=0.8$. The symbol $log$ signifies a base-10 logarithm. \begin{deluxetable*}{lllllllll} \tabletypesize{\scriptsize} \tablewidth{0pt} \tablecaption{Properties of the AGN Samples \label{tbl-1}} \tablehead{ \colhead{Sample} & \colhead{ N } & \colhead{ $\Sigma$Pdf$_j$(z$\geq$2.9)} & \colhead{$\langle z \rangle$} & \colhead{$log \langle L_{bol} \rangle$} & \colhead{$b$} & \colhead{logM$_h$} & \colhead{$b$} & \colhead{logM$_h$}\\ \colhead{} & \colhead{} & \colhead{} & \colhead{} & \colhead{erg s$^{-1}$} & \colhead{Eq. 7} & \colhead{h$^{-1}$M$_{\odot}$} & \colhead{Eq. 14} & \colhead{h$^{-1}$M$_{\odot}$} } \startdata Spec-zs + Phot-z Pdfs & 457 & 221.6 & 3.36\tablenotemark{a} & 45.99$\pm$0.53 & 6.6$^{+0.6}_{-0.55}$ & 12.83$^{+0.12}_{-0.11}$ & 6.53$^{+0.52}_{-0.55}$ & 12.82$^{+0.11}_{-0.13}$\\ Spec-zs + Best-fit Phot-zs & 212 & 212 & 3.34 & 45.93$\pm$0.17 & 6.48$^{+1.27}_{-1.36}$ & 12.81$^{+0.24}_{-0.35}$ & 6.96$^{+0.72}_{-0.73}$ & 12.90$^{+0.15}_{-0.15}$\\ Spec-zs only & 107 & 107 & 3.35 & 45.92$\pm$0.34 & 7.5$^{+1.6}_{-1.7}$ & 13.0$^{+0.25}_{-0.35}$ & 7.98$^{+1.4}_{-1.5}$ & 13.08$^{+0.22}_{-0.25}$\\ \enddata \tablenotetext{a}{Mean redshift of the sample weighted by the Pdfs.} \end{deluxetable*} | \label{sec:conc} We use the new CCL catalog to probe the projected and redshift-space 2pcf of X-ray selected AGN for the first time at 2.9$\leq$z$\leq$5.5, using the generalized clustering estimator based on phot-z Pdfs in addition to any available spec-z. We model the clustering signal with the 2-halo model and we derive the bias factor and the typical mass of the hosting halos. Our key results are: \begin{enumerate} \item At z$\sim$3.4, CCL AGN have a bias b = 6.6$^{+0.60}_{-0.55}$, which corresponds to a typical mass of the hosting halos of log M$_h$ = 12.83$^{+0.12}_{-0.11}$ h$^{-1}$ M$_{\odot}$. A similar bias is derived using the z-space 2pcf, modelled including the typical phot-z error $\sigma_z$ = 0.052 of our sample. This confirms that the convergence of the projected 2pcf observed only at large scales ($\pi_{max}\geq200$ h$^{-1}$ Mpc) is due to large phot-z errors. \item A slightly larger bias b = 7.5$^{+1.6}_{-1.7}$ (but consistent within the error bars) is found using a sample of 107 CCL AGN with known spec-z. The modelling of $\xi(s)$ suggests that this larger bias can be explained assuming that spec-zs are affected by errors of the order of $\sigma_z = 0.02-0.025$. This would explain the convergence of the projected 2pcf surprisingly observed only at $\pi_{max}\geq200$ h$^{-1}$ Mpc, even when phot-zs are not included in the analysis. However, given the low statistics smaller spec-z errors and then bias can not be ruled-out. \item We estimate the bias factor for the sample of 346 XMM and \textit{Chandra} AGN used in Allevato et al. (2014) using $\pi_{max}=200$ h$^{-1}$ Mpc in estimating the projected 2pcf and then accounting for the large phot-z errors. In particular we found b = 5.8$^{+0.61}_{-0.55}$, which is significantly larger than the AGN bias measured in Allevato et al. (2014) and corresponds to logM$_h$ = 12.92$^{+0.13}_{-0.18}$ at z=2.8. \item Our results suggest only a slight increase of the bias factor of COSMOS AGN at z$\gtrsim$3, with the typical hosting halo mass of moderate luminosity AGN almost constant with redshift and equal to logM$_h$ = 12.92$^{+0.13}_{-0.18}$ at z=2.8 and log M$_h$ = 12.83$^{+0.12}_{-0.11}$ at z$\sim$3.4, respectively. \item The observed redshift evolution of the bias of COSMOS AGN implies that moderate luminosity AGN still inhabit group-sized halos, but slightly less massive than observed in different independent studies using X-ray AGN at z$\leq2$. \item Theoretical models presented in Shen (2009) and Hopkins et al. (2007b) that assume an AGN activity mainly triggered by major mergers of host halos underpredict our results at z$\sim$3.4 for CCL AGN with mean L$_{bol} \sim 10^{46}$ erg s$^{-1}$. A similar tension is also observed when comparing to the semi-empirical models presented in Conroy \& White (2013). In the latter model, this disagreement can be explained if AGN have a duty cycle approaching unity at z$>$3. On the other hand, following the semi-analytic models presented in Gatti et al. (2016), in both galaxy interaction and disk instability models only galaxies with stellar masses above 10$^{11}$ M$_{\odot}$ would be able to host AGN with luminosity of L$_{bol}\sim$10$^{46}$ erg/s and highly biased such as COSMOS AGN at z$>$2-3. \end{enumerate} Only future facilities, like the X-ray Surveyor (Vikhlinin 2015) and Athena (PI K. P. Nandra), will be able to collect sizable samples ($\sim$1000s) of low luminosity (L$_X<$10$^{43}$ erg/s) AGN at z$>$3 (Civano 2015), allowing to explore the clustering for significantly less luminous source and to test AGN triggering scenarios at different AGN luminosities. | 16 | 9 | 1609.02553 |
1609 | 1609.03469_arXiv.txt | { Mass-loss in massive stars plays a critical role in their evolution, although the precise mechanism(s) responsible - radiatively driven winds, impulsive ejection and/or binary interaction - remain uncertain. In this paper we present ALMA line and continuum observations of the supergiant B[e] star Wd1-9, a massive post-Main Sequence object located within the starburst cluster Westerlund 1. We find it to be one of the brightest stellar point sources in the sky at millimetre wavelengths, with (serendipitously identified) emission in the H41$\alpha$ radio recombination line. We attribute these properties to a low velocity ($\sim100$\,\kms) ionised wind, with an extreme mass-loss rate $\gtrsim$6.4$\times10^{-5}(d/5kpc)^{1.5}$\,\solmasyr. External to this is an extended aspherical ejection nebula indicative of a prior phase of significant mass-loss. Taken together, the millimetre properties of Wd1-9 show a remarkable similarity to those of the highly luminous stellar source MWC349A.% We conclude that these objects are interacting binaries evolving away from the main sequence and undergoing rapid case-A mass transfer. As such they - and by extension the wider class of supergiant B[e] stars - may provide a unique window into the physics of a process that shapes the life-cycle of $\sim70$\% of massive stars found in binary systems.} | The evolution of massive OB stars is dominated by mass-loss throughout their life-cycle, with the nature of their death - core-collapse supernova versus prompt collapse - and ultimate fate (neutron star or black hole) dependent on how much matter they are able to shed. Historically, line-driven winds have been thought to mediate this mass-loss. However, there is a significant uncertainty in our understanding of this phenomenon; depending on whether such winds are smooth or clumped, the resultant mass-loss rates of massive stars are in question at the order-of-magnitude level \citep[e.g.][]{fullerton06,puls06}. Given this discord, other physical mechanisms have been invoked to explain how the outer mantle of O stars may be stripped to yield H-depleted Wolf-Rayet (WR) stars, with impulsive mass-loss driven by instabilities in the luminous blue variable (LBV) or yellow hypergiant/red supergiant phase a prime candidate \citep[e.g.][]{humphrey79,abr14}. However, recent analysis has shown that fully 70\% of massive stars are in binaries that will interact at some point in their life-cycle \citep{demink14}, also resulting in significant mass-loss from both the primary to the secondary as well as from the binary system as a whole. In order to place observational constraints on these disparate processes we obtained millimetre continuum observations of the massive Galactic starburst cluster Westerlund 1 \citep[Wd1;][]{clark05} with the Atacama Large Millimetre/Submillimeter Array (ALMA). Observations at such wavelengths are particularly sensitive to wind-clumping via thermal Bremsstrahlung emission \citep[e.g.][]{blomme03}, while ejection nebula formed via impulsive mass-loss events or binary interaction may also be identified and resolved at millimetre-radio wavelengths. With an age of $\sim5$\,Myr, Wd1 is the ideal laboratory to undertake such a study since it contains a uniquely rich population of massive, evolved single and binary stars at a distance that enables the detection of several tens of objects (Fenech et al. in prep.). In this study we focus on the millimetre properties of the supergiant (sg)B[e] star Wd1-9 (R.A. 16$^{{h}}$47$^{m}$4.162$^{s}$ Dec. -45$^{o}$50$'$31.373$''$ ). SgB[e] stars are characterised by an aspherical circumstellar environment, comprising a high velocity polar wind and a warm, dusty and gaseous equatorial disc/torus \citep{zickgraf85}, leading to a pronounced near-IR continuum excess and rich emission line spectrum. The origin of the disc is uncertain, but recently interest has grown in binarity as a physical driver of the phenomenon \citep{pods06,clark13a}; if this is the case then the circumstellar environment may encode information on the mass-loss history of the binary and critically the physics driving it. Wd1-9 is the only sgB[e] star known to reside in a star cluster and hence one of the handful of Galactic objects for which a post-MS (main sequence) evolutionary state may be unambiguously determined \citep{clark13b}. Interest in it was reignited by the discovery that it is one of the most luminous stellar radio sources known \citep{clark98,dough10}. A synthesis of the multi-wavelength and epoch dataset compiled as a result of this finding was presented in \cite{clark13b}, which revealed an emission spectrum arising from the circumstellar environment that apparently entirely veils the nature of the central source, although the IR excess clearly indicates the presence of a hot dusty torus. Additionally, the presence of emission lines from high-excitation species (e.g. [S\,{\sc iv}], [O\,{\sc iv}]) and an unexpected X-ray luminosity \citep{clark08} provides persuasive evidence for binarity. | We have presented new ALMA 3mm observations of Wd1-9 that reveal it to be exceptionally luminous, suggesting an extreme mass-loss rate in comparison to other cluster members. We identify both compact and extended thermal components and unexpected line emission in the H41$\alpha$ transition associated with the former. Analysis of the line centroids implies a velocity gradient of $\sim200$\,\kms perpendicular to the bright N-S component of the extended nebula. \cite{clark13b} have previously highlighted the close similarity between Wd1-9 and the B[e] star MWC349A at optical and IR wavelengths; these observations extend such a comparison to the mm and radio regime. Both systems are thought to host compact circumstellar torii which, in the case of MWC349A drives a massive, slow wind perpendicular to its surface. By analogy to MWC349A we suggest that the H41$\alpha$ line emission associated with Wd1-9 likewise arises in such an outflow. {\em If} this hypothesis is correct, the terminal velocity and mass-loss rates of the winds in both systems are comparable to within a factor of $\sim4$. Moreover, if Wd1-9 and MWC349A host similar circumstellar environments we would expect that future optimised observations will reveal higher transitions (e.g. H29-31$\alpha$) that are double peaked \citep[cf.][]{weintroub08,baez13} and delineate the dusty torus (running N-S) perpendicular to our putative outflow. Assuming Wd1-9 and MWC349A are similar systems we may draw several far reaching conclusions. Firstly MWC349A would be unambiguously a post-MS object, resolving a decades-old debate \citep[cf.][ and refs. therein]{gvar12}. Secondly both systems would form prime laboratories for the study of disc-wind launching, of considerable interest to the study of massive protostars as well as post-MS evolution. Similarly, both would provide a window into the poorly understood physics of massive interacting binaries. The mass-loss rates inferred for both Wd1-9 and MWC349A are sufficiently large that they cannot be attributed to binary interaction on the nuclear ($\sim10^6$yr) timescale since this would yield a total mass lost to the systems that was directly comparable to the initial masses inferred for the primaries. Instead we suppose that we are observing rapid case A mass transfer in both systems occuring on the much shorter thermal timescale (few $10^4$yr). In such a scenario a percentage of this mass is accreted by the binary, with the remainder lost to the binary system \citep{petrovic05}; both MWC349 and Wd1-9 suggest that the latter process is mediated via a low-terminal-velocity disc wind. Extending this conclusion, we note that while formally classifiable as sgB[e] stars following the criteria of e.g. \cite{zickgraf85}, both Wd1-9 and MWC349A appear to differ from other examples in the Magellanic Clouds. In both cases this is due to (i) the origin of the polar wind (disc-driven rather than a classical stellar wind) and (ii) the high temperature and compact nature of the dusty toroids. We speculate that these differences result from the stage of binary interaction in which we observe these systems, with extensive mass transfer and mass-loss occurring within Wd1-9 and MWC349A on a thermal timescale, while archetypal sgB[e] stars such as R126 \citep{kastner10} with colder more extended discs may represent post mass-transfer systems. Finally, as noted before, the high mass loss rates and low wind velocities that characterise both Wd1-9 and MWC349A are similar to the outflow properties of LBVs - both inferred directly and from analysis of ejection nebulae \citep{clark14}. Given the LBV-like variability observed in some sgB[e] binaries \citep{clark13a}, one might speculate that the extreme mass loss rates of some LBV-like systems are also driven by binarity (cf. RY Scuti). | 16 | 9 | 1609.03469 |
1609 | 1609.04280_arXiv.txt | {% Nowadays large spectroscopic surveys, like the Gaia-ESO Survey (GES), provide unique stellar databases for better investigating the formation and evolution of our Galaxy. Great attention must be paid to the accuracy of the basic stellar properties derived: large uncertainties in stellar parameters lead to large uncertainties in abundances, distances and ages. Asteroseismology has a key role in this context: when seismic information is combined with information derived from spectroscopic analysis, highly precise constraints on distances, masses, extinction and ages of Red Giants can be obtained. In the light of this promising joint-action, we started the CoRoT-GES collaboration. We present a set of 1,111 CoRoT stars, observed by GES from December 2011 to July 2014, these stars belong to the CoRoT field LRc01, pointing at the inner Galactic Disk. Among these stars, 534 have reliable global seismic parameters. By combining seismic informations and spectroscopy, we derived precise stellar parameters, ages, kinematic and orbital parameters and detailed element abundances for this sample of stars. We also show that, thanks to asteroseismology, we are able to obtain a higher precision than what can be achieved by the standard spectroscopic means. This sample of CoRoT Red Giants, spanning Galactocentric distances from 5 to 8 kpc and a wide age interval (1-13 Gyrs), provides us a representative sample for the inner disk population.} | Galactic Archeology, the study of how Milky Way formed and evolved, is nowadays entering in a golden era. The forthcoming Gaia mission data releases \citep{Perryman2001} and the large spectroscopic Galactic surveys, like RAVE \citep{Steinmetz2006}, GES \citep{Gilmore2012}, APOGEE \citep{Majewski2015}, GALAH \citep{Freeman2010} and the future 4MOST \citep{deJong2012}, are offering wide, unique and promising stellar datasets for testing the modern chemo-dynamical models. By comparing the main observables of the different components of our Galaxy with those predicted by models, like the age-metallicity relation, chemical gradients and kinematics, we will be able to understand the mechanisms that led to the actual Milky Way (i.e. \citet{Minchev2014}, and Minchev, Famay, Gerhard contributions to this conference). This comparison with models requires high precision and accuracy in distance, velocity, element abundances and ages. Typically, for chemo-dynamical investigations, accuracies in velocity better than 1 km/s, few \% in distance and lower than 0.1 dex for element abundances are needed, in addition to an information on age, with an error lower than 20\%. The accuracy in distance can be achieved thanks to Gaia, while the high precision on abundances can be addressed thanks to high resolution spectroscopy on high SNR spectra (e.g. 0.08 dex for GES and 0.05 dex for APOGEE). Most of the modern stellar spectroscopic surveys are targeting Red Giants, since they are the perfect tracers for Galactic investigations, thanks to their intrinsic brightness and incidence. However, for Red Giants stars, the atmospheric parameters determination from spectroscopy, especially for surface gravity, log(g), can be difficult. The surface gravity, in fact, can be affected by systematics up to 0.2 dex for this kind of stars (\citealp{Morel2012, Hekker2013, Heiter2015}). Since the abundances determination is coupled with atmospheric parameters, such systematics can lead to systematics of the same magnitude in the element abundances, compromising the quality of the data sample. Regarding the age determination, while it can be computed with a reasonable accuracy only for few targets (i.e. clusters) or dwarfs (with the standard method of isochrone fitting), it remains still precluded for field Red Giants: when using the commonly used isochrone fitting technique, age uncertainty can be up to 80\% \citep{Bergemann2014}, due to the degeneracies affecting the Red Giants locus. While waiting for Gaia data releases, asteroseismology can help in improving atmospheric parameters, abundances, age and distance for field Red Giants, required for Galactic Archaeology investigations. The CoRoT and Kepler space missions revolutionised the view on Red Giants, showing that it is possible to directly link the two main seismic observables, $\Delta\nu$ and $\nu_{\rm{max}}$, to the stellar mass and radius. Thank to these scaling relations, it is therefore possible to determine a very precise and accurate log(g), with an error of only 0.03 dex (\citealp{Morel2012, Thygesen2012}). Fixing the gravity to the very precise log(g) provided by asteroseismology, abundances with a precision of 0.05 dex can be measured, as showed, for few stars, in \citet{Morel2014} and \citet{Batalha2011}. Since the seismic scaling relations provides a very precise value of the star's mass and radius (typical errors of 10\% and 3\% respectively), it is also possible to derive the stellar age (since the age of a Red Giant star is directly linked to its mass), even though always using models, and distance. Asteroseismology have been already successfully applied for better investigating disk population in \citet{Miglio2013} and in identifying a new population of young alpha-enhanced stars, see \citet{Chiappini2015} (CoRoT data) and \citet{Martig2015} (Kepler data). Nowadays asteroseismology has been included in the main spectroscopic surveys as a calibration tool,as in GES \citep{Pancino2012}, APOGEE \citep{Pinsonneault2014}, and LAMOST \citep{Wang2016}, where a benchmark of Red Giants possessing very good seismic parameters have been used for better testing, and eventually calibrating, the measured log(g), or as training set for their pipelines. In this contribution we present how we analysed the spectra of the sample of CoRoT solar-like oscillating stars observed by the Gaia-ESO Survey (GES). These stars belong to the LRc01 field of CoRoT, pointing at the inner part of the Galactic disk. In section 2 we present how the sample of CoRoT Red Giants were selected and observed, in Section 3 we present how spectra have been analysed using asteroseismic information on gravity and how distances, ages, reddening and orbit parameters have been computed. Finally, in section 4, we present our conclusions. | \begin{table} \caption{Typical errors on atmospheric parameters, abundances, mass, radius, age and distance obtained with classic techniques (e.g. spectroscopy, isochrone fitting) and the errors on the same values obtained using seismic information.} \begin{tabular}{llcc} \hline \hline $\sigma$ & & Spectroscopy & Spectroscopy + \\ & & & Asteroseismology \\ \hline & & & \\ $T_{\rm eff}$ & GIR. & 100 & 65 \\ $[\rm{K}]$ & UVES & 70 & 55 \\ & & & \\ log(g) & GIR. & 0.20 & 0.03 \\ $[\rm{dex}]$ & UVES & 0.12 & 0.03 \\ & & & \\ $[\rm{Fe/H}]$ & GIR. & 0.10 & 0.08 \\ $[\rm{dex}]$ & UVES & 0.09 & 0.05 \\ & & & \\ $[\rm{elem./Fe}]$& GIR. & 0.20 & 0.08 \\ $[\rm{dex}]$ & UVES & 0.08 & 0.05 \\ & & & \\ \hline $\sigma$ & & & Asteroseismology \\ \hline & & & \\ Mass & & - & 7\% \\ Radius & &- & 3\% \\ Age & & $>$80\% & 21\% \\ Dist. & & - & 2\% \\ \hline \hline \end{tabular} \end{table} We derived, using a pipeline that implements seismic gravity in the analysis, refined atmospheric parameters and abundances of a sample of 534 CoRoT Red Giants observed by GES. In the analysis we fixed the log(g) to the seismic value, and we iteratively derived T$_{\rm eff}$~and overall metallicity [M/H], abundances of alpha-elements (O, Mg,Al, Si, Ca, Ti), n-capture elements (Y, Sr, Zr), Fe-peak elements (Sc, V, Cr, Mn, Ni, Fe), Na, Li and K. The typical errors on the atmospheric parameters and abundances, and the comparison with the errors obtained by using only spectroscopy, are reported on Tab.~3. The method was tested on the Gaia Benchmark stars and on a set of 77 CoRoT stars in common between GES and APOGEE, and resulting in reliable atmospheric parameters and abundances. We finally obtained a sample of 498 stars, possessing not only precise abundances (typical error on element abundances $<$ 0.10 dex), but also distances and ages with an error of 2\% and 21\% respectively (see Table~3), in a more precise way than what can be obtained with the classic methods. Our sample is distribute along a beam pencil, spanning 5 - 8 Kpc in Galactocentric distance (see Fig.~1), and covering a wide age interval, from $\sim$1 Gyr to 12 Gyr. The use of this sample for Galactic Archaeology purposes is discussed in Valentini et al. (2016, in prep). | 16 | 9 | 1609.04280 |
1609 | 1609.09655.txt | Stray light in X-ray telescopes is a well-known issue. Unlike rays focused via a double reflection by usual grazing-incidence geometries such as the Wolter-I, stray rays coming from off-axis sources are reflected only once by either the parabolic or the hyperbolic segment. Although not focused, stray light may represent a major source of background and ghost images especially when observing a field of faint sources in the vicinities of another, more intense, just outside the field of view of the telescope. The stray light problem is faced by mounting a pre-collimator in front of the mirror module, in order to shade a part of the reflective surfaces that may give rise to singly-reflected rays. Studying the expected stray light impact, and consequently designing a pre-collimator, is a typical ray-tracing problem, usually time and computation consuming, especially if we consider that rays propagate throughout a densely nested structure. This in turn requires one to pay attention to all the possible obstructions, increasing the complexity of the simulation. In contrast, approaching the problems of stray light calculation from an analytical viewpoint largely simplifies the problem, and may also ease the task of designing an effective pre-collimator. In this work we expose an analytical formalism that can be used to compute the stray light in a nested optical module in a fast and effective way, accounting for obstruction effects. | \label{sec:intro} Optical modules for X-ray telescopes are usually double reflection systems, like the widespread Wolter-I design\cite{VanSpey}. Reflecting X-rays twice halves the focal length and largely suppresses the coma aberration, enabling more compact spacecrafts and larger field of views, at the sole cost of some reduction of the effective area with respect to a single reflection. When a Wolter-I mirror module is illuminated by an on-axis source at infinite distance, all the rays that are reflected by the parabolic segment also impinge onto the hyperbolic segment. But things may change when the source is off-axis or at finite distance: some X-rays can make a single reflection on the parabola, while others can directly impinge on the hyperbola (Fig.~\ref{fig:straylight}). In both cases, they can reach the focal plane without being focused and increase the background or generate "ghost" images. These rays, usually referred to as "stray light", are a well-known problem in X-ray astronomy as they can seriously hamper the observation of faint targets by contamination from intense X-ray objects just {\it outside} the telescope field of view, like e.g., the Crab nebula or even the Sun if the detectors are not shielded against the visible or the infrared light\cite{Peterson1997}. When designing X-ray optical modules, the assessment of the stray light contamination is an important step in order to study possible countermeasures. Mirror modules consist of densely nested mirror shells, hence part of stray rays is blocked by the rear (usually non-reflective) side of the inner shells. However, the mirror nesting cannot be too tight, or also doubly-reflected rays from off-axis sources will be obstructed, at the expense of the effective area for imaged sources within the field of view. For this reason, other solutions have been devised out, like X-ray precollimators (see, e.g.\cite{Cusumano2007, Mori2012}) to prevent X-rays from reaching the optical surfaces from directions that may generate stray rays. These auxiliary items are carefully designed, manufactured and aligned to the mirror aperture\cite{eROSITAbaffle}, but at the same time to minimize the obstruction of the effective area for double reflection. Because of the complexity of the possible paths followed by X-rays (stray or not) throughout an optical module, the design and the performance verification of an X-ray optical module and of the pre-collimator is usually done via ray-tracing. This task can be computationally intense and time consuming, especially because the design performance has to checked until an optimal solution is found. In contrast, approaching the same problem from an analytical viewpoint would be useful. Not only to compute in an easy and fast way the stray light impact on a given X-ray module design, but also to find the optimal configuration for the mirror module and the precollimator without the need of writing complex ray-tracing routines. \begin{figure}[hbt] \centering \includegraphics[width = 0.85\textwidth]{straylight_fig1.eps} \caption{Origin of stray light in an unobstructed Wolter-I mirror shell. Off-axis rays reflected by the primary segment above the $V(\varphi)$ line (Eq.~\ref{eq:V_coef}) do not make the double reflection. Stray light also stems from a direct reflection below the $z$~=0 plane.} \label{fig:straylight} \end{figure} In previous papers\cite{Spiga2009, Spiga2010} we had already faced the problem of off-axis effective area computation for mirror shells, using analytical expressions. The results could be extended to the case of shells nested in mirror modules, accounting for the mutual obstruction\cite{Spiga2011}. The effective area, as a function of the off-axis angle and the X-ray wavelength, can be expressed by integral equations, with results in excellent agreement with the findings of ray-tracing. In addition, the numerical integration of the analytical formulae requires a time that is orders of magnitudes lower than ray-tracing, and without being affected by statistical uncertainties. In this paper we show that the same formalism can be easily extended to the computation of the {\it effective area for stray light} for an off-axis source. In Sect.~\ref{sec:recall} we briefly recall the analytical theory of the effective area\cite{Spiga2009, Spiga2011}. This will introduce us to some concepts that we can use to write analytical expressions of the effective area for doubly-reflected rays and, in Sect.~\ref{sec:straylight}, of the effective area for stray light. The expressions for the stray light off the primary and the secondary segment are different and, for brevity, we refer to the former as "primary stray light" and to the latter as "secondary stray light". In Sect.~\ref{sec:valid} we show examples of computation, validating the results by comparison with the findings of a ray-tracing routine. In Sect.~\ref{sec:geometric} we solve the integral equations for the ideal case of a mirror with 100\% reflectivity, obtaining algebraic expressions for the stray light geometric area, in a completely similar way as we did for the double-reflection area\cite{Spiga2009, Spiga2011}. Results are briefly summarized in Sect.~\ref{sec:conclusions}. We explicitly remark that the formalism for the focused effective area and for stray light are both developed in {\it double cone approximation}. As discussed in detail\cite{Spiga2009}, we are allowed to do this owing to the shallow incidence angles. While the longitudinal curvature of the Wolter-I profile is crucial to concentrate X-ray to a focus, this affects the effective area only to a small extent. For example, in a double cone geometry the incidence angle for a infinitely distant source on-axis is a constant, $\alpha_0$, throughout the entire surface of the primary and the secondary segments. In a real Wolter-I profile, they exhibit a small variation $\Delta \alpha$, related to the curvature of the profile. Fortunately, in grazing incidence geometry it is possible to prove that $\Delta \alpha/\alpha_0 \approx L/4f$, where $L$ is the length of the single segment of the mirror and $f$ its focal length. In practice, the error we make assuming a constant incidence angle is below a few percent in real cases. It can also be proven\cite{Spiga2009} that also the estimation of factors affecting the effective area (e.g., the vignetting) are still on the order of the $L/f$ ratio; therefore, the double cone approximation can be safely applied in the effective area computation. In this work we assume that the double cone approximation can be applied to the stray light theory within a relative error of $L/2f$, which in worst cases amounts to a few percent. Finally, for simplicity we assume the shells to be continuous at the intersection plane. The theory is easy to extend to the case of primary and secondary segments separated by a gap, but we are not reporting it here, in order to avoid a complication of the expressions. | \label{sec:conclusions} In this paper we have reviewed the possible sources of obstruction for focused and stray rays in nested modules of Wolter-I mirrors. We have thereby found integral formulae to compute, in addition to the already known expression for the double-reflection (focused) intensity, the effective area for stray light off the primary (Eq.~\ref{eq:Asl_p}) and the secondary (Eq.~\ref{eq:Asl_h}) mirror segments, also accounting for the finite size of the detector (Eqs.~\ref{eq:Asl_p_lim} and~\ref{eq:Asl_h_lim}). The predictions are in very good agreement with the ray-tracing findings. In the ideal case of a mirror with constant reflectivity, algebraic expressions for the geometric area could be provided in different ranges of off-axis angles. The formalism provided here can be useful in designing a mirror module maximizing the focused effective area and, at the same time, minimizing the stray light impact. In fact, the solution of designing a completely obstruction-free mirror module within the field of view\cite{Spiga2011} might leave too much spacing for the stray light to propagate throughout the mirror nesting. In contrast, the formalism provided here enables, from a given mirror module design, not only a fast assessment of the stray light impact from off-axis sources; it also returned useful relations between the tolerable stray-light magnitude and the obstruction parameters. They therefore provide a way to establish the optimal obstruction to minimize the effective area for stray light while preserving the required effective area in the field of view. Should these formulae be solved numerically for $\Phi$, $\Psi$, and $\Sigma$, the complex task of mirror module design problem could be solved easily without the need to run a complex ray-tracing program. Finally, the same method might be applied to the problem of designing an X-ray baffle, only by a simple re-definition of the obstruction parameters. As a future development of this work, the formalism might be extended to include the case of segmented mirrors, as the ones foreseen for the ATHENA telescope. However, this kind of optics usually include stiffening ribs that represent a further source of obstruction, and the problem becomes more complicated to treat analytically. | 16 | 9 | 1609.09655 |
1609 | 1609.08298_arXiv.txt | We obtain an analytic solution for accretion of a gaseous medium with a adiabatic equation of state ($P=\rho$) onto a Reissner-Nordstr\"{o}m black hole which moves at a constant velocity through the medium. We obtain the specific expression for each component of the velocity and present the mass accretion rate which depends on the mass and the electric charge. The result we obtained may be helpful to understand the physical mechanism of accretion onto a moving black hole. | Accretion of matter onto astronomical objects is a long-standing interesting phenomenon for astrophysicists. Pressure-free gas being dragged onto a star moving at constant velocity was first discussed qualitatively in \cite{Bondi:1944jm}. In the case of spherical accretion onto a stationary black hole, exact solutions have been found in the context of Newtonian gravity \cite{Bondi:1952ni} and in the framework of general relativity \cite{michel1972accretion}. Thereafter accretion has been analyzed in literatures for various black holes, such as a Schwarzschild \citep{Babichev:2004yx}, a charged black hole \citep{michel1972accretion,Jamil:2008bc}, a Kerr-Newman black hole \citep{JimenezMadrid:2005rk,Babichev:2008dy,Bhadra:2011me}, a Reissner-Nordstrom black hole \citep{Babichev:2008jb}, a Kiselev black hole \citep{Yang:2016sjy}, a higher dimensional black hole \citep{Giddings:2008gr,Sharif:2011ih,John:2013bqa,Debnath:2015yva}, a black hole in a string cloud background \citep{Ganguly:2014cqa}, and a cosmological black hole or a Schwarzschild-(anti-)de Sitter black hole \citep{Mach:2013fsa, Mach:2013gia, Karkowski:2012vt, Gao:2008jv}. An exact solution was derived for one or two dust shells collapsing towards a black hole \cite{Liu:2009ts}. Quantum gravity corrections to accretion onto a Schwarzschild black hole were investigated in \cite{Yang:2015sfa}. In general, numerical analysis are required to dealt with nonspherical accretion for either Newtonian or relativistic flow. One exact, fully relativistic, nonspherical solution had been found for a Schwarzschild or Kerr black hole moving through a medium obeying a stiff $P=\rho$ equation of state \citep{petrich1988accretion,1989ApJ}. This solution provides valuable insight into the more general cases and serves as a benchmark for the test of numerical codes. However, because of the mathematical complexity, until now analogous exact solutions still had not been obtained for a moving Reissner-Nordstr\"{o}m black hole. Here we consider accretion onto such a black hole, and find an exact, fully relativistic, nonspherical solution. The result obtained here may be helpful to understand the physical mechanism of accretion onto a moving black hole. The rest of the paper is organized as follows. In next section, we will present the fundamental equations for accretion on to moving black hole. In section III and IV, we will determine the mass accretion rate for a moving Reissner-Nordstr\"{o}m black hole. Finally, we will briefly summarize and discuss our results in section V. | We have considered accretion onto a Reissner-Nordstr\"{o}m black hole. The black hole moves at a constant velocity through the medium which obeys a stiff $P=\rho$ equation of state. We obtained the mass accretion rate which depends on the mass and the electric charge. We obtain the specific expression for each component of the velocity which implies the flow is two spatial dimension. The results obtained here may provide valuable physical insight into the more complicated cases and can be generalized to other types black hole. | 16 | 9 | 1609.08298 |
1609 | 1609.04413_arXiv.txt | Due to its proximity, SN~1987A offers a unique opportunity to directly observe the geometry of a stellar explosion as it unfolds. Here we present spectral and imaging observations of SN~1987A obtained \hbox{$\sim\ 10,000$} days after the explosion with HST/STIS and VLT/SINFONI at optical and near-infrared wavelengths. These observations allow us to produce the most detailed 3D map of H$\alpha$ to date, the first 3D maps for [Ca~II]$\ \lambda \lambda 7292,\ 7324$, [O~I]$\ \lambda \lambda 6300,\ 6364$ and Mg~II$\ \lambda \lambda 9218,\ 9244$, as well as new maps for [Si~I]$+$[Fe~II]$\ 1.644\ \mu$m and He I $2.058~\mu$m. A comparison with previous observations shows that the [Si~I]$+$[Fe~II] flux and morphology have not changed significantly during the past ten years, providing evidence that it is powered by $^{44}$Ti. The time-evolution of H$\alpha$ shows that it is predominantly powered by X-rays from the ring, in agreement with previous findings. All lines that have sufficient signal show a similar large-scale 3D structure, with a north-south asymmetry that resembles a broken dipole. This structure correlates with early observations of asymmetries, showing that there is a global asymmetry that extends from the inner core to the outer envelope. On smaller scales, the two brightest lines, H$\alpha$ and [Si~I]$+$[Fe~II]$\ 1.644\ \mu$m, show substructures at the level of $\sim 200 - 1000\ \kms$ and clear differences in their 3D geometries. We discuss these results in the context of explosion models and the properties of dust in the ejecta. | Supernova (SN) 1987A, located in the Large Magellanic Cloud, is the only modern SN where the ejecta are spatially resolved in optical/NIR imaging observations. This makes it a unique target for studying ejecta composition, energy sources and asymmetries. In this work, we focus on asymmetries by studying the spatial distribution of the inner ejecta. These ejecta have not been affected by the reverse shock and are therefore still in the homologous expansion phase, which was reached about one week after the explosion (\citealt{Gawryszczak2010}). The spatial distribution reflects the conditions at the time of explosion and hence carries information about the progenitor and explosion mechanism (e.g., \citealt{Wongwathanarat2015}). The ejecta of \8 are surrounded by a triple ring system that was created approximately $20,000$ years before the explosion, possibly as a result of a binary merger (\citealt{Morris2007}). The outermost ejecta have been interacting with the inner, equatorial ring since day $\sim 3000$, resulting in a number of hotspots appearing in the optical images as well as a sharp increase in flux across the electromagnetic spectrum (e.g, \citealt{Groningsson2008a,Helder2013,Ng2013}). The optical emission from the ring peaked around $8000$ days and is now decaying as the ring is being destroyed by the shocks (\citealt{Fransson2015}). The same behavior is seen in the infrared (\citealt{Arendt2016}), while the soft X-ray light curve leveled off at $\sim 9500$ days (\citealt{Frank2016}). The X-ray emission from the ring is also affecting the inner ejecta. In particular, it is most likely responsible for the increase in optical emission seen after day $\sim 5000$ (\citealt{Larsson2011}). In the freely-expanding ejecta the observed Doppler shifts are directly proportional to the distance along the line of sight to the centre of the explosion ($v_{\rm{obs}} = z/t$, where $t$ is the time since the explosion on 1987 February 23 and $z$ is the distance). This means that the combination of imaging and spectroscopy makes it possible to infer the three-dimensional (3D) distribution of ejecta. This technique has previously been used for \8 in \cite{Kjaer2010} and \cite{Larsson2013} (L13 from here on). \cite{Kjaer2010} used observations obtained in 2005 with the integral field spectrograph SINFONI at the Very Large Telescope (VLT) to study the [Si~I]$+$[Fe~II]$\ 1.644\ \mu$m and He I $2.058~\mu$m lines. A clearly asymmetric distribution was found, with the ejecta being predominantly blueshifted in the north and redshifted in the south. In L13 we confirmed these results for the $1.644\ \mu$m line with more recent SINFONI data, and also used Hubble Space Telescope (HST) spectra and imaging to compare with the distribution of H$\alpha$. The large-scale 3D emissivity of H$\alpha$ was found to be similar to the $1.644\ \mu$m line, except for H$\alpha$ extending to higher blueshifted velocities, although it should be noted that the 3D information for H$\alpha$ had significantly lower resolution and suffered from strong contamination by scattered light from the ring. In addition, we studied the temporal evolution of the ejecta morphology in the HST images, finding that it changes from an approximately elliptical shape before $\sim 5000$ days, to an edge-brightened more irregular morphology thereafter. This transition coincides with and can be explained by the change in the dominant energy source powering the ejecta (from radioactive decay of $^{44}$Ti to X-rays from the ring, see also \citealt{Fransson2013}). Here we extend the work on the 3D distribution of ejecta using new observations obtained in 2014, approximately 10,000 days after the explosion, with HST/STIS and SINFONI. The STIS observations cover the whole ejecta with narrow (0$\farcs{1}$) slits. This represents a major improvement compared to our previous observations. The last STIS observations covering the whole ejecta were obtained using 0$\farcs{2}$ slits in 2004, when the ejecta were $\sim 2/3$ of the current size. The new observations allow us to create the most detailed map of H$\alpha$ to date, as well as the first spatially-resolved maps of the ejecta in [Ca II]~$\lambda \lambda 7292,\ 7324$, [O~I]~$\lambda \lambda 6300,\ 6364$ and Mg~II~$\lambda \lambda 9218,\ 9244$. From the SINFONI observations, we present new maps of the [Si~I]$+$[Fe~II]$\ 1.644\ \mu$m and He I $2.058~\mu$m lines, which we compare to the STIS data. We also compare with our previous SINFONI observations in order to assess the time evolution of these two lines. This paper is organized as follows: we describe the observations and data reduction in section \ref{obs}, present the analysis and results in section \ref{analysis}, and discuss our findings in section \ref{discussion}. We summarize our conclusions and discuss future prospects in section \ref{conclusions}. | \label{conclusions} We have analyzed HST/STIS observations of \8 obtained $10,000$ days after the explosions. From this we present the most detailed 3D map of the H$\alpha$ emissivity to date, as well as the first 3D information for [O I] $\lambda \lambda 6300, 6364$, [Ca II] $\lambda \lambda 7292,\ 7324$ and Mg~II~$\lambda \lambda 9218,\ 9244$. We have also analyzed SINFONI observations of the 3D emissivity of [Si~I]$+$[Fe~II]~$1.644\ \mu$m and He~I $2.058\ \mu$m from the same epoch, which we compare with the STIS results, as well as previous SINFONI observations. The interpretation of the 3D emissivities depend on the energy sources powering the different lines, as well as the properties of dust in the ejecta. We therefore first summarize our conclusions regarding these issues, followed by the conclusions for the morphology. \begin{itemize} \item The integrated H$\alpha$ profile has changed dramatically since the last spectral observations covering the full ejecta from day 6355, brightening significantly in the range $[-2000,0]\ \kms$. We show that this is primarily due to the brightening of a region in the western ejecta. This is in agreement with our previous finding that H$\alpha$ is primarily powered by the X-ray emission from the ring, which also gives rise to an edge-brightened morphology. We also observe flux increases for [O I], [Ca II], Mg~II and He~I, showing that these lines are also affected by the X-rays. However, the centrally peaked morphology of the He~I line indicates that it is powered by both $^{44}$Ti and X-rays. \item The [Si~I]$+$[Fe~II] emission does not show any significant changes in flux, line profile or morphology compared to previous observations obtained during the past ten years (days 8714 and 6816). This shows that it is powered by the radioactive decay of $^{44}$Ti. In agreement with this, the integrated line profile is consistent with the NuSTAR profile of the $^{44}$Ti decay lines, albeit given large uncertainties. \item The Br$\gamma$ line shows a similar low-surface brightness region (or ``hole") at the centre as H$\alpha$, supporting our previous arguments that the dust is likely to be in in optically thick clumps that affect the optical and NIR in the same way. Furthermore, the constant [Si~I]$+$[Fe~II] line profile argues against dust becoming optically thin during the last ten years. \item A similar large-scale geometry is seen in all lines that have sufficient signal to noise (i.e H$\alpha$, [Ca II], [Si~I]$+$[Fe~II] and He~I). As projected on the sky, the ejecta extend in the north-east to south-west direction. In 3D, the ejecta emission is concentrated between the line of sight and the plane of the ring in the north, but closer to the plane of the sky in the south, resembling a broken dipole. This is consistent with the structure seen in previous SINFONI observations and indicated by previous low-resolution H$\alpha$ observations (\citealt{Kjaer2010}, L13). Remarkably, the large-scale structure correlates with very early observations of asymmetries. We therefore conclude that there is a global asymmetry that extends from the inner metal core to the outer H envelope. \item In [Si~I]$+$[Fe~II] and H$\alpha$, which offer the best signal, substructure is seen on a scale of $\sim 200 - 1000\ \kms$, close to the level of resolution in the different directions. As seen in previous observations, the [Si~I]$+$[Fe~II] is concentrated to two main regions. The H$\alpha$ 3D map differs from this, with, for example, an arc being seen around the ``hole", and the brightest emission being concentrated to the clump in the west. These differences are due to a combination of the true distribution of the different elements and the energy sources powering the two lines. \item The centre of the bright, western clump in H$\alpha$ has a space velocity of $\sim 1800\ \kms$. Both [O I] and Mg~II also have their brightest emission concentrated to this region. This shows that there is high-abundance of O in the H-region that is most affected by X-rays. For Mg~II, a correlation with H$\alpha$ is expected since it is powered by Ly~$\alpha$ fluorescence. The H$\alpha$ map also shows that H is mixed in to about $450\ \kms$. For [Si~I]$+$[Fe~II], the space velocities of the centers of the brightest regions are $\sim 2000\ \kms$ and $\sim 2700\ \kms$ in the north and south, respectively. \end{itemize} \8 is unique among SNe and allows us to probe the explosion physics during the first seconds through its radioactivity, nucleosynthesis and morphology. Our observations have provided additional information about the latter. As has been shown by \cite{Wongwathanarat2015} and other simulations, there is a direct connection between the large-scale instabilities created in the explosion and those in the homologous phase for at least for some progenitors. However, this needs to be probed in more detail for different progenitor models and different explosion physics. In addition, rotational effects and magnetic fields remain to be included in the models. On the observational side, there will be additional important information from different molecules and dust from ALMA. A taste of this has already been seen \citep{Kamenetzky2013, Indebetouw2014,Matsuura2015}, but the full ALMA will have a resolution comparable or higher than that of HST. We also hope to get 3D maps of both the He I $\lambda 2.058 \ \mu$m line and the H$_2$ lines in the NIR, using deeper SINFONI observations. The connection between these different molecules and atomic lines with the mass distribution, however, require a detailed modeling of the emission, including the excitation by the positrons \citep{Jerkstrand2011} and the X-rays from the ring collision. Finally, the compact object in the centre of \8 remains to be revealed. HST, ALMA, NuStar and JWST may provide new opportunities for this. | 16 | 9 | 1609.04413 |
1609 | 1609.03233_arXiv.txt | Virial shocks at edges of cosmic-web structures are a clear prediction of standard structure formation theories. We derive a criterion for the stability of the post-shock gas and of the virial shock itself in spherical, filamentary and planar infall geometries. When gas cooling is important, we find that shocks become unstable, and gas flows uninterrupted towards the center of the respective halo, filament or sheet. For filaments, we impose this criterion on self-similar infall solutions. We find that instability is expected for filament masses between $10^{11}-10^{13}\Msun\,\Mpc^{-1}.$ Using a simplified toy model, we then show that these filaments will likely feed halos with $10^{10}M_{\odot}\lesssim M_{halo}\lesssim 10^{13}M_{\odot}$ at redshift $z=3$, as well as $10^{12}M_{\odot}\lesssim M_{halo}\lesssim 10^{15}M_{\odot}$ at $z=0.$ The instability will affect the survivability of the filaments as they penetrate gaseous halos in a non-trivial way. Additionally, smaller halos accreting onto non-stable filaments will not be subject to ram-pressure inside the filaments. The instreaming gas will continue towards the center, and stop either once its angular momentum balances the gravitational attraction, or when its density becomes so high that it becomes self-shielded to radiation. | The thermodynamic state of gas in cosmic web filaments has important implications for observations and theoretical predictions. It will affect halos which are fed by those filaments, as well as small halos that accrete onto those filaments as part of the cosmic hierarchical growth. Spherical virialization of gas in halos has been a prediction of galaxy formation models for decades. In particular, it has been shown \citep{ro77,silk77,binney77,white91} that a comparison between cooling times and ages of galactic halos can predict the transition from galaxies to groups and clusters. \citet[hereafter BD03]{bd03} derived a stability criterion against gravitational collapse of the gas in the presence of significant cooling. They find that for halos below $M_{crit}\simeq 10^{12}\Msun$ a hot gaseous halo is not expected to form, and gas will free-fall until it reaches the disk, at which point it will stop, radiating its kinetic energy abruptly at that point. This has been confirmed in multiple hydrodynamical simulations \citep[e.g.\@ ][]{keres05,ocvirk08,cafg11} and successfully reproduces star forming galaxies at high-z \citep{db06,dekel09} and the color-magnitude bi-modality \citep{db08,croton06,cattaneo06}. Observational indications of this scenario are gradually accumulating \citep[e.g.\@][]{dijkstra09,kimm10,martin15}. In this letter we derive a criterion for the stability of virial shocks around filaments and sheets that form the cosmic web, analogous to the BD03 criterion for halos. Following \citet[hereafter FG84]{fg84}, we construct self-similar density profiles of filaments. We apply our stability analysis to these profiles to identify filaments around which a stable virial shock is expected to form. This criterion is translated to a more useful form by identifying which halos are expected to be fed by these filaments. We find that filament instability influences a large portion of halos in the universe throughout cosmic age. The stability of filaments has been addressed before, numerically \citep{harford11} and analytically \citep{freundlich14,breysse14}, but without taking into account cooling, and by analyzing the stability of initially static filaments, ignoring the effects of the shock at the filaments' edge. In \S~2 we derive the stability criterion for the existence of virialized gas in 1,2 and 3 dimensional collapse. In \S~3 we relate our local criterion to cosmic filaments according to the self-similar solutions of FG84. In \S~4 we relate these filaments to typical halo masses that will likely be fed by them. In \S~5 we summarize and conclude. | \label {sec:summary} We have shown that in the presence of significant cooling, the accretion process of gas onto cosmic-web structures will not always proceed according to the standard virialization scenario of the infall-heating-cooling sequence. The analysis shown here can be applied for accretion onto spherical halos, cylindrical filaments and planar sheets. For filament of $10^{11}-10^{13}\Msun\,\Mpc^{-1}$, we show that gas is expected to fall without ever passing a shock, resulting in dense, thin filaments with low entropy. This is in complete analogy to spherical cold accretion onto halos that have been shown in BD03 and demonstrated in observations and simulations. Using a simplified toy model for the relation between halo mass and redshift to typical filaments that feed it, we show that throughout cosmic history galaxies and clusters are affected by that instability. In particular, high-z star forming galaxies ($M_{halo}>10^{10}\Msun$ at $z=3$), and low redshift groups and clusters ($10^{12}\Msun\lesssim M_{halo}\lesssim 10^{15}\Msun$ at $z=0$) will be fed by filaments for which the gas is unstable. The process that eventually stops the infall is still unclear, and we postulate that it is either angular momentum support from an original helicity of the filament, or by reduction of the cooling rate due to self-shielding of the gas. A prediction of this work is thus that filament gas, in the non-stable regime, will be highly rotating and angular momentum supported. Both processes are hard to identify in simulations, and have not been examined so far. In their absence, gas in simulations will flow towards the center of the filament until it approaches the numerical gravitational smoothing length, at which point the force will diminish. This indicates that the density and entropy of gas in unstable filaments are a numerical artifact and will not converge to the right values. This problem will be examined in future work. \refereebf{The lack of virialized gas in filaments is expected to significantly affect the outcome of galaxies falling onto the filament, and of halos fed by the filament. Halos falling onto filaments are expected to looe gas through ram pressure stripping, and to enrich the filament with metals. Both these processes will be suppressed when galaxies fall into filaments with no stable atmosphere. Penetration of cosmic-web gas directly to galaxies affects the ISM state, and the gas available for star formation and AGNs, as well as their feedback efficiencies. \citet{mandelker16} analyze the Kelvin-Helmholtz stability of supersonic filaments. They find that filaments lose stability via bulk modes, that correspond to standing waves reflecting through edges of the filament. These results do not account for the effects of gravitational attraction towards the center of the filament, and to angular momentum support, both expected to stabilize the filament further. These effects will be addressed in future work. Observationally, the temperature of the filament could affect its detectability through Lyman$-\alpha$ absorption \citep{narayanan10,wakker15} and emission \citep{martin15}. The temperature of the filaments will also affect the soft X-ray background and the total amount of gas in the ``warm phase'' \citep{cen99,dave01}. All these effects are left for analysis in future work.} | 16 | 9 | 1609.03233 |
1609 | 1609.05415_arXiv.txt | Based on the recently completed {\it Chandra}/ACIS survey of X-ray point sources in nearby galaxies, we study the X-ray luminosity functions (XLFs) for X-ray point sources in different types of galaxies and the statistical properties of ultraluminous X-ray sources (ULXs). Uniform procedures are developed to compute the detection threshold, to estimate the foreground/background contamination, and to calculate the XLFs for individual galaxies and groups of galaxies, resulting in an XLF library for 343 galaxies of different types. With the large number of surveyed galaxies, we have studied the XLFs and ULX properties across different host galaxy types, and confirm with good statistics that the XLF slope flattens from lenticular ($\alpha\sim1.50\pm0.07$) to elliptical ($\sim1.21\pm0.02$), to spirals ($\sim0.80\pm0.02$), to peculiars ($\sim0.55\pm0.30$), and to irregulars ($\sim0.26\pm0.10$). The XLF break dividing the neutron star and black hole binaries is also confirmed, albeit at quite different break luminosities for different types of galaxies. A radial dependency is found for ellipticals, with a flatter XLF slope for sources located between $D_{25}$ and 2$D_{25}$, suggesting the XLF slopes in the outer region of early-type galaxies are dominated by low-mass X-ray binaries in globular clusters. This study shows that the ULX rate in early-type galaxies is $0.24\pm0.05$ ULXs per surveyed galaxy, on a $5\sigma$ confidence level. The XLF for ULXs in late-type galaxies extends smoothly until it drops abruptly around $4\times10^{40}$ erg s$^{-1}$, and this break may suggest a mild boundary between the stellar black hole population possibly including 30 $M_\odot$ black holes with super-Eddington radiation and intermediate mass black holes. | X-ray observations of more than forty years have revealed a variety of X-ray point sources beyond the solar system in the Milky Way and Magellanic Clouds. Although the bright X-ray sources in the Milky Way are easily studied, many of them are not observable due to the heavy obscuration of the galactic disk, and one must correct for completeness when studying the statistical properties of Galactic X-ray binaries \citep[e.g.,][]{Grimm2002}. Studies of X-ray binaries in more distant galaxies may be free from the obscuration problem, and more importantly, provide us uniform samples of X-ray binaries in different environments. Much work has been done to study the X-ray source populations in distant galaxies since the lunch of {\it Einstein} \citep{Fabbiano1989} and {\it ROSAT} \citep{Roberts2000}, and a quantum leap in this field has been achieved with the {\it Chandra} mission \citep[see][for a review]{FW2006, Fabbiano2006}. New classes of X-ray sources not previously seen in the Milky Way have emerged from studies of distant galaxies, such as the ultraluminous X-ray sources (ULXs), which were first detected by {\it Einstein} \citep{Long1983, Fabbiano1989}. ULXs are defined as extranuclear sources with an observed luminosity of at least 10$^{39}$ erg/s, in excess of the Eddington limit of a neutron star. A large amount of both observational and theoretical work has focused on the nature of ULXs, including the high luminosity, the X-ray spectra and short time variability, the association with active star-forming regions, and the formation and evolution of ULXs \citep[see][for a review]{Fabbiano2005, Fabbiano2006}. Recently, it was generally believed that ULXs in nature appear in two categories: the long-sought intermediate mass black holes (IMBHs) as the seeds of supermassive black holes \citep{Heger2002, Miller2004}; or stellar mass black holes in a new ultraluminous accretion state \citep{Makishima2000, King2001, Zezas2002}. The X-ray luminosity function (XLF) is a powerful tool for the characterization of the populations of discrete X-ray sources detected in nearby galaxies. \citet{Kilgard2002} studied three starburst and five spiral galaxies, and found that starburst galaxies have flatter XLF slopes than do normal spirals. \citet{Grimm2003} studied the X-ray sources (mainly high-mass X-ray binaries; HMXBs) in a dozen of late-type and starburst galaxies, and found that the total X-ray luminosity and the XLF scale with the star formation rate (SFR). While HMXB XLFs are typically described by straight power-laws, XLFs are typically described by broken or cut-off power-laws for low-mass X-ray binaries in early-type galaxies or regions lacking current star formation. \citet{Mineo2013} reported the numbers of ULXs (in one colliding galaxy pair) are strongly correlated with the local SFR densities, but the luminosities of these sources show weak correlation with SFR densities. In the general picture, starbursting galaxies have flatter XLFs than spiral or early-type galaxies, disagreements, however, exist for some important details of the XLFs. For low mass X-ray binaries (LMXBs) in galaxies lacking current star formation, the XLFs are typically described by broken or cut-off power-laws, with several possible but controversial knees: ($i$) The knees at a few $10^{37}$ erg s$^{-1}$ \citep{Grimm2003, Gilfanov2004} may be explained by accretion on neutron star from Roche lobe overflow, which is driven by gravitational wave emission when below $\sim~2\times10^{37}$ erg s$^{-1}$, and by magnetic stellar winds at higher luminosities \citep{Postnov2005}. Some studies proposed that systems below $10^{37}$ erg s$^{-1}$ may be transients \citep{Bildsten2004}. However, \citet{Kim2006} did not find such a break in the LMXB XLF in two normal elliptical galaxies. ($ii$) The knees at a few $10^{38}$ erg s$^{-1}$ \citep{Sarazin2000,Sarazin2001} may be consistent with the Eddington luminosity of normal neutron star (NS) binaries, the luminosity of massive neutron stars (NSs), He-enriched NS binaries \citep{Ivanova2006}, or low-mass BH binaries. ($iii$) The knees at a few $10^{39}$ erg s$^{-1}$ \citep{Jeltema2003} could be attributable to ULXs or small number statistics. In some sense, these discrepancies are expected, because those studies use different galaxy samples that are usually small (on the order of a dozen), and they adopt different methods to construct XLFs and different methods to correct for the survey incompleteness. We have embarked on an effort to study the X-ray point sources, especially ULXs, in nearby galaxies with uniform procedures using the wealth of {\it Chandra} Data Archive after eight years' accumulation \citep[][hereafter Paper I]{Liu2011}. As detailed in Paper I, 383 galaxies within 40 Mpc with isophotal major axis above 1 arcminute have been observed by 626 ACIS observations, and our uniform analysis of these observations has led to 11,824 point sources within the 2$D_{25}$ isophotes of 380 galaxies, by far the largest extragalactic X-ray point source catalog of such. There are a large number of galaxies observed for each galaxy morphological type, making it possible to compare XLFs among galaxies of each type and between galaxies of different types with good statistics. Meticulous efforts have been made to identify the nuclear X-ray sources, so that we can excise them from the catalog and study the remaining ordinary X-ray binary populations. In this paper, we study XLFs and statistical properties for ULXs for different samples of galaxies with uniform procedures. In section 2, we describe our treatments of the detection threshold and the foreground/background contamination estimates for extragalactic point source surveys. In section 3, we describe our procedures to construct a stitched survey for individual galaxies aided by an example, and present a library of XLFs for 343 galaxies. In section 4, we describe how to construct surveys of a sample of galaxies, and present the XLFs and statistical properties for different samples of galaxies. We summarize our results and discuss the significance of the XLF break at $4\times10^{40}$ erg s$^{-1}$ in section 5. | Some often-debated questions of XLFs for nearby galaxies could be addressed, with the large number of surveyed galaxies and resulting X-ray point sources, combined with the uniform procedures devised for this survey. \subsection{XLF Breaks} \label{sec.break} Different break luminosity may reflect the differences in the binary formation mechanisms \citep{Voss2009}, the binary configurations and consequently the accretion rates of the X-ray binaries in different types of galaxies. Some studies of X-ray point sources in early-type galaxies have identified an XLF shape that required two power-laws with a break around $3\times10^{38}$ erg s$^{-1}$ \citep{Sarazin2000}. However, such a break may result from biases affecting the detection threshold of the data \citep{KF2003}. \citet{KF2004} analyzed 14 early-type galaxies with varying sizes of point source samples and found that, although XLFs of individual galaxies may not require a broken power law, the combined differential XLF shows a statistically significant break around $L_b=(5\pm1.6)\times10^{38}$ erg s$^{-1}$ with a slope of $\alpha=1.8\pm0.2$ below $L_b$ and $\alpha=2.8\pm0.6$ above $L_b$. Using a {\it Chandra} survey of LMXBs in 24 early-type galaxies, \citet{Humphrey2008} acquired a best-fit power law for the composite XLF, with a break at ($2.21^{+0.65}_{-0.56}$)$\times10^{38}$ erg s$^{-1}$, and slopes being $1.40^{+0.10}_{-0.13}$ and $2.84^{+0.39}_{-0.30}$ below and above it. It is suggested that the lower end of the XLF is mainly populated by neutron star X-ray binaries \citep[e.g., ultracompact binaries that form predominantly in star clusters;][]{Bildsten2004}, while the upper end of the XLF is mainly populated by black hole X-ray binaries \citep{Ivanova2006}. The break would correspond to the Eddington luminosity of massive neutron stars; if always true for individual galaxies, locating this break in the XLF for an early-type galaxy may allow us to determine the distance of the galaxy. This survey confirms that the break is present in the composite XLFs for both early-type galaxies and late-type galaxies \citep[e.g.,][]{KF2004, Humphrey2008}. The composite differential XLF for 130 early-type galaxies shows a significant break around $L_b=(8.9\pm0.2)\times10^{38}$ erg s$^{-1}$ (or $4.6~L_{\rm Edd}$ of a 1.5 $M_\odot$ neutron star) with a slope of $\alpha=1.0\pm0.03$ below $L_b$ and $\alpha=2.3\pm0.2$ above $L_b$. The break is also present in the XLFs for the sample of 55 elliptical galaxies at a break luminosity of $L_b=(4.5\pm0.2)\times10^{38}$ erg s$^{-1}$ and for the sample of 75 lenticular galaxies at a break luminosity of $L_b=(7.4\pm0.3)\times10^{38}$ erg s$^{-1}$. The composite XLF for early-type galaxies in this survey are flatter than those in \citet{KF2004} and \citet{Humphrey2008}, because we have excluded the background/foreground contaminating sources that are present in larger numbers at lower luminosities. The higher break luminosity of early-type galaxies in this paper is examined. First, we study the systematic effect due to the observation bias since we are not dealing with a truly serendipitous survey. A simulation ($10^4$ runs) is performed that 100 galaxies are randomly selected from the 130 early-type galaxies and the composite XLF is fitted with a broken power law. Figure \ref{fig.test} shows the distribution of the fitting parameters compared to the results in this paper. The 1$\sigma$ systematic uncertainty is 0.08$\times10^{38}$ erg s$^{-1}$ for the break luminosity, which is too small to affect the fitting results. Another possible bias is due to that there may be more nearby lower optical luminosity galaxies, which preferentially populate the low $L_X$ part of the XLF, while the higher optical luminosity galaxies contribute more higher $L_X$ objects. For distant galaxies ($D \gtrsim$ 30--50 Mpc), sometimes $Chandra$ resolution becomes insufficient and only the total luminosity of the galaxy can be measured \citep{Gilfanov2004c}. To examine this bias, we divide the early-type galaxies into three samples as 0-20 Mpc (68 galaxies), 10-30 Mpc (95 galaxies), and 20-40 Mpc (62 galaxies), then perform the XLF fittings. The fitted breaks are 8.1$\times10^{38}$, 9.8$\times10^{38}$, and 1.3$\times10^{39}$ erg s$^{-1}$, respectively. A small bias does exist, however, it can not explain the discrepancy between this paper and previous studies. Actually, both the samples in \citet{Kim2004} and \citet{Humphrey2008} included galaxies with different distances, and these studies may also suffer from this bias. Therefore, we think the higher break luminosity in this paper is real. Finally, we investigate the bias induced by the $B$-band luminosity correction. One elliptical sample (NGC1399, NGC1407, NGC3379, NGC4365, NGC4374, NGC4472, NGC4621, NGC4636, NGC4649, and NGC4697) and one spiral sample (NGC253, NGC628, NGC891, NGC2403, NGC3031, NGC3628, NGC4631, NGC4945, NGC6946, and NGC7793) are chosen to produce a composite XLF, respectively. Both of the XLFs are corrected by the $B$-band and $K$-band luminosity, and then fitted with a broken power law. The $K$-band luminosity for these galaxies are from \citet{Jarrett2003} and \citet{KF2004}. The fitted break luminosity is $2.1\times10^{38}$ erg s$^{-1}$ for the ellipticals with $B$-band correction, below that ($2.7\times10^{38}$ erg s$^{-1}$) with $K$-band correction. For the spirals, the break luminosities are $7.2\times10^{38}$ and $1.2\times10^{39}$ erg s$^{-1}$ for the XLFs with $B$- and $K$-band luminosity correction, respectively. These suggest that the blue light correction may be slightly insufficient for correcting the XLF in the low luminosity side. The dependency of the break luminosity on galaxy age is then examined. We collect (averaged) ages of elliptical galaxies from previous studies \citep{Trager2000, Terlevich2002, Thomas2005, Sanchez2006, Annibali2007}, and obtain seven young ($<$ 5 Gyr) and twenty old ($\ge$ 5 Gyr) elliptical samples with detection thresholds below $10^{38}$ erg s$^{-1}$. The fitted break luminosity is 7.08$\times10^{38}$ and 4.67$\times10^{38}$ erg s$^{-1}$ for the young and old elliptical samples, respectively. The age effect indicates possible presence of binaries with massive neutron stars or He-enriched neutron stars \citep{Podsiadlowski2002, KF2004} in young elliptical galaxies. This may account for the higher break luminosity in this paper, since there are a group of young early-type galaxies in our sample, while in previous studies, few such galaxies were included. In addition, \citet{Fragos2008} reported at early times ($<$ 5--6 Gyr) of one galaxy, the XLF receives notable contributions from intermediate-mass X-ray binaries, and that more luminous sources at earlier times makes the shape of the XLF flatter. The composite differential XLF for 213 late-type galaxies shows a break, although less significant, around $L_b=(6.3\pm0.3)\times10^{38}$ erg s$^{-1}$ with a slope of $\alpha=0.6\pm0.03$ below $L_b$ and $\alpha=1\pm0.05$ above $L_b$. Note that these breaks are the results of combining galaxies with different XLFs, and are not necessarily present in individual galaxies, even if the break is universal and the composite XLF is representative of the sample XLF. For instance, a galaxy may not sample the population enough (i.e., too few sources) to unveil the break. In this sense, this break can not be used to determine the distance of host galaxies. \subsection{XLF Normalizations and Slopes} The large number of surveyed galaxies in this study allows to reveal how the XLFs change over different galaxy types of different galaxy properties. In Section 4.2, we normalize the best-fit XLF models by the surveyed blue light to estimate the differential source density, and one obvious trend shown in Figure 12 is that the $\pounds_B$ normalized X-ray point source density at luminosities below $10^{38}$ erg s$^{-1}$ decreases from elliptical to lenticular, to spiral, to peculiar and to irregular galaxies. For instance, the source density around $10^{38}$ erg s$^{-1}$ are 1.99$\pm$0.35/1.51$\pm$0.65/1.07$\pm$0.14/ 1.06$\pm$0.75/0.60$\pm$0.14 for these galaxies. While the blue light $\pounds_B$ is intended as a mass indicator, there is a trend for the mass-to-blue light ratio to decrease along the sequence of galaxy types because the stars in these galaxies become progressively younger, bluer, and brighter on average. In addition, the same amount of mass corresponds to decreasing number of stars/binaries along the sequence because the stars in these galaxies become progressively more massive on average. Another contributing factor is the globular cluster (GC) specific frequency, which decreases from elliptical to spiral galaxies. This leads to less X-ray binaries along the sequence because large fractions (20\%-70\%) of X-ray binaries are formed through stellar interactions in GCs as confirmed by observations \citep[e.g.,][]{Sarazin2003}. Another obvious trend revealed in this study is that the XLF slopes become flatter from early-type to late-type galaxies, as already revealed by previous studies of smaller samples. For example, \citet{Kilgard2002} presented a comparison of XLFs of four nonstarburst spiral galaxies and three starburst galaxies, with the slopes determined from $\sim$1.3 to 0.5. Using $Chandra$ observations of 32 nearby spiral and elliptical galaxies, \citet{Colbert2004} determined the power-law index as $\sim$1.4 for elliptical galaxies, and 0.6-0.8 for spiral and starburst galaxies. The large number of galaxies in this study also allows to confirm this trend for subtypes of galaxies, that is, the XLF slope flattens from lenticular ($\alpha\sim1.50\pm0.07$) to elliptical ($\sim1.21\pm0.02$), to spirals ($\sim0.80\pm0.02$) of S0/a-Sa ($\sim0.98\pm0.03$), Sab-Sb ($\sim0.78\pm0.03$), Sbc-Sc ($\sim0.70\pm0.04$), Scd-Sd ($\sim0.54\pm0.04$) to Sdm-Sm ($\sim0.34\pm0.17$) subtypes, to peculiars ($\sim0.55\pm0.30$), and to irregulars ($\sim0.26\pm0.10$). A flatter XLF slope means a larger population of X-ray sources at high luminosities (i.e., a larger ULX population), and such an XLF slope evolution suggests a connection between the ULX population and the star formation activities, which are believed to be more prominent in later type/subtype galaxies \citep{Zezas1999, Roberts2000, Fabbiano2001, Colbert2004}. Most ULXs found in regions of star formation are HMXBs \citep{King2004}, and a significant correlation has been found between these ULXs and young OB associations \citep{Swartz2009}. Indeed, studies of galaxy samples with different SFRs do show a monotonic increase of the ULX rate with the SFR \citep{Mineo2013}, and the XLF slope becomes monotonically flatter for higher SFRs. Different shape of XLFs (e.g., breaks) of HMXBs and LMXBs can be used to diagnose on-going star formation, while non-detection of luminous sources immediately constrains the star formation rate of the galaxy \citep{Gilfanov2004}. The XLF evolution with different galaxy types may reflect a flatter mass distribution for the compact objects in later type galaxies, meaning more massive black holes in later type galaxies. Alternatively, this may reflect more binaries at higher Eddington luminosities in later type galaxies, due to higher accretion rates from more massive and younger secondaries in later type galaxies. \subsection{Radial Dependency} As mentioned in Section 3.3, for each galaxy, we determine XLFs for sources within elliptical annuli of the galaxy from galactic center to 2$D_{25}$ in steps of 0.1 elliptical radii, which can be used to study the radial dependency of XLF for different galaxy types. Figure \ref{fig.Cr} presents the cumulative curves for sources in various parts of early-type and late-type galaxies. Only galaxies with most of the 2$D_{25}$ region having been observed are used, and this leads to a sample of 101 early-type and 189 late-type galaxies. It seems that the XLF slopes of late-type galaxies are similar for different parts of galaxies. In contrast, The XLF slopes of early-type galaxies become flatter for sources between $D_{25}$ and 2$D_{25}$, indicating a larger proportion of bright sources in the outer part of elliptical galaxies. The cumulative curves are fitted with a single power-law of $N(>L) = A L_{39}^{-\alpha}$, with results shown in Figure \ref{fig.Cr}b. The radial distribution of sources, which is normalized to the actually observed area of each individual annulus in unit of pixel, is determined for different galaxy types (Figure \ref{fig.radius}). Here we select the galaxies with detection thresholds below $10^{38}$ erg s$^{-1}$, and use the sources brighter than $10^{38}$ erg s$^{-1}$ to avoid selection effects. Also, only galaxies with most of the 2$D_{25}$ region having been observed are used, and this leads to a sample of 66 early-type and 120 late-type galaxies. The radial distribution of X-ray sources in early-type galaxies is slightly flatter than the $B$-band surface brightness. Previous studies showed that in elliptical galaxies, the spatial distribution of GC LMXBs is more extended than field LMXBs \citep{Kim2006, Kundu2007}, and the surface density of (blue) GC LMXBs exceeds the stellar surface brightness when the radius increases \citep{Paolillo2011}. Therefore, the slight excess of X-ray surface brightness in Figure \ref{fig.radius}a may suffer from the GC LMXBs \citep{Mineo2014a}. Another feature is that the distribution of luminous sources ($10^{39} \leq L_{X} < 10^{49}$ erg s$^{-1}$) is nearly uniform, which is quite different with less luminous ones. Because GCs host relatively more bright sources than filed LMXBs \citep{Kim2009, Peacock2016}, we conclude the XLF slopes in the outer region of early-type galaxies are dominated by GC LMXBs \citep{Irwin2005}. For late-type galaxies, the X-ray surface brightness significantly extends past the optical, suggesting strong effects from bright young populations. \subsection{Age and Luminosity Dependency of Luminous X-ray Sources} \citet{Kim2010} showed that young ($<$ 5 Gyr) ellipticals host a larger fraction of luminous X-ray sources than old ellipticals. With the young and old elliptical samples (Section \ref{sec.break}), we define the fraction ($F_{LX}$) of luminous X-ray sources (within the $D_{25}$ ellipse) following \citet{Kim2010}, \begin{equation} F_{LX} = N(L_{X} > 5\times10^{38}\ {\rm erg~s^{-1}})/N(L_{X} > 10^{38}\ {\rm erg~s^{-1}}). \end{equation} The luminous X-ray source fractions ($F_{LX}$) are 0.32$\pm$0.04 (68 out of 211) and 0.18$\pm$0.02 (199 out of 1085) for young and old ellipticals, respectively. Therefore, the number of luminous X-ray sources is higher by a factor of $\sim$ 1.8 in the young sample compared to the old sample, similar to that ($\sim$ 2) of \citet{Kim2010}. In addition, no dependency of $F_{LX}$ on the stellar luminosity of the galaxy is found (Figure \ref{fig.fll}b), which is in agreement with \citet{Kim2010}. A further examination, about the radial distribution of X-ray sources in young and old ellipticals, is given in Figure \ref{fig.yoradius}. As shown in \citet{Brassington2008, Brassington2009}, the radial distribution of X-ray sources follows that of the optical light, especially for old ellipticals. No clear discrepancy is found for the young and old elliptical samples. Although young early-type galaxies contain more bright field LMXBs and flatter field-LMXB XLFs, there is no significant difference of the (GC-LMXB and field-LMXB) combined XLFs between young and old early-type galaxies \citep{Lehmer2004}. \subsection{Stellar Mass and SFR Indicators} Previous studies of nearby passive early-type galaxies and star-forming late-type galaxies have shown that the X-ray point-source emission from old LMXBs and younger HMXBs correlates well with galaxy stellar mass and SFR, respectively \citep{Colbert2004}. The total number of LMXBs and their combined luminosity are proportional to the stellar mass of the host galaxy, and the X-ray luminosity of a galaxy due to LMXBs can be used as an independent stellar mass indicator \citep{Gilfanov2004}; on the other hand, the number of HMXBs and their collective luminosity scale with the SFR, therefore HMXBs are a good tracer of the recent star-formation activity in the host galaxy \citep{Bauer2002,Grimm2003,Gilfanov2004a,Gilfanov2004b,Lehmer2008,Mineo2012,Mineo2014b,Mineo2014c}. This is explained by that at high X-ray energies (2--10 keV), the emission power from young stars and hot interstellar gas decreases steeply, while HMXBs start to dominate the galaxy-wide X-ray intensity and therefore correlate strongly with SFR \citep{Persic2002,Persic2004,Persic2007,Ranalli2003,Lehmer2010}. Here we make a simple examination on the population of X-ray binaries and its relation to the stellar mass and SFR of the host galaxy. The masses of galaxies in our sample is collected from \citet{Tully2015}, which are calculated from the $K$-band luminosity with assumed light to mass conversion factor. Galaxies with detection thresholds below $10^{37}$ erg s$^{-1}$ are selected, and sources with luminosities $L_X~>~10^{37}$ erg s$^{-1}$ are chosen following \citet{Gilfanov2004}. However, early-type galaxies typically have higher detection threshold/completeness limit, therefore few early-type galaxies are selected. Figure \ref{fig.massVSlx} displays the number of sources with luminosities $L_X~>~10^{37}$ erg s$^{-1}$ and their collective X-ray luminosity versus stellar mass. A clear trend can be seen, but the distribution is more diffuse than that reported by \citet{Gilfanov2004}. This could be due to two reasons: (1) the masses of the inner part of galaxies (e.g., galaxy nucleus) are not excluded, while nuclear X-ray sources have been excluded; (2) more late-type galaxies are included in the sample, which contains a number of HMXBs and therefore causes some scatter. Figure \ref{fig.sfrVSlx} displays the number of sources with luminosities $L_X~>~2\times10^{38}$ erg s$^{-1}$ and their collective X-ray luminosity versus SFR. Only galaxies with SFR above 0.2 $M_{\odot}$/yr are plotted. The trend is clear, but more diffuse than those from previous studies \citep[e.g.,][]{Gilfanov2004a}. These AGNs (e.g., LINERS, Seyfert galaxies) in our sample may partly contribute to the diffuse distribution \citep{Ranalli2003,Lehmer2010}. \subsection{ULXs in Elliptical Galaxies} This study, with its large number of galaxies and X-ray sources, also reveals some highly significant new results. The radial distribution of the ULXs detected in elliptical galaxies in this paper and previous studies \citep{Liu2005, Swartz2011} is presented in Figure \ref{fig.ulx}. The ULXs have a wide distribution in galaxies, although more ULXs above $10^{39}$ erg s$^{-1}$ are detected between 0.7 $D_{25}$ and $D_{25}$ isophote. The normalized radial distribution (Figure \ref{fig.ulx}b) shows a second peak located out of the 0.5 $D_{25}$ isophote, which may indicate a different population compared with the ULXs in the inner parts of galaxies. Previous studies suggested that there are very few ULXs in early-type galaxies, e.g., $0.1\pm0.1$ ULXs per surveyed early-type galaxy in a $ROSAT$ HRI survey of nearby galaxies \citep{Liu2006}, and most ULXs above $2\times10^{39}$ erg s$^{-1}$ in early-type galaxies are probably foreground or background objects not physically associated with the host galaxies \citep{Irwin2004}. Recently, \citet{Swartz2011} reported a ULX rate of 0.23 ULXs per elliptical galaxy, but only on a $2\sigma$ confidence level. This study, with the foreground/background objects and galactic nuclear sources carefully removed, finds a total of 49 ULXs above $2\times10^{39}$ erg s$^{-1}$ from the stitched survey of 130 early-type galaxies. Given 17.3 foreground/background objects predicted for the stitched survey of these galaxies, the ULX rate is $0.24\pm0.05$ per surveyed early-type galaxy, which is in good agreement with \citet{Swartz2011}; this is five (4$\sim$7) times lower than the rate for late-type galaxies, but nonetheless a $5\sigma$ non-zero result. There is still a significant population of ULXs above $4\times10^{39}$ erg s$^{-1}$, with 17 ULXs and 7.3 foreground/background objects, and a ULX rate of $0.07\pm0.03$ ULXs per surveyed early-type galaxy. Most of these ULXs reside in elliptical galaxies, which has a ULX rate 3-9 times higher than lenticular galaxies, consistent with a much steeper and lower XLF for lenticular galaxies as shown in Figure \ref{fig.12}b. Given the old ages of the stellar populations in early-type galaxies, these ULXs must have rather old and low mass companions to their primary black holes, similar to the ULXs located in the bulges of early-type spirals \citep{Swartz2009}. Such systems have low accretion rates, and are most likely soft X-ray transients with small duty cycles for outbursts \citep{King2002}. An examination of the age dependency of ULXs is also performed using the young and old elliptical examples. The fraction ($F_{ULX}$) of ULXs (within the $D_{25}$ ellipse) is defined as, \begin{equation} F_{ULX} = N(L_{X} > 2\times10^{39}\ {\rm erg~s^{-1}})/N(L_{X} > 10^{38}\ {\rm erg~s^{-1}}). \end{equation} \citet{Kim2010} showed that young ($<$ 5 Gyr) elliptical galaxies hosts more ($\sim$ 5 times) ULX type LMXBs ($L_X > 2\times10^{39}$ erg s$^{-1}$) than old ($>$ 7 Gyr) elliptical galaxies. The ULX fractions ($F_{ULX}$) are 0.06$\pm$0.03 (3 out of 52) and 0.03$\pm$0.01 (19 out of 763) for young and old ellipticals, respectively. Compared to \citet{Kim2010}, a weaker relation is seen between $F_{ULX}$ and age (Figure \ref{fig.fllULX}a). Also, no dependency of $F_{ULX}$ on the stellar luminosity of the galaxy is found (Figure \ref{fig.fllULX}b). \subsection{ULX Classification from A New XLF Break} The multitude of ULXs from this survey has enabled detailed studies of XLFs at the high luminosity end. In a previous study of the X-ray point sources in nearby galaxies with $ROSAT$/HRI, the XLF in late-type galaxies is shown to be a smooth power-law with a break at $10^{40}$ erg s$^{-1}$ \citep{Liu2006}. Such a break was also present in a study of X-ray point sources in star forming galaxies \citep{Grimm2003}. This break, if real, suggests that ULXs below $10^{40}$ erg s$^{-1}$ belong to the same population of stellar mass black hole binaries, while the handful ULXs above $10^{40}$ erg s$^{-1}$ belong to another population of IMBHs with different binary formation mechanisms. However, this break may be an artifact caused by the scarcity of ULXs above $10^{40}$ erg s$^{-1}$ in these studies. Indeed, with 75 ULXs above $10^{40}$ erg s$^{-1}$ from this survey, the XLFs extend smoothly to $4\times10^{40}$ erg s$^{-1}$ without any breaks at $10^{40}$ erg s$^{-1}$ for the all-galaxy sample, the late-type galaxy sample, and the spiral-galaxy sample, confirming that the XLF break at $10^{40}$ erg s$^{-1}$ from previous studies are simply caused by small number of ULXs above $10^{40}$ erg s$^{-1}$. There is a break around $4\times10^{40}$ erg s$^{-1}$ in the new XLFs from this survey as shown in Figure \ref{fig.9}a. Again, this break may be the dividing line between the population of the stellar mass black hole binaries and another population of IMBH binaries. Note that this break luminosity would correspond to massive stellar black holes \citep[as massive as 30 $M_\odot$ as in the case of IC 10 X-1;][]{Silverman2008} with mildly super-Eddington radiation \citep{King2009}, while the best IMBH candidate HLX-1 can have luminosities as low as a few $\times10^{40}$ erg s$^{-1}$ \citep{Farrell2009}. This break may also be an artifact caused by the small number (three) of ULXs above $4\times10^{40}$ erg s$^{-1}$ in this survey. To test whether the break is real, we are carrying out another {\it Chandra}/ACIS survey of $\sim1800$ galaxies within 200 Mpc, the results of which will be reported in another paper. | 16 | 9 | 1609.05415 |
1609 | 1609.02373_arXiv.txt | Charged particles in a magnetosphere are spontaneously attracted to a planet while increasing their kinetic energy via inward diffusion process. A constraint on particles' micro-scale adiabatic invariants restricts the class of motions available to the system, giving rise to a proper frame on which particle diffusion occurs. We investigate the inward diffusion process by numerical simulation of particles on constrained phase space. The results reveal the emergence of inhomogeneous density gradient and anisotropic heating, which is consistent with spacecraft observations, experimental observations, and the recently formulated diffusion model on the constrained phase space. | 16 | 9 | 1609.02373 |
||
1609 | 1609.07283_arXiv.txt | We present three-dimensional (3D) non-local thermodynamic equilibrium (non-LTE) radiative transfer calculations for silicon in the solar photosphere, using an extensive model atom that includes recent, realistic neutral hydrogen collisional cross-sections. We find that photon losses in the \SiI~lines give rise to slightly negative non-LTE abundance corrections of the order $-0.01\,\mathrm{dex}$. We infer a 3D non-LTE based solar silicon abundance of $\lgeps{Si\odot}=7.51$. With silicon \markaschanged{commonly chosen to be} the anchor between the photospheric and meteoritic abundances, we find that the meteoritic abundance scale remains unchanged compared with the \citet{2009ARA&A..47..481A}~\markaschanged{and \citet{2009LanB...4B...44L}} results. | \label{introduction} Silicon is one of the most abundant metals, and has many astrophysical applications. With a solar abundance and \markaschanged{ionisation} energy comparable to those of iron, it is a significant electron donor in the atmospheres of cool stars, and a key source of opacity in the interiors of solar-type stars. This has direct implications on, for example, the predicted solar neutrino flux \citep{2009ApJ...705L.123S,2016EPJA...52...78S}. As an $\alpha$-capture element, patterns in abundance ratios such as $\xfe{Si}$~against $\feh$, in the Milky Way disk \citep[e.g.][]{2002A&A...390..225C}, bulge \citep[e.g.][]{2016MNRAS.460..884H}, and halo \cite[e.g.][]{2007ApJ...659L.161C, 2009A&A...503..533S,2013ApJ...762...26Y}, provide insight into stellar nucleosynthesis and the chemical evolution of the Galaxy. Finally, \markaschanged{silicon is commonly used \citep[e.g.][]{2015A&A...573A..25S,2015A&A...573A..26S, 2015A&A...573A..27G} to set the meteoritic abundances \citep{2009LanB...4B...44L}~on the same absolute scale as the solar photospheric abundances \citep{2009ARA&A..47..481A}, because silicon is the reference element in meteorites where hydrogen is depleted. (Others \citep[e.g.][]{2009LanB...4B...44L} prefer to use a selection of elements to determine the scale factor but in practice the outcome is basically the same as when only employing silicon for the purpose.)} It is therefore important to have accurate stellar silicon abundance determinations, \markaschanged{and particularly for} the Sun. \markaschanged{Unfortunately, errors can enter spectroscopic abundance analyses from a number of different places. Often, errors in the transition probabilities of the spectral lines used to carry out the abundance analysis have a large effect. This is an issue for silicon \citep[e.g.][]{2012ApJ...755..176S}, for which few laboratory measurements have been made within the past thirty years, while theoretical calculations typically have relatively large uncertainties \citep[see for example the critical compilation of][]{2008JPCRD..37.1285K}. Assuming such errors are not systematic, which however is often the case, they can be circumvented by basing the abundance analysis on some weighted mean inferred from many spectral lines. Once given reliable transition probabilities for hopefully many spectral abundance diagnostics, the main systematic errors in the classic spectroscopic methodology arise from the use of one-dimensional (1D) hydrostatic model atmospheres and from the assumption that the material is in local thermodynamic equilibrium \citep[LTE; e.g.][]{2005ARA&A..43..481A}.} \markaschanged{The problems with 1D hydrostatic model atmospheres stem from their unrealistic treatment of convection; since they neglect fluid motions and time evolution, 1D model atmospheres must therefore rely on the Mixing-Length Theory \citep[MLT;][]{1958ZA.....46..108B,1965ApJ...142..841H}, which comes with a number of free parameters that need to be calibrated. Furthermore, spectral lines generated from 1D model atmospheres are too narrow compared to observed line profiles because they neglect the Doppler shifts associated with the convective velocity field and temperature inhomogeneities, so two more free parameters, microturbulence and macroturbulence, must also be invoked in order to fit observed spectra \citep[e.g][Chapter 17]{2008oasp.book.....G}. In contrast, 3D hydrodynamical model solar and stellar atmospheres successfully reproduce the observations to exquisite detail, including the line shapes, shifts and asymmetries \citep[e.g.][]{2000A&A...359..729A, 2009LRSP....6....2N,2013A&A...554A.118P}.} There have been several detailed investigations into the departures from LTE in \SiI~lines in the solar photosphere. Non-LTE \markaschanged{calculations} based on 1D model atmospheres by \citet{2001A&A...373..998W}~found non-LTE abundance corrections ($\lgeps{NLTE}-\lgeps{LTE}$)~that are typically very slightly negative, and of the order $-0.01\,\mathrm{dex}$, a result later consolidated by \citet{2008A&A...486..303S}~using \markaschanged{a more extensive} model atom. The non-LTE calculations by \citet{2012KPCB...28..169S}~using 1D hydrostatic model atmospheres, and by \cite{2012ApJ...755..176S}~using a 3D hydrodynamic model atmosphere and treating each column of the model atmosphere independently (i.e.~the so-called 1.5D~approximation), suggest slightly more severe abundance corrections of the order $-0.05\,\mathrm{dex}$. Beyond the Sun, there is evidence of much larger non-LTE effects in \SiI~lines~\citep[e.g.][]{2009A&A...503..533S, 2011A&A...534A.103S, 2012ApJ...755...36S,2013ApJ...764..115B,2016ApJ...823...36T}. Recent 1D non-LTE calculations by \citet{2016AstL...42..366M}~suggest that the severity of the non-LTE effects in the afore-mentioned studies may have been overestimated. \citet{2016AstL...42..366M}~utilized for the first time the collisional cross-sections of \citet{2014A&A...572A.103B}~for excitation and charge-transfer with neutral hydrogen, which were calculated using the Born-Oppenheimer formalism. Using the model atom of \citet{2008A&A...486..303S}, and 1D \marcs~hydrostatic model atmospheres \citep{2008A&A...486..951G}, \citet{2016AstL...42..366M}~commented briefly that the new collisional data reduce the non-LTE effects to vanishingly small levels in metal-poor turn-off stars. This is a significant result because, prior to these cross-sections becoming available, the semi-empirical recipe of \citet{1968ZPhy..211..404D,1969ZPhy..225..483D}, as formulated by \citet[][]{1984A&A...130..319S}~or \citet[][]{1993PhST...47..186L}, was used, typically with a global scaling factor $\sh=0.1$. This recipe does not provide a \markaschanged{realistic} description of the physics of the collisional interactions, being based on the classical \cite{thomson1912xlii}~electron \markaschanged{ionisation} cross-section; it is typically in error by several orders of magnitude \citep[e.g.][]{2016A&ARv..24....9B}. The canonical solar photospheric silicon abundance itself has seen a slight downwards revision, from $\lgeps{Si\odot}=7.55$~\citep{1989GeCoA..53..197A,1998SSRv...85..161G}, to $\lgeps{Si\odot}=7.51$~\citep{2000A&A...359..755A, 2005ASPC..336...25A,2009ARA&A..47..481A, 2015A&A...573A..25S}. The most recent of these was based on a 3D LTE analysis of nine \SiI~lines and one \SiII~line \markaschanged{in the solar disk-centre intensity spectrum} and adopting 1D non-LTE abundance corrections \markaschanged{to the solar flux spectrum} from \citet{2008A&A...486..303S}. Correcting 3D LTE abundances with 1D non-LTE abundance corrections \markaschanged{in this way} is not consistent; furthermore these abundance corrections were obtained prior to the calculations of \citet{2014A&A...572A.103B}~\markaschanged{for the neutral hydrogen collisional cross-sections.} A detailed investigation using a 3D hydrodynamic model solar atmospheres and 3D non-LTE radiative transfer, with the best available atomic data, is therefore highly desirable. In this paper we study the 3D non-LTE \SiI~line formation in a 3D hydrodynamic model solar atmosphere. To obtain accurate results we construct a realistic model atom that includes recent neutral hydrogen collision data from \citet{2014A&A...572A.103B}. We use the same \stagger~model solar atmosphere as that used by \citet{2015A&A...573A..25S}. This enables us to directly apply our derived abundance corrections to their 3D LTE results, and thereby obtain a consistent 3D non-LTE solar photospheric silicon abundance. | 16 | 9 | 1609.07283 |
|
1609 | 1609.02145_arXiv.txt | We present a 1.4~GHz Karl G. Jansky Very Large Array (VLA) study of a sample of early-type galaxies (ETGs) from the volume- and magnitude-limited \atlas\ survey. The radio morphologies of these ETGs at a resolution of $\theta_{\mathrm{FWHM}} \approx$ 5$^{\prime \prime}$ are diverse and include sources that are compact on sub-kpc scales, resolved structures similar to those seen in star-forming spiral galaxies, and kpc-scale radio jets/lobes associated with active nuclei. We compare the 1.4~GHz, molecular gas, and infrared (IR) properties of these ETGs. The most CO-rich \atlas\ ETGs have radio luminosities consistent with extrapolations from H$_{2}$ mass-derived star formation rates from studies of late-type galaxies. These ETGs also follow the radio-IR correlation. However, ETGs with lower molecular gas masses tend to have less radio emission relative to their CO and IR emission compared to spirals. The fraction of galaxies in our sample with high IR-radio ratios is much higher than in previous studies, and cannot be explained by a systematic underestimation of the radio luminosity due to the presence extended, low-surface-brightness emission that was resolved-out in our VLA observations. In addition, we find that the high IR-radio ratios tend to occur at low IR luminosities, but are not associated with low dynamical mass or metallicity. Thus, we have identified a population of ETGs that have a genuine shortfall of radio emission relative to both their IR and molecular gas emission. A number of mechanisms may conspire to cause this radio deficiency, including a bottom-heavy stellar initial mass function, weak magnetic fields, a higher prevalence of environmental effects compared to spirals and enhanced cosmic ray losses. | Early-type (elliptical and lenticular) galaxies (ETGs) were once considered a homogeneous class of ``red and dead" systems devoid of cold gas and young stars, archetypes of the end point of hierarchical galaxy formation and evolution. However, evidence is mounting that a significant fraction of nearby ETGs are in fact still continuing to form stars. We now know that ETGs commonly host neutral hydrogen (H{\tt I}) distributed in discs, rings, or disturbed structures, with masses ranging from $\sim10^{6} - 10^{8}$ M$_{\odot}$ (e.g., \citealt{morganti+06, oosterloo+10}). Recent statistical searches for H{\tt I} have reported detection rates of $\sim$40\% in field ETGs, and $\sim$10\% in ETGs in more densely populated environments \citep{serra+14}. In addition to cold atomic gas, CO studies have found that many ETGs also harbor substantial reservoirs of molecular gas (e.g., \citealt{knapp+96, welch+03, combes+07}). Recently, the first statistically-complete single-dish CO survey of molecular gas in the \atlas\ galaxies quantified the prevalence of molecular gas in ETGs, reporting a detection rate of 22\% $\pm$ 3\% \citep{young+11}. Interferometric molecular gas imaging studies have shown that ETG molecular gas reservoirs span a range of diverse morphologies and kinematics \citep{young+08, crocker+11, alatalo+13, davis+13}. While secular processes such as stellar mass loss from asymptotic giant branch (AGB) or post-asymptotic giant branch (pAGB) stars may be responsible for the presence of the molecular gas in ETGs in some cases \citep{faber+76, knapp+92, mathews+03, temi+07}, the disturbed morphologies and kinematics of the gas in other cases point to an external origin (i.e., mergers; \citealt{sarzi+06, young+08, duc+15, davis+11, davis+16}). Other authors have suggested that molecular gas in massive ETGs galaxies may originate from cooled gas from the hot X-ray halos in which these galaxies typically reside \citep{werner+14}. While it has become clear that many ETGs contain significant cold gas reservoirs, the ultimate fate of this gas has remained a subject of debate. Whether the gas is actively engaged in star formation (SF), and the efficiency of that SF compared to spiral galaxies, is still unclear. The difficultly in addressing these questions largely arises from the fact that common SF tracers, such as ultraviolet (UV) and infrared (IR) emission, may be contaminated by emission from the underlying evolved stellar population in ETGs \citep{jeong+09, temi+09, sarzi+10, davis+14}. Emission from active galactic nuclei (AGNs) in ETGs can also contaminate many standard SF tracers. Nevertheless, recent studies have argued in favor of the presence of ongoing SF in ETGs. The detection of young stellar populations through UV observations with the {\it Galaxy Evolution Explorer} and the {\it Hubble Space Telescope}, especially in gas-rich ETGs, has provided support for this scenario \citep{yi+05, kaviraj+07, ford+13}. UV emission re-processed by dust in star-forming galaxies and re-emitted in the IR provides another avenue for SF studies of ETGs, and is less susceptible to dust extinction compared to SFR tracers at shorter wavelengths. Although the possibility of contamination from old stars complicates the use of IR emission as a SFR tracer in ETGs, techniques for isolating the portion of IR emission associated with SF have shown promising results (e.g., \citealt{davis+14}). Another potential ETG SFR tracer is radio continuum emission. Unlike other tracers, such as optical or UV emission, centimeter-wave radio continuum emission is virtually unaffected by extinction or obscuration \citep{condon+92}. Recent upgrades at the Karl G. Jansky Very Large Array (VLA) offer the ability to obtain sensitive measurements over relatively short timespans, making radio continuum observations an efficient means of detecting even weak SF in ETGs. Although radio continuum emission may be contaminated by AGNs, strong AGNs can be readily identified based on their radio morphologies (e.g., \citealt{wrobel+91b}) and through comparisons with other SF and AGN diagnostics (e.g., \citealt{nyland+16}). Radio continuum emission is well-established as a SF tracer in late-type galaxies. Studies of the relationship between radio continuum and IR emission have demonstrated a tight correlation between these two quantities that extends over at least three orders of magnitude among ``normal" star-forming galaxies (e.g., \citealt{helou+85, condon+92, yun+01}). This so-called ``radio-IR" relation is believed to be driven by SF in the host galaxy. The radio continuum emission is generated by massive stars as they end their lives as supernovae, accelerating cosmic rays and subsequently producing non-thermal synchrotron emission. Dusty H{\tt II} regions in turn re-radiate optical and UV light emitted by young stars at IR wavelengths. Numerous studies of the radio-IR relation for samples of star-forming spiral galaxies using IR data at both far-infrared (FIR) and mid-infrared (MIR) wavelengths (e.g., \citealt{yun+01, condon+02, appleton+04, sargent+10}) have been performed. However, detailed studies of the radio-IR correlation in ETGs have been rare. Some authors have reported that ETGs closely follow the same tight radio-IR correlation as spiral galaxies \citep{walsh+89, combes+07}, while others have found that ETGs as a class tend to be systematically ``radio faint" \citep{wrobel+91b, lucero+07, crocker+11}. A large, sensitive study of the radio continuum emission on kpc-scales of a statistical sample of ETGs is therefore needed to improve our understanding of the incidence and efficiency of SF in bulge-dominated galaxies. Here, we present new 1.4~GHz VLA observations at 5$^{\prime \prime}$ spatial resolution of a subset of the statistically-complete \atlas\ survey. We combine these new VLA data with existing archival 1.4~GHz measurements to study the global relationship between the radio continuum and IR emission in ETGs. We also compare the radio continuum emission properties to those of the molecular gas in our sample galaxies, all of which have single-dish CO observations available, to study the SF efficiency in ETGs. In Section~\ref{sec:sample}, we describe the \atlas\ survey. We explain the selection, observations, data reduction, and results of our new VLA observations in Section~\ref{sec:data}. Ancillary molecular and infrared data are discussed in Section~\ref{sec:multiwav}. In Section~\ref{radio_corr}, we describe the radio-CO, radio-IR, and IR-CO relations and discuss potential explanations for the observed deficit of radio emission in Section~\ref{sec:discussion}. We summarize our results and provide concluding remarks in Section~\ref{sec:summary}. | \label{sec:summary} We have presented new, sensitive 1.4~GHz VLA observations of the kpc-scale radio continuum emission in 72 ETGs from the volume- and magnitude-limited \atlas\ survey. Combined with data from FIRST, we have studied the 1.4~GHz properties of 97\% of the \atlas\ ETGs. We detected radio continuum emission in 71\% of our new 1.4~GHz VLA observations on scales ranging from $\approx$200 to 900~pc in compact sources to as large as 18~kpc in the most extended source. For the majority of the ETGs in our sample, the 1.4~GHz emission has a morphology that is similar in appearance to the discs of radio emission associated with SF in spiral galaxies. In at least two cases, the radio morphology is characterized by extended jets, and is clearly associated with an active nucleus rather than SF. We compared these radio data with existing molecular gas and IR observations to study the CO-radio and IR-radio relations in the largest sample of nearby ETGs to date. The main conclusions from this study are as follows:\\ \begin{enumerate} \renewcommand{\theenumi}{(\arabic{enumi})} \item The most molecular gas-rich \atlas\ ETGs have radio luminosities consistent with expectations from radio-SFR calibrations and SFRs derived from molecular gas masses \citep{gao+04, murphy+11}. The gas-rich ETGs in our sample also follow the radio-IR correlation. These ETGs may be in the process of efficiently forming stars, and SF likely proceeds in a manner similar to that in typical star-forming spiral galaxies. The radio-IR relation in these systems likely arises from SF, but for some sources harboring low-luminosity radio AGNs, the correlation may be driven by AGN activity. \item ETGs with lower H$_2$ masses tend to emit less radio continuum emission than expected based on standard H$_2$-SFR relations. This population of ETGs is also characterized by high IR-radio ratios compared to ``normal" star-forming galaxies. Correlations between the radio continuum and IR emission are similar for both FIR and MIR emission. High q-values persist in the MIR even after correction for the contribution to the 22$\,\mu$m emission made by an underlying dusty, evolved stellar population. \item The incidence of high q-values is much higher in this sample than in previous studies of the IR-radio relation in samples dominated by late-type galaxies. About 19\% of our sample ETGs have high q-values and are candidate FIR-excess sources. Considering \atlas\ ETGs with only upper limits the level of radio continuum emission, this fraction may even be as high as $\approx$50\%. \item By comparing to lower-resolution archival radio data, we conclude that the amount of large-scale radio emission that would have been resolved-out by our higher-resolution data is modest. While there are some ETGs in our study that have normal star-forming q-values when measurements are made using the lower-resolution radio data, the high q-values persist in other ETGs even when data more sensitive to extended, low-surface-brightness emission are included. \item The high q-values in our sample tend to occur at low-IR luminosities but are not associated with low dynamical mass or metallicity. This is in contrast to previous studies, which were dominated by late-type star-forming galaxies. \item Possible explanations that could explain both the high CO-radio and IR-radio ratios in our sample of ETGs include bottom-heavy IMFs, weak magnetic fields, and a higher prevalence of environmental effects leading to enhanced cosmic ray electron escape compared to spirals. \end{enumerate} Although our data indicate that some ETGs are deficient in their overall radio continuum emission compared to their CO and IR emission, further studies are needed to verify the underlying cause. Improved estimates of SF rates, SF efficiencies, ISM conditions, and galactic magnetic fields in ETGs will also help improve our understanding and interpretation of the correlations discussed in this work. Examples of future studies include spectral energy distribution modeling, deep high-resolution imaging of denser molecular gas species with the Atacama Large Millimeter Array, and deep radio continuum polarization studies capable of tracing the strength and structure of the weak magnetic fields of nearby ETGs. | 16 | 9 | 1609.02145 |
1609 | 1609.02927.txt | When a star passes within the tidal radius of a supermassive black hole, it will be torn apart\cite{Rees88}. For a star with the mass of the Sun ($M_\odot$) and a non-spinning black hole with a mass $<10^8 M_\odot$, the tidal radius lies outside the black hole event horizon\cite{Hills75} and the disruption results in a luminous flare\cite{vanVelzen11, Gezari12, Arcavi14, Holoien14ae}. Here we report observations over a period of 10 months of a transient, hitherto interpreted\cite{Dong16} as a superluminous supernova\cite{Quimby11}. Our data show that the transient rebrightened substantially in the ultraviolet and that the spectrum went through three different spectroscopic phases without ever becoming nebular. Our observations are more consistent with a tidal disruption event than a superluminous supernova because of the temperature evolution\cite{Holoien14ae}, the presence of highly ionised CNO gas in the line of sight\cite{Cenko16} and our improved localisation of the transient in the nucleus of a passive galaxy, where the presence of massive stars is highly unlikely\cite{Lunnan14,Leloudas15}. While the supermassive black hole has a mass $> 10^8 M_\odot$ \cite{ReinesVolonteri15,McConnellMa13}, a star with the same mass as the Sun could be disrupted outside the event horizon if the black hole were spinning rapidly\cite{Kesden12a}. The rapid spin and high black hole mass can explain the high luminosity of this event. | 16 | 9 | 1609.02927 |
||
1609 | 1609.09075_arXiv.txt | We examine the feasibility of detecting auroral emission from the potentially habitable exoplanet Proxima Centauri b. Detection of aurorae would yield an independent confirmation of the planet's existence, constrain the presence and composition of its atmosphere, and determine the planet's eccentricity and inclination, thereby breaking the mass-inclination degeneracy. If Proxima Centauri b is a terrestrial world with an Earth-like atmosphere and magnetic field, we estimate the power at the 5577\AA\ OI auroral line is on the order of 0.1~TW under steady-state stellar wind, or ${\sim} 100 {\times}$ stronger than that on Earth. This corresponds to a planet-star contrast ratio of $10^{-6}-10^{-7}$ in a narrow band about the 5577\AA\ line, although higher contrast ($10^{-4}-10^{-5}$) may be possible during periods of strong magnetospheric disturbance (auroral power $1-10$~TW). We searched the Proxima Centauri b HARPS data for the 5577\AA\ line and for other prominent oxygen and nitrogen lines, but find no signal, indicating that the OI auroral line contrast must be lower than $2\times 10^{-2}$ (with power $\lesssim$ 3,000~TW), consistent with our predictions. We find that observations of 0.1~TW auroral emission lines are likely infeasible with current and planned telescopes. However, future observations with a space-based coronagraphic telescope or a ground-based extremely large telescope (ELT) with a coronagraph could push sensitivity down to terawatt oxygen aurorae (contrast $7\times 10^{-6}$) with exposure times of ${\sim} 1$ day. If a coronagraph design contrast of $10^{-7}$ can be achieved with negligible instrumental noise, a future concept ELT could observe steady-state auroral emission in a few nights.\\[0in] | } The discovery of Proxima Centauri b (henceforth `Proxima Cen b'), only 1.3pc distant from the Sun \citep{Anglada-Escude2016}, ushers in a new era of characterization of nearby potentially habitable exoplanets. Although Proxima Cen b is not known to transit---making transmission spectroscopy impossible---it is an ideal candidate for high-contrast direct spectroscopy using an extremely large coronagraph-equipped telescope. However, even with the enhancement in angular resolution provided by the proximity of its host star, Proxima Cen b's close-in orbit \citep[$a = 0.0485$ AU;][]{Anglada-Escude2016} precludes imaging with current coronagraphs, such as the Gemini Planet Imager \citep[GPI;][]{Macintosh2014} and the Very Large Telescope's Spectro-Polarimetric High-contrast Exoplanet REsearch facility \citep[VLT-SPHERE;][]{Beuzit2008}, which operate primarily in the near-infrared. This is in part due to the poorer Strehl ratios currently achievable at visible wavelengths with ground-based adaptive optics (AO) systems\footnote{See, e.g., \url{https://www.eso.org/sci/facilities/paranal/instruments/sphere/overview.html}}. Consequently, in advance of larger diameter ground- and space-based telescopes, and improvements in visible AO systems, we must initially consider observations that do not rely on transits or current coronagraphy to search for and characterize the atmosphere of Proxima Cen b. Phase curves may offer one of the first means to study the atmosphere of Proxima Cen b \citep{Turbet2016,Kreidberg2016,Meadows2016} by potentially showing the reduction in day-night thermal emission contrast associated with an atmosphere. Phase curves have proven to be a successful means to characterize the atmospheres of planets larger and hotter than Proxima Cen b \citep{Cowan2007, Knutson2007, Knutson2008, Crossfield2010, Brogi2012, Zellem2014, Stevenson2014}, including ones that do not transit \citep{Selsis2011, Faigler2011, Maurin2012, Brogi2014}. However, the expected planet-star contrast ratio in the visible and NIR due to reflected stellar radiation is likely to be below the anticipated systematic noise floor for JWST/NIRSpec \citep{Meadows2016}, and although the planet-star contrast ratio becomes quite favorable beyond $10\mu$m, where the planetary thermal emission peaks, mid-IR phase curves will require JWST/MIRI. To complement the anticipated JWST thermal phase curve measurements, in this work we explore the possibility of directly detecting optical auroral emission from the atmosphere of Proxima Cen b using high-resolution optical spectroscopy. Numerous studies have investigated exoplanet aurorae in the radio due to cyclotron and synchrotron emission to constrain the planetary magnetic field \citep[e.g.][]{Bastian2000,Grie2007,Zarka2007,Hess2011,Driscoll2011,Griessmeier2015}. Others have explored the detectability of optical and UV auroral emission from hot Jupiters, including \citet{France2010}, who searched for far-UV auroral and dayglow H$_2$ emission from the hot Jupiter HD 209458b and placed upper limits on its magnetic field strength, and \citet{Menager2013}, who studied the detectability of Lyman $\alpha$ auroral emission from HD 209458b and HD 189733b. Some studies have also investigated auroral emission from terrestrial planets, including \citet{Smith2004}, who modeled the role of aurorae in redistributing high energy incident stellar flux to the surface of rocky exoplanets, and \citet{Bernard2014}, who investigated how the detection of the green oxygen airglow line could be used to infer the presence of planetary hydrogen coronae on CO$_2$-dominated planets. Finally, \citet{SparksFord2002} suggested that exoplanet airglow and/or aurorae could be detected with a combination of high contrast imaging and high dispersion spectroscopy. However, a detailed calculation of the expected auroral signal strength on a nearby terrestrial exoplanet and the feasibility of its detection has not yet been fully performed. \begin{deluxetable}{lcr} \tablewidth{\linewidth} \tablecaption{Proxima Centauri b properties} \tablenum{1} \tablehead{\colhead{Property} & \colhead{Value$^{\dagger}$} & \colhead{1$\sigma$ Interval}} \startdata Distance from Earth (pc) & 1.295 & \\ Host spectral type & M5.5V & \\ Host mass, $M_\star$ (M$_\odot$)& 0.120 & [0.105 -- 0.135] \\ Period, $P$ (days) & 11.186 & [11.184 -- 11.187]\\ Semi-major axis, $a$ (AU) & 0.0485 & [0.0434 -- 0.0526]\\ Minimum mass, $m_p\sin{i}$ (M$_\oplus$)& 1.27 & [1.10 -- 1.46] \\ Radius, $R_p$ (R$_\oplus$) & Unknown & [0.94 -- 1.40]$^{\ddagger}$ \\ Eccentricity, $e$ & $<0.35$ & \\ Mean longitude, $\lambda$ ($^\circ$)& 110 & [102 -- 118] \\ Inclination, $i$ ($^\circ$) & Unknown & [0 -- 90] \enddata \tablenotetext{$\dagger$}{Values from \citet{Anglada-Escude2016} unless otherwise noted.} \tablenotetext{$\ddagger$}{Plausible range from \citet{Brugger2016}, assuming $m_p = 1.27\mathrm{M}_\oplus$.} \label{tab:sysparams} \end{deluxetable} Detecting optical auroral emission from the possible atmosphere of Proxima Cen b is likely much more favorable for this system than for an Earth-Sun analog. This is due to both planetary and stellar characteristics that favor auroral production and improve detectability (see Table~\ref{tab:sysparams}). In particular, Proxima Cen b's intrinsic planetary properties may favor production of aurorae from oxygen atoms. If Proxima Cen b is Earth-like in composition, recent dynamical/planetary interior modeling results by \citet{Barnes2016} and \citet{Zuluaga2016} suggest that the planet may have a magnetic field, potentially increasing the likelihood of atmospheric retention and of auroral emission. Atmospheres rich in oxygen-bearing molecules, including O$_2$ and CO$_2$, have been predicted for Proxima Cen b \citep{Meadows2016} as a result of the evolutionary processes for terrestrial planets orbiting M dwarfs \citep{LugerBarnes2015,Barnes2016}. On Earth, the oxygen (OI) auroral line at 5577\AA\ provides the distinctive green glow observed in both the Aurora Borealis and the Aurora Australis, and is the brightest (i.e., highest photon emission rate) auroral feature \citep{Chamberlain1961, Dempsey2005}. For emissions from the upper atmosphere, only the 1.27$\mathrm{\mu}$m O$_2$ airglow and combined near-infrared OH night glow features are brighter \citep{Hunten1967}. The oxygen green line is seen in both the Earth's O$_2$-rich atmosphere \citep{Chamberlain1961} and Venus' CO$_2$-dominated atmosphere \citep{Slanger2001}, where it has been observed to increase in brightness after CME events \citep{Gray2014}. The stellar properties and the planet-star separation are also likely to enhance the auroral power on Proxima Cen b relative to an Earth-Sun analog. Proxima Centauri is an active flare star with a magnetic field ${\sim} 600\times$ stronger than that of the Sun \citep{Reiners2008,Davenport2016}. Since stellar activity drives auroral emission for an Earth-like magnetosphere, such features may be much stronger on planets orbiting active M dwarfs. Additionally, with a close-in orbit of 0.0485 AU, Proxima Cen b is about $20\times$ closer to Proxima Centauri than the Earth is to the Sun \citep{Anglada-Escude2016}. This proximity further increases particle fluxes incident on the planetary atmosphere that drive ionization and the subsequent recombination radiation. In addition to increasing the likelihood and strength of the aurora, the characteristics of the Proxima Centauri system may also enhance its detectability. Since the Proxima system is only 1.3pc away, it is perhaps the best-case scenario for the detection of the faint auroral signal from a terrestrial exoplanet. Even though the planet-star contrast ratio in reflected visible light is poor \citep[${\lesssim} 10^{-7}$; see][]{Turbet2016,Kreidberg2016,Meadows2016}, if Proxima Cen b exhibits auroral emission, this will brighten the planet and potentially boost the planet-star contrast by one or more orders of magnitude at the wavelengths of the auroral emission features. The short wavelength of the oxygen green line also improves the contrast of the planet relative to the star due to the star's cool temperature and TiO absorption, which strongly suppresses the brightness of the star in the visible. This improvement in contrast is significantly less for the near-infrared O$_2$ 1.27$\mathrm{\mu}$m and OH airglow lines. In addition to increasing the contrast, the small semi-major axis of Proxima Cen b results in an orbital velocity of ${\sim} 50$ km/s, which will cause its auroral emission to be Doppler-shifted by as much as 1\AA\ over the course of its orbit, making it easier to disentangle it from stellar features via high resolution spectroscopy. An additional advantage of the short wavelength of the OI feature is the smaller inner working angle and point-spread function that may be achieved with a coronagraph at that wavelength \citep{Agol2007}. These factors all improve the chance of detection with ground-based telescopes. The detection of the oxygen auroral line at 5577\AA\ would provide an important diagnostic for planetary properties. Its detection would not only confirm the existence of the planet, but would point to the presence of an atmosphere with abundant oxygen atoms, which is more likely to indicate a terrestrial body. Additionally, the detection of the line would yield a measurement of the radial velocity (RV) of the planet, which combined with the RV measurements of the star \citep{Anglada-Escude2016} would enable the measurement of the eccentricity and inclination of the orbit, ultimately yielding the mass of the planet \citep[see, e.g.,][]{LovisFischer2010}. Detection of the oxygen auroral line would therefore provide several key planetary parameters that could be used to constrain Proxima Cen b's potential habitability \citep{Barnes2016,Meadows2016}. This paper is organized as follows: in \S\ref{sec:signal} we calculate the expected auroral emission strength of Proxima Cen b under different assumptions of stellar and planetary properties. In \S\ref{sec:detect} we model the planet-star contrast ratio in a narrow band centered on the OI 5577\AA\ line and calculate the integration times required to detect the feature with different instruments. In \S\ref{sec:search} we conduct a preliminary search for auroral emission in the HARPS high-resolution, ground-based spectroscopy used by \citet{Anglada-Escude2016} for the RV detection of Proxima Cen b. Finally, in \S\ref{sec:disc} we discuss our results and present our conclusions. | \label{sec:disc} Our calculations above assume that Proxima Cen b is a terrestrial planet with an Earth-like atmosphere. Although its radius is not known --- making an estimate of its density impossible --- the planet is statistically likely to be rocky. The \emph{a priori} probability distribution for the inclination of an exoplanet is $\mathrm{P}(i)\mathrm{d}i = \sin(i)\mathrm{d}i$. With 68\% confidence, the inclination of Proxima Cen b is greater than 47$^\circ$, and with 95\% confidence it is greater than 18$^\circ$. Given $m_p\sin i = 1.27\mathrm{M}_\oplus$, this corresponds to true planet masses smaller than 1.7M$_\oplus$ (68\% confidence) and 4.1M$_\oplus$ (95\% confidence). Recent exoplanet population studies suggest that the transition between rocky and gaseous exoplanets occurs at a radius of about ${\sim}1.6\mathrm{R}_\oplus$ \citep{Rogers2015, Wolfgang2015}. However, a corresponding value for the transition \emph{mass} is still uncertain, and since the radius of Proxima Cen b has not been measured, we cannot argue for its terrestrial nature based on its mass alone. Nevertheless, we can obtain predictions for its radius under certain assumptions. Assuming it is a rocky planet with Earth-like composition, we may use the scaling law from \citet{Fortney2007} to obtain radii of 1.16R$_\oplus$ ($m_p = 1.7\mathrm{M}_\oplus$) and 1.45R$_\oplus$ ($m_p = 4.1\mathrm{M}_\oplus$). Using the mass-radius grids of \citet{Lopez2012} and assuming instead that Proxima Cen b is a super-Earth/mini-Neptune with a thin H/He envelope with mass equal to 1\% the planet mass, the radii jump to 1.7R$_\oplus$ ($m_p = 1.7\mathrm{M}_\oplus$) and 2.1R$_\oplus$ ($m_p = 4.1\mathrm{M}_\oplus$). Planet occurrence rate calculations for cool M dwarfs \citep{Dressing2013} suggest that there is a steep drop in the number of short-period planets per star with radii above 1.4R$_\oplus$ ($0.19^{+0.07}_{-0.05}$), compared to those with radii below 1.4R$_\oplus$, which are more than twice as common ($0.46^{+0.09}_{-0.06}$). This suggests that Proxima Cen b is more likely to be terrestrial than Neptune-like. A thin ($\lesssim 1\%$ by mass) H/He veneer is still possible, but given the extended pre-main sequence phase of the host star, past hydrodynamic escape is likely to have blown it off \citep{Luger2015, Barnes2016}. Nevertheless, we cannot definitively rule out the possibility that Proxima Cen b has an atmosphere dominated by H/He, in which case we would not expect OI auroral emission. A search for Lyman-Werner H$_2$ emission in the UV would be more appropriate in this case. Although broader than the lines we consider here, this emission is likely stronger, and is unlikely to be confused with stellar emission, given that it is molecular in origin. But perhaps more importantly, a robust \textit{non}-detection of this and other H/He features could rule out a large gaseous envelope and confirm the terrestrial nature of the planet. That said, we are currently unable to efficiently probe near-face-on orbits due to the much smaller Doppler shift of the planetary lines. Observations made exclusively at quadrature, when the planet RV is highest, may help with this in the future. Alternatively, Proxima Cen b could be terrestrial but be significantly larger than the Earth, with mass as high as ${\sim}4\mathrm{M}_\oplus$ and radius ${\sim}1.5\mathrm{R}_\oplus$. Since the scaling methods used in \S\ref{sec:signal} implicitly assume Proxima Centauri b is similar to Earth in size, the auroral strength could be different than what we estimate. Assuming a global field, a larger planetary radius (and therefore core radius) could increase the magnetospheric cross-section to the stellar wind, leading to an increase in the emitted power. Furthermore, assuming an Earth-like atmospheric composition, the higher surface gravity would decrease the ionospheric scale height, which could lead to larger magnetic field parallel potential drops in auroral acceleration regions. This would increase the upward flowing current and therefore the downward accelerated electron beams into the upper atmosphere, which could also enhance the auroral signal. Moreover, an increased atmospheric density at the depth where precipitating electrons deposit their energy could also change recombination rates and alter the energy level distribution of the O atoms, which would in turn affect the auroral strength in different lines. A quantitative estimate of these effects is beyond the scope of this study, as it would require both modeling the changes to the atmospheric structure and solving the Boltzmann kinetic transport equation. Assuming Proxima Cen b is terrestrial, our HARPS search constrains its auroral power to be $< 3\times 10^{3}$~TW. This is consistent with the calculations in \S\ref{sec:signal}, which suggest the OI auroral power on Proxima Cen b is likely ${\sim} 0.1$~TW, or ${\sim} 100\times$ that of the Earth during steady-state solar wind conditions. Those calculations, however, ignore transient increases in stellar magnetic activity, which can enhance the auroral signal and the diffuse airglow emission of the planet. As discussed in \S\ref{sec:signal}, transient magnetospheric activity could result in auroral power for the 5577\AA\ line up to $10-100$~TW, lasting from $10-10^3$ minutes (\S\ref{sec:signal_m2}). In addition, for a planet that is under near constant CME activity, it is possible that the storm conditions last for weeks or longer \citep{Gonzalez1994,Gonzalez1999}. Spectra taken during periods of vigorous stellar activity could thus enhance the chances of detecting auroral emission. Even if Proxima Cen b is terrestrial, an auroral signal is not guaranteed to be present. The existence of an atmosphere is still an open question, owing to vigorous past hydrodynamic escape \citep{LugerBarnes2015, Barnes2016}, observed persistent stellar activity \citep{Davenport2016} and an observationally unconstrained planetary magnetic field. If an atmosphere is in fact present, it may not be Earth-like; instead, it could be dominated by CO$_2$ \citep[e.g.][]{Meadows2016}. Airglow and auroral 5577\AA\ emission are still expected for such an atmosphere, since atomic oxygen is produced by photodissociation of CO$_2$; in fact, OI 5577\AA\ emission has been observed at both Mars and Venus, both of which are CO$_2$-dominated \citep[e.g.][]{Bertaux2005, Slanger2001}. In particular, \citet{Slanger2001} and \citet{Slanger2006} found that the Venusian airglow strength is comparable to that of Earth. Given that Venus receives about twice the solar flux Earth receives, airglow and/or auroral emission from a CO$_2$-rich Proxima Cen b could be a factor of ${\sim}2$ weaker than the values we predict in this paper, although detailed photochemical modeling is required to accurately model this scenario. In the case that the atmosphere is oxygen-rich, nitrogen may need to be present to enhance the auroral signal. On Earth, the OI green line emission results primarily from O$_2^+$ dissociative recombination, as well as collisions with excited N$_2$ and direct electron impact \citep{Strickland2000}. It is unclear whether or not other molecular species could play a similar role if N$_2$ is not the bulk atmospheric constituent. However, the detection of the 3914\AA\ N$_2^+$ band could be a good diagnostic in the UV, where the star is even fainter. If we assume that the power of the 3914\AA\ nitrogen band is comparable to that of the OI line (which is typical for higher energy magnetospheric particle populations), then the planet-star contrast in the N$_2^+$ band would be an order of magnitude greater than at 5577\AA. Since the strength of the N$_2^+$ band scales with magnetospheric parameters, stellar activity could cause strong transient features in the UV, which may be observable. Note, however, that limitations in UV detector efficiencies may complicate the detection of nitrogen and other UV aurorae. Our integration times for the predicted steady-state 5577\AA\ OI auroral line render its detection infeasible for current facilities. However, if key design goals are met for future coronagraphs, steady-state aurorae may be more easily detected. As shown in Fig.~\ref{fig:contrast}, achieving the optimal star-planet contrast ratio at the emission line requires that the width of a spectral element (resolving power) is smaller (greater) than the line's equivalent width. For our predicted steady-state auroral emission (${\sim}0.1$~TW), this requires future spectrographs to achieve $R \gtrsim 10^5$. High-resolution spectroscopy is also needed to resolve the Doppler shift of the planetary auroral emission (${\sim}1$\AA) and place strong constraints on the eccentricity and mass/inclination of the planet. Such constraints would lead to greater confidence in the terrestrial nature of Proxima Cen b. Furthermore, since read noise and dark current dominate the coronagraph instrumental noise budget, the development of low-noise detectors, e.g. MKIDS \citep{Mazin2012,Mazin2015}, would significantly help the detection sensitivity and would allow such high-resolution spectroscopy to be downbinned to the lower resolution typically considered for direct exoplanet spectroscopy. For instance, if future detectors render read noise and dark current negligible, a LUVOIR concept telescope could observe a $0.1$~TW OI $5577$\AA\ auroral emission feature in 100 hours as opposed to the $2 \times 10^3$ hours for the noised observations considered in \S\ref{sec:detect}. In addition, low-noise detectors would make stacking short observations (required to mitigate broadening of the line due to the planet's orbital motion) more feasible. If TMT is built with a coronagraph that can achieve a design contrast of $10^{-7}$ and negligible instrumental noise, it could observe steady-state auroral emission (${\sim}0.1$~TW) in a few nights. However, such an observation would require the development of an effective AO system in the optical. Alternatively, observations made during periods of vigorous stellar activity may enhance the detectability of exo-aurorae on Proxima Cen b. Transient magnetospheric activity could increase auroral power to $1-100$~TW, depending on the planetary magenetic dipole strength, allowing future coronagraph-equipped TMT and LUVOIR telescopes to detect auroral emission in 1 hour or less. This is comparable to estimated CME timescales \citep{Khodachenko2007} and much shorter than the timescales of long-lasting solar storm conditions \citep{Gonzalez1994,Gonzalez1999}. Future observing missions similar to the MOST campaign \citep[e.g.][]{Davenport2016} could be used to characterize and monitor Proxima Centauri's activity levels to constrain the star's activity cycles. Such missions could aid in scheduling spectroscopic observations of Proxima Centauri. Observing the star following a CME-like event or long duration fast solar streams could enhance detectability of the planetary auroral signal. The methods of exo-auroral detection discussed here are not limited to Proxima Cen b, but may be applicable to any exoplanet orbiting a nearby late-type star or brown dwarf. For example, the recently discovered TRAPPIST-1 system \citep{Gillon2016} consists of three planets orbiting an active late M8 ultracool dwarf only 12 pc away; one planet in the system, TRAPPIST-1d, potentially lies in the habitable zone. Since TRAPPIST-1 is a later type star than Proxima Centauri, it is likely more active \citep[e.g.,][]{West2008} and hence could generate larger particle fluxes and a stronger interplanetary magnetic field than Proxima Centauri, leading to more powerful aurorae on its planets. Additionally, the redder blackbody spectrum of TRAPPIST-1 makes it a factor of about 6 dimmer than Proxima Centauri at the OI $5577$\AA\ line, resulting in far more favorable contrast ratios. However, due to its distance, auroral emission from this system will be ${\sim} 100\times$ dimmer than that from Proxima Centauri, likely making its detection infeasible. For coronagraphic observations, the distance to the TRAPPIST system would require an inner working angle smaller than the diffraction limit to extend as long as 5577\AA\ for all known TRAPPIST-1 planets observed with a 10m class telescope, further complicating the observation. Another planet to consider is GJ1132b, which orbits a M3.5 star 12 pc away \citep{Berta2015}. Since it receives ${\sim} 19\times$ the Earth's flux and may have an O$_2$ rich atmosphere \citep{Schaefer2016}, it could display strong auroral emission. However, as with the TRAPPIST-1 system, its distance makes auroral characterization difficult. Moreover, the earlier type host emits a larger fraction of its light in the optical, resulting in a poorer auroral contrast ratio. Finally, exoplanets orbiting nearby brown dwarfs may be prime targets for exo-auroral searches. Early-type brown dwarfs display significant magnetic activity \citep{West2008} and are significantly fainter than M dwarfs in the optical, potentially enhancing the detectability of the 5577\AA\ signal from planets in orbit around them. Although no short-period exoplanets are currently known to orbit nearby brown dwarfs \citep{He2016}, the methods described in this paper may be used as means of exoplanet detection, as suggested by \citet{SparksFord2002}. Since the stack-and-search method described in \S\ref{sec:search} does not require previous RV observations of a system, a long baseline of spectroscopic observations of nearby M dwarfs and brown dwarfs could be used to search for Doppler-shifted 5577\AA\ OI emission. Our method is particularly sensitive to short-period terrestrial planets, whose auroral power (if an atmosphere is present) is large and whose large RV will Doppler-shift the signal by one or more \AA. However, since the stack-and-search method is best suited to detect steady-state emission, exoauroral searches will likely have to wait for a future generation of space-based telescopes or noiseless ground-based ELTs capable of detecting sub-TW aurorae. These searches may someday reveal the presence of unknown nearby terrestrial exoplanets, including ones in the habitable zone. \clearpage All code used to generate the tables and figures in this paper is open source and available at \url{https://github.com/rodluger/exoaurora}. A static version of the code is archived at \url{https://doi.org/10.5281/zenodo.192459}.\\[0in] | 16 | 9 | 1609.09075 |
1609 | 1609.00362_arXiv.txt | We analyze eight epochs of \emph{Hubble Space Telescope} H$\alpha$+[\ion{N}{ii}] imaging of $\eta$ Carinae's outer ejecta. Proper motions of nearly 800 knots reveal that the \btxt{detected} ejecta are divided into three \btxt{apparent} age groups, \btxt{dating to around 1250 A.D., to around 1550 A.D., and to during or shortly before the Great Eruption of the 1840s.} Ejecta from these groups reside in different locations and provide a firm constraint that $\eta$ Car experienced multiple major eruptions prior to the 19\textsuperscript{th} century. The 1250 and 1550 events did not share the same axisymmetry as the Homunculus; the 1250 event was particularly asymmetric, even one-sided. In addition, the ejecta in the S ridge, which have been associated with the Great Eruption, appear to predate the ejection of the Homunculus by several decades. \btxt{We detect essentially ballistic expansion across multiple epochs.} We find no evidence for large-scale deceleration of the \btxt{observed} knots that could power the soft X-ray shell by plowing into surrounding material, suggesting that the observed X-rays arise instead from fast, rarefied ejecta from the 1840s overtaking the older dense knots. \btxt{Early deceleration and subsequent coasting cannot explain the origin of the older outer ejecta---significant episodic mass loss prior to the 19\textsuperscript{th} century is required.} The timescale and geometry of the past eruptions provide important constraints for any theoretical physical mechanisms driving $\eta$ Car's behavior. Non-repeating mechanisms such as the merger of a close binary in a triple system would require additional complexities to explain the observations. | \label{sec:intro} One of the most remarkable stars in our galaxy, $\eta$ Carinae has been puzzling astronomers for over 150 years. In the mid-nineteenth century, it became increasingly variable, then peaked temporarily as the second brightest star in the sky \citep{innes1903,davidsonhumphreys1997,frew2004,smithfrew2011} before slowly fading over more than a decade. During this Great Eruption, $\eta$ Car ejected an estimated 10--15 M$_{\sun}$ into the well-known bipolar Homunculus nebula \citep{smith2003}. A second, Lesser Eruption followed in 1890 \citep{innes1903,humphreys1999,frew2004}, but only ejected $\sim0.1$ M$_{\sun}$ \btxt{\citep{ishibashi2003,smith2005a}}. $\eta$ Car belongs to a class of stars known as luminous blue variables \citep[LBVs,][]{humphreysdavidson1994}, very massive, unstable, post-main-sequence stars characterized by luminous mass-loss events. Even among LBVs, $\eta$ Car is unusual and its parameters are extreme. In its current quiescent state, $\eta$ Car is substantially more luminous than most other known LBVs \citep[e.g.,][]{vangenderen2001,smithtombleson2015}. It is one of only two Galactic LBVs that has been observed in a giant eruption. The other is P Cygni, whose largest eruption involved significantly less energy and mass loss, similar to $\eta$ Car's Lesser Eruption \citep{smithhartigan2006}. Moreover, $\eta$ Car has a massive binary companion in an eccentric 5.5-year orbit \citep{damineli1996,damineli1997,damineli2000,corcoran2001,whitelock2004}, and is located in the rich cluster Trumpler 16, home to dozens of O-type stars \citep{smith2006a}. In contrast, most LBVs are relatively isolated and lack O-type neighbors \citep{smithtombleson2015}. The mechanism of $\eta$ Car's Great Eruption---which released roughly $10^{50}$ ergs of kinetic energy \citep{smith2003,smith2008}---remains a mystery. Many theories treat it as part of single-star evolution, invoking super-Eddington radiation-driven winds \citep{davidson1971,maeder1983,dejager1984,lamersfitzpatrick1988,stotherschin1993,glatzelkiriakidis1993,glatzel1994,humphreysdavidson1994,shaviv2000,owocki2004}. However, the source of the increased bolometric luminosity in these scenarios is unclear. Alternatively, the eccentric orbit of $\eta$ Car's companion has been taken to imply that the Great Eruption was influenced by periastron interactions between the two binary components. Based on nineteenth-century observers' estimates of the primary star's color and brightness, its radius must have been much larger than at present, large enough that its companion would significantly interact or even physically collide \citep{iben1999,smith2011}. \btxt{A speculative idea is that} the collision mixed fresh nuclear fuel to greater depths, causing a sudden burst of increased nuclear burning \citep{smith2011}. It has also been proposed that periastron tidal interactions spun up the primary to unstable rates, leading to a burst of mass loss \citep{cassinelli1999}, or that the Great Eruption was fueled by accretion from the primary onto its companion \citep{soker2007,kashisoker2010}. Still other theories postulate a hierarchical triple system, in which the close inner pair either merged \citep{gallagher1989,iben1999,morrispodsiadlowski2009,podsiadlowski2010,portegieszwartvandenheuvel2016} or underwent a dynamical exchange with the outer companion \citep{liviopringle1998}. Models of the driving cause of the Great Eruption must also \btxt{incorporate} the outer ejecta, a collection of irregular condensations found out to nearly half a parsec outside the Homunculus \citep{thackeray1950,walborn1976,meaburn1996a,smithmorse2004,weis2012}. These outer ejecta (Figure \ref{fig:labels}) are highly nitrogen-rich, suggesting a substantial degree of CNO processing \citep{davidson1986,smithmorse2004}. They contain a minimum mass of 2--4 M$_{\sun}$ \citep{weis2012}, with dust observations suggesting a much larger total mass \citep{gomez2010}. The various models for the Great Eruption produce different explanations for the outer ejecta. The merger model of \citet{portegieszwartvandenheuvel2016}, for instance, predicts that the outer ejecta were formed after the formation of the Homunculus. The proper motions of the outer ejecta provide concrete constraints on $\eta$ Car's mass-loss history. The bright S condensation and the ``jet''-shaped N bow (see Figure \ref{fig:labels}) have motions consistent with having been ejected during the Great Eruption \citep{walborn1978,ebbets1993,currie1996,morse2001}. Some results have suggested, however, that the extended S ridge is up to one hundred years older \citep{walborn1978,morse2001}. The age of the E condensations is even less clear: \citet{walborn1978} found transverse velocities of 300--400 km s$^{-1}$, indicating ejection dates in the mid-1400s, but \citet{walbornblanco1988} determined ten years later that the same features had slowed dramatically, suggesting they were from the Great Eruption after all. While the motions of the outer condensations have hinted at prior mass-loss events, a single ejection date around the time of the Great Eruption could not be ruled out. \begin{figure} \includegraphics[width=\columnwidth]{f1} \caption{\emph{HST} WFPC2 image of $\eta$ Car in the F658N filter, which captures intrinsic and scattered [\ion{N}{ii}] $\lambda$6584 emission along with redshifted H$\alpha$ emission \citep{morse1999,morse2001}. Prominent features are labeled according to the convention of \citet{walborn1976} and \citet{weis2012}.} \label{fig:labels} \end{figure} In this paper, we measure the proper motions of $\eta$ Car's outer ejecta to unprecedented accuracy, using 16 different baselines over 21 years of \emph{Hubble Space Telescope} (\emph{HST}) data. The depth and resolution of the \emph{HST} images allow us to re-evaluate the origins of the N, E, and S features, and, for the first time, measure motions of the fainter NNE, NW, and SE condensations. We find no evidence of widespread deceleration, and show that \btxt{while some of} the outer ejecta come from the Great Eruption (or the lead-up to it), \btxt{many features require} at least one prior mass-loss event centuries earlier. Our data, image registration, and approach to measuring proper motions are described in Section \ref{sec:obs}; the results are presented in Section \ref{sec:results}. We discuss the implications of our results on models of $\eta$ Car in Section \ref{sec:disc} and conclude with a summary in Section \ref{sec:conc}. | \label{sec:conc} We have aligned eight epochs of \emph{HST} imaging (both WFPC2 and ACS) of $\eta$ Car's outer ejecta to the same distortion-corrected reference frame and measured the proper motions of 792 ejecta features, many for the first time. We achieve unprecedented time coverage, with each feature measured in up to 16 baselines over 21 years, as well as unprecedented velocity precision (few km s$^{-1}$) and spatial resolution for these features. All 792 features measured in $\eta$ Car's outer ejecta have transverse velocities pointing nearly directly away from the star. The majority \btxt{have proper motions} of 300--600 km s$^{-1}$, although some are as fast as 1500 km s$^{-1}$. The fastest-moving material is found in the large feature known as the S ridge and in the broadly jet-shaped N bow. Both date back to $\eta$ Car's Great Eruption in the 1840s or to a few decades prior. Over the 21 years of data, we see no evidence for large-scale acceleration or deceleration of any of the outer ejecta: 94\% of the knots are consistent with moving at constant velocity over that time. \btxt{Comparison to images from 1949--1950 support ballistic motion over a longer time period.} Under the assumption of constant velocity, we find that the material in and around the E and NNE condensations was ejected in the mid-1200s A.D., give or take 50--100 years. With the exception of three small knots to the far south of $\eta$ Car and one to the northwest, the ejecta dating to the mid-1200s are all found to one side of the central star and are blueshifted. We also see evidence of a third, intermediate eruption that occurred in the sixteenth century. Ejecta dating to the mid-1500s are found in the SE arc, the W condensation, and in and around the NW condensation. From proper motions alone, we cannot rule out that this intermediate date peak is the result of newer ejecta from the Great Eruption hitting the older material from the 1200s. However, the radial velocities of these features place them in a different part of three-dimensional space from the thirteenth-century ejecta. The lack of X-ray emission over the SE arc and the features to the far north also indicates a lack of strong interaction between ejecta at those spots. In summary, we have shown with distance-independent measurements that $\eta$ Car erupted at least once, likely twice, before its Great Eruption in the 1800s. Models for this still-enigmatic star must therefore explain the recurrence of these major mass-loss events, along with their \btxt{several-hundred-}year timescale and their various asymmetries. | 16 | 9 | 1609.00362 |
1609 | 1609.07297_arXiv.txt | We report on the evidence of highly blue-shifted resonance lines of the singly ionised isotope of \beviiii ~ in high resolution UVES spectra of Nova Sagittarii 2015 No.\,2 (V5668 Sgr). The resonance doublet lines \beviiii ~ at $\lambda\lambda$313.0583,\,313.1228 nm are clearly detected in several non saturated and partially resolved high velocity components during the evolution of the outburst. The total absorption identified with \beii ~ has an equivalent width much larger than all other elements and comparable to hydrogen. We estimate an atomic fraction $N(\mbox{\bevii})/N(\mbox{Ca})$ $\approx$ 53-69 from unsaturated and resolved absorption components. The detection of $^{7}$Be in several high velocity components shows that $^{7}$Be has been freshly created in a thermonuclear runaway via the reaction $^{3}\mbox{He}(\alpha,\gamma)^{7}\mbox{Be}$ during the Nova explosion, as postulated by \citet{arn75}, however in much larger amounts than predicted by current models. \beviiii ~ decays to \liviiii\ with a half-life of 53.22 days, comparable to the temporal span covered by the observations. The non detection of \liviii\ requires that \livii\ remains ionised throughout our observations. The massive \beii~ ejecta result into a \livii\ production that is $\approx$ 4.7-4.9 dex above the meteoritic abundance. If such a high production is common even in a small fraction ($\approx$5\%) of Novae, they can make all the {\it stellar} \livii ~ of the Milky Way. | \livii ~ is a unique element that shows a large variety of production processes. These include primordial nucleosynthesis, spallation processes by high energy cosmic rays in the interstellar medium, stellar flares in low mass stars, Cameron-Fowler mechanism in Asymptotic Giant Branch (AGB) stars and Novae, and neutrino induced nucleosynthesis in SNae explosions. Observations show that \livii ~ has a constant abundance among metal-poor stars and begins to rise at [Fe/H] $\approx$ -1 to reach the meteoritic value at solar metallicities \citep{reb88} requiring a net \livii~ production \citep{rom99}. The rate of the Li increase favours AGB stars and Novae as the most significant $\it stellar$ sources. Although \livii ~ has been observed in AGB stars the observational evidence for Novae has only recently been found by \citet{izz15} with the first detection of the \liviii ~ $\lambda\lambda$6708 line in the spectra of Nova Centauri 2013 (V1369 Cen) and by \citet{taj15} with the first detection of $^{7}$Be in the post-outburst spectra of the classical Nova Delphini 2013 (V339 Del). Here, we report a study of the \beii ~ by means of UVES observations of Nova Sagitarii 2015 No.\,2 (V5668 Sgr). A spectrum from the High Dispersion Spectrograph of the Subaru Telescope taken at day 63 after maximum has been discussed by \citet{taj16} who reported the presence of \beviiii ~ in this Nova, and also in V2944 Oph. | We have analysed UVES high resolution observations of V5668 covering six outburst phases from day 58 to day 89 from maximum. The evolution of the absorption offers clear evidence in support of the identification of \beviiii ~ by \citet{taj16}. In particular, the weakening of the \beviiii\ absorptions at a late epoch shows that the iron-peak species are a minor contaminant. By means of unsaturated \beii ~ components we derived an abundance of $N(\mbox{\bevii})/N(\mbox{Ca})$ $\approx$ 53-69 when the $^{7}$Be decay is taken into account. Assuming all the \bevii~ goes into \livii~ this corresponds to a \livii~ overproduction of 4.7 - 4.9 dex over the solar-meteoritic value. We then argue that a rate of 2 yr$^{-1}$ of such events in a Galaxy lifetime, i.e. only a small fraction of all Novae, could be responsible for the production of the whole \livii~ required from {\it stellar} sources. We also notice that such a high \bevii\ production should increase the probability of detecting the 478-keV $\gamma$-ray photons emitted in the \bevii\ to \livii\ reaction which have been so far elusive despite several $\gamma$-ray searches. | 16 | 9 | 1609.07297 |
1609 | 1609.02367_arXiv.txt | We explore the possibility of the formation of globular clusters under ultraviolet (UV) background radiation. One-dimensional spherical symmetric radiation hydrodynamics (RHD) simulations by Hasegawa et al. have demonstrated that the collapse of low-mass ($10^{6-7}\solmass$) gas clouds exposed to intense UV radiation can lead to the formation of compact star clusters like globular clusters (GCs) if gas clouds contract with supersonic infall velocities. However, three-dimensional effects, such as the anisotropy of background radiation and the inhomogeneity in gas clouds, have not been studied so far. In this paper, we perform three-dimensional RHD simulations in a semi-cosmological context, and reconsider the formation of compact star clusters in strong UV radiation fields. As a result, we find that although anisotropic radiation fields bring an elongated shadow of neutral gas, almost spherical compact star clusters can be procreated from a ``supersonic infall'' cloud, since photo-dissociating radiation suppresses the formation of hydrogen molecules in the shadowed regions and the regions are compressed by UV heated ambient gas. The properties of resultant star clusters match those of GCs. On the other hand, in weak UV radiation fields, dark matter-dominated star clusters with low stellar density form due to the self-shielding effect as well as the positive feedback by ionizing photons. Thus, we conclude that the ``supersonic infall'' under a strong UV background is a potential mechanism to form GCs. | According to the concordant cosmology, the formation of low-mass sub-galactic objects are thought to have been the prime mode of the star formation in the early Universe. Considering the fact that stars are born in the form of star clusters in present-day galaxies \citep[e.g.,][]{Lada&Lada2003,Meurer+95,Fall+05}, it is of great importance to explore the formation of star clusters in such sub-galactic objects, to reveal the structure formation history in the Universe. Globular clusters (GCs) are significant tracers of early star formation history, since they are low-metal, oldest star clusters in the Universe. GCs are relatively massive ($10^{4-6}\solmass$) and stellar-dominated systems in which stars are tightly distributed in color-magnitude diagram. Thus, GCs are thought to be of a single stellar population. Their ages can be evaluated by isochrone fitting. Although there are some uncertainties in the distances to GCs, the metallicity, and the stellar evolution models, the typical age is evaluated to be $\gtrsim 10$ Gyr with an uncertainty of $\sim$ Gyr \citep[e.g.,][]{Krauss2003, Dotter+07, VandenBerg+13}. Recently, \citet{PLANCK2016} have reported the reionization redshift as $7.8< z_{\rm r} <8.8$ from the Thomson scattering optical depth of the cosmic microwave background (CMB). Based on the comparison between ages of GCs and the reionization epoch, most of old GCs seem to have formed under the influence of UV background radiation fields after the cosmic reionization. The internal dynamics of GCs is quite distinctive from other systems with comparable luminosities such as dwarf spheroidal galaxies (dSphs). GCs are very compact systems, the half-light radii ($r_{\rm h}$) of which are around 1-10 pc, regardless of their luminosity \citep{McConnachie12}. The velocity dispersions ($\sigma$) of GCs are as high as 10~km/s, and show steep dependence on luminosity ($L$) as $\sigma \propto L^{1/2}$ \citep[e.g.,][]{McLaughlin2000,Drinkwater,Hasegan,Forbes08}, % which is insensitive to their radii and masses. These characteristic features of GCs imply that they formed in their inherent environments. The formation scenarios for GCs have been proposed by many authors, but still under debate. For instance, \cite{Kravtsov&Gnedin} have performed high-resolution cosmological simulations to explore the formation of GCs in a Milky Way (MW)-sized galaxy. They have found that cold metal-poor gas is supplied to the center of the galaxy by direct gas accretion along dark matter (DM) filaments during minor mergers of smaller galaxies. The collisions of accreting gas spawn dense molecular clouds, which may be able to evolve to GCs. Although the spatial resolution of the simulations was not sufficient to resolve the internal structure of each star cluster, their result suggests that GCs possibly form in the cosmological context. The formation of giant molecular clouds can be expected also in major mergers of galaxies. \citet{Saitoh+09} have performed $N$-body/SPH simulations of major mergers to explore the evolution of the interstellar medium (ISM). As a result, they have shown that the formation of GC-sized massive star clusters is triggered at high dense filamentary regions compressed by shocks. Besides, some high-resolution cosmological $N$-body simulations have revealed that the radial distribution of sub-halos originating from relatively rare peaks resembles the distribution of the Galactic GCs \citep{Diemand+05, Moore2006}. This result implies that GCs may stem from DM sub-halos, but GCs are usually observed as stellar-dominated systems. To reconcile this inconsistency, \cite{Saitoh+06} have shown, using a semi-cosmological hydrodynamic simulation, that the tidal force by a host galaxy effectively strips DM halos surrounding the star clusters. However, no previous work has not succeeded in accounting for the characteristic internal properties of GCs. As stated above, the formation of GCs is likely to be intimately related to the UV background radiation. Many observations have shown that cosmic reionization took place around the GC formation epoch. For instance, \citet{ASPC_Umemura+01} have estimated the reionization epoch to be $6< z_{\rm r} <10$, by confronting the radiative transfer simulations on reionization to Ly$\alpha$ absorption systems seen in high-$z$ quasar spectra. Also, \citet{Fan+06} have estimated neutral hydrogen fractions at $z\sim 5-6$ from QSO Ly$\alpha$ absorption lines and concluded that reionization is almost completed by $z\gtrsim 6$. Besides, Gamma-Ray Bursts (GRBs) are also available to probe neutral hydrogen at high redshifts, because of their cosmological distances. \citet{Totani+06} have analyzed the Ly$\alpha$ damping wing in the optical afterglow spectrum of GRB050904 at $z=6.3$, and concluded that a large fraction of intergalactic hydrogen seems to be ionized at $z=6.3$. \citet{Ouchi+10} have investigated the evolution of high-$z$ Ly$\alpha$ luminosity functions, and constrained the neutral hydrogen fraction in the intergalactic space as $f_{\rm HI}<0.2$ at $z=6.6$. UV radiation ionizes gas clouds and heats them up to $T \sim10^4$~K. As a result, the gravitational contraction of clouds is suppressed if their virial temperatures are lower than $\sim 10^4$K. Moreover, UV photons dissociate $\HH$ molecules that are the most important coolant at $T\lesssim 10^4$~K under metal-poor environments in the early Universe. Thus, in order for stars to form in the low-mass gas clouds exposed to UV background radiation, the clouds should be self-shielded from a UV background \citep[][]{Tajiri&Umemura}. \cite{Hasegawa2009} (Hereafter HUK09) have performed spherically symmetric radiation hydrodynamics (RHD) simulations to explore the possibility of the star cluster formation under UV background radiation. As a result, they have found that the star cluster formation processes branch off into three paths according to the timing of the self-shielding. If the self-shielding occurs in the stage of supersonic contraction of a cloud, it leads to the formation of very compact star clusters like GCs. (The details of physical processes are described in \S~\ref{sec:physical_model}). However, in the simulations by HUK09, only isotropic irradiation of UV was investigated. In realistic situations, we should consider three-dimensional effects. First, background radiation fields are usually expected to be anisotropic. Under an anisotropic UV background, the self-shielded regions also become anisotropic. Hence, the contraction of clouds is thought to proceed in a different fashion from the spherical symmetric collapse. Furthermore, if the density distributions in clouds are inhomogeneous, the self-shielding is subject to shadowing effects. In the context of the cosmic reionization, \citet{Nakamoto+01} have shown, by six-dimensional radiative transfer simulations, that the reionization process in an inhomogeneous media is considerably delayed compared to a homogeneous medium case due to the shadowing effects. Such inhomogeneity also increases an effective recombination rate in gas clouds, since the local recombination rate is proportional to the square of density \citep{Madau+1999}. These three-dimensional radiation hydrodynamic effects may bring significant impacts on the star formation in the early Universe. In this paper, we perform three-dimensional RHD (3D-RHD) simulations, where the six-dimensional radiative transfer is coupled with 3D hydrodynamics, and investigate how the three-dimensional effects have impacts on the formation processes of star clusters under UV background radiation. This paper is organized as follows. In Section 2, the physical models of star cluster formation are described based on HUK09. Section 3 is devoted to the numerical method of the present study. The numerical results are presented in Section 4, where the evolution of gas clouds exposed to external UV radiation and resultant stellar dynamics are shown. Also, we compare the properties of simulated star clusters to those of velocity dispersion-supported systems such as GCs, dSphs, and ultra compact dwarfs (UCDs). Finally, we discuss and conclude our results in Section 5 and 6, respectively. Throughout this paper, we assume a CDM cosmology neglecting the dark energy, since it is less important in the early Universe. We work with cosmological parameters; $\Omega_{\rm M} = 1.0$, $h=0.6777$, and $\Omega_{\rm b} = 0.1564$ \citep{Planck14}. | \label{sec:dis} \subsection{Formation Sites of Globular Clusters} \begin{table} \begin{center} \caption{Ionizing photon flux required for supersonic infall star formation} \label{table:UVbackground} \begin{tabular}{cccc} \hline \hline $z_{\rm c}$ & $z_{\rm UV}$ & $M_{\rm ini} $ & $F_{\rm ion} $ \\ & & [$10^6M_\odot$] & [${\rm photons}\: {\rm cm^{-2}}\: {\rm s^{-1}}$] \\ \hline 6 & 6.8 & 2.5 & $7.3\times 10^8$ \\ 6 & 6.9 & 5.0 & $7.6\times 10^8$ \\ 9 & 10.3 & 2.5 & $1.8 \times 10^9$ \\ 9 & 10.5 & 5.0 & $4.5 \times 10^8$ \\ 9 & 10.5 & 10.0 & $1.4 \times 10^9$ \\ 12 & 13.8 & 2.5 & $3.3 \times 10^9$ \\ 12 & 14.0 & 5.0 & $2.2 \times 10^9$ \\ 12 & 14.0 & 10.0 & $4.0 \times 10^9$ \\ \hline \end{tabular} \end{center} \end{table} As we have seen in the previous sections, strong UV background radiation is one of the essential conditions to produce GC-like compact star clusters. In addition, we can recognize in Figs.~\ref{fig:Rh-Mv}-\ref{fig:sigma-Mv} that the timescale $t_{\rm rise}$ of UV intensity rise is a significant factor for the cluster formation. If $t_{\rm rise} \lesssim$10 Myr, then the background UV intensity reaches the maximum value before the cloud undergoes the extensive star formation. As a result, the evolution of gas clouds differs little from the case of constant UV background, resulting in the ``supersonic infall''. On the other hand, if the rise of the background UV is as slow as $t_{\rm rise} >$10 Myr, the ``prompt star formation'' proceeds instead of ``supersonic infall'', since the self-shielding is effective in an early phase of contraction. Consequently, the results deviate from ``supersonic infall'' , as shown by gray diamonds in Figs.~\ref{fig:Rh-Mv}-\ref{fig:sigma-Mv}. Therefore, the rise of UV radiation should be faster than the cloud contraction to form GCs. We argue the formation sites of GCs from viewpoints of UV radiation intensity and its variation timescale. For the purpose, we firstly evaluate the photon number flux required. We define $F_{\rm ion}$ as \begin{equation} F_{\rm ion} \equiv \frac{\dot{N}_{\rm ion}}{\pi r_{\rm UV,in}^2}, \label{F_ion} \end{equation} where $r_{\rm UV,in}$ denotes the radius of a cloud at the irradiation epoch. Although this estimation is higher by a factor of 2-4 than the flux we actually assumed in the simulations, we make an order estimation here with this evaluation. We summarize the evaluated fluxes using Eq. (\ref{F_ion}) in Table \ref{table:UVbackground}. This shows that the required ionizing photon number flux is of the order of $\sim 10^9~{\rm photons}~{\rm cm^{-2}}~{\rm s^{-1}}$, which roughly corresponds to $J_{21}\sim 100-1000$, where $J_{21}$ is the mean intensity at the hydrogen Lyman limit frequency in units of $10^{-21}$ erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$ sr$^{-1}$. This value seems to be much higher than $J_{21}$ expected for the global UV background radiation during the epoch of reionization. Thus, we consider the possibilities of local sources. The first possibility is Population III (Pop III) stars. The $\Lambda$CDM cosmology predicts that Pop III stars form in low-mass mini-halos with the masses of $\sim 10^{5-6}\solmass$, which collapse typically at $z\sim$10-30 \citep[e.g.,][]{Tegmark+97,Yoshida+03}. Although the initial mass spectrum of Pop III stars is still controversial, several theoretical studies have shown that Pop III stars are typically massive as $\sim100\solmass$ \citep[e.g.,][]{Nakamura&Umemura01,Susa+14,Hirano+14,Hirano+15}. Therefore, strong UV radiation can be expected in the vicinity of a Pop III halo. If we assume a Pop III star with the mass of $100-1000 \solmass$ and the ionizing photon emissivity of $10^{50-51}~\per{s}{1}$ \citep{Schaerer02}, the ionizing photon number flux is $\sim10^{8-9}~\per{cm}{2}~\per{s}{1}$ at 100~pc, which roughly corresponds to the virial radius of a mini-halo. Thus, the ionizing photon number flux to allow the formation of compact star clusters can be easily accomplished if a Pop III star as massive as $>100\solmass$ forms at $\approx 100$pc from a collapsing cloud. Also, the Kelvin-Helmholtz timescale of a Pop III star with $\approx 100\solmass$ is $\sim 10^5$yr \citep{O'Shea&Norman07}, and therefore the star reaches the main sequence faster than the cloud contraction. However, the lifetime of a Pop III star is a few $10^6$yr. Thus, the formation of GCs by Pop III radiation is realized only for clouds contracting within $10^6$yr. We note that the ionizing photon number flux possibly changes with time according to the stellar motion, if the Pop III star formation takes place during the hierarchical merging process \citep[e.g.,][]{Johnson+08}. The variation timescale of UV radiation is thought to be roughly the infall timescale in the GC-host halo, which is $\sim 100~$Myr. Since this timescale is longer than the cloud contraction time, the variation of UV radiation due to the virial motion does not affect the cloud evolution. The second possibility is young star-forming galaxies, e.g., Lyman $\alpha$ emitters (LAEs). \citet{Wise&Cen09} have numerically simulated the high-$z$ young dwarf galaxies and traced the star formation histories. They have shown that the starburst rises up within a few times 10 Myr and the burst-phase continues for $\sim100~$ Myr, if the virial masses of the halo are as massive as $10^9~\solmass$. Although the SFR varies with the timescale of $\leq 10~$Myr, the luminosity changes are within a factor of three. Hence, the timescale condition for the formation of compact star clusters is likely to be satisfied. Recently, \citet{Yajima+14} have calculated the emissivities of ionizing photons of young star-forming galaxies. According to their result, the ionizing photon number emissivities of the galaxies at $z>6$ correspond to $\sim 10^{52-53}\: \per{s}{1}$,which is translated into the ionizing photon number flux of $\sim 10^{9} ~\per{cm}{2}~\per{s}{1}$ at 1 kpc from the galactic center. Thus, if the star forming regions in LAEs are as compact as $\sim 1$~kpc, compact star clusters may form in sub-halos of the LAEs. The third possibility is active galactic nuclei (AGNs). Recent studies have pointed out the possibility that high-$z$ quasars and faint AGNs bring large contribution to cosmic reionization \citep{Glikman+11,Giallongo+15,Madau&Haardt15,Yoshiura16}. Therefore, it is reasonable to consider UV radiation from AGNs. As for faint AGNs, their typical luminosity is $10^{43}$ erg/s in the range of 2-10 keV \citep{Giallongo+15}. If we assume a simple power-law of the spectrum energy distribution as $L_\nu \propto \nu^{-1}$, the ionizing photon number emitted by the AGN $\dot{N}_{\rm ion}$ is roughly estimated as $\sim 10^{53}\;\per{s} {1}$. Thus, even a faint AGN provides the ionizing photon number flux at 1~kpc away as $\gtrsim 10^9 \: {\rm photons}\: \per{cm}{2}\: \per{s}{1}$. Several authors have argued that the duty cycle of the AGN activity is in the timescale of $10^{8}~$Myr \citep[e.g.,][]{Haehneit+98}, which is comparable to the Eddington timescale. If the mass accretion on to a central black hole is driven by a nuclear starburst, the accretion timescale can be as short as $10^7$yr \citep{Umemura+97}. If this gives the rise time of luminosity, then the situation is favorable for the ``supersonic infall''. Therefore, we can expect the formation of compact star clusters, if the rise time of the AGN luminosity is shorter than 10~Myr. \subsection{Effect of Tidal Field} As shown in \S~\ref{sec:comparison}, low-mass ($\lesssim 10^5~\solmass$) GCs are not formed in our simulations. Here, we assess the effect of tidal stripping by host galaxies, which might work so as to reduce the masses of GCs. Assuming a host galaxy as a point-mass for simplicity, the tidal radius $r_{\rm t}$ of a star cluster orbiting a host galaxy is roughly given by \begin{equation} \label{eq:rt_difinition} \frac{Gm(r<r_{\rm t})}{r_{\rm t}^2} \sim 2\frac{GM_{\rm gal}m(r<r_{\rm t})r_{\rm t}}{r_{\rm gal}^3}, \end{equation} where $m(r<r_{\rm t})$, $M_{\rm gal}$, and $r_{\rm gal}$ denote the cumulative cluster mass within the tidal radius $r_{\rm t}$, the host galaxy mass, and the distance from the galactic center to the cluster, respectively. Supposing a high-$z$ low-mass galaxy of $M_{\rm gal} = 10^9\solmass$ and $r_{\rm gal}=$ 0.3-1 kpc and using the simulated mass profiles (Fig.~\ref{fig:cumulative_mass}), the tidal radii are estimated to be a few $\times$ 10 pc to $\sim$ 100 pc. This estimation implies that the star-dominant parts of the compact star clusters likely to gradually lose their masses in the tidal fields according as the two-body relaxation proceeds, while the diffuse dark matter components would be totally stripped away as shown by \citet{Saitoh+06}. Furthermore, the variety of the orbits of star clusters possibly leads to the variety of mass-loss rates of the clusters. Hence, it seems important to take the tidal stripping into consideration for more quantitative comparison between simulations and observations. \subsection{Internal feedback processes} Throughout this paper, we have concentrated on the impacts of the external background radiation but neglected internal feedback processes. Actually, stars formed in self-shielding regions are expected to emit UV radiation, which ionizes the self-shielded regions internally. Besides, type II supernova (SN) explosions pose dynamical impacts on the gas in the star forming regions. These internal feedbacks may play a significant role to regulate the subsequent star formation \citep[e.g.,][]{Kitayama+04,Kitayama&Yoshida05,Hasegawa&Semelin13}. \citet{Kitayama+04} have explored the impact of the internal UV radiation feedback by a massive Pop III star in a low-mass halo with $10^6\solmass$, and found that the feedback reduces the ambient gas density by photo-evaporation and suppresses the subsequent star formation. In our simulations, the mass resolution is $\sim 10^3\solmass$ (\S~\ref{sec:method}), which correspond to the mass of stars formed simultaneously. Then, the emitted ionizing photon number is evaluated as $\sim 10^{50}~\per{s}{1}$ by utilizing STARBURST99 \citep{STARBURST99} assuming an instantaneous starburst model for $Z/Z_{\odot} = 0.02$ and the Salpeter IMF. Therefore, the argument by \cite{Kitayama+04} is partially applicable to our simulations, and the subsequent star formation in the gas clouds is expected to be suppressed by the internal UV feedback. However, in the present situation, UV radiation from stars surrounding the cloud center might positively work to compress the central star-forming region. Since such a complicated behavior is expected, it is hard to assess how much the internal UV feedback quantitatively affects our results, before the internal UV feedback is actually incorporated. On the other hand, we expect that SN feedbacks are unlikely to affect on the formation of compact star clusters, since the star formation is quickly quenched by $\sim 5$Myr (see Fig.~\ref{fig:SFR}). | 16 | 9 | 1609.02367 |
1609 | 1609.00538_arXiv.txt | Cadmium-Zinc-Telluride Imager (CZTI) is one of the five payloads on-board recently launched Indian astronomy satellite {\astrosat}. CZTI is primarily designed for simultaneous hard X-ray imaging and spectroscopy of celestial X-ray sources. It employs the technique of coded mask imaging for measuring spectra in the energy range of 20 - 150 keV. It was the first scientific payload of {\astrosat} to be switched on after one week of the launch and was made operational during the subsequent week. Here we present preliminary results from the performance verification phase observations and discuss the in-orbit performance of CZTI. | \label{intro} {\astrosat} is India's first dedicated satellite mission for multi-wavelength Astronomy~\cite{singh_astrosat}. It was launched on September 28, 2015 by Polar Satellite Launch Vehicle PSLV C-30 of the Indian Space Research Organization (ISRO) into 650 km orbit with inclination of 6 degrees. {\astrosat} carries five instruments: (i) Large Area Xenon Proportional Counter (LAXPC), (ii) Soft X-ray Telescope (SXT), (iii) Cadmium Zinc Telluride Imager (CZTI), (iv) Ultra-violet Imaging Telescope (UVIT) and (v) Scanning Sky Monitor (SSM). The first four instruments are co-aligned to provide simultaneous multi-wavelength observations of astrophysical sources. The SSM is mounted on a rotating platform and scans the sky for detecting new transient X-ray sources as well as to monitor the variability of known X-ray sources. {\astrosat} also carries an auxiliary instrument - Charge Particle Monitor (CPM) to detect the presence of South Atlantic Anomaly (SAA). Figure~\ref{astrosat_deployed} shows the location of all instruments of {\astrosat} in deployed condition. \begin{figure} \begin{center} \includegraphics[width=0.7\textwidth]{astrosat} \caption{ \label{astrosat_deployed} Deployed configuration of {\astrosat} depicting the locations of various scientific instruments (image credit ISRO).} \end{center} \end{figure} All instruments on-board {\astrosat} as well as all satellite subsystems are functioning normally since the launch. The period of first six months after launch was designated as performance verification (PV) phase, which was successfully completed in March 2016. The observations during the PV phase were used to carefully characterize all instruments and to obtain actually realized values for various instrument parameters. The period of next six months (i.e. till September 2016) is reserved for the scientific observations by the instrument teams, which is successfully progressing at present. From October 2016 onwards, {\astrosat} will be available for guest observations, initially for the Indian astronomers for one year, and subsequently open to all. During this period, {\astrosat} will be operated as a proposal driven observatory. The nominal lifetime of {\astrosat} is planned to be five years. Thus it will provide a great opportunity to carry out observations of a variety of astrophysical sources in wavebands ranging from visible to hard X-rays. Here we present the results of in-orbit characterization of the CZTI instrument carried out with the observations during the PV phase. Section~\ref{inst_desc} provides the detailed description of the instrument, a brief description of pre-flight calibration is presented in section~\ref{calib}, section~\ref{inorbit} discusses the in-orbit operations and in section~\ref{results} preliminary results from CZT Imager are presented. | The Cadmium-Zinc-Telluride Imager (CZTI) on-board {\astrosat} is designed for simultaneous hard X-ray imaging and spectroscopy in energy range of 20 - 150 keV. The full instrument, in terms of both hardware and software, is functioning exactly as designed and is performing very satisfactorily. The enhanced capabilities of CZTI i.e. hard X-ray monitoring and polarimetry in the extended energy range of 100 - 300 keV, are also functioning as per expectation. Detailed characterization of the CZTI imaging, spectroscopic and timing performance has been carried out during the first six months of performance verification observations. During the initial observations, it was found that CZTI is highly sensitive to the `particle shower' like events due to the large number of pixels, resulting in much larger data volume mostly consisting of events useless for the scientific objectives of CZTI. A remedy for this problem has been implemented in the onboard software as well as in the CZTI data processing pipeline. The full data analysis software, including the calibration database (CALDB) and the higher level data products generation software, is fully functional and is released for wider use/testing. Overall the CZTI instrument is performing as per the expectations. | 16 | 9 | 1609.00538 |
1609 | 1609.00012_arXiv.txt | The total contribution of diffuse halo gas to the galaxy baryon budget strongly depends on its dominant ionization state. In this paper, we address the physical conditions in the highly-ionized circumgalactic medium (CGM) traced by \ovi\ absorption lines observed in COS-Halos spectra. We analyze the observed ionic column densities, absorption-line widths and relative velocities, along with the ratios of \nvovi\ for 39 fitted Voigt profile components of \ion{O}{6}. We compare these quantities with the predictions given by a wide range of ionization models. Photoionization models that include only extragalactic UV background radiation are ruled out; conservatively, the upper limits to \nvovi\ and measurements of N$_{\rm OVI}$ imply unphysically large path lengths $\gtrsim$ 100 kpc. Furthermore, very broad \ovi\ absorption (b $>$ 40 km s$^{-1}$) is a defining characteristic of the CGM of star-forming L$^{*}$ galaxies. We highlight two possible origins for the bulk of the observed \ovi: (1) highly structured gas clouds photoionized primarily by local high-energy sources or (2) gas radiatively cooling on large scales behind a supersonic wind. Approximately 20\% of circumgalactic \ion{O}{6} does not align with any low-ionization state gas within $\pm$50 km s$^{-1}$ and is found only in halos with M$_{\rm halo}$ $<$ 10$^{12}$ M$_{\odot}$. We suggest that this type of unmatched \ion{O}{6} absorption traces the hot corona itself at a characteristic temperature of 10$^{5.5}$ K. We discuss the implications of these very distinct physical origins for the dynamical state, gas cooling rates, and total baryonic content of L$^*$ gaseous halos. | \label{sec:intro} Quasar absorption-line techniques have established that present-day galaxies are enveloped by a highly-ionized and enriched plasma extending to hundreds of kpc. Although largely invisible in its faint emission, this `halo gas' or `circumgalactic medium' (CGM) is revealed by observations of the \ovi\ doublet at $\lambda\lambda 1031,1037$ in far-UV absorption-line spectroscopy of background quasars \citep[e.g.][]{tripp08, wakker09,pwc+11,tumlinson11, savage14}. Dedicated studies have assessed the covering fraction, surface density, and radial extent of this \ovi-bearing gas, including its relationship to galaxy stellar mass and star formation rate. While this \ovi\ is common around $\sim L^*$ star-forming galaxies, it appears less common around non-star-forming galaxies \citep{tumlinson11} and low-mass dwarf galaxies \citep{pwc+11}. It may \citep{tripp00, shull03, chen09, stocke14, johnson15} or may not \citep{wakker09} be common in galaxy-group environments and/or intracluster gas \citep[e.g.][]{bowen01}. Recent observational studies have placed lower limits on the highly-ionized metal mass in the CGM of L$^{*}$ galaxies traced by \ovi. It most likely exceeds 10$^{7}$ M$_{\odot}$ \citep{tumlinson11}, comparable to or greater than the metal mass within their ISM \citep{peeples14}. The chief uncertainty in CGM metal-mass estimates arises from the largely unknown ionization conditions of the gas. Typically, lower limits on oxygen budgets in the highly ionized gas are derived by simply assuming that \ovi\ is at its maximum ionization fraction in typical low-density conditions. However, the true metal budgets vary greatly depending on the model used to explain the ionization, which in turn significantly affects conclusions about the fate and origin of the highly ionized gas around galaxies. Several authors have recently addressed the origin and total baryonic content of the highly ionized gas in the CGM by comparing simulations with observations. Strong feedback processes (from both star-formation and AGN) give rise to significant \ovi\ absorption in the CGM of both cosmological zoom-in and hydrodynamical simulations \citep[e.g.][]{shen13, hummels13, cen13, suresh15, liang16, ford16, oppenheimer16, rahmati16, cen16}. However, the column density of the simulated \ovi\ absorption depends upon an assumed ionization mechanism of the gas. Just as observers must model their absorption-line data with a radiative transfer code like Cloudy (Ferland et al. 2013) to determine the physical characteristics of the gas (e.g. Werk et al. 2014), so must simulators reproduce the physical conditions in their gas cells using radiative transfer models that provide corresponding values of gas column density of a given ion. Under the standard assumption of ionization equilibrium and a combination of photo- and collisional ionization, simulators compare their model-derived radial distributions of \ovi\ absorption with the \ovi\ column density distributions from observations (though see Oppenheimer et al. 2016). These comparisons routinely reveal a deficit of \ovi\ column density in the simulated gas cells compared to observations, perhaps hinting at some unaccounted for ionization process \citep[e.g.][]{suresh15}. It has thus far been difficult to generalize the physical conditions of intergalactic and/or circumgalactic gas bearing \ovi\ despite its frequent detection and broad characterization in UV spectroscopic observations. Detailed studies have been carried out on an absorber-by-absorber basis in high-resolution and high-S/N QSO spectra that trace low-density foreground gas in the IGM and galaxy halos \citep{richter04, sembach04, tripp06, tripp08, howk09, narayanan10, narayanan11, savage11, narayanan12, savage14, muzahid15, pachat16}. These studies have revealed the ionization mechanisms of \ovi\ and other `high-ions' like \nv\ to be both varied and complex over a wide range of environments \citep[e.g.][]{sembach04}. Line diagnostics from low, intermediate, and high ions, including ionic column density ratios and absorption-line profiles, sometimes support a similar, photoionized origin for \ovi, \nv, and low-ionization state gas \citep[e.g.][]{tripp08, muzahid15}, and sometimes require \ovi\ to be ionized by collisions of electrons with ions in a $\sim$10$^{5.5}$ K plasma \citep[e.g.][]{tumlinson05, fox09, savage11,tripp11, wakker12, narayanan12, meiring13}. Often, the multiple components for a single absorber show both narrow and broad absorption lines consistent with both scenarios. A similar challenge arises from the diffuse gas known as high velocity clouds (HVCs) in the halo of the Milky Way \citep{sembach03, fox04,fox05,fox06, lehner09, wakker12}. In this case, the complex mixing and shocking of cooler gas within a hot, ambient medium is often consistent with the absorption-line diagnostics. \begin{figure*}[t!] \begin{centering} \hspace{0.1in} \includegraphics[width=0.91\linewidth]{fig1.pdf} \end{centering} \caption{Ionic species stacks drawn from the {\emph{HST}}/COS absorption-line spectra, centered on the transitions relevant to this study: HI Ly$\alpha$, \ion{N}{2} or \ion{N}{3}, NV, and \ion{O}{6}. We show data for two representative galaxies, J1016$+$4706: 274\_6 at $z=0.252$, and J1330$+$2813: 289\_28 at $z = 0.192$. In each panel, v $=$ 0 corresponds to the transition wavelength at the galaxy systemic redshift. Generally, the associated CGM absorption falls within $\pm 300$ km s$^{-1}$ of the galaxy systemic redshift. In both of these examples, we see strong absorption from HI, the low ionization state transitions of nitrogen, and OVI, but NV is undetected. The light shaded area on each NV panel marks a feature that is not due to NV in either case, but absorption from a system at some other redshift. For example, in the NV panel associated with J1016$+$4706: 274\_6, this absorption line is part of the Lyman Series (\ion{H}{1} $\lambda$= 930 \AA) for an absorption-line system at z $=$ 0.665.} \label{fig:exspec} \end{figure*} \begin{figure*}[t!] \begin{centering} \includegraphics[width=0.75\linewidth, angle =90.0]{fig2.pdf} \end{centering} \caption{Average COS-Halos absorption-line spectra for NII, NIII, SiIII, SiIV, NV, and OVI, stacked at v $=$ 0 km s$^{-1}$ corresponding to the galaxy systemic redshift. Only the COS-Halos star-forming galaxies are included in these stacks. The resultant stacks may include coincident absorption from unrelated gas at different redshifts (or, in the case of \ion{N}{3}~989, a neighboring transition from \ion{Si}{2}), but this is a minor effect. The number of spectra averaged in each panel is shown in the lower-left corners. We use a spectrum if there is coverage of the transition. We note that strong absorption features (SiIII, OVI) are not driven by a small set of events, since a median stack yields qualitatively similar results. } \label{fig:stacks} \end{figure*} Here, we focus on the physical conditions of the \ion{O}{6}-bearing gas around z$\sim$0.2 star-forming $L \approx \mlstar$ galaxies. Previous studies have focused on the Milky Way itself \citep[e.g.][]{fox04, lehner11, wakker12}, on single sightlines with exquisite, high-S/N UV spectra \citep[e.g.][]{narayanan10, narayanan11, tripp11, narayanan12, meiring13}, z$\sim$ 2$-$3 Lyman limit systems or damped Lyman $\alpha$ systems \citep{fox09, lehner14}, or on absorption lines originating in wide variety of galaxy or group environments \citep{heckman02, grimes09, wakker09, bordoloi16}. The COS-Halos dataset provides a uniform sample of absorbers with well-characterized host galaxy properties that allow us to generally constrain ionization processes affecting \ovi\ in the star-forming galaxy halo environment and potentially relate them to galaxy properties. While the COS-Halos spectra have only moderate S/N ($\sim$10 at \ovi ), they cover a wide range of transitions at z$\sim$ 0.2, including (but not limited to) \sii, \siii, \siiv, \nii, \niii, \nv, and \ovi. The coverage of such a variety of ionic species allows us to uniquely assess the multiphase nature of the gas within 150 kpc of a star-forming galaxy for 24 distinct sightlines. Ultimately, our goal is to apply the best available diagnostics from the COS-Halos dataset \citep{tumlinson13, werk13} to the origins of the highly ionized gas in gaseous galactic halos, and from these diagnostics to draw conclusions about the gas flows driving evolution in these galaxies. In Section 2 we describe the data and review the relevant properties of the low-ionization state gas. In Section 3, we present a detailed Voigt profile-based kinematic analysis, joint with the low-ions, of the individual \ion{O}{6} absorbers, and examine several significant correlations between the gas kinematics and COS-Halos host galaxy properties. Section 4 presents an analysis of the many ionization processes capable of producing a highly ionized plasma and compares their predictions with the COS-Halos data for component column densities, gas velocities, line-widths, and ratios of \siovi\ and \nvovi. In Section 5, we summarize our results. Finally, in Section 6 we address several recent analytical and phenomenological models and comment on their applicability to our findings and implications for the co-evolution of the galaxy and its CGM. | Conclusions regarding the origin and fate of circumgalactic gas are inextricably linked to the initial assumptions we make about the physical processes that determine its ionization state. In this work we have demonstrated that the kinematics of the highly-ionized gas in addition to gas column density ratios, contain a wealth of information useful for constraining physical models of the CGM. We have examined a number of equilibrium and non-equilibrium ionization models that predict gas characterized by strong \ion{O}{6} absorption. The constraints from the COS-Halos absorption-line column density and kinematics measurements strongly disfavor many of the models considered. At least some fraction of \ion{O}{6} appears to represent halo gas at the virial temperature, while most of the total column of \ion{O}{6} may result from either gas photoionized primarily by local high-energy sources or gas radiatively cooling on large scales behind a multiphase, fast wind. The latter two models imply very different ionization states and physical origins for the gas. Each model has its own set of strengths and weaknesses. Successful models of the CGM must account for: (1) the velocity correspondence between the low-ionization state, photoionized gas and the \ion{O}{6} absorption; (2) \ion{O}{6} column densities $\gtrsim$ 10$^{14}$ cm$^{-2}$ and highly variable line widths, 10 km s$^{-1}$ $<$ $b$ $<$ 100 km s$^{-1}$; and (3) the absence of \ion{O}{6} around non-star-forming galaxies and related tight correspondence between $\novi$ and SFR/R$^2$. For photoionization models to progress, future studies must focus on a detailed treatment of both the origin and long-term survival of these richly structured clouds and the sources of ionizing radiation. A different issue faces the analytic, phenomenological, and hydrodynamical models that generate significant \ion{O}{6} by collisional ionization. These models often provide a highly detailed physical treatment of the gas dynamics, origin, and ionization state, but make limited predictions with respect to the full range of possible observational measurements. Future insight on the origins of \ion{O}{6} will come from comparing detailed observational analyses of gas kinematics with the model predictions for the line-of-sight kinematics. | 16 | 9 | 1609.00012 |
1609 | 1609.07318_arXiv.txt | We study the evolution of G2 in a \textit{Compact Source Scenario}, where G2 is the outflow from a low-mass central star moving on the observed orbit. This is done through 3D AMR simulations of the hydrodynamic interaction of G2 with the surrounding hot accretion flow. A comparison with observations is done by means of mock position-velocity (PV) diagrams. We found that a massive ($\dot{M}_\mathrm{w}=5\times 10^{-7} \;M_{\odot} \; \mathrm{yr^{-1}}$) and slow ($v_\mathrm{w}=50 \;\mathrm{km\; s^{-1}}$) outflow can reproduce G2's properties. A faster outflow ($v_\mathrm{w}=400 \;\mathrm{km\; s^{-1}}$) might also be able to explain the material that seems to follow G2 on the same orbit. | In year 2012, \citet{Gillessen12} discovered a small cloud, later named ``G2'', at few thousands Schwarzschild radii from SgrA*. G2 has both a dust component, visible in the near infrared $L'$ band, and a gaseous component, visible in Br$\gamma$ and other recombination lines. The cloud lies on a very eccentric orbit ($e \approx 0.98$) and reached its pericenter in early 2014, with a distance from SgrA* of $\approx 2400$ Schwarzschild radii ($R_S$). The line emission shows an increasing spatial extent and a broadening in the velocity space, interpreted as tidal stretching of the cloud by the tidal field of SgrA* \citep{Gillessen13a,Gillessen13b}. Observed position-velocity (PV) diagrams and Br$\gamma$ maps also show the presence of a tail (G2t), following G2 on roughly the same orbit \citep{Pfuhl15}. \begin{figure}[h] \begin{center} \includegraphics[scale=0.25]{Fig1.eps} \caption{a) Density distribution for the low velocity model described in the text. The white circles show the outflow reforming after pericenter. b) sketch of the outflow's evolution at late time, for the high velocity model described in the text. c) Simulated PV diagrams for the high velocity model. The black contours mark the observed G2 and G2t.} \label{fig1} \end{center} \end{figure} | 16 | 9 | 1609.07318 |
|
1609 | 1609.00224_arXiv.txt | {In spite of the numerous studies of low-luminosity galaxies in different environments, there is still no consensus about their formation scenario. In particular, a large number of galaxies displaying extremely low-surface brightnesses have been detected in the last year, and the nature of these objects is under discussion. } {In this paper we report the detection of two extended low-surface brightness (LSB) objects ($\mu_{\rm eff_{g'}}\simeq27$ mag) found, in projection, next to NGC\,3193 and in the zone of the Hickson Compact Group (HCG) 44, respectively.} {We analyzed deep, high-quality, GEMINI-GMOS images with ELLIPSE within IRAF in order to obtain their brightness profiles and structural parameters. We also search for the presence of globular clusters (GC) in these fields.} {We have found that, if these LSB galaxies were at the distances of NGC\,3193 and HCG\,44, they would show sizes and luminosities similar to those of the ultra-diffuse galaxies (UDGs) found in the Coma cluster and other associations. In that case, their sizes would be rather larger than those displayed by the Local Group dwarf spheroidal (dSph) galaxies. We have detected a few unresolved sources in the sky zone occupied by these galaxies showing colors and brightnesses typical of blue globular clusters.} {From the comparison of the properties of the galaxies presented in this work, with those of similar objects reported in the literature, we have found that LSB galaxies display sizes covering a quite extended continous range ($r_{\rm eff}\sim 0.3-4.5$ kpc), in contrast to {\it ``normal''} early-type galaxies, which shows $r_{\rm eff}\sim 1.0$ kpc with a low dispersion. This fact might be pointing to different formation processes for both types of galaxies.} | \label{introduccion} Despite the numerous studies of low-mass galaxies performed through observations and numerical simulations, their formation scenario is still hotly debated. As an example, in the case of low-mass early-type galaxies located in rich groups and clusters, several authors have claimed that they might be gas-rich disk galaxies entering the cluster and being transformed through the interaction with the intracluster medium (e.g. \citealp{2009ApJS..182..216K,2014ApJ...786..105J,2015ApJ...799..172T}). Some others, however, based on the strong photometric relations defined by them, propose a scenario of in situ formation within the cluster environment \citep{2013ApJ...772...68S,2015MNRAS.454.2502S}. In addition, there is observational and theoretical evidence that support a third scenario in which low-mass galaxies might arise from the interaction of massive gas-rich galaxies, as bound structures in the filaments or bridges that are formed as a consecuence of the encounters (e.g. \citealp{1996ApJ...462...50H,2013MNRAS.429.1858D,2014MNRAS.440.1458D}). The very faint end of the early-type galaxy population is defined by dwarf spheroidal (dSph) galaxies, which are extended objects displaying extremely low-surface brightnesses and no evidence of star formation. In the last decade, the number of studies of low-surface brightness (LSB) galaxies outside the Local Group (LG) has increased significantly, thanks to the development of detection surveys with amateur telescopes (e.g. \citealp{2016A&A...588A..89J}), and the access to 8-m class telescopes capable to obtain extremely high-quality deep images that allow to follow up the objects identified with small telescopes (e.g. \citealp{2016MNRAS.457L.103R}). The interest in the identification of new examples of such extremely faint galaxies resides in the fact that they can be used as test-beds for constraining models predictions. As an example, $\Lambda$-Cold Dark Matter ($\Lambda$CDM) models produce a larger number of dSph satellites around bright galaxies than observed (the so-called Missing Satellite Problem, e.g., \citealp{1999ApJ...522...82K,1999ApJ...524L..19M}), as well as an isotropic distribution around the brightest galaxies which is not detected in the LG and M\,81 \citep{2013MNRAS.435.1928P,2013AJ....146..126C}. In addition, hierarchical models predict DM dominated dSph galaxies, while faint galaxies of tidal origin would not contain DM at all (\citealp{2013MNRAS.429.1858D} and references therein). HCG\,44 was originally classified as a Hickson Compact Group dominated by an elliptical galaxy (NGC\,3193 or HCG44b), with three additional bright galaxy members: an Sa galaxy (NGC\,3189/3190 or HCG44a), an SBc galaxy (NGC\,3185 or HCG44c) and an Sd galaxy (NGC\,3187 or HCG44d) \citep*{1982ApJ...255..382H,1989ApJS...70..687H}. Later, \citet{1991AJ....101.1957W} identified the dwarf like galaxy [WMv91]\,1015+2203 at the redshift of the group, and \citet*{2010ApJ...710..385B} detected four additional members in their HI study of the group, increasing the HCG\,44 galaxy population to nine members. However, \citet{2001ApJ...546..681T} obtained a surface brightness fluctuation (SBF) distance of 34 Mpc to NGC\,3193 and, more recently, \citet{2013MNRAS.428..370S} have considered [WMv91]\,1015+2203 as a background galaxy due to its radial velocity ($V_{\rm r}=$1940 km s$^{-1}$), larger than that of the group ($V_{\rm r}=$1379 km s$^{-1}$). In addition, in their spectroscopic study of the AGN activity in Hickson Compact Groups, \citet{2010AJ....139.1199M} reconsidered NGC\,3193 as a member of HCG\,44. Therefore, the real galaxy content of HCG\,44 is still under discussion. In this paper we present a photometric study of two low-surface brightness galaxies detected in the zone of HCG\,44. One is located, in projection, within the halo of NGC\,3193, and the other, between the galaxies NGC\,3189/3190 and NGC\,3185. We analyze their structural properties considering the distance modulii reported in the literature for the bright galaxies placed in this region of the sky (Table\,\ref{distancias}), without attempting to clarify the membership status of the latters to the HCG\,44. The paper is organized as follows. In Section\,\ref{observaciones} we describe the photometric data, in Section\,\ref{resultados} we present the results obtained from the analysis of these data, in Section\,\ref{Globulars} we analyze the existence of globular clusters associated with these galaxies, and in Section\,\ref{conclusions}, we present our conclusions. \begin{table*} \caption{Distance information of the bright galaxies present in the region of the HCG\,44.} \label{distancias} \centering \begin{tabular}{lccccccc} \hline\hline \multicolumn{1}{c}{Galaxy} & \multicolumn{1}{c}{R.A.} & \multicolumn{1}{c}{DEC} & \multicolumn{1}{c}{(m-M)} & \multicolumn{1}{c}{Distance} & \multicolumn{1}{c}{Method} & \multicolumn{1}{c}{$\rm H_0$} & \multicolumn{1}{c}{Reference} \\ \multicolumn{1}{c}{} & \multicolumn{1}{c}{(J2000)} &\multicolumn{1}{c}{(J2000)} & \multicolumn{1}{c}{(mag)} & \multicolumn{1}{c}{(Mpc)} & \multicolumn{1}{c}{} & \multicolumn{1}{c}{(km s$^{-1}$ Mpc$^{-1}$)} & \multicolumn{1}{c}{} \\ \hline NGC\,3185 & 10:17:38.5 & 21:41:18 & 32.00 & 25.1 & TF & 74.4 & (1) \\ NGC\,3187 & 10:17:47.8 & 21:52:24 & 32.08 & 26.1 & T est & 75.0 & (2) \\ & & & 32.62 & 33.3 & TF & & (3) \\ NGC\,3189/3190 & 10:18:05.6 & 21:49:56 & 31.73 & 22.2 & SNI & 74.4 & (1) \\ NGC\,3193 & 10:18:24.9 & 21:53:38 & 32.63 & 33.6 & SBF & 74.4 & (1) \\ \hline \end{tabular} \tablefoot{TF: Tully-Fisher; T est: Tully estimation; SNI: Supernova Type I; SBF: Surface Brightness Fluctuations; (1): \citet{2013AJ....146...86T}; (2): Nearby Galaxy Cataloge (1988, Tully R. B.); (3): \citet{2007A&A...465...71T}. } \end{table*} | \label{conclusions} From deep GEMINI-GMOS images, we have detected two extended LSB objects in the region of the HCG\,44. From our photometric analysis we have found the following results: \begin{itemize} \item Both galaxies display smooth morfologies. The sustraction of models of the galaxies built from elliptical isophotes with fixed parameters, left no residuals. This might be evidence of the absence of inner structures, star-forming regions and interactions with the brighter galaxies.\\ \item The colors of these objects are in agreement with those reported for the dSphs studied by \citet{2015A&A...581A..82M}. However, they are much bluer than those of the LSB galaxies detected in M\,83 \citep{2015A&A...583A..79M} and those of the dSph galaxies analyzed by \citet{2012AJ....144..190L} around NGC\,7331. \\ \item If both galaxies were placed at the distances of NGC\,3193 and HCG\,44, the effective radius and luminosities of both objects would be similar to those of the ultra-diffuse galaxies (UDGs) reported in different environments. These sizes would be significantly larger than those of the dSph galaxies of the LG \citep{2012AJ....144....4M} and of the LSB galaxies recently identified in the M\,83 subgroup in Centaurus \citep{2015A&A...583A..79M}. \\ \item If dSph1 were at the distance of NGC\,3193, it would be one of the most extended LSB objects reported up to date.\\ \item If the LSB structure that seems to link dSph2 with NGC\,3189/3190 is real, it would be evidence of a tidal origin for dSph2. \\ \item We have detected several unresolved sources in the frames of dSph1 and dSph2 consistent with being blue GCs, but we have not detected any overdensity that could be associated to any of the two galaxies. However, if dSph1 and dSph2 were low-mass galaxies at the distances of NGC\,3193 and/or HCG\,44, it would not be expected a high number of GCs belonging to them. If the unresolved objects were indeed GCs, and dSph1 and dSph2 were at such distances, dSph1 might contain massive nuclear star clusters like those reported by \citet{2009MNRAS.392..879G} in several dSph galaxies, and the distribution of the GCs of dSph2 would be assymetric like that shown by the IKN dwarf in the M\,81 group \citep{2015A&A...581A..84T}.\\ \item As we are not able to establish the distances to dSph1 and dSph2, it could not be discarded that these objects were isolated LSB galaxies. \end{itemize} According to the {\it Dual Dwarf Galaxy Theorem} postulated by \citet{2012PASA...29..395K}, two types of dwarf galaxies exist: {\it primordial dwarf galaxies}, formed within low-mass dark matter halos, that, as a consequence, are DM dominated; and {\it tidal/ram-pressure dwarfs}, formed in the encounters of galaxies, which are devoided of DM. An interesting general result obtained from Figure\,\ref{mueff_dSphs}, is that LSB galaxies display a quite extended range of effective radius, in comparison with {\it ``normal''} early-type galaxies which show $r_{\rm eff}\sim1.0$ kpc with a low dispersion. This different behaviour might be pointing to different formation scenarios for both types of galaxies. In the context of the discussion about UDGs being failed luminous galaxies or genuine dwarfs, \citep{2015ApJ...798L..45V,2016arXiv160408024B}, if the unresolved objects detected in the frames of dSph1 and dSph2 are indeed GCs belonging to these galaxies, both galaxies would be in agreement with the scenario of \citet{2016arXiv160408024B}. That is, they would be quenched dwarfs, as all detected unresolved sources would be blue GCs, and failed luminous galaxies are expected to present both red and blue GCs and a more rich GC system. If both dSph galaxies belong to HCG\,44, and considering HCG\,44 as composed by four galaxies (NGC\,3185, NGC\,3187, NGC\,3189/90 and [WMv91]\,1015+2203), then both LSB galaxies would be located in a high-density environment, similar to that of the center of rich galaxy clusters. In addition, they would constitute the $\sim33$\% of the total galaxy population of the group, and the $\sim66$\% of the dwarf galaxy population. These fractions would be in agreement with the estimations made by \citet{1996ApJ...462...50H} about the amount of dwarf galaxies of tidal origin expected in Hickson Compact Groups. We hope that the obtention of spectroscopic data of both dSph galaxies and their GC candidates in the near future, help to disentangle the real nature of all these objects. | 16 | 9 | 1609.00224 |
Subsets and Splits
No community queries yet
The top public SQL queries from the community will appear here once available.