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1607 | 1607.02577_arXiv.txt | Intermediate polars with known rates of spin period changes are not numerous because such tasks require measurements performed for a long time. To measure a spin period change, MU~Cam is a good candidate because it has a spin oscillation with a large amplitude enabling measurements with high precision. Fortunately, in the past the spin period of MU~Cam was measured with high precision. To measure the spin period anew, in 2014--2015 we performed extensive photometric observations of MU~Cam, spanning a total duration of 208~h within 46 nights. We found that the spin, sideband and orbital periods are equal to $1187.16245\pm0.00047$~s, $1276.3424\pm0.0022$~s and $4.71942\pm0.00016$~h, respectively. Comparing the measured spin period with the spin period of MU~Cam in the past, we detected the spin period change with ${\rm d}P/{\rm d}t=-(2.17\pm0.10)\times 10^{-10}$. This rate of the spin period change was not stable and varied in a time scale of years. During four nights in 2014 April--May MU~Cam was fainter than usual by 0.8 mag, and the amplitude of the sideband oscillation was five times larger, denoting significant fraction of disc-overflow accretion. The sideband oscillation showed a double-peaked pulse profile in the normal brightness state. When the star brightness was decreased by 0.8 mag, the sideband oscillation showed a single-peaked pulse profile. In contrast, the spin pulse, which was quasi-sinusoidal, remained remarkably stable both in profile and in amplitude. Moreover, the spin pulse was also remarkably stable in a time scale of years and even decades. MU~Cam is of great interest because it represents a distinctive object with a large and unstable rate of the spin period change and exhibits a distinctive behaviour of the pulse profiles. | Magnetic cataclysmic variables, polars and intermediate polars (IPs), are interacting binary stars, in which accretion occurs onto a magnetic white dwarf. In contrast with polars, intermediate polars (IPs) contain a magnetic white dwarf that spins strongly non-synchronously with the orbital period of the system. Because the magnetic axis does not coincide with the spin axis of the white dwarf, this causes an oscillation with the spin period, which can be observable in optical light, in X-rays and in polarimetry. One more periodic oscillation can appear with the beat period, $1/P_{\rm beat} = 1/P_{\rm spin} - 1/P_{\rm orb}$. The oscillation with the beat period is called the orbital sideband. A natural reason for the sideband oscillation, when it occurs in optical light, is the reprocessing of X-rays at some part of the system that rotates with the orbital period. A review of IPs is given in \citet{Patterson94}. The usual criteria for IP classification are optical and X-ray spin oscillations. These oscillations must show a high degree of coherence to distinguish them from quasi-periodic oscillations. Especially it is important for optical oscillations because X-ray observations by their nature cannot be very long and because quasi-periodic oscillations in X-rays are less probable. Many theoretical works make assumption that IPs are in spin equilibrium. This, however, is questionable. Therefore long-term tracking of the spin period is an important task because allows an observational test of spin equilibrium due to alternating spin-up and spin-down \citep{Patterson94}. When an IP is not in spin equilibrium, the rate of the spin period change gives understanding of the angular momentum flows within the system \citep{King99}. MU Cam was identified as an intermediate polar by \citet{Araujo03} and \citet{Staude03}. Optical photometry showed variability with two periods, which were tentatively identified with the orbital period of the system and the spin period of the white dwarf, $P_{\rm orb}= 4.719\pm0.006$~h and $P_{\rm spin}=1187.246\pm0.004$~s, respectively \citep{Staude03}. The latter period was the dominant signal in the hard X-rays. This proves that this period is indeed the spin period of the white dwarf but not the sideband period \citep{Staude08}. \citet{Yun11} analysed times of maxima of the spin oscillation of MU~Cam, which were presented in \citet{Araujo03}, \citet{Staude03}, \citet{Kim05} and their own times of maxima obtained in 2005--2006. The (O--C) diagram revealed the significant variation of the spin period of MU~Cam. Although, from this diagram, the behaviour of the spin period in the future was unclear, \citeauthor{Yun11} reported ${\rm d}P/{\rm d}t \approx -4.08 \times 10^{-8}$ (see their page 11). This extremely large ${\rm d}P/{\rm d}t$ seems entirely impossible because an oscillation with such ${\rm d}P/{\rm d}t$ must reveal a low degree of coherence. None the less, the (O--C) diagram presented in Yun et al. suggests a measurable change of the spin period. Therefore we can attempt to find the spin period change by using high-precision measurements of the spin period obtained in different moments, which are separated by large time intervals. Fortunately, in 2002 \citet{Staude03} measured the spin period with high precision. To measure the spin period in 2014--2015, we performed extensive photometric observations of MU~Cam. In addition, these observations allowed us to find out interesting details of the behaviour of the spin pulse profile and of the sideband pulse profile. In this paper we present results of our extensive photometric observations, which have a total duration of 208~h and cover 15 months. | We performed extensive photometric observations of MU~Cam over 46 nights in 2014 and 2015 and obtained the following results: \begin{enumerate} \item Due to the large observational coverage and low relative noise level, we evaluated the spin period of the white dwarf with high precision. The spin period is equal to $1187.16245\pm0.00047$~s. The semi-amplitude of the spin oscillation is 90~mmag. \item Comparing this spin period and the spin period, which was found by \citet{Staude03} 12 years ago, we discovered that the spin period of MU~Cam decreases with ${\rm d}P/{\rm d}t=-(2.17\pm0.10)\times 10^{-10}$. The rate of the spin period change is not stable and fluctuates in a time scale of years. \item The pulse profile of the spin oscillation is asymmetric with a slow rise and a rapid decline. This pulse is remarkably stable both in profile and in amplitude in time scales of years and even decades. \item The profile and amplitude of the spin pulse remained practically unchanged when MU~Cam temporarily decreased its brightness by 0.8 mag. This means that the accretion regions emitting optical light do not change their structure when their light intensity is changed by two times. It also means that the light intensity of the disc and the light intensity of these regions change in equal proportion. \item In addition to the spin oscillation, we detected the sideband oscillation with a period of $1276.3424\pm0.0022$~s and with an average semi-amplitude of 25~mmag. \item Generally, the amplitude of the sideband oscillation was 3.5 times less than the amplitude of the spin oscillation. During the four latter observing nights of 2014 the brightness of MU~Cam was decreased by 0.8 mag, and the amplitude of the sideband oscillation was increased by five times, denoting that during these four nights we observed MU~Cam in the state of significant fraction of disc-overflow accretion. \item During the normal brightness of MU~Cam, the sideband oscillation showed a double-peaked pulse profile. During the low brightness of MU~Cam, which denotes significant fraction of disc-overflow accretion, the sideband oscillation showed a single-peaked pulse profile. \end{enumerate} | 16 | 7 | 1607.02577 |
1607 | 1607.01024_arXiv.txt | We study the linear and non-linear bias parameters which determine the mapping between the distributions of galaxies and the full matter density fields, comparing different measurements and predictions. Associating galaxies with dark matter haloes in the MICE Grand Challenge N-body simulation we directly measure the bias parameters by comparing the smoothed density fluctuations of haloes and matter in the same region at different positions as a function of smoothing scale. Alternatively we measure the bias parameters by matching the probability distributions of halo and matter density fluctuations, which can be applied to observations. These direct bias measurements are compared to corresponding measurements from two-point and different third-order correlations, as well as predictions from the peak-background model, which we presented in previous articles using the same data. We find an overall variation of the linear bias measurements and predictions of $\sim 5 \%$ with respect to results from two-point correlations for different halo samples with masses between $\sim 10^{12} - 10^{15}$ $h^{-1}M_\odot$ at the redshifts $z=0.0$ and $0.5$. Variations between the second- and third-order bias parameters from the different methods show larger variations, but with consistent trends in mass and redshift. The various bias measurements reveal a tight relation between the linear and the quadratic bias parameters, which is consistent with results from the literature based on simulations with different cosmologies. Such a universal relation might improve constraints on cosmological models, derived from second-order clustering statistics at small scales or higher-order clustering statistics. | \label{sec:introduction} The increase of data from upcoming and next generation of galaxy surveys is pulling down errors on the observed statistics of the large-scale galaxy distribution. Thus, the inferences on cosmological models from these statistics requires a modelling of cosmological fluids and their statistical properties with an accuracy of at least the same order of magnitude as the observational errors. One of the largest uncertainties comes from the modelling of the mapping between the observed fluctuations of galaxies to the fluctuation of the underlying matter distribution (hereafter referred to as $\delta_g$ and $\delta_m$ respectively). These fluctuations are defined as normalised deviations of the density $\rho$, smoothed typically with a top-hat window of characteristic scale $R$, from the mean density of the universe $\bar \rho$ at the position $\bf r$, \begin{equation} \delta(\bf r) \equiv \frac{\rho(\bf r) - \bar \rho}{\bar \rho}. \label{eq:delta} \end{equation} The mapping from $\delta_g$ to $\delta_m$ is described by the so-called bias function, $\delta_g=F[\delta_{m}, \nabla_{ij}\Phi_v]$, where $\nabla_{ij}\Phi_v$ are second-order derivatives of the velocity potential. The latter relate $\delta_g$ to the matter distribution beyond the smoothing scale $R$ and are therefore referred as {\it non-local} contributions to the bias model \citep{chan12, baldauf12}. For sufficiently large smoothing scales such non-local contributions may be negligible, which allows for a local description of biasing. Thus, the bias function can be modelled as a Taylor expansion in terms of the matter fluctuations \citep{FG} \begin{equation} \delta_h=F[\delta_{m}] \simeq \sum_{i=0}^{N}\frac{b_i}{i!}\delta_{m}^i, \label{eq:biasmodel_local} \end{equation} where $b_i$ are the bias parameters. However, for smaller smoothing scales, which are typically used for studying the two- and three-point statistics of the galaxy distribution, non-local contributions need to be considered. When fluctuations in the density and velocity potential of the matter distribution are sufficiently small the non-local bias model may be described by its second-order expansion, which is hereafter referred to as non-local quadratic bias model \begin{equation} \delta_g = b_1 \biggl\{ \delta_m + \frac{1}{2}[ c_2(\delta_m^2 - \langle \delta_m^2 \rangle) + g_2 \mathcal{G}_2] \biggr\}, \label{eq:biasmodel_quadratic_nonlocal} \end{equation} where $b_1$ and $c_2 \equiv b_2/b_1$ are referred to as linear and quadratic (or second-order) bias parameters respectively. The non-local contribution to the bias function consists in that case of the product of the second-order non-local bias parameter $g_2$ and the smoothed second-order Gallileon \begin{equation} \mathcal{ G}_2({\bf r})= \int \beta_{12}\theta_v({\bf k}_1) \theta_v({\bf k}_2) \ \hat W[k_{12}R]e^{i {\bf k}_{12}\cdot {\bf r} }d^3 {\bf k}_1 d^3 {\bf k}_2, \label{eq:G2} \end{equation} where ${\bf k}_i$ and ${\bf k}_{12} \equiv {\bf k}_2 - {\bf k}_1$ are wave vectors of density oscillations, $\beta_{12} \equiv 1 - ({\hat{\bf k}}_1 \cdot {\hat{\bf k}}_2)^2$ represents the mode-coupling between density oscillations which describes tidal forces, $\theta_v \equiv \nabla^2 \Phi_v$ is the divergence of the normalised velocity field (${\bf v}/\mathcal{H}/f$) and $ \hat W[k_{12}R]$ is the window function in Fourier space. Note that in the case $g_2=0$ equation (\ref{eq:biasmodel_quadratic_nonlocal}) corresponds to the local quadratic bias model. The bias parameters are highly relevant for constraining cosmological models via the growth of matter fluctuations, derived from second-order statistics of the observed galaxy distribution. In particular at large scales, the linear bias factor $b_1$ is completely degenerate with the linear growth factor. Hence, growth-independent measurements of $b_1$ can strongly tighten cosmological constraints from galaxy surveys. Third-order statistics probes the linear and quadratic bias parameters independently of the growth and can be used to break the growth-bias degeneracy in the second-order statistics. Furthermore, the second-order bias measurements from third-order statistics allow growth measurements from second-order statistics at small scales, where non-linear and non-local terms contribute significantly to the signal. Such a small scale analysis would strongly improve the cosmological constraining power of galaxy surveys. However, the value of combining second- with third-order statistics for constraining cosmological models with high precision arises from a detailed understanding of how exactly the bias parameters enter these statistics at different scales, redshifts and for different samples of tracers. We therefore investigated in previous works \citep[][where the latter is hereafter referred to as \paperII]{HBG15-1, BHG15} how accurately the linear bias can be measured from different third-order statistics. These studies were based on the large cosmological MICE Grand Challenge (hereafter referred to as MICE-GC) N-body simulation in which haloes were detected as tracers of the cosmic web and associated with galaxies. The fact that the dark matter distribution is accessible in simulations allows for reliable measurements of the linear bias via second-order statistics, which can then be used as a reference for validating linear bias measurements from third-order statistics as well as theory predictions. Note that a reliable reference for validating higher order bias is currently only provided by running {\it separate universe} N-body simulations \citep{Wagner2015, Lazeyras15}, which is beyond the scope of this article. Alternatively to growth-independent bias measurements from third-order statistics one can employ bias predictions from the peak-background split model for breaking the growth-bias degeneracy. In \citet[][hereafter referred to \paperIII]{HBG15-2} we tested linear bias predictions in the MICE-GC simulations, confirming reports on inaccuracies of these predictions in the literature on a larger mass range thanks to the large volume and resolution of the MICE-GC simulation. As in case of the third-order statistics no reliable measurements for validating the predicted non-linear bias parameters were available. However, an interesting outcome of this analysis was a simple universal relation between the linear and non-linear bias parameters in the peak-background split model, which is independent of redshift and cosmology for halo samples with $b_1 \gtrsim 2$. In the present work we conclude our series of articles, presenting results from an other method for allowing to measure bias parameters in the MICE-GC simulation. This method is based on a direct comparison of the halo and matter density fluctuations and may therefore be seen as the most direct way of measuring bias parameters. We also study a variant of this method, which is based on abundance matching of tracers and matter density contrasts. The interest of this alternative method is that it can be applied in galaxy surveys to estimate to bias function. Both methods deliver linear, as well as non-linear bias measurements. We compare these new measurements with the most reliable results from our previous work (paper I and II, see Table \ref{table:references}), which allows for the validation of how well bias can be measured with each approach. The different measurements of linear and non-linear bias parameters furthermore allow for a validation of the universal relation between the bias parameters, which we found for peak-background split predictions in \paperIII. \rev{The strength of our comparison emerges from the large halo mass range of roughly $10^{12}-10^{15}$ \Msun, probed by the MICE-GC simulation as well as the fact that we use the same halo mass samples throughout the whole comparison project.} The remainder of this article is organised as follows. In Section \ref{sec:simulation} we describe the MICE-GC simulation as well as the halo samples on which our analysis is based on. The different methods for obtaining the bias parameters are briefly reviewed in Section \ref{sec:bias_estimators}, while details on bias measurements from the comparison of density contrasts are given in the appendix. We present our results in Section \ref{sec:results} which we summarise and discuss in Section \ref{sec:conclusion}. \begin{table} \centering \caption{Abbreviations for previous articles of this series.} \label{table:references} \begin{tabular}{c c c} name in the text & reference \\ \hline \paperII & \citet{BHG15} \\ \paperIII & \citet{HBG15-2} \end{tabular} \end{table} | \label{sec:conclusion} This analysis is the last part of a series of articles on the accuracy of bias parameters derived from the MICE Grand Challange (MICE-GC) simulation using clustering statistics and peak-background split (referred to as PBS) predictions \citep[][\paperII, \paperIII]{HBG15-1}. In the present analysis we studied bias parameters derived from the relation between matter and halo density contrasts ($\delta_h,\delta_m$ respectively) as an additional method for measuring bias. These measurements are compared to a selection of our most robust previous results \rev{using the same four mass samples M0-M3 at redshift $0.0$ and $0.5$. Thanks to the large volume and resolution of the MICE-GC simulation these samples span a large mass range from between roughly $10^{12}$ and $10^{15}$ \Msun, which corresponds to Milky Way like haloes and massive galaxy clusters respectively.} Our previous have been derived from two-point halo-matter cross-correlations ($\xi^{\times}$), a combination of three-point halo-matter auto- and cross-correlations ($\Delta Q$) as well as a combination of the halo-matter cross-skewness and cross-correlators ($\tau^\times$). We thereby employ leading-order modelling of clustering statistics, at which the linear bias parameters from these estimators are not affected by non-local contributions to the bias model. We therefore obtain the linear bias from $\xi^{\times}$ and the linear and quadratic bias from $\Delta Q$ and $\tau^\times$. The PBS predictions are based on MICE-GC mass function fits from \paperIII, using different mass function models, while we study bias parameters up to order three, for which we have corresponding measurements from the $\delta_h-\delta_m$ relation for validation. We studied in this work bias measurements from the $\delta_h-\delta_m$ relation in two ways. The more common method is to fit a polynomial to the $\delta_h-\delta_m$ relation, measured in the simulation. Alternatively we explore bias measurements obtained from the probability distribution function (referred to as PDF) of $\delta_h$ and $\delta_m$. The latter method has the advantage that it can be directly applied to observations, since the PDF of matter density contrasts can be modelled with theory or simulations \citep[][]{bel2016,DiPorto2015, mar05, bernardeau02}. However, it has not been tested so far how accurately the bias parameters can be determined with this method. The results of our bias comparison are summarized in Fig. \ref{fig:bias_comparison}. In the case of the linear bias $b_1$ we consider results from $\xi^\times$ ($b_\xi^\times$) as the most reliable since non-linear local and non-local term can safely be neglected. Furthermore, $b_{\xi}$ is highly relevant for cosmology since it is weakening constraints on cosmological parameters inferred from $\xi$ due to its degeneracy with the growth of matter fluctuations. Linear bias measurements and predictions from all other methods are in a $\sim 5$ percent agreement with $b_\xi^\times$, while $\Delta Q$ delivers the most accurate measurements with an overall percent level accuracy. Standard measurements from the $\delta_h-\delta_m$ relation are in a slightly better overall agreement with $b_\xi^\times$ than those from the $\delta$PDF-method. The strongest deviations from $b_\xi^\times$ shows the linear bias predictions from the different PBS model. These predictions are consistently up to $7$ percent below $b_\xi^\times$ at high masses as reported in the literature \citep[e.g.][]{M&G11, Pollack12}. \rev{The fact that we find similar results for different mass function models indicates shortcomings in the standard PBS modelling, i.e. the assumption of a constant matter density threshold for gravitational collapse \citep{Paranjape13b, Paranjape13a, Lazeyras15}.} We do not find a clear change of the variation between the different measurements and predictions with mass or redshift. In the case of the quadratic bias $c_2 \equiv b_2/b_1$ we find consistent results from the different measurements and predictions as $c_2$ increases from negative values of $\gtrsim -0.5$ at low halo masses ($\sim 10^{12} M_{\odot}/h$) to positive values of up to $\sim 2$ at high masses ($\gtrsim 10^{14} M_{\odot}/h$). However, the variation between the different results tends to be larger than in the case of $b_1$. In the case of the measurements this effect is presumably caused by the strong assumptions such as the validity of tree-level perturbation theory ($\Delta Q$ and $\tau^\times$), Poissonian shot-noise ($\tau^\times$), or a local deterministic bias model ($\delta_h - \delta_m$). We therefore do not consider any of these measurements to be sufficiently reliable for being a reference, such as the results from $\xi^\times$ in the case of the linear bias. However, the fact that the $\Delta Q$ method delivers highly accurate measurements of the linear bias suggests that also the quadratic bias is measured reliably with this approach. We find the PBS predictions for $c_2$ to be consistently above (below) all measurements in the low (high) mass range as they show overall weaker mass dependence. The latter finding lines up with reports from \citet{M&G11} and \citet{Pollack12}. Both measurements of the third-order bias $c_3\equiv b_3/b_1$ from the $\delta_h-\delta_m$ relation agree overall mutually at the $1\sigma$ level. These measurements allow for a validation of the corresponding PBS predictions. Results from both methods are similar as $c_3$ is positive below unity in the low mass range and decreases to negative values of down to $\sim -3$ in the high masses range. However, deviations are significant, as the PBS predictions tend to be below the measurements at low halo masses while results based on different mass function fits are consistent with each other. We use our various linear and non-linear bias measurements for validating the universal polynomial relation between linear and non-linear bias ($b_2(b_1)$ and $b_3(b_1)$), which we deduced in \paperIII\ from PBS predictions \rev{based on the \citet{PS74} mass function. Since this expression is independent of the peak-height we do not expect a strong dependence on halo mass definition.} Our measurements show an overall agreement with the universal behaviour predicted by the PBS model. Furthermore they agree with results from the literature derived via the Bispectrum in Fourier space or the separate universe approach from simulations with cosmologies different to the one of MICE \citep[i.e.][]{chan12,Lazeyras15}. We fit a second-order polynomial to $b_2(b_1)$ measurements from $\Delta Q$, which we consider as the most reliable $c_2$ estimator as mentioned above. Such a universal relation between linear and non-linear bias can be useful for reducing errors on the linear bias and the growth from clustering analysis when the latter is affected by $c_2$, for instance in the case of three-point correlations or two-point correlations at small scales. For applying universal polynomial relations between bias parameters in the analysis of galaxy surveys it would be interesting to show that their universality also holds for halo samples, which are selected by galaxy properties, such as luminosity and colour instead of halo mass. A correlation between the linear and quadratic bias from the 3pc and Bispectra has been reported by for SDSS galaxy samples and mock HOD catalogues by \citet{Kayo04} and \citet{Nishimichi07} and compared with PBS predictions. The limited accuracy of their measured $b_1-c_2$ relations could be strongly improved using the methods studied in the present analysis in combination with data from upcoming galaxy surveys. \FloatBarrier | 16 | 7 | 1607.01024 |
1607 | 1607.01738_arXiv.txt | An innovative field-particle correlation technique is proposed that uses single-point measurements of the electromagnetic fields and particle velocity distribution functions to investigate the net transfer of energy from fields to particles associated with the collisionless damping of turbulent fluctuations in weakly collisional plasmas, such as the solar wind. In addition to providing a direct estimate of the local rate of energy transfer between fields and particles, it provides vital new information about the distribution of that energy transfer in velocity space. This velocity-space signature can potentially be used to identify the dominant collisionless mechanism responsible for the damping of turbulent fluctuations in the solar wind. The application of this novel field-particle correlation technique is illustrated using the simplified case of the Landau damping of Langmuir waves in an electrostatic 1D-1V Vlasov-Poisson plasma, showing that the procedure both estimates the local rate of energy transfer from the electrostatic field to the electrons and indicates the resonant nature of this interaction. Modifications of the technique to enable single-point spacecraft measurements of fields and particles to diagnose the collisionless damping of turbulent fluctuations in the solar wind are discussed, yielding a method with the potential to transform our ability to maximize the scientific return from current and upcoming spacecraft missions, such as the \emph{Magnetospheric Multiscale} (\emph{MMS}) and \emph{Solar Probe Plus} missions. | A grand challenge problem at the forefront of space physics and astrophysics is to understand how the energy of turbulent plasma flows and electromagnetic fields is converted into energy of the plasma particles, either as heat or some other form of particle energization. Under the typically low-density and high-temperature conditions of turbulent plasmas in the heliosphere, such as the solar wind, the turbulent dynamics is weakly collisional, requiring the application of six-dimensional (3D-3V) kinetic plasma theory to follow the evolution of the turbulence, where the damping of the turbulent fluctuations occurs due to collisionless interactions between the electromagnetic fields and the individual plasma particles. Although \emph{in situ} spacecraft measurements in the solar wind provide detailed information about the electromagnetic and plasma fluctuations, these measurements are typically limited to one point (or, at most, a few points) in space. Of great benefit to plasma turbulence research would be a scheme to use single-point measurements of the electromagnetic fields and particle velocity distribution functions (VDFs) to diagnose the collisionless damping of the turbulent fluctuations and to characterize how the damped turbulent energy is distributed to particles with different velocities. Here we present an innovative technique to identify and characterize the collisionless mechanisms that govern the net transfer of energy from the electromagnetic fields to the plasma particles by correlating measurements of the electromagnetic fields and particle VDFs at a single point in space. These \emph{field-particle correlations} yield a local estimate of the rate of particle heating, and further provide a characteristic \emph{velocity-space signature} of the collisionless damping mechanism that leads to the energization of the plasma particles. Early attempts to explore wave-particle interactions using spacecraft measurements sought the spatial or temporal coincidence of wave fields with enhanced particle fluxes \citep{Gough:1981,Park:1981,Kimura:1983}. Later, wave-particle correlators were flown on rockets and spacecraft to identify the phase-bunching of electrons by correlating the counts of electrons in a single energy and angle bin with the phase of the dominant wave \citep{Ergun:1991a,Ergun:1991b,Muschietti:1994,Watkins:1996,Ergun:1998,Ergun:2001,Kletzing:2005,Kletzing:2006}. Motivated by modern particle instrumentation with improved temporal and phase-space resolution, the field-particle correlation technique described here takes a significant leap forward by recovering the correlation as a function of particle velocity, generating a much more detailed velocity-space signature of the collisionless interactions. Although the novel field-particle correlation technique devised here is intended for use in diagnosing the damping of turbulent fluctuations in the weakly collisional solar wind, to illustrate the concept in a simplified framework, we present here its application to the 1D-1V Vlasov-Poisson system to explore the collisionless damping of electrostatic fluctuations in an unmagnetized plasma. After this demonstration of the fundamental concept of using field-particle correlations to investigate collisionless damping of fluctuations, we discuss the application of this technique to spacecraft observations of solar wind turbulence. | Here we present a novel field-particle correlation technique that requires only single-point measurements of the electromagnetic fields and particle VDFs to return an estimate of the net rate of energy transfer between fields and particles. Furthermore, this innovative method yields valuable information about the distribution of this energy transfer in velocity space, providing a vital new means to identify the dominant collisionless mechanisms governing the damping of the turbulent fluctuations beyond that provided by measurements of velocity-integrated quantities such as $\V{j} \cdot \V{E}$. This field-particle correlation technique fully exploits the vast treasure of information contained in the \emph{fluctuations} of the particle VDFs, potentially enabling new discoveries using single-point spacecraft measurements. We believe this very general technique of field-particle correlations will transform our ability to maximize the scientific return from current, upcoming, and proposed spacecraft missions, including the \emph{Magnetospheric Multiscale} (\emph{MMS})\citep{Burch:2016}, \emph{Solar Probe Plus} \citep{Fox:2015}, \emph{Turbulent Heating ObserveR} (\emph{THOR}), and \emph{ElectroDynamics and Dissipation Interplanetary Explorer} (\emph{EDDIE}) missions. Further testing and refinement of this technique will characterize its sensitivity to the noise, limited velocity-space resolution, and limited cadence of spacecraft measurements, as well as its ability to extract a meaningful velocity-space signature of the collisionless damping mechanism in the presence of the broadband spectrum of fluctuations that is characteristic of a turbulent plasma. This work was supported by NSF AGS-1331355, NSF CAREER Award AGS-1054061, and DOE DE-SC0014599. | 16 | 7 | 1607.01738 |
1607 | 1607.04601_arXiv.txt | We use the SDSS and WISE surveys to investigate the real nature of galaxies defined as LINERs in the BPT diagram. After establishing a mid-infrared colour $\ww = 2.5$ as the optimal separator between galaxies with and without star formation, we investigate the loci of different galaxy classes in the \wha versus \ww space. We find that: (1) A large fraction of LINER-like galaxies are emission-line retired galaxies, i.e galaxies which have stopped forming stars and are powered by hot low-mass evolved stars (HOLMES). Their \ww colours show no sign of star formation and their \Ha\ equivalent widths, \wha, are consistent with ionization by their old stellar populations. (2) Another important fraction have \ww indicative of star formation. This includes objects located in the supposedly `pure AGN' zone of the BPT diagram. (3) A smaller fraction of LINER-like galaxies have no trace of star formation from \ww and a high \wha, pointing to the presence of an AGN. (4) Finally, a few LINERs tagged as retired by their \wha but with \ww values indicative of star formation are late-type galaxies whose SDSS spectra cover only the old `retired' bulge. This reinforces the view that LINER-like galaxies are a mixed bag of objects involving different physical phenomena and observational effects thrusted into the same locus of the BPT diagram. | Since 1981, inferring the main excitation sources of galaxies generally relies on diagrams involving emission line ratios \citep*{Baldwin1981,Veilleux1987}, the most widely used diagram being $\oiii/\Hb$ versus $\nii/\Ha$\footnote{Throughout this work we will use the denomination \oiii and \nii to refer to \Oiii and \Nii, respectively.}, now referred to as BPT. According to these diagrams, galaxies are divided into star-forming (SF) and active galactic nuclei (AGN) hosts, with subdivisions as LINERs, Seyferts, and composite \citep[][hereafter K01, and \citealt{Kauffmann2003c}]{Kewley2001}. \citet{Stasinska2008} showed that the BPT diagram is not able to identify a category of galaxies whose existence was inferred from stellar evolution theory: The retired galaxies, which have the same location as LINERs in the BPT diagram. Retired galaxies are objects where star formation has stopped long ago. If they contain gas, they show emission lines which are the result of photoionization by hot, low-mass evolved stars (HOLMES). \citet[][hereafter CF10 and CF11]{CidFernandes2010, CidFernandes2011} proposed a new diagram, the WHAN diagram (the \Ha equivalent width, $W_{\Ha}$, versus $\nii/\Ha$), which is able to discriminate between weak AGN and retired galaxies. Since then, many studies have found evidence for the existence of retired galaxies \citep[e.g.][]{Sarzi2010, Singh2013, Belfiore2015MNRAS449867B, Penny2015arXiv150806186P}, although numerous studies still ignore them. As shown by \citet{Stasinska2015MNRAS}, disregarding this category leads to an erroneous census of galaxy types. The main aim of this paper is to show that a large fraction of BPT-LINERs, which are often considered as galaxies with a scaled-down nuclear activity, is actually composed of objects whose emission lines are naturally explained by their old stellar populations and do not require any special `activity', be it due to an accreting black hole or to shocks or to whatever was suggested in the first place to explain LINERs \citep{Heckman1980}. An independent point raised by several authors is that diagnostics based on emission-line diagrams, when applied to spectroscopic data, as for example those from the Sloan Digital Sky Survey \citep[SDSS;][]{York2000}, lead to a classification that is aperture-dependent and may not reflect the real nature of the galaxies \citep{Gomez2003ApJ584210G, Brinchmann2013MNRAS4322112B, Papaderos2013, Stasinska2015MNRAS, Gomes2015arXiv151101300G}. These two issues, the real nature of the BPT-LINER galaxies and the aperture-dependence of the spectral classification methods, lead us to consider a new quantity, not related to emission lines: the infrared colour. Studies dealing with the mid-infrared photometry of galaxies already noticed the bimodality in infrared colours, separating objects containing warm dust (attributed to star formation) from those devoid of it \citep[e.g.][]{daCunha2008, daCunha2010MNRAS403, Alatalo2014ApJ794L}. Here we take advantage of this fact by considering data from the Wide-field Sky Survey Explorer \citep[WISE;][]{Wright2010}. WISE photometry has another advantage: the measurements refer to the entire galaxies and not only to the parts sampled by a spectroscopic fibre. This paper is organised as follows: Section \ref{sec:data} describes the data and proposes a mid-infrared criterion to separate SF galaxies from those not forming stars. It also introduces the WHAN and BPT-based spectral classes used throughout this work. Section \ref{sec:whaw} presents a new diagram based on $\mathrm{W_{H\alpha}}$ and the WISE colour \ww. Section~\ref{sec:discussion} discusses the aperture effects and some properties of the LINER-like galaxies. Section~\ref{sec:summary} summarises our results. | \label{sec:discussion} Armed with the optimal \ww separator of star-forming and retired systems obtained in Section~\ref{sec:divisor}, the BPT-based LINER samples defined in Section \ref{sec:liners}, and the WHAW diagram introduced in Section \ref{sec:whaw}, this section revisits the nature of conventionally defined LINERs and the aperture effects involved. Fig.~\ref{fig:ewhaise} shows our three BPT-based LINER sub-samples (S06, K01 and KW) in the WHAW diagram. Each subsample is divided according to the value of $R_f$, the galactic radius corresponding to the 1.5 arcsec radius of the SDSS fibre. Three bins are considered: (i) $R_f < R_{50}/4$, (ii) $R_{50}/4 < R_f < R_{50}/2$, and (iii) $R_f > R_{50}/2$, where $R_{50}$ is the Petrosian photometric radius within which 50 per cent of the $r$-band light of a galaxy originate. In this figure we follow the colour code of Fig.~\ref{fig:divisor} to identify the WHAN spectral classes. To the right of our \ww division line are the galaxies with ongoing star formation (they have warm dust and may or may not host an AGN), while to the left are objects without ongoing star formation. In all panels we indicate the number of objects in each WHAN class to the left and to the right of the $\ww = 2.5$ line. The \wha axis indicates if the gas ionisation requires source besides HOLMES. \begin{figure*} \centering \includegraphics[width=0.95\textwidth]{figure05} \caption{\ww versus $W_{\Ha}$ (the WHAW diagram) for our three BPT-LINER subsamples. Top, middle and bottom panels show S06-, K01- and KW-LINERs, respectively. Panels from left to right correspond to different bins in SDSS fibre radii ($R_f$) in units of the corresponding $r$-band Petrosian radius ($R_{50}$). The colours follow the same scheme as in Fig. \ref{fig:divisor}: green points are sAGN, orange are wAGN, red are ELR, and black are LLR galaxies. The vertical line marks our optimal separation between SF and LLR galaxies. The numbers inside the panels indicate the number of objects in the different WHAN classes on each side of the vertical line.} \label{fig:ewhaise} \end{figure*} \subsection{Impact of aperture effects in the BPT and WHAN classifications} \label{sec:aperture} Let us first explore the fact that the WISE photometry does not suffer from aperture effects to study how the limited size of the SDSS fibre affects the spectral classification. From Fig.~\ref{fig:ewhaise} we see that a fraction of ELR galaxies have high values of \ww, indicating the presence of warm dust. The incidence of these conflicting signs ($W_{\Ha} < 3$ \AA\ pointing to a retirement regime but \ww values signalling ongoing star formation) is clearly related to fibre coverage. Let us look at the central row of Fig.~\ref{fig:ewhaise}, which plots galaxies above the K01 line in bins of SDSS fibre coverage. From the numbers listed in the panels, the proportion of ELR galaxies with $\ww > 2.5$ decreases from 23 to 12 and 7 percent for fibres covering $< 1/4$, 1/4 to 1/2, and $> 1/2$ of $R_{50}$, respectively. This pattern holds true for the other LINER subsamples, although with different fractions. This all happens because the WHAN classification relies on fibre spectra, which, for small values of $R_f/R_{50}$, cover only the old bulge, while external regions rich in star-formation are not seen. It is for this same reason that in Section~\ref{sec:divisor} it was found that galaxies being predominantly of late-type can be tagged as retired when using the WHAN diagram \citep[as recently observed by][]{Gomes2015arXiv151100744G}. For the redshift range that we consider ($z < 0.2$), the small fibre coverage problem will affect all galaxies. This is specially true for massive face-on disc galaxies, in which the fibre misses the spiral arms and their young stellar populations. In fact, even for the highest distances in the sample, the coverage of the galaxies will be only of a few kpc, so the BPT and WHAN spectral classification are necessarily aperture-dependent. As $R_f$ increases, of course, the SDSS spectra refer to larger portions of the galaxies, and the signatures of the AGN and old bulge become weaker. For early-type galaxies the aperture effect will be less dramatic, given that they have a more spatially uniform distribution of stellar populations \citep{Gonzalez-Delgado2015AA581A103G}. \subsection{The nature of LINER-like galaxies in the WHAW diagram} \label{sec:ewhaise} Let us now inspect Fig.~\ref{fig:ewhaise} along its vertical direction and examine how the different sub-samples of BPT-LINERs populate the WHAW space. Qualitatively, the results are the same whatever the range of apertures chosen. In the following, we give numbers for the first column of plots only. The figure shows that most S06-LINERs (top panels) exhibit \ww\ values indicative of ongoing star formation. The occurrence of star-formation in these galaxies is not surprising given that the S06 line is built to mark the upper boundary of \textit{pure} SF systems in the BPT plane. `Composite' objects, where both star-formation and nuclear activity or other line excitation processes operate, are thus expected to be present in such a sample. Nearly 80 percent of our S06-LINERs have $\ww > 2.5$. The \ww\ values of many K01- (and even KW-LINERs) also point to star formation in spite of the fact that these objects are commonly considered as `pure AGNs'. This is in agreement with the photoionization models of \citetalias{Stasinska2006} which show that as much as 70 percent of the \Ha emission in objects on the K01 line can be due to massive stars. The incidence of SF-like \ww colours among these presumably purer-AGN subsamples is nonetheless smaller than for S06-LINERs. In fact, the distribution of \ww changes systematically along the S06-K01-KW sequence. Among S06-LINERs, 27 percent have $\ww < 2.5$, while this fraction increases to 69 and 86 percent for K01- and KW-LINERs, respectively. In other words, the further up the right wing in the BPT diagram, the smaller the fraction of galaxies with ongoing star-formation. At first sight this result fits nicely with the idea that galaxies should have emission lines increasingly dominated by nuclear activity as they move away from the star-forming sequence in the BPT plane \citep[e.g.][]{Kauffmann2009}. However, this widely-held opinion was forged prior to the awareness of the demographic relevance of retired galaxies (\citetalias{CidFernandes2011}, \citealt{Stasinska2015MNRAS}), whose emission lines are powered not by accretion onto a supermassive black hole but by HOLMES. It is therefore useful to evaluate how the importance of HOLMES varies along the S06-K01-KW sequence of BPT-LINERs. This is where the y-axis of our WHAW diagram comes into play. The fraction of galaxies in the (empirically and theoretically motivated) $\wha < 3$ \AA\ and $\ww < 2.5$ HOLMES-dominated regime increases from 69 percent in the case of S06-LINERs to as much as 78 and 89 percent for K01 and KW-LINERs, respectively. (These fractions are only slightly affected by the aperture effects discussed in the previous section.) In other words, {\it the further up the right wing in the BPT diagram, the higher the fraction of retired galaxies among BPT-LINERs.} These numbers reveal that, as far as BPT-LINER systems are concerned, galaxies traditionally taken to represent `pure' AGN (i.e., K01- and KW-LINERs) are, statistically speaking, more likely to represent systems whose emission-line properties can be fully understood in terms of old stellar population properties. Let us be clear that we are not saying that objects with low nuclear activity do not exist. Such objects belong to our weak and strong AGN classes (orange and green points in Fig.~\ref{fig:ewhaise}). What we argue is that many BPT-LINERs are in fact ELR galaxies, as first shown by the WHAN diagram and now confirmed by the WHAW. \subsection{The negligible contribution of AGNs to the mid-IR emission of LINERs} Our whole analysis neglects the potential contribution of AGNs to the WISE fluxes. To some extent, the result that the \ww colours of all AGN considered in this work lie between those of SF and LLR galaxies (see Fig.~\ref{fig:divisor}a), neither of which is suspected to host AGN, vindicates this assumption. It is nevertheless worth addressing this issue in light of previous work which indicates that AGN emission in the mid-IR can be significant in some cases \citep{Mateos2012MNRAS4263271M,Mateos2013MNRAS434941M,Yan2013,Caccianiga2015}. To evaluate the potential contribution of AGN to the mid-IR emission of our LINERs, Fig.~\ref{fig:wise_liners} presents a $W1-W2$ versus $W2-W3$ colour-colour diagram. The points are coloured as a function of the concentration index. In this figure, we also show the optimal divisor line obtained in this work, $W2-W3 = 2.5$, the mid-IR criterion to select AGNs proposed by \citet{Stern2012}, $W1-W2 > 0.8$, and the locus expected for a power-law spectrum of varying slope, above which the infrared emission could be explained purely by nuclear activity \citep{Caccianiga2015}. Dashed red lines mark the `AGN wedge' defined by \citet{Mateos2012MNRAS4263271M}. As it is clear from the plot, only a few objects in our sample lie in the AGN region. For the overwhelming majority of the BPT-LINERs, the host galaxy dominates the mid-IR emission, explaining their loci on the $W1-W2$ versus $W2-W3$ diagram. The black, purple, and blue lines in Fig.~\ref{fig:wise_liners} mark the 90 percent contours of S06, K01, and KW-LINERs, respectively. A gradual shift towards smaller \ww and slightly smaller $W1-W2$ along the S06-K01-KW sequence is observed. If AGN emission contributed significantly to the mid-IR, this sequence towards purer-AGN should move sources towards the top-right of this diagram, which is not what is seen. Seyfert galaxies (green contours) do stretch towards the AGN region, indicating some contribution of the AGN to the mid-IR emission. For BPT-LINERs, the focus of this paper, however, we conclude that this contribution is negligible. \begin{figure} \centering \includegraphics[width=0.92\columnwidth, trim=20 30 20 20]{figure06} \caption{ $W1-W2$ versus $W2-W3$ colour-colour diagram for our BPT-LINER sample. Points are coloured as a function of concentration index $CI$. The contours indicate the 90th percentiles for each LINER sub-sample selected by using the S06, K01, and KW lines (black, purple, and blue lines, respectively). The 90th-percentile contours in cyan and green mark the loci for SF galaxies and Seyfert hosts selected with the S06 line in the BPT diagram (note that we do not plot the points for those galaxies). The vertical solid line indicates our best divisor between objects with ongoing star formation and retired galaxies. The dashed horizontal line is the mid-IR criterion to select AGNs proposed by \citet{Stern2012}. The black solid line marks the location of a simple power-law spectrum with varying slope as proposed by \citet{Caccianiga2015}, and the dashed red lines delineate the `AGN wedge' as defined by \citet{Mateos2012MNRAS4263271M}.} \label{fig:wise_liners} \end{figure} In this work we have selected a sample of galaxies from the SDSS and WISE surveys and studied their mid-infrared colour \ww, a powerful indicator of the presence or absence of ongoing star formation. We find that $\ww = 2.5$ optimally separates SF from retired galaxies. We focused our attention on galaxies commonly tagged as LINERs in the BPT plane and constructed a new diagram, the WHAW diagram, plotting \wha versus \ww. We found that this diagram confirms that most galaxies classified as retired by the WHAN diagram (i.e. those with $\wha < 3$~\AA, which includes many BPT-defined LINERs) indeed do not present any sign of ongoing star formation in the WISE data. \citet{Stasinska2008} and \citetalias{CidFernandes2011} showed that the optical emission lines of these systems are explained by their populations of HOLMES and do not require the presence of an AGN. In addition, we warn against interpreting the entire right wing in the BPT diagram as a SF--AGN mixing line. By studying subsamples of BPT-LINERs farther away from the SF sequence, we conclude that the tip of the BPT-LINER wing is dominated by emission-line retired galaxies powered by HOLMES, and not by AGNs. The WISE data also allowed us to tackle aperture effects, which may bias galaxy classifications based on fibre spectroscopy. For galaxies with a small covering fraction, a non-negligible fraction of retired galaxies have a \ww colour indicative of star formation. In these cases, the WHAN diagram relates to a `retired' bulge, while the \ww colour refers to the entire galaxy and indicates the presence of SF regions in the galactic disk. In conclusion, while WHAN-retired SDSS galaxies may be counterfeiters of retired bulges in SF galaxies, WISE-retired galaxies are truly retired. | 16 | 7 | 1607.04601 |
1607 | 1607.08872_arXiv.txt | Radio relics in galaxy clusters are associated with powerful shocks that (re)accelerate relativistic electrons. It is widely believed that the acceleration proceeds via diffusive shock acceleration. In the framework of thermal leakage, the ratio of the energy in relativistic electrons to the energy in relativistic protons should should be smaller than $K_{\rm e/p} \sim 10^{-2}$. The relativistic protons interact with the thermal gas to produce $\gamma$-rays in hadronic interactions. Combining observations of radio relics with upper limits from $\gamma$-ray observatories can constrain the ratio $K_{\rm e/p}$. In this work we selected 10 galaxy clusters that contain double radio relics, and derive new upper limits from the stacking of $\gamma$-ray observations by \emph{Fermi}. We modelled the propagation of shocks using a semi-analytical model, where we assumed a simple geometry for shocks and that cosmic ray protons are trapped in the intracluster medium. Our analysis shows that diffusive shock acceleration has difficulties in matching simultaneously the observed radio emission and the constraints imposed by \emph{Fermi}, unless the magnetic field in relics is unrealistically large ($\gg 10 ~\rm \mu G$). In all investigated cases (also including realistic variations of our basic model and the effect of re-acceleration) the mean emission of the sample is of the order of the stacking limit by \emph{Fermi}, or larger. These findings put tension on the commonly adopted model for the powering of radio relics, and imply that the relative acceleration efficiency of electrons and protons is at odds with predictions of diffusive shock acceleration, requiring $K_{\rm e/p} \geq 10-10^{-2}$. | \label{sec:intro} Radio relics are steep-spectrum radio sources that are usually detected in the outer parts of galaxy clusters, at distances of $\sim 0.5-3$ Mpc from their centres. They are very often found in clusters with a perturbed dynamical state. There is good evidence for their association with powerful merger shocks, as early suggested by \citet[][]{1998A&A...332..395E}. Among them are {\it giant double-relics} that show two large sources on opposite sides of the host cluster's centre \citep[e.g., ][]{fe12,2012MNRAS.426...40B,fdg14}. These relics are associated with shocks in the intracluster medium that occur in the course of cluster mergers. Only a tiny fraction of the kinetic power dissipated by typical cluster merger shocks ($\ll 10^{-3}$) is necessary to power the relics, and diffusive shock acceleration (DSA, \citealt[e.g.][for modern reviews]{2012JCAP...07..038C,kr13}) has so far been singled out as the most likely mechanism to produce the relativistic electrons and to produce the observed power-law radio spectra \citep[e.g.][]{hb07}. However, if the standard DSA model is correct, the same process should also lead to the acceleration of cosmic-ray (CR) protons. Indeed, the process should be much more efficient for protons, owing to their larger Larmor radius {\footnote{Instead, since the Larmor radius of thermal electrons is much smaller than the typical shock thickness, thermal electrons {\it cannot} be easily accelerated to relativistic energies by DSA. This so called injection problem for electrons is still largely unresolved \citep[e.g.][]{2014IJMPD..2330007B}}}. \\ To date, high-energy observations of nearby galaxy clusters have not revealed any diffuse $\gamma$-ray emission resulting from the interaction between relativistic protons and thermal particles of the intracluster medium (ICM) \citep[][]{re03,aha09,al10,alek12,arl12,2014MNRAS.440..663Z}. Recently, the non-detection of diffuse $\gamma$-ray emission from clusters by \emph{Fermi} has put the lowest upper limits on the density of CRs in the ICM, $\leq$ a few percent of the thermal gas energy within the clusters virial radius \citep{ack10,fermi14}. Moreover, the stacking of subsets of cluster observations leads to even lower upper-limits \citep[][]{2013A&A...560A..64H,2014ApJ...795L..21G}. These low limits on the energy content of CRs can be used to constrain shock-acceleration models. Recently, \citet{va14relics} have already suggested that the present statistics of radio observations, combined with available upper limits by \emph{Fermi} places constraints on DSA as the source of giant radio relics. In \citet{va14relics}, we assumed that the population of clusters with radio relics was similar to the population of (non-cool-core) clusters for which the stacking of \emph{Fermi} clusters was available. In the present paper, we repeat a similar analysis by comparing to a more realistic stacking of the \emph{Fermi} data. Our method is outlined in Sec.~\ref{subsec:algo}, while our results are given in Sec. \ref{sec:results}. In the latter Section, we also discuss on the role played by the several open parameters in our modeling. We find (Sec.~\ref{sec:results}) that the present upper limits from \emph{Fermi} imply energy densities of CR-protons that are too low to be explained by standard DSA: if DSA produces the electrons in relics, then we should have already detected hadronic $\gamma$-ray emission in some clusters, or in stacked samples. In our conclusions (Sec.~\ref{sec:conclusions}), we discuss possible solutions to this problem, as suggested by recent hybrid and particle-in-cell simulations of weak, collisionless shocks. | \label{sec:conclusions} Most of the evidence from X-ray and radio observations suggests a link between radio relics and merger shocks: merger axes of clusters and relic orientations correlate \citep[][]{2011A&A...533A..35V}, the power of radio relics scales with X-ray luminosities \citep[][]{2012MNRAS.426...40B,fe12} and mass \citep[][]{fdg14} of the host cluster. Cosmological simulations produce emission patterns consistent with observed radio relics just using a tiny fraction of the kinetic energy flux across shock waves \citep[e.g.][]{ho08,2008MNRAS.385.1242P,2009MNRAS.393.1073B,sk11,va12relic,2012MNRAS.420.2006N,sk13}.\\ Still, a number of recent observations have revealed some open issues, including uncertain merger scenarios \citep[e.g.][]{2013MNRAS.433..812O}, departures from power-law spectra \citep[e.g.][]{2014MNRAS.445.1213S,2014arXiv1411.1113T}, missing associations between radio emission and X-ray maps \citep[e.g.][]{2011MNRAS.417L...1R,2014MNRAS.443.2463O}, efficiencies problems \citep[e.g.][]{2011ApJ...728...82M,2012MNRAS.426...40B}, inconsistencies between Mach numbers derived from X-ray and radio observations \citep[e.g.][]{2012MNRAS.426.1204K,2013PASJ...65...16A,2013MNRAS.433.1701O} and apparent connections to radio galaxies \citep[][]{2014ApJ...785....1B}.\\ In this work, we used a simple semi-analytical model of expanding merger shocks in clusters to reconstruct the propagation history of shocks leading. We used a spherically symmetric model and assumed that cosmic ray protons are trapped in the intracluster medium on all relevant timescales. A range of realistic scenarios for the acceleration of relativistic electrons and protons via DSA, varying the upstream gas conditions, the shock parameters and the budget of pre-existing cosmic rays, gives very similar results. In all realistic scenarios, a significant fraction of our objects ($\sim 1/2-1/3$) has difficulties in matching at the same time the observed radio emission and the constraints imposed by the \emph{Fermi} limits, unless the magnetic field in all problematic objects is much larger than what usually considered realistic ($\gg 10 ~\rm \mu G$). The scenario in which radio emitting electrons comes from the re-acceleration of pre-existing electrons \citep[][]{ka12,pinzke13} can alleviate the tension with \emph{Fermi} if the pre-existing electrons are not the result of previous injection by shocks as we investigated here, but are instead released by mechanisms that mostly inject leptons (e.g. leptonic-dominated jets from AGN), as already discussed in \citet{va14relics}. Based on our semi-analytical model, the standard DSA scenario with thermal leakage that predicts that $E_{\rm CRe} \ll E_{\rm CRp}$ cannot simultaneously explain radio relics and produce less $\gamma$-radiation than the upper limits from \emph{Fermi}, unless unrealistically large magnetic fields are assumed at the position of relics (e.g. $B_{\rm relic} \ge 10-100 ~\rm \mu G$). This result is very robust, at least in the statistical sense, against all investigated variations of our fiducial parameters for the modeling of the shock acceleration of CRs. Additional effects that go beyond our idealized modeling of cluster mergers, e.g. a clumpy ICM, a succession of mergers and the additional acceleration of CRs by AGN, supernovae, turbulence or reconnection exacerbate this discrepancy. \\ Despite its obvious degree of simplification, a semi-analytical method is useful to tackle the case of double relic systems. In these systems it is reasonable to assume that most of the energetics is related to the observed pair of giant merger shocks. Their shape and location is rather regular and symmetric with respect to the cluster centre, suggesting that one can make reasonable estimates for their propagation history. This setup allows us to run very fast testing of different possible acceleration scenarios, and as we showed in our various tests it generally gives a {\it lower limit} on the expected $\gamma$-ray emission. This method is meant to be complementary to fully cosmological numerical simulations where the effects of multiple shocks, particle advection and cooling, as well as inputs from galaxy formation and other mechanisms can be taken into account at run-time \citep[e.g.][]{pf07,scienzo}. However, a thorough exploration of models is computationally demanding because of the required high resolution and the complexity of the numerics. Also the agreement between different numerical techniques on this topic is still unsatisfactory \citep[see discussion in][]{va11comparison}. Another way of illustrating our result is found by rescaling the efficiency for proton acceleration, $\eta(M)$, such that the upper limits from \emph{Fermi} are not violated. In the case without pre-existing CRs, this is a simple exercise as we only need to rescale $\eta(M)$ for each relic separately, and compute the average of the efficiencies for each bin of the Mach number (here we chose a bin size of $\Delta M=0.6$ to achieve a reasonable sampling of the sparse distribution of Mach numbers in the dataset). Here, we keep the magnetic field fixed at $B=2 \mu$G as suggested by recent observations \citep[][]{fdg14}. The result is shown in Fig.~\ref{fig:last}, where we show the maximally allowed acceleration efficiency for CR-electrons, protons, as well as $K_{\rm e/p}$ as a function of $M$. This relation results from a somewhat coarse simplification of the problem but it is a rough estimate of the acceleration efficiencies in weak ICM shocks. For shocks with $M \leq 2$, the flux ratio of injected electrons is larger than that in protons, $K_{\rm e/p} \sim 1-100$, at odds with standard DSA (even including re-accelerated electrons). For $M \geq 2.5$ the acceleration efficiency of protons can become significant ($\sim 10^{-3}-10^{-2}$) while the acceleration efficiency of electrons flattens and $K_{\rm e/p} \sim 10^{-2}$. The functional shape of the acceleration efficiency for protons is consistent with the \citep[][]{kr13} model, but the absolute normalisation is lower by a factor $\sim 10-100$. \begin{figure} \includegraphics[width=0.495\textwidth]{images/sausage_tabulated_models_kr13_Bfix.dat_efficiency_SDA.eps} \caption{Acceleration efficiency of CR-protons (green, rescaled by a factor $\times 10$ down) and CR-electrons (blue), and electron to proton acceleration ratio (red) allowed by our combined radio and $\gamma$-ray comparison with observations. In this case, we assumed a fixed magnetic field of $B=2 \mu G$ for all relics.} \label{fig:last} \end{figure} A possible solution has been suggested by \citet{2014ApJ...788..142K}, who assumed that electrons and protons follow a $\kappa$-distribution near the shock transition. A $\kappa$-distribution is characterized by a power-law rather than by an exponential cutoff at high energies, thus ensuring a more efficient injection of high-energy particles into the DSA cycle. This distribution is motivated by spacecraft measurements of the solar wind as well as by observations of HII and planetary nebulae \citep[e.g.][and references therein]{2012ASSP...33...97L}. \citet{2014ApJ...788..142K} explored the application of the $\kappa$-distribution to $M \leq 2$ shocks in the ICM, and concluded that the distribution can have a different high-energy tail as a function of the shock obliquity and of the plasma parameters. In the ICM, the distribution might be more extended towards high energies for electrons than for protons thus justifying a higher acceleration efficiency for electrons than for protons. However, in order to explain the origin of these wider distributions, one must resort to detailed micro-physical simulations of collisionless shocks.\\ The most promising explanation for the non-observation of $\gamma$-rays has been suggested by \citet[][]{2014arXiv1409.7393G} who studied the acceleration of electrons with particle-in-cell (PIC) simulations under conditions relevant to merger shocks. They showed that $M \leq 3$ shocks can be efficient accelerators of electrons in a Fermi-like process, where electrons gain energy via shock drift acceleration (SDA). The electron gain energy from the motion of electric field and scatter off oblique magnetic waves that are self-generated via the firehose instability. They found that this mechanism can work for high plasma betas and for nearly all magnetic field obliquities. However, these simulations have been performed in 2D, and could not follow the acceleration of electrons beyond a supra-thermal energy because of computing limitations. At the same time, hybrid simulations of proton acceleration by \citet{2014ApJ...783...91C} have shown that the acceleration efficiency is a strong function of the obliquity angle. If indeed the magnetic field in radio relics is predominantly perpendicular to the shock normal, as found e.g. in the relic in the cluster CIZA 2242.2+5301, then the prediction is that the acceleration efficiency of protons is strongly suppressed, thus explaining the non-detection of hadronic emission. It remains to be seen if the results of these simulations hold in 3D, with realistic mass ratios between electrons and protons and coupled to a large scale MHD flow. It is also not clear whether the magnetic field is quasi-perpendicular in all the relics of this sample and how the alignment of the magnetic fields with the shock surface observed on large scales can be scaled down to scales of the ion gyro radius. | 16 | 7 | 1607.08872 |
1607 | 1607.01953_arXiv.txt | Spiral structure in disk galaxies is modeled with nine collisionless N-body simulations including live disks, halos, and bulges with a range of masses. Two of these simulations make long-lasting and strong two-arm spiral wave modes that last for $\sim5$ Gyr with constant pattern speed. These two had a light stellar disk and the largest values of the Toomre $Q$ parameter in the inner region at the time the spirals formed, suggesting the presence of a Q-barrier to wave propagation resulting from the bulge. The relative bulge mass in these cases is about 10\%. Models with weak two-arm spirals had pattern speeds that followed the radial dependence of the Inner Lindblad Resonance. | \label{sec:intro} Spiral structure in disk galaxies results from gravitationally amplified growth of density perturbations (\citealt{GoldreichLyndenBell65,Carlberg85,ToomreKalnajs91, Huber02}, see reviews in \citealt{Athanassoula84,Bertin96}) that range in scale from interstellar clouds \citep{DOnghiaetal2013}, to other spirals \citep{Masset1997}, to bars \citep{Salo10} and passing galaxies \citep{ToomreToomre72,Salo00}. Simulated spiral arms are usually individually short-lived (\citealt{Sellwood2011}; however, see \citealt{Elmegreen1993}), although the presence of spiral structure in one form or another may be long-lived \citep{Sellwood2014}. \cite{Lindblad1962,Lindblad1963} considered spirals ``quasi-stationary'' when he proposed they might result from synchronized epicyclic motions over a wide range of radii as viewed in a rotating coordinate system. This was a better model than material arms which would wrap up too quickly to explain the high fraction of disk galaxies with spirals. However, the correspondence noted by Lindblad was not perfect and the proposed arms would still deform over time. \cite{LinShu1964} offered a solution to this problem by noting that arm self-gravity would perturb the epicycles in the right sense and cause them to lock into phase. Their solution was only a start, though, as \cite{Toomre1969} pointed out that Lin-Shu waves also wrap up because their group velocity is inward. The second part of this solution proposed that incoming waves refract or reflect in the central regions and move back out \citep{Lin70,Mark1976a}. Then they can amplify by disk gravity at corotation (CR) and send another trailing wave back in as well as a trailing wave out \citep{Mark1976b,Mark1976c}. The result is a growing spiral wave mode, which is a standing wave with components moving in both directions \citep{Mark1977,Bertin1989a,Bertin1989b}. In the WASER type II mode discussed by \cite{Bertin1983} and \cite{LinBertin1985} an inward-moving trailing spiral wave reflects off a high-Q barrier in the inner region, such as a bulge, and returns to CR as a weak leading wave. The leading wave then swings around into a stronger trailing wave in analogy to the swing amplifier proposed by \cite{Toomre1981}, and the trailing wave moves inward again. The outward moving trailing wave beyond CR extends to the outer Lindblad resonance where it resonates with the thermal motions of stars. For such a wave mode, the pattern speed is determined by the propagation condition for a wave to move in to the reflection radius and back out to CR in the time it takes the pattern to rotate to another arm \citep{Bertin1989b}. Astronomical evidence for long-lived global modes is weak. \cite{EE1983} suggested that two-arm spirals live for at least $\sim2$ Gyr considering an observed increase in the fraction of such ``grand-designs'' with galaxy group crossing rate and galaxy-galaxy collision rate. This duration corresponds to $\sim5$ rotations in the outer parts of a typical galaxy and is reasonably consistent with a spiral mode. Morphological evidence for spiral modes was suggested in \cite{Elmegreen1989} and \cite{Puerari2000} by the presence of symmetric spiral arm amplitude variations from interfering inward and outward moving waves \citep[see also][]{Bertin1993}. The arm variations observed for M81 were fit to the modal theory by \cite{Lowe1994}. Further evidence for modes was shown in \cite{Elmegreen1992} by the presence of 2, 3, and 4-arm symmetric spirals in 18 galaxies, with regular amplitude variations along the arms and arm endpoints at the appropriate Lindblad resonances. These latter two studies used computer-enhanced and symmetrized images to show features that are not obvious to the eye. Modern simulations always have transient spiral features, even when they are wave modes \citep{Sellwood2011}. This transience is partly because the modes adjust the radial distribution of stars and their velocity dispersions to change the basic state \citep{Sellwood2012}, and they also trigger new modes which spring up at resonances of the old ones \citep{Sygnet1988}. Still, one wonders if a simulation with the right initial conditions can make a single, long-lasting wave mode in the sense discussed by \cite{Bertin1989b} and others. One essential component is a bulge that shields the inner Lindblad resonance (ILR) from incoming waves so they can reflect back to the amplification zone at CR. Otherwise the ILR will absorb the wave \citep{Lynden-BellKalnajs1972}. The present paper shows several examples of particle simulations that have a live disk, bulge, and halo and that appear to make single, long-lasting wave modes with constant pattern speeds. They have the strongest amplitude when the bulge is optimal for shielding the ILR. The models are in a sequence of increasing bulge mass, which also increases the strength of the ILR. For low-mass bulges, the ILR is weak but the bulges are also weak, and the spirals end up weak themselves. For high-mass bulges, the ILR absorption and bulge shielding are both strong but the ILR apparently wins again, making the waves weak. But for intermediate mass bulges, the conditions are optimal to make a strong and persistent two-arm wavemode, which circulates at constant pattern speed from $\sim 4$ Gyr to $\sim8$ Gyr in the simulation. Eventually the large arm amplitudes that result from this continuous growth can distort and destroy the spiral mode. \begin{figure} \begin{flushleft} \rotatebox{0}{\includegraphics[height=6.3 cm]{Fig1-10may16.eps}} \caption{Initial set up and Q-bump during the growth of spiral arms. Upper panel(left): circular velocity curves. Lower panel(left): radial variation of $\Omega, \Omega - \kappa/2,\Omega - \kappa/4 $ for different model galaxies. $S_{\rm ILR}$ denotes the ILR strength defined in section~\ref{sec:simulation}. Upper panel (right): Initial Toomre Q profiles for all the model galaxies. The bulge mass ($M_b$) is in units of $10^{9}M_{\odot}$. Lower panel(right): The Toomre Q parameter versus radius for models b20p8(black), b6p8 (magenta) and b1p8 (green). Q for model b6p8 develops a peak in the inner region over time, and grows strong spiral arms when the peak is largest (Fig.~\ref{fig:A2rt}). Model b1p8 has a smaller Q peak later and grows a weak spiral later. Model b20p8 has no significant Q peak and very weak arms throughout the simulation. Only the models with a Q peak in the inner parts developed strong two-arm spirals, and they did this after the Q peak appeared.} \label{fig:omgkapa} \end{flushleft} \end{figure} | \label{sec:discuss} Two of our particle simulations with a live halo, bulge and a light disk were found to generate spiral wave modes that grew slowly at first over a period of 1 to 2 Gyr and then quickly to a relatively large amplitude over the next 1 Gyr, after which they maintained a constant pattern speed for another 5 Gyr. These two simulations differed from the others, which did not make strong spiral arms, in the relative height of the maximum value of the Toomre Q parameter in the inner region. This inner Q peak results from the bulge. We interpret this result as evidence that wave reflection off a classical bulge can lead to the formation of a long-lasting spiral wave mode, as proposed by \cite{Bertin1989a}. \medskip \noindent{\bf Acknowledgement:} The authors thanks the anonymous referee for several useful comments including the Fourier power spectrum (shown in Fig.~5). | 16 | 7 | 1607.01953 |
1607 | 1607.01448_arXiv.txt | Astrophysical environments that reach temperatures greater than $\sim 100$ keV can have significant neutrino energy loss via both plasma processes and nuclear weak interactions. We find that nuclear processes likely produce the highest-energy neutrinos. Among the important weak nuclear interactions are both charged current channels (electron capture/emission and positron capture/emission) and neutral current channels (de-excitation of nuclei via neutrino pair emission). We show that in order to make a realistic prediction of the nuclear neutrino spectrum, one must take nuclear structure into account; in some cases, the most important transitions may involve excited states, possibly in both parent and daughter nuclei. We find that the standard technique of producing a neutrino energy spectrum by using a single transition with a Q-value and matrix element chosen to fit published neutrino production rates and energy losses will not accurately capture important spectral features. | In this paper, we calculate energy spectra for neutrinos produced in nuclear weak interaction processes that occur in pre-collapse massive stars. A key motivation for this work is the possibility of detecting a neutrino signal from a massive star, perhaps even months before collapse \cite{omk:2004a,omk:2004b,oh:2010,asakura-etal:2016}. Patton \& Lunardini (hereafter P\&L) \cite{pl:2015} have studied the neutrino emissivity physics in this problem and the associated prospects for detection. In this paper, we build on the work of P\&L, but we concentrate on the nuclear physics which determines important aspects of the neutrino energy spectra, especially at high neutrino energy. Higher neutrino energies are, of course, key to detection. Our nuclear structure considerations and our shell model calculations allow us to illuminate features of specific sd-shell nuclei which are likely to be key contributors to the high-energy end of the expected neutrino spectrum. Beginning with core carbon burning, neutrino production dominates the energy loss of massive stars. Depending on the mass of the star and its stage of burning, these neutrinos can be produced through electron-positron pair annihilation, the photo process (wherein a photon interacts with an electron and produces a neutrino pair), electron neutrino pair bremsstrahlung, electron capture, and other processes. In low mass stars and in massive stars prior to core collapse, the neutrinos stream unimpeded through stellar material, removing entropy from the core and greatly accelerating the evolution of the star \cite{whw:2002}. In the final stages before collapse of a massive star, the core is hot and dense, but the entropy is low \cite{bbal:1979}. The temperature is $\sim 0.5$ MeV, but the electron Fermi energy can be $\sim 5$ MeV, implying very electron-degenerate conditions \cite{arnett:1977,bw:1982,bw:1985,bmd:2003,bm:2007,sjfk:2008,bbol:2011,ajs:2007,liebendorfer-etal:2008,liebendorfer-etal:2009,hjm:2010,hix-etal:2010,bdm:2012}. The electron degeneracy relatively suppresses neutrino production processes with electrons in the final state, processes with intermediate electron loops, and electron-positron annihilation. At the same time, the high Fermi energy relatively enhances electron capture (figure \ref{fig:feyn_cap}), while the high temperature gives a population of excited nucleons that can de-excite by emission of a neutrino pair (figure \ref{fig:feyn_nc_decay}) \cite{btz:1974,fm:1991,mbf:2013}. In many cases, excited nuclei can also more readily decay by electron or positron emission \cite{ffn:1982a}, which is always accompanied by an anti-neutrino or neutrino, respectively (figure \ref{fig:feyn_dec}). High temperatures allow the nuclei to access excited parent states which may have large Q-values and large weak interaction matrix elements for charged current transitions. Large Q-values imply larger phase space factors for weak interactions, but against this, Boltzmann population factors for these highly excited initial states can be small. However, ameliorating the effect of small Boltzmann factors is the near-exponential increase in nuclear level densities with increasing excitation energies. In the end, the balance between all these factors must be evaluated on a case-by-case basis for individual nuclei and particular thermodynamic conditions in the star. \begin{figure}[here] \includegraphics[scale=.6]{feyn_cap.pdf} \caption{Electron capture on a nucleus of mass number A, proton number Z, and excitation energy E, producing a nucleus of mass number A, proton number Z-1, and excitation energy E$'$. The electron and neutrino may be exchanged in this diagram for their antiparticles, yielding a final nucleus with proton number Z+1.} \label{fig:feyn_cap} \end{figure} \begin{figure}[here] \includegraphics[scale=.6]{feyn_nc_decay.pdf} \caption{Neutral current neutrino pair emission from a nucleus of mass number A with initial excitation energy E and final excitation E$'$.} \label{fig:feyn_nc_decay} \end{figure} \begin{figure}[here] \includegraphics[scale=.6]{feyn_dec.pdf} \caption{Electron decay from a nucleus of mass number A, proton number Z, and excitation energy E to a nucleus of mass number A, proton number Z+1, and excitation energy E$'$. The electron and antineutrino may be exchanged in this diagram for their antiparticles, yielding a final nucleus with proton number Z-1.} \label{fig:feyn_dec} \end{figure} This situation has a profound effect on the neutrino spectrum, as energetic electrons can capture onto excited parent nuclei, which might have a \emph{less} excited final state in the daughter nucleus; this results in an unusually high energy neutrino. Furthermore, these excited nuclei may directly produce neutrino pairs. When excited nuclei de-excite, the usual channel is gamma ray emission; however, they may also emit a virtual Z$^0$ boson that decays into a neutrino anti-neutrino pair, shown schematically in figure \ref{fig:feyn_nc_decay}. In fact, this can be the dominant source of neutrino pairs in a collapsing stellar core \cite{gershtein-etal:1975,km:1979,fm:1991,mbf:2013,flm:2013}. If the nucleus de-excite from a highly excited state, it can produce an energetic neutrino pair of any flavor, and these neutrinos can make a substantial contribution at the high energy end of the neutrino spectrum. One final process that we will not discuss but which falls under the general purview of nuclear neutrinos is neutral current inelastic neutrino scattering on nuclei (figure \ref{fig:feyn_nc_scatter}) \cite{wh:1988,woosley-etal:1990,fm:1991}. Scattering does not produce neutrinos, but it can alter the neutrino spectrum. During the event, the nucleus can either gain internal energy from the neutrino in a subelastic scatter, or the nucleus can give up energy to the neutrino in a superelastic scatter. The former will shift the neutrino spectrum down in energy, while the latter will shift it up. Under most circumstances, there will be greater strength for a nuclear ``up-transition'' \cite{fm:1991}, meaning a subelastic scatter that lowers neutrino energy. However, in supernova environments, there may be a sufficient population of excited nuclei to shift part of the neutrino spectrum up, lengthening the high energy tail of the spectrum, with possible implications for detection. \begin{figure}[here] \includegraphics[scale=.6]{feyn_nc_scatter.pdf} \caption{Neutral current neutrino scattering from a nucleus of mass number A with initial excitation energy E and final excitation E'.} \label{fig:feyn_nc_scatter} \end{figure} Sec. II details the calculation of the charged current process neutrino spectra and shows some results of high temperature shell model calculations. In Sec. III we discuss neutral current nuclear de-excitation neutrino production and spectra, and in Sec \ref{sec:discussion}, we go over the results and their implications. | \label{sec:discussion} Detecting neutrinos from highly evolved pre-collapse stars could give key insights into stellar evolution. This is an exciting prospect. P\&L astutely point out the importance of nuclear neutrinos in understanding late stellar neutrino spectra, and we build on that by examining the effects of nuclear structure. To that end, we draw specific attention to $^{32}$P, shown in figure \ref{fig:32p_cc} of this work and figure 3 of P\&L \cite{pl:2015}. In P\&L figure 3, there is a small bump in the anti-neutrino spectrum at $\sim 4$ MeV that the authors say is due positron capture on $^{32}$P. At that point in P\&L's simulation, the mass fraction of $^{32}$P is $\sim 10^{-4}$ (personal communication). Using the mass fraction and the density of the core, we convert the P\&L y-axis and find that the height of the P\&L $^{32}$P 4 MeV anti-neutrino peak is $\sim 8.5\times 10^{-10}$ neutrinos/second/baryon/MeV. This corresponds roughly with the height of the $\sim 1$ MeV positron capture neutrino peak in our figure \ref{fig:32p_cc}. By design, the single Q-value technique will give the correct total neutrino output with the correct average energy, but in this case, the energetics of the positron capture neutrinos are incorrect. In this particular case, the published rates are dominated by electron emission from the first excited parent state, but most captures occur between the parent ground state and the first excited state of the daughter, pushing the positron capture neutrino energy down; this results in erroneous conclusions from the single Q-value method. Similarly, the single Q-value method fails to capture the significant contribution of 3-4 MeV anti-neutrinos from the $E_i=2.23$ MeV state. Finally, comparing figure \ref{fig:nc_spectra} in this work with the final plot in figure 4 of P\&L indicates that in late silicon burning, neutral current de-excitation may be a leading source of anti-neutrinos with energies greater than 10 MeV. This is contradicted by the $^{28}$Al spectrum in this work's figure \ref{fig:28al_cc}, however, so we must be circumspect in drawing conclusions. Nevertheless, it is clear that the production rates of $>10$ MeV neutrinos increase dramatically with temperature (due to the exponential dependence of the Boltzmann factor for excited states), and these rates are entirely unaffected by density and the associated baggage (such as electron blocking). This implies that in this neutrino energy range, the effectiveness of this process relative to charged current processes is highly sensitive to the ambient temperature, and no solid conclusions can be drawn until realistic nuclear neutrino spectra are included in a simulation. At higher temperatures--approaching the onset of collapse--neutral current de-excitation may be the dominant source of $>10$ MeV neutrinos. \begin{figure} \includegraphics[scale=.42]{56fe_cc_043.pdf} \includegraphics[scale=.42]{56fe_cc_059.pdf} \caption{$^{56}$Fe charged current process neutrino spectra computed from the FFN prescription. The $E_i=11.44$ MeV line corresponds to the isobaric analog of the $^{56}$Mn ground state. In the upper panel, the electron chemical potential is less than the Gamow-Teller resonance energy, while in the lower panel, it is greater than the GT resonance energy. Because of this, the peak in the lower panel in more than 4 orders of magnitude greater, despite the comparatively small increase in temperature.} \label{fig:56fe_cc} \end{figure} During core collapse, the electron chemical potential ($\mu_e$) climbs as density increases, with the consequence that the average energy of a captured electron is very high. When $\mu_e$ reaches the energy of the Gamow-Teller resonance of a typical nucleus, the capture rate takes off, producing neutrinos prodigiously. Following precisely the method of FFN, we computed electron capture and positron decay neutrino spectra for $^{56}$Fe. Using the FFN prescription, $^{56}$Fe has a GT resonance at $\sim8$ MeV. Figure \ref{fig:56fe_cc} (same line designations as in figure \ref{fig:27si_cc}) shows the spectra for two points leading up to and during collapse. The upper panel has $\mu_e=2.22$ MeV (less than the resonance), and the lower has $\mu_e=9.66$ MeV (greater than the resonance). The increase in temperature is not large, but bringing $\mu_e$ above the resonance energy increases the peak in the spectrum by more than 4 orders of magnitude. Figure \ref{fig:56fe_cc_multi} shows the $^{56}$Fe electron neutrino energy spectra computed using the FFN prescription at several points during collapse. The solid lines are for electron capture, the dotted lines are for positron decay, and the colors correspond to different temperature and density conditions. The results of figures \ref{fig:56fe_cc} and \ref{fig:56fe_cc_multi} are qualitative (strength is unquenched, delta function resonance, etc.) but indicate that at high $\mu_e$, the distribution of the bulk of the strength dominates the effects of precise structure. \begin{figure}[here] \includegraphics[scale=.42]{56fe_cc_multi.pdf} \caption{$^{56}$Fe charged current process neutrino spectra computed from the FFN prescription at various points during collapse. The enormous jump in neutrino production between the lowest two temperatures is due to the chemical potential in the higher temperature point being greater than the Gamow-Teller resonance energy.} \label{fig:56fe_cc_multi} \end{figure} Given the obvious importance of nuclear contributions to neutrinos with detectable energies, we will move forward in generating tabulated nuclear neutrino energy spectra in the same vein as the neutrino production and energy loss rates of earlier works. | 16 | 7 | 1607.01448 |
1607 | 1607.08043_arXiv.txt | It is postulated in Einstein's relativity that the speed of light in vacuum is a constant for all observers. However, the effect of quantum gravity could bring an energy dependence of light speed. Even a tiny speed variation, when amplified by the cosmological distance, may be revealed by the observed time lags between photons with different energies from astrophysical sources. From the newly detected long gamma ray burst GRB~160509A, we find evidence to support the prediction for a linear form modification of light speed in cosmological space. | 16 | 7 | 1607.08043 |
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1607 | 1607.06097_arXiv.txt | We present the large-scale 3-point correlation function (3PCF) of the SDSS DR12 CMASS sample of $777,202$ Luminous Red Galaxies, the largest-ever sample used for a 3PCF or bispectrum measurement. We make the first high-significance ($4.5\sigma$) detection of Baryon Acoustic Oscillations (BAO) in the 3PCF. Using these acoustic features in the 3PCF as a standard ruler, we measure the distance to $z=0.57$ to $1.7\%$ precision (statistical plus systematic). We find $\DV = 2024\pm29\;{\rm Mpc\;(stat)}\pm20\;{\rm Mpc\;(sys)}$ for our fiducial cosmology (consistent with {\it Planck} 2015) and bias model. This measurement extends the use of the BAO technique from the 2-point correlation function (2PCF) and power spectrum to the 3PCF and opens an avenue for deriving additional cosmological distance information from future large-scale structure redshift surveys such as DESI. Our measured distance scale from the 3PCF is fairly independent from that derived from the pre-reconstruction 2PCF and is equivalent to increasing the length of BOSS by roughly 10\%; reconstruction appears to lower the independence of the distance measurements. Fitting a model including tidal tensor bias yields a moderate significance ($2.6\sigma)$ detection of this bias with a value in agreement with the prediction from local Lagrangian biasing. | \label{sec:intro} Determining the nature of dark energy is one of the most pressing problems of modern cosmology. Efforts have focused on measuring the dark energy equation of state $w$, which is $-1$ if dark energy is a cosmological constant; any other value of $w$ means the dark energy density evolves in time (Copeland, Sami \& Tsujikawa 2006). The dark energy density dictates the Universe's expansion through the Friedmann equation, and so measuring the Universe's size as a function of time or redshift constrains the equation of state. A number of techniques exist to do this (Weinberg et al. 2013), one of the most prominent being the Baryon Acoustic Oscillation (BAO) method. The BAO method exploits a preferred scale imprinted on the baryon density at decoupling ($z\sim1020$). Prior to decoupling, the Universe is a hot, dense, ionized plasma in which electrons couple to photons by Thomson scattering and protons follow electrons under the Coulomb force. Primordially overdense regions are overpressured, and the radiation pressure, dominant at high redshift, launches spherical pressure-density (sound) waves of baryons and photons outwards from each overdensity at roughly $c/\sqrt{3}$ (Sakharov 1966; Peebles \& Yu 1970; Sunyaev \& Zel'dovich 1970; Bond \& Efstathiou 1984, 1987; Holtzmann 1989; Hu \& Sugiyama 1996; Eisenstein \& Hu 1998; Eisenstein, Seo \& White 2007). These are the BAO, and the sound waves travel outwards until decoupling, where they halt as the photons precipitously release them since Thomson scattering is no longer effective. At the wavefront the baryon velocity is maximal and the late-time growing mode inherits the spatial structure of the velocity. The BAO thus correlate the original overdensity with a sharp excess density of baryons a sound horizon (roughly $\rs \approx 100\Mpch$ comoving) away. Once the Universe is neutral, on large scales the baryons and the dark matter experience only gravity, so the two components converge and the excess density of baryons imprints on the total matter density (Hu \& Sugiyama 1996; Eisenstein \& Hu 1998; Slepian \& Eisenstein 2016a). When galaxies begin to form, they trace the matter density field and so the BAO produce a slight excess of galaxy pairs separated by $\sim 100\Mpch$. This excess translates to a sharp, localized BAO bump in the 2-point correlation function (2PCF) of galaxies, which measures the excess probability over random of finding one galaxy at a given separation from another; there are analogous BAO features in the 2PCF's Fourier-space analog the power spectrum. Since the BAO signal is produced by large-scale, pre-decoupling physics, it is frozen into the comoving distribution of galaxies. Consequently measuring the BAO scale from galaxy clustering in different redshift slices provides a differential history of the Universe's expansion rate (Eisenstein, Hu \& Tegmark 1998; Blake \& Glazebrook 2003; Hu \& Haiman 2003; Linder 2003; Seo \& Eisenstein 2003). The BAO scale is also imprinted on the temperature anisotropies in the Cosmic Microwave Background (CMB), since the density structure at that epoch determined the temperature. The CMB therefore offers an absolute scale for the BAO method. Thus far, the BAO method has used the 2PCF of galaxies as well as the galaxy power spectrum to measure the cosmic distance scale to high precision. Since the original detections of the BAO bump in the 2PCF of galaxies (Cole et al. 2005; Eisenstein et al. 2005), large-scale redshift surveys such as the Sloan Digital Sky Survey (SDSS) and Baryon Oscillation Spectroscopic Survey (BOSS) have yielded ever-increasing precision via the BAO method. The current precision on the distance scale from the 2PCF/power spectrum is of order $1\%$ (Anderson et al. 2014; Cuesta et al. 2016; Gil-Mar\'in et al. 2016), and future surveys such as Dark Energy Spectroscopic Instrument (DESI) (Levi et al. 2013) and {\it Euclid} (Laureijs et al. 2011) should achieve a factor of five improvement in precision. The Lyman-$\alpha$ forest has also been used for BAO measurements, with the most recent results in Delubac et al. (2015). The first detection of BAO in voids has also recently been made, offering an additional possible avenue to the distance scale (Kitaura et al. 2016). Until now the BAO method has not explicitly used higher correlations of the galaxy density field. As earlier noted, the BAO produce an excess of pairs of galaxies separated by $100\Mpch$, but the BAO also imprint on triplets of galaxies, creating a slight excess of triangles where one or more triangle side is of the BAO scale. Triplets develop correlations both due to non-linear structure formation and non-linear bias. Slepian \& Eisenstein (2016b; hereafter SE16b) shows that there are distinctive BAO features in the 3-point correlation function (3PCF) of galaxies. Detecting these features would enable a measurement of the cosmic distance scale from the 3PCF alone. Thus far, only two previous works have measured the 3PCF on physical scales large enough to access the BAO scale. Gazta\~{n}aga et al. (2009) used a sample of $\sim 40,000$ Luminous Red Galaxies (LRGs) from SDSS DR7. They find a $2-3\sigma$ detection of the BAO using all opening angles of a single triangle configuration with side lengths $r_1 = 33 \Mpch$ and $r_2 = 88\Mpch$. Slepian et al. (2015; hereafter S15) used $777,202$ LRGs from the CMASS sample within SDSS-III BOSS to measure the 3PCF in a compressed basis where many triangle configurations were used but one of the two sides was integrated out over a wedge set by the remaining free side. That work found a $2.8\sigma$ detection of the BAO. Given the larger sample of S15, by comparison to Gazta\~naga et al. (2009) a higher significance BAO detection might be expected, suggesting that there is BAO information the compressed basis does not exploit. On the other hand, Gazta\~naga et al. (2009) did find an anomalously high baryon fraction (roughly double the presently-accepted value), which would increase the significance of a BAO detection. Neither work used these moderate-significance BAO detections to measure the cosmic distance scale. In this work, we use the same dataset and 3PCF measurement as S15. However, we do not compress by integrating out one triangle side. The compression scheme of S15 was motivated by avoiding any triangle side's becoming small and two galaxies becoming close, where linear perturbation theory is likely a poor model. Here, we avoid this limit by choosing triangle sides such that the smallest side never is below $20\Mpch$. We again use the novel algorithm of Slepian \& Eisenstein (2015b,c; hereafter SE15b,c), which computes the 3PCF's multipole moments in $\oO(Nn\Vmax)$ time using spherical harmonic decompositions, where $\Vmax$ is the volume of a sphere of radius $\Rmax$, the maximum triangle side length to which correlations are measured. The covariance matrix also turns out to be tractable in the multipole basis (SE15b). The main outcomes of this work are: \indent{\bf 1)} The first high-significance $(\sim 4.5\sigma)$ detection of the BAO in the 3PCF.\\ \indent {\bf 2)} A measurement of the cosmic distance scale at redshift $0.57$ to $1.7\%$ precision from the 3PCF.\\ \indent {\bf 3)} High precision ($\sim 1\%$) determination of the linear bias at fixed $\sigma_8$ for this sample from the 3PCF.\\ An interesting subsidiary result of this work is that the tidal tensor bias $b_t$ (further detailed in \S\ref{sec:models}) of the dataset agrees well with the theoretically predicted relation with linear bias $b_1$, $b_t = -(2/7)[b_1-1]$ (Baldauf et al. 2012; Chan, Scoccimarro \& Sheth 2012), offering mild evidence for the validity of local Lagrangian biasing. In contrast, $b_t$ for the {\textsc PATCHY} mock catalogs for SDSS DR12 (described further in \S\ref{sec:data_randoms_mocks}) does not agree with this theoretical relation. With our work's error bars on $b_t$, the tension between mocks and data is only mild, but this possible misfit between the data and the {\textsc PATCHY} mocks as well as the {\textsc PATCHY} mocks and the theory may warrant further investigation. The paper is laid out as follows. \S\ref{sec:data_randoms_mocks} details our dataset, the random catalogs used for edge correction, and the mock catalogs used to obtain parameters within the covariance matrix as well as to verify our pipeline. \S\ref{sec:method} summarizes the multipole basis we use for the 3PCF as well as the algorithm of SE15b used for the measurement, while \S\ref{sec:covar} discusses our covariance matrix. In \S\ref{sec:models} we outline the two different bias models we use to analyze the data, a ``minimal'' model that includes linear and non-linear biasing, and a ``tidal tensor'' model that includes these elements and also tidal tensor biasing. \S\ref{sec:fitting_procedure} details our parameter-fitting procedure, and \S\ref{sec:BAO_fit} presents our BAO detection and best-fit parameters for the data and mocks. \S\ref{sec:distance_scale} gives our distance scale measurement in physical units and compares with other recent works, while \S\ref{sec:bias_parameters} discusses our measured bias parameters. We conclude in \S\ref{sec:concs}. | \label{sec:concs} In this work, we have used the novel 3PCF algorithm of SE15b to compute the 3PCF of $777,202$ LRGs from the CMASS sample of SDSS DR12. This is the largest sample used for the 3PCF or bispectrum to date. Using full triangles for a set of bins selected to avoid the regime where linear perturbation theory breaks down, we make the first high-significance detection of the BAO (respectively $4.5\sigma$ and $4.4\sigma$ for the two bias models we fit). Our previous work S15 measured the 3PCF of this sample in a compressed basis that integrated one triangle side over a wedge set by the other; while this approach had several physical motivations, it was likely sub-optimal for detecting BAO features. That previous work found a $2.8\sigma$ preference for the BAO. With the present work's high significance BAO detection, we use the 3PCF to constrain the cosmic distance scale $\DV=2024\;{\rm Mpc}+29\;{\rm Mpc\;(stat)}+20\;{\rm Mpc\;(sys)}$ to redshift $z=0.57$ with $1.7\%$ precision (statistical plus systematic). This distance measurement is the first use of the BAO method with the 3PCF. We briefly explore the independence of the distance scale measured from the 3PCF relative to that measured from the 2PCF prior to reconstruction, and find they are essentially entirely independent. However, reconstruction is known to introduce some distance information from the 3PCF back into the 2PCF, and further work will explore the impact of reconstruction on the independence of the distance scales. For the moment, we note that adding our distance scale to that measured from the unreconstructed 2PCF would produce a distance measurement with $1\%$ precision, comparable to the most recent precision of the distance scale measured from the reconstructed 2PCF (Cuesta et al. 2016) or power spectrum (Gil-Mar\'in et al. 2016). Holding $\sigma_8$ fixed, we also place an extremely precise constraint on the linear bias, measuring $b_1$ to sub-percent accuracy in our minimal bias model. This constraint is the most precise placed on the linear bias from either the 3PCF or the bispectrum, and is competitive with the most precise constraints on $b_1$ placed using any other techniques. Finally, we have made a moderate-confidence detection ($2.6\sigma$) of tidal tensor bias in agreement with the prediction of local Lagrangian biasing. Our error bars on the tidal tensor bias remain large because the tidal tensor bias is highly degenerate with the linear bias as they enter the 3PCF. In future it will be worthwhile to explore avenues for breaking this degeneracy. The mocks do not appear to match the tidal tensor model well, but we do not expect this affects the BAO significance. In particular, the BAO significances are similar between our minimal and tidal tensor models, so the tidal tensor biasing is likely not a substantial driver of the BAO significance. We note that the bispectrum has been used before to test tidal tensor biasing on observational data; Feldman et al. (2001) found very mild evidence in favor of Eulerian biasing ($b_t = 0$) for the IRAS PSCz galaxy catalog. This galaxy population is very different from the LRGs considered in this work and so there is no conflict with the moderate-confidence detection of tidal tensor bias we report here. In a companion paper (Slepian et al. 2016b), we use the same 3PCF dataset analyzed here with a different bias model to look for the imprint of baryon-dark matter relative velocities (Tseliakhovich \& Hirata 2010) in the late-time clustering of galaxies. This imprint can be an important possible systematic for BAO measurements from the 2PCF or power spectrum (Yoo, Dalal \& Seljak 2011), and as discussed in Slepian \& Eisenstein (2015a) has a unique signature in the 3PCF. The companion paper finds that the 3PCF offers a $\sim 0.3\%$ constraint on any possible shift the relative velocity induces in the BAO scale inferred from the 2PCF, arguing for the 3PCF's utility in protecting future redshift surveys from this possible bias. In sum, the same spectroscopic data sets currently used for 2PCF analyses with the BAO method can be used for 3PCF analyses with the BAO method. We hope the 3PCF will offer a new avenue to the cosmic distance scale. Used in conjunction with the 2PCF, we believe the 3PCF can increase the cosmological leverage of a given survey. With upcoming efforts such as DESI providing of order $30$ million spectra (Levi et al. 2013), taking full advantage of the BAO information in the 3PCF as well as that in the 2PCF will be highly desirable. | 16 | 7 | 1607.06097 |
1607 | 1607.08275_arXiv.txt | Dark | \label{sec:int} In Section 2, we briefly describe the MOCCA code for star cluster simulations, the COCOA and SISCO codes used to simulate observational data for our simulated cluster model. We provide the details of the dark cluster model that we simulated with MOCCA in Section 3. The section also shows the results of the simulated observations of our cluster model from COCOA and SISCO. The simulated observations are compared to the Galactic GC NGC 6535. The last section deals with the conclusion and discussion of our results. | The last numbered section should briefly summarise what has been done, and describe the final conclusions which the authors draw from their work. | 16 | 7 | 1607.08275 |
1607 | 1607.07744_arXiv.txt | We present a detailed analysis of the ionization and thermal structure of the intergalactic medium (IGM) around a high-redshift QSO using a large suite of cosmological, multi-frequency radiative transfer (RT) simulations, exploring the contribution from galaxies as well as the QSO, and the effect of X-rays and secondary ionization. We show that in high-$z$ QSO environments both the central QSO and the surrounding galaxies concertedly control the reionization morphology of hydrogen and helium and have a non-linear impact on the thermal structure of the IGM. A QSO imprints a distinctive morphology on $\HII$ regions if its total ionizing photon budget exceeds that of the surrounding galaxies since the onset of hydrogen reionization; otherwise, the morphology shows little difference from that of $\HII$ regions produced only by galaxies. In addition, the spectral shape of the collective radiation field from galaxies and QSOs controls the thickness of the I-fronts. While a UV-obscured QSO can broaden the I-front, the contribution from other UV sources, either galaxies or unobscured QSO, is sufficient to maintain a sharp I-front. X-rays photons from the QSO are responsible for a prominent extended tail of partial ionization ahead of the I-front. QSOs leave a unique imprint on the morphology of $\HeII/\HeIII$ regions. We suggest that, while the physical state of the IGM is modified by QSOs, the most direct test to understand the role of galaxies and QSOs during reionization is to perform galaxy surveys in a region of sky imaged by 21 cm tomography. | In the current paradigm of extragalactic astronomy, the Universe has undergone two major reionization epochs. $\HI$ and $\HeI$ reionization are thought to occur at $z \sim 6-20$, driven by early star forming galaxies and/or quasars (QSOs) (e.g. \citealt{2015ApJ...802L..19R,2015ApJ...813L...8M}). This first reionization epoch contains information about the formation of the first luminous objects in the Universe, a period called Cosmic Dawn. The second reionization epoch, from $\HeII$ to $\HeIII$, occurs later time, at $z\sim2-4$, and it is likely driven by the high quasar activity near the peak of cosmic star formation history. Understanding the full reionization process of intergalactic hydrogen and helium will provide a milestone in investigating the nature of high-redshift galaxies and QSOs and of their interaction with the intergalactic medium (IGM). Current observations of the $\HI$ $\LyA$ Gunn-Peterson trough suggest that hydrogen reionization is largely completed by $z \sim 6$ (e.g. \citealt{2006AJ....132..117F}). However, the nature of the sources that drive hydrogen reionization is still unknown. A concordance picture is one in which (undetected) faint star forming galaxies with an escape fraction of ionizing photons as large as $\sim20$ per cent are the main reionizing agents (e.g. \citealt{2015ApJ...802L..19R}). At lower redshift, observations of the $\HeII$ $\LyA$ Gunn-Peterson trough suggest that $\HeII$ reionization is mostly completed by $z\sim2.5$ (\citealt{2011ApJ...733L..24W,2011ApJ...726..111S}). Unlike $\HI$ reionization, $\HeII$ reionization is likely driven by the already detected population of QSOs (\citealt{2015A&A...578A..83G}). In fact, simulations calibrated with the observed QSO luminosity function can grossly reproduce the observed properties of the $\HeII$ $\LyA$ forest (\citealt{2009ApJ...694..842M,2013MNRAS.435.3169C,2014MNRAS.445.4186C}). As hydrogen and helium reionization have a strong impact on the thermal state of the IGM (\citealt{2006ApJ...644...61L,2011MNRAS.415..977M,2015MNRAS.447.2503G}), constraints on the reionization process can be obtained by measurements of the IGM temperature. Such measurements from the $\LyA$ forest (\citealt{2011MNRAS.415..977M,2014ApJ...788..175L}) suggest that the end of $\HI$ reionization should be at a redshift lower than $z\simeq9$ (\citealt{2002ApJ...567L.103T,2010MNRAS.406..612B,2012MNRAS.421.1969R}), while $\HeII$ reionization should be extended over $2\leq z\leq4.8$ (\citealt{2011MNRAS.410.1096B}). Furthermore, the IGM temperature in QSO near zones at $z\sim6$ suggests that the onset of $\HeII$ reionization may be as early as the formation of the first QSOs at $z>6$ (\citealt{2012MNRAS.419.2880B}). Although many authors have studied the impact of QSOs on the ionization and thermal state of the IGM (e.g. \citealt{1987ApJ...321L.107S,1993ApJ...412...34M,1999ApJ...514..648M,2000ApJ...542L..75C,2000ApJ...530....1M,2003ApJ...586..693W,2004ApJ...604..484M,2005ApJ...620...31Y,2005ApJ...623..683Y,2006ApJ...648..922S,2007ApJ...657...15K,2007MNRAS.376L..34M,2007MNRAS.380L..30A,2007MNRAS.380.1369T,2007ApJ...670...39L,2008MNRAS.384.1080T,2008MNRAS.385.1561K,2008ApJ...686...25F,2012MNRAS.424..762D,2013MNRAS.429.1554F,2016MNRAS.455.2778F,2015MNRAS.454..681K}), the role of QSOs during $\HI$ reionization at $z>6$ still remains unclear. While it has been argued that QSOs alone cannot be responsible for hydrogen reionization based on constraints from the unresolved X-ray background (\citealt{2004ApJ...613..646D,2005MNRAS.362L..50S,2007MNRAS.374..761S,2015A&A...575L..16H}) and the decreasing number density of high-$z$ QSOs (\citealt{2005MNRAS.356..596M}), recent re-investigations suggest a possible important contribution (e.g. \citealt{2011ApJ...728L..26G,2012MNRAS.425.1413F,2015A&A...578A..83G,2015ApJ...813L...8M}). Furthermore, theoretical predictions about how the QSOs impact on the ionization and thermal state of the local environment do not always agree. For example, there seems to be no consensus about whether a QSO imprints a distinctive ionization and thermal structure on the IGM by producing a very large ionized region (e.g. \citealt{2013MNRAS.429.1554F,2007MNRAS.375.1269Z,2008MNRAS.384.1080T}), if it can be distinguished from one produced only by galaxies (\citealt{2007MNRAS.380L..30A,2007ApJ...670...39L,2012MNRAS.424..762D}), if the spectra of galaxies and QSOs induce a different shape of the ionization front (I-front) (\citealt{2005MNRAS.360L..64Z,2008MNRAS.385.1561K}), or whether the sphericity of an ionized region may serve as discriminator of the type of source that created it (\citealt{2012MNRAS.424..762D}). These theoretical discrepancies must be resolved to interpret QSO absorption spectra and upcoming 21 cm observations. In other words, a correct theoretical understanding is vital to place constraints on the role of galaxies and QSOs in driving reionization. Understanding the impact of ionization from galaxies and/or QSOs on the physical state of the IGM relies on an accurate modeling of the effects of (\rmnum{1}) the spectral shape of QSOs and galaxies (i.e. multi-frequency radiative transfer [RT]), (\rmnum{2}) galaxies surrounding a QSO (i.e. cosmological $N$-body/hydrodynamical simulations), (\rmnum{3}) an anisotropic propagation of the I-front (i.e. 3D simulation), and (\rmnum{4}) a coherent treatment of both UV and X-ray photons as well as secondary ionization (i.e. UV/X-ray physics). Previous works addressed the above aspects separately. For example, the 1D RT simulations by \cite{2008MNRAS.384.1080T} and \cite{2008MNRAS.385.1561K} focused on (\rmnum{1}) and (\rmnum{4}), whereas the 3D RT simulation by \cite{2012MNRAS.424..762D} including both galaxies and QSOs, but no thermal structure, focused on addressing (\rmnum{1}), (\rmnum{2}) and (\rmnum{3})\footnote{\cite{2013MNRAS.429.1554F} and \cite{2015MNRAS.454..681K} have also conducted RT simulation post-processing cosmological hydrodynamical simulation, but without simultaneously accounting for both galaxies and QSOs.}. At the time of writing, no radiative transfer calculation addressing all four points in a single simulation is reported. In this work we therefore investigate the ionization state of both hydrogen and helium, and the thermal state of the IGM in the environment of a high-$z$ QSO by performing a suite of multi-dimensional, multi-frequency radiative transfer simulations post-processing a cosmological hydrodynamical simulation. This is to our knowledge the most detailed calculation of this kind to date. This work distinguishes itself from previous investigations, as we present the case-by-case analysis of hydrogen and helium reionization, as well as the thermal state of the environment of a QSO at $z=10$, including the effect of surrounding galaxies, X-ray and secondary ionization. This large suite of RT simulations aims at providing insights into the underlying physical mechanisms responsible for controlling the physical state of the IGM around QSOs, and on how all processes collectively shape it. Our work is an important first step towards full cosmological reionization simulations including both galaxies and multiple QSOs. The paper is organized as follows. First we describe the simulation setup in \S~\ref{sec:sim}. In \S~\ref{sec:H_reion}, we present the results of hydrogen and helium reionization. In \S~\ref{sec:thermal}, we discuss the thermal state in the QSO environment. We compare our results with previous works in \S~\ref{sec:comparison}. Observational implications are briefly discussed in \S~\ref{sec:obs}. Conclusions are then presented in \S~\ref{sec:conclusion}. | \label{sec:conclusion} We have presented a detailed analysis of hydrogen and helium reionization scenarios in high-$z$ QSO environments including both galaxies and a central QSO using a suite of cosmological multi-frequency radiative transfer simulations. This allows us to understand the array of physical mechanisms that shapes the ionization and thermal states of the IGM. We find that: \begin{itemize} \item if the integrated number of ionizing photons emitted by the surrounding galaxies since the onset of reionization exceeds or is comparable to that emitted by the QSOs during its active phase, the morphology of the $\HII$ regions in galaxies only and galaxies+QSO models is similar. On the other hand, if the ionization budget is dominated by the QSO, distinctive features can be observed in the $\HII$ region. \item The hard spectrum of the QSO always leaves a unique signature in the morphology of the $\HeII$ and $\HeIII$ regions. \item Soft X-ray photons from the QSOs produce extended partially ionized tails of hydrogen and helium around the central $\HII$ and $\HeII/\HeIII$ regions. Ionization from secondary electrons enhances the contribution of photoionization by X-rays in the partially ionized tail. On the other hand, neither the morhopology of the fully ionized regions nor the thinckness of the I-fronts are sensitive to the X-ray photons. \item UV-obscured QSOs can broaden the I-fronts significantly. A noticeable change in the ionization morphology occurs when soft X-ray photons (from the QSO), rather than UV photons, dominate the growth of I-fronts. \item The thermal state of the IGM is strongly affected in a complex, non-linear, way by photoionization heating of both hydrogen and helium. A larger amount of total ionizing photons does not necessarily increase the gas temperature. The highest temperature is attained in the region ionized only by the QSO, as the gas temperature is primarily determined by the spectral type of the sources that first ionized the medium. \item The X-ray photons from the QSO heat the partially ionized gas ahead of the I-fronts to $T\sim10^3\rm~K$. The net X-ray heating is slightly reduced when secondary ionization is included because it adds another channel to deposit the energy into ionization. The X-ray pre-heating is a clean signature of QSOs activity. \end{itemize} In summary, in the environment surrounding a high-$z$ QSO, {\it the physical state of the IGM in terms of hydrogen and helium reionization as well as gas temperature, is determined by a complex and non-linear interplay of the galaxies and the QSO}. This picture emphasizes the importance of correctly modelling hydrogen, helium and temperature evolution with a multi-frequency treatment of the radiative transfer in 3D numerical simulations. The large suite of multi-frequency RT simulations presented here provides a highly valuable theoretical resource to understand the physical mechanisms responsible for shaping the ionization and thermal structures of the IGM. This will facilitate the use of RT simulations in interpreting QSO spectra and upcoming 21 cm observations to better understand the role of galaxies and QSOs in driving reionization. | 16 | 7 | 1607.07744 |
1607 | 1607.07869_arXiv.txt | I show that the Reduced Speed of Light (RSL) approximation, when used properly (i.e.\ as originally designed - \emph{only for the local sources but not for the cosmic background}), remains a highly accurate numerical method for modeling cosmic reionization. Simulated ionization and star formation histories from the ``Cosmic Reionization On Computers'' (CROC) project are insensitive to the adopted value of the reduced speed of light for as long as that value does not fall below about 10\% of the true speed of light. A recent claim of the failure of the RSL approximation in the Illustris reionization model appears to be due to the effective speed of light being reduced in the equation for the cosmic background too, and, hence, illustrates the importance of maintaining the correct speed of light in modeling the cosmic background. | 16 | 7 | 1607.07869 |
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1607 | 1607.02950_arXiv.txt | { We present simulated observations of galaxies at $z=2$ and $z=3$ to probe the capabilities of next-generation telescopes (E-ELT and JWST) to measure the structural and photometrical properties of high-redshift galaxies. We carry out an extensive set of simulations of high-redshift galaxies adopting the specifications of the E-ELT first light instrument MICADO. The main parameters (sizes, Sersic index, and magnitudes) of the galaxies are measured using GALFIT and the accuracy of the determinations is assessed by comparing the input values to the measurements from many runs with different statistical noise. We also address the effects on the accuracy of the measurements of possible spatial variation of the point spread function (PSF) in the field. We find that from $3 h$ exposure E-ELT near-infrared (IR) images of galaxies at $z \sim 2$ and $z \sim 3$ it will be possible to measure the size, total magnitude, and galaxy morphology with an accuracy of 2-5\% for objects as faint as $H\sim$ 25 and half-light size of 0.2 arcsec. The effective radius of compact, early-type galaxies is also recovered with $\sim 5$\% accuracy, provided that their half-light size exceeds 20 mas. These results are compared with those expected from simulated observations obtained with NIRCam on board the JWST.} | \label{sec:intro} Understanding the assembly history of galaxies is of paramount importance for answering fundamental questions about the processes of formation and evolution of galaxies and their associations (groups, clusters, superclusters, etc). To this end, it is of extreme relevance to be able to characterise the properties of galaxies at high redshift to probe their evolution over a significant interval of cosmic time. Their global structure and colours yield insight into the conditions of star formation and the subsequent merger events and/or secular processes \citep[e.g.][]{dalcanton+1997,mo+1998}. Although galaxies usually comprise multiple different components (i.e. bulges, disks, substructures) fitting their surface brightness profile with a single Sersic law can provide relevant basic data \citep[e.g.][]{kelvin+2012}. This is in particular needed for the study of distant galaxies that are poorly resolved with current ground-based telescopes and partially available with Hubble Space Telescope \cite[HST; e.g.][]{vand+2012}. Unfortunately the characterisation of photometrical and structural properties of high-$z$ objects is hampered by the faintness of the targets and by their very small angular size. Hubble Space Telescope observations have shown that galaxies at high-$z$ are significantly smaller in size than galaxies of similar mass at low $z$ \citep[e.g.][]{dadd+2005,truj+2006,buit+2008,vulcani+2014,kennedy+2015} and this size evolution ($R_e\sim (1+z)^{-1}$) with redshift, furthermore, hinders the capability of studying the structural properties of these galaxies. In fact present estimates of galaxy size at $z\gtrsim2$ yield apparent sizes of $\sim 100-200$ mas or less. At higher redshift the situation is even more challenging since the angular size of the galaxies become smaller than 50 mas \citep{ono+2013}. As a consequence, it is now possible to characterise these galaxies only via HST observations, but this is limited to galaxies with mass $\gtrsim10^{10}M_\sun$\citep[][and refs therein]{bram+2012,vand+2012,vand+2014}. In spite of the diffraction limit images obtained by HST, the size of these high-redshift galaxies is so small that a significantly better spatial resolution is needed to be able to properly characterise their structure. A major improvement in this direction is expected by future observations gathered by the James Webb Space Telescope (JWST), which also has the important advantage of an extremely low background in the near-infrared (IR). Because of the relative small aperture (6.5m), however, the JWST resolution is limited to few hundredths of arcsec. A more significant step is expected by the future generation of extremely large ground-based optical/near-IR telescopes. In fact these 30-40 m aperture telescopes assisted by adaptive optics can reach resolution of few mas in the near-IR. In this paper we aim to address the imaging capabilities of the planned future instrumentation for Extremely Large Telescopes (ELTs) using the expected performances of the Multi-Adaptive Optics Imaging Camera for Deep Observations \citep[MICADO;][]{davi+2010} at the European Extremely Large Telescope (E-ELT)\footnote{\url{http://www.eso.org/sci/facilities/eelt/}} as a reference. We quantify the accuracy with which it will be possible to characterise the properties of high-redshift galaxies with simulated MICADO observations. In particular these simulations are used to investigate the possibility to accurately measure both the size and morphology of the galaxies. Moreover using multi-band observations we estimate the accuracy that could be achieved in the measurements of the colour gradient; this is a key tool to test models of galaxies formation. In fact the colour gradients trace variations of the properties of stellar populations, and are linked to the star formation history of galaxies \citep[see e.g][]{sagl+2000,tamu+2000,koba+2004,garg+2011}. Finally, the results obtained with simulated E-ELT observations are compared with those expected for the same targets from observations secured with the Near-Infrared camera (NIRcam) on board the JWST\footnote{\url{http://www.stsci.edu/jwst/instruments/nircam}}. Throughout this paper, we adopt a concordance cosmology with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_m = 0.3$, and $\Omega_\Lambda= 0.7$. All magnitudes are in AB system. | In this paper we presented a detailed analysis of the expected performances of next-generation ELTs in the characterisation of the properties of high-$z$ galaxies. We evaluated the accuracy in the measurements of the structural and photometric parameters that will be obtained from future observations using an extensive set of simulated observations. As a reference, we used the specifications of the MICADO camera to be mounted at the E-ELT and the NIRCAM at the JWST. The extraordinary large diameter of the E-ELT will provide imaging of high-redshift galaxies with unprecedented spatial resolution. Our results confirm that MICADO will provide extremely accurate measurements of the structural parameters of high-redshift galaxies; systematic uncertainties in the measurements of structural parameters are expected to be negligible for galaxies with effective radii larger than $\sim20$ mas. The overall uncertainties in the measurements of the structural parameters of late-type galaxies result to be smaller than $5\%$ and the uncertainties in the total magnitude of a few hundredths of magnitude, even for the smallest simulated galaxies ($M=10^9 M_\sun$, $H\sim$24.5 mag). For early-type galaxies that are more compact and have a steeper profile ($n=4$) than late-type galaxies, we found that MICADO would provide measurements of the effective radius and the Sersic index with uncertainties of just 10--20\% even for the smallest and faintest galaxies considered in our study ($M=10^9 M_\sun$, $H\sim$26.5 mag). These results confirm that MICADO observations will represent a real breakthrough as they will facilitate results far beyond the capabilities of present-day instrumentation. Today, this accuracy is obtained with the HST for galaxies $\sim 3$ mag brighter and at least 10 times more massive and/or bigger \citep[see e.g.][]{vand+2014}. We also estimate that, from ELTs observations, it will be possible to obtain reliable measures of colour gradients as small as 0.3 mag/dex in galaxies at $z=3$. Therefore it will be possible to study the population gradients in high-redshift galaxies. This will represent a fundamental step to understand the assembly history and physical processes regulating the formation, the growth and, evolution of the galaxies. This study is nowadays possible only for nearby galaxies; the advent of the ELTs will make it possible to extend these studies at $z=2$ and 3, probing galaxies at the epoch in which the star formation rate density peaked \citep[$z\simeq2$; e.g.][]{mada+2014} and beyond. We also provided a quantitative estimate of the performances expected for the JWST. We found that NIRCam will provide excellent measurements for late-type galaxies down to $M\sim10^{10} M_\sun$ at $z$=2. The size of low-mass early-type galaxies ($M\lesssim10^{10}M_\sun$) at $z=3$ would be of the order of (or even smaller than) the spatial resolution of the JWST and therefore it would be extremely difficult to derive reliable measurements of their spatial and photometric parameters. The first opportunity to explore the details of the structure of distant compact galaxies, or the very inner region of galaxies at lower redshift, will be offered by the advent in the next decade of the imagers at ground-based extremely large telescopes and, in particular, of MICADO at the E-ELT. We finally performed simulations of high-$z$ galaxies using high-resolution HST images of galaxies at low redshift. These simulations manifest how the combined capabilities of excellent angular resolution and sensitivity will allow us to investigate in great detail the small substructures (spiral arms, clumpy star-forming regions, dust lanes, etc.) of distant galaxies that could not be investigated by any other present or actually planned ground- or space-based facilities. | 16 | 7 | 1607.02950 |
1607 | 1607.07058_arXiv.txt | Spectator field models such as the curvaton scenario and the modulated reheating are attractive scenarios for the generation of the cosmic curvature perturbation, as the constraints on inflation models are relaxed. In this paper, we discuss the effect of Hubble induced masses on the dynamics of spectator fields after inflation. We pay particular attention to the Hubble induced mass by the kinetic energy of an oscillating inflaton, which is generically unsuppressed but often overlooked. In the curvaton scenario, the Hubble induced mass relaxes the constraint on the property of the inflaton and the curvaton, such as the reheating temperature and the inflation scale. We comment on the implication of our discussion for baryogenesis in the curvaton scenario. In the modulated reheating, the predictions of models e.g.~the non-gaussianity can be considerably altered. Furthermore, we propose a new model of the modulated reheating utilizing the Hubble induced mass which realizes a wide range of the local non-gaussianity parameter. | The large scale structure of the universe and the fluctuation of the cosmic microwave background (CMB) are naturally explained by the perturbation generated during cosmic inflation~\cite{Mukhanov:1981xt,Hawking:1982cz,Starobinsky:1982ee,Guth:1982ec,Bardeen:1983qw}. In the simplest scenario, a scalar field called an inflaton is responsible both for the inflation and the generation of the cosmic perturbation. In general, any light scalar field obtains a fluctuation of its field value during inflation and can be the origin of the cosmic perturbation, even if their energy density is negligible during inflation. We refer to such field as a spectator field. In the curvaton scenario~\cite{Mollerach:1989hu,Linde:1996gt,Enqvist:2001zp,Lyth:2001nq,Moroi:2001ct}, the fluctuation turns into a fluctuation of the energy density of the universe as the spectator field dominates the energy density of the universe. In the modulated reheating~\cite{Dvali:2003em,Kofman:2003nx}, the decay rate of the inflaton (or any fields dominating the energy density of the universe) depends on the field value of the spectator field, and hence the timing of the reheating is modulated. Spectator field models are attractive since the inflaton field itself is no more responsible for the cosmic perturbation, and hence constraints on inflation models are relaxed. For example, the chaotic inflation~\cite{Linde:1983gd} is one of the most attractive scenarios as it is free from the initial condition problem~\cite{Linde:2005ht}. The chaotic inflation, however, generically predicts a large tensor fraction and the simplest model is disfavored by the observations of the CMB~\cite{Ade:2015lrj}. In spectator field models, this constraint is evaded In this paper, we discuss the influence of a Hubble induced mass after inflation on the dynamics of the spectator field. We focus on the curvaton and the modulated reheating scenario. A Hubble induced mass of a field is given by the coupling of the field to the potential or kinetic energy of the inflaton. As the potential and the kinetic energy terms are neutral under any symmetry, the Hubble induced mass is in general unsuppressed, unless the field has an approximate shift symmetry. Actually, the Hubble induced mass is a generic feature in supersymmetric models~\cite{Ovrut:1983my,Holman:1984yj,Goncharov:1983mw,Coughlan:1984yk}. The Hubble induced mass of the spectator field during inflation, namely the coupling to the potential energy of the inflaton, must be suppressed to explain the almost scale invariant spectrum of the curvature perturbation. However, the coupling of the spectator field to the kinetic energy of the inflaton is expected to be unsuppressed. Such coupling generates a sizable Hubble induced mass after inflation. In the curvaton scenario, the magnitude of the cosmic perturbation strongly depends on the evolutions of the inflaton and the spectator field (curvaton) after inflation. The Hubble induced mass of the curvaton after inflation can enhance the field value of the curvaton, and the curvaton field can more easily dominate the energy density of the universe. As a result, constraints on inflation models and curvaton models, e.g.~those on the inflation scale and the reheating temperature, are considerably relaxed. This was pointed out in Refs.~\cite{McDonald:2003xq,McDonald:2004by} in the context of the leptogenesis by the sneutrino curvaton. Here we extend the discussion to generic curvaton models, and comment on an implication to the scenarios of the baryogenesis. In the modulated reheating scenario, as the spectator field moves during the reheating, the dependence of the timing of the reheating on the spectator field value is modified. We find that the magnitudes of the curvature perturbation and the non-gaussianity are affected. We also propose a new model of the modulated reheating which is enabled by the time evolution of the spectator field due to the Hubble induced mass. The model can realize a wide range of the local non-gaussianity parameter. The organization of this paper is as follows. In Sec.~\ref{sec:curvaton}, we discuss the influence of the Hubble induced mass on the curvaton scenario. We also comment on the implications of our findings to the scenarios of baryogenesis. In Sec.~\ref{sec:modulated}, we investigate the modulated reheating scenario and propose a new model. Sec.~\ref{sec:summary} is devoted to summary and discussion. | \label{sec:summary} In this paper, we have discussed the effect of a Hubble induced mass after inflation on the dynamics of spectator field models of the generation of the cosmic perturbation. Although the Hubble induced mass of the spectator field must be suppressed during inflation, it can be sizable after inflation. Actually, in supersymmetric inflation models, it is in general unsuppressed. We have found that the Hubble induced mass can enhance the field value of the curvaton, which helps the curvaton to dominate the energy density of the universe. As a result, the constraints on inflation models and curvaton models are considerably relaxed. We have also discussed the implications of our finding to the scenarios of baryogenesis. The possibility of the baryogenesis is extended. For the modulated reheating scenario, we find that the magnitudes of the curvature perturbation and the non-gaussianity are affected by the Hubble induced mass. We also propose a new model of the modulated reheating utilizing the Hubble induced mass, which can produce an arbitrary value of the local non-gaussianity parameter. | 16 | 7 | 1607.07058 |
1607 | 1607.07328_arXiv.txt | We present the results of the search for decaying dark matter with particle mass in the 6-40~keV range with NuSTAR deep observations of COSMOS and ECDFS empty sky fields. We show that main contribution to the decaying dark matter signal from the Milky Way galaxy comes through the aperture of the NuSTAR detector, rather than through the focusing optics. High sensitivity of the NuSTAR detector, combined with the large aperture and large exposure times of the two observation fields allow us to improve previously existing constraints on the dark matter decay time by up to an order of magnitude in the mass range 10-30~keV. In the particular case of the $\nu$MSM sterile neutrino dark matter, our constraints impose an upper bound $m<20$~keV on the dark matter particle mass. We report detection of four unidentified spectral lines in our data set. These line detections are either due to the systematic effects (uncertainties of calibrations of the NuSTAR detectors) or have an astrophysical origin. We discuss different possibilities for testing the nature of the detected lines. | \label{sec:intro} A range of particle models of the Dark Matter (DM) considers light weight (much below the proton and electron masses) DM particles which are unstable and could decay into the Standard Model particles with production of photons. The most known examples are sterile neutrinos (appearing e.g. in the $\nu$MSM model ~\cite{numsm1,numsm11,numsm,review,review1}) and axion-like particles \cite{axions}. The most clear observational signature of these models is the monoenergetic photon flux at the energy $E=m_{DM}/2$ ($m_{DM}$ is the DM particle mass) expected from all massive DM halos~\cite{pal,barger}. The strongest signal is generically expected to come from the Milky Way galaxy and is detectable from all the directions on the sky \cite{strategy}, with a moderate excess in the direction toward the inner galaxy, compared to the Galactic anti-centre direction. This means that all telescopes sensitive to photons with energies close to the DM particle mass could potentially be used as the DM detectors. If the DM particles are fermions, their mass is constrained to be heavier than $\sim 1$~keV, a constraint imposed by the phase space density of DM in the compact low mass DM halos~\cite{tremainegunn,dodelson,tg1,gorbunov08}. Somewhat tighter bounds arise from the non-observation of small scale structures suppression in Ly$\alpha$ forest data~\cite{lya1,lya2,lya3}. The decay signal from the fermionic DM is detectable with X-ray telescopes. A range of constraints on the lifetime of decaying DM (or on the mixing angle $\theta$ of sterile neutrino DM) have been previously derived from non-observation of the DM decay line by the X-ray telescopes \cite{constr1,constr2,constr3,constr4,constr5,constr6,constr7,spi,strategy,dima_thesis,riemer:14,horiuchi:14,dm_nustar_bullet,gbm,sekiya16}. An unidentified X-ray line at the energy 3.55~keV has been recently reported in the staked spectrum of galaxy clusters and in M31 galaxy~\cite{line35_1,line35_2}. Interpretation of this line as a DM decay line is in tension with the non-observation of the line in nearby dwarf spheroidal galaxies~\cite{dsphs,dsphs2}, galaxy groups~\cite{anderson15} and in the X-ray background~\cite{sekiya16}. The DM decay line signal in an X-ray telescope appears on top of an unrelated astrophysical instrumental background which typically consists of continuum emission and a set of atomic lines. The signal-to-noise ratio (SNR) could be maximised with a suitable choice of the observation target. The strongest Milky Way signal typically occupies the entire field-of-view (FoV) of an X-ray telescope. Larger FoV provides higher signal statistics. The FoVs of existing telescopes operating in the 0.1-10~keV band (XMM-Newton, Chandra) are limited to a fraction of a degree. In this case a further boost of the signal could be achieved by choosing an observation direction which contains, apart from the Milky Way, also a signal from a nearby DM halo, such as e.g. a dwarf spheroidal galaxy or a galaxy cluster \cite{strategy}. To the contrary, the DM decay signal detectable by large FoV telescopes operating in the hard X-ray band above 15~keV, like e.g. ISGRI and SPI telescopes on board of INTEGRAL satellite \cite{spi}, or of the GBM detector on board of Fermi satellite \cite{gbm} is completely dominated by the Milky Way flux. \nus telescope \cite{nustar} provides a large effective collection area (compared to INTEGRAL/SPI), of the order of $10^3$~cm$^2$ in the energy band above 10~keV. This is achieved with the focusing optics, as opposed to the coded mask optics of SPI. The focusing optics also provides an advantage of low background for the observations of point or mildly extended sources (compared to the coded mask optics). The large effective area an low background make \nus competitive as a DM detector which is able to provide a higher sensitivity probe of the DM decay line signal above $10$~keV, in spite of much lower energy resolution, compared to SPI. \nus data have already been used to derive constraints on the DM decay line signal from the direction of Bullet galaxy cluster \cite{dm_nustar_bullet}. The upper bound on the DM decay time stemming from this observation is comparable to the bound previously derived from the INTEGRAL/SPI \cite{spi} and Fermi/GBM data \cite{gbm}. The DM signal considered in \cite{dm_nustar_bullet} was the signal collected through the focusing optics of the telescope. In what follows we show that the \nus detector perceives a much stronger DM signal (compared to that of a galaxy cluster, or of a dwarf spheroidal galaxy) in any astronomical observation. This signal originates from the Milky Way galaxy and is collected not through the focusing optics of the telescope, but rather through the aperture of the X-ray detector unit. We show that account of this signal allows to improve the sensitivity of the DM search with \nus by up to an order of magnitude in the energy range around 10~keV. The Milky Way signal is distributed all over the sky and all the \nus pointings could, in principle, be used for the search of the DM signal. We demonstrate this by analysing two deep, several Msec long, observations of outskirts of our galaxy (COSMOS and ECDFS fields) which were previously used for the analysis of the background of \nus \cite{wik14}. We show that non-detection of the DM decay line in these observations imposes tight constraints on the mixing angle of sterile neutrino DM with masses in the range 10-30~keV. The \nus bound rules out the $\nu$MSM sterile neutrino DM with the mass higher than 20~keV. | Our search of narrow lines in the \nus background spectrum has resulted in detection of a set of lines listed in Table \ref{tab:lines}. Some of these lines are clearly of instrumental origin, such as e.g. the lines at 4.5, 6.4 and 8 keV excited by illumination of the spacecraft by the Sun. The origin of other lines is less clear. The lines in the energy range 14-20~keV are likely to be residuals from improper modelling of the \nus RMF. This RMF has a low-energy "wing" which could result in "ghost" lines at the low-energy wings of the strong instrumental lines. The \nus\ background in the 20-30~keV band is dominated by a blend of strong instrumental lines \cite{nustar}. Our testing of the hypothesis of "ghost" origin of the lines has resulted in suggestive association of the lines above 14~keV with "ghosts" of the stronger instrumental lines at higher energies. The line at 10.4~keV is most probably also of instrumental origin. The effective area of \nus experience a large jump at the L-edge of tungsten. The fit of the background model might favour inclusion of an additional line at the energy of the edge to compensate for the defects of the modelling of the edge shape in the effective area. Two lines: at 3.5 and 4.7~keV have less obvious origin. Although the 3.5~keV is mentioned as being of solar origin in the Ref. \cite{wik14}, we do not find the expected variation of the line flux between the ``sun-illuminated'' and ``no-sun'' parts of the data set. The most probable origin of this line is the imperfection of the modelling of the energy dependence of the \nus effective area close to the low energy threshold of the instrument (nominally at 3~keV). At the same time, we could not exclude the possibility of the dark matter origin of the line. We also have not found a plausible explanation for the 4.7~keV line.\footnote{Note however, that $Cs L\beta$ line can be produced at similar energy and stronger Cs $K\alpha$, $K\beta$ lines are present in instrumental model at $\gtrsim 30$~keV band which may indicate the possible instrumental origin of 4.7~keV line.} The alternative hypotheses of instrumental and astrophysical (dark matter decay) origin of the detected lines could be tested using the strategy described in the Ref. \cite{spi}. The intensity of the line originating from the dark matter halo of the Milky Way is expected to vary across the sky exhibiting an excess toward the inner Galaxy and deficit of the flux from the outer Galaxy. To the contrary, the instrumental line is not expected to show clear large scale variability pattern on the sky (spurious variations could still be induced by the changes of observational conditions during exposures of different parts of the sky). A test of the dark matter origin of the detected lines could therefore be done with a set of additional "deep field" exposures at different off-Galactic-centre directions. An alternative approach for the testing of the dark matter origin of the lines is to obtain deep exposures of selected dwarf spheroidal galaxies (e.g. Draco). This should provide a factor-of-two boost in the expected dark matter line signal for the real dark matter decay line (as derived in Ref. \cite{we_athena}) and no variation of the signal for the instrumental line. The improvement of sensitivity of the search of the dark matter decay line provided by the \nus exposures of COSMOS and ECDFS fields is important in the context of testing the reference $\nu$MSM model of sterile neutrino dark matter. Fig. \ref{fig:constraints} shows that the \nus data limit the mass of the dark matter neutrino to be below 20~keV, otherwise production of sufficient amount of the dark matter in the Early Universe would require too high lepton asymmetry. Exclusion of the 20-30~keV mass range closes a sensitivity gap of future X-ray telescope Athena, which will have an X-ray spectrometer XIFU able to detect the dark matter decay line in the energy range below 10~keV (neutrino mass range below 20~keV) \cite{we_athena}. \label{sec:conclusions} | 16 | 7 | 1607.07328 |
1607 | 1607.02507_arXiv.txt | We show how the interplay between active galactic nuclei (AGN) and merger history determines whether a galaxy quenches star formation at high redshift. We first simulate, in a full cosmological context, a galaxy of total dynamical mass $\Mvir = 10^{12}\,\Msol$ at $z=2$. Then we systematically alter the accretion history of the galaxy by minimally changing the linear overdensity in the initial conditions. This ``genetic modification'' approach allows the generation of three sets of $\Lambda$CDM initial conditions leading to maximum merger ratios of 1:10, 1:5 and 2:3 respectively. The changes leave the final halo mass, large scale structure and local environment unchanged, providing a controlled numerical experiment. Interaction between the AGN physics and mergers in the three cases lead respectively to a star-forming, temporarily-quenched and permanently-quenched galaxy. However the differences do not primarily lie in the black hole accretion rates, but in the kinetic effects of the merger: the galaxy is resilient against AGN feedback unless its gaseous disk is first disrupted. Typical accretion rates are comparable in the three cases, falling below $0.1\,\Msol\,\yr^{-1}$, equivalent to around $2\%$ of the Eddington rate or $10^{-3}$ times the pre-quenching star formation rate, in agreement with observations. This low level of black hole accretion can be sustained even when there is insufficient dense cold gas for star formation. Conversely, supernova feedback is too distributed to generate outflows in high-mass systems, and cannot maintain quenching over periods longer than the halo gas cooling time.\vspace{0.5cm} | Reproducing the population of galaxies observed in the Universe within a $\Lambda$CDM cosmological paradigm requires significant energetic feedback \citep{WhiteFrenk91Review} to prevent over-cooling and excessive star formation. In fact star formation must be inefficient in both low and high mass dark matter halos, with a peak efficiency corresponding approximately to the luminosity function turnover $L_\star$ \citep[][]{Bower06AGNQuenching,Guo10AbundanceMatching,Moster13AbundanceMatch,Behroozi13AbundanceMatch}. Regulation of galaxies with luminosities $L<L_{\star}$ can be attributed to ultraviolet background radiation slowing the accretion of gas into shallow potential wells \citep[][]{Efstathiou92UV,Bullock00UV,Somerville02UV}, and subsequently energy input from young stars \cite[][]{WhiteRees78,Cole94Semianalytic,EfstathiouSNfeedback,Governato07,Pontzen08DLAs}. The relatively shallow potential wells mean that the typical speeds achieved by supernova- and radiation-driven galactic winds can exceed the escape velocity \citep{MacLow99outflows,Christensen15outflows}. However many lines of reasoning suggest that these processes must become increasingly ineffective as dynamical masses approach $10^{12}\,\Msol$ \citep[][]{Benson03,Bower06AGNQuenching,Hopkins14FIRE,KellerWadsley16}. Accounting for reduced star formation at high masses has proved more contentious. Early scaling arguments and semi-analytic models suggested that long cooling times associated with the raised virial shock temperature could be a simple explanation \citep{Binney77,ReesOstriker77,WhiteRees78,Kauffmann93semianalytic,Cole94Semianalytic,Somerville99semianalytic}. Unfortunately the resulting suppression is only effective for low baryon fractions, making it insufficient to account for observed luminosity functions given today's concordance $\Lambda$CDM parameters \citep{Benson03}. A related possibility is offered by starvation of galaxies simply because accretion slows down at late times \citep{Feldman2015AccretionQuenching,Feldman2016AccretionQuenching}. However the host halo masses of red galaxies are observed to be systematically lower than those of star-forming galaxies, suggesting that accretion continues after star formation shuts down \citep{Mandelbaum16LensingMasses}. Semi-analytic models also suggest that without additional mechanisms, it is very difficult to quantitatively explain the division of the population into actively star-forming disks and ``quenched'', red ellipticals \citep{Baldry04Bimodality,Bower06AGNQuenching}. This division appears to be in place even at high redshift \citep{Ilbert10COSMOSbimodality,Brammer11bimodality}. Environmental effects, especially in high-density group and cluster regions can strip an infalling galaxy of its gas reservoir which leads to a more natural bimodality \citep[e.g.][]{QuilisMoore2000EnvQuenching,Gomez03SFRvsEnv,Boselli06,vdBosch08SatelliteQuenching,McCarthy08SatelliteQuenching,Bahe15GroupQuenching}. However, quenching also occurs outside of these environments \citep{vdBosch08SatelliteQuenching,Guo09GroupGalaxyProperties,Peng10Quenching,Wijesinghe12gamaEnvironmentSFR}. We are therefore led to re-examine the role of feedback. One route to increasing the efficacy of stellar feedback at intermediate and high masses is to induce intense, compact starbursts through major mergers or violent disk instabilities \citep{DekelBurkert14,Ceverino15BulgeFormation,Zolotov15Quenching,Tacchella16Compaction}. However quenching in this scenario remains fairly gradual, with star formation declining over several gigayears \citep{Zolotov15Quenching}, whereas a number of lines of evidence suggest that a significant fraction of early type galaxies are formed by much more rapid quenching \citep[e.g.][]{Thomas05,Schawinski14GreenValley,Belli15,Barro16RapidQuenching}. Moreover, in this picture, quenching can be maintained over long time periods only in high-mass galaxies where the virial shock prevents rapid cooling of material accreted after the starburst event \citep{WhiteFrenk91Review,BirboimDekel03}. By definition, quenching star formation for longer than the cooling time of halo gas requires an energy source other than stellar feedback. Many studies suggest that the crucial input comes from active galactic nuclei (AGN) powered by central supermassive black holes; see, for example, \cite{tdMatteo2005Quasars,Hopkins05BHgalaxyMergers,Bower06AGNQuenching,Croton06,Sijacki07AGNfeedback,diMatteo2008QuasarsCosmo,Cattaneo09BHreview,Johansson09,Fabian12BHreview,Dubois13AGNGalFormation,Dubois16} and references therein. AGN drive rapid outflows which can be directly observed in post-starburst galaxies \citep[e.g.][]{Tremonti07OutflowsPostStarburst,RupkeVeilleux11}, suggesting that black holes are able to suppress star formation by removing the supply of gas. Adding weight to the connection between a galaxy and its central black hole (BH) is the strong observed correlation between BH mass and stellar mass in the bulge \citep{Cattaneo09BHreview,Volonteri12BHreview,Kormendy13BH_annrev}, which can be interpreted either as evidence that the BH and SF are fed from the same supply of cold gas; that feedback from BH regulates the star formation rate; or even, for low-mass galaxies, that feedback from SF regulates the BH accretion \citep{Dubois15BHSNinteraction}. It is possible that the true explanation is a combination of all three effects. But whatever the regulatory role that BHs play, observations also show that AGN activity is common in highly star-forming galaxies \citep[e.g.][]{Nandra07AGNhostgalaxies,Simmons12AGN,Mullaney2012,Rosario12SFRofAGN,Rosario13AGNandSFcorrelate,ForsterSchreiber14SinsOutflows,Mullaney15SFRofAGN,Carniani16QuasarsDontAlwaysQuench}, suggesting that the precise role of the BH is strongly dependent on other, unidentified factors. In this paper, we will identify those factors by studying the interaction between AGN feedback and mergers in a realistic cosmological environment, discussing how it can lead to quenching that is maintained for periods longer than the halo cooling time. To address this question we require an approach that offers control over environment, accretion history, and feedback models. The basis for our study is a $z=2$ galaxy with a dynamical mass of $10^{12}\,\Msol$, simulated using the \textsc{ChaNGa} code \citep{Changa15}. Uniquely, we are able to modify the accretion history of the system by making minimal modifications to the large-scale structure. This ability to ``genetically modify'' the system that we study, while keeping the cosmological conditions consistent with the $\Lambda$CDM inflationary scenario, is introduced by \cite{Roth15GM}. Our main aim here is to study how the BHs respond to changing the significance of the most major merger. The combination of being able to control feedback and history with a fixed local environment (including the precise directions of the filaments that feed cool gas to the galaxy) allows us to isolate and identify the conditions required for quenching. In particular, all simulations are run first with and then without BHs to quantify the effect of the AGN feedback. This paper is structured as follows. In Sec. \ref{sec:simulations} we describe the simulations in more detail. The results are discussed in Sec. \ref{sec:results}. Finally we conclude in Sec. \ref{sec:concl-disc}. | \label{sec:concl-disc} In this paper, we investigated how galaxy mergers and black hole (BH) feedback work together to quench star formation. We used ``genetically modified'' initial conditions to set up three galaxies which live in the same environment and large scale structure, but that differ in merger history. Coupled to a new implementation of BH physics \citep{Tremmel16Romulus}, we were able to isolate how mergers lead to quenching. Our results show that quenching in field galaxies with virial masses larger than a few times $10^{11}\,\Msol$ can only be achieved using BH feedback (Fig. \ref{fig:ssfr}). Supernova feedback becomes ineffective at driving galactic winds at these masses \citep{Christensen15outflows,KellerWadsley16}. However the addition of BH feedback does not increase the total energy budget. Instead, its strong central concentration generates low-density outflows at high temperatures and velocities, minimising radiative energy losses. As the wind moves outwards, it sweeps up halo material, increasing the outflow density only after the deepest part of the gravitational potential has already been overcome. In this way the total energy loss due to purely gravitational effects is minimised. Mergers trigger quenching without invoking a near-Eddington quasar phase for the central AGN. The average BH accretion rate agrees with constraints from stacking \citep{Mullaney2012}. All of our galaxies, regardless of the merger history, have an active BH at their centre, in keeping with observed AGN in both quenched and unquenched systems \citep{Simmons12AGN, Mullaney2012, Rosario15, Mancini15}, driving a large-scale galactic wind \citep{Genzel14, Harrison14,ForsterSchreiber14SinsOutflows}. The cold disk of gas in the unquenched galaxies slows BH accretion due to angular momentum support. It also confines the effect of the BHs and directs it outwards in a funnel, allowing star formation to proceed despite the rapid central outflows \citep{CanoDiaz12QuasarQuenchingPlusSFElsewhere,Carniani16QuasarsDontAlwaysQuench}. Mergers start the quenching process through mechanical disruption of the cold disk. Subsequently the AGN feedback is able to have far greater impact in disrupting the existing star-forming gas and cool gas inflows. Acting together in this way, the merger and central BH push the galaxy into a self-sustaining, long-term quiescent state. The drop-off from the main sequence is rapid, occurring over around $250\,\Myr$ in both our reference and enhanced-merger simulations. Many clues suggest that quenching is indeed rapid in real galaxies \citep{Barro13,Barro15,Mancini15}. In the simulations, quenching is followed by a slow decline in BH activity 0.5 -- 1 Gyr after the galaxy leaves the main sequence, in agreement with observational evidence that AGN in star forming galaxies contribute the majority of the X-ray luminosity density \citep{Mancini15, Rodighiero15}. The quenching mode is then maintained by the turbulent remnants of the disk sweeping a low level of infalling cool gas into the BH, which allows BH accretion to continue as seen in recent observations by \cite{Tremblay16AGNColdAccretion}. One prediction of this scenario is that quiescent galaxies will have an offset BH-stellar mass relation compared to star-forming galaxies \citep{Terrazas16SemiAnalyticAGNQuenching}. The BH feedback has a significant dynamical impact on the galaxy. Gas that is accreted into the disk at early times tends to carry little angular momentum; in the absence of a strong outflow or cycling mechanism, this piles up, creating a strong central bar (Fig. \ref{fig:portraits}, third row of panels). The maximum circular speed in the disk plane reaches $\vmax \simeq 500\,\kms$ in all three merger scenarios when only stellar feedback is available (Fig. \ref{fig:vmax}, top panel). The stronger outflows of the BH+SNe simulations prevent this pile-up from occurring, using a combination of a galactic fountain and outright ejection to set a limit of $\vmax \simeq 250 \kms$. {There is considerable interest in integral field spectroscopy of high-redshift galaxies \citep[e.g.][]{ForsterSchreiber06SINS,Newman15,Wisnioski15Kmos3d},} and we anticipate that the kinematics of galaxies can therefore be used as an indicator for the balance between BH and stellar feedback. Our study hints at the broader potential of the ``genetic modification'' (GM) technique \citep{Roth15GM} to shed light on problems in galaxy formation. Because our three scenarios (reference, enhanced and suppressed) share the same environment and large scale structure (Fig. \ref{fig:portraits}), with the central galactic disk assembling from gas streaming along the same filaments, we have been able to construct a clean test of the effect of merger ratio. The GM approach can search systematically over histories, {providing controlled tests which retain the benefits of idealised simulations in pinning down the effect of merger ratio \cite[e.g.][]{Johansson09} while also including cosmological accretion}. We illustrated this point by constructing a scenario with a merger ratio intermediate between the reference and enhanced cases (Fig. \ref{fig:semi-enhanced-ssfr}). The results allowed us to estimate that a quarter of systems in our mass range at $z\simeq 2$ should appear quenched, and a further quarter will have recovered from an earlier episode of quenching -- assuming that merger ratio is the significant variable controlling the outcome. Future work could extend this study by changing, for example, the mean density of the region within which the galaxy is hosted to directly interrogate the important role of environment \citep[e.g.][]{Peng10Quenching,Wijesinghe12gamaEnvironmentSFR}. | 16 | 7 | 1607.02507 |
1607 | 1607.05111_arXiv.txt | We evaluate the effects of a distant planet, commonly known as planet 9, on the dynamics of the giant planets of the Solar System. We find that, given the large distance of planet 9, the dynamics of the inner giant planets can be decomposed into a classic Lagrange-Laplace dynamics relative to their own invariant plane (the plane orthogonal to their total angular momentum vector) and a slow precession of said plane relative to the total angular momentum vector of the Solar System, including planet 9. Under some specific configurations for planet 9, this precession can explain the current tilt of $\sim 6^{\circ}$ between the invariant plane of the giant planets and the solar equator. An analytical model is developed to map the evolution of the inclination of the inner giant planets' invariant plane as a function of the planet 9's mass, inclination, eccentricity and semimajor axis, and some numerical simulations of the equations of motion of the giant planets and planet 9 are performed to validate our analytical approach. The longitude of the ascending node of planet 9 is found to be linked to the longitude of the ascending node of the giant planets' invariant plane, which also constrain the longitude of the node of planet 9 on the ecliptic. Some of the planet 9 configurations that allow explaining the current solar tilt are compatible with those proposed to explain the orbital confinement of the most distant Kuiper belt objects. Thus, this work on the one hand gives an elegant explanation for the current tilt between the invariant plane of the inner giant planets and the solar equator and, on the other hand, adds new constraints to the orbital elements of planet 9. | The gradual discovery of increasingly distant trans-Neptunian objects (TNOs) has allowed new tests for the existence of a yet undiscovered distant planet in the solar system. \cite{gomesetal2015} analyzed the large semimajor axis centaurs and concluded that they are produced continually by the decrease of perihelia of scattered disk objects, induced by the perturbation of a distant planet. \cite{trujillo2014}, as they announced the discovery of the distant TNO 2012 VP113, also noted that distant TNOs not perturbed by close encounters with Neptune show a remarkable alignment of their arguments of perihelia and proposed that a distant planet is responsible for this alignment. More recently, \cite{BaBr2016} studied more deeply the orbital alignment of those distant TNOs, showing that the six most distant objects exhibit also a clustering in their longitudes of node; they estimated that the probability that this double alignment in argument of perihelion and longitude of the node is just fortuitous is $0.007$\%. Moreover they showed that a planet 9 (hereafter named just pl9) could account for said alignment if it had a mass of about $10 M_{\oplus}$ and an orbit with semimajor axis between $300$ and $900$ au, perihelion distance between $200$ and $350$ au, and orbital inclination of about $30^{\circ}$ to the ecliptic plane. The Batygin-Brown approach based on secular dynamics is able to determine an approximate orbit for the distant planet that could explain the said alignment, but not the planet's position on that orbit. \cite{fienga2016} use a typical orbit among those proposed by \cite{BaBr2016} and determined the range in true longitude of pl9 on that orbit that decreases the residuals in INPOP ephemerids of Saturn, relative to the Cassini data. \cite{hol-payne2016} obtained a similar result using JPL ephemerids. \cite{BrBa2016} refined their previous results by further constraining the mass and orbital elements of pl9 that are compatible with the observed TNOs orbital alignment. They now argue for pl9's semimajor axis in the range $380–980$au, perihelion distance in the range $150–350$ au and a mass between $5$ and $20 M_{\oplus}$, for an orbital inclination of $30^{\circ}$. \cite{malhotra2016} looked for extra constraints on pl9 orbit by analyzing the orbital periods of the four longest period TNOs. Their approach is based on the supposition that pl9 is in mean motion resonances with those TNOs. \cite{beust2016}, however, showed that a mean motion resonant configuration is not necessary to explain the orbital confinement. Here we study the precession of the plane orthogonal to the total angular momentum of the four giant planets due to the perturbation of pl9. We find that, given the large distance of pl9, the dynamics of the giant planets can be decomposed into a classic Lagrange-Laplace dynamics relative to their own invariant plane (the plane orthogonal to their total angular momentum vector, hereafter named iv4) and a slow precession of said plane relative to the total angular momentum vector of the Solar System, including pl9. Planetary system formation predicts that planets are formed from a disk of gas and dust and this disk rotates on the same plane of the star's equator. The final planetary orbits, if no mutual close encounters take place, must be approximately coplanar and coincident with the star's equator. We thus suppose that the giant planets and the solar equator were initially on the same plane. We assume that pl9 was scattered away from the region of the other giant planets when the disk was still present and the solar system was still embedded in a stellar cluster \citep{izidoroetal2015}. The stellar cluster is needed, so that the perihelion of the orbit of pl9 can be lifted and pl9 can decouple from the other planets \citep{brasseretal2008}. Because most of the angular momentum is in the protoplanetary disk, it is likely that the ejection of pl9 onto an inclined orbit did not significantly change the inclination of the disk and of the other giant planets. Notice also that the inclination of pl9 might have been increased by the action of the cluster, while lifting the perihelion in a Lidov-Kozai like dynamics \citep{brasseretal2008}. Thus, we assume that, at the removal of the protoplanetary disk and of the birth cluster of the Sun, the 4 major giant planets were on orbits near the solar equator, while pl9's orbit was off-plane. At this point, a slow precession of iv4 started to take place relative to the total angular momentum vector of the Solar System, including pl9, keeping however the orientation of the solar equator plane unchanged. Thus the current angle between the solar equator and iv4 (about $6^{\circ}$ - see below) must be a signature of pl9 perturbation and we aim at finding ranges of orbital elements and mass for pl9 that can explain quantitatively the present tilt of iv4 relative to the solar equator. The solar equator with respect to the ecliptic is identified by an inclination $I_S=7.2^{\circ}$ and a longitude of the ascending node $\Omega_S=75.8^{\circ}$ \citep {beck&giles}. The invariant plane with respect to the ecliptic is defined by an inclination $I_i=1.58^{\circ}$ and a longitude of the ascending node $\Omega_i=107.58^{\circ}$ \citep {souami2012}. Employing two rotations we can find the invariant plane angles with respect to the solar equator to be $I_v=5.9^{\circ}$ and $\Omega_v=171.9^{\circ}$. We will use these parameters throughout the rest of the paper. We also consider that iv4, as above defined, is equivalent to the invariant plane of the solar system mentioned in the works above, which take into account also the inner planets. In Section 2 we develop two analytical approaches aimed at determining the tilt experienced by iv4 due to the perturbation of pl9. We also perform some numerical integrations of the full equations of motion to validate our analytical approaches. In Section 3, we apply our analytical method to determine the range of masses and orbital elements of pl9 that can account for the observed tilt of iv4 to the solar equator plane. In Section 4, we draw our conclusions. | Some ideas had been put forward to possibly explain the inclination of the invariant plane of the known planets relative to the solar equatorial plane \citep{batyginetal2011} but in view of the convincing case presented in \cite{BaBr2016} for the existence of pl9, it is quite natural to suppose that such a tilt was caused by a slow precession of iv4 around the total angular momentum vector of the solar system (including pl9). In this paper we constrained possible masses and orbital elements for pl9 that can account for the present tilt of iv4 with the solar equator. Our results are usually compatible with those of \cite {BaBr2016} and \cite{BrBa2016} though with somewhat larger eccentricities. We also determine a range of possible longitudes of the ascending node for pl9 which often overlaps with the range given in \cite{BaBr2016} except for smaller masses and inclinations of pl9. For instance, for $I_9=30^{\circ}$ we need a mass larger than $\sim 4 \times 10^{-5} M_{\odot} \sim 13 M_{\oplus}$ to match the range of the longitudes of the ascending node for pl9 proposed by \cite{BaBr2016}. | 16 | 7 | 1607.05111 |
1607 | 1607.07863_arXiv.txt | Clusters of galaxies have the potential of providing powerful constraints on possible deviations from General Relativity. We use the catalogue of Sunyaev-Zel'dovich sources detected by \emph{Planck} and consider a correction to the halo mass function for a $f(R)$ class of modified gravity models, which has been recently found to reproduce well results from $N$-body simulations, to place constraints on the scalaron field amplitude at the present time, $f_{R}^0$. We find that applying this correction to different calibrations of the halo mass function produces upper bounds on $f_{R}^0$ tighter by more than an order of magnitude, ranging from $\log_{10}(-f_{R}^0) < -5.81$ to $\log_{10}(-f_{R}^0) < -4.40$ ($95\%$ confidence level). This sensitivity is due to the different shape of the halo mass function, which is degenerate with the parameters used to calibrate the scaling relations between SZ observables and cluster masses. Any claim of constraints more stringent that the weaker limit above, based on cluster number counts, appear to be premature and must be supported by a careful calibration of the halo mass function and by a robust calibration of the mass scaling relations. | 16 | 7 | 1607.07863 |
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1607 | 1607.02394_arXiv.txt | We present an updated constrained hyperbolic/parabolic divergence cleaning algorithm for smoothed particle magnetohydrodynamics (SPMHD) that remains conservative with wave cleaning speeds which vary in space and time. This is accomplished by evolving the quantity $\psi / c_{\rm h}$ instead of $\psi$. Doing so allows each particle to carry an individual wave cleaning speed, $c_{\rm h}$, that can evolve in time without needing an explicit prescription for how it should evolve, preventing circumstances which we demonstrate could lead to runaway energy growth related to variable wave cleaning speeds. This modification requires only a minor adjustment to the cleaning equations and is trivial to adopt in existing codes. Finally, we demonstrate that our constrained hyperbolic/parabolic divergence cleaning algorithm, run for a large number of iterations, can reduce the divergence of the field to an arbitrarily small value, achieving $\nabla \cdot {\bf B}=0$ to machine precision. | Accurately evolving the equations of magnetohydrodynamics (MHD) in numerical simulations is crucial in astrophysical fluid dynamics. In smoothed particle magnetohydrodynamics (SPMHD) \citep{gm77, pm85, pm04a, pm04b, pm05, price12}, upholding the divergence-free constraint of the magnetic field has been the main technical difficulty. The usual approach is to evolve the magnetic field directly by the induction equation (as in \citealt{pm85}), but this preserves a divergence-free magnetic field only to truncation error. These errors cause more harm than just yielding an unphysical field. They introduce spurious monopole accelerations, which have to be carefully handled in SPMHD in order to ensure numerical stability, at the price of no longer exactly conserving momentum \citep{pm85, morris96, bot01}. Handling the divergence-free constraint on the magnetic field is therefore one of the most important aspects of accurate SPMHD simulations. One option is to define the magnetic field in a way that manifestly enforces the divergence-free constraint. Use of the Euler potentials, ${\bf B} = \nabla \alpha \times \nabla \beta$ where $\alpha$ and $\beta$ are passive scalars, was proposed as early as \citet{pm85}, and recently the potentials have been used in simulations of protostar formation \citep{pb07}, star cluster formation \citep{pb08, pb09} and magnetised galaxies \citep{dp08, kotarbaetal09}. However, the Euler potentials cannot represent winding motions, prevent dynamo processes by construction \citep{brandenburg10}, and it is not clear how to incorporate non-ideal dissipation. A vector potential implementation, ${\bf B} = \nabla \times {\bf A}$, was tested for SPMHD by \citet{price10}, but was found to be numerically unstable. \citet{se15} recently proposed that the vector potential could be used, if one added numerical diffusion to the potential, enforced the Coulomb gauge condition on the vector potential ($\nabla \cdot {\bf A} = 0$) and smoothed the resulting magnetic field, though it is not clear how robust this approach is in practice. The second option to handle the divergence-free constraint in SPMHD is to directly evolve the magnetic field with the induction equation, but then `clean' errors out of the field. For example, parabolic diffusion terms can be used to smooth the magnetic field at the resolution scale \citep{morris96}. The artificial resistivity formulation of \citet{pm04a, pm05} has been used for this purpose (e.g., \citealt{burzleetal11b}), however, artificial resistivity is intended for shock capturing and dissipates physical as well as unphysical components of the field. A similar idea is to periodically smooth the magnetic field to remove fluctuations below the resolution limit \citep{bot01}, but this adds computational expense, is time resolution dependent, and reduces the spatial resolution of the magnetic field. At present, the best option for divergence cleaning in SPMHD is the `constrained' hyperbolic/parabolic divergence cleaning method of \citet{tp12}, an improved version of the method by \citet{dedneretal02}. The original idea from \citet{dedneretal02} was to couple an additional scalar field, $\psi$, to the induction equation according to \begin{align} \frac{\partial {\bf B}}{\partial t} & = \nabla \times ({\bf v} \times {\bf B}) - \nabla \psi, \label{eq:dbdt-gradpsi-dedner} \\ \frac{\partial \psi}{\partial t} & = - c_{\rm h}^2 (\nabla \cdot {\bf B}) - \frac{\psi}{\tau}, \label{eq:psievolution-dedner} \end{align} where ${\bf B}$ is the magnetic field and ${\bf v}$ is the velocity. These may be combined to produce a damped wave equation for the divergence of the magnetic field, \begin{equation} \frac{{\partial^2} (\nabla \cdot {\bf B})}{\partial t^2} - c_{\rm h}^2 \nabla^2 (\nabla \cdot {\bf B}) + \frac{1}{\tau} \frac{\partial (\nabla \cdot {\bf B})}{\partial t} = 0. \label{eq:divbwave} \end{equation} From Equation~(\ref{eq:divbwave}), we see that Equation~(\ref{eq:dbdt-gradpsi-dedner}) and the first term on the right hand side of Equation~(\ref{eq:psievolution-dedner}) represent hyperbolic transport of divergence errors at a characteristic speed, $c_{\rm h}$, which we refer to as the `wave cleaning speed'. This is typically chosen to be the fast MHD wave speed so that it obeys the local Courant condition and does not impose any additional timestep constraint. The second term on the right hand side of Equation~(\ref{eq:psievolution-dedner}) produces parabolic diffusion on a timescale defined according to \begin{equation} \tau \equiv \frac{h}{\sigma c_{\rm h}}, \end{equation} where $h$ is the smoothing length (resolution scale) and $\sigma$ is a dimensionless constant with empirically determined optimal values of $0.3$ and $1.0$ in 2D and 3D, respectively \citep{tp12}. The combination of hyperbolic and parabolic terms in Equations~(\ref{eq:dbdt-gradpsi-dedner})--(\ref{eq:psievolution-dedner}) spreads the divergence of the magnetic field over a larger area, reducing the impact of any single large source of error, while also allowing the diffusion to be more effective. In \citet{tp12}, we showed that the original \citet{dedneretal02} approach could be unstable at density jumps and free surfaces, leading to exponential growth of magnetic energy. To remedy this, we derived a version of the cleaning equations under the constraint that the hyperbolic transport should conserve energy. Though $\psi$ is not a physical variable, conservation of energy for the hyperbolic term between the magnetic and $\psi$ fields ensures that, when the parabolic term is included, magnetic energy can only ever be removed by divergence cleaning, never added, guaranteeing numerical stability. The `constrained' or `conservative' cleaning equations we derived in \citet{tp12} are given by \begin{align} \frac{{\rm d}{\bf B}}{{\rm d}t} & = ({\bf B}\cdot\nabla) {\bf v} - {\bf B}(\nabla\cdot{\bf v}) - \nabla \psi, \label{eq:dbdt-gradpsi-original} \\ \frac{{\rm d}\psi}{{\rm d}t} & = - c_{\rm h}^2 (\nabla \cdot {\bf B}) - \frac{\psi}{\tau} - \frac{1}{2} \psi (\nabla \cdot {\bf v}) , \label{eq:psievolution-original} \end{align} where ${\rm d}/{\rm d}t \equiv \partial/\partial t + {\bf v}\cdot\nabla$ is the Lagrangian time derivative. The formulation of the induction equation (Equation~(\ref{eq:dbdt-gradpsi-original})) in the absence of the $\nabla\psi$ term follows the `divergence preserving scheme' of \citet{powelletal99} (see also \citealt{janhunen00,dellar01}), meaning that divergence errors are preserved by the flow in the absence of cleaning. The third term in Equation~(\ref{eq:psievolution-original}) was introduced by \citet{tp12} to account for changes in $\psi$ from compression or rarefaction of the gas, and is necessary to ensure total energy conservation in the absence of damping. The practical advantage of this algorithm for SPMHD is that it adds no additional timestep constraint, is simple to implement, computationally efficient, and has been successfully used to enforce the divergence-free constraint in simulations of jets and outflows during protostar formation \citep{ptb12, btp14, lbp15, wpb16}. However, our original method was derived assuming that the cleaning speed, $c_{\rm h}$, is constant in both space and time, but this is not true in practice and presents a source of non-conservation of energy. Furthermore, source terms are added to the right hand side of Equation~(\ref{eq:divbwave}) when $c_{\rm h}$ or $\tau$ are time or spatially variable, by the addition of the $\tfrac{1}{2} \psi (\nabla \cdot {\bf v})$ term, and by solving the cleaning equations in the Lagrangian frame of motion. How these source terms change the propagation of divergence errors is not properly understood, but will be addressed in this work. In this paper, we derive an improvement to constrained hyperbolic/parabolic divergence cleaning such that the hyperbolic evolution equations remain conservative even in the presence of a variable cleaning speed (Section~\ref{sec:cleaningv2}). We demonstrate that these equations create a generalised wave equation which naturally incorporates the source terms (Section~\ref{sec:sourceterms}). Aspects of the method are tested in Section~\ref{sec:idealised-tests} using a series of test problems. In particular, we will show that, if the time variability of the cleaning wave speed is not properly accounted for, the non-conservation of energy introduced may reduce the effectiveness of the divergence cleaning, and, worst case scenario, lead to runaway energy growth and numerical instability. In Section~\ref{sec:practical-tests}, the original and updated versions of the method are compared using standard MHD tests to quantify how much of an improvement the new scheme confers. Finally, in Section~\ref{sec:divbzero}, we demonstrate that, by iterating the divergence cleaning equations, it is possible to clean the magnetic field until $\nabla \cdot {\bf B}=0$ to machine precision in the chosen divergence operator. We summarise in Section~\ref{sec:summary}. While our focus in this paper is on improved divergence cleaning methods for SPMHD, our analysis and in particular our reformulation of the cleaning equations should apply equally to implementations of hyperbolic/parabolic cleaning in grid-based MHD codes, particularly in the context of adaptive mesh refinement (AMR) where jumps in the cleaning speed may occur at refinement boundaries. Application to Eulerian MHD codes is beyond the scope of this paper but would be an interesting and worthwhile extension to our work. | \label{sec:summary} We have developed a new formulation of hyperbolic/parabolic divergence cleaning for SPMHD which takes account of the variability in the wave cleaning speed. This is accomplished by evolving $\psi / c_{\rm h}$ instead of $\psi$ as the primary variable. In Section~\ref{sec:cleaningv2}, cleaning equations were derived in terms of this quantity. Using this set of equations ensures that divergence cleaning cannot lead to increases in magnetic energy, as the parabolic damping can only remove magnetic energy and the hyperbolic terms are guaranteed to exactly conserve $e_\psi$ and magnetic energy. The new cleaning equations remain similar to the previous equations, differing only by factors of $c_{\rm h}$, but permit the wave cleaning speed to evolve in time without needing an explicit expression for the time derivative of $c_{\rm h}$. In Section~\ref{sec:sourceterms}, the generalised wave equation was derived demonstrating that the propagation of divergence errors remains hyperbolic/parabolic, but occurs in the co-moving frame and naturally accounts for changes in the density, wave speed and parabolic damping term. The new method was tested using a series of idealised tests (Section~\ref{sec:idealised-tests}) and standard MHD test problems (Section~\ref{sec:practical-tests}). The issue related to variable wave cleaning speeds was demonstrated in Section~\ref{sec:tests-timevary} using a simplified test of the advection of a divergence blob. When the wave cleaning speed was varied in time, it led to exponential increases of magnetic energy in the form of increased divergence error. This occurred both for purely hyperbolic cleaning ($\sigma = 0$) and mixed hyperbolic/parabolic cleaning. No such errors were found when the test was repeated for wave cleaning speeds that were constant in time but which had spatial variations (Section~\ref{sec:tests-xvary}), nor for time or spatial discontinuities in the parabolic damping parameter (Section~\ref{sec:tests-disctau}). In Section~\ref{sec:tests-advection}, the effect of advecting $\psi/c_{\rm h}$ was tested. The motivation for this test was that the original \citet{dedneretal02} formulation used Eulerian derivatives, i.e.\ no advection of $\psi$, however, the constraint of energy conservation requires the use of Lagrangian derivatives, adding advection of $\psi/c_{\rm h}$ to our scheme. Using the divergence advection test, we found that if the cleaning equations are implemented using Eulerian derivatives, the average divergence error increased by $30\%$ when the background velocity of the fluid increased from $\mathcal{M}=0.45$ to $\mathcal{M}=10$. By contrast, our Lagrangian implementation produced equivalent results for all flow velocities. Furthermore, computing Eulerian derivatives require `reverse advection' terms be added to counteract the Lagrangian nature of SPMHD, adding a velocity dependence into the Courant timestep constraint. For these reasons, we conclude that the cleaning equations should be implemented with Lagrangian derivatives. Our final idealised test was to confirm that the $\tfrac{1}{2} (\psi / c_{\rm h}) (\nabla \cdot {\bf v})$ term added to account for compression and rarefaction is indeed necessary to exactly conserve energy (Section~\ref{sec:tests-divv}). To investigate this, supersonic compressional motions were added to the divergence advection test. As the errors due to time-stepping were reduced through reductions of the Courant factor, the total energy of the simulations with the compression term converged to a constant value in time, whereas the simulations without the term did not. Thus, the compression term resolves a source of non-conservation of energy, and we conclude that this term is strictly required to exactly conserve energy, though we note that the errors introduced by its absence are smaller than those from the time-stepping algorithm in general simulations. In Section~\ref{sec:practical-tests}, the new cleaning method was applied to simulations of a blast wave in a magnetised medium (Section~\ref{sec:tests-mhdblast}), the Orszag-Tang vortex (Section~\ref{sec:tests-orszag}), and the MHD rotor problem (Section~\ref{sec:tests-rotor}). In general, using the new cleaning method provided reductions of average divergence error of $1$--$2\%$, to a maximum of $5\%$ occurring when the wave cleaning speed underwent its most rapid changes. For the blast wave test, using the new cleaning equations led to less overall dissipation of magnetic energy, dissipating magnetic energy at a rate $5\%$ less than the original method. We note that the dissipation of magnetic energy from divergence cleaning is $\lesssim 10$\% of that from artificial resistivity, meaning that it is only a minor contribution to the total dissipation. Finally, in Section~\ref{sec:divbzero}, we demonstrated that it is possible to clean a magnetic field to arbitrarily small values of $\nabla \cdot {\bf B}$ in SPMHD, albeit with a large number of iterations. We found that using a lower value for the damping parameter ($\sigma=0.02$--$0.03$ in 2D) was optimal for reducing long wavelength divergence modes, though higher values ($\sigma=0.3$ in 2D) remained optimal for removing short wavelength errors and therefore for general use in simulations. Sub-cycling the divergence cleaning equations between timesteps was not found to have any adverse effect on the quality of the solution of the Brio-Wu shocktube test (Section~\ref{sec:divbzero-accuracy}), indeed only leading to further reductions in divergence error. In summary, we recommend that our new divergence cleaning method be universally adopted over the previous method. The previous method had a numerical issue which could cause, in certain circumstances, an increase in magnetic energy and divergence error that would reduce the effectiveness of divergence cleaning. Though this effect is likely small in practical simulations, adopting the new method removes this source of energy growth, potentially yielding improvements in the reduction of average divergence error with lower associated numerical dissipation. It is trivial to adapt existing codes to evolve $\psi / c_{\rm h}$, and doing so provides a more robust, numerically stable method at no additional computational expense. | 16 | 7 | 1607.02394 |
1607 | 1607.00672_arXiv.txt | {Until recently it was thought that high Galactic latitude clouds were a non-star-forming ensemble. However, in a previous study we reported the discovery of two embedded clusters (ECs) far away from the Galactic plane ($\sim5$ kpc). In our recent star cluster catalogue we provided additional high and intermediate latitude cluster candidates. } {This work aims to clarify if our previous detection of star clusters far away from the disc represents just an episodic event or if the star cluster formation is currently a systematic phenomenon in the Galactic halo. We analyse the nature of four clusters found in our recent catalogue and report the discovery of three new ECs with unusually high latitude and distance from the Galactic disc midplane. } {The analysis is based on 2MASS and WISE colour-magnitude diagrams (CMDs), and stellar radial density profiles (RDPs). The CMDs are built by applying a field-star decontamination procedure, which uncovers the cluster's intrinsic CMD morphology.} { All of these clusters are younger than 5 Myr. The high-latitude ECs C 932, C 934, and C 939 appear to be related to a cloud complex about 5 kpc below the Galactic disc, under the Local arm. The other clusters are above the disc, C 1074 and C 1100 with a vertical distance of $\sim3$ kpc, C 1099 with $\sim2$ kpc, and C 1101 with $\sim1.8$ kpc.} { According to the derived parameters there occur ECs located below and above the disc, which is an evidence of widespread star cluster formation throughout the Galactic halo. Thus, this study represents a paradigm shift, in the sense that a sterile halo becomes now a host of ongoing star formation. The origin and fate of these ECs remain open. There are two possibilities for their origin, Galactic fountain or infall. The discovery of ECs far from the disc suggests that the Galactic halo is more actively forming stars than previously thought and since most ECs do not survive the \textit{infant mortality} it may be raining stars from the halo into the disc, and/or the halo harbours generations of stars formed in clusters like those hereby detected.} | \label{sec:1} Embedded clusters (ECs) are the first evolutionary stage of open clusters and provide a means to explore the stellar content in gas and dust enshrouded complexes \citep[e.g.][]{Tutukov78, Lada03, Camargo11, Camargo12}. In particular distances and ages can be constrained more accurately. In the Galaxy most young open clusters and embedded clusters are essentially located within the thin disc below $\sim250$ pc from the Galactic plane \citep[e.g.][]{Camargo13, Camargo15c}. However, we discovered recently two ECs (Camargo 438 and Camargo 439) within a high-latitude cloud \citep[][hereafter Paper I]{Camargo15b} using WISE \citep{Wright10}. These clusters appear to be related to the high latitude cloud HRK 81.4-77.8 \citep{Heiles88}. Subsequently, in an extended version of our cluster catalogue \citep{Camargo16} we found some other ECs projected close to or on high and intermediate Galactic latitude clouds. Such results suggested ongoing star formation in the Galactic halo. The Galactic halo is populated by HI clouds known as intermediate and high-latitude clouds (HLCs), which are traced by the 21 cm hyperfine transition line \citep{Blitz84, Dickey90, Magnani96}. HLCs as a rule present kinematics inconsistent with Galactic rotation, and are designated as high velocity (HVCs) and intermediate velocity clouds (IVCs), in contrast to low velocity ones (LVCs) that populate the Galactic disc \citep{Muller63, Kuntz96, Martin15}. HLCs appear to be common in disc galaxies such as the Milky Way, and some of them show signs of recent or ongoing mergers \citep{Fraternali01, Thilker04, Battaglia06, Heald07, Oosterloo07, Sancisi08, Cresci10}. Regarding their origin, there is no consensus and both Galactic and extragalactic sources have been proposed. \begin{figure*} \centering \begin{minipage}[b]{0.8\linewidth} \begin{minipage}[b]{0.49\linewidth} \includegraphics[width=\textwidth]{aa3.eps} \put(-180.0,195.0){\makebox(0.0,0.0)[5]{\fontsize{14}{14}\selectfont \color{green} C 1074}} \end{minipage}\hfill \vspace{0.02cm} \begin{minipage}[b]{0.49\linewidth} \includegraphics[width=\textwidth]{aa4.eps} \put(-180.0,195.0){\makebox(0.0,0.0)[5]{\fontsize{14}{14}\selectfont \color{green} C 939}} \end{minipage}\hfill \vspace{0.02cm} \begin{minipage}[b]{0.49\linewidth} \includegraphics[width=\textwidth]{aa1.eps} \put(-180.0,195.0){\makebox(0.0,0.0)[5]{\fontsize{14}{14}\selectfont \color{green}C 1099}} \end{minipage}\hfill \vspace{0.02cm} \begin{minipage}[b]{0.49\linewidth} \includegraphics[width=\textwidth]{aa6.eps} \put(-180.0,195.0){\makebox(0.0,0.0)[5]{\fontsize{14}{14}\selectfont \color{green}C 934}} \end{minipage}\hfill \vspace{0.02cm} \caption[]{WISE ($15'\times15'$) multicolour images centred on the central coordinates of the embedded clusters C 1074, C 939, C 1099, and C 934. North is to the top and east to the left. Circles encompass more probable cluster stars (Sect. 2).} \label{f1} \end{minipage}\hfill \end{figure*} \begin{table}[!hb] \centering {\footnotesize \caption{Positions of the present star clusters or candidates.} \vspace{-0.3cm} \label{tab1} \renewcommand{\tabcolsep}{2.6mm} \renewcommand{\arraystretch}{1,2} \begin{tabular}{lrrrrr} \hline \hline Target&$\alpha(2000)$&$\delta(2000)$&$\ell$&$b$\\ &(h\,m\,s)&$(^{\circ}\,^{\prime}\,^{\prime\prime})$&$(^{\circ})$&$(^{\circ})$ \\ ($1$)&($2$)&($3$)&($4$)&($5$)\\ \hline C 932 &2:05:02&-18:09:26&188.89&-70.83\\ C 934 &2:05:54&-17:56:17&188.65&-70.54\\ C 939 &2:07:08&-18:13:15&189.83&-70.43\\ C 1074 &10:39:27&-2:00:39&250.15&46.89\\ C 1099 &11:49:55.7&-32:41:42.8&288.23&28.41\\ C 1100 &12:11:39.9&-34:44:45.5&293.73&27.41\\ C 1101 &12:14:24.5&-35:02:04.2&294.41&27.22\\ \hline \end{tabular} \begin{list}{Table Notes.} \item Cols. $2-3$: Central coordinates. Cols. $4-5$: Corresponding Galactic coordinates. \end{list} } \end{table} HLCs may be produced by an internal engine based on the stellar feedback. In this process winds from OB stars and supernovae blow away gas and dust from the disc in a chimney-like scenario, which subsequently fall back onto the disc as \textit{Galactic fountains} \citep{Shapiro76, Bregman80}. During this phase they can merge to form molecular clouds. \textit{Chimneys} powered by multiple supernovae within OB associations may blow superbubbles, which can throw gas/dust away on kiloparsec scales \citep{Quilis01, Pidopryhora07}. In this sense, \citet{Melioli08} point out that typical Galactic OB associations with 100 SNe may eject gas/dust up to $\sim2$ kpc \citep{Quilis01, Pidopryhora07}, but \citet{Melioli09} argue that even multiple OB association cannot throw dust beyond 3.5 kpc. There are two possibilities for extragalactic origin, \textit{(i)} primordial cold dark-matter encapsulated clouds accreted directly from the intergalactic medium or \textit{(ii)} infall of dark-matter free clouds remaining from tidal disruption of dwarf galaxies and galaxy collisions \citep{Oort66, Blitz99, Putman04, Keres05, Kaufmann06, Oosterloo07, Sancisi08, Kaufmann10, Hammer15, Wolfe15, Tepper15}. Gas accretion is needed to provide the low-metallicity material required by chemical evolution models \citep{Chiappini01}, since in the \textit{Galactic fountain} apparently the gas falls back onto the Galactic disc close to its original locus \citep{Melioli08, Melioli09} and does not affect the metal abundance \citep{Spitoni09}. Besides, the Galaxy on large timescales is apparently forming stars at a constant rate \citep{Binney00, Fraternali14}, which implies that its gaseous content is being continuously replenished by infall of low metallicity gas \citep{Fraternali12, Joung12}. \begin{figure*} \centering \begin{minipage}[b]{0.8\linewidth} \begin{minipage}[b]{0.49\linewidth} \includegraphics[width=\textwidth]{supernova3.eps} \end{minipage}\hfill \begin{minipage}[b]{0.49\linewidth} \includegraphics[width=\textwidth]{supernova4.eps} \end{minipage}\hfill \caption[]{WISE multicolour and W1 images ($7'\times7'$) centred on C 1100. C 1100 shows a dust emission shell. North is to the top and east to the left.} \label{f2} \end{minipage}\hfill \end{figure*} \begin{table*} {\footnotesize \begin{center} \caption{Fundamental parameters and Galactocentric components for the ECs in this work.} \renewcommand{\tabcolsep}{3.4mm} \renewcommand{\arraystretch}{1.0} \begin{tabular}{lrrrrrrr} \hline \hline Cluster&$A_V$&Age&$d_{\odot}$&$R_{GC}$&$x_{GC}$&$y_{GC}$&$z_{GC}$\\ &(mag)&(Myr)&(kpc)&(kpc)&(kpc)&(kpc)&(kpc)\\ ($1$)&($2$)&($3$)&($4$)&($5$)&($6$)&($7$)&($8$)\\ \hline C 932 &$1.40\pm0.03$&$2\pm1$&$5.7\pm0.53$&$10.55\pm0.29$&$-9.07\pm0.17$&$-0.29\pm0.03$&$-5.38\pm0.50$\\ C 934 &$1.46\pm0.06$&$2\pm1$&$5.31\pm0.51$&$10.27\pm0.27$&$-8.97\pm0.17$&$-0.27\pm0.03$&$-5.01\pm0.48$\\ C 939 &$1.30\pm0.06$&$3\pm2$&$5.40\pm0.50$&$10.34\pm0.27$&$-9.00\pm0.17$&$-0.31\pm0.03$&$-5.09\pm0.47$\\ C 1074 &$0.93\pm0.06$&$3\pm1$&$4.14\pm0.39$&$9.12\pm0.15$&$-8.18\pm0.09$&$-2.66\pm0.25$&$3.02\pm0.28$\\ C 1099 &$0.71\pm0.06$&$5\pm1$&$4.32\pm0.61$&$7.32\pm0.30$&$-6.03\pm0.17$&$-3.61\pm0.51$&$2.05\pm0.28$\\ C 1100 &$0.93\pm0.06$&$1\pm1$&$6.87\pm0.36$&$8.00\pm0.23$&$-4.76\pm0.13$&$-5.59\pm0.29$&$3.16\pm0.16$\\ C 1101 &$0.96\pm0.06$&$3\pm1$&$3.91\pm0.55$&$6.83\pm0.27$&$-5.78\pm0.20$&$-3.16\pm0.44$&$1.78\pm0.25$\\ \hline \end{tabular} \begin{list}{Table Notes.} \item Col. 2: $A_V$ in the cluster central region. Col. 2: age, from 2MASS photometry. Col. 3: distance from the Sun. Col. 4: $R_{GC}$ calculated using $R_{\odot}=8.3$ kpc as the distance of the Sun to the Galactic centre. Cols. 5 - 8: Galactocentric components. \end{list} \label{tab2} \end{center} } \end{table*} \begin{table*} {\footnotesize \begin{center} \begin{minipage}[b]{0.7\linewidth} \caption{Structural parameters for the high latitude embedded cluster C 939.} \renewcommand{\tabcolsep}{1.9mm} \renewcommand{\arraystretch}{1.6} \begin{tabular}{lrrrrrrr} \hline \hline Cluster&$(1')$&$\sigma_{0K}$&$R_{core}$&$R_{RDP}$&$\sigma_{0K}$&$R_{core}$&$R_{RDP}$\\ &($pc$)&($*\,pc^{-2}$)&($pc$)&($pc$)&($*\,\arcmin^{-2}$)&($\arcmin$)&($\arcmin$)\\ ($1$)&($2$)&($3$)&($4$)&($5$)&($6$)&($7$)&($8$)\\ \hline C 939 &$1.36$&$4.1\pm1.2$&$1.8\pm0.4$&$13.6\pm2.0$&$7.65\pm2.15$&$1.32\pm0.29$&$10.0\pm1.5$\\ \hline \end{tabular} \begin{list}{Table Notes.} \item Col. 2: arcmin to parsec scale. To minimise degrees of freedom in RDP fits with the King-like profile, $\sigma_{bg}$ was kept fixed (measured in the respective comparison field) while $\sigma_{0}$ and $R_{core}$ were allowed to vary. \end{list} \label{tab3} \end{minipage} \end{center} } \end{table*} There is evidence that IVCs ($|v_{LSR}|=50-100\,km/s$) arise from \textit{Galactic fountains}, while HVCs ($|v_{LSR}|>100\,km/s$) appear to be related to infalling gas \citep{Putman12}. However, if not disrupted on their infall, massive dark-matter free HVCs may be decelerated to LVCs mainly by ram-pressure stripping, Rayleigh-Taylor and Kelvin-Helmholtz instabilities, and dragging forces, leading to substructured clouds \citep{Benjamin97, Maller04, Heitsch09, Plockinger12, Hernandez13}. Given the implications of our recent results, a new study was necessary to verify if star formation in the Galactic halo is systematic or an episodic event. Thus, in this work we derive parameters and discuss the properties of seven high and intermediate Galactic latitude ECs. This paper is organized as follows. In Sect.~\ref{sect2} we present the cluster sample and derive the fundamental parameters. In Sect.~\ref{sect3} we discuss the results, and in Sect.~\ref{sect4} we provide the concluding remarks. \begin{figure} \resizebox{\hsize}{!}{\includegraphics{aa5.eps}} \put(-230.0,230.0){\makebox(0.0,0.0)[5]{\fontsize{14}{14}\selectfont \color{green} C 932}} \caption[]{WISE multicolour image ($15'\times15'$) centred on C 932. C 932 is deeply embedded. North is to the top and east to the left.} \label{f2b} \end{figure} \begin{figure} \resizebox{\hsize}{!}{\includegraphics{aa2b.eps}} \put(-230.0,300.0){\makebox(0.0,0.0)[5]{\fontsize{14}{14}\selectfont \color{green} C 1101}} \caption[]{WISE ($15'\times15'$) multicolour images centred on the central coordinates of the embedded cluster C 1101. North is to the top and east to the left.} \label{f3} \end{figure} \begin{figure} \centering \begin{minipage}[b]{0.85\linewidth} \resizebox{\hsize}{!}{\includegraphics{C1064.eps}} \end{minipage} \vspace{-0.2cm} \caption[]{C 1074: example of decontamination procedure. Top panels: 2MASS observed CMDs. Middle panels: equal area comparison field. Bottom panels: field-star decontaminated CMDs fitted with PARSEC isochrones for MS and PMS stars. We also show the reddening vector for $A_v=1$ to $5$.} \label{cmd1a} \end{figure} | \label{sect4} In Paper I we reported two ECs (C 438 and C 439) within a high-latitude cloud, which were the first high-latitude embedded clusters discovered. In this study we find new results about star cluster formation in high-latitude clouds and further infer that the Galactic halo is currently forming stars within ECs. Using 2MASS and WISE photometry we analysed the nature of seven ECs at high and intermediate Galactic latitudes, three of them first reported in this work (C 1099, C 1100, and C 1101). The age of these clusters are in the range of 1 to 5 Myr. C 932, C 934, and C 939 are high-latitude ECs projected within the newly identified cloud complex including CBB 188.13-70.84. These clusters are located at a vertical distance of about 5 kpc below the Galactic disc. C 1074, C 1099, C 1100, and C 1101 are in the range 1.7 to 3.2 kpc above the disc. The clusters show decontaminated CMDs with the typical pattern of MS and PMS stars in embedded clusters (Paper I and references therein). Their spatial distribution is given in Fig.~\ref{Milky1}. We have gathered a significant collection of very young star clusters in the halo. Paper I and the present additional study point to a paradigm shift in the halo, which becomes an ongoing site of star formation in the Galaxy. \vspace{0.8cm} \textit{Acknowledgements}: We thank an anonymous referee for valuable comments and suggestions. This publication makes use of data products from the Two Micron All Sky Survey (2MASS) and Wide-field Infrared Survey Explorer (WISE). The 2MASS is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Centre/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. WISE is managed and operated by NASA’s Jet Propulsion Laboratory (JPL) in Pasadena, California and is a project of the JPL/California Institute of Technology, funded by the National Aeronautics and Space Administration. The spacecraft scanned the entire sky twice. E. Bica and C. Bonatto acknowledge support from CNPq (Brazil). | 16 | 7 | 1607.00672 |
1607 | 1607.04784_arXiv.txt | We present optical spectra of the nearby Type~Ia supernova SN~2011fe at 100, 205, 311, 349, and 578 days post-maximum light, as well as an ultraviolet spectrum obtained with \textit{Hubble Space Telescope} at 360 days post-maximum light. We compare these observations with synthetic spectra produced with the radiative transfer code \texttt{PHOENIX}. The day +100 spectrum can be well fit with models which neglect collisional and radiative data for forbidden lines. Curiously, including this data and recomputing the fit yields a quite similar spectrum, but with different combinations of lines forming some of the stronger features. At day +205 and later epochs, forbidden lines dominate much of the optical spectrum formation; however, our results indicate that recombination, not collisional excitation, is the most influential physical process driving spectrum formation at these late times. Consequently, our synthetic optical and UV spectra at all epochs presented here are formed almost exclusively through recombination-driven fluorescence. Furthermore, our models suggest that the ultraviolet spectrum even as late as day +360 is optically thick and consists of permitted lines from several iron-peak species. These results indicate that the transition to the ``nebular'' phase in Type~Ia supernovae is complex and highly wavelength-dependent. | \label{sec:intro} The Type~Ia supernova (SN~Ia) SN~2011fe was discovered on 2011 August 24, just 11~hr after explosion \citep{nugent11}. It is among the nearest ($\sim 6.9$~Mpc) and youngest ($\sim 11$~hr) SNe~Ia ever discovered. Extensive spectroscopic and photometric studies of SN~2011fe indicate that it is ``normal'' in nearly every sense: in luminosity, spectral and color evolution, abundance patterns, etc. \citep{parrent12,richmond12,roepke12,vinko12,munari13,pereira13}. Its unremarkable nature coupled with the wealth of observations made over its lifetime render it an ideal laboratory for understanding the physical processes which govern the evolution of normal SNe~Ia. Indeed, these data have allowed observers to place numerous and unprecedented constraints on the progenitor system of a particular SN~Ia \citep[e.g.,][]{li11,nugent11,bloom12,chomiuk12,horesh12,margutti12}. Equally as information-rich as observations taken at early times are those taken much later, when the supernova's photosphere has receded and spectrum formation occurs deep in the SN core. For example, \citet{shappee13} used late-time spectra to further constrain the progenitor system of SN~2011fe, namely that the amount of hydrogen stripped from the putative companion must be $< 0.001~M_\odot$. \citet{mcclelland13} found that the luminosity from SN~2011fe in the 3.6~$\mu$m channel of \textit{Spitzer}/IRAC fades almost twice as quickly as in the 4.5~$\mu$m channel, which they argue is a consequence of recombination from doubly ionized to singly ionized iron peak elements. In addition, \citet{kerzendorf14} used photometric observations near 930~d post-maximum light to construct a late-time quasi-bolometric light curve, and showed that the luminosity continues to trace the radioactive decay rate of $^{56}$Co quite closely, suggesting that positrons are fully trapped in the ejecta, disfavoring a radially combed or absent magnetic field in this SN. \citet{graham15} presented an optical spectrum at 981~d post-explosion and used constraints on both the mass of hydrogen as well as the luminosity of the putative secondary star as evidence against a single-degenerate explosion mechanism. \citet{taubenberger15} presented an optical spectrum at 1034~d post-explosion, and speculated about the presence of [\ion{O}{1}] lines near 6300~\AA, which, if confirmed, would provide strong constraints on the mass of unburned material near the center of the white dwarf progenitor of SN~2011fe. Non-detections of the H$\alpha$ line at both of these very late epochs also strengthened the constraints on the presence of hydrogen initially posed by \citet{shappee13}. Finally, \citet{mazzali15} used spectrum synthesis models of SN~2011fe from 192 to 364 days post-explosion to argue for a large central mass of stable iron and a small mass of stable nickel -- about 0.23~$M_\odot$ and 0.01~$M_\odot$, respectively. We complement these various late-time analyses with a series of radiative transfer models corresponding to a series of optical and ultraviolet (UV) spectra of SN~2011fe. | We extended \texttt{PHOENIX} to calculate radiative transfer models well into the late-time epochs of SNe~Ia, with an eye toward obtaining good fits to the high-quality optical and UV spectra of SN~2011fe. Doing so required similar methods to those discussed in \citet{friesen14}, in particular using an alternative method to that of Uns\"old-Lucy for calculating the temperature structure of the gas, as well as accounting for the collisional and radiative rate data for forbidden lines, which behave quite differently than permitted lines. The resulting synthetic spectra, ranging from +100 to +578 days post-maximum light, vary in degrees of fidelity to corresponding observed spectra of SN~2011fe, with the earlier epochs fitting quite well and the later epochs less so. At day +100 we found that radiative transfer calculations which neglect forbidden lines and those which include them can produce remarkably similar optical spectra, but with quite different atomic species and combinations of lines forming the various features. We found that, at least as late as day +360, permitted lines such as \ion{Ca}{2} H \& K and IR3 continue to influence spectrum formation in the optical, and permitted lines of \ion{Fe}{2} form much of the spectrum in the UV. In addition, these models indicate that some emission features from permitted lines are replaced by other emission features of forbidden lines at nearly the same wavelength as the SN evolves. For example, the emission from \ion{Ca}{2} H \& K at 4000~\AA\ is replaced around day +205 by [\ion{Fe}{3}]~$\lambda 4008$~\AA, and the emission from \ion{Ca}{2} IR3 at around 8600~\AA\ is replaced by [\ion{Fe}{2}]~$\lambda 8617$~\AA. | 16 | 7 | 1607.04784 |
1607 | 1607.08755_arXiv.txt | { We perform a forecast analysis of the capability of the eLISA space-based interferometer to constrain models of early and interacting dark energy using gravitational wave standard sirens. We employ simulated catalogues of standard sirens given by merging massive black hole binaries visible by eLISA, with an electromagnetic counterpart detectable by future telescopes. We consider three-arms mission designs with arm length of 1, 2 and 5 million km, 5 years of mission duration and the best-level low frequency noise as recently tested by the LISA Pathfinder. Standard sirens with eLISA give access to an intermediate range of redshift $1\lesssim z \lesssim 8$, and can therefore provide competitive constraints on models where the onset of the deviation from $\Lambda$CDM (i.e.~the epoch when early dark energy starts to be non-negligible, or when the interaction with dark matter begins) occurs relatively late, at $z\lesssim 6$. If instead early or interacting dark energy is relevant already in the pre-recombination era, current cosmological probes (especially the cosmic microwave background) are more efficient than eLISA in constraining these models, except possibly in the interacting dark energy model if the energy exchange is proportional to the energy density of dark energy. } | % \label{sec:introduction} With the first direct detection of gravitational waves (GWs) by the LIGO/Virgo collaboration \cite{Abbott:2016blz,Abbott:2016nmj,TheLIGOScientific:2016pea} the era of GW astronomy has begun. The information gathered from present and future GW observations will improve our understanding of the astrophysical objects emitting the GW signal, of the origin and evolution of the universe and its structure, and of the gravitational interaction. Earth-based detectors, such as the advanced LIGO \cite{ligo} and Virgo \cite{virgo} interferometers, target the GW frequency window $10 - 1000$ Hz, while pulsar timing arrays (PTA), such as the ones united under the International Pulsar Timing Array (IPTA) collaboration \cite{2010CQGra..27h4013H}, probe much lower frequencies around $10^{-9} - 10^{-8}$ Hz. In order to fill the gap in frequency between Earth-based interferometers and PTA, space-born GW observatories have been proposed, which will be able to reach high sensitivity in the frequency band $10^{-4} - 10^{-1}$ Hz. Such range of frequencies in the GW landscape is supposed to be rich of astrophysical sources and it is yet completely unexplored. In particular the strong GW signal emitted by merging massive black hole binaries (MBHBs) from $10^4$ to $10^7$ solar masses is expected to fall exactly within the frequency band targeted by space-born detectors. Since such massive black holes are believed to reside at the centre of galaxies, observing the GW signal they emit will help to better understand the formation and evolution of galaxies and cosmic structures. In 2013, the European Space Agency has approved a GW observatory in space as the L3 mission of its ``Cosmic Vision Program'' scheduled for launch around 2030-2034, for which the ``evolved LISA'' (eLISA) space-based interferometer is the main candidate \cite{elisaweb,Seoane:2013qna}. eLISA is designed to probe the GW landscape around the mHz region where the signal produced by MBHBs is expected to be the loudest. The final design of the mission, which is composed by three satellites orbiting around the Sun in an equilateral triangular formation, has not been decided yet, and some variables are still under considerations (see e.g. \cite{Klein:2015hvg}): the number of laser links between the satellites (four or six), corresponding to the number of active arms (two or three); the arm-length of the triangle (from one to five million km); and the duration of the mission (two or five years). The eLISA low-frequency noise level (another of the variables previously considered) has been recently tested by the LISA Pathfinder mission \cite{pathfinderweb}, and according to the first results \cite{Armano:2016bkm} the expected noise is almost one hundred times better than the original requirement for the instrument. In an earlier work, Ref.~\cite{Tamanini:2016zlh}, we have studied the capability of eLISA in probing the acceleration of the universe by means of MBHB mergers as {\it standard sirens}, i.e.~as sources of known distance \cite{schutz,Holz:2005df,Cutler:2009qv}. We have derived eLISA constraints on standard cosmological models: $\Lambda$CDM, dynamical dark energy, non-zero spatial curvature and so on. In the present paper, we specifically consider alternative scenarios to explain the acceleration of the universe: in particular, we study early and interacting dark energy models, see sections \ref{sec:early_DE} and \ref{sec:interacting_dark_energy}. The principle of standard sirens is the following. The measured GW waveform depends directly on the luminosity distance of the source, and thus parameter estimation allows to infer the distance to the source for every GW event detected. If subsequent electromagnetic (EM) observations are able to identify an EM counterpart, then one is able to obtain a measure of the source redshift and thus a point in the distance-redshift space. Once a sufficient number of standard sirens is observed, the theoretically predicted distance-redshift relation can be compared against the data and constraints on the cosmological parameters can be statistically inferred. We assume spatial flatness throughout the paper, so that \begin{equation} d_L(z) = c \left(1+z\right) \int_0^z \frac{1}{H(z')} dz' \,, \label{eq:dist_red_rel} \end{equation} where $c$ is the speed of light and $H(z)$ is the Hubble rate. The analysis presented here is meant to complete the one performed in \cite{Tamanini:2016zlh}. We concentrate on early (EDE) and interacting (IDE) dark energy, but use the same standard siren catalogues that have been obtained in \cite{Tamanini:2016zlh} starting from simulated rates of MBHB mergers detectable by different eLISA configurations, and considering realistic scenarios for the observation of the EM counterparts, based on the capabilities of future EM telescopes (LSST, SKA, ELT). Here we also consider the same three astrophysical models of MBHB formation and evolution appearing in \cite{Tamanini:2016zlh}, namely a light seeds model (popIII), a heavy seeds model with delay (Q3d) and a heavy seeds model without delay (Q3nod) (see also \cite{Klein:2015hvg} and references therein for more information). We present separate results for all these models. The number of standard sirens for each MBHB formation scenario has been selected under the hypothesis that the sky localisation of the event can be achieved using also the merger and ringdown phases of the signal. In this procedure the telescopes can be pointed only after the merger to look for a distinctive signature, therefore one implicitly assumes that there is a delay between the merger and the flare, or that the electromagnetic signal is persistent and peculiar enough that it can be confidently identified also minutes to hours after merger. This procedure was labelled the ``optimistic scenario'' in \cite{Tamanini:2016zlh}. Moreover, the statistical methods employed to handle the simulated data coincide with the ones adopted in \cite{Tamanini:2016zlh}: in particular we perform a Fisher matrix analysis and obtain constraints and contour plots following the procedures exposed in \cite{Tamanini:2016zlh}. Differently from \cite{Tamanini:2016zlh}, in what follows we consider only three eLISA configurations, letting the arm-length to vary as one (A1), two (A2) and five (A5) million km, but fixing the number of laser links to six (L6), the mission duration to five years (M5) and the low-frequency noise to the LISA Pathfinder ``expected'' one (N2) (see \cite{Tamanini:2016zlh} and \cite{Klein:2015hvg} for details). The reasons for this choice are the following: \begin{itemize} \item The aim of the present paper is not to carefully analyse all possible eLISA configurations to understand the science return of each of them (as it was in \cite{Tamanini:2016zlh}), but rather to investigate simple extensions of the $\Lambda$CDM model in order to understand the pros and cons of the eLISA mission in probing alternative cosmological models. \item According to the results of \cite{Tamanini:2016zlh}, four-link (two arms) configurations perform much worse than six-link (three arms) configurations in providing a sufficiently high number of MBHB standard sirens for cosmology. We therefore ignore four-link configurations since we expect that they will not be able to give meaningful constraints on the parameters of alternative cosmologies beyond $\Lambda$CDM. \item The number of detections, and thus the number of standard sirens, scales linearly with the mission duration: the longer the mission, the higher the number of datapoints. In analogy with the investigation of \cite{Tamanini:2016zlh} we thus only focus on a mission of five years\footnote{A method to estimate how the cosmological constraints change as the mission duration changes has been outlined in \cite{Tamanini:2016zlh}.}. \item Finally we only consider the ``expected'' low-frequency noise from LISA Pathfinder, called N2 \cite{Tamanini:2016zlh,Klein:2015hvg}. According to the first results of the mission \cite{Armano:2016bkm} this requirement has been met at frequencies $f>1$ mHz, while at lower frequencies the situation is still open: however one can optimistically forecast that the N2 noise level, if not a better one, will be finally achieved over the whole frequency spectrum. \end{itemize} The configuration with two million km arms, denoted N2A2M5L6, will be taken as our reference design for eLISA upon which the majority of subsequent results are based. When we need to pick a specific MBHB formation model, we choose the popIII scenario which is the one providing an intermediate number of standard sirens. In what follows we first investigate EDE in section \ref{sec:early_DE} and then analyse two different models of IDE in section \ref{sec:interacting_dark_energy}. We have chosen EDE and IDE as alternative cosmological models because they are simple one-parameter extensions of $\Lambda$CDM and they allow to better expose the advantages of eLISA in probing the expansion of the Universe at high redshift. In the following, we set the fiducial values of the parameters to $\Omega_m^0=0.3$,~$w_0=-1$,~$h=0.67$,~ $\Omega_{de}^e=0$, $\epsilon_1=0$, $\epsilon_2=0$. Section \ref{sec:discussion_and_conclusion} contains discussions and conclusions. | % \label{sec:discussion_and_conclusion} We have presented a forecast analysis of the capabilities of the eLISA mission to constrain two alternative models of dark energy, namely early and interacting dark energy. These models have been widely studied in the literature, and previous analyses have shown the advantages of testing them using data at various redshifts, combining different observational probes such as CMB, BAO, SNIa, LSS, weak lensing, RSD. The motivation of this work resides in the fact that standard sirens with eLISA can provide access to an intermediate range of redshift $1\lesssim z \lesssim 8$, higher than what can be reached with SNIa and matter structure data. Furthermore, the measurement of the luminosity distance with standard sirens being given by gravitational waves, it provides an independent access to the distance-redshift relation, partly complementary to electromagnetic observations. In the present analysis we have used the same procedure developed in \cite{Tamanini:2016zlh}: we have started from simulations of the event rates of MBHB in three different models for the BH seeds, and we have used realistic scenarios for the occurrence and detection of the EM counterparts. Equipped with catalogues of standard sirens, we have selected those which are visible by different eLISA configurations setting the thresholds of SNR$>8$ and sky localisation better than 10 ${\rm deg}^2$. These have been achieved including both the inspiral and merger and ringdown phases of the GW event (the ``optimistic scenario'' in \cite{Tamanini:2016zlh}). Since in Ref.~ \cite{Tamanini:2016zlh} we demonstrated that eLISA configurations with four links are not very powerful in probing the expansion of the universe, here we concentrated only on six-link configurations. We have also fixed the duration of the mission to five years (as done in \cite{Tamanini:2016zlh}) and the noise level to the Pathfinder expected one (N2), which is justified after the success of the LISA Pathfinder mission \cite{Armano:2016bkm}. We have therefore considered three eLISA configurations: N2A1M5L6, N2A2M5L6, N2A5M5L6. The main result of the present analysis is that standard sirens with eLISA can be competitive in constraining EDE and IDE models if the onset of the deviation from $\Lambda$CDM (i.e.~the epoch when EDE starts to be non-negligible, or when the interaction with DM begins) occurs relatively late, at $z \lesssim 6$. Models for which the deviation from $\Lambda$CDM starts far in the past, typically before recombination, are well constrained by current cosmological probes, in particualr by the CMB; the present constraints are way better than those that can be achieved with even the best configurations of eLISA (except perhaps for the IDE2 case, depending on the analysis one compares with: c.f. discussion at the end of Sec.~\ref{sec:IDE2}). On the other hand, if the deviation starts relatively late, the present observational constraints on both EDE and IDE models are highly degraded, and eLISA becomes therefore competitive in testing these scenarios. This happens because the redshift distribution of standard sirens peaks between $2\leq z \leq 4$, and very few standard sirens are available at redshift larger than six. The errors on the EDE and IDE parameters beyond $\Lambda$CDM (namely $w_0$, $\Omega_{de}^e$, $\epsilon_{1}$, $\epsilon_2$), when all the other cosmological parameters are held fixed, decrease with the increase of the redshift at which the deviation from $\Lambda$CDM starts. However, this occurs up to a redshift for the onset of the deviation of about six; for higher deviation redshift, the eLISA errors stabilize and do not change appreciably up to far in the radiation era. This behaviour of the errors follows the redshift distribution of the standard sirens (see appendix~\ref{sec:redshift_distribution_of_standard_sirens}) which are available for the measurement: after a redshift of about six, the number of standard sirens detected does not increase sufficiently any longer to provide a better measurement of the cosmological parameters. We have also demonstrated, however, that this behaviour can be affected by degeneracies among the parameters, in particular among $\Omega_m^0$ and $\Omega_{de}^e$ or $\epsilon_{1}$, $\epsilon_2$ (a degeneracy which is expected in both EDE and IDE models due to the way these parameters enter in the distance-redshift relation). If $\Omega_m^0$ is not set to its fiducial value, the errors on $\Omega_{de}^e$ or $\epsilon_{1}$, $\epsilon_2$ increase when the redshift of the onset of the deviation from $\Lambda$CDM increases. Once again this reflects the fact that the peak of the standard siren distribution resides in the interval $2\leq z \leq 4$ (see appendix~\ref{sec:redshift_distribution_of_standard_sirens}): for low deviation redshift, the bulk of the MBHB standard sirens detected by eLISA effectively probes $\Lambda$CDM or dynamical dark energy (i.e.~with $w_0\neq -1$), and this in turns leads to a reduction of the errors on $\Omega_{de}^e$ or $\epsilon_{1}$, $\epsilon_2$. This happens especially for the models where either all four, or three parameters are free to vary. Therefore, without setting an exact prior on $\Omega_m^0$, eLISA can only constrain models where the redshift of the onset of the deviation from $\Lambda$CDM is sufficiently low. Note however, that whenever the deviation from $\Lambda$CDM starts well after recombination, from the point of view of the eLISA analysis one can confidently use the very precise CMB measurement of $\Omega_m^0$ as an exact prior. We can therefore conclude that eLISA with six-link configurations will serve as an independent mean to test alternative models for the acceleration of the universe such as EDE and IDE, and will be able to improve the present constraints, in particular for the EDE and IDE2 models, if one considers low values of the redshift at which the deviation from $\Lambda$CDM starts. | 16 | 7 | 1607.08755 |
1607 | 1607.01397_arXiv.txt | % Approximately $0.2 \pm 0.2$ of white dwarfs (WDs) show signs of pollution by metals, which is likely due to the accretion of tidally disrupted planetary material. Models invoking planet-planet interactions after WD formation generally cannot explain pollution at cooling times of several Gyr. We consider a scenario in which a planet is perturbed by Lidov-Kozai oscillations induced by a binary companion and exacerbated by stellar mass loss, explaining pollution at long cooling times. Our computed accretion rates are consistent with observations assuming planetary masses between $\sim 0.01$ and $1\,M_\mathrm{Mars}$, although nongravitational effects may already be important for masses $\lesssim 0.3 \, M_\mathrm{Mars}$. The fraction of polluted WDs in our simulations, $\sim 0.05$, is consistent with observations of WDs with intermediate cooling times between $\sim 0.1$ and 1 Gyr. For cooling times $\lesssim 0.1$ Gyr and $\gtrsim 1$ Gyr, our scenario cannot explain the high observed pollution fractions of up to 0.7. Nevertheless, our results motivate searches for companions around polluted WDs. | \label{sect:introduction} The atmospheres of cool white dwarfs (WDs) are expected to consist entirely of hydrogen or helium due to efficient gravitational settling of metals \citep{1945AnAp....8..143S}. However, in $0.2 \pm 0.2$ of white dwarfs \citep{2006A&A...453.1051K,2014A&A...566A..34K}, spectra have revealed emission lines from a large range of metals, suggesting that these `polluted' WDs have recently accreted metal-rich material (see \citealt{2014AREPS..42...45J,2016RSOS....3.0571V,2016NewAR..71....9F} for reviews). Observations indicate that the pollution rate is approximately independent of cooling time \citep{2014A&A...566A..34K}, requiring a continuous pollution process. Accretion from the interstellar medium \citep{1993ApJS...87..345D} has been ruled out \citep{2003ApJ...596..477Z,2006A&A...453.1051K,2007ApJ...663.1291D,2008AJ....135.1785J}. WD pollution could instead originate from accreting tidally disrupted rocky planetary material (e.g. \citealt{1986ApJ...302..462A,1993AJ....105.1033A,2002ApJ...572..556D,2003ApJ...584L..91J}) with a composition similar to Earth's (see e.g. \citealt{2014AREPS..42...45J}, and references therein), originating from planetesimals of mass $\sim 10^{20}\,\mathrm{kg}$ to planets as massive as Mars \citep{2009ApJ...699.1473J}. This is supported by the observation that all WDs with discs are polluted, and by the observed transiting planetesimals in tight orbits around WD 1145+017 \citep{2015Natur.526..546V}. Polluted WDs are therefore a probe for planetary systems around WDs (see \citealt{2016RSOS....3.0571V} for a review). Bodies in tight orbits are engulfed by the star as it expands along the red giant branch (RGB; \citealt{2009ApJ...705L..81V,2011ApJ...737...66K,2014ApJ...794....3V}) and asymptotic giant branch (AGB; \citealt{2012ApJ...761..121M}) phases. At larger distances, stellar mass loss, tides, interactions with stellar ejecta and nongravitational effects are important. Early after WD formation, dynamical instabilities arising from planet-planet interactions and mass loss could lead to the disruption of planetary material and WD pollution \citep{2002ApJ...572..556D,2011MNRAS.414..930B,2012ApJ...747..148D,2016MNRAS.458.3942V}. These instabilities typically occur on short time-scales, and cannot explain continued pollution of WDs with cooling times of several Gyr. \citet{2015MNRAS.454...53B} proposed a scenario independent of the WD cooling time, in which the WD planetary system is perturbed by a wide binary companion whose orbit is driven to high eccentricity due to Galactic tides. We investigate a related scenario in which the WD and planet are orbited by a secondary star. We focus on planets with radii $\gtrsim 1000 \, \mathrm{km}$, for which nongravitational effects are not important (e.g. \citealt{2016RSOS....3.0571V}). Mass loss of the primary star triggers adiabatic expansion of both the inner (planet's) and outer (secondary's) orbits. The importance of Lidov-Kozai (LK) oscillations \citep{1962P&SS....9..719L,1962AJ.....67..591K} in the inner orbit then typically increases \citep{2012ApJ...760...99P,2013ApJ...766...64S,2013MNRAS.430.2262H,2014ApJ...794..122M}. Consequently, the inner orbit can be driven to high eccentricity for the planet to be tidally disrupted by the WD, polluting the latter. Pollution can be prolonged to several Gyr after the WD formed. | \label{sect:conclusions} We considered a scenario for WD pollution by planets triggered by LK oscillations induced by a binary companion. Our computed accretion rates are consistent with observations for planetary masses between $\sim 0.01$ and $1\,M_\mathrm{Mars}$. The fraction of polluted WDs is consistent with observations of WDs with intermediate cooling times ($0.1 \, \mathrm{Gyr} \lesssim t_\mathrm{cool} \lesssim 1\,\mathrm{Gyr}$). For short and long cooling times, our scenario cannot explain the high observed pollution fractions of up to 70 per cent. Our scenario may also apply to planetesimals, but further work is needed to incorporate nongravitational effects. | 16 | 7 | 1607.01397 |
1607 | 1607.08080_arXiv.txt | { The deviations of the mid-transit times of an exoplanet from a linear ephemeris are usually the result of gravitational interactions with other bodies in the system. However, these types of transit timing variations (TTV) can also be introduced by the influences of star spots on the shape of the transit profile. \\ Here we use the method of unsharp masking to investigate the photometric light curves of planets with ambiguous TTV to compare the features in their $O-C$ diagram with the occurrence and in-transit positions of spot-crossing events. This method seems to be particularly useful for the examination of transit light curves with only small numbers of in-transit data points, i.e., the long cadence light curves from {\it Kepler} satellite.\\ As a proof of concept we apply this method to the light curve and the estimated eclipse timing variations of the eclipsing binary KOI-1452, for which we prove their non-gravitational nature. Furthermore, we use the method to study the rotation properties of the primary star of the system KOI-1452 and show that the spots responsible for the timing variations rotate with different periods than the most prominent periods of the system's light curve. We argue that the main contribution in the measured photometric variability of KOI-1452 originates in g-mode oscillations, which makes the primary star of the system a $\gamma$-Dor type variable candidate.} | One of the most efficient ways to study the physical parameters of an exoplanet is by studying their transits in front of their parent star. As shown by \citetads{2005MNRAS.359..567A}, it is even possible to infer the existence of additional low-mass and otherwise unobservable companions with the measurement of the mid-transit times and their deviations from linear period ephemeris. Depending on the actual planetary configuration, these transit timing variations (TTV) can reach differences up to a few hours with respect to the linear period ephemeris of the planet, especially for systems close to low order mean motion resonances. Although additional system components can also be inferred from precise radial velocity (RV) measurements, the importance of the TTV method increases with the number of planets around fainter stars, where the quality of RV measurements is decreasing. While the measurements of TTV are possible with ground-based observations, the high accuracy, continuous light curves from the {\it CoRoT} \citepads{2006ESASP1306...33B} and {\it Kepler} \citepads{2010ApJ...713L..79K} space missions provide evidence for statistically significant TTV in more than 60 cases.\footnote{According to data list acquired by the web page http://exoplanet.org } Despite its success, the study of exoplanets with the method of transits using high precision photometry can be somewhat problematic in the case of active host stars. The transit profile of planets orbiting around spotted stars become distorted when the planet is passing over dark star spots during its transit (e.g., \citetads{2009A&A...494..391R}, \citetads{2009A&A...504..561W}). The severity of this effect can be such that quite large systematic errors are introduced in the calculation of the physical parameters of the planets (see \citetads{2009A&A...505.1277C} and \citetads{2013A&A...556A..19O}). Furthermore, \citetads{2016A&A...585A..72I} show that the influence of these spot-crossing events on the estimates of the mid-transit times of a planet can lead to statistically significant false positive TTV detections under certain conditions, and there are several cases where the detected low-amplitude TTV are actually believed to be caused by spot-crossing events (e.g., \citetads{2011ApJ...733..127S}, \citetads{2012ApJ...750..114F}, \citetads{2013A&A...553A..17S}, \citetads{2013ApJS..208...16M}). {Likewise, \citetads{2002A&A...387..969K} predict that eclipse timing variations (ETV) may be caused by spot-crossing events in the eclipses of close binaries with at least one photospherically active component (e.g., \citeads{2013ApJ...774...81T}) } The identification of the spot-crossing events and their correlation with the TTV of a planet can be cumbersome, especially for planets with small periods and thus a large number of transits. Here we show that the method of unsharp masking, i.e., examining the transit light curve produced by subtracting out the mean transit model (see Sect.~\ref{Sect:data_pre}) is a suitable tool to compare the measured $O-C$ mid-transit time variations with the transit light curves and to identify correlations between the occurrence of spot-crossing events and the measured TTV. To demonstrate the power of this method, we present a case study of the TTV of KOI-1452 in Sect.~\ref{Sect:koi1452cs}, and finally we discuss our results in Sect.~\ref{sect:1542disc}, and summarize with our conclusions in sect.~\ref{Sect:Conc_trj}. \begin{figure}[t] \begin{center} \includegraphics[width=\linewidth]{oldnnew} \caption{Top: Consecutive simulated transits of a planet with an orbital period $P_\mathrm{orb}$ = 1.0132 $P_{\star}$. Each light curve is affected by spot-crossing events; see text for details.\\ Middle: Estimated mid-transit times of the transits in the top panel. The sinusoidal shape of the $O-C$ diagram is the result of spot-crossing anomalies in the transits \citetads{2016A&A...585A..72I}.\\ Bottom: Unsharp masking transit residuals of the planet; note the correlation with the sinusoidal O-C diagram in the middle panel. } \label{fig:trad_meth} \end{center} \end{figure} | \label{sect:1542disc} \subsection{The different periods of KOI-1452} As far as the observed light curve modulations of KOI-1452 are concerned, the most straightforward interpretation would be to attribute them to starspots and one would interpret the different peaks in the L-S diagram as indicators of differential rotation \citepads[e.g.,][]{2013A&A...560A...4R}. {While F-type stars do not usually show high levels of photospheric activity, the inflated state of the KOI-1452 main component - because of its youth - allows us to consider such a possibility.} Also, if KOI-1452 is a very active star, as suggested by its {\it Kepler} light curve (see Fig.~\ref{fig:koi1452lc}), one would also expect spot-crossing events as observed, for example, for CoRoT-2 (e.g., \citeads{2009A&A...504..561W}, \citeads{2009A&A...505.1277C} and \citeads{2010A&A...514A..39H}) and, indeed, our unsharp masking analysis (see Figs.~\ref{fig:koi1452} and~\ref{fig:multi_cross}) does suggest spot-crossing events. Considering the spots that are involved in spot-crossing events and using the average of the $D_\mathrm{sp}$ values of the spot crossing features shown in Fig.~\ref{fig:multi_cross}, we estimate $D_\mathrm{sp}\simeq35\pm5$ days; we note that the sign of $D_\mathrm{sp}$ is positive since we observe the spots drifting with a direction from ingress to egress (see Sect.~\ref{Sect:data_pre}). Using Eq.~\ref{eq:spotdur}, we then compute the period of the spots responsible for the spot-crossing features in the eclipses of KOI-1452 as $P_\mathrm{rot}=\,1.1357\pm0.003$ days, assuming a prograde orbit, or $P_\mathrm{rot}=\,1.1687\pm0.003$ days, assuming a retrograde orbit. {Despite the relatively young age of the system (see Table~\ref{tab:par-1452}) we estimate - using the formulation derived by \citetads{1977A&A....57..383Z} - that the synchronization and circularization processes should occur much earlier in the lifetime of the system (i.e., $\tau_\mathrm{circ}\lesssim$1 Myr). In combination with the light curve modeling results (i.e., the eclipses have equal durations and they are separated by half orbital period $P_\mathrm{orb}$ as shown in Fig.~\ref{fig:koi1452lc}), it is reasonable to assume that the two stars are tidally locked}, i.e., the rotational periods of (actually both) stars coincide with the orbital period of the system, and thus the rotational periods ought to be close to 1.15~days. However, the highest peak in the LS periodogram is found at $P_\mathrm{rot}\simeq$~1.512 days (cf., Fig~\ref{fig:laprompeaks}). If this peak was produced by star spots rotating with this period, these spots cannot be responsible for the spot-crossing events observed in the transits of KOI-1452, since for this period ratio the difference between the transit positions of the spots in consecutive transits would be much greater. Using Eq.~\ref{eq:spotdur_tr}, we calculate that a spot with a period of $P_\mathrm{rot}\simeq$~1.512 days requires $D_\mathrm{sp\_tr}=-1.6$ rotations to cover the total transit profile, which would lead to two problems: first, these spots would not be visible continuously from one rotation to the next, resulting in completely different ETV patterns and spot-crossing features, and second, the spots should occur with a direction from egress to ingress, opposite to the direction in which they appear to drift (see Fig.~\ref{fig:multi_cross}). \begin{figure}[t] \begin{center} \includegraphics[width=\linewidth]{multi_cross} \caption{Collection with the clearest spot-crossing features from the light curve of KOI-1452; white regions are due to data absence. The black lines indicate the first and fourth contact of the eclipse. } \label{fig:multi_cross} \end{center} \end{figure} \subsection{Differential rotation~?} In Fig.~\ref{fig:koi1452_path}, we show an artist's impression of the KOI-1452 system during a primary eclipse, using the estimated orbital characteristics shown in Table~\ref{tab:par-1452} and assuming that rotational and orbital angular momenta are aligned. The semi-transparent black band represents the eclipsed part of the primary in the stellar latitude range between $25\degree\lesssim\phi_\star\lesssim50\degree$. The equatorial regions at $0\degree$ are assumed to rotate with a period of 1.15~days. \begin{figure}[t] \includegraphics[trim = 0mm 20mm 0mm 20mm, clip,width=1\linewidth]{koi1452_path} \caption{An impression of the KOI-1452 system during a primary eclipse. The horizontal lines indicate the latitudes of the system's primary, which appears tilted according to our hypothesis of zero system obliquity. The semi-transparent band represents the latitudes occulted during an eclipse. } \label{fig:koi1452_path} \hspace{-30pt} \end{figure} It is hard to see where regions rotating with 1.51~days should be located on the star, and also the level of differential rotation required (period difference of $\sim$30\% between neighboring latitudes) is very large. \citetads{2010MNRAS.404.1263B} discuss the possibility of zonal stellar surface flows, i.e., flows similar to those observed on the surfaces of the giant planets in our solar system. Following this hypothesis would lead to accounting for the non-linear evolution of the differential rotation along the stellar latitude. As a result, this might explain the required rapid alternations in the rotation period. Also, in the case of the giant planets in the solar system, the velocity differences observed in the zonal bands are in the percent range, but not 30\%. An even more exotic hypothesis would be that the differential rotation of the star is opposite to that of the Sun, i.e., the star rotates more slowly in the equator. Anti-solar differential rotation would explain the fact that the spots occulted by the companion spots, which are slower, appear in high latitudes. In addition, it would explain the dominance of the slow spots in the modulations of the light curve, as they would appear closer to the equator and thus larger (as a result of the projection effect). Although this is a plausible explanation, anti-solar differential rotation is proposed for evolved stars \citepads{2009ApJ...702.1078B}. \subsection{Gravity-mode pulsations~?} Both the area of the Hertzsprung-Russell diagram, where the main component of the KOI-1452 system is located, and the nature of the periodicities found in the {\it Kepler} light curve, make it a candidate for a $\gamma$-Dor type variable star. These stars are {F-type stars with effective temperatures of between $\sim$6 900 K and $\sim$7$\,$500 K, which lie on or just above the main sequence. They are} slightly more massive than the Sun (1.4-2.5$M_\odot$) and show multi-periodic oscillations with periods between 0.3 and 3 days, which are the result of gravity-mode (g-mode) pulsations \citepads[e.g.,][]{1999PASP..111..840K}. The mechanism responsible for the g-mode pulsations is related to the fluctuations of the radiative flux at the base of the convective envelope of the stars \citepads{2000ApJ...542L..57G}. Assuming a non-rotating star with a chemically homogeneous radiative envelope, one should expect pulsations with identical spherical degree $l$ and a variety of radial orders $n$ with equidistantly spaced periods \citepads{1980ApJS...43..469T}. Furthermore, stellar rotation introduces frequency splitting, which results in separated period spacing patterns, relative to the azimuthal order $m$. \citetads{2015ApJS..218...27V} present a survey of several dozens of $\gamma$-Dor variable stars. While the majority of these stars have slightly shorter pulsation periods and have period spacings again somewhat shorter than the period spacing of P$~=~$0.1063~days observed for KOI-1452 (see bottom panel of Fig.~\ref{fig:laprompeaks}), the observed parameters for KOI-1452 do not fall completely outside the range observed for $\gamma$~Dor variable stars. While an in-depth asteroseismic analysis of KOI-1452 is beyond the scope of this paper, the hypothesis that stellar pulsations and not stellar activity is responsible for the observed light curve modulations of KOI-1452, appears to be the simplest hypothesis and would not require the introduction of somewhat exotic differential rotation scenarios. Yet, the spot-crossing events visible in the primary eclipses of the binary suggest that, in addition, activity modulations must exist which are probably described by the part of the L-S periodogram in the period range between 1.13 days and 1.17 days. \subsection{Retrograde orbit~?} The assumption that the most dominant features in the L-S periodogram of the KOI-1452 system are the result of g-mode pulsations could explain the large discrepancy between the measured periodicities, i.e., the difference between $P_\mathrm{orb}$ and $P_{1}$ in Fig.~\ref{fig:laprompeaks}. However, the fact that the spot-crossing features in the eclipses of KOI-1452 rotate with $P_\mathrm{rot}=\,1.1357\pm0.003$ days or $P_\mathrm{rot}=\,1.1687\pm0.003$ days, assuming a prograde or retrograde orbit respectively, requires some additional discussion. There is no feature in the L-S periodogram close to period with value $P_\mathrm{rot}=\,1.1357\pm0.003$ days. Yet there is a peak in the L-S periodogram at a period $P_\mathrm{rot}=\,1.1687\pm0.003$ days. Since the value of this peak is very close to the suggested rotational period of the occulted spots $P_\mathrm{rot}=\,1.1687\pm0.003$ days, it is tempting to identify the two. However, this requires us to assume that the primary star rotates in the opposite direction of the orbital motion of its companion, which is in contradiction to the hypothesis of tidal locking. | 16 | 7 | 1607.08080 |
1607 | 1607.08339_arXiv.txt | We describe a new phenomenon of `bar damping' that may have played an important role in shaping the Milky Way bar and bulge as well as its spiral structure. We use a collisionless $N$-body simulation of a Milky Way-like galaxy initially composed of a dark matter halo and an exponential disk with Toomre parameter slightly above unity. In this configuration, dominated by the disk in the center, a bar forms relatively quickly, after 1 Gyr of evolution. This is immediately followed by the formation of two manifold-driven spiral arms and the outflow of stars that modifies the potential in the vicinity of the bar, apparently shifting the position of the $L_1/L_2$ Lagrange points. This modification leads to the shortening of the bar and the creation of a next generation of manifold-driven spiral arms at a smaller radius. The process repeats itself a few times over the next 0.5 Gyr resulting in further substantial weakening and shortening of the bar. The time when the damping comes to an end coincides with the first buckling episode in the bar which rebuilds the orbital structure so that no more new spiral arms are formed. The morphology of the bar and the spiral structure at this time show remarkable similarity to the present properties of the Milky Way. Later on, the bar starts to grow rather steadily again, weakened only by subsequent buckling episodes occurring at more distant parts of the disk. | Bars are quite common morphological features among late-type galaxies. Depending on the exact definition about half of them can be considered as barred \citep[see e.g.][and references therein]{Buta2015}. Bars formed in isolation are believed to originate from instabilities in axisymmetric disks \citep[for a review see][]{Athanassoula2013} and their evolution is controlled by a number of parameters, including the presence and properties of a live halo \citep{Athanassoula2002, Debattista2006, Sellwood2016} and the initial velocity dispersion of the disk \citep{Athanassoula2003}. In simulations, bars tend to grow in time in terms of length and strength while decreasing their pattern speed. This rather steady growth can be slowed down by episodes of buckling instability that lead to the vertical thickening of the bar and the formation of boxy/peanut shapes or pseudobulges \citep{Combes1990, Raha1991, Martinez2006}. Bars can also form as a result of tidal interactions with other galaxies or environment \citep{Miwa1998, Lang2014, Lokas2014, Lokas2016, Martinez2016}. The Milky Way galaxy seems to be a typical barred spiral \citep[for a recent review of its properties see][]{Bland2016}. The bar in the Milky Way is composed of a shorter (the length of 3 kpc) and a longer (5 kpc) component. The shorter one is thicker and known to possess a boxy/peanut shape since the observations in the infrared by the COBE satellite \citep{Weiland1994}. The density distributions of the red clump stars within 1 kpc are approximately exponential with axis ratios 10:6.3:2.6 \citep{Wegg2013}. The bar pattern speed is not accurately known and an average over different measurements gives values around $43 \pm 9$ km s$^{-1}$ kpc$^{-1}$. From our position near the Sun we see the bar at an angle of about 27 deg. Recent kinematic measurements of the bar from the APOGEE survey reveal that it is characterized by cylindrical rotation at the level of 170 km s$^{-1}$ at 35 deg of galactic longitude and the central velocity dispersion of 120 km s$^{-1}$ \citep{Ness2016}. The spiral structure of the Milky Way is much less constrained, mainly due to our unfavorable observational position close to the disk plane. Although a number of spiral arms were identified and named, their detailed properties remain poorly known. In a recent study \citet{Hou2014} combined data for thousands of different spiral tracers and found them to be well fitted by models of both three- and four-arm logarithmic spirals. To further complicate the picture, they also found that polynomial-logarithmic spirals (with variable pitch angles) are able to better match the observed tangential directions. The origin of Milky Way spiral arms is equally vague and many scenarios have been proposed, starting from the classical density wave theory of \citet{Lin1964} to spiral features seeded by density inhomogeneities \citep{D'Onghia2013} and tidal effects of nearby dwarf galaxies \citep{Purcell2011}. It has also been proposed that the spiral structure formation is driven by the bar through manifolds \citep{Athanassoula2009b}. In this Letter we report on a discovery of a new phenomenon in the simulation of a bar in a Milky Way-like galaxy, which we refer to as `bar damping'. In the configuration we consider here it occurs early on, soon after the formation of the bar and seems to be tightly related to the manifold-driven spiral arms. | We interpret the observed behavior of the bar as due to the formation and evolution of manifold-driven spirals. Manifolds were proposed as a possible origin of spiral arms and rings in barred galaxies \citep{Romero2006, Romero2007, Athanassoula2009a, Athanassoula2009b, Athanassoula2010}. According to this theory, spiral arms are formed by stars on orbits confined to manifolds associated with the periodic orbits around the saddle points of the potential in the frame of reference corotating with the bar \citep{Binney2008}. In such a frame there are five equilibrium (Lagrange) points at which the derivative of the effective potential vanishes. The most important for us here are the $L_1$ and $L_2$ points located along the bar major axis, near its ends. These points are saddle points and unstable causing the stars in their vicinity to escape from the neighborhood. Manifolds can be interpreted as tunnels along which this escape can take place. The properties of manifolds so far were mostly studied by tracing orbits in analytically set potentials, however, \citet{Athanassoula2012} demonstrated that they actually occur in full, self-consistent $N$-body simulations and lead to the formation of spiral arms. Here we have witnessed a similar phenomenon. However, in our case the formation of manifold-driven spirals is not a single event but seems to repeat itself a few times. Around $t=1.135$ Gyr (see the upper left panel of Figure~\ref{surden}) the manifold-driven spirals form for the first time. From the comparison between the circular frequency of the galaxy and the pattern speed of the bar, $\Omega = \Omega_{\rm p}$, we can estimate the corotation radius where the $L_1/L_2$ are located. For this time we obtain the value of $R_{\rm CR}=9$ kpc which agrees with the distance from the center of the galaxy, where the spiral arms emanate from the bar. This confirms our interpretation of the spirals as manifold-driven since according to the theory in cases of barred galaxies with no rings and a relatively weak spiral structure the $L_1/L_2$ points are located at the end of the bar, where the arm joins the bar \citep{Athanassoula2010}. Later on the picture becomes more complicated. In addition to the outflow of stars along the outer manifolds (spiral arms) there is a flow along the inner manifolds located near the bar \citep[see fig. 4 in][]{Athanassoula2010} that rebuilds its structure so that for a short period of time it looks like a double bar (see the middle right panel of Figure~\ref{surden}). The pattern speed of such a bar cannot be reliably measured so the position of the $L_1/L_2$ points is difficult to determine. However, these flows seem to modify the potential near the bar sufficiently in order to move the $L_1/L_2$ points towards the center of the galaxy and thus decrease the length and strength of the bar. This process could involve the stabilization of the $L_1/L_2$ points by the additional mass present in the spiral arms and the creation of a new set of unstable $L_1/L_2$ points nearby as described in section 5 of \citet{Athanassoula2009a}. After the bar rebuilds itself it `emits' a next generation of spiral arms along outer manifolds now located closer to the center and the bar is restructured. This shifts the position of the $L_1/L_2$ points again and the whole process is repeated. One cycle of such evolution is shown by the four panels corresponding to times $t=1.32$-1.425 Gyr in Figure~\ref{surden}. As illustrated in the lower panel of Figure~\ref{a2modestime} and the first four panels of Figure~\ref{propdetail}, between 1 and 1.7 Gyr there are at least five such cycles distinguishable. At about $t=1.6$ Gyr the bar starts to buckle for the first time so its orbital structure is substantially rebuilt and no well-defined bar-driven spiral arms form later on. Interestingly, right after the damping period is finished the bar is quite short, with the length of the order of 5 kpc, as can be estimated from the drop of the $A_2(R)$ profile to half its maximum value (see the orange line for $t=1.8$ Gyr in Figure~\ref{a2profiles}). In addition, it is surrounded by a family of rather irregular spiral arms of different length originating from the multiple generations of manifold-driven double spirals (see the lower right panel of Figure~\ref{surden}). This morphology is qualitatively similar to the present spiral structure of the Milky Way as far as we know it \citep{Vallee2008, Hou2014}. We note that this transition from two- to multiple-arm structure cannot be explained by arguments based on the swing amplification theory and the disk stability criteria \citep{Athanassoula1987, DOnghia2015}. According to this theory the number of arms grows with decreasing disk mass fraction within $2.2 R_{\rm D}$ while in our case this fraction remains approximately constant (within 2\%) during the whole damping period and close to the initial value of 0.6. We conclude that the mechanism of bar damping we described here may have contributed to the formation of the present-day structure of the Milky Way. It could influence the shaping of the bar and spiral structure, as these seem to be intimately related in this process, but may also have an effect on the formation of Milky Way (pseudo)bulge if it formed via buckling instability as the first episode of this phenomenon takes place immediately after damping. Although in our simulation the damping occurs early in the evolution of the bar, it is quite possible that with different initial conditions and/or addition of gas physics it could be a much more recent phenomenon. | 16 | 7 | 1607.08339 |
1607 | 1607.08613_arXiv.txt | The MaNGA Survey (Mapping Nearby Galaxies at Apache Point Observatory) is one of three core programs in the Sloan Digital Sky Survey IV. It is obtaining integral field spectroscopy (IFS) for 10K nearby galaxies at a spectral resolution of $R\sim2000$ from $3622-10,354{\rm \AA}$. The design of the survey is driven by a set of science requirements on the precision of estimates of the following properties: star formation rate surface density, gas metallicity, stellar population age, metallicity, and abundance ratio, and their gradients; stellar and gas kinematics; and enclosed gravitational mass as a function of radius. We describe how these science requirements set the depth of the observations and dictate sample selection. The majority of targeted galaxies are selected to ensure uniform spatial coverage in units of effective radius ($R_e$) while maximizing spatial resolution. About 2/3 of the sample is covered out to $1.5R_e$ (Primary sample), and 1/3 of the sample is covered to $2.5R_e$ (Secondary sample). We describe the survey execution with details that would be useful in the design of similar future surveys. We also present statistics on the achieved data quality, specifically, the point spread function, sampling uniformity, spectral resolution, sky subtraction, and flux calibration. For our Primary sample, the median r-band signal-to-noise ratio is $\sim73$ per $1.4{\rm \AA}$ pixel for spectra stacked between 1--1.5 R$_{e}$. Measurements of various galaxy properties from the first year data show that we are meeting or exceeding the defined requirements for the majority of our science goals. | Large spectroscopic galaxy surveys, such as the Sloan Digital Sky Survey \citep{York00}, and the Two-degree Field Galaxy Redshift Survey \citep{Colless01}, have revolutionized the way we study galaxy evolution. The huge statistical power brought in by targeting a large number of galaxies using the same instrument with excellent calibration enabled huge progress. Not only have these efforts quantified accurately with great precision those trends and scaling relations that were previously known, such as the color-bimodality \citep{Strateva01, BaldryGB04}, the color-density relation \citep{Hogg03, BlantonEH05}, the mass-metallicity relation for gas \citep{TremontiHK04} and stars \citep{Thomas10,Johansson12}, and the Fundamental Plane \citep{Bernardi03}, they have also discovered many new relations and trends, such as the dependence of star formation history on stellar mass \citep{KauffmannHW03}, the star formation rate vs. stellar mass relation \citep{Brinchmann04, Salim07, Wuyts11}, the strong mass dependence of the radio-loud AGN fraction \citep{Best05}, large scale galactic conformity \citep{Kauffmann13}, and many others. They also connected large scale structure studies and galaxy evolution studies thanks to environmental measurements enabled by dense and uniform sampling of complete galaxy samples (see \citealt{BlantonM09} and references therein). However, these massive surveys lacked spatial coverage in individual galaxies. The single 3\arcsec\ fibers used by SDSS, for instance, cannot cover most of the light in nearby galaxies. For example, comparing the flux incident on the SDSS 3\arcsec\ fibers with the total flux of all main sample galaxies in SDSS, 80\% of galaxies have less than 36\% of their light covered by the fiber. The spectra provide a lot of information, about both stellar and gaseous components, but they only sample the center of the galaxies and can give a strongly biased picture. Nearly all studies based on SDSS have to take this aperture effect into account in their analysis. Many studies combining spectroscopic information with photometry also need to make corrections, extrapolations, or use simplified assumptions. For example, to obtain the total star formation rate in a galaxy, one either has to apply large aperture corrections to the spectroscopically-derived star formation rate based on the central region \citep{Brinchmann04}, or turn to broadband photometry which suffers more from dust extinction and degeneracies in stellar population modeling \citep{Salim05, Salim07}. Furthermore, a full kinematic description is impossible with single-fiber observations. Past long-slit surveys are also inefficient at obtaining the spatial information as one only probes a narrow elongated region and the signal-to-noise is poor in galaxy outskirts. Integral field spectroscopy (IFS) solves these problems. Several IFS surveys have made great progress in recent years (see \citealt{Cappellari16} for a review). SAURON \citep{Bacon01} and ATLAS${\rm 3D}$ \citep{Cappellari11} surveyed 260 early-type galaxies in the nearby universe using a lenslet array intergral field instrument, SAURON, on the William Herschel Telescope on La Palma. They had a relatively narrow wavelength coverage ($4800-5380{\rm \AA}$) and focused exclusively on early-type galaxies. % The DiskMass survey \citep{Bershady10} used two fiber-fed Integral Field Unit (IFU), SparsePak on WIYN and PPak on the Calar Alto 3.5m Telescope. It targeted 146 nearly face-on disk galaxies to study stellar and gas kinematics. For the purpose of kinematic measurements, this survey utilized high spectral resolution in three narrow wavelength windows around 515, 660, and 860 nm. % The VENGA \citep{Blanc13} survey used a fiber-fed integral field spectrograph, VIRUS-P on the 2.7-m Telescope at McDonald Observatory, and targeted 30 nearby spiral galaxies. Recently, the CALIFA survey used the PPak instrument and targeted 600 nearby galaxies selected to sample a wide variety of stellar mass and star formation rate. With the improved sample size and wide wavelength coverage, CALIFA has produced numerous results, such as the universal metallicity gradient among star-forming galaxies \citep{Sanchez14}, the nature of LINER-like galaxies \citep{Kehrig12, Singh13, Gomes16}, the spatially-resolved growth history \citep{Perez13, Sanchez-Blazquez14}, the spatially-resolved stellar mass-metallicity relation \citep{GonzalezDelgado14}, and the resolved star formation main sequence \citep{Cano-Diaz16, GonzalezDelgado16}. However, if one were to do an SDSS-like study of galaxies by binning galaxies by stellar mass, environment, and morphology, one quickly loses sample size for significant statistics \citep[e.g.][]{GonzalezDelgado16}. The main limitation for the sample size is that all these surveys are targeting galaxies one by one and are inefficient at building up a large statistical sample. To address this issue, two large IFS surveys of the general galaxy population targeting thousands of galaxies are ongoing right now. Both utilize multiple fiber bundles to target multiple galaxies at the same time, enabling much more efficient observing. One of them is the SAMI Galaxy survey \citep{Bryant15, Allen15} using the fiber-fed SAMI instrument \citep{Croom12} on the 3.5m Anglo-Australian Telescope at Siding Spring Observatory. SAMI will eventually target 3400 galaxies and has already produced results on many topics, including the kinematic morphology-density relation \citep{Fogarty14}, outflows and extraplanar gas \citep{Ho14,Ho16}, dynamical scaling relations \citep{Cortese14}, dynamical M/L ratio of disk galaxies \citep{Cecil16}, and aperture corrections for star formation rates \citep{Richards16}. The other large IFS survey is the SDSS-IV/MaNGA galaxy survey operating at the 2.5m Sloan Foundation Telescope. Given the large 3$^\circ$ field of view of the SDSS telescope and sizeable detector real estate, MaNGA uses multiple fiber bundles to target 17 galaxies (and 12 standard stars) at the same time. This allows us to build a 10K galaxy IFS sample with much wider and continuous wavelength coverage than other surveys, enabling powerful statistical studies of the spatially-resolved properties of nearby galaxies. This paper complements MaNGA's other descriptive publications by providing a complete picture of the survey's design and execution, and an evaluation of the resulting data quality. In Section~\ref{sec:requirements}, we describe the science requirements of our survey, and how they flow down to specific decisions on the sample design and observing strategy. We summarize the hardware in Section~\ref{sec:hardware} and the sample design in Section~\ref{sec:sample}. In Section~\ref{sec:execution}, we describe the execution of the survey, including the observing strategy, setting of the completeness thresholds, choice of the fields, plate design, observing procedure, and the optimization of the instrument focus. In Section~\ref{sec:progress} we dscribe our survey progress and projection. In Section~\ref{sec:quality}, we provide an evaluation of the initial data quality: PSF, sampling uniformity, spectral resolution, sky subtraction accuracy, and flux calibration accuracy. In the end (Section~\ref{sec:verification}), we present a series of tests checking whether we are meeting the science requirements. We summarize in Section~\ref{sec:summary}. | \label{sec:summary} MaNGA is an integral field spectroscopic survey of 10K nearby galaxies with wide wavelength coverage at medium resolution with uniform spatial coverage in units of \Reff. Up to the time of writing, we have already obtained observations for more than 2550 galaxies and are on track to finish $\sim10{\rm K}$ by summer of 2020. In this paper, we have detailed the survey science requirements, both in terms of random and systematic errors, and how the high-level science requirements flow down in an interconnected way to the low level requirements on the hardware, sample selection, observations, and analysis. In this context we have described in detail how the sample selection is carried forward to generating a survey footprint on the sky, how this footprint is parsed into tiles, how these tiles are targeted with plates, and how these plates are designed, fabricated and scheduled for observation. The observing procedures are likewise detailed at a level necessary for a complete and reliable reconstruction of the survey execution. Finally, as proof of practice, we have given a complete demonstration of the data quality in both basic data products and high level derived science products across the full first year of data. The basic data quality of the survey is excellent. We have reached the S/N target while staying on track to finish observing 10K galaxies by 2020. We obtain a spatial resolution about 2.5\arcsec FWHM with a carefully characterized profile with uniform and near-critical sampling from multiple dithered observations. The sky subtraction is nearly Poisson even at near-infrared wavelengths. Both the absolute and relative flux calibrations are better than 5\%. The spectral resolution is a function of wavelength and is characterized for each fiber in each exposure. Exposure-to-exposure variations should be taken into account if the science case warrants it. % The high level derived science products are also of high quality. We have met the majority of the science requirements set forth, such as the precision on the star formation rate surface density, the gas metallicity gradient, the stellar population age and metallicity gradient. On the several kinematics requirements, such as the specific angular momentum, the enclosed mass, and the dark matter fraction, the systematic errors due to simplified modeling assumptions dominate the precision of the measurements. The formal errors appear to meet the science requirements, but whether the scienctific goals on kinematics could be reached awaits further analysis facilitated by detailed simulations. The first year data will be released in SDSS Data Release 13\ in summer 2016. | 16 | 7 | 1607.08613 |
1607 | 1607.05731_arXiv.txt | The optical spectra of Seyfert galaxies are often dominated by emission lines excited by both star formation and AGN activity. Standard calibrations (such as for the star formation rate) are not applicable to such composite (mixed) spectra. In this paper, we describe how integral field data can be used to spectrally and spatially separate emission associated with star formation from emission associated with accretion onto an active galactic nucleus (AGN). We demonstrate our method using integral field data for two AGN host galaxies (NGC~5728 and NGC~7679) from the Siding Spring Southern Seyfert Spectroscopic Snapshot Survey (S7). The spectra of NGC~5728 and NGC~7679 form clear sequences of AGN fraction on standard emission line ratio diagnostic diagrams. We show that the emission line luminosities of the majority ($>$~85 per cent) of spectra along each AGN fraction sequence can be reproduced by linear superpositions of the emission line luminosities of one AGN dominated spectrum and one star formation dominated spectrum. We separate the \Ha, \Hb, \NII$\lambda$6583, \mbox{\SII$\lambda \lambda$6716, 6731}, \OIII$\lambda$5007 and \mbox{\OII$\lambda \lambda$3726, 3729} luminosities of every spaxel into contributions from star formation and AGN activity. The decomposed emission line images are used to derive the star formation rates and AGN bolometric luminosities for NGC~5728 and NGC~7679. Our calculated values are mostly consistent with independent estimates from data at other wavelengths. The recovered star forming and AGN components also have distinct spatial distributions which trace structures seen in high resolution imaging of the galaxies, providing independent confirmation that our decomposition has been successful. | The impact of AGN activity on the evolution of the host galaxy is a topic of significant debate. In the local universe, AGN reside in galaxies at a range of evolutionary stages, from spiral galaxies with strong circumnuclear star formation \citep[e.g.][]{GonzalezDelgado01, Joguet01, StorchiBergmann01, Raimann03, CidFernandes04, Gu06, Davies07}, to `green valley' galaxies (whose optical colours suggest that they may be transitioning from the star-forming blue cloud to the quiescent red sequence) \citep[e.g.][]{Ka03,Baldry04,Schawinski10, Leslie15} and quiescent elliptical galaxies \citep[e.g.][]{Olsen13}. The evolved nature of many AGN host galaxies may be a direct result of AGN feedback quenching star formation \citep[e.g.][]{DiMatteo05,Nandra07, Schawinski09, Schawinski10, Leslie15} or may simply be a consequence of the increase in both the fraction of passive galaxies and the fraction of galaxies hosting AGN with increasing stellar mass \citep[e.g.][]{Ka03,Baldry04}. Global colours and integrated star formation rates (SFRs) provide limited insight into the physical processes impacting the gas reservoirs within AGN host galaxies, making it difficult to distinguish between these scenarios. The combination of imaging and spectroscopy can provide a more direct view of the impact of AGN activity on the surrounding interstellar medium (ISM). For example, \citet{Cresci15} used near infrared integral field spectroscopy (IFS) to identify both positive and negative AGN feedback in a radio quiet quasar at \mbox{z = 1.59}. They find that star formation is being suppressed due to the entrainment of molecular gas within an AGN driven outflow, but is also being triggered at the edges of the outflow due to the pressure imposed on the surrounding ISM. Several other spatially resolved and/or multi-wavelength studies have found direct evidence for star formation being suppressed due to strong outflows \citep[e.g.][]{CanoDiaz12} or triggered due to the compression of molecular clouds by radio jets \citep[e.g.][]{Croft06, Elbaz09, Rauch13, Salome15}. These studies highlight the power of spatial information for providing insights into the connection between star formation and AGN activity in galaxies. Modern IFS surveys (such as S7 \citep{Dopita15}, the Sydney AAO Multi-Object Integral Field Spectrograph (SAMI) Survey \citep{Croom12, Bryant15}, the Calar Alto Legacy Integral Field Area (CALIFA) Survey \citep{Sanchez12} and the Mapping Nearby Galaxies at Apache Point Observatory (MaNGA) Survey \citep{Bundy15}) are providing optical spectra for many tens of spatially resolved regions across hundreds to thousands of galaxies. These data can potentially be used to map the SFR and the strength of the AGN ionizing radiation field across large samples of AGN host galaxies. However, extracting this information from the spectra of Seyfert galaxies (which often have significant contributions from both star formation and AGN activity) is non-trivial. The \Ha\ luminosity of an \HII\ region is directly proportional to the SFR \citep[e.g.][]{Kennicutt98}. However, \Ha\ can be collisionally excited in the presence of an AGN ionizing radiation field, and therefore the \Ha\ luminosity is not a valid diagnostic of the SFR in AGN narrow line regions (NLRs). Similarly, the \OIII\ $\lambda$5007 luminosity is directly proportional to the strength of the local AGN ionizing radiation field within AGN dominated regions \citep[e.g.][]{Heckman04, Lamastra09}, but is not a valid diagnostic of the radiation field strength in regions with active star formation. It is therefore necessary to separate the composite spectra of Seyfert galaxies into contributions from star formation and AGN activity before calculating SFRs or AGN luminosities. Many techniques exist for separating emission associated with star formation and AGN activity. In the mid infrared, these two ionization mechanisms have very different spectral signatures and can be separated using template fitting \citep[e.g.][]{Nardini08, AlonsoHerrero12, Kirkpatrick14, HernanCaballero15}. However, mid infrared observations typically have low spatial resolution and therefore cannot provide detailed maps of the emission associated with star formation and AGN activity across galaxies. On the other hand, optical IFS surveys represent a very promising avenue for studying the connection between star formation and AGN activity on scales of hundreds of parcsecs to kiloparsecs. The optical spectra of AGN host galaxies are dominated by strong emission lines which can be excited by both emission from massive stars and the AGN ionizing radiation field. The ratios of forbidden to recombination lines are dependent on the relative contributions of star formation and AGN activity and can therefore be used to estimate the fraction of the line emission excited by AGN activity (`AGN fraction'). The \NIIHa\ vs. \OIIIHb\ diagnostic diagram is a valuable tool for separating emission associated with star formation and AGN activity \citep{Baldwin81, Veilleux87, Ke01a}. Spectra dominated by star formation fall along the star-forming sequence which traces variations in the ionized gas abundance \citep{Dopita86, Dopita00}. Spectra with contributions from more energetic ionization mechanisms lie along the AGN branch of the diagnostic diagram which spans from the enriched end of the star-forming sequence towards larger \NIIHa\ and \OIIIHb\ ratios. The presence of a harder ionizing radiation field increases the collisional excitation rate in the nebula and therefore increases the strengths of forbidden lines such as \NII\ and \OIII\ (produced by radiative transitions from collisionally excited metastable states) relative to the \Ha\ and \Hb\ recombination lines. The greater the AGN fraction, the greater the enhancement in the \NIIHa\ and \OIIIHb\ ratios and the further along the AGN branch a galaxy will lie. The Sloan Digital Sky Survey (SDSS) provided single fibre optical spectra for hundreds of thousands of galaxies in the local universe, leading to several pioneering works in the separation of star formation and AGN activity. \citet{Heckman04} corrected the \OIII\ luminosities of observed composite spectra for the contribution of star formation using average AGN fractions calculated for synthetic composite spectra (generated by summing observed star formation and AGN dominated spectra) in bins of \OIII\ luminosity. \citet{Ke06} established the `star-forming distance' ($d_{SF}$, the distance of a galaxy spectrum from the star-forming sequence of the \NIIHa\ vs. \OIIIHb\ diagnostic diagram) as a metric for the relative contribution of star formation to the line emission. Following this, \citet{Kauffmann09} used the positions of galaxies along the AGN branch of the diagnostic diagram to estimate AGN fractions and correct the \OIII\ luminosities for the contribution of star formation. These pioneering techniques facilitated the first large statistical studies of black hole accretion rates and Eddington ratios as a function of host galaxy properties. With the advent of integral field spectropscopy, it is now possible to use similar techniques to separate star formation and AGN activity in individual spatially resolved regions within Seyfert host galaxies. \citet{Davies14b,Davies14a} showed that spectra extracted from individual spectral pixels (spaxels) of AGN host galaxies sometimes fall along tight sequences spanning the full range of line ratios observed along the SDSS global mixing sequence. The line ratios vary smoothly with galactocentric distance, from AGN-like line ratios in the galaxy nuclei to \HII-like line ratios at larger radii, providing strong evidence that the line ratio variations are primarily driven by variations in the AGN fraction \citep[see also][]{Scharwaechter11, Dopita14, Belfiore15, Davies16RP}. We note that the diagnostic line ratios are sensitive to the physical conditions of the ionized gas \citep[see e.g.][]{Ke03, Groves04, Dopita13} and therefore radial variations in the local ionization parameter and/or gas pressure would also drive radial variations in the line ratios \citep[e.g.][]{Greene11,Liu13}. However, the presence of spectra in the star forming, composite and AGN dominated regions of the \NIIHa\ vs. \OIIIHb\ diagnostic diagram within a single galaxy is inconsistent with ionization by a single mechanism and can only be explained by mixing. The observation of spectra at a range of AGN fractions makes it possible to directly constrain the shape of the starburst-AGN mixing curve for individual galaxies, without any \textit{a priori} knowledge of the ISM conditions. In this paper, we present a new method for separating emission associated with star formation and AGN activity in individual emission lines within individual spaxels of IFS datacubes. We describe our data and introduce our galaxy sample in Section \ref{sec:obs}. We discuss the motivation for our method and outline the steps in Section \ref{sec:method}. In Section \ref{sec:results} we present the results of the decomposition for each galaxy, and compare our results to independent tracers of star formation and AGN activity at other wavelengths. We discuss the advantages and limitations of our decomposition method and compare it to existing methods in Section \ref{sec:discussion}. We summarise our main conclusions in Section \ref{sec:conc}. Throughout this paper we adopt cosmological parameters $H_{0} = 70.5 \, {\rm kms}^{-1}{\rm Mpc}^{-1}$, ${\Omega}_{\Lambda} = 0.73$, and $\Omega_{M}=0.27$ based on the 5-year \emph{Wilkinson Microwave Anisotropy Probe} (WMAP) results by \citet{Hinshaw09} and consistent with flat $\Lambda$-dominated cold dark matter ($\Lambda$ CDM) cosmology. | \label{sec:conc} In this paper we have explained how emission associated with star formation and AGN activity can be separated both spatially and spectrally in integral field observations of AGN host galaxies. Our method is applicable to any galaxy with a clear mixing sequence on the \NIIHa\ vs. \OIIIHb\ diagnostic diagram, and requires only the \Ha, \Hb, \NII, \SII\ and \OIII\ fluxes for each spectrum along the mixing sequence. We demonstrated our method using integral field data for two AGN host galaxies from the S7 survey, NGC~5728 and NGC~7679. We showed that: \begin{itemize} \item Many of the spectra extracted from the datacubes of NGC~5728 and NGC~7679 lie along clear mixing sequences between star formation and AGN activity on the \NIIHa\ vs. \OIIIHb\ diagnostic diagram, and \item The emission line luminosities of $>$~85 per cent of the spectra along each mixing sequence can reproduced by linear superpositions of the emission line luminosities of one AGN dominated basis spectrum and one star formation dominated basis spectrum. \end{itemize} We separated the luminosity of each strong emission line in each spaxel into contributions from star formation and AGN activity, and compared our decomposed emission line maps to independent tracers of star formation and AGN activity at other wavelengths. The star formation component of NGC~5728 traces the star forming ring seen in the HST F336W image of this galaxy, and the SFR calculated from the star formation component of the \Ha\ emission is consistent with the SFR calculated from the 8$\mu$m emission over the WiFeS FOV. The AGN component is asymmetric and its alignment matches the position angle of the ionization cone identified in HST narrow band imaging of NGC~5728. The AGN bolomeric luminosity calculated from the AGN component of the \OIII\ emission also matches the bolometric luminosity calculated from the \mbox{2-10 keV} luminosity. The star formation component of NGC~7679 traces \HII\ regions in a disk. The SFR calculated from the star formation component of the \Ha\ emission is a factor of 1.9 (0.27 dex) higher than the SFR calculated from the 8$\mu$m emission. This discrepancy cannot be due to under-correcting the \Ha\ emission for the contribution of the AGN, and may be driven by differences in the timescales traced by the different SFR indicators, inaccuracies in accounting for dust extinction, and/or intrinsic scatter in the SFR calibrations. The AGN component reveals a clear ionization cone, and the AGN bolometric luminosity calculated from the AGN component of the \OIII\ emission matches the bolometric luminosity calculated from X-ray observations of NGC~7679. Our decomposed emission line maps are consistent with independent tracers of star formation and AGN activity. We therefore conclude that our decomposition method allows us to robustly separate emission associated with star formation and AGN activity in NGC~5728 and NGC~7679. The ability to separate emission associated with star formation and AGN activity will provide unique insights into the impact of AGN feedback on star formation in AGN host galaxies. The \Ha\ luminosities of the star formation component can be directly converted to SFRs, facilitating analysis of SFR gradients in AGN host galaxies. The spatial variations in the SFR surface densities can be compared to the spatial variations in the strength of the AGN ionizing radiation field to search for evidence of positive and/or negative AGN feedback (modulo the limitations of beam smearing and dust attenuation, as discussed in Section \ref{subsec:dust_and_beam_smearing}). Our results indicate that the combination of integral field spectroscopy and emission line ratio diagnostic diagrams is a powerful avenue by which to gain insights into the connection between star formation and AGN activity in galaxies. | 16 | 7 | 1607.05731 |
1607 | 1607.07444_arXiv.txt | {} {This study aims to probe the thermodynamic properties of the hot intragroup medium (IGM) plasma in the core regions of the NGC 4636 galaxy group by detailed measurements of several emission lines and their relative intensities. } {We analyzed deep XMM-\emph{Newton} Reflection Grating Spectrometer (RGS) data in five adjacent spectral regions in the central parts of the NGC 4636 galaxy group. We examined the suppression of the \ion{Fe}{xvii} resonance line (15.01 Å) as compared to the forbidden lines of the same ion (17.05 Å and 17.10 Å). The presence and radial dependence of the cooling flow was investigated through spectral modeling. Parallel analysis with deep Chandra Advances CCD Imaging Spectrometer (ACIS) data was conducted to gain additional information about the thermodynamical properties of the IGM.} {The plasma at the group center to the north shows efficient \ion{Fe}{xvii} ion resonant scattering, yielding $(I_{\lambda17.05}+I_{\lambda17.10})/I_{\lambda15.01}$ line ratios up to 2.9$\pm$0.4, corresponding $\text{to about}{}$ twice the predicted line ratio. In contrast, no resonant scattering was detected at the south side. The regions featuring resonant scattering coincide with those embodying large amounts of cool ($kT\lesssim0.4\mathrm{keV}$) gas phases, and the spectral imprints of cooling gas with a total mass deposition rate of $\sim0.8$ M$_{\sun}$ yr$^{-1}$ within the examined region of $2.4\arcmin\times 5.0\arcmin$.} {We interpret the results as possible evidence of asymmetric turbulence distribution in the NGC 4636 IGM: Turbulence dominates the gas dynamics to the south, while collective gas motions characterize the dynamics to the north. X-ray images show imprints of energetic AGN at both sides, yet we find evidence of turbulence heating at the south and gas cooling at the north of the core. We infer that the observed asymmetry may be the result of the specific observation angle to the source, or arise from the turbulence driven by core sloshing at south side.} | {\label{intro}} The intergalactic space within galaxy clusters and groups is filled with diffuse highly ionized plasma with virial temperatures ranging from sub-keV to $\sim 10$ keVs. X-ray spectra of such intergalactic plasmas consist of a thermal bremsstrahlung continuum and line emission of several high-Z elements, whose intensities and line widths carry information about the thermodynamical state and composition of the emitting gas. In this paper we focus on detailed spectroscopic studies of properties of the X-ray halo surrounding a nearby giant elliptical galaxy, NGC 4636, the dominant galaxy of a group located at the outer parts of the Virgo galaxy cluster. The NGC 4636 galaxy group is approximately 15 Mpc distant and its apparent X-ray luminosity is one of the brightest of all group luminosities. Dynamically the group consists of a concentrated dark matter halo (\citealt{schuberth06}) filled with baryonic gas, with a temperature of $kT\sim0.5$ keV in the cool dense core region and $kT\sim0.8$ keV in the outer parts (e.g., \citealt{loewenstein02,sullivan05,finoguenov06,johnson09}). It is plausible that the core region of the intragroup medium (IGM) halo contains multiphase plasma, as discussed in \cite{sullivan05}, for example. The NGC 4636 X-ray halo has a complex morphology. The most prominent structures are two arms extending toward opposite sides of the galaxy, both of which connect to hot ellipsoidal bubble-shaped X-ray cavities several kpc from the central galaxy (see Fig. \ref{figure:morphology}). As scrutinized in \cite{baldi09}, for instance, these structures are related to a large, ancient AGN outburst into the surrounding IGM gas. The outburst took place $\sim2-3\cdot10^6$ years ago (\citealt{baldi09}), releasing a total energy of $\sim6\cdot10^{56}$ ergs into the IGM (\citealt{jones02}) through relativistic jets. Observations of the radio jets directed toward the X-ray cavities were published in the multiwavelength study by \cite{giacintucci11}, indicating that the interaction between the central black hole (BH) and IGM is currently weaker. Such feedbacks can heat the IGM through turbulence (e.g. \citealt{zhuravleva14}), sound wave (\citealt{fabian12}), or weak shock heating (\citealt{randall15}). In addition to the AGN feedbacks, often quoted mechanisms capable of driving turbulence and heating the IGM include galactic merging and core sloshing (\citealt{aschasibar06}). Theory and simulations suggest that IGM plasma heating occurs mainly through the dissipation of turbulent kinetic energy into thermal plasma energy (see, e.g., \citealt{peterson06,hillel14} and references therein). Turbulence can also mix multiphase gas, reducing the temperature phase distribution, hence slowing down the processes leading to cooling flows, as discussed in \cite{banerjee14}, for example. It is currently not fully understood to what extent AGN heating counterbalances radiative cooling in individual cool-core systems such as the NGC 4636. Nevertheless, if the cool gas phases are formed by cooling from the hot phase and the central supermassive black hole remains in a quiescent state for a sufficiently long period of time, large cooling flows should develop that might trigger a new cycle of high-energy AGN outbursts (see, e.g., \citealt{jones02,werner13}). As the core region of the NGC 4636 group has been exposed to a large AGN outburst in the relatively recent past, it is a suitable source for studying the consequences for present gas thermodynamics in different regions of the X-ray halo. The radiative cooling time in the core is considerably shorter than the Hubble time (e.g., \citealt{mathews03,chen07}), and therefore cooling flows toward the center of gravity may be present as a consequence of decreased thermal pressure of radiatively cooling IGM gas. Predictions of the traditional cooling flow model (\citealt{fabian94}) suggest total mass deposition rates of approximately $\dot{M}\sim1-2$ M$_{\sun}$ yr$^{-1}$ for the NGC 4636 group (e.g., \citealt{bertin95,chen07}). The NGC 4636 group's core IGM temperature gives rise to efficient emission of the \ion{Fe}{xvii} ion line complex (\citealt{doron02}). The XMM-\emph{Newton} RGS instrument with the designed wavelength band of $\sim 6-38$ Å ($\sim 0.3-2.5$ keV) is currently the most powerful observational tool to resolve and study these lines. For typical IGM conditions the \ion{Fe}{xvii} forbidden lines are optically thin, but the resonance line can become optically thick in dense regions with low turbulent velocities. Consequently, detailed measurements of the \ion{Fe}{xvii} emission can give important information on the IGM properties. Resonant scattering in galaxy clusters was first discovered by \citet{gilfanov87}, while observations of \ion{Fe}{xvii} resonant scattering in the NGC 4636 group core region has previously been published by \citet{xu02} and \citet{werner09}. We investigated the \ion{Fe}{xvii} resonant scattering by comparing the intensities of the unresolvable blend of the forbidden line doublet at 17.05 and 17.10 Å to that of the resonance line (15.01 Å). Observations of the $(I_{\lambda17.05}+I_{\lambda17.10})/I_{\lambda15.01}$ ratio yields information on the gas dynamics, since the emission line optical thicknesses depend on both the ion column density and gas velocity distributions. Therefore spatial observations of resonant scattering may be used in identifying regions of turbulent and collective gas motions. However, since the RGS measurement technique relies on the use of grating elements, the spatial information is only obtained in the cross-dispersion direction at the instrument's focal plane, limiting its applicability in such measurements. In this study we have taken full advantage of this spatial dimension by analyzing long-exposure data of NGC 4636 and using five adjacent spectral regions extracted in the RGS cross-dispersion direction. We also fit the spectra with a cooling flow model to study the magnitude and radial characteristics of the flow. In addition, independent thermal analyses were conducted using the Chandra ACIS instrument to gain relevant 2D data to interpret the results of the RGS analyses. The paper is organized as follows: In Sect. \ref{data} we present the observational data and in Sect. \ref{preparation} the data preparation methods and analyses. In Sect. \ref{discussions} we discuss the results, and we conclude in Sect. \ref{concl}. | {\label{concl}} Observations of \ion{Fe}{xvii} resonant scattering revealed an asymmetric velocity field distribution in the core regions of the NGC 4636 group. In particular, we observed an asymmetry between the southern and northern parts of the X-ray halo; the analyses suggest that the IGM gas is turbulent in the south side, while the center and north side spectra express characteristics of collective gas motions. The X-ray images show that both N and S sides have been subjected to energetic AGN outbursts in the relatively recent past, and spectral maps show evidence of core sloshing with motion toward the NE. Both of these processes are capable of driving turbulence in the IGM gas. We suggest three alternatives to interpret our observation results: \newline 1) The IGM gas is turbulent in both sides, but the inclination of the AGN jet axis and 3D distribution of turbulent and laminar gas volumes are such that differences in the resonant scattering are observed in our line of sight. \newline 2) The IGM gas is turbulent in both sides, but the microturbulence velocity distributions are anisotropic between the N and S sides, so that in our line of sight, the radial microturbulence velocity distribution is smaller in the N side than in the S side. \newline 3) The core-sloshing gas flows drive turbulence to the S side and are the dominating source of microturbulence in the core region, leading to the distinct thermodynamical properties observed in the two sides. The modeling of cooling the flow yields a cumulative mass deposition rate of $\dot{M}\sim0.8$ M$_{\sun}$yr$^{-1}$ within the $2.4\arcmin\times 5.0\arcmin$ solid angle, a magnitude that generally agrees with the predictions of the traditional cooling flow model for NGC 4636 and \ion{O}{vi} observations made in the FUV band. Nevertheless, we find more than twice the mass deposition rate in the N side than the S side, instead of $\dot{M}$ being proportional to the X-ray surface brightness profile. Overall, we find that the \ion{Fe}{xvii} and \ion{O}{vii} resonant scattering, the magnitude of the cooling flow, and the presence of multiphase gas, including the high concentrations of $kT\lesssim0.4$ keV gas phases, are emphasized in regions C, N1, and N2. The analysis suggests that concurrent cooling and heating of IGM gas occurs in the core regions of the NGC 4636 galaxy group. | 16 | 7 | 1607.07444 |
1607 | 1607.00787_arXiv.txt | We present CARMA CO ($J=1\rightarrow0$) observations and \textit{Herschel} PACS spectroscopy, characterizing the outflow properties toward extremely young and deeply embedded protostars in the Orion molecular clouds. The sample comprises a subset of the Orion protostars known as the PACS Bright Red Sources (PBRS) (Stutz et al.). We observed 14 PBRS with CARMA and 8 of these 14 with \textit{Herschel}, acquiring full spectral scans from 55~\micron\ to 200~\micron. Outflows are detected in CO ($J=1\rightarrow0$) from 8 of 14 PBRS, with two additional tentative detections; outflows are also detected from the outbursting protostar HOPS 223 (V2775 Ori) and the Class I protostar HOPS 68. The outflows have a range of morphologies, some are spatially compact, $<$10000 AU in extent, while others extend beyond the primary beam. The outflow velocities and morphologies are consistent with being dominated by intermediate inclination angles (80\degr~$\ge$~$i$~$\ge$20\degr). This confirms the interpretation of the very red 24~\micron\ to 70~\micron\ colors of the PBRS as a signpost of high envelope densities, with only one (possibly two) cases of the red colors resulting from edge-on inclinations. We detect high-J (J$_{up}$~$>$~13) CO lines and/or H$_2$O lines from 5 of 8 PBRS and only for those with detected CO outflows. The far-infrared CO rotation temperatures of the detected PBRS are marginally colder ($\sim$230~K) than those observed for most protostars ($\sim$300~K), and only one of these 5 PBRS has detected [OI] 63~\micron\ emission. The high envelope densities could be obscuring some [OI] emission and cause a $\sim$20~K reduction to the CO rotation temperatures. | The earliest stage of the star formation process is characterized by a dense, infalling envelope of gas and dust surrounding a nascent protostar. This early phase, in particular, is known to be associated with powerful outflows \citep{arce2007,frank2014}. These outflows may ultimately play a role in halting the mass infall process and dispersing the envelope \citep{arce2006}, thereby contributing to the overall low efficiency of the star formation process \citep{offner2014}. These outflows develop rapidly and with velocities of $\sim$10 - 100 \kms\ the outflows may propagate by 0.1 pc in 10,000 yr - 1,000 yr timescales. Therefore, outflows are important to characterize at the youngest possible ages in order to understand their early evolution. The youngest identified protostars are known as Class 0 sources \citep{andre1993}; they are distinguished from more-evolved Class I sources by their cold bolometric temperatures (T$_{bol}$ $<$ 70 K; \citep{myers1993}) and/or ratio of submillimeter luminosity (L$_{submm}$) to bolometric luminosity ($L_{bol}$) being $>$ 0.5\%. These diagnostics indicate that Class 0 sources typically have denser and more massive infalling envelopes than Class I sources. In addition to the Class 0 sources, an earlier phase of the star formation process has been postulated, the first hydrostatic cores \citep[FHSC; e.g.,][]{larson1969}. A number of candidate FHSCs have been identified \citep{enoch2010,chen2010, pineda2011, schnee2012}; moreover, candidate FHSCs have quite low luminosities and bear some similarity to the \textit{Spitzer}-identified very low-luminosity sources \citep[VeLLOs][]{young2004,dunham2006}. The exact nature of the VeLLOs and candidate FHSCs remains unclear as it is difficult to distinguish bonafide FHSCs from sources that will go on to form very low mass stars \citep{dunham2014}. As part of the \textit{Herschel} Orion Protostar Survey (HOPS) \citep[e.g.,][]{fischer2010,stanke2010,ali2010,manoj2013,furlan2016}, a sample of 19 protostars with bright 70 \micron\ and 160 \micron\ emission and correspondingly faint or undetected (8 sources) 24 \micron\ emission were detected in the Orion star forming region \citep[][ hereafter ST13]{stutz2013}. We refer to these protostars as the PACS Bright Red Sources (PBRS); of the 19 PBRS, 12 were first identified as protostars by \textit{Herschel} and 7 \textit{Spitzer}-identified protostars also fulfilled the 24~\micron\ to 70~\micron\ color criteria (ST13). The PBRS are \textit{not} low-luminosity like the VeLLOs and candidate FHSCs; they have bolometric luminosities (L$_{bol}$) ranging between 0.65 L$_{\sun}$ and 30.6 L$_{\sun}$, with a median L$_{bol}$ of $\sim$3 L$_{\sun}$. Thus, the PBRS are the largest sample of extremely young protostars with typical luminosities; the median luminosity of Class 0 protostars is 3.5~L$_{\sun}$ in Orion and 1.4 L$_{\sun}$ in the nearby clouds \citep{dunham2014}. While the PBRS have only been well-characterized in Orion, similar examples are present in more nearby clouds (e.g., VLA 1623, IRAS 16293-2422), and \citet{sadavoy2014} identified several protostars in Perseus that were not classified as protostars in \textit{Spitzer} or undetected at 24~\micron\ \citep[i.e., HH211-mms][]{rebull2007}. We further characterized the envelopes of 14 PBRS using observations of the 2.9 mm dust continuum \citep{tobin2015}; that study, hereafter Paper I, focused specifically on the most deeply embedded and \textit{Herschel}-identified sources. The observed PBRS were all detected and found to have among the largest 2.9 mm luminosities of known Class 0 protostars. We also found that 6 out of 14 have visibility amplitudes that are flat within increasing uv-distance. The flat visibility amplitudes indicate that the 2.9 mm emission is very concentrated, and this finding, together with the high 2.9 mm luminosities, confirms that most PBRS have dense envelopes. This corroborates the interpretation of the spectral energy distribution (SED) model comparisons in ST13. The characterization of the PBRS from both the SEDs and millimeter continuum have led us to conclude that the PBRS may be among the youngest Class 0 objects. If the PBRS represent a distinct portion of early Class 0 evolution, as suggested by ST13, then the relative numbers of PBRS to Class 0 sources in Orion indicates that a `PBRS phase' could last $\sim$25,000 yr. This estimate assumes that the Class 0 phase lasts $\sim$150,000 yr \citep{dunham2014}. A remaining source of uncertainty in the interpretation of the PBRS as the youngest Class 0 protostars is their unknown disk/envelope inclination angles with respect to the plane of the sky. There is a degeneracy between high envelope densities versus high (nearly edge-on) inclinations that could not be mitigated due to the lack of emission shortward of 10 \micron\ toward most PBRS \citep[e.g.,][]{whitney2003,furlan2016}. Assuming that outflows are perpendicular to the disk or envelope midplanes, observations of outflows to constrain their orientations (e.g., in molecular lines) are an excellent way to estimate disk/envelope inclinations and further constrain the envelope properties. Furthermore, if the PBRS are among the youngest Class 0 protostars, then the sample as a whole represents an opportunity to examine the outflow properties of the youngest protostars. The jets and outflows from protostars are detected with a variety of complementary methods and the types of outflows and the ways to detect them also vary with evolution. Collimated jets detected in optical or near-infrared line emission are typically associated with more evolved Class I or Class II sources \citep[e.g., HH111]{reipurth1997,reipurth2010}, while Class 0 protostars typically have a molecular outflow observable in only millimeter lines of CO and other molecules \citep{arce2007, frank2014}. However, this does not mean there is no collimated jet emission, just that it may be undetectable due to high levels of obscuration. The molecular outflow emission toward some low-mass protostars has an angular dependence of velocity, with low-velocity material at the edges of the outflow cavity and velocities as high as $\sim$ 100 \kms\ along the main axis of the outflow \citep[e.g.,][]{santiago2009,hirano2011}. Jet-like features can also be seen in shock-tracing molecules such as SiO and SO \citep[e.g.,][]{lee2008,lee2009}. The velocity gradients along the outflow axis also offer crucial information of disk-protostar orientation \citep[e.g.,][]{cabrit1986, lee2000}. Far-infrared spectroscopy with the \textit{Infrared Space Observatory} and the \textit{Herschel Space Observatory} has also been found to be an excellent probe of the physical conditions of outflows from young stars. The high-J CO (J$_u$ $>$ 13) and H$_2$O transitions, in addition to OH and [OI] transitions, probe the warm and hot outflow conditions on scales very near the protostar and the jet driving source \citep[e.g., ][]{vankempen2010,karska2013, green2013,manoj2013}. The lines are thought to be excited primarily by shocks \citep{manoj2013}, with UV radiation photo-dissociating H$_2$O, causing lower abundances relative to non-irradiated shock models \citep{karska2014}. The initial development of the outflows and their subsequent breakout from their surrounding envelopes are still quite uncertain. Outflows have also been detected from VeLLOs and candidate FHSCs \citep{dunham2011, pineda2011, schnee2012, tobin2015}. Theory has predicted that such young objects can indeed produce the slow outflows ($\sim$2 - 7 \kms) that have been observed \citep{price2012}, and the outflows may develop prior to the formation of a rotationally-supported accretion disk \citep[e.g.,][]{li2013,li2014}. However, it is still uncertain how quickly more powerful outflows emerge in protostars; do the outflows have a steady growth in power as the source luminosity (from accretion) increases or do they only become powerful once a certain threshold in luminosity is reached? In order to examine the outflow conditions from the youngest known Class 0 protostars, we have obtained interferometric observations of the CO ($J=1\rightarrow0$) molecular line and far-infrared spectroscopy with the \textit{Herschel Space Observatory} toward the PBRS in the Orion A and B molecular clouds. The youth and number of PBRS sources in Orion offers an unique opportunity to examine the properties of outflows toward objects that are consistent with being among the youngest protostars. Furthermore, spectrally and spatially resolved observations of the molecular outflows toward these protostars will enable us to constrain the range of possible inclination angles of the protostellar sources, ensuring that their characterization as the youngest protostars is not strongly influenced by orientation. We have observed 14 PBRS (from the full sample of 19 cataloged by \citet{stutz2013} and Paper I) with the Combined Array for Research in Millimeter-wave Astronomy (CARMA), focusing on the \textit{Herschel}-detected PBRS sample. We observed the protostars in both the dust continuum and spectral line emission to examine the envelope and outflow properties of these sources. We discuss the observations in Section 2, our outflow results from CO ($J=1\rightarrow0$) and \textit{Herschel} spectroscopy are presented in Section 3, we discuss the results in Section 4, and summarize our main conclusions in Section 5. | We have presented an observational study of both the cold and warm/hot molecular gas in outflows from the youngest known protostars in the Orion molecular clouds, the PACS Bright Red Sources (PBRS). The cold gas was probed toward 14 out of 19 PBRS using observations of the CO ($J=1\rightarrow0$) transition from CARMA, and the warm/hot gas was examined for 8 out of the 19 PBRS using full spectral scans (55~\micron\ to 200~\micron) from the \textit{Herschel} PACS far-infrared spectrometer. Finally, we also examined \textit{Spitzer} 4.5 \micron\ imaging to look for evidence of both compact and extended outflow activity from both scattered light and shocked H$_2$ emission. The results from the follow-up work done in this study and Paper I demonstrate the critical need for complementary data in the determining the nature of protostellar sources that are otherwise only characterized by their SEDs. Our main conclusions are as follows. 1. We detect clear outflows toward 8 out of 14 PBRS (119019, 090003, 093005, 135003, HOPS~373, 082012, and 019003) in the CO ($J=1\rightarrow0$) molecular transition. There is tentative evidence for outflows toward an additional three PBRS (HOPS~372, 302002, and 061012). We also detect outflows from two non-PBRS HOPS 223, a FU Ori-like outbursting protostar \citep{fischer2012} and HOPS 68 \citep{poteet2011}; the HOPS 68 outflow also appears to be quadrupolar. No detectable outflow activity is found toward the PBRS 097002, 082005, 091015, and 091016 in CO ($J=1\rightarrow0$), 4.5 \micron\ emission, or far-infrared spectroscopy (only 091015 and 091016). 2. The outflows toward 090003 and 093005 are the most compact, subtending less than 20\arcsec\ (8400 AU) in total extent, having dynamical ages $\le$2,500 yr. These outflows are also found to have momenta, energies, and forces that are at the low end for Class 0 protostars. This observation, in addition to the lack of detectable outflows toward several other PBRS, leads us to suggest that outflows may start out weak in protostellar sources and become more energetic with time. These sources are also the only ones with flat visibility amplitudes to have detected outflows and we find a tentative tendency for the sources with flat visibility amplitudes in the 2.9~mm continuum (see Paper I) to either have no detected outflow activity or the most spatially compact outflows. This is further evidence for the sources with flat visibility amplitude being among the youngest protostars and the youngest PBRS. 3. The outflow from 082012 is extremely powerful, with red-shifted emission detected out to $+$40 \kms\ from line center and extent greater than the CARMA primary beam. Its total energy is in excess of any individual outflow in the NGC 1333 star forming region \citep{plunkett2013} and comparable to some of the most powerful known outflows from Class 0 protostars \citep[e.g.,][L1448C]{hirano2011}. 4. We detect far-infrared CO emission lines toward 6 out of the 8 PBRS observed. H$_2$O lines are detected toward 5 out of 8 PBRS, and OH and [OI] are detected toward 1 PBRS. The far-infrared CO, H$_2$O, and [OI] lines do not reveal outflows in the absence of outflow detections from other diagnostics. The CO luminosities and [OI] detections/upper limits are consistent with the results from larger samples of Class 0 protostars. However, the CO rotation temperatures tend to be lower than the typically observed 300~K CO rotation temperature for protostars; however, given the uncertainties the PBRS are consistent with the larger samples. Nevertheless, with a simple calculation of envelope opacity to a radius of 1000~AU, we find that the observed rotation temperatures of the PBRS could appear $\sim$20~K lower due to envelope opacity, given that the PBRS seem to have denser envelopes than typical Class 0 protostars. We wish to thank the anonymous referee for excellent suggestions which have significantly improved the quality of the manuscript. The authors also wish to acknowledge fruitful discussions with M. Dunham, L. Kristensen, and J. Mottram regarding this work. J.J.T. is currently supported by grant 639.041.439 from the Netherlands Organisation for Scientific Research (NWO). J.J.T acknowledges past support provided by NASA through Hubble Fellowship grant \#HST-HF-51300.01-A awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. The work of A.M.S. was supported by the Deutsche Forschungsgemeinschaft priority program 1573 ('Physics of the Interstellar Medium'). AK acknowledges support from the Foundation for Polish Science (FNP) and the Polish National Science Center grant 2013/11/N/ST9/00400. This work is based in part on observations made with Herschel, a European Space Agency Cornerstone Mission with significant participation by NASA. Support for this work was provided by NASA through an award issued by JPL/Caltech. We are very grateful to have had the opportunity to conduct these follow-up observations with the CARMA array in California. The discontinuation of support for this productive facility is a loss that will continue to be felt into the future. Support for CARMA construction was derived from the states of Illinois, California, and Maryland, the James S. McDonnell Foundation, the Gordon and Betty Moore Foundation, the Kenneth T. and Eileen L. Norris Foundation, the University of Chicago, the Associates of the California Institute of Technology, and the National Science Foundation. Ongoing CARMA development and operations are supported by the National Science Foundation under a cooperative agreement, and by the CARMA partner universities. {\it Facilities:} \facility{CARMA}, \facility{\textit{Herschel}}, \facility{\textit{Spitzer}}, \facility{Magellan} \appendix | 16 | 7 | 1607.00787 |
1607 | 1607.07672_arXiv.txt | We constructed a comprehensive model atom for C\ione\ -- C\ii\ using the most up-to-date atomic data available and evaluated the non-local thermodynamic equilibrium (NLTE) line formation for C\ione\ and C\ii\ in classical 1D models representing the atmospheres of A and late B-type stars. Our NLTE calculations predict the emission that appears at effective temperature of 9250 to 10\,500~K depending on log~$g$ in the C\ione\ 8335, 9405\,\AA\ singlet lines and at \Teff~$>$~15\,000~K (log~$g$ = 4) in the C\ione\ 9061 -- 9111\,\AA\,, 9603 -- 9658\,\AA\, triplet lines. A prerequisite of the emission phenomenon is the overionization-recombination mechanism resulting in a depopulation of the lower levels of C\ione\ to a greater extent than the upper levels. Extra depopulation of the lower levels of the transitions corresponding to the near-infrared lines, is caused by photon loss in the UV lines C\ione\ 2479, 1930, and 1657\,\AA. We analysed the lines of C\ione\ and C\ii\ in Vega, HD~73666, Sirius, 21~Peg, $\pi$~Cet, HD~22136, and $\iota$ Her taking advantage of their observed high-resolution spectra. The C\ione\ emission lines were detected in the four hottest stars, and they were well reproduced in our NLTE calculations. For each star, the mean NLTE abundances from lines of the two ionization stages, C\ione\ and C\ii, including the C\ione\ emission lines, were found to be consistent. We show that the predicted C\ione\ emission phenomenon depends strongly on whether accurate or approximate electron-impact excitation rates are applied. | Rapid development of observational techniques over the last decade has resulted in dramatic improvement of the quality of spectral observations in astronomy. Thanks to Echelle spectrographs, we are able to cover wide spectral regions with one exposure, obtaining spectra with a high spectral resolution of up to $R$~=~$\lambda$/$\Delta\lambda$~=~120\,000 and high signal-to-noise ratio. Huge amounts of high quality spectra are collected in different archives with the open access for the astronomical community. High quality spectra require immediately an adequate improvement in theoretical methods of spectrum analysis, model atmospheres, line formation scenarios, etc. \citet{2009A&A...503..945F} studied three apparently normal A and B sharp-lined stars and found the C\ione\ emission lines at 8335, and 9406\,\AA\, in the hottest of them, $\pi$~Cet. The absorption C\ione\ lines in the 7111--7120\,\AA\, range are rather shallow and allow us to determine only an upper limit for the abundance. In local thermodynamic equilibrium (LTE) analysis, $\pi$~Cet reveals a disparity between C\ione\ and C\ii, in line with \citet{1990ApJS...73...67R}. For the cooler star, 21~Peg, \citet{2009A&A...503..945F} could not obtain consistent abundances from different C\ione\ lines, namely the abundances obtained from C\ione\ 4932, 5052, 5380~\AA\, were found to be significantly smaller than those from other C\ione\ lines. \citet{2009A&A...503..945F} discussed several mechanisms of the C\ione\ emission in $\pi$~Cet. Emission lines may form, if the star has a chromosphere, however, there is no evidence for chromospheric activity in $\pi$~Cet. This star was classified as a Herbig Ae/Be star due to small infrared (IR) excess \citep{1998A&A...331..211M}, suggesting the existence of a circum-stellar disc. This is supported by detection of an emission signature in H$_\alpha$ \citep{2009A&A...503..945F}. However, half-widths of the C\ione\ emission lines do not differ from those of the C\ione\ absorption lines, giving evidence for their common origin in the star's atmosphere. \citet{2009A&A...503..945F} proposed that the emission can be caused by the departures from LTE in the C\ione\ line formation. In the literature there are examples of successfully reproducing the observed emission lines by non-local thermodynamic equilibrium (NLTE) calculations, for example, Mg\ione\ 12~$\mu$m in the Sun \citep{Carlsson92}, Mn\ii\ 6122-6132\,\AA\ in the three late type B stars \citep{2001A&A...377L..27S}, C\ii\ 6151, 6462\,\AA\ in $\tau$~Sco (B0V) and C\ii\ 6462\,\AA\ in HR~1861 (B1V) \citep{2008A&A...481..199N}. There are few NLTE studies of C\ione/C\ii\ in the early A to late B stars. \citet{1996A&A...312..966R} computed a negative and small absolute value of less than 0.05~dex NLTE abundance corrections for lines of C\ii\ in A-type stars. \citet{2001AA...379..936P} obtained very small NLTE corrections for the C\ione\ and C\ii\ lines in the visible region using a Vega model atmosphere with an effective temperature \Teff\ = 9550~K. The NLTE carbon abundance analysis of 20 sharp-lined early B stars in the effective temperature range of 16000--33000~K, including two stars with \Teff\ $<$ 18000~K, was performed by \citet{2012A&A...539A.143N}. None of the carbon NLTE papers investigated the NLTE effects for the C\ione\ near-IR lines in the early A to late B stars. This paper aims to understand mechanism(s) of the C\ione\ emission in B-type stars and to treat the method of accurate abundance determination from different lines of C\ione\ and C\ii\ based on NLTE line formation. We construct a comprehensive model atom for C\ione\ -- C\ii\ using the most up-to-date atomic data available so far and analyse lines of C\ione\ and C\ii\ in high-resolution spectra of reference A and B-type stars with well-determined stellar parameters. The paper is organized as follows. Section\,\ref{Sect:atom} describes an updated model atom of C\ione-C\ii and discusses departures from LTE for C\ione\ -- C\ii\ in the model atmospheres of A-B stars and mechanisms driving the C\ione\ emission lines. In Sect.\,\ref{Sect:Stars}, we analyse the C\ione\ near-IR emission lines observed in the four B-type stars and determine the C abundance of the selected A-B stars. We inspect the abundance differences between different lines of a common species and between two ionisation stages, C\ione\ and C\ii. Our conclusions are summarized in Sect.\,\ref{Sect:Conclusions}. \begin{figure*} \begin{minipage}{160mm} \includegraphics[scale=0.6]{Figure1.eps} \caption{The term diagram for singly-ionized carbon. The dashed lines indicate the seven transitions, where the investigated spectral lines arise. } \label{Grot} \end{minipage} \end{figure*} \begin{figure*} \begin{minipage}{175mm} \parbox{0.3\linewidth}{\includegraphics[scale=0.28]{Figure2.eps}\\ \centering} \hspace{0.2\linewidth} \parbox{0.3\linewidth}{\includegraphics[scale=0.28]{Figure3.eps}\\ \centering} \hfill \\[0ex] \parbox{0.3\linewidth}{\includegraphics[scale=0.28]{Figure4.eps}\\ \centering} \hspace{0.2\linewidth} \parbox{0.3\linewidth}{\includegraphics[scale=0.28]{Figure5.eps}\\ \centering} \hfill \caption{NLTE and LTE fractions of C~I, C~II, and C~III in the model atmospheres of different effective temperature. } \label{balance} \end{minipage} \end{figure*} | \label{Sect:Conclusions} The motivation of this work was to solve two problems in stellar astrophysics: the search for explanation of the appearance of emission lines in C\ione\ in the near-IR spectral region and a reliable determination of carbon abundances for AB-type stars. We constructed a comprehensive model atom for C\ione\ -- C\ii\ using the most up-to-date atomic data and evaluated the NLTE line formation for C\ione\ and C\ii\ in classical 1D models representing the atmospheres of A and late B-type stars. Our NLTE calculations predict that some lines of C\ione\ in the near IR spectral range may appear as emission lines depending on the atmospheric parameters. The emission appears first in the C\ione\ 8335\,\AA\ and 9405\,\AA\ singlet lines at effective temperature of 9250~K to 10\,500~K depending on the value of log~$g$. It is strengthened toward higher \Teff, reaches a maximal level at \Teff\ = 16\,000~K (log~$g$ = 3) and almost disappears at \Teff\ = 22\,000~K. The C\ione\ triplet lines at 9061--9111\,\AA\ and 9603--9658\,\AA\ come into emission at \Teff\ $>$ 15\,000~K (log~$g$ = 4). The mechanisms driving the C\ione\ emission can be understood as follows. A prerequisite of the emission phenomenon is the overionization-recombination mechanism resulting in a depopulation of the lower levels of C\ione\ to a greater extent than the upper levels. Extra depopulation of 3s$^{1}$P$^{\circ}$ and 3s$^{3}$P$^{\circ}$, which are the lower levels of the transitions corresponding to the listed near-IR lines, can be caused by photon loss in the UV lines C\ione\ 2479, 1930, and 1657\,\AA. In the models with \Teff\ of about 10\,000~K only, C\ione\ 2479\,\AA\ plays a role draining population of the 3s$^{1}$P$^{\circ}$ singlet level effectively in the layers, where the C\ione\ 8335 and 9405\,\AA\ lines form, while these layers are optically thick for radiation at 1930 and 1657\,\AA. With increasing \Teff, formation depths of all the C\ione\ UV lines shift inwards, resulting in a depopulation of not only 3s$^{1}$P$^{\circ}$, but also 3s$^{3}$P$^{\circ}$ and the lower levels of the emission triplet lines at 9061 -- 9111\,\AA\ and 9603 -- 9658\,\AA. Our theoretical results were confirmed with observations of the reference stars. The stellar sample consists of seven bright and apparently slow-rotating A- and late B-type stars in the solar neighbourhood. Our analysis is based on high S/N and high-resolution spectra with a broad wavelength coverage. The C\ione\ emission lines were measured in the four hottest stars, with \Teff\ $\ge$ 10\,400~K, and they were well reproduced in our NLTE calculations. For each star, the mean NLTE abundances from lines of the two ionisation stages, C\ione\ and C\ii, including the C\ione\ emission lines, were found to be consistent. Thus, we settled the dispute on whether the C\ione\ emission in the late-B stars is produced by the circumstellar disc, or vertical stratification / horizontal inhomogeneity of the atmosphere, or the NLTE effects in the atmosphere. The six of our stars reveal highly uniform and close-to-solar carbon abundance. We confirm a significant underabundance of carbon in Sirius, with [C/H] = $-0.72$. We show an importance of applying accurate atomic data to the statistical equilibrium calculations. In particular, the C\ione\ emission phenomenon turns out to be extremely sensitive to varying electron-impact excitation rates. If the latter are not accounted for properly, the stellar C\ione\ emission lines cannot be reproduced. The results obtained in this study favour the collisional data for C\ione\ from predictions of \citet{2013PhRvA..87a2704W}. It is worth noting, WELs of Mg\ii, Si\ii, P\ii, Ca\ii, Cr\ii, Fe\ii, Ni\ii, Cu\ii, and Hg\ii\ were detected in the near-UV and visible spectral regions for several stars not showing any trace of chromosphere \citep{2000ApJ...530L..89S, 2000AA...362L..13W, 2004AA...418.1073W,2007AA...475.1041C}. According to \citet{2008CoSka..38..279W}, WELs of some metals are observed in sharp-lined spectra of mid- to late-B type stars. The WELs are detected over a range of element abundance and are found among both chemically-normal and chemically-peculiar stars. The NLTE line--formation calculations for these specific ions are highly desirable to understand mechanisms of the observed emission. {\it Acknowledgements.} This research is based on observations obtained with MegaPrime/MegaCam, a joint project of CFHT and CEA/IRFU, at the Canada--France--Hawaii Telescope (CFHT) which is operated by the National Research Council (NRC) of Canada, the Institut National des Science de l$'$Univers of the Centre National de la Recherche Scientifique (CNRS) of France, and the University of Hawaii. We thank Oleg Zatsarinny for providing us the data on effective collision strengths for the C\ione\ transitions published in \citet{2013PhRvA..87a2704W}. This work was supported in part by the Russian Foundation for Basic Research (grants 16-32-00695 and 15-02-06046). | 16 | 7 | 1607.07672 |
1607 | 1607.05677_arXiv.txt | {We test the theoretical predictions of several cosmological models against different observables to compare the indirect estimates of the current expansion rate of the Universe determined from model fitting with the direct measurements based on Cepheids data published recently.} {We perform a statistical analysis of type Ia supernova (SN Ia), Hubble parameter, and baryon acoustic oscillation data. A joint analysis of these datasets allows us to better constrain cosmological parameters, but also to break the degeneracy that appears in the distance modulus definition between $H_0$ and the absolute B-band magnitude of SN Ia, $M_0$.} {From the theoretical side, we considered spatially flat and curvature-free $\Lambda$CDM, $w$CDM, and inhomogeneous Lema\^{i}tre-Tolman-Bondi (LTB) models. To analyse SN Ia we took into account the distributions of SN Ia intrinsic parameters.} {For the $\Lambda$CDM model we find that $\Omega_m=0.35\pm0.02$, $H_0=(67.8 \pm 1.0)\,$km$\,$s$^{-1}/$Mpc, while the corrected SN absolute magnitude has a normal distribution ${\cal N}(19.13,0.11)$. The $w$CDM model provides the same value for $\Omega_m$, while $H_0=(66.5\pm1.8)\,$km$\,$s$^{-1}/$Mpc and $w=-0.93\pm0.07$. When an inhomogeneous LTB model is considered, the combined fit provides $H_0=(64.2 \pm 1.9)\,$km$\,$s$^{-1}/$Mpc.} {Both the Akaike information criterion and the Bayes factor analysis cannot clearly distinguish between $\Lambda$CDM and $w$CDM cosmologies, while they clearly disfavour the LTB model. For the $\Lambda$CDM, our joint analysis of the SN Ia, the Hubble parameter, and the baryon acoustic oscillation datasets provides $H_0$ values that are consistent with cosmic microwave background (CMB)-only Planck measurements, but they differ by $2.5\sigma$ from the value based on Cepheids data.} | \label{sec:int} Since the early determination by Hubble \citep{Hubble29}, the Hubble constant was for a long time believed to be between 50 and 100 km$\,$s$^{-1}/$Mpc \citep{Kirschner03}. Recent findings are obtained by means of space facilities, improved control of systematics, and the use of different calibration techniques, as in the Hubble Space Telescope Key Project, which estimated $H_0=(72 \pm 8)\,$km$\,$s$^{-1}/$Mpc \citep{HST01}. \citet{Riess16} provided the most recent direct estimate of the expansion rate of the Universe: $H_0=(73.0 \pm 1.8)\,$km$\,$s$^{-1}/$Mpc. Together with these extraordinary improvements in the direct determination of the distance ladder, there are by now different classes of observations that allow an indirect estimate of the Hubble constant. Among others, the observations of the cosmic microwave background (CMB) anisotropy by WMAP \citep{WMAP9} and \citet{Planck15} satellites yielded values of $H_0=(70.0 \pm 2.2)$ km$\,$s$^{-1}/$Mpc and $H_0=(67.27 \pm 0.66)$ km$\,$s$^{-1}/$Mpc, respectively. In addition to the CMB anisotropy measurements, other observables have been crucial to constrain the cosmological parameters, such as type Ia supernovae (SN Ia). The high-z supernova search team led by Adam Riess together with Brian P. Schmidt \citep{Riess98} and the supernova cosmology project led by Saul Perlmutter \citep{Perlmutter99} reported the first evidence for an accelerated cosmic expansion. Since then, the number of observed SN Ia increased by about an order of magnitude. Different publicly available compilations have been used to constrain cosmological models: Union2 \citep{Amanullah10}, Union2.1 \citep{Suzuki12}, Constitution set \citep{Hicken09}, and JLA \citep{Betoule14}. The results confirm the need for a late accelerated expansion of the Universe, consistent with the findings of the WMAP and Planck missions. Unfortunately, the observations of SN Ia by themselves are not able to provide a value for the local expansion rate of the Universe, $H_0$, since this parameter is degenerate with the SN absolute magnitude. However, there are other cosmological observables that are more directly sensitive to the value of the Hubble constant. On one hand, passively evolving red galaxies, which are dominated by the older stellar population, whose age can be accurately estimated from a spectroscopic analysis (also known as cosmic chronometers), can be used to provide the redshift dependence of the expansion rate, $H(z)$, as suggested by \citet{Jimenez02}. Fitting these observational Hubble data (OHD), \citet{Liu15} found a value of $H_0=67.6\,$km$\,$s$^{-1}/$Mpc. On the other hand, the baryon acoustic oscillation (BAO) data have been used to constrain the cosmological parameters, providing results that agree with the most recent findings of the Planck Collaboration. In particular, a recent estimate of the Hubble constant provides $H_0 = (68.11 \pm 0.86)$ km$\,$s$^{-1}/$Mpc \citep{Cheng15}.\\ \indent It is clear that the indirect estimates of the Hubble constant lead to lower values of $H_0$ compared to the direct measurements. Even the earlier estimate of $H_0=(73.8\pm2.4)\,$km$\,$s$^{-1}/$Mpc by \citet{Riess11} and the latest one \citep{Riess16} contradict the most recent result from Planck (TT, TE, EE + lowP) at the $2.6\sigma$ and $3.0\sigma$ level, respectively. The question now is whether this difference hides new physics beyond what is by now commonly called the concordance model. This point has been addressed by \citet{Efstathiou14}, who reanalysed the Cepheid data used by \citet{Riess11}. He obtained a value $H_0=(72.5\pm2.5)$ km$\,$s$^{-1}/$Mpc, reducing the difference to Planck to only $2 \sigma$ and concluding that there is no evidence for new physics (see also \citet{Chen11} and \citet{Marra13}). We here extend this discussion to determine whether any difference is present when observables other than CMB are considered. To do so, we perform a separate and a joint analysis of SN Ia, OHD, and BAO data. The joint analysis promises to provide more stringent constraints on the cosmological models, and to break the degeneracy between the SN absolute magnitude and the Hubble constant, which is peculiar to the SN analysis. Several SN datasets (such as Union and Constitution) provide cosmological distance moduli that are derived assuming a flat $\Lambda$CDM model. Hence, these datasets need to be treated with caution when used to constrain cosmological models that are different from $\Lambda$CDM. We used the JLA dataset, which provides model-independent apparent magnitudes instead of model-dependent distance moduli. Moreover, the increase in the amount of data and the improvement in systematics imply that a more complete statistical analysis is necessary. We therefore followed the approach proposed by \citet{Nielsen15} for the SN data analysis. For the theoretical models we considered the standard flat $\Lambda$CDM model and its extensions, which include the curvature-free $k\Lambda$CDM model and a dark energy model characterised by an equation of state (EoS) $p=w\rho c^2$, with $w=const$. In addition, we also considered a different class of models, based on the Lema\^{i}tre-Tolman-Bondi (LTB) metric, which describes an isotropic but inhomogeneous Universe \citep{Lemaitre33,Tolman34,Bondi47,Krasinski97}, to stress the dependence of the Hubble constant estimates on the assumed theoretical model. The plan of the paper is as follows. In \Cref{sec:th} we review the theoretical models we considered. In \Cref{sec:analysis} we review the observables and datasets used in our analysis. In \Cref{sec:res} we show the results of our comparison between theory and observations. Finally, in \Cref{sec:con} we summarise our findings and conclusions. | \label{sec:con} The combined analysis of JLA, OHD and BAO datasets allowed us to reach more stringent constraints on the cosmological parameters\footnote{Our analysis has been implemented in {\it Mathematica} 10. The code is available upon request.} and to break the degeneracy between the SN absolute magnitude and the cosmic expansion rate. Our main findings can be summarised as follows. By fitting the cosmological and the SN intrinsic parameters to the combined set of JLA, OHD, and BAO data, we constrained the distributions of SN absolute magnitude, stretch, and colour. The resulting values are cosmological-model independent, with the exception of the $M_0$ value obtained for LTB. The method we used can in principle be extended to include the effects related to mass and metallicity of the host galaxies. We studied the $\Lambda$CDM model and its extensions to consider non-vanishing spatial curvature and different assumptions for the DE component. The combined analysis clearly prefers the concordance model, as it forces the curvature to vanish and the DE EoS to be consistent with $w=-1$. We also studied an LTB model with a Gaussian profile, which is strongly disfavoured with respect to the concordance model by information criteria, such as AIC analysis or Bayes factor. For the $\Lambda$CDM model, the JLA+OHD+BAO analysis provides a value of $H_0=(67.8 \pm 1.0) \,$km$\,$s$^{-1}/$Mpc that is fully consistent with the Planck (TT, TE, EE + lowP) result. This means that the difference with the direct measurements by \citet{Riess16} is very likely not due to systematics in the Planck CMB measurements. It also seems difficult to reconcile direct and indirect $H_0$ measurements by considering an additional source of dark radiation in the early Universe, as this would not affect the JLA+OHD fit, which is still consistent with Planck. Therefore, it is still unclear wether it is necessary to extend the concordance $\Lambda$CDM model. | 16 | 7 | 1607.05677 |
1607 | 1607.06397_arXiv.txt | Balloon experiments are an economically feasible method of conducting observations in astronomy that are not possible from the ground. The astronomical payload may include a telescope, a detector, and a pointing/stabilization system. Determining the attitude of the payload is of primary importance in such applications, to accurately point the detector/telescope to the desired direction. This is especially important in generally unstable lightweight balloon flights. However, the conditions at float altitudes, which can be reached by zero pressure balloons, could be more stable, enabling accurate pointings. We have used the Inertial Measurement Unit (IMU), placed on a stratospheric zero pressure balloon, to observe 3-axis motion of a balloon payload over a flight time of $\sim 4.5$ hours, from launch to the float altitude of 31.2 km. The balloon was launched under nominal atmospheric conditions on May 8th 2016 from a Tata Institute of Fundamental Research Balloon Facility, Hyderabad. | High-altitude balloon platforms are an economical alternative to space missions for testing instruments as well as for specific classes of observations, particularly those that require a rapid response such as comets, or other transients. Telescopic platforms at high altitudes have significant advantages over operations from the ground enabling observations at forbidden wavelengths. A UV telescope (200--400 nm) in stratosphere with aperture of just 6 inch in diameter with sufficient pointing stability/accuracy and a 1K$\times 1$K CCD array could provide wide-field images with FWHM better than $1^{\as}$ approaching the diffraction limit (Fesen \& Brown, 2015), similar to that of space observatories but at a much lower cost. We have initiated a high-altitude balloon program at the Indian Institute of Astrophysics to develop low-cost instruments for use in atmospheric and astronomical studies (Nayak et al. 2013, Safonova et al. 2016). We have developed a number of payloads which operate in the near-ultraviolet (NUV), but we are limited to weights under 6 kg for regulatory reasons which constrains our payload size (Sreejith et al. 2016a). Our first experiments were of atmospheric lines (Sreejith et al. 2016b) where the pointing stability is less important, but we plan to observe astronomical sources for which a pointing mechanism is required. Unlike in space missions, in stratospheric balloon pointing systems the payload is attached to the balloon by a long flight train. Besides transmitting the balloon's buoyant force, the flight train is the source of disturbances that the pointing control must reject. A typical stratospheric balloon has a train of several meters in length comprised of a recovery parachute beneath with flight termination systems, which is connected to a gondola consisting of scientific equipment with communication and associated electronics. The balloon flight, planned by the Tata Institute of Fundamental Research Balloon Facility (TIFR-BF) on 6--8th May 2016, provided a timely opportunity to piggyback attitude instrumentation on the gondola solely to measure its natural motion during the float phase of the flight. We expect that this information would be helpful in efforts to characterize disturbances that could be expected on any balloon-borne pointing system. | The flight data derived from our x-IMU and from measurements by the TIFR-BF main payload are consolidated in Table~\ref{table: flight_summary}. \begin{table}[hb!] \caption{Summary of the flight ($\sim 2.5$ hrs at float altitude)} \label{table: flight_summary} \begin{center} \begin{tabular}{|l|c|} \hline \rule[-1ex]{0pt}{3.5ex} Maximum height reached by payload & 31.4 Km \\ \hline \rule[-1ex]{0pt}{3.5ex} Total duration of the flight & 5 hrs 56 mins \\ \hline \rule[-1ex]{0pt}{3.5ex} Average ascent rate & 4.59 m/s \\ \hline \rule[-1ex]{0pt}{3.5ex} Float altitude & 31.2 km \\ \hline \rule[-1ex]{0pt}{3.5ex} Float reached (time) & 08:35 am IST \\ \hline \rule[-1ex]{0pt}{3.5ex} \begin{tabular}{@{}c@{}} RMS motion of payload \\ during float in azimuth ($^{\circ}/s$) \end{tabular} & 0.5865 \\ \hline \rule[-1ex]{0pt}{3.5ex} \begin{tabular}{@{}c@{}} RMS motion of payload \\ during float in elevation ($^{\circ}/s$) \end{tabular} & 0.0104 \\ \hline \rule[-1ex]{0pt}{3.5ex} \begin{tabular}{@{}c@{}} RMS motion of payload \\ during float in tilt ($ ^{\circ}/s$) \end{tabular} & 0.01 \\ \hline \rule[-1ex]{0pt}{3.5ex} \begin{tabular}{@{}c@{}} Average temperature \\ inside payload\end{tabular} & 34.35$^{\circ}$C \\ \hline \rule[-1ex]{0pt}{3.5ex} \begin{tabular}{@{}c@{}} Average acceleration \\ of payload (X axis)\end{tabular} & 0.176 g \\ \hline \rule[-1ex]{0pt}{3.5ex} \begin{tabular}{@{}c@{}} Average acceleration \\ of payload (Y axis)\end{tabular} & 0.0036 g\\ \hline\rule[-1ex]{0pt}{3.5ex} \begin{tabular}{@{}c@{}} Average acceleration \\ of payload (Z axis)\end{tabular} & 1.026 g\\ \hline \end{tabular} \end{center} \end{table} \begin{enumerate} \item During the flight the payload reached the maximum height of 31.2 km, where the outside temperature was $-32^{\circ} $C. The temperature inside the payload stayed above $0^{\circ} $C. This shows that the electronic components inside the payload were thermally insulated. \item The stratospheric conditions during the TIFR flight at float are more stable than the near-surface conditions we have experienced during our previous tethered launches at the IIA, and comparable to our previous stratospheric flights. The full analysis is presented in the forthcoming paper (Nirmal et al. 2016). \end{enumerate} | 16 | 7 | 1607.06397 |
1607 | 1607.03784_arXiv.txt | The effects of a nonminimally coupled curvature-matter model of gravity on a perturbed Minkowski metric are presented. The action functional of the model involves two functions $f^1(R)$ and $f^2(R)$ of the Ricci scalar curvature $R$. This work expands upon previous results, extending the framework developed there to compute corrections up to order $O\left(1\slash c^4\right)$ of the 00 component of the metric tensor. It is shown that additional contributions arise due to both the non-linear form $f^1(R)$ and the nonminimal coupling $f^2(R)$, including exponential contributions that cannot be expressed as an expansion in powers of $1/r$. Some possible experimental implications are assessed with application to perihelion precession. | Dark matter and dark energy are key contemporary concepts used to account, for instance, for the astrophysical problem of the flattening of galactic rotation curves and the cosmological issue of the accelerated expansion of the universe, respectively. Dark energy accounts for 69$\%$ of the energy budget of the universe \cite{Planck}; among several other proposals, it has been the object of several so-called "quintessence" models \cite{quintessence}, which posit the existence of scalar fields with negative pressure, as an alternative to a suitably adjusted Cosmological Constant, which presents the eponymous problem of reconciling the large order of magnitude difference between its observed and predicted values \cite{Weinberg}. Dark matter searches focus on the characterization of additional matter species arising from extensions to the Standard Model of particles, collectively dubbed as weak-interacting massive particles (WIMPS) such as, for instance, neutralinos or axions \cite{DM}. As an alternative, some proposals assume that both dark components may be described in a unified fashion \cite{scalarfield,Chaplygin}. Other models assume that, instead of additional matter species, the fundamental laws of General Relativity (GR) may be incomplete, prompting {\it e.g.} for corrections and alternatives to the Einstein-Hilbert action. Among such theories, those involving a nonlinear corrections to the geometric part of the action via the scalar curvature, aptly called $f(R)$ theories, have gained much attention (see Ref. \cite{felice} for a thorough discussion). These can be extended also to include a nonminimal coupling (NMC) between the scalar curvature and the matter Lagrangian density, leading to an even richer phenomenology and implying that the energy-momentum tensor may not be (covariantly) conserved \cite{BBHL} (see also Ref. \cite{Lobo} for a more general model). NMC models have yielded several interesting results, including the impact on stellar observables \cite{stelobserv}, energy conditions \cite{energcondit}, equivalence with multi-scalar-tensor theories \cite{multiscalar}, possibility to account for galactic \cite{drkmattgal} and cluster \cite{drkmattclus} dark matter, cosmological perturbations \cite{cosmpertur}, a mechanism for mimicking a Cosmological Constant at astrophysical scales \cite{mimlambda}, post-inflationary reheating \cite{reheating}, dark energy \cite{curraccel,Friedmann,Ribeiro}, dynamical impact of the choice of the Lagrangian density of matter \cite{dynimpac1,dynimpac2}, gravitational collapse \cite{gravcollapse} and black hole solutions \cite{BertolamiCadoni}, its Newtonian limit \cite{newtlimit}, the existence of closed timelike curves \cite{closedtimecurve} and the modified Layzer-Irvine equation \cite{LayzerIrvine} (see Ref. \cite{review} for a review and Refs. \cite{puetzobukiorio} for other NMC gravity theories and their potential applications). Recently, the impact of NMC gravity on the spacetime metric surrounding a spherical central body was considered in Ref. \cite{BMP}, where the additional degree of freedom arising from a non-trivial $f(R)$ function is light, thus yielding a long-range additional force which requires considering the background cosmological setting; following the procedure set out in Ref. \cite{CSE} for $f(R)$ gravity, the Parameterized Post-Newtonian (PPN) parameter $\gamma $ was computed, provided that a set of requirements for $f(R)$ and the NMC function are obeyed. Then the compatibility has been assessed between a NMC model which accounts for the observed accelerated expansion of the Universe and Solar System experiments. Conversely, the case where the former is short-ranged enables one to neglect the background cosmological setting and derive the ensuing corrections to the gravitational potential \cite{CPM}, which are shown to be of the Yukawa-type --- as previously reported in Ref. \cite{NJ} for $f(R)$ gravity. In particular, it is found that the range of this Yukawa potential is given solely by $f(R)$, with the NMC affecting only its strength: this is a natural result, since the effect of the latter vanishes in vacuum, but affects the gravitational source. The purpose of this work is thus to further examine those findings, extending the formalism used in Ref. \cite{CPM} to include terms up to order $O\left(1/c^4\right)$ in the 00 component of the metric tensor. The nonlinear correction to the geometry part of the action is represented by a function $f^1(R)$, and the NMC is represented by a function $f^2(R)$ which multiplies the matter Lagrangian density. Both functions are assumed analytic at $R=0$ and the coefficients of the Taylor expansions around $R=0$ are considered as the parameters of the model. This work is organized as follows: In section II, the NMC model is presented and in section III its nonrelativistic limit is derived. Section IV computes the post-Newtonian and Yukawa corrections to the metric tensor by considering matter as a perfect fluid (without assumptions of symmetry). In Section V the metric around a static, spherically symmetric body is computed. Section VI addresses the ensuing Solar System constraints, namely through perturbations to perihelion precession. Recent observations of Mercury, including data from the Messenger spacecraft, are used to constrain the parameters of the model. Finally, conclusions are drawn in Section VII. | In this work we have computed the metric solutions for a NMC gravity model around a Minkowski background. It is shown that, up to order $O(1/c^4)$, the corrections depend on the $f^1(R)$ and $f^2(R)$ functions and cannot be expressed in terms of powers of $1/r$: indeed, it is found that the obtained solutions must be expressed in the PPNY approximation, as first proposed in Ref. \cite{CPM}. This opens up the possibility of addressing a wider class of physical situations with great accuracy. Furthermore, the results obtained in this work might be relevant for distinguishing between GR, $f(R)$ and non minimally coupled theories from the analysis of detailed observations data in the future. \appendix | 16 | 7 | 1607.03784 |
1607 | 1607.03267_arXiv.txt | We report the discovery of a microlensing planet OGLE-2012-BLG-0950Lb with the planet/host mass ratio of $q \simeq 2 \times 10^{-4}$. A long term distortion detected in both MOA and OGLE light curve can be explained by the microlens parallax due to the Earth's orbital motion around the Sun. Although the finite source effect is not detected, we obtain the lens flux by the high resolution Keck AO observation. Combining the microlens parallax and the lens flux reveal the nature of the lens: a planet with mass of $M_{\rm p} = 35^{+17}_{-9} M_{\oplus}$ is orbiting around a M-dwarf with mass of $M_{\rm host} = 0.56^{+0.12}_{-0.16} M_{\odot}$ with a planet-host projected separation of $r_{\perp} =2.7^{+0.6}_{-0.7}$ AU located at $D_{\rm L} = 3.0^{+0.8}_{-1.1}$ kpc from us. This is the first mass measurement from only microlens parallax and the lens flux without the finite source effect. In the coming space observation-era with {\it Spitzer}, {\it K2}, {\it Euclid}, and {\it WFIRST}, we expect many such events for which we will not be able to measure any finite source effect. This work demonstrates an ability of mass measurements in such events. | Gravitational microlensing is a technique by which planets can be detected without measurements of light from the host star \citep{mao91,gouloe92,gau12}. Microlensing can detect planets that are difficult to detect by other methods such as planetary systems in the Galactic Bulge (e.g., Batista et al. 2014), planets around late M-dwarfs or brown dwarfs \citep{ben08, sum16}, and even free floating planets not hosted by any stars \citep{sum11}. Compared to other techniques, microlensing is sensitive to Earth mass planets \citep{ben96} orbiting just outside of the snow line where the core accretion theory \citep{ida04} predicts is the most active planet formation region. Microlensing observations so far have revealed a population of planets beyond the snow line \citep{gou10, sum10, cas12, shv16, suz16}. \citet{suz16} finds a steeper slope with $dN/d\log {q} \sim q^{-0.9}$ and a break (and possible peak) in the mass ratio function at $q \sim 1.0 \times 10^{-4}$. We are capable of studying the distance distribution of planets in our Galaxy via microlensing. \citet{pen16} suggests the possibility of a lack of planets in the Galactic bulge. The detection of extra solar planets by gravitational microlensing presents a number of challenges. Firstly, gravitational microlensing is an extremely rare phenomenon with a probability of one per one million stars and a planetary deviation lasts for only hours or a few days. For these reasons, microlensing observations for exoplanets are conducted towards the Galactic bulge, the most crowded field in our Galaxy. Whereas hundreds of planets are detected by the radial velocity (RV) method \citep{but06,bon13} and thousands of planetary candidates are detected by the {\it Kepler} telescope \citep{bor10} to date, the microlensing method has been used to detect about 50 exoplanets so far. Several survey groups have been conducting high cadence survey observations using their telescopes with wide FOV cameras in different time zones. The Microlensing Observations in Astrophysics (MOA; Bond et al. 2001, Sumi et al. 2003) group uses the 2.2-deg$^2$ FOV MOA-cam3 \citep{sak08} CCD camera mounted on the 1.8 m MOA-II telescope at the Mt.\ John University Observatory in New Zealand and alerts the community about 600 microlensing events per year. The Optical Gravitational Lensing Experiment group (OGLE; Udalski 2003) upgraded their camera to the 1.4-deg$^2$ FOV OGLE-IV camera in 2010 \citep{uda15a} and discovered more than 2000 microlensing events per year in the last few years with the camera mounted on the 1.3 m Warsaw telescope at the Las Campanas Observatory, Chile. The Wise observatory group in Israel also conducts microlensing observations \citep{gor10}. In 2015, the Korean Microlensing Telescope Network (KMTNet; Kim et al. 2016) also started their survey observations. Now the equipment requirements for second-generation microlensing surveys \citep{gau09,gau12} are fulfilled and the number of planet detections is increasing over the next few years. Measuring the mass of a lens $M_{\rm L}$ and the distance to the lens system $D_{\rm L}$ is challenging. There are three observables in microlensing which can yield a mass-distance relation of the lens system: the angular Einstein radius $\theta_{\rm E}$, the microlens parallax $\pi_{\rm E}$ and the lens flux. The first two of these yield each mass-distance relation by combining the following relationship between them; \begin{equation} M_{\rm L} = \frac{\theta_{\rm E}}{\kappa \pi_{\rm E}} \end{equation} with the definitions of $\pi_{\rm E}$, $\pi_{\rm E} \equiv \pi_{\rm rel} / \theta_{\rm E}$, where $\kappa$ is a constant and $\pi_{\rm rel} \equiv {\rm AU} (1/D_{\rm L} - 1/D_{\rm S})$. One can calculate the mass and distance of the lens system if we can measure any two of these quantities. We can measure $\theta_{\rm E}$ in the following manner. Included in most models explaining planetary microlensing light curve data is the source star radius in units of $\theta_{\rm E}$: $\rho \equiv \theta_*/\theta_{\rm E}$. By estimating the angular radius of the source star, $\theta_*$, by an analysis of the source star's color and magnitude, and using our modeled value of $\rho$, we arrive at an estimate of $\theta_{\rm E}$. Microlens parallax can be observed only in relatively rare events and lens flux measurements need follow-up observations with high resolution imaging by an 8-m class telescope. Therefore only half of planetary events published so far are detected with lens mass measurements and masses of the other half planetary systems are just given their probability distributions by a Bayesian analysis (e.g., Beaulieu et al. 2006; Bennett et al.2014; Koshimoto et al.2014; Skowron et al 2015). In the microlensing planetary events published so far, there are events with the mass measurements from the angular Einstein radius and microlens parallax (e.g., Bennett et al. 2008; Gaudi et al. 2008; Muraki et al. 2011), from the angular Einstein radius and the lens flux (e.g., Bennett et al. 2006; Batista et al. 2015; Bennett et al. 2015), and from all three relations (e.g., Dong et al. 2009; Bennett et al. 2010; Beaulieu et al. 2016; Bennett et al. 2016), but events with mass measurement from only microlens parallax and the lens flux have not been published to date. This is simply because the angular Einstein radius is observed much commonly than microlens parallax as mentioned above. However, it has been possible to measure precise microlens parallax by observing simultaneously from space and ground thanks to the {\it Spitzer} microlensing campaign \citep{cal15,uda15b,yeeet15,zhu15}. Also, {\it K2} campaign 9 ({\it K2}C9), started in April 2016, has surveyed the Galactic bulge for three months to date. By combining {\it K2}C9 data and ground-based survey data, it is expected to measure microlens parallax for more than 120 events \citep{hen15}. These next generation space- and ground-based simultaneous observations for microlensing can measure microlens parallax almost regardless of the event timescale. Microlens parallax should become a more common observable rather than the angular Einstein radius in coming next generation, thus the mass measurement without the angular Einstein radius should be important \citep{yee15}. This paper reports an analysis of a microlensing planetary event OGLE-2012-BLG-0950, which is the first event where a mass measurement is possible from only the measurements of the microlens parallax and lens flux. The survey observations of this event are described in Section \ref{sec-obs}. Section \ref{sec-reduction} explains our data reduction procedure. Section \ref{sec-model} shows our modeling results. % We show the constraint on the angular Einstein radius by the source angular radius derived from the color and light curve modeling in Section \ref{sec-color}. In Section \ref{sec-KECK} we describe our Keck observations, the constraints on the excess flux and calculate the probability of the contamination to the excess flux. In Section \ref{sec-lens}, we derive the lens properties by combining microlens parallax and the lens flux. Finally, Section \ref{sec-disc} discusses and concludes the results of this work. | \label{sec-disc} We analyzed the microlensing event OGLE-2012-BLG-0950. A negative perturbation in the microlensing light curve consistent with a low-mass planet was detected \citep{abe13}. All the models we analyzed have a planetary mass ratio, $q \simeq 2 \times 10^{-4}$. We could not detect a significant finite source effect because the source did not cross any caustic, but we did detect a parallax signal. The parallax solutions indicate a Neptune/sub-Saturn mass planet with mass of $M_{\rm p} = 35^{+17}_{-9} M_{\oplus}$ around an M/K-dwarf host with mass of $M_{\rm host} = 0.56^{+0.12}_{-0.16} M_{\odot}$. We measured the lens mass by combining microlens parallax and the lens flux obtained by Keck AO observations. This is the first case in which the lens mass was measured using only microlens parallax and the lens flux. The planet orbits outside of the snow line of the host star and has a mass between that of Neptune and Saturn, $M_{\rm p} = 35_{-9}^{+17} M_{\oplus}$. Planets with this mass range (intermediate mass, hereafter) are predicted to be rare inside the snow line, but to be common like Neptune- or Saturn- mass planets outside the snow line according to the core accretion theory \citep{ida04, ida13}. A paucity of intermediate mass planets orbiting close to their metal-poor host stars is confirmed \citep{bea13}. On the other hand, the predicted relative abundance outside the snow line has not been confirmed yet. Figure \ref{fig-exo} shows the distribution of the exoplanets\footnote{http://exoplanet.eu} discovered so far. The solution of our parallax model is indicated as the purple filled circle located just around the valley of the bimodal mass distribution histogram on the left side of the figure. Note that this distribution is not corrected for detection efficiency. Only a few intermediate mass planets orbiting outside the snow line have been discovered by the RV and microlensing methods. The parallax model of this work could be the second such intermediate mass exoplanet with mass measurement, following OGLE-2012-BLG-0026Lb \citep{han13,bea16}. The mass of other intermediate mass planets are estimated by Bayesian analysis \citep{miy11,pol14,sko16b}. However the Bayesian estimates depend on the choice of prior \citep{ben14,sko15}. In a future space-based microlensing survey by {\it WFIRST} \citep{spe15} or {\it Euclid} \citep{pen13}, and in the survey of Campaign 9 of the {\it K2} Mission \citep{hen15} conducted from April 2016 to July 2016, or the {\it Spitzer} microlensing campaign from 2014 \citep{yeeet15}, it is important and easier to determine the lens mass for each event by combining microlens parallax and lens flux as pointed out by \citet{yee15} for the following reasons. First, space- and ground-based simultaneous observations are expected to obtain microlens parallax for a significant fraction of events regardless of number of the lens bodies, in contrast to the finite source effect which can be obtained only by observing the peak of high-mag event or caustic crossing. Second, for low-mass and nearby lenses, the mass-distance relations derived from flux and from $\theta_{\rm E}$ are partially degenerate (see Figure 2 in Yee 2015 or Figure 7 in Fukui et al. 2015) although we can obtain $\theta_{\rm E}$ by the measurement of astrometric microlensing effects with the precision of {\it WFIRST} \citep{gou14}. Third, the cases without detection of $\theta_{\rm E}$ like this event are expected to increase even for planetary events because a higher precision and higher cadence survey can detect more subtle planetary signals including cases without crossing caustics \citep{zhu14}. Finally, it is possible to measure the lens fluxes even after the events by follow-up observations with high angular resolution, and ultimately, {\it WFIRST} and {\it Euclid} can routinely measure the lens fluxes as part of the survey observations. Our analysis for the parallax model is the first demonstration of the mass measurement from only microlens parallax and the lens flux, and thus it has particular significance for the developing era of space-based microlensing. We acknowledge the following support: Work by N.K. is supported by JSPS KAKENHI Grant Number JP15J01676. The MOA project is supported by the grant JSPS25103508 and 23340064. N.J.R is a Royal Society of New Zealand Rutherford Discovery Fellow. OGLE Team thanks Profs.\ M.~Kubiak and G.~Pietrzy{\'n}ski, former members of the OGLE team, for their contribution to the collection of the OGLE photometric data over the past years. The OGLE project has received funding from the National Science Centre, Poland, grant MAESTRO 2014/14/A/ST9/00121 to AU. V.B. was supported by the CNES and the DIM ACAV, R\'egion \^lle-de-France. V.B., J.P.B., and J.B.M. acknowledge the support of PERSU Sorbonne Universit\'e, the Programme National de Plan\'etologie and the labex ESEP. \appendix | 16 | 7 | 1607.03267 |
1607 | 1607.01054_arXiv.txt | A debate has arisen concerning the fundamental nature of luminous blue variables (LBVs) and their role in stellar evolution. While Smith \& Tombleson proposed that their isolated environments indicate that LBVs must be largely the product of binary evolution, Humphreys et al.\ have recently expressed the view that the traditional single-star view still holds if one appropriately selects a subsample of LBVs. This paper finds the claim of Humphreys et al.\ to be quantitatively unjustified. A statistical test of ``candidate'' as opposed to ``confirmed'' LBVs shows no significant difference ($<$1$\sigma$) between their environments. Even if the sample is further subdivided as proposed, the three most luminous LBVs are spatially dispersed similar to late O-type dwarfs, which have much longer median lifetimes than expected for classical LBVs. The lower-luminosity LBVs have a distribution associated with red supergiants (RSGs), but these RSGs are dominated by stars of 10-15 $M_{\odot}$ initial mass, with much longer lifetimes than expected for those lower-luminosity LBVs. If one's view is restricted to the {\it highest-luminosity} LBVs, then the appropriate comparison is with {\it early} O-type stars that are their presumed progenitors; when this is done, it is clear that even the high-luminosity LBVs are more dispersed than expected. Humphreys et al.\ also suggest that velocities of LBVs support the single-star view, being inconsistent with runaways. A quantitative analysis of the radial velocity distribution of LBVS in M31 and M33 contradicts this; modest runway speeds expected from mass gainers in binary evolution are consistent with the observed velocities, although the data lack the precision to discriminate. | The eruptive mass loss exhibited by luminous blue variables (LBVs) is thought to be an important ingredient in stellar evolution (see, e.g., \citealt{so06,smith14}), but the way LBVs actually figure in this evolution and the physical mechanisms of their outbursts remains challenging to understand. Moreover, LBVs are thought to be related to some extragalactic non-supernova (SN) transients \citep{smith11,vdm12} and their mass loss is reminiscent of extreme pre-SN eruptions \citep{smith14}. In a recent study, Smith \& Tombleson (2015; ST15 hereafter) analyzed the projected spatial distribution of LBVs on the sky and found them to be surprisingly isolated from O-type stars. ST15 concluded that the results were inconsistent with expectations for the traditional picture of LBVs in single-star evolution (e.g., \citealt{hd94}), wherein LBVs are descended from very massive main sequence O-type stars, and where LBVs are the key agent that provides the required mass loss to drive them into the Wolf-Rayet (WR) phase. In particular, ST15 found that LBVs were more dispersed from O stars on the sky than WR stars, making it impossible for the observed population of LBVs to turn into the observed population of WR stars. ST15 concluded that many LBVs are likely to be the product of binary evolution, where stars are spun-up, chemically enriched, and rejuvinated by mass transfer, and possibly kicked by their companion's SN explosion. In this view, LBVs are evolved massive blue stragglers. Humphreys et al. (2016; H16 hereafter) present a contrasting viewpoint, arguing that environments of LBVs are instead consistent with the traditional view if one divides and culls the sample of LBVs in the way they prefer. They also claim that observed kinematics indicate that none of the confirmed LBVs are runaway stars. The discussion below critically examines these claims, since the role played by LBVs and their mass loss is fundamental to our understanding of massive star evolution and the origin of WR stars. Essentially, it is found that the claims made by H16 are not quantitatively justifiable based on the data, even if one permits the selective subdivision of the LBV sample as they envision. In some cases the quantitative implication of the data yields the opposite of their qualitative interpretation. This paper undertakes a critique of the claims made by H16. Before examining their analysis, Section 2 first corrects some errors and points out misconceptions that influence the data and expectations in the H16 paper. Section 3 concentrates on the claimed difference between confirmed and candidate LBVs (each corresponding to about half the sample in ST15), showing that their environments are indistinguishable with available data. Section 4 reasseses the analysis of H16; here we permit the subjective subdivision of the sample as preferred by H16, and we ignore small number statistics. Even with these accomodations, we find that the central conclusions found by H16 were not the correct conclusions indicated by their suggested division of the data. This is because if one wishes to split a cumulative distribution, one must also split the distribution of comparison objects in order to draw a meaningful conclusion. Finally, Section 5 provides a quantitative analysis of the claims by H16 regarding radial velocities, finding them to be invalid. Runaway LBVs are certainly allowed by the radial velocity data, although not clearly required. | In conclusion, a quantitative look at the data of LBVs and their environments shows that even if one adopts the selective criteria advocated by H16, the results do not support their claims. When only a few objects from the tail of a distribution are extracted, there remains no statistical power to discriminate between that sample and the remainder. Moreover, the analysis above shows that even if those selections are permitted, the interpretation arrived at by H16 was incorrect, because they did not consider the quantitative implications of the comparison stars. Namely, one finds that the most luminous confirmed LBVs have environments similar to late O-type stars, which have median ages about twice as long as the presumed ages of those LBVs in a single-star scenario. Similarly, the lower-luminosity confirmed LBVs have a distribution similar to RSGs, the bulk of which have initial masses (10$-$15 $M_{\odot}$), less than half that of the low-luminosity LBVs (25$-$40 $M_{\odot}$). This discrepancy rules out the traditional single-star view of LBVs, and requires instead that they are massive blue stragglers. A statistical test of the confirmed and candidate LBVs shows no statistical difference between their environments, contrary to the central claim that motivated H16's reanalysis. Thus, it is not clear that it is appropriate to separate them. Even if one does separate the confirmed and candidate LBVs, and if one further separates the low and high luminosity group of LBVs, the central results of ST15 remain the same -- that LBVs are more isolated from O-type stars than they should be in the traditional single star view of stellar evolution. Rejuvenation by mass transfer and mergers, and possibly runaway motion from a companion's SN, are required to explain their isolation. Available kinematics of LBVs do not argue against SN-induced runaways, mostly because the expected motion is slow and the precision of available data cannot clearly discriminate. It will be interesting to see if there are other indications that LBVs do not have anomalous motion; if they do not, then this will point to rejuvenation in binary evolution (i.e. massive blue stragglers) as the main explanation for their isolated environments. \smallskip\smallskip\smallskip\smallskip \noindent {\bf ACKNOWLEDGMENTS} \smallskip \footnotesize I thank an anonymous referee for helpful suggestions. Support was provided by NSF awards AST-1312221 and AST-1515559, and by the National Aeronautics and Space Administration (NASA) through HST grant AR-14316 from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS5-26555. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France. | 16 | 7 | 1607.01054 |
1607 | 1607.04995_arXiv.txt | We report on the first deep optical observations of two $\gamma$-ray pulsars, both among the very first discovered by the {\em Fermi} Gamma-ray Space Telescope. The two pulsars are the radio-loud PSR\, J1907+0602 in the TeV pulsar wind nebula (PWN) MGRO\, J1908+06 and the radio-quiet PSR\, J1809$-$2332 in the "Taz" radio/X-ray PWN. These pulsars are relatively young and energetic and have been both detected in the X-rays by \xmm, which makes them viable targets for optical observations. We observed the pulsar fields in the B and V bands with the Very Large Telescope (VLT) in June/July 2015 to search for their optical counterparts. Neither of the two pulsars has been detected down to $3\sigma$ limiting magnitudes of $m_{\rm v} \sim 26.9$ and $m_{\rm v} \sim 27.6$ for PSR\, J1907+0602 and PSR\, J1809$-$2332, respectively. We discuss these results in the framework of the multi-wavelength emission properties of pulsars. | Pulsars are rapidly rotating isolated neutron stars, powered by their rotational energy (see Kaspi \& Kramer 2016 for a recent review). Prevalently observed at radio wavelengths, they are also detected at X and $\gamma$-ray energies and in the optical band, where they are challenging targets owing to their intrinsic faintness (Mignani 2011). While in the 1980s/1990s pulsar searches in the optical were mainly driven by their X-ray detection, since the launch of the {\em Fermi} Gamma-ray Space Telescope in 2008 the wealth of pulsar $\gamma$-ray detections (see, Caraveo 2014 and Grenier \& Harding 2015 for recent reviews) have spurred their search both at X-ray and optical wavelengths. With over 200 $\gamma$-ray pulsars now identified by {\em Fermi}\footnote{See {\texttt https://confluence.slac.stanford.edu/display/GLAMCOG/} for a continually updated list.}, the number of those detected in the optical (or with at least a candidate optical counterpart) is still tiny (Abdo et al.\ 2013; Moran et al.\ 2013) owing to the paucity of sensitive optical observations (see, e.g. Mignani et al.\ 2016a for a summary). Recently, we detected a candidate optical counterpart to the middle-aged $\gamma$-ray pulsar PSR\, J1741$-$2054 (Mignani et al.\ 2016b), with the ESO Very Large Telescope (VLT). In the same run, we observed other two $\gamma$-ray pulsars discovered by {\em Fermi}, PSR\, J1907+0602 and PSR\, J1809$-$2332, as part of a dedicated pilot survey. The characteristics of these two pulsars are summarised in Table \ref{psr}. \begin{table} \begin{center} \caption{Coordinates, proper motion, position reference epoch (Kerr et al.\ 2015) for the two {\em Fermi} pulsars discussed in this work, together with their spin period P$_{\rm s}$, period derivative $\dot{P}_{\rm s}$, and inferred values of the characteristic age $\tau_c \equiv P_{\rm s}/ 2 \dot{P}_{\rm s}$, rotational energy loss $\dot{E}_{\rm rot}$ and surface dipolar magnetic field $B_{\rm s}$. The latter two values have been derived from the standard formulae $\dot{E}_{\rm rot} = 4 \times 10^{46} \dot{P}_{\rm s}/P_{\rm s}^{3}$ erg s$^{-1}$ and $B_{\rm s} = 3.2 \times 10^{19} \sqrt{P_{\rm s} \dot{P}_{\rm s}}$ G, derived by assuming for the neutron star a moment of inertia $I = 10^{45}$ g cm$^{2}$ (e.g., Kaspi \& Kramer 2016). The values have been obtained from the ATNF pulsar catalogue (Manchester et al.\ 2005).} \label{psr} \begin{tabular}{lll} \hline & PSR\, J1907+0602 & PSR\, J1809$-$2332 \\ \hline $\alpha$ (J2000) & $19^{\rm h} 07^{\rm m} 54\fs76$ (0\fs05) & $18^{\rm h} 09^{\rm m} 50\fs249$ (0\fs030) \\ $\delta$ (J2000) & $ +06^\circ 02\arcmin 14\farcs6$ (0\farcs7) & $-23^\circ 32\arcmin 22\farcs67$ (0\farcs10) \\ $\mu_{\alpha}$ (mas yr$^{-1}$) & - & $+12\pm8$ \\ $\mu_{\delta}$ (mas yr$^{-1}$) & - & $-24\pm6$ \\ Epoch (MJD) & 55555 & 55555 \\ P$_{\rm s}$ (s) & 0.106 & 0.146 \\ $\dot{P}_{\rm s}$ ($10^{-14}$ s s$^{-1}$) & 8.68 & 3.44 \\ $\tau_c$ (kyr) & 19.5 & 67.6 \\ $\dot{E}_{\rm rot}$ ($10^{36}$ erg s$^{-1}$) & 2.8& 0.43\\ $B_{\rm s}$ ($10^{12}$ G) & 3.08 & 2.27 \\ \hline \end{tabular} \end{center} \end{table} PSR\, J1907+0602 was discovered as a $\gamma$-ray pulsar during a blind search for pulsations in unidentified {\em Fermi}-LAT sources (Abdo et al.\ 2009a; 2010). These characteristics make PSR\, J1907+0602 quite similar to the slightly younger (11.2 kyr) Vela pulsar (Manchester et al.\ 2005), one of the historical $\gamma$-ray pulsars (e.g., Abdo et al.\ 2009b). Very faint radio pulsations from PSR\, J1907+0602 were detected with the Arecibo telescope at 1.5 GHz (Abdo et al.\ 2010; Ray et al.\ 2011), soon after its discovery as a $\gamma$-ray pulsar. The radio dispersion measure (DM=82.1$\pm$1.1 cm$^{-3}$ pc) puts the pulsar at a nominal distance of 3.2$\pm$0.6 kpc (Abdo et al.\ 2010), according to the model of the Galactic free electron density (Cordes \& Lazio 2002). This would make PSR\, J1907+0602 one of the faintest known radio pulsars, with a 1.4 GHz radio luminosity of 0.035 mJy kpc$^{2}$. No radio parallax measurement has been obtained for this pulsar. PSR\, J1907+0602 has been searched for but not detected in a pulsar survey at 34 MHz (Maan \& Aswathappa 2014). Both the DM distance and spin-down age suggest that PSR\, J1907+0602 was probably born at the centre of the supernova remnant (SNR) G40.5$-$0.5 (Abdo et al.\ 2010). After a preliminary detection by \chan\ (Abdo et al.\ 2010; Marelli et al.\ 2011), the pulsar has been observed by \xmm\ (Abdo et al.\ 2013). No X-ray pulsations have been detected yet. PSR\, J1907+0602 is likely associated with a pulsar wind nebula (PWN) detected at TeV energies by MILAGRO, HESS, VERITAS, and HAWC (Abdo et al.\ 2010; Abeysekara et al.\ 2016). PSR\, J1809$-$2332 is older than PSR\, J1907+0602, ideally half way between the young, Vela-like pulsars and the middle-aged ones (100 kyr--1 Myr). Like PSR\, J1907+0602, PSR\, J1809$-$2332 has been identified as a $\gamma$-ray pulsar by the {\em Fermi}-LAT during a blind search for pulsations from unidentified sources (Abdo et al.\ 2009a). The LAT source (3FGL\, J1809.8$-$2332; Acero et al.\ 2015) associated with the pulsar is identified with the {\em Compton}/EGRET source 3EG\, J1809$-$2338 (Hartmann et al.\ 1999). The latter was found to be spatially coincident with the dark nebula Lynds 227 and a PWN candidate detected in the X rays by {\em ASCA} (Oka et al.\ 1999). The PWN, later dubbed the "Tasmanian devil" ("Taz" for short), was then observed by \chan\ (Braje et al.\ 2002; Roberts \& Brogan 2008), which also resolved the point-like source CXOU\, J180950.2$-$233223 that Abdo et al.\ (2009) identified with the PSR\, J1809$-$2332 X-ray counterpart. The "Taz" PWN is also detected in radio and located within the shell SNR G7.5$-$1.7 (Roberts \& Brogan 2008). PSR\, J1809$-$2332 has been also observed by \xmm\ (Marelli et al.\ 2011) but X-ray pulsations have not been detected yet. The pulsar was searched for radio emission (Ray et al.\ 2011) but it was not detected down to a flux limit of 26 $\mu$Jy. A pulsar search at 34 MHz (Maan \& Aswathappa 2014) also resulted in a non detection. Since PSR\, J1809$-$2332 is not detected in radio there is no direct measurement of its distance. This is estimated to be 1.7$\pm$1.0 kpc from the distance to the dark nebula Lynds 227, which Oka et al.\ (1999) associated with the "Taz" PWN. Using \chan, Van Etten et al.\ (2012) measured a proper motion for PSR\, J1809$-$2332, which confirms the association with the SNR G7.5$-$1.7, as proposed by Roberts \& Brogan (2008). The optical emission of young pulsars ($\tau_{\rm c} \la 0.1$ Myr) is ascribed to synchrotron emission from energetic electrons in the pulsar magnetosphere (e.g., Pacini \& Salvati 1983) and the spectrum is characterised by a power-law (PL). For older pulsars, the thermal emission from the cooling neutron star surface also contributes to the optical emission and the spectrum is the combination of both a PL and a Rayleigh-Jeans (see, e.g. Mignani 2011 for a review). Nothing is known about the optical emission properties of PSR\, J1907+0602 and PSR\, J1809$-$2332. Recently, Brownsberger \& Romani (2014) carried out observations of the two pulsar fields in H$_{\alpha}$ with the 4.2m SOAR telescope to search for bow-shock nebulae. No wide-band imaging observations of the pulsar fields have been performed, though. Here, we report the results of our VLT observations of PSR\, J1907+0602 and PSR\, J1809$-$2332, the first carried out with a 10m-class telescope. In Sectn.\ 2 we describe the observations and data analysis, whereas we present and discuss the results in Sectn.\ 3 and 4, respectively. \begin{figure*} \centering \begin{tabular}{cc} \subfloat[PSR\, J1809-2332]{\includegraphics[width=8cm,bb=0 0 796 791,clip=]{psrj1809_fc_new.eps}} \subfloat[PSR\, J1907+0602]{\includegraphics[width=8cm,bb=0 0 796 791,clip=]{psrj1907_fc_new.eps}} \end{tabular} \caption{\label{fc} $10\arcsec \times 10\arcsec$ VLT/FORS2 $v_{\rm HIGH}$-band images of the pulsar fields. The pulsar positions determined by \chan\ (Kerr et al.\ 2015) are marked by the ellipses. The size of the ellipses accounts for statistical uncertainties only and not for the systematic uncertainty associated with the astrometry calibration of the VLT images ($\sim$0\farcs1). The PSR\, J1809$-$2332 position has been corrected for the pulsar proper motion (Van Etten et al.\ 2012). Quite surprisingly, in the case of PSR\, J1907+0602 no stars are detected within the entire $10\arcsec \times 10\arcsec$ region around the reference pulsar position. } \end{figure*} | We observed the two $\gamma$-ray pulsars PSR\, J1907+0602 and PSR\, J1809$-$2332 with the VLT. With them, there are now six isolated $\gamma$-ray pulsars discovered by {\em Fermi}, for which the VLT obtained the first deep optical observations. The others are: PSR\, J1357$-$6429 (Mignani et al.\ 2011), PSR\, J1028$-$5819 (Mignani et al.\ 2012), PSR\, J1048$-$5832 (Razzano et al.\ 2013; Danilenko et al.\ 2013), and PSR\, J1741$-$2054 (Mignani et al.\ 2016b). Of them, a candidate optical counterpart has been found for PSR\, J1741$-$2054 (Mignani et al.\ 2016b), whereas for PSR\, J1357$-$6429 a candidate counterpart might have been identified in the near infrared (Zyuzin et al.\ 2016). Thus, one can certainly say that the VLT leads the optical follow-up of isolated $\gamma$-ray pulsars, at least in the southern hemisphere. The VLT yielded the only optical counterpart identifications obtained so far for the new southern $\gamma$-ray pulsars discovered by {\em Fermi}\footnote{The optical counterpart to PSR\, B0540$-$69 in the Large Magellanic Cloud (Caraveo et al.\ 1992) was identified long before its detection as a $\gamma$-ray pulsar (Ackermann et al.\ 2015).}. In the northern hemisphere, a candidate optical counterpart to the $\gamma$-ray pulsar PSR\, J0205+6449 was discovered in archival data taken with the Gemini telescope (Moran et al.\ 2013). Unfortunately, in this case we were not successful in detecting the optical counterparts to the target pulsars, and we could only set $3\sigma$ upper limits to their optical brightness of $m_{\rm v} \sim 26.9$ and $m_{\rm v} \sim 27.6$ for PSR\, J1907+0602 and PSR\, J1809$-$2332, respectively. These are the deepest constraints on the optical fluxes ever obtained for these two pulsars, which had never been observed with 10m-class telescopes so far. | 16 | 7 | 1607.04995 |
1607 | 1607.03051_arXiv.txt | Current detectors for Very-High-Energy $\gamma$-ray astrophysics are either pointing instruments with a small field of view (Cherenkov telescopes), or large field-of-view instruments with relatively large energy thresholds (extensive air shower detectors). In this article, we propose a new hybrid extensive air shower detector sensitive in an energy region starting from about 100 GeV. The detector combines a small water-Cherenkov detector, able to provide a calorimetric measurement of shower particles at ground, with resistive plate chambers which contribute significantly to the accurate shower geometry reconstruction. A full simulation of this detector concept shows that it is able to reach better sensitivity than any previous gamma-ray wide field-of-view experiment in the sub-TeV energy region. It is expected to detect with a $5\sigma$ significance a source fainter than the Crab Nebula in one year at $100\,$GeV and, above $1\,$TeV a source as faint as 10\% of it. As such, this instrument is suited to detect transient phenomena making it a very powerful tool to trigger observations of variable sources and to detect transients coupled to gravitational waves and gamma-ray bursts. | High energy gamma rays are important probes of extreme, non thermal, events taking place in the universe. Being neutral, they can cover large distances without being deflected by galactic and extragalactic magnetic fields. This feature enables the direct study of their emission sources. The gamma emission is also connected to the acceleration of charged cosmic rays and to the production of cosmic neutrinos. Gamma-rays can also signal the existence of new physics at the fundamental scales, namely by the annihilation or decay of new types of particles, as it is the case for dark matter particles in many models. This motivation, associated to the advances of technology, has promoted a vigorous program of study of high energy gamma rays, with important scientific results (see \cite{Pimenta:2015ab,Degrange:2016ab,Funk:2015ab,Hillas:2013am} for a summary of the main achievements). The detected sources of cosmic gamma-rays above 30 MeV are concentrated around the disk of the Milky Way; in addition there is a set of extragalactic emitters. About 3000 sources emitting above 30 MeV were discovered, mostly by the Large Area Telescope (LAT) detector \cite{Acero:2015hja} onboard the $Fermi$ satellite, and some 200 of them emit as well above 30 GeV \cite{Ackermann:2016abc} (see Fig. \ref{fig:gammamap}) - the region which is labeled the Very High Energy (VHE) region. Our Galaxy hosts about half of the VHE gamma-ray emitters~\cite{Abeysekara:2016ab} and most of them are associated to supernova remnants of various classes (shell supernova remnants, pulsar-wind nebulae, etc.). The remaining emitters are extragalactic. The angular resolution of current detectors, which is slightly better than 0.1$^\circ$, does not allow to assign the identified extragalactic emitters to any particular region in the host galaxies; however, there is some consensus that the signals detected from the Earth must originate in the proximity of supermassive black holes at the center of the galaxies \cite{Fuhrmann:2014ab}. \begin{figure}[ht] \begin{center} \includegraphics[width=0.49\textwidth]{figs/tevcatjune13.png} \end{center} \caption {\label{fig:gammamap}Sources of VHE emission displayed in galactic coordinates. The background represents the high-energy gamma-rays detected by $Fermi$-LAT. The clear blue (dominant on the left side) area corresponds to the visible region within 30$^\circ$ of the zenith from a detector at a latitude of 22 degrees in the Northern hemisphere while the pink (dominant on the right side) shows the corresponding region in the Southern hemisphere. From http://tevcat.uchicago.edu/, June 2016.} \end{figure} Still, many problems remain open, of which we may mention: \begin{itemize} \item {\it The origin of cosmic rays} -- supernova remnants (SNRs) are thought to be the sites for the acceleration of protons up to few PeV. However, the mechanism of acceleration of particles to energies of that order is still to be established experimentally. The study of the photon yield from Galactic sources for energies larger than 100~GeV and all the way up to PeV, might solve the problem (see for example \cite{Aharonian:2016ab}). Actually, photons, which come from $\pi^0$ decay, correspond to hadronic cascades initiated at energies at least an order of magnitude larger. \item {\it The propagation of gamma-rays} -- tells us about their interaction with the cosmic background radiation and is a probe to cosmology themes. \item {\it New physics} -- the ultimate nature of matter and of physics beyond the Standard Model, dark matter or new particles in general, the energy density of the vacuum or even quantum gravity may leave imprints in the spectrum of VHE gamma rays. High-energy gamma-ray astrophysics is sensitive to energy scales important for particle physics. For instance, cold dark matter is expected to be found in the 100 GeV scale; supersymmetric particles could appear at the TeV scale; and the Planck scale (an energy $\sim 10^{19}$~GeV, corresponding to a mass $\sqrt{\hbar c/G}$) could be probed indirectly (for discussion see for example \cite{Pimenta:2015ab}). \item {\it Transients} -- Many VHE sources are characterised by variability, and it is important to be able to detect and measure the corresponding flares. Such flares have a duration that can go from few seconds -- like for the short gamma-ray bursts or the expected counterparts of gravitational waves -- to minutes for the long gamma-ray bursts, to minutes, hours or even days for the accretion flares of blazars (see e.g. \cite{loeb} and references therein). \end{itemize} The layout of this paper is the following. In the Introduction we have briefly presented the field of very-high-energy gamma-ray astrophysics. In Section 2 we outline the characteristics of the existing detectors, and we explain why a large field-of view detector is needed. In Section~\ref{sec:project} we make a case study with a possible design for such a detector. In Section~\ref{sec:data} we describe the characteristics of the gamma-ray signal and of its background, and the various Monte Carlo samples used in the analysis of the performance of the proposed detector. In Section~\ref{sec:results} we evaluate the performance of such a new detector using the simulation. We conclude the paper with final remarks and a summary in Section~\ref{sec:concl}. | \label{sec:concl} We have proposed a novel, large field-of-view, hybrid, extensive air shower detector sensitive to gamma rays of energies starting from 100 GeV (or even less for special high flux transients), based on individual units made (from top to bottom) by: \begin{itemize} \item a thin slab of lead, to allow conversion of secondary photons in showers into electron-positron pairs; \item a position-sensitive detector like a RPC; \item a water Cherenkov detector (or any other electromagnetic calorimeter with some muon identification capabilities). \end{itemize} We have shown that such a detector, if deployed on a surface of some 20~000 m$^2$, can reach a sensitivity that can provide the missing link between $Fermi$ and HAWC, between 100 and 350 GeV. Such detector would be able to detect in one year with a $5\sigma$ significance a source as faint as the Crab Nebula at $100\,$GeV. Above $1\,$TeV it could reach sensitivities better than 10\% of the Crab Nebula flux. The instrument is able to survey half of the sky, with enhanced ability to detect transient phenomena making it a very powerful tool to trigger observations of variable sources, in particular flares by active galactic nuclei, and to detect transients coupled to gravitational waves and gamma-ray bursts. An external sparse array of units could be easily achieved, due to the modular nature of the detector. This would not only allow to extend the energy range but also would further improve the sensitivity of the core array at lower energies. The presented hybrid detector concept is currently the baseline design of the core array of the LATTES project, foreseen as a new EAS gamma-ray detector in the Southern hemisphere, complementary to the planned Cherenkov Telescope Array. | 16 | 7 | 1607.03051 |
1607 | 1607.03321_arXiv.txt | RR Lyrae stars are old, usually metal-poor, population II stars, currently evolving on the horizontal branch, and crossing the classical instability strip. Most of them, the RRab stars, pulsate in the fundamental mode and have well recognisable light curves, with a very short and steep ascending branch and a long descending branch that often, but not always, exhibits a strong bump feature before reaching minimum light. First-overtone pulsators, members the RRc subclass, display a comparatively more sinusoidal light curve that is usually less asymmetric, and often features a notable depression, or hump, before maximum light. (Quasi-continuous observations of space photometric missions provided numerous textbook examples of RR Lyrae light curves: we refer the reader to the works by Benk\H{o} et~al.\ 2010, 2014, Moskalik et~al.\ 2015, Nemec et~al.\ 2011, Szab\'o et~al.\ 2014, and references therein.) The light curve shape of RR Lyrae stars strongly depends on the physical parameters of the star, therefore photometric data can be exploited to derive those properties. Jurcsik \& Kovacs (1996, JK96 hereafter) derived a formula to calculate the [Fe/H] indices of RRab stars from the pulsation period and the $\phi_{31}$ Fourier coefficients, defined as the $\phi_{31} = \phi_3-3\,\phi_1$ relative phase difference by Simon \& Lee (1981). A similar relation was determined for RRc stars by Morgan et~al.\ (2007), and later modified by Nemec et~al.\ (2013), who also provided an updated formula for RRab stars observed in the passband of the \textit{Kepler} space telescope. We note that the accuracy of the JK96 formula is limited for very low- and high-metallicity stars (see, e.g., Nemec et~al.\ 2013). While most RR Lyrae stars are metal-poor, with negative [Fe/H] indices (see, e.g.\ Feast et~al.\ 2008), some may reach metallicities close to solar, such as AV Peg (--0.08), RW TrA (--0.13) from the HIPPARCOS sample (Feast et~al.\ 2008) or V839 Cyg ($-0.05\pm0.14$) and V784 Cyg ($-0.05\pm0.10$) in the original \textit{Kepler} sample (Nemec et~al.\ 2013). Another interesting example is the short-period, high-metallicity star TV Lib (with ($P$ = 0.26962~d, and [Fe/H] = --0.17) that nevertheless displays very characteristic RRab light variations (Clube et~al.\ 1969, Kov\'acs 2005). During the target search for the K2 mission of the \textit{Kepler} space telescope (Howell et~al.\ 2014; Plachy et~al.\ 2016) we encountered one star, V397 Gem, whose photometric metallicity appeared to be extremely high, [Fe/H] = 0.42, based on the NSVS light curve (Northern Sky Variability Survey, Hoffmann et~al.\ 2009) and the original formula of JK96. V397 Gem also displayed a very short pulsation period of $P = 0.294$ d, considering that it was classified as an RRab star by Wils et~al.\ (2006). After a literature search, we found three more examples in the measurements of the All-Sky Automated Sky Survey (ASAS), based on the work of Szczygie\l{} et~al.\ (2009). All three stars were classified as RRab ones, have short periods (between 0.3--0.35 d) and [Fe/H] indices consistently above +0.5, based on methods of both JK96 and Sandage (2004). The folded light curves of the four stars are displayed in Fig.\ 1. The apparently extremely high metallicity suggests that these stars cannot be RRab stars, and the calculated values are simply erroneous. Alternate possibilities include high-amplitude $\delta$ Scuti stars; their rare, high-mass variants, the AC And-type pulsators; and RRc stars with unusually asymmetric, sawtooth-like light curves that mislead the (semi-)automated classification schemes of the surveys. In this paper we examine whether these objects could be misclassified RRc stars. \IBVSfig{11cm}{6175-f1.eps}{Folded light curves of the four RR Lyrae stars. Data were obtained by the NSVS survey for V397 Gem, and by the ASAS survey for the other three stars.} \IBVSfigKey{6175-f1.eps}{MT Tel; V397 Gem; ASAS J075127-4136.3; ASAS J091803-3022.6}{light curve} | 16 | 7 | 1607.03321 |
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1607 | 1607.03117_arXiv.txt | The pressure exerted by massive stars' radiation fields is an important mechanism regulating their formation. Detailed simulation of massive star formation therefore requires an accurate treatment of radiation. However, all published simulations have either used a diffusion approximation of limited validity; have only been able to simulate a single star fixed in space, thereby suppressing potentially-important instabilities; or did not provide adequate resolution at locations where instabilities may develop. To remedy this we have developed a new, highly accurate radiation algorithm that properly treats the absorption of the direct radiation field from stars and the re-emission and processing by interstellar dust. We use our new tool to perform three-dimensional radiation-hydrodynamic simulations of the collapse of massive pre-stellar cores with laminar and turbulent initial conditions and properly resolve regions where we expect instabilities to grow. We find that mass is channeled to the stellar system via gravitational and Rayleigh-Taylor (RT) instabilities, in agreement with previous results using stars capable of moving, but in disagreement with methods where the star is held fixed or with simulations that do not adequately resolve the development of RT instabilities. For laminar initial conditions, proper treatment of the direct radiation field produces later onset of instability, but does not suppress it entirely provided the edges of radiation-dominated bubbles are adequately resolved. Instabilities arise immediately for turbulent pre-stellar cores because the initial turbulence seeds the instabilities. Our results suggest that RT features are significant and should be present around accreting massive stars throughout their formation.\\ | \label{sec:intro} Massive stars live fast and die young. They are the major contributors to heavy element production in the Universe through their explosive deaths enriching the interstellar medium (ISM). Massive stars are rare, representing only $\sim 1\%$ of the stellar population by number, yet they dominate the energy budget in the Milky Way and other star-forming galaxies because of their strong radiation fields, stellar winds, and supernova explosions. This stellar feedback -- the injection of energy and momentum by stars into the ISM -- limits their masses thereby affecting nuclear yields, slows down nearby star formation, and affects galaxy evolution. Recent studies suggest that the pressure exerted by massive stars' radiation fields may be the dominant feedback mechanism during their formation \citep{Krumholz2009a, Kuiper2011a, Kuiper2012a, Klassen2016a}. Massive stars have short Kelvin-Helmholtz timescales (the time required for a star to radiate away its gravitational binding energy) and contract to the main-sequence while they are accreting \citep{Palla91a, Palla92a, Behrend01a, Hosokawa2009a}. Therefore they attain their main sequence luminosities while they are still actively accreting and the radiation pressure associated with their high luminosities can oppose gravity and halt accretion \citep{Larson71a, Yorke79a, Yorke95b, Wolfire86a, Wolfire87a, Yorke1999a}. The relative importance of the radiative force ($f_{\rm rad}$) and the gravitational force ($f_{\rm grav}$) can be described in terms of the Eddington ratio, $f_{\rm edd} = f_{\rm rad}/f_{\rm grav}$, which simplifies to \begin{equation} \label{eqn:feddh} f_{\rm edd} = 7.7 \times 10^{-5} \left( 1 + f_{\rm trap} \right) \left( \frac{L_{\rm \star}}{M_{\rm \star}} \right)_{\rm \odot} \left( \frac{\Sigma}{1 \; \rm{g \; cm^{-2}}} \right)^{-1} \end{equation} where $\Sigma$ is the surface density of the optically thick infalling material and $\left(L_{\rm \star}/M_{\rm \star}\right)_{\rm \odot}$ is the stellar light-to-mass ratio in solar units. The factor $\left( 1 + f_{\rm trap} \right)$ included in $f_{\rm rad}$ denotes the combined contribution from the direct radiation pressure associated with the first absorption of the stellar radiation field and the reprocessed thermal, diffuse radiation pressure associated with the re-emission by interstellar dust, respectively. Here $f_{\rm trap}$ denotes the trapping factor at which the radiation field is enhanced by the subsequent absorption and re-emission by interstellar dust. For spherically symmetric accretion, Equation (\ref{eqn:feddh}) exceeds unity for stars with masses above $\sim 15-20 \; M_{\rm \odot}$ \citep{Pollack1994a,Krumholz2009a}. If accretion onto the star were isotropic then stars with masses in excess of this limit should not form, a problem commonly known as ``the radiation pressure barrier problem." However, recent studies suggest that massive stars with initial masses well in excess of $150 \; M_{\rm \odot}$ exist and can have a dramatic impact on their environments \citep{Crowther2010a, Crowther2016a}. Given the existence of massive stars, a number of solutions to the radiation pressure problem have been proposed in the literature. \citet{Nakano89a} and \citet{Jijina96a} present analytic models suggesting that accretion through a disk could circumvent the radiation pressure barrier, while \citet{McKee2003a} suggest that high accretion rates could provide sufficient ram pressure even in spherical symmetry. \citet{Krumholz2005c} showed that escape of radiation through outflow channels could ease the radiation pressure problem. Numerical simulations within the last several decades generally support these hypotheses. Most of these simulations model the collapse of isolated, slowly rotating, and initially laminar pre-stellar massive cores \citep{Yorke1999a, Yorke2002a, Krumholz2009a, Kuiper2011a, Kuiper2012a, Klassen2016a}. In these idealized simulations, the radiation pressure barrier is circumvented by the formation of an optically thick accretion disk that surrounds the massive star. With this anisotropy, the radiative flux easily escapes along the polar directions of the star, launching radiation pressure dominated bubbles both above and below the star. This ``flashlight" effect allows material to be funneled to the star by the accretion disk and gravitational instabilities present in the disk can enhance the accretion rate onto the star \citep{Yorke2002a, Krumholz2009a, Kuiper2011a, Kuiper2012a, Klassen2016a}. Whether material is supplied to the star via disk accretion alone has been heavily debated in the literature \citep{Krumholz2009a, Kuiper2011a, Kuiper2012a, Klassen2016a}. \citet{Krumholz2009a} performed the first adaptive mesh refinement (AMR) 3D radiation-hydrodynamic simulation of the formation of a massive stellar system and found that the dense shells that surround the radiation pressure dominated bubbles become radiative Rayleigh-Taylor (RT) unstable. In this configuration, the dense shells that surround the rarefied radiation pressure dominated bubbles develop perturbations at the interface that grow exponentially, leading to ``fingers" in the heavier fluid (the accreting gas) that sink into the lighter, more buoyant fluid (represented by the radiation field; \citet{Jacquet2011a}). These RT ``fingers" can reach the star-disk system if they are not pushed back by radiation pressure, and deliver a significant amount of mass to the accretion disk that can then be incorporated into the star. The presence of these instabilities can allow stars to grow beyond their Eddington limit but their development and growth is sensitive to how the radiation pressure is treated. \citet{Krumholz2009a} only included the dust-reprocessed radiation pressure, which was modeled with the gray flux limited diffusion (FLD) approximation, and assumed that the stellar radiation energy was depositied within the vicinity of the star, which underestimated the true radiation pressure. If the radiation pressure, especially the component of the radiative force that is anti-parallel to the gravitational force, is underestimated then the gas is less likely to be pushed away by radiation. Furthermore, an anisotropic radiation field can lead to density perturbations in the dense shells of the radiation pressure dominated bubbles that can then amplify and become RT unstable. These instabilities can grow and deliver material to the star-disk system. To better represent the true radiation field in massive star formation simulations \citet{Kuiper2010a} developed a hybrid radiation algorithm that included a multi-frequency raytracer, in which a series of rays travel radially away from the star and transfer energy and momentum to the absorbing dust, coupled to gray FLD to model the diffuse dust-reprocessed radiation field. With this method, \citet{Kuiper2011a, Kuiper2012a} performed a series of 3D simulations of the formation of massive stars from the collapse of laminar pre-stellar cores on a non-adaptive spherical non-uniform grid with resolution increasing logarithmically towards the center. The authors find that the star is fed through disk accretion only and that the radiation pressure dominated bubbles do not become RT unstable. They conclude that inclusion of the direct radiation pressure is responsible for maintaining stability of the expanding bubble shells. The work of \citet{Krumholz2009a} and \citet{Kuiper2011a, Kuiper2012a} both have their advantages and disadvantages. AMR simulations with a general Cartesian geometry, such as the simulation presented in \citet{Krumholz2009a}, can handle an arbitrary number of moving stars. The resulting gravitational interaction of the massive star with its accretion disk can induce gravitational instabilities leading to disk fragmentation. In addition, movement of the massive star within the accretion disk can lead to shielding of the stellar radiation field resulting in a greater asymmetry in the direct radiation pressure, potentially seeding RT instabilities. One key advantage in AMR simulations, as compared to a non-adaptive grid, is that instabilities that may develop in the dense bubble shells can be resolved dynamically throughout the bubble evolution. In classical RT theory, the smallest perturbations grow fastest in the linear regime and these perturbations can only grow if they are resolved. The bubble shells in the work of \citet{Krumholz2009a} are resolved to the finest level, likely allowing for small RT instabilities to grow large enough to deliver material to the star-disk system. In contrast, the bubble shells in the work of \citet{Kuiper2011a, Kuiper2012a} are poorly resolved because they use a non-adaptive spherical grid. Furthermore, the star is artificially held at the origin of the grid, thereby suppressing potentially-important instabilities that could seed RT instabilities. However, these simulations included a much better treatment of the radiation field by incorporating a multi-frequency raytracer to model the direct radiation field. In such a geometry raytracing becomes trivial because the rays travel radially from the non-moving star, but this geometry can not support additional stars or disk asymmetries induced by stellar movement. Hence, the next generation of massive star formation simulations must include the advantages of both methods to better understand how massive stars can overcome the Eddington limit by including hybrid radiative transfer on adaptive grids. The question of whether RT instability is important for massive star formation has been muddied further by studies of radiation pressure-driven instabilities in the context of galactic winds. \citet{Krumholz2012b, Krumholz2013a} study the ability of radiation to drive galactic winds using the same FLD methods as \citet{Krumholz2009a}, and find that RT instabilities arise and prevent the onset of winds entirely. \citet{Rosdahl15a} reach the same conclusion using an M1 closure to treat the radiation. \citet{Davis2014a}, using a variable Eddington tensor method on a fixed grid, and \citet{Tsang15a}, using implicit Monte Carlo, concur that RT instability occurs, but find that it does not prevent a wind from being launched, contrary to the results of \citeauthor{Krumholz2012b} and \citeauthor{Rosdahl15a}. Moreover, none of these calculations included a treatment of the direct radiation field. The conflicting results discussed thus far have motivated the implementation of a new generation of hybrid radiation solvers in AMR simulation codes. Both \citet{Klassen2014a} and \citet{Rosen2016a} developed novel hybrid radiation schemes in the \flash\ and \orion\ AMR simulation codes, respectively. Both implementations model the direct radiation field with a raytracer while the diffuse component is handled by a FLD solver, and can be used with an arbitrary number of moving stars. The raytracer employed in the Hybrid Adaptive Ray-Moment Method (\harm) algorithm developed by \citet{Rosen2016a} uses the method of long characteristics, which traces rays on a cell by cell basis thus providing maximum possible accuracy. Their method is adaptive, in which rays are allowed to split as they travel away from their source, greatly reducing the computational cost; and is capable of representing multi-frequency stellar irradiation \citep{Abel2002a, Wise2011a, Rosen2016a}. The multi-frequency treatment is ideal for stars since they have color temperatures much higher than the absorbing medium. The raytracer employed in \citet{Klassen2014a} models only single frequency irradiation and uses hybrid characteristics, which is a combination of long characteristics within individual grids and short characteristics between grids (i.e., in which only neighboring grid cells are used to interpolate incoming intensities; \citet{Rijkhorst2006a}). The method of short characteristics is typically faster but more diffusive than long characteristics. Because of this limitation the long characteristics method employed in \citet{Rosen2016a} has been highly optimized. To revisit the problem of massive star formation and whether or not mass is delivered to the star via RT instabilities, \citet{Klassen2016a} simulated the collapse of initially laminar pre-stellar cores with the new hybrid radiation algorithm presented in \citet{Klassen2014a}. Like the work of \citet{Kuiper2011a, Kuiper2012a} they find that their radiation pressure dominated bubbles remain stable and that the massive star is fed by disk accretion alone. However, the authors employ poor refinement criteria in their simulations, which results in the bubble shells being poorly resolved, potentially suppressing RT instabilities that are not resolved. To address this, we perform similar simulations of the collapse of a laminar massive pre-stellar core in which we choose to resolve the bubble shells, like that of \citet{Krumholz2009a}, and use the \harm\ hybrid radiation algorithm to determine if RT instabilities are a real effect or if the direct radiation pressure inhibits their growth. As we will show, the development of RT instabilities is resolution dependent and therefore we find that authors can arrive at conflicting results if the bubble shells are not properly refined. The simulations discussed thus far were highly idealized. To date only the collapse of initially laminar massive pre-stellar cores have been studied numerically with a detailed treatment of the direct and diffuse radiation fields, yet observations of star forming regions show that star-forming cores are turbulent \citep{Tatematsu2008a,Sanchez-Monge2013a}. In such a configuration, the initial turbulence should act as seeds for RT instabilities. Furthermore, the asymmetric gas distribution in turbulent cores can yield low-density channels where radiation can easily escape, even in the absence of channels cut by outflows. The purpose of this paper is to study how radiation pressure affects the formation of massive stars via direct numerical simulation. For this work, we use the new highly accurate \harm\ algorithm described in \citet{Rosen2016a}, which treats the direct radiation field from stars and the indirect radiation field associated with the re-emission and processing by interstellar dust. In this work, we simulate the collapse of both initially laminar and turbulent pre-stellar cores to determine how massive stars attain their mass. For the laminar cores, we also examine how resolution and treatment of radiation pressure can affect the onset of RT instabilities. We simulate the collapse of an initially turbulent core to model a more realistic setup of how massive stars form to show that RT instabilities are a common occurrence in their formation. The simulations presented in this work are still highly idealized since we do not include magnetic fields or outflows. This paper is organized as follows: we describe our numerical methodology and simulation design in Section \ref{sec:numeth}, we present and discuss our results in Sections \ref{sec:results} and \ref{sec:discussion}, respectively, and conclude in Section \ref{sec:conclusion}. \begin{figure} \centerline{\includegraphics[trim=0.2cm 0.2cm 0.2cm 0.2cm,clip,width=0.95\columnwidth]{binOpacity10bins.pdf}} \centering \caption{ Specific frequency dependent dust opacities (per gram of dust) from \citet{Weingartner2001a} for their $R_{\rm v}=5.5$ extinction curve (teal line) with black body weighted binned opacities (pink diamonds) over-plotted for ten frequency bins used in the simulations presented in this work. } \label{fig:opacity} \end{figure} | \label{sec:conclusion} In this paper, we have used our powerful new hybrid radiation transfer tool, \harm, in a suite of radiation hydrodynamic simulations that followed the collapse of initially laminar and turbulent massive pre-stellar cores to study the formation of massive stars. \harm\ uses a multi-frequency adaptive long-characteristics ray tracing scheme to capture the first absorption of the direct radiation from stars by the intervening interstellar dust and molecular gas, and flux limited diffusion to model the diffuse radiation field associated with the subsequent re-emission by interstellar dust. Our method is highly optimized and can run efficiently on hundreds of processors, works on adaptive grids, can be coupled to any moment method, and can be used for an arbitrary number of moving stars \citep{Rosen2016a}. The primary goal of our work is to determine how massive stars attain their mass when radiation pressure is the only feedback mechanism considered (i.e., in the absence of magnetic fields, outflows, and photoionization). Do massive stars obtain their mass through disk accretion alone? Or do radiative Rayleigh Taylor instabilities that develop in the radiation pressure dominated bubble shells, which are launched by the stars' intense radiation fields, deliver material directly via collapse onto the stars or star-disk systems? Or is it a combination of both of these processes? For initially laminar cores, we find that the majority of mass delivered to the massive star is due to disk accretion, but that RT instabilities are responsible for bringing material onto the disk before it is subsequently incorporated into the star. For initially weakly turbulent cores, in contrast, we find that dense filaments and RT unstable material supply most of the mass to the accreting massive star directly, without mediation by an extended disk (i.e., an accretion disk with a radius larger than the 80 AU accretion zone radius of the sink particle) for the run time considered. However, we find that once an extended disk formed, disk accretion supplies material to the primary star. Our results show that the radiation escapes through low density channels that are not necessarily located along the polar directions of the star and that sustained radiation pressure dominated bubbles do not appear until late times when a significant accretion disk develops. For stronger turbulence at the level seen in many massive cores, we would expect this effect could be enhanced. Our results suggest that the ``flashlight" effect that occurs in our laminar runs, which allows the radiative flux to escape along the polar directions of the star and material to be accreted onto the star by an optically thick accretion disk, is not required for massive stars that form from turbulent cores. Instead, the asymmetric density distribution allows the radiation to escape through the path(s) of least resistance, allowing the dense infalling material to fall onto the star regardless of its location. Our results also demonstrate that RT instabilities are a natural occurrence in the formation of massive stars regardless of whether the star-forming core is initially turbulent or laminar. These instabilities arise immediately for turbulent cores because the initial turbulence seeds the instabilities. RT instabilities develop later for laminar cores because the initially symmetric gas distribution must be perturbed. These perturbations can then seed RT instabilities that grow in time and can eventually deliver material to the star-disk system. We find that the development of an accretion disk and gravitational torques induced within the disk destroy the symmetry of the gas distribution and seed the initial perturbations that lead to RT instabilities in the bubble shells as first shown by \citet{Krumholz2009a}. Our work suggests that the seeds for RT instabilities that arise in initially laminar pre-stellar cores are asymmetries induced by the shielding of the direct radiation field by the accretion disk and the non-symmetric distribution of material within the bubbles. These asymmetries arise from disk flaring, disk fragmentation, and the gravitational interaction of the massive star with the accretion disk and companions. Previous work that simulated the collapse of initially laminar cores concluded that the direct radiation field inhibited the growth of RT instabilities \citep{Kuiper2012a, Klassen2016a}. Contrary to their results we find that inclusion of the direct radiation field only suppresses the non-linear growth of these instabilities at early times. As the asymmetry in the system grows, these instabilities can grow non-linearly and become dense enough to overcome the radiation-pressure barrier and deliver material to the star-disk system. We argue instead that poor shell resolution is the likely culprit as to why \citet{Kuiper2012a} and \citet{Klassen2016a} do not obtain bubble shells that become RT unstable. We check this hypothesis directly by conducting a resolution study where we intentionally de-resolve the bubble shell to the point where our resolution is comparable to that used in earlier work, and we show that doing so both delays the onset of instability and reduces its intensity. We further find in the work of \citet{Kuiper2012a}, that limitations of their fixed grid approach with a star that is centrally fixed in a spherical grid does not permit the movement of the star-disk system that would naturally allow asymmetries to arise and lead to seeding the RT instability. We find that both turbulent and laminar cores lead to hierarchical star systems that consist of a massive primary star and several low-mass companions. We find that our multiplicity results are sensitive to the physics included, radiative transfer treatment used, and sink creation and merging criteria employed. Inclusion of the direct radiation pressure leads to cooler disks that are prone to greater fragmentation when compared to our comparison run that neglected the direct radiation field and assumed that the stellar radiation was immediately absorbed within the vicinity of the stars. However, given the idealized nature of our simulations, we cannot address the true multiplicity properties of massive stars. Despite this limitation, our work settles a long-debated question in massive star formation: how does radiation pressure limit the masses of stars? We find that radiation pressure is still an important feedback mechanism that must be considered in massive star formation, but RT instabilities can overcome the radiation pressure barrier, at least in the context of the idealized numerical experiments performed thus far. However, there are still many other physical processes at play that we neglect. These include collimated outflows, fast stellar winds, and magnetic fields. Future work will include these other feedback mechanisms to determine a more complete picture of how massive stars form and how their associated feedback can limit stellar masses. | 16 | 7 | 1607.03117 |
1607 | 1607.05526.txt | We present an analysis of the gas-phase oxygen abundances of a sample of 28 galaxies in the local Universe ($z<0.02$) hosting Type Ia Supernovae (SNe~Ia). The data were obtained with the 4.2m William Herschel Telescope (WHT). We derive local oxygen abundances for the regions where the SNe~Ia exploded by calculating oxygen gradients through each galaxy (when possible) or assuming the oxygen abundance of the closest \Hii region. The sample selection only considered galaxies for which distances not based on the the SN~Ia method are available. Then, we use a principal component analysis to study the dependence of the absolute magnitudes on the color of the SN~Ia, the oxygen abundances of the region where they exploded, and the stretch of the SN light curve. We demonstrate that our previous result suggesting a metallicity-dependence on the SN~Ia luminosity for not-reddened SNe~Ia \citep{2016ApJ...818L..19M} can be extended to our whole sample. These results reinforce the need of including a metallicity proxy, such as the oxygen abundance of the host galaxy, to minimize the systematic effect induced by the metallicity-dependence of the SN~Ia luminosity in future studies of SNe~Ia at cosmological distances. | \label{Section1} Type Ia Supernovae (SNe~Ia) are claimed to be thermonuclear explosions of carbon-oxygen white dwarfs (CO WD) \citep{1960ApJ...132..565H}. Their origin is not well-established, since there is still an open discussion about the different possibles progenitor scenarios. The single-degenerate scenario \citep[SD;][]{1982ApJ...253..798N, 1973ApJ...186.1007W} occurs when a WD in a binary system accretes mass from its non-degenerate companion until the Chandrasekhar mass limit ($\sim$~1.44 $M_{\odot}$) is reached. At that moment, the degenerate-electron pressure is not longer supported, and the thermonuclear explosion occurs. On the other hand, the double-degenerate (DD) scenario \citep{1984ApJS...54..335I,1984ApJ...277..355W} consists of two CO WDs gravitationally bounded that lose angular momentum and merge \citep{1976Afz....12..521T, 1979AcA....29..665T}. At that moment, the SN explodes resulting no fossil but the SN remnant \citep{2012Natur.489..533G}. SNe~Ia are very bright (M$_{\rm B}\sim$-19.4 mag at peak), and show very low intrinsic luminosity dispersion \citep[around 0.36 mag,][]{1993ApJ...405L...5B} so they are considered extraordinary tools for measuring distances in cosmological scales. Although SNe~Ia are called \textit{standard candles}, they are not pure \textit{standard}, but \textit{standarizable}. \cite{1993ApJ...413L.105P}, \cite{1996AJ....112.2391H,1996AJ....112.2438H}, and \cite{1999AJ....118.1766P} proved that a correlation between the absolute magnitude at maximum brightness and the luminosity decline after maximum, lately parametrized as the light-curve (LC) width. \citet{1996ApJ...473...88R} also found a relation between the peak magnitude and the SN color. In this way, the distance to these objects can be estimated from their distance modulus $\mu = m_{B} - M_{B}$ (where $m_{B}$ is the apparent magnitude and $M_{B}$ the absolute magnitude, both in band $B$) by just studying SNe~Ia multiwavelength LCs. These calibrations allowed to reduce the scatter of distances in the \textit{Hubble Diagram} (HD), in which $\mu$ is represented as function of redshift, $z$. In fact, diverse standardization techniques have been developed to standardize SN~Ia LCs and obtain the absolute magnitudes at maximum, and simultaneously the parameters that better reduce the scatter in the HD. Modern techniques, such as SALT2 \citep{2007A&A...466...11G} adjust SN LC templates to the observed LC and determine the SN~Ia color at maximum brightness ($C$), the stretch applied to the LC template ($s$), the apparent magnitude at maximum brightness in the $B$ band ($m_{B}$), and the epoch of the maximum brightness ($t_{max}$). Then, the distance modulus $\mu_{SALT}$ can be calculated using the equation \begin{equation} \mu_{SALT}=m_{B}-(M_{B}-\alpha\,(s-1)+\beta\,C), \label{eq:salt} \end{equation} where $\alpha$, $\beta$ and $M_{B}$ are obtained by minimizing the HD residuals. With similar techniques the SNe~Ia-based cosmology projects discovered that the Universe is in accelerated expansion \citep{1999ApJ...517..565P,1998AJ....116.1009R}. However after this method, there still exists a certain inhomogeneity in SNe~Ia at peak. A plausible source of inhomogeneity is a dependence of the properties of the SN~Ia on the characteristics of its environment. Since the average properties of host galaxies evolve with redshift, any such dependence not included in the standardization techniques will impact on the cosmological parameter determination. Many recent studies have indeed analyzed the dependence of SNe~Ia properties on global characteristics of their hosts \citep{2006ApJ...648..868S,2008ApJ...685..752G, 2009ApJ...691..661H, 2009ApJ...700..331H,2010ApJ...715..743K, 2010MNRAS.406..782S, 2010ApJ...722..566L, 2011ApJ...743..172D, 2011ApJ...740...92G, 2011ApJ...734...42N,2010MNRAS.406..782S, 2012ApJ...755..125G, 2013MNRAS.435.1680J,2013ApJ...770..108C, 2014A&A...568A..22B, 2014MNRAS.438.1391P}. % In summary, all found that SNe~Ia are systematically brighter in more massive galaxies than in less massive ones {\sl after LC shape and color corrections}. Through the mass-metallicity relation \citep{2010MNRAS.406..782S} this would lead to a correlation between SNe~Ia magnitudes and the metallicities of their host galaxies: more metal-rich galaxies would host brighter SNe~Ia {\sl after corrections}. However the cause of these correlations is not well-understood. % In addition, due to the metal enrichment of galaxies with time, a change in chemical abundances with redshift \citep{2006ApJ...644..813E,2009A&A...505..529L} is expected. All these SNe~Ia calibrations are based on local objects mostly having around solar abundances\footnote{Here we use the terms metallicity, total abundance in metals, Z, (being X+Y+Z=1 in mass), and oxygen abundances indistinctly, assuming that $\log(\rm Z/Z_{\odot})=\log({\rm O/H})-\log({\rm O/H})_{\odot}$, $12+\log({\rm O/H})_{\odot}=8.69$, and $Z_{\odot}=0.019$ being the solar values \citep{2009ARA&A..47..481A}.}. Therefore, a standard calibration between the LC shape and the M$_{B}$ of SNe~Ia might not be completely valid for objects with chemical abundances which are different to those for which the calibration was made. Therefore, the metallicity may be one source of systematic errors when using these techniques. The dependence of SNe~Ia luminosity on metallicity was studied by \citet{2005ApJ...634..210G}, who estimated elemental abundances using emission lines from host-galaxy spectra following the \cite{2002ApJS..142...35K} method. They found that most metal-rich galaxies have the faintest SNe~Ia. \citet{2008ApJ...685..752G} analyzed the spectral absorption indices in early-type galaxies, also finding a correlation between magnitudes and the metal abundance of their galaxies, in agreement with the above trend observed for late-type galaxies reported. These results are however not precise enough: \citet{2005ApJ...634..210G} based their conclusion on the analysis of the Hubble residuals (see their Figure 15a), which implies the use of the own SN~Ia LC to extract the information, while \citet{2008ApJ...685..752G} used theoretical evolutive synthesis models which still have many caveats, since predictions are very dependent on the code used technique and input spectra, with the extra bias included by the well-known age-metallicity degeneracy in theoretical evolutive synthesis models. Theoretically, there is a predicted dependence between the maximum luminosity of the SN~Ia and the metallicity of the binary system: assuming the progenitor mass (WD) is constant, the parameter which leads the relation between the light curve width and its maximum magnitude is the opacity of the outer part of the ejected material \citep{1996ApJ...457..500H,2001ApJ...547..988M}, which depends on temperature and, thus, on the heating due to the radioactive $\rm ^{56}Ni$ decay. Then the luminosity of the supernova depends basically on the $\rm ^{56}Ni$ mass ejected from the explosion \citep{1982ApJ...253..785A}: \begin{equation} L \propto M(^{56}\rm Ni)\hspace{0.13cm} erg\,s^{-1}. \label{timmes} \end{equation} \citet{2003ApJ...590L..83T} showed that the neutron excess, which controls the radioactive ($\rm ^{56}Ni$) to non-radioactive (Fe-peak elements) abundance ratio, in the exploding WD is a direct function of the initial metallicity of the WD progenitor. This acts upon the maximum luminosity of the explosion \citep[see][for detailed calculations]{2005A&A...443.1007T,2006astro.ph..8324P}. The maximum luminosity of the SN~Ia depends thus on the initial abundances of C, N, O, and Fe of the progenitor WD. Models by \citet{2003ApJ...590L..83T} predicted this dependence, suggesting that a variation of a factor 3 in the metallicity may cause a variation up to $\sim 25\%$ in the mass of $^{56}$Ni synthesized during the explosion for initial metallicities higher than solar. More recently, \cite{2010ApJ...711L..66B} computed a series of SNe~Ia scenarios, finding an even stronger dependence on metallicity (see their Figure 1 and Eq.~2) than that estimated using \cite{2003ApJ...590L..83T}, \begin{equation} \label{bravo2} M(^{56}{\rm Ni}) \sim f({\rm Z})=1-0.075\frac{\rm Z}{\rm Z_{\odot}}. \end{equation} Following \citet{2007ApJ...655L..93C} prescriptions, \cite{2010ApJ...711L..66B} also explored the dependence of the explosion parameters on the local chemical composition, C mass fraction, and neutronization. These authors found a non-linear relation between the synthesized mass of $^{56}$Ni and the metallicity of the progenitor binary system (see their Figure 1 and Eq.~3): \begin{equation} M(^{56}Ni) \sim f({\rm Z})=1-0.18\frac{\rm Z}{\rm Z_{\odot}}\Bigg(1-0.10\frac{\rm Z}{\rm Z_{\odot}}\Bigg). \label{bravo3} \end{equation} This dependence on Z translates into different bolometric LC luminosity-width relationships for different metallicities. Their Figure 3, which plots the LC luminosity-width relationship for three initial metallicities (Z/Z$_{\odot}$=0.1, 1, and 3) and the same LC width, clearly shows this effect: the luminosity is {\it smaller} at higher Z than at lower Z. \cite{2010ApJ...711L..66B} results imply that SNe~Ia located in galaxies with metallicity higher than solar {\it might be dimmer} than expected as compared to those with solar and subsolar abundances. Since the number of SNe~Ia detections will extraordinarily increase in the forthcoming surveys, statistical errors will decrease while systematic errors will dominate, limiting the precision of SNe~Ia as indicators of extragalactic distances. Hence the importance of characterizing a possible dependence of the SN~Ia luminosity on the metallicity. The final purpose of this project is to seek if a dependence between the SNe~Ia maximum luminosity and the metallicity of its host galaxy (provided by the gas-phase oxygen abundance) does exist. For this, we perform a careful analysis of a sample of galaxies of the local Universe hosting SNe~Ia to estimate the oxygen abundances in the regions where those SNe~Ia exploded. Our aim is to perform this analysis in a very basic way, just searching for a simple dependence of the magnitude $M_{B}$ of the SNe~Ia in their LC maximum with the oxygen abundance without using any standardization technique. Therefore, we build our sample considering local SNe~Ia host galaxies that have distances well determined by methods which are different of those following SNe~Ia techniques. This way we estimate the absolute peak magnitude for each SN~Ia using the classic equation: $M_{B}=m_{B}-5\log{D} +5$, eliminating possible problems coming from the use of cosmological techniques. % On the other hand, since galaxies are nearby enough, gas-phase abundances may be estimated in several \Hii~regions across the galaxies, and at different galactocentric distances (GCDs), thus allowing us to derive, in many cases, metallicity gradients. We then use the corresponding value of the oxygen abundance at the same GCD the SN~Ia exploded as a proxy of its metallicity. This method has been already used in \citet{2012A&A...545A..58S, 2016arXiv160307808G}, who also studied galaxies hosting SNe Ia and estimated the oxygen abundances in the regions where the explosions took place, using radial gradients. The technique is different, though. They use the PPAK/PPMAS Integral Field Spectrograph (IFS) mounted in the 3.5m telescope at the Calar Alto Observatory, as part of the CALIFA collaboration project. They obtained oxygen abundances at every position along the galaxy disk of each galaxy. It also allows to obtain azimuthal averaged values and estimate the oxygen abundance at each SN Ia location. We have three galaxies in our sample that are in common with (NGC\,0105, UGC\,04195 and NGC\,3982), and we use these data to check and improve the accuracy of our results. % In \citet{2016ApJ...818L..19M} we presented our results obtained for non-reddened SNe~Ia ($z\le 0.02$), thus eliminating possible dependences of luminosities on the color of the objects, and found that this dependence on metallicity seems to exist. Our data show a trend, with an 80\% of chance not being due to random fluctuation, between SNe~Ia absolute magnitudes and the oxygen abundances of the host galaxies, in the sense that luminosities tend to be higher for galaxies with lower metallicities. Our result agrees with the theoretical expectations and with other findings suggested in previous works. In this paper we present all the details about the \mbox{analysis} of the oxygen abundances derived for our low-redshift SN~Ia host galaxies. Section~2 discusses the sample selection and the data reduction process. The analysis of the spectra and the determination of oxygen abundances and absolute magnitudes for the SN~Ia are described in Sect~3. Section~4 presents our results, which are discussed in Sect.~5. For this we are taking into consideration the SNe~Ia color and stretch parameters used in the classic Supernova Cosmology, performing a principal component analysis to seek dependences among observed parameters, including the oxygen abundance. Our conclusions are given in Sect.~6. \begin{figure} \centering \includegraphics[width=\linewidth]{figure_1.pdf} \caption{Distances to galaxies having both Tully-Fisher an Cepheids distances. Grey dashed line represents identity. When both measurements are available, a mean distance is calculated considering as many values as possible.} \label{distance_TF_Ceph} \end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 16 | 7 | 1607.05526 |
|
1607 | 1607.03874_arXiv.txt | During eruptive solar flares and coronal mass ejections, a non-pot{\-}ential magnetic arcade with much excess magnetic energy goes unstable and reconnects. It produces a twisted erupting flux rope and leaves behind a sheared arcade of hot coronal loops. We suggest that: the twist of the erupting flux rope can be determined from conservation of magnetic flux and magnetic helicity and equipartition of magnetic helicity. It depends on the geometry of the initial pre-eruptive structure. Two cases are considered, in the first of which a flux rope is not present initially but is created during the eruption by the reconnection. In the second case, a flux rope is present under the arcade in the pre-eruptive state, and the effect of the eruption and reconnection is to add an amount of magnetic helicity that depends on the fluxes of the rope and arcade and the geometry. | \label{sec1} The standard understanding of eruptive solar flares (e.g., \opencite{schmieder12}; \opencite{priest14a}; \opencite{aulanier14}; \opencite{janvier15}) is that excess magnetic energy and magnetic helicity build up until a threshold is reached at which point the magnetic configuration either goes unstable or loses equilibrium, either by breakout \cite{antiochos99a,devore08} or magnetic catastrophe \cite{demoulin88,priest90a,forbes91b,lin00,wang09} or kink instability \cite{hood79a} or by torus instability \cite{lin98,kliem06,demoulin10a,aulanier10,aulanier12}. Some solar flares (known as {\it eruptive flares}) are associated with the eruption of a magnetic structure containing a prominence (observed as a coronal mass ejection) and typically produce a two-ribbon flare, with two separating H$\alpha$ ribbons joined by a rising arcade of flare loops. Others are contained and exhibit no eruptive behaviour. While some coronal mass ejections are associated with eruptive solar flares, others occur outside active regions and are associated with the eruption of a quiescent prominence. Coronal mass ejections outside active regions do not produce high-energy products, because their magnetic and electric fields are much smaller than in eruptive solar flares, but their magnetic origin and evolution may well be qualitatively the same. Magnetic helicity is a measure of the twist and linkage of magnetic fields, and its basic properties were developed by \opencite{woltjer58}, \opencite{taylor74}, \opencite{moffatt78}, \opencite{berger84b}, \opencite{berger00} and \opencite{demoulin06b}. It was first suggested to be important in coronal heating, solar flares and coronal mass ejections by \opencite{heyvaerts84}, who proposed that, when the stored magnetic helicity is too great, it may be ejected from the Sun in an erupting flux rope (see also \opencite{rust95}; \opencite{low03}; \opencite{kusano02b}). The flux of magnetic helicity through the photosphere, its buildup in active regions and its relation to sigmoids has been studied by \opencite{pevtsov95}, \opencite{canfield98}, \opencite{canfield99,canfield00}, \opencite{pevtsov00b}, \opencite{pevtsov02a}, \opencite{green02a}, \opencite{pariat06b}, and \opencite{poisson15}. Indeed, the measurement of magnetic helicity in the corona is now a key topic with regard to general coronal evolution \cite{chae01a,demoulin02a,mandrini04a,zhang06,zhang08,mackay11,mackay14,gibb14}. Also, since magnetic helicity is well-conserved on timescales smaller than the global diffusion time, measuring it in interplanetary structures such as flux ropes and magnetic clouds \cite{gulisano05,qiu07,hu14,hu15} allows us to link the evolution of CMEs in the solar wind with their source at the Sun \cite{nindos03,luoni05}. The common scenario described above for an eruptive solar flare or coronal mass ejection is that, after the slow buildup of magnetic helicity in a magnetic structure, the eruption is triggered and drives three-dimensional reconnection which adds energy to post-flare loops. Originally, it was thought that all of the magnetic energy stored in excess of potential would be released during a flare, and therefore that the final post-flare state would be a potential magnetic field. However, the modern realisation is that magnetic helicity conservation provides an extra constraint that produces a different final state. It is also observed that flare loops in an eruptive flare do not relax to a potential state, since the low-lying loops remain quite sheared \cite{asai04,warren11c,aulanier12}. Our aim in this paper is to determine two key observational consequences of the reconnection process, namely, the amount of magnetic helicity, and therefore twist, in the erupting flux rope, and in the shear of the underlying flare loops. \begin{figure}[h] {\centering \includegraphics[width=9cm]{fig1.eps} \caption{(a) A simple sheared arcade that reconnects to produce (d) an erupting flux rope plus an underlying less-sheared arcade. In the initial state (a), footpoints A$_+$, B$_+$, C$_+$, D$_+$, E$_+$, F$_+$, are joined to footpoints A$_-$, B$_-$, C$_-$, D$_-$, E$_-$, F$_-$, respectively. In the first stage, shown in (b), loops A$_+$A$_-$ and F$_+$F$_-$ have reconnected to give new loops A$_+$F$_-$ and F$_+$A$_-$. In the second stage, shown in (c), loops B$_+$B$_-$ and E$_+$E$_-$ have reconnected to give new loops B$_+$E$_-$ and E$_+$B$_-$. Then in the final stage, shown in (d), loops C$_+$C$_-$ and D$_+$D$_-$ have reconnected to give a twisted flux rope C$_+$D$_-$ that may erupt and a set of underlying loops D$_+$C$_-$. (e) and (f) show a projection of the field lines onto a vertical section through the configuration, looking from the left, in the initial (left) and final state (right) during the eruption.} \label{fig1}} \end{figure} \begin{figure}[h] {\centering \includegraphics[width=12cm]{fig2.eps} \caption{(a) A sheared arcade overlying a flux rope reconnects to produce (b) an erupting flux rope plus an underlying less-sheared arcade.} \label{fig2}} \end{figure} We determine the magnetic helicity in the erupting flux rope and underlying flare loops in two cases. In the first (Fig.\ref{fig1}), no flux rope is present in the initial arcade, but a flux rope is created during the process of the reconnection. In the second case (Fig.\ref{fig2}), the initial state consists of a magnetic arcade that overlies a flux rope, and the effect of the eruption and reconnection is to enhance the flux and twist of the flux rope. Our aim is to produce the simplest model that preserves the core physics of the process. For example, we neglect the internal structure of the coronal arcade and model it simply as a flux rope whose photospheric footpoints are stretched out in two lines either side of the polarity inversion line. The second case (with the initial flux rope) is much more likely to go unstable and erupt. We calculate the twist in the erupting flux rope from the properties of the initial pre-eruptive state by making three assumptions, namely: (i) conservation of magnetic flux; (ii) conservation of magnetic helicity; (iii) and equipartition magnetic helicity. The third assumption implies that the same amount of magnetic helicity is transferred by reconnection to the erupting flux rope and underlying arcade \cite{wright89}. We have considered an alternative possibility, namely, preferential transfer of magnetic helicity to the flux rope during the reconnection, which would imply the flaring loops have vanishing self-helicity. However, this seems less likely, because, although the flare loops are generally seen as non-twisted structures, they are also observed to be nonpotential, since the low-lying loops appearing early in the flare possess more shear than the high-lying loops \cite{aulanier12}. The reconnection of twisted tubes has been studied in landmark papers by \opencite{linton01}, \opencite{linton02}, \opencite{linton03}, \opencite{linton05} in which they consider also the extra constraint of energy both numerically and analytically. Straight tubes of a variety of twists and inclinations are brought together by an initial stagnation-point flow and allowed to reconnect self-consistently by either bouncing, slingshot, merging or tunnelling reconnection. Although energy is not considered in detail here, we make initial comments about energy considerations in Section \ref{sec3.3} and hope to develop them further in future numerical treatments. The structure of the paper is as follows. In Section \ref{sec2} we summarise the basic properties of magnetic helicity that are needed for our analysis. Then in Section \ref{sec3} we present our simple model for a sheared magnetic arcade in which no flux rope is present initially but is created during the eruption by the reconnection. This is followed in Section \ref{sec4} by a different model in which the initial arcade contains a flux rope, which erupts and leaves behind a less-sheared arcade. Finally, suggestions for follow-up are given in Section \ref{sec5}. | \label{sec5} We have set up a simple model for estimating the twist in erupting prominences, in association with eruptive two-ribbon flares and/or with coronal mass ejections. It is based on three simple assumptions, namely, conservation of magnetic flux, conservation of magnetic helicity and equipartition of magnetic helicity. While the first and second are well established, the third is more of a reasonable conjecture. In future, it would be interesting to test the model and the conjecture with both observations and computational experiments. During the main phase of a flare, the shear of the flare loops is observed to decrease in time, so that they become oriented more perpendicular to the polarity inversion line. This is a natural consequence of our model, where flux is added to the flux rope first from the innermost parts of the arcade (i.e., closest to the polarity inversion line), so that, as can be seen in Figs.\ref{fig1}b, \ref{fig2}b and \ref{fig9}b, the final shear of the arcade is smaller than the initial shear. In other words, the change in shear is a consequence of the geometry of the three-dimensional reconnection process. The cause of the eruption is a separate topic that has been discussed extensively elsewhere (e.g., \opencite{priest14a}), and includes either nonequilibrium, kink instability, torus instability or breakout. One puzzle is what happens with confined flares, where the flare loops and H$\alpha$ ribbons form but there is no eruption. A distinct possibility is that the overlying magnetic field and flux are too strong to allow the eruption, but this needs to be tested by comparing nonlinear force-free extrapolations with observations (e.g., \opencite{wiegelmann08b,mackay11,mackay12b}). Another puzzle is the cause of preflare heating. One possibility is the slow initiation of reconnection before a fast phase (Yuhong Fan, private communication), but another is that the flux rope goes unstable to kink instability which spreads the heating nonlinearly throughout the flux rope in a multitude of secondary current sheets \cite{hood09a}; if the surrounding field is stable enough, such an instability can possibly occur without an accompanying eruption. The present simple model can be developed in several ways, which we hope to pursue in future. One is to conduct computational experiments, in which the energies before and after reconnection will be calculated in order to check which states are energetically accessible. A second way is to extend the model to more realistic initial configurations with more elements, in which the reconfiguration is by quasi-separator or separator reconnection \cite{priest96a,longcope96,longcope05a,longcope2008b,parnell10b} and the internal structure of the flux rope and arcade are taken into account. In particular, the distribution of magnetic flux within the arcade will be included, both normal to and parallel to the polarity inversion line. \appendix | 16 | 7 | 1607.03874 |
1607 | 1607.01795_arXiv.txt | The Quintuplet cluster is one of the most massive star clusters in the Milky Way, situated very close to the Galactic center. We present a new search for variable stars in the vicinity of the cluster, using the five-year database of the Vista Variables in the Via Lactea (VVV) ESO Public Survey in the near-infrared. A total of 7586 objects were identified in the zone around $2'$ from the cluster center, using 55 $K_S$-band epochs. Thirty-three stars show $K_S$-band variability, 24 of them being previously undiscovered. Most of the variable stars found are slow/semiregular variables, long-period variables of the Mira type, and OH/IR stars. In addition, a good number of our candidates show variations in a rather short timescale. We also propose four Young Stellar Object (YSO) candidates, which could be cluster members. | The Galactic Center (GC) is an exceptional laboratory to test the formation and evolution of stars under extreme conditions. Three massive and young star clusters have been found in this region: Arches, Quintuplet and the central cluster surrounding Sagittarius A*, a supermassive black hole. Each is supposed to contain about $10^4$ stars. Among these, Quintuplet has been known to be rich in massive stars, such as Wolf-Rayet and OB supergiants \citep{fig99}. Even though this is a well-studied cluster, not many variability searches have been carried out, and the few available are mostly restricted to bright stars. This could be a consequence of the observational challenges this field poses. The first known variable was the \lq Pistol star\rq \, \citep{fig98}, a luminous blue variable (LBV), and one of the brightest stars in the Milky Way; a second LBV was discovered by \citet{geb99}. \citet{gla99} found several variable candidates brighter than $K_S=11$, including some large-amplitude asymptotic giant branch (AGB) variables. Later on, \citet[hereafter MKN09]{mat09} performed a near-infrared survey of Miras in the GC, which covered the Quintuplet area. It is clear that a better understanding of the stellar population of this cluster can provide valuable information about the life cycle of stars in the central regions of our Galaxy. The irruption of near-infrared surveys, like the Vista Variables in the Via Lactea ESO Public Survey \citep[VVV;][]{min10, sai12}, can help us to achieve that goal, due to its multi-epoch and multi-band nature. For example, \citet{dek15} recently discovered a young thin stellar disk population across the Galactic bulge, using data from this survey. In this paper we analyze $K_S$-band data from the VVV survey around the Quintuplet star cluster, in order to find variable sources that allow us to get more insights about the evolution of stars in the cluster. \subsection{Quintuplet} The Quintuplet cluster ($\alpha$: $17^h46^m15^s$, $\delta$: $-28^{\circ}49'41''$; J2000), named after the prominent presence of five bright stars \citep{nag90, oku90}, is located at a projected distance of about 30~pc from Sgr A*. With an age of $3.0\pm0.5$~Myr \citep{lie12, lie14}, and a tidal radius of $\approx1$~pc \citep{por02}, the central part of the cluster presents a flat present-day mass function with a slope of $-1.68$ \citep{hus12}, which may be caused by the fast dynamical evolution of the cluster within the strong Galactic tidal field. A similar result has been found for the inner region of Arches \citep{hab13}. Given that massive clusters in the GC are supposed to dissolve within a few tens of Myr \citep{por02}, we expect that these clusters significantly contribute to the isolated massive stars population \citep{hab14}. Furthermore, \citet{hus12} determined cluster membership through proper motion analysis. And finally, \citet{sto14} studied the orbital motion of the Quintuplet, and suggested that this cluster and Arches may have a common origin, a situation that we intend to address in a subsequent paper. | \subsection{Comparison with Previous Works} As mentioned in Section~1, only a few variability searches have been carried out in Quintuplet. From our sample, nine candidates had been previously discovered. Hence, 24 new variable candidates have been found by the present work. The results are summarized in Table~\ref{tbl-2}. The first column contains the name used in the present paper. Column 2 includes the number used in the respective catalog: [MKN09] stands for \citet{mat09}, and [MFK13] for \citet{mat13}. Columns 3 and 4 compare periods found by our present work and previous papers, respectively, while column 5 delivers the variability type of the candidate. In Figure~\ref{fig:comp} we combined our data with the information available from MKN09 and MFK13. The agreement with the two variables from \citet{mat13} is remarkable, with both periods being almost equal. For the Miras of MKN09 we found a good agreement, specially for VC19. Additionally, we managed to estimate periods for VC03 and VC31, while MKN09 could not find any. VC03 is worth discussing, since this is the candidate with the most notorious discrepancy with MKN09. We cannot attribute these differences to photometric errors, thus a possible explanation is that the star has undergone changes in the nature of its variability during the past years. As a consequence, we tentatively classify this variable as a slow irregular or semiregular. The other three candidates (VC23, VC25 and VC32) do not show a clear periodicity, even with the combined data, so we also classify them as slow irregular or semiregular variables. \begin{table} \centering \caption{Previously discovered variables} \label{tbl-2} \begin{tabular}{rcccl} \hline ID & Prev. ID & P & Prev. P & Type\\ \hline VC03 & [MKN09] 1172 & 213.3 & & Mira\\ VC04 & [MFK13] 37 & 0.94255 & 0.94255 & Ecl \\ VC07 & [MKN09] 1112 & 103.1 & 102 & Mira \\ VC16 & [MFK13] 39 & 31.29 & 31.279 & Ceph(II) \\ VC19 & [MKN09] 1250 & 110.9 & 108 & Mira \\ VC23 & [MKN09] 1194 & & & Mira \\ VC25 & [MKN09] 1286 & & & Mira \\ VC31 & [MKN09] 1260 & 525.1 & & Mira \\ VC32 & [MKN09] 1236 & & & Mira \\ \hline \end{tabular} \end{table} \begin{figure*} \includegraphics[width=\textwidth]{variables_matsunaga_new} \caption{Combination of epochs from VVV (circles) and MKN09, MKF13 (crosses) for six variable candidates. Except by VC03 and VC31, periods were fitted to the combined data, and they are similar to those found individually in table 2. For those two particular cases, we kept the original period found by this work, since there is no agreement between both data sets.\label{fig:comp}} \end{figure*} \subsection{GLIMPSE data} Since one of our long-term goals is to detect young stellar objects, we cross-checked our sample with the magnitudes of GLIMPSE \citep{ben03,chu09}, in order to search for infrared excess. The data was obtained from the GLIMPSE Source Catalog (I + II + 3D), directly from the IRSA website\footnote{http://irsa.ipac.caltech.edu/data/SPITZER/GLIMPSE}. The cross-match found magnitudes for 17 of our objects, with a tolerance of $1''$. The resulting color-color diagram is shown in Figure~\ref{fig:glimpse}. The point in the upper right corner corresponds to VC17, an extremely red star. This is a peculiar object, with a very large $K_S$ amplitude. Unfortunately, it does not appear in our $H$ or $J$ photometry, since at the epoch of these measurements the star was in a faint phase, hence no additional information could be obtained from the CMD. All candidates with a matched GLIMPSE photometry show some degree of infrared excess. Five of these objects have $[3.6]-[4.5]>1$, these are VC06, VC15, VC19 and VC23, besides the already mentioned VC17. \begin{figure} \includegraphics[width=\columnwidth]{glimpse_bw} \caption{Color-color diagram for the variable candidates found within the GLIMPSE database. Periodic variables are plotted as open circles, while black circles correspond to non-periodic stars. The dashed lines enclose the region of MS stars.\label{fig:glimpse}} \end{figure} \subsection{Analysis of individual candidates} \label{disc} As mentioned in Section~\ref{cmd}, our variable candidates do not seem to belong to the cluster population. In addition, an analysis of the CMD position and shape of their light curves led us to conclude that the majority of these candidates correspond to long-period variables (LPV), including Miras and semiregular variables, which are typically oscillating red giants \citep[see, e.g.,][]{cat15}. If we take into account the young age of the Quintuplet cluster, we do not expect to find a significant number of red giant members. In addition to the Miras already discovered by MKN09, VC06 and VC17 can be classified as Miras too. For VC01 and VC07, periodicity is not well defined, thus they are classified as semiregular variables. Other variables like VC05, VC08 and VC12 either have very long periods (over $\sim 900$ days) or are semiregular, meaning they are probably OH/IR stars. A larger time coverage would be required to answer this question. In any case, since these are evolved stars, they are most likely not related to the cluster population. Given its position on the CMD, VC02 seems a very strong YSO candidate. Located in the PMS branch of the 4~Myr isochrone, its light curve shows some sudden brightness increases, the most important being around $\sim 1000$ days. These episodes are consistent with the bursts in the accretion rate found in PMS stars. If we assume it is a cluster member, the position in the CMD suggests a mass of $\sim 2 M_{\odot}$. Since we lack $J$-band and GLIMPSE photometry for this object, no color-cut could be performed to classify it. VC13 and VC26 also qualify as YSO candidates, given their positions in the CMD and the shape of the light curves. A similar situation occurs with VC21. Its light curve show strong, irregular bursts that resemble a YSO, and its $K_S$ magnitude and position in the field makes us think it is a young Quintuplet member. However, no $H$ {\bf or} $J$ information is available for that candidate. For all of these objects, spectroscopic follow-up is required to confirm them as YSOs. The case of VC11 is not as clear as for the other candidates discussed. As observed in the right panel of Figure~\ref{cmd}, the object lies too far from the PMS tracks, even from the youngest one. The light curve is classified as that of an irregular variable star. VC01, VC09, VC22 and VC28 are classified as semiregular variables, since there is a resemblance of periodicity in all of their light curves. The remainder of the non-periodic variables are difficult to classify. Some of them, for example VC15 and VC24, closely resemble classical T-Tauri light curves \citep[see e. g.][]{mcg15}, but are too bright to be YSOs belonging to the cluster. However, they could be pre-main sequence field stars. | 16 | 7 | 1607.01795 |
1607 | 1607.05663_arXiv.txt | \vspace{.3in} \noindent We study two-field bouncing cosmologies in which primordial perturbations are created in either an ekpyrotic or a matter-dominated contraction phase. We use a non-singular ghost condensate bounce model to follow the perturbations through the bounce into the expanding phase of the universe. In contrast to the adiabatic perturbations, which on large scales are conserved across the bounce, entropy perturbations can grow significantly during the bounce phase. If they are converted into adiabatic/curvature perturbations {\it after} the bounce, they typically form the dominant contribution to the observed temperature fluctuations in the microwave background, which can have several beneficial implications. For ekpyrotic models, this mechanism loosens the constraints on the amplitude of the ekpyrotic potential while naturally suppressing the intrinsic amount of non-Gaussianity. For matter bounce models, the mechanism amplifies the scalar perturbations compared to the associated primordial gravitational waves. | The currently accepted paradigm for the formation of all structure in the universe is that stars and galaxies collapsed from small primordial density perturbations. Direct evidence for the existence of such early density perturbations is provided by the temperature fluctuations in the cosmic microwave background. But what caused these primordial density perturbations? An intriguing possibility (and currently the theoretically best studied option) is that these density fluctuations arose from the amplification of quantum perturbations via particular cosmological dynamics. The most popular version of this scenario involves a period of accelerated expansion, inflation, preceding the hot big bang cosmological evolution \cite{Mukhanov:1981xt, Hawking:1982cz, Guth:1982ec, Starobinsky:1982ee, Bardeen:1983qw}. But this is not the only possibility. Quantum perturbations can also be amplified during phases of cosmological contraction, in particular during an ekpyrotic phase of slow high-pressure contraction \cite{Notari:2002yc, Finelli:2002we, Tsujikawa:2002qc}, or during a period of matter-dominated pressure-free contraction \cite{Wands:1998yp, Finelli:2001sr}. In all of these scenarios, expanding and contracting, if more than one field is present then in addition to density/curvature perturbations there can occur an amplification of entropy/isocurvature perturbations. What is more, entropy perturbations can act as a source for curvature perturbations, and in this manner the curvature perturbations present at the onset of the hot big bang phase can have an involved pre-history. (Note in this context that once/if the universe reaches thermal equilibrium during the hot big bang phase, the entropy perturbations disappear \cite{Weinberg:2004kf}, which implies that it is not surprising that none have been seen to date in the CMB data \cite{Ade:2015xua}.) The vast majority of cosmological models studied to date have assumed that the curvature perturbations of interest have either been fully created only during an expanding or a contracting phase of cosmological evolution. In the present paper, we want to explore a different possibility: namely that entropy perturbations are created during a contracting phase of the universe, and that they act as an important source for the curvature perturbations only at the beginning of the current expanding phase. We show that such a scenario offers interesting new possibilities, due to a possible non-trivial evolution of entropy perturbations during the bounce phase joining the contracting and expanding phases together. In particular, we will show with the example of a specific bounce model that entropy perturbations can grow significantly during the bounce phase, with the consequence that these enhanced entropy perturbations can easily provide the dominant contribution to the final post-bounce amplitude of the curvature perturbations. We will consider two types of models: ekpyrotic models in which essentially only entropy perturbations are amplified during the contraction phase, and matter-dominated contraction models in which both curvature and entropy perturbations are created in roughly equal measure during the contraction phase. We then follow these perturbations across a non-singular ghost condensate bounce, where the long-wavelength curvature perturbations of interest remain conserved, while the entropic fluctuations are amplified further. In the ensuing expansion phase the entropy perturbations are then converted into curvature fluctuations. This has several broad consequences: \begin{itemize} \item The amplitude of the final curvature perturbations tends to be significantly enhanced. Conversely, this means that the entropy perturbations at the end of the ekpyrotic or matter-dominated contracting phases can be smaller than currently assumed in these models. In other words, the energy scale of the contraction phase can be smaller. \item The enhancement of the amplitude of the entropy perturbations implies that the importance of its intrinsic non-Gaussianity is reduced. This is because the most important contribution to the non-Gaussianity of the final curvature perturbation comes from the non-linearity of the conversion process, and not from the intrinsic non-Gaussianity that can already be present in the entropy perturbations. \item In matter-dominated contraction phases primordial gravitational waves are also produced, with an amplitude that is comparable to that of the curvature perturbations that are produced simultaneously \cite{Wands:1998yp}. This leads to a tension with current observations. In the models that we study here, the entropy perturbations grow to become even larger than the curvature perturbations produced during contraction, and the final amplitude of curvature perturbations (after conversion during the expanding phase) can thus be significantly enhanced compared to the tensor perturbations with the consequence that the tensor-to-scalar ratio is typically reduced by one or more orders of magnitude, in some cases bringing it close to current observational bounds. \end{itemize} We will begin with a review of non-singular ghost condensate bounce models in Sec.~\ref{section:bounce}, where we also discuss the evolution of both curvature and entropy perturbations in bouncing space-times. These results can then be used in specific ekpyrotic (Sec.~\ref{section:ekpyrotic}) or matter-dominated models (Sec.~\ref{section:matter}). We end with a discussion in Sec.~\ref{section:discussion}. | \label{section:discussion} We have analysed ekpyrotic and matter-dominated contracting cosmologies in which entropy perturbations provide the dominant source of the final curvature perturbations, and thus the dominant source of the eventual temperature fluctuations observed in the cosmic microwave background. In these models entropy perturbations are generated during the contracting phase. However, in contrast to earlier models, these entropy perturbations are not converted into curvature perturbations at the end of the contracting phase, but only during the subsequent expanding phase, following a bounce. The crucial difference in having the entropy perturbations around for longer is that they become increasingly important. We have shown that this is due to a combination of effects: (i) they grow by a factor of a few during the bounce itself, i.e.~during the period when the null energy condition is violated, (ii) they grow logarithmically during the kinetic phase that immediately follows the bounce (and will also grow logarithmically during any kinetic phase that may precede the bounce), and (iii) entropy perturbations can potentially grow significantly more if there is an unstable transverse potential at any time during the ekpyrotic, kinetic or bounce phases. An important point is that although (in the absence of an unstable potential) the growth is only logarithmic during the kinetic-dominated phases, the cumulated growth can be very significant and can in fact increase the amplitude of the entropy perturbations by up to several orders of magnitude. Thus if a conversion event of entropy into curvature fluctuations occurs after the bounce, then typically one may expect the entropy perturbation to be the main originator of the late-time density fluctuations. Note that there is no conflict with the fact that the temperature fluctuations in the CMB are observed to be adiabatic, as in the present models the entropy perturbations are converted into curvature perturbations during the expanding phase, while we assume that the universe subsequently reaches thermal equilibrium at which time the remaining entropy perturbations vanish. The extra amplification of entropy perturbations implies that the fluctuations generated during the contracting phase can have a smaller initial amplitude than what is currently assumed in such models, i.e.~the contracting phase can end at a lower energy scale. The main observational signature of the conversion process are the associated non-Gaussianities. While they can be small as long as the conversion process is efficient, they are certainly not negligible. Our numerical studies suggest that efficient conversions typically lead to \be -5 \lesssim f_{NL} \lesssim +5\,, \ee which is consistent with the current observational bounds $f_{NL} = 0.8 \pm 10.0$ at the $95\%$ confidence level from the Planck collaboration \cite{Ade:2015xua}, but within the reach of not-too-distant future experiments such as the SKA \cite{Gaensler}. A further distinguishing feature of these models is that one expects them to also have a non-trivial running of the spectral index, as shown in \cite{Lehners:2015mra}. Therefore, these models provide interesting alternatives to existing models, with the benefit that they can be tested by near-future observations. As for most cosmological models, a challenge remains to embed the present models in a more fundamental framework. There are several aspects to this challenge: an obvious one is to find a model that contains the appropriate fields and potentials. While it should be straightforward to find potentials that lead to a conversion process \cite{Lehners:2007nb}, it remains to be demonstrated that one can derive potentials from a more fundamental theory that lead to efficient conversions. As we have discussed, the efficiency of the conversion process is paramount in obtaining reliable and phenomenologically interesting predictions. What is more, non-singular bounces are extremely interesting effective theories, but it remains an open problem to derive them as effective descriptions of say loop quantum gravity bounces, or from a string theory setting. (Bounces naturally arise in loop quantum cosmology \cite{Ashtekar:2006wn} as well as in condensate states of loop quantum gravity \cite{Oriti:2016ueo}, but it is not known if loop quantum gravity effects can cause an efficient conversion process near the bounce point; bounces can also be embedded in supergravity \cite{Koehn:2013upa}, but a possible relation to a more fundamental theory remains unknown --- these are both promising starting points, but much remains to be done.) Also, as we mentioned in discussing a braneworld inspired setting, the conversion process may have interesting links to the formation and abundance of dark matter. Given the large number of cosmological models that can potentially explain the origin of the primordial fluctuations, it might well be the case that our trust or distrust of various models will come not only from more detailed observations (which will certainly play a leading role), but also from additional correlations with altogether different observations, such as the dark matter abundance. This particular point thus deserves further study. | 16 | 7 | 1607.05663 |
1607 | 1607.00270_arXiv.txt | \noindent Data-driven model-independent reconstructions of the dark energy equation of state $w(z)$ are presented using Planck 2015 era CMB, BAO, SNIa and Lyman-$\alpha$ data. These reconstructions identify the $w(z)$ behaviour supported by the data and show a bifurcation of the equation of state posterior in the range $1.5{<}z{<}3$. Although the concordance $\Lambda$CDM model is consistent with the data at all redshifts in one of the bifurcated spaces, in the other a supernegative equation of state (also known as `phantom dark energy') is identified within the $1.5 \sigma$ confidence intervals of the posterior distribution. To identify the power of different datasets in constraining the dark energy equation of state, we use a novel formulation of the Kullback--Leibler divergence. This formalism quantifies the information the data add when moving from priors to posteriors for each possible dataset combination. The SNIa and BAO datasets are shown to provide much more constraining power in comparison to the Lyman-$\alpha$ datasets. Further, SNIa and BAO constrain most strongly around redshift range $0.1-0.5$, whilst the Lyman-$\alpha$ data constrains weakly over a broader range. We do not attribute the supernegative favouring to any particular dataset, and note that the $\Lambda$CDM model was favoured at more than $2$ log-units in Bayes factors over all the models tested despite the weakly preferred $w(z)$ structure in the data. | \label{sec:intro} The nature of dark energy (DE) remains a significant outstanding problem in cosmology. The $\Lambda$CDM model considers a constant equation of state (EoS) parameter $w{=}{-}1$ motivated by vacuum energy. The most frequent generalisation of the $\Lambda$CDM dark energy EoS is to allow an alteration of the time-independent EoS parameter so that $w\ne{-}1$ (hereafter referred to as $w$CDM). Allowing $w$ to vary in time $w=w(z)$ results in quintessence DE models. Many quintessence models~\citep{Ratra1988,Caldwell1998,Tsujikawa2013}, including phantom DE~\citep{Caldwell2002,Sahni2004}, as well as modified GR theories~\citep{Sahni2004} make predictions for the behaviour of $w(z)$ which may be tested against cosmological datasets~\citep{PlanckCollaboration2015_DE}. Time-dependent behaviour can also be investigated by choosing equations that are simple or mathematically appealing, to test as a DE model. These phenomenological models include the CPL~\citep{Chevallier:2000qy, Linder:2002et}, JPB~\citep{Jassal2004} and FNT~\citep{Felice2012} models. Lastly, free-form approaches attempt to avoid any commitment to particular equations and instead aim to allow the observational data to define any time-dependent features in $w(z)$~\citep{Huterer2003, Zunckel2007, Zhao2008, Serra2009, Lazkoz2012, Vazquez2012}. Other free-form reconstruction methods include gaussian processes~\citep{Holsclaw2010a, Holsclaw2010b, Seikel2012}. We refer the reader to an older review by \cite{Sahni2006} which describes the general reconstruction process and new results by \cite{PlanckCollaboration2015_DE} for further reading on dark energy constraints. In this paper, we use Bayes factors combined with a `nodal' free-form method, which reconstructs a function using a spline between nodes whose amplitude and position can vary, as first proposed by~\cite{Vazquez2012c}, to investigate the constraints on $w(z)$. This approach has subsequently been used by~\cite{Vazquez2012, Aslanyan2014, PlanckCollaboration2015_inflation, Hee2015} and has the benefit of remaining general and focussing on the cosmological datasets rather than a specific model. The first aim of this paper is to investigate potential deviations from the $\Lambda$CDM constant dark energy equation of state using Bayesian model selection. The second aim is to analyse the constraining power on $w(z)$ of the datasets using the Kullback-Leibler divergence ($\DKL$). Observational data are improving in quality with many upcoming missions promising to increase our ability to understand DE models. Assessing the datasets in the manner this paper proposes provides a robust, quantitative measure of DE information that may easily be compared with past or future missions. The paper is structured as follows: We first identify the datasets and computational techniques in Section~\ref{sec:method}. An analysis of $w(z)$ constraints from Planck satellite era cosmological datasets is presented in Section~\ref{sec:wz} and the analysis of these additional datasets using the $\DKL$ approach is presented in Section~\ref{sec:dkl}. We conclude in Section~\ref{sec:conclusions}, considering the findings in relation to $\Lambda$CDM and constraints on $w(z)$ and comment on the efficacy of the techniques used for quantifying dataset constraining power and information content. | \label{sec:conclusions} We have presented a detailed Bayesian model selection analysis applied to the nodal reconstruction of $w(z)$, concluding that the Bayes factors on the Jeffreys scale `slightly favour' $\Lambda$CDM when compared to $w$CDM and `significantly disfavour' the $t$CDM, $1$CDM, $2$CDM and $3$CDM models, with an error on the Bayes factors of around $0.29$. Despite this favouring, a model averaging approach presents a bifurcation of the $P(w|z)$ plane reconstruction space which shows that, whilst $w{=}{-}1$ for all redshift is strongly favoured, a supernegative $w(z)$ equation of state at redshift $z{>}1.5$ within the $1.5 \sigma$ confidence intervals of the posterior on $w(z)$ is supported by the data. To understand this possible deviation we analysed the constraining power of the datasets using the Kullback-Leibler divergence ($\DKL$). We calculated a novel function $\DKL(z)$ to analyse the information gained when moving from the prior distribution of $w(z)$ to the posterior distribution, in slices of constant $z$, as well as a single $\DKL$-value for the whole plane. For each we used both \codeF{CosmoMC} priors and flat priors to observe information gain due to the data and the overall constraining power respectively, and we analysed each permutation of datasets using the $2$CDM model. We observed that the $BAO$ and $JLA$ datasets constrained the $w(z)$ plane much more strongly than the $\Lya$ datasets used. These two datasets had a strong peak at redshifts ${<}0.5$ whilst the $\Lya$ datasets peaked more broadly at $z{=}1$. As expected, the combination of all datasets had the greatest constraining power, specifically the $Planck$ dataset alone had $\DKL{=}0.33 \nats$, the combination with $BAO$ and $JLA$ datasets had $\DKL{=}0.82 \nats$ and the combination $Planck+BAO+JLA+\Lya$ had $\DKL{=}0.91 \nats$. The same dataset combination had a maximum information gain at redshift $0.2$ of $1.5 \nats$. Reviewing the plane reconstructions and $\DKL(z)$ functions showed that the $\Lya$ datasets provided additional constraints at $z{>}1.5$ that favours a supernegative equation of state, with $\Lambda$CDM disfavoured at $1 \sigma$ significance. Generally, many of the dataset combinations disfavoured $\Lambda$CDM at $1 \sigma$ significance around $1.5{<}z{<}2$, with higher redshifts being too poorly constrained to draw conclusions. For redshifts below $1.5$, the $\Lya$ datasets favoured a supernegative $w(z)$, the $JLA$ dataset typically agrees with $\Lambda$CDM and the $BAO$ dataset tends towards $w{>}{-}1$ values (around $1 \sigma$ significance at $z{=}0.25$). Concluding on the higher redshift deviations, we do not attribute this supernegative favouring to a particular dataset, but note that the inclusion of $\Lya$ data adds prominence as it provides a small amount of much needed constraining power over that range. In the future, the conclusions of an analysis with these techniques will strengthen as data quality improves. The nodal reconstruction has again been shown to be useful in constraining cosmological models and developing a model independent data driven analysis~\citep{Vazquez2012c,Vazquez2012,Aslanyan2014,PlanckCollaboration2015_inflation,Hee2015}. In addition, the novel formalism introduced here of the Kullback-Leibler divergence as a function of redshift provides a quantitative analysis of dataset information content applied to specific cosmological problems. Future applications of this method with upcoming mission and survey data or for forecasting with mock-data will provide useful insights into the value of datasets in constraining our cosmological models. | 16 | 7 | 1607.00270 |
1607 | 1607.02043_arXiv.txt | The radio nebula W\,50 harbours the relativistic binary system SS\,433, which is a source of the powerful wind and jets. The origin of W\,50 is wrapped in the interplay of the wind, supernova remnant and jets. The evolution of the jets on the scales of the nebula is a Rosetta stone for its origin. To disentangle the roles of these components, we study physical conditions of the jets propagation inside W\,50, and determine deceleration of the jets. The morphology and parameters of the interior of W\,50 are analyzed using the available observations of the eastern X-ray lobe, which traces the jet. In order to estimate deceleration of this jet, we devised a simplistic model of the viscous interaction, via turbulence, of a jet with the ambient medium, which would fit mass entrainment from the ambient medium into the jets of the radio galaxy 3C\,31, the well studied case of continuously decelerating jets. X-ray observations suggest that the eastern jet persists through W\,50 as hollow one, and is recollimated to the opening $\sim 30^\circ$. From the thermal emission of the eastern X-ray lobe, we determine a pressure of $P \sim 3\cdot 10^{-11}$\,erg/cm$^3$ inside W\,50. In the frame of a theory of the dynamics of radiative supernova remnants and stellar wind bubbles, this pressure in combination with other known parameters restricts W\,50's origin to a supernova happened $\sim 100\,000$\,yr ago. Also, this pressure in our entrainment model gives a deceleration of the jet by $\sim 60\%$ in the bounds of W\,50's spherical component, of radius $\sim 40$\,pc. In this case, the age of the jet should be $\ll 27\,000$\,yr so as to satisfy the sphericity of W\,50. The entrainment model comes to the viscous stress in a jet of a form $\sigma = \alpha P$, where the viscosity parameter $\alpha$ is predefined by the model. | Astrophysical jets derive a significant part of energy released in accretion disks, and influence radically their environment. The powerful jets of SS\,433 (\cite[1979]{AM79}) are so intimately interconnected with the surrounding shell-type radio nebula W\,50, that the nebula is thought to be not entirely of a supernova's origin. These jets are radiatively inefficient, that suggests their huge kinetic energy, of flux $L_\mathrm{k0} \sim 10^{39}$~erg/s, transforms into the thermal and mechanical energies of W\,50, and the latter is inflated partly by them. The revealed by \cite[(2015)]{Bor15} gamma-ray emission, located within the W\,50 area, points one more item of the jets expenses: up to 10s of percentages of the jets energy might transform into relativistic particles. Moreover, it may indicate a region where the jets decelerate and transmit energy to the nebula: in the interior of W\,50, not at the nebula shell. The role of the jets of SS\,433 in formation of the peculiar W\,50 is unknown (see \cite[(2016)]{Far16} for a comprehensive review of the origin of W\,50). The morphology of W\,50 in radio emission closely resembles an achatina, the giant African snail (\cite[1998]{Dub98}; see also Fig.~\ref{Xcone}). W\,50 looks like consisting of coils, narrowing to the tips of the nebula like a pyramid. The torque and elongated appearance of W\,50 might be due to the jets, which, however, are not explicitly resolved at the scales of W\,50 ($1^\circ \times 2^\circ$). By a convention, one discerns in W\,50 a spherical component, of radius $\sim 29^\prime$ (\cite[1998]{Dub98}), and two protrusions, the so called ears. \begin{figure*}[ht] \centering \resizebox{\hsize}{!}{\includegraphics{fig1}} \caption{ The 1--2\,keV image of the bright knot region (the XMM-Newton observatory), laid over W\,50's radio image (\cite[1998]{Dub98}), depicts the geometry of the eastern X-ray lobe in the hard emission. The images are given at the epoch J2000, north is up, east is left, SS\,433 ($\mathrm{RA} = 19^\mathrm{h}11^\mathrm{m}49.\!\!^\mathrm{s}57$, $\mathrm{Dec.} = 04^\circ 58^\prime 57.\!\!^{\prime\prime}9$) is on the right. There are indicated the spherical component of W\,50, the precession cone and its axis, ticked every $15^\prime$ from the beginning at SS\,433. The borders of \cite['s\ (1983)]{Wat83} cut, for the radial profile of brightness, are delineated by arrows. } \label{Xcone} \end{figure*} W\,50 is elongated along the jets, which are observed closely to SS\,433 as outflows in X-ray, optical and radio bands. At distances from the jets source up to $\sim 6''$, or at distances $r\lesssim 0.13$\,pc\footnote{$1'' = 0.67\cdot 10^{17}$~cm at the observer distance of $D=4.5$\,kpc, accepted hereinafter (\cite[2002]{St02}; \cite[2010]{Pan10}; \cite[2013]{Mar13}; see also \cite[2007]{Loc07}, and \cite[2014]{Pan14} for discussion on the distance); data from the literature are compiled to this distance.}, the jets show a regular precession with period $162.250$\,days, under angle $19.\!\!^{\circ}75$, around an axis whose inclination to the line of sight is $78.\!\!^{\circ}8$ (\cite[2008]{Dav08}) and position angle on the plane of the sky is $98.\!\!^{\circ}2$ (\cite[2002]{St02}). At these distances the jets are ballistic, unchanging (e.g. Fig.~2 of \cite[2010]{Rob10}), except may be the predicted 10\% deceleration and twist (the shift of the precession phase by $\sim -0.1$) in the innermost $\sim 0.\!\!^{\prime\prime}5$ (\cite[2014]{Pan14}; see also \cite[2004]{St04}). At larger distances the jets become unobservable, possibly because of weakening of the interaction of the jet with the ambient medium. The jet there would look like a hollow cone of tightly wound coils, probably blending. The cone of the precession is indicated on the radio image (\cite[1998]{Dub98}) of the eastern part of W\,50 in Fig.~\ref{Xcone} -- evidently, the cone fits the orientation and transverse size of the ear. The overlaid X-ray image\footnote{The HEASARC data archive, http://heasarc.gsfc.nasa.gov/FTP/xmm/ data/rev0//0075140401/PPS/, the EPIC PN camera of the XMM-Newton observatory.} in Fig.~\ref{Xcone} suggests that the lobe of the X-ray emission, observed at distances $> 15^\prime$ from SS\,433 (\cite[1983]{Wat83}; \cite[1996]{Bri96}), shapes the jet. However, the angular extension of the lobe in the hard X-ray band, $>1$\,keV, is much smaller than the opening $40^{\circ}$ of the precession cone. Are the jets recollimated? On the contrary, the optical filaments at the contact between the ears and sphere of W\,50 (\cite[2007, Fig.~1 and 2]{Bou07}), which possibly trace the interaction of the jet and spherical shell of W\,50, subtend an angle a little more than $40^{\circ}$ at SS\,433. The eastern lobe, which is more exposed in observations, has sharp edges in the hard X-ray band and is almost perfectly axisymmetrical and smooth, except the bright knot (Fig.~\ref{Xcone}; \cite[2007, Fig.~3 and 8]{Bri07}). The latter is probably the segment of the nearly spherical shell, heated by the jet (\cite[1983]{Wat83}). More than 30 years as the problem of the recollimation of the SS\,433 jets inspires its investigation. \cite[(1983)]{Eic83} has proposed that a precessing jet merges in a smooth hollow cone and inevitably undergoes focusing by the ambient pressure. \cite[(1990)]{Koc90} have ascertained the problems staying before the hydrodynamical simulations of Eichler's mechanism and of the pointed form of the W\,50 ears for the case of a hollow conical jet. Later, from a series of hydrodynamical simulations targeted at the ears geometry, \cite[(2011b)]{Go11b} have devised a history of the jet evolution intermittent in the speed and precession cone opening; moreover, they got the recollimation mechanism at work, although inefficient for the formation of the ears of the observed narrowness. However, there are not seen the traits of the intermittency, known for intermittent astrophysical jets: in W\,50 the X-ray lobes are rather smooth at large scales. In particular, \cite[(2011b)]{Go11b} resume that in the case of a permanency the SS\,433 jet should decelerate only at the terminal shock, i.e. in the ear. On other hand, from the non-observation of the proper motion of the radio filaments in the ears, \cite[(2011a)]{Go11a} received that the jets velocity in the ears is at least eight times smaller (for the distance $D=4.5$\,kpc) than the optical jets velocity $v_\mathrm{j0} = 0.2581c$, $c$ being the speed of light (\cite[2008]{Dav08}). Besides, the brightness of the X-ray ring-like structure at distances $\sim 60^\prime$, thought to be a terminal shock, coincident with the radio filaments in the eastern ear, is much smaller than one of the X-ray lobe at 15--$45^\prime$ from SS\,433 (\cite[2007, Fig.~1 and 2]{Bri07}). It is hardly to explain unless the jets decelerate in the interior of W\,50. Thus, the questions about the evolution of the jets of SS\,433 at W\,50's scale and their role in the inflation of W\,50 remain. This paper solves the key question: Could the jets decelerate in the interior of W\,50, before the termination in the ears? Firstly, in Sect.~\ref{PhC} we characterize the physical conditions of the jets propagation in the eastern X-ray lobe using the available X-ray observations, namely the density and pressure in the surroundings of the jet. On the basis of the found pressure, the age of the nebula W\,50 is evaluated, and the roles of a supernova remnant (SNR), SS\,433 system's wind, and the jets in the nebula formation are clarified in Sect.~\ref{Age}. Encouraged by the studies of the decelerating relativistic jets of the Fanaroff-Riley class I radio galaxies, in Sect.~\ref{Entr1} we construct the model which would fit mass entrainment from the environment into the jets of the radio galaxy 3C\,31, the well known case of continuously decelerating jets. The entrainment model is applied to the SS\,433 jet, and results on the jet behaviour at scales of dozens parsecs are presented in Sect.~\ref{SS}. The conclusions are given in Sect.~\ref{Iss}. | \label{Iss} The morphology of the eastern X-ray lobe in W\,50 evidences a continuous proceeding of SS\,433's jet through the nebula on scales of dozens parsecs. At that, the jet is found to be recollimated from the opening $40^\circ$ to the $\sim 30^\circ$. The X-ray brightness distributions in the soft and hard energy bands of the XMM-Newton observations are consistent with the hollow structure of the jet. This picture suggests that the jet bypasses the X-ray brightest knot, which resides at the place of the shell existed even before the jet. From the X-ray observations presented in (\cite[1983]{Wat83}; \cite[1996]{Bri96}; \cite[2007]{Bri07}), we derived an electron density of $\sim 0.1$\,cm$^{-3}$ and a pressure of $\sim 3 \cdot 10^{-11}$\,erg/cm$^3$ for the thermal gas in the spherical component of W\,50. We note that only the XMM-Newton observations have allowed \cite[(2007)]{Bri07} to separate surely the thermal components of the X-ray emission of the bright knot and, less surely, of the region to the west of the knot. However, they have not characterized the physical conditions in the X-ray lobe in whole. Solely a model of SNR dynamics predicts the observed pressure in conjunction with the reliable ages of SS\,433 10\,000--100\,000\,yr (\cite[2000]{Kin00}), not a model of stellar wind bubble. Though, this is not conclusive in the case of an inhomogeneous ambient medium, when the wind origin of W\,50 is possible too (\cite[1983]{Kon83}). Such small observed pressure pinpoints W\,50 as an old SNR of age $\sim 100\,000$\,yr, which is obtained accounting for radiative losses of the SNR in the pressure driven phase, unlike all previous estimations of W\,50's age, based on the Sedov model. In the case of an inhomogeneous interstellar medium, the age would be smaller. An influence of the jets on the spherical component would be restricted to prevent a large distortion of the sphericity, therefore the jets exhaust most of their energy into the ears or are much younger than the SNR. The thermal energy content in the W\,50 nebula is $\sim (\Gamma-1) P_\mathrm{W50} (V_\mathrm{SNR} + V_\mathrm{E} + V_\mathrm{W}) = 4.4\cdot 10^{50}$\,erg, where $V_\mathrm{SNR} = 6.6\cdot 10^{60}$\,cm$^3$, $V_\mathrm{E} = 2.4\cdot 10^{60}$\,cm$^3$, and $V_\mathrm{W} = 6.9\cdot 10^{59}$\,cm$^3$ are the volumes of the spherical component, the eastern and western ears, respectively. Supposedly, the same amount of energy is contained in the kinetic energy of the nebula's shell, then the total energy of W\,50 is $\sim 10^{51}$\,erg. The simulations of \cite[(2011b)]{Go11b} provide an important view that the jet-SNR interaction doesn't influence the expansion of W\,50's spherical shell, and the jets are a relatively young phenomenon, $\lesssim 20\,000$\,yr, the age they give to the nebula W\,50. The latter issues naturally from an amount of the jets momentum imparted to the ears, and hence from the ears size. Besides clarification of the role of the jets in the origin of W\,50, this justifies the application of theories of the spherical wind bubbles and SNRs to the evolution of W\,50. However, so small age of W\,50's spherical component, obtained in the simulations, contradicts not only to our result, based on a theory for the shock dynamics of SNRs, but also to the simulations on SNRs elsewhere (e.g. \cite[2015]{Kim15}; \cite[2015]{Kor15}). \cite['s\ (2011b)]{Go11b} simulations invoke the intermittent activity of the jets for the observed shape of the ears. At that, the continuous activity, coupled with the recollimation and deceleration, appeals more urgently as the observed non-unique properties of the broad class of jets, of radio galaxies of Fanaroff-Riley Class I (\cite[2014]{Lai14}). There seems the sameness between them and SS\,433's jets: the pressure jump in the surroundings (afforded possibly by the wind in the case of W\,50), beyond which the jets brighten, recollimate and decelerate (the latter two are obvious only for FR\,I jets), and the faintness of terminal shock at the jet head (contrary to the prominent hotspots of FR\,II jets) -- the characteristic of decelerated jets. Indeed, the derived pressure of W\,50's interior turns out to be so large as to cause deceleration of the jet via the viscous jet-surroundings interaction. There is as yet no theory of this interaction for relativistic astrophysical jets. The studies of FR\,I jets suggest that these jets decelerate continuously via entrainment of the ambient medium. We devised the model of the entrainment, which successfully fits the semiempirical profile of entrainment rate for the jets of the radio galaxy 3C\,31, found by \cite[(2009)]{Wan09}. By this model, the entrainment results from turbulence at the jet surface, that exerts in effect a viscous tension in the jet $\sigma = \alpha P$, of the well known form in the theory of accretion disks, where the viscosity parameter $\alpha$ is defined by the model. This model predicts that the eastern jet of SS\,433 decelerates by $\sim 60$\% in bounds of the circular shell of W\,50. Thus, the decelerated jet would inject $\sim$ a half of its energy into W\,50's spherical component. For the sphericity, the age of the jets, Eq.\,(\ref{tjet}), should be much smaller than 27\,000\,yr, that is consistent with the age delimitation in (\cite[2011b]{Go11b}). Also, there seems to be a balance between the jet power and the luminosity of the X-ray lobe. During tentative lifetime $t_\mathrm{j} \sim 10^4$\,yr the jet has put approximately a half the ejected energy $W_\mathrm{j} = L_\mathrm{k0} t_\mathrm{j} \sim 3\cdot 10^{50}$\,erg in the expansion of W\,50, another part has been deposited into the thermal energy of the interior gas. While an estimation $W_\mathrm{th} = L_\mathrm{th} t_\mathrm{th} \sim 3\cdot 10^{49}$\,erg, where $L_\mathrm{th} \sim 10^{35}$\,erg/s is the thermal luminosity of the X-ray lobe (Sect.~\ref{PhC}), and $t_\mathrm{th} \sim kT/n_\mathrm{e} \Lambda \sim 10^7$\,yr the time of the radiative cooling of the lobe, gives for the thermal energy $\sim 1/10$ of the $W_\mathrm{j}$ vs. the above proportion $\sim 1/2$. The discrepancy by a factor of 5 is possibly attributable to roughness of our estimations. Further studies of the distribution of physical parameters in the nebula W\,50, in particular on the basis of the available X-ray data, would improve the above picture of the evolution of SS\,433's jets on scales of dozens parsecs. | 16 | 7 | 1607.02043 |
1607 | 1607.07450_arXiv.txt | Binary neutron star mergers are promising sources of gravitational waves for ground-based detectors such as Advanced LIGO. Neutron-rich material ejected by these mergers may also be the main source of r-process elements in the Universe, while radioactive decays in the ejecta can power bright electromagnetic post-merger signals. Neutrino-matter interactions play a critical role in the evolution of the composition of the ejected material, which significantly impacts the outcome of nucleosynthesis and the properties of the associated electromagnetic signal. In this work, we present a simulation of a binary neutron star merger using an improved method for estimating the average neutrino energies in our energy-integrated neutrino transport scheme. These energy estimates are obtained by evolving the neutrino number density in addition to the neutrino energy and flux densities. We show that significant changes are observed in the composition of the polar ejecta when comparing our new results with earlier simulations in which the neutrino spectrum was assumed to be the same everywhere in optically thin regions. In particular, we find that material ejected in the polar regions is less neutron rich than previously estimated. Our new estimates of the composition of the polar ejecta make it more likely that the color and timescale of the electromagnetic signal depend on the orientation of the binary with respect to an observer's line-of-sight. These results also indicate that important observable properties of neutron star mergers are sensitive to the neutrino energy spectrum, and may need to be studied through simulations including a more accurate, energy-dependent neutrino transport scheme. | \label{sec:intro} The detection of gravitational waves from binary black hole mergers by Advanced LIGO~\cite{Abbott:2016blz,Abbott:2016nmj} just opened a new window through which to observe the universe. In the coming years, Advanced LIGO~\cite{aLIGO2} and its European and Japanese counterparts, Advanced VIRGO~\cite{aVirgo2} and KAGRA~\cite{kagra}, are expected to detect neutron star-neutron star (NSNS) and neutron star-black hole (NSBH) mergers~\cite{Abadie:2010cfa}. Compact binary mergers in the presence of at least one neutron star could put strong constraints on the equation of state of nuclear matter in the extreme conditions reigning in the core of neutron stars~\cite{Read2009b,DelPozzo:13,Lackey2014}. They are also likely to be the progenitors of short gamma-ray bursts~\cite{moch:93,LK:98,Janka1999,Fong2013}, and are followed by bright radioactively powered optical/infrared~\cite{1976ApJ...210..549L,Li:1998bw,Roberts2011,Kasen:2013xka,Tanaka:2013ana} and radio transients~\cite{Nakar:2011cw,Hotokezaka:2016} which could provide us with useful information about the merging objects and their environment. Finally, matter ejected during a neutron star merger is a prime candidate for the so-far unknown site of r-process nucleosynthesis, where many heavy elements (e.g. gold, uranium) are produced~\cite[e.g.][]{korobkin:12,Wanajo2014}. Numerical simulations are a critical tool to understand the gravitational wave and electromagnetic signals powered by NSBH and NSNS mergers. Yet, the complexity of the physical processes playing an important role in these mergers places significant limitations on the realism of existing simulations. General relativity, magnetohydrodynamics, neutrino radiation, and nuclear physics all influence at least some important observables of these systems. Considering the high computational cost of including each of these components, simulations have generally focused on a subset of these physical processes, either by improving the microphysics with approximate treatments of gravity, or using full general relativity with much simpler physics (see~\cite{Duez:2009yz,faber:12,2014ARA&A..52..661L,BaiottiReview2016} for reviews of numerical simulations). Recent general relativistic simulations have only begun to partially resolve the effects of magnetic fields~\cite{Rezzolla:2011da,Kiuchi2014,Kiuchi2015,2016arXiv160402455R,2015arXiv150202021D}, to include approximate treatments of the neutrinos and better equations of state for dense matter~\cite{Sekiguchi:2011zd,Neilsen:2014hha,Sekiguchi:2015,Foucart:2015,2016arXiv160300501L,RadiceLeak:2016,Sekiguchi:2016}, or both (with sub-grid models for the growth of magnetic fields)~\cite{Palenzuela2015}. In this paper, we focus on the treatment of the neutrinos and their impact on the post-merger properties of NSNS mergers. Neutrinos are particularly important as the main source of cooling in the post-merger remnant. They also play a critical role in setting the relative number of neutrons and protons in the remnant and in the material ejected from the system. The composition of the fluid is needed to predict the properties of optical/infrared transients powered by r-process nucleosynthesis in material ejected by the merger~\cite{Kasen:2013xka,2013ApJ...775...18B}, as well as the relative abundances of the r-process elements produced in the ejecta~\cite{Wanajo2014,Lippuner2015}. Finally, neutrinos can drive strong winds from the post-merger remnant~\cite{Dessart2009,Just2014,Perego2014,Sekiguchi:2015,FoucartM1:2016}. Neutrinos were first included in general relativistic simulations of neutron star mergers through a simple gray (i.e. energy-integrated) leakage scheme~\cite{Sekiguchi:2011zd}, based on approximate methods developed for Newtonian simulations~\cite{1997A&A...319..122R,Rosswog:2003rv}. A leakage scheme uses the local properties of the fluid and an estimate of the neutrino optical depth to determine the amount of energy lost locally to neutrino-matter interactions, and the associated change in the composition of the fluid. Leakage schemes provide an order-of-magnitude accurate estimate of neutrino cooling in the post-merger remnant, and have thus been used to capture the first-order effect of neutrino-matter interactions in general relativistic simulations of compact binary mergers~\cite{Sekiguchi:2011zd,Deaton2013,Foucart:2014nda,Neilsen:2014hha,Palenzuela2015,2016arXiv160300501L,RadiceLeak:2016}. In most implementations, they do not account for irradiation of low-density regions by neutrinos emitted from hot, dense regions. This potentially leads to large errors in the composition of the outflows, mostly by underestimating the number of protons~\cite{FoucartM1:2015,FoucartM1:2016}. Accordingly, the simplest leakage schemes are very inaccurate when attempting to predict the properties of post-merger electromagnetic signals. More complex leakage schemes have been developed to attempt to include neutrino absorptions in low-density regions, either by assuming propagation of the neutrinos along the radial direction~\cite{RadiceLeak:2016}, or through a more expensive global procedure (only used in Newtonian physics so far) to estimate where neutrinos are transported~\cite{Perego:2016}. The latter scheme also includes a discretization of the neutrinos in energy space. The only general relativistic simulations going beyond neutrino leakage use a moment formalism with an analytic closure to approximate the Boltzmann equation~\cite{1981MNRAS.194..439T,shibata:11}. In particular, neutron star merger simulations have been performed with a gray M1 scheme~\cite{Sekiguchi:2015,FoucartM1:2015,FoucartM1:2016,Sekiguchi:2016}, in which the energy density and flux density of each neutrino species are evolved. In NSNS mergers, the use of this moment formalism showed that a wide range of compositions, and thus of nucleosynthesis outcomes, exists in the material ejected by the merger~\cite{Wanajo2014}. Comparisons with leakage schemes for BHNS~\cite{FoucartM1:2015} and NSNS~\cite{FoucartM1:2016} mergers clearly show that irradiation of the polar outflows by neutrinos emitted by the post-merger remnant causes those outflows to be significantly less neutron-rich than predicted by a leakage scheme which does not account for neutrino absorption. The gray M1 scheme is far from perfect. One obvious limitation is the impact of the analytical closure, which causes unphysical ``shocks'' in regions in which neutrinos converge. This occurs in the polar regions of post-merger remnants, putting into question the accuracy with which we can recover the composition of the polar outflows in those systems. Another limitation is the lack of information about the energy spectrum of the neutrinos, or even their average energy. In~\cite{FoucartM1:2015,FoucartM1:2016}, for example, neutrinos in optically thick regions are assumed to be in equilibrium with the fluid, which is reasonable, but neutrinos in optically thin regions are assumed to follow everywhere a blackbody spectrum with a temperature determined from the average properties of the neutrino radiation predicted by the simpler leakage scheme. This neglects potentially important spatial variations in the neutrino spectrum, deviations from a blackbody spectrum, and the effects of relativistic beaming on the neutrino energies. These approximations could easily affect our ability to predict the composition of the ejected material, as many neutrino-matter cross-sections scale as the square of the neutrino energy. Additionally, the transport method used in~\cite{FoucartM1:2015,FoucartM1:2016} does not guarantee conservation of the total lepton number. Performing a full merger simulation with an energy-dependent transport scheme, even in the relatively simple M1 approximation, is too costly with our current code. In this paper, we take an alternative route to assess the impact of some of the missing information about the neutrino energies. In addition to the neutrino energy density and flux density, we now evolve the neutrino number density. This does not provide us with any information about the shape of the neutrino spectrum, but does provide a local estimate of the average neutrino energy, and accounts for relativistic beaming. By evolving the neutrino number density, we can also guarantee conservation of the total lepton number. We consider in particular a low-mass neutron star merger ($1.2M_\odot-1.2M_\odot$) already studied with our previous M1 and leakage schemes~\cite{FoucartM1:2016} (hereafter Paper I), to facilitate comparisons. We show that relativistic beaming, spatial variations in the average neutrino spectrum, and an improved treatment of the diffusion rate of the neutrino number density can play a significant role in the composition of the ejected material and of the post-merger remnant. We organize the paper as follow. The numerical methods and detail of the improved M1 scheme are provided in Appendix~\ref{sec:form}. The physical system under consideration and initial conditions are discussed in Sec.~\ref{sec:setup}. The impact of the neutrino scheme on the emitted neutrino radiation is presented in Sec.~\ref{sec:neutrinos}. Finally, we discuss consequences on the properties of the ejected material and associated electromagnetic signal in Sec.~\ref{sec:outflows}, and summarize our results in Sec.~\ref{sec:conclusions}. | \label{sec:conclusions} We have presented a detailed study of a NSNS merger with a general relativistic hydrodynamics code and two variations of an approximate, gray neutrino transport scheme. We considered in particular the impact of the method used to approximate the neutrino energy spectrum on the post-merger evolution of the system. In previous simulations (Paper I), we estimated the average neutrino energy in all optically thin regions using a single neutrino temperature for each neutrino species, taken from the prediction of a simple leakage scheme~\cite{FoucartM1:2015,FoucartM1:2016}. In this work, we instead evolved the neutrino number and energy density, and used those evolved variables to estimate a spatially-varying average neutrino energy. The new scheme has the advantages of exactly conserving the total lepton number, taking into account spatial variations in the neutrino energy, and being sensitive to the impact of relativistic beaming on the average neutrino energies. It generally predicts higher neutrino energies, particularly immediately after merger and in the polar regions, and neutrino luminosities differing by $\lesssim 40\%$. These differences do not appear to affect the dynamics of the post-merger remnant, or to have a significant impact on its temperature. However, they do have important consequences for the evolution of the composition of the fluid. Material unbound in the polar regions as a neutrino-driven wind absorbs fewer electron antineutrinos and more electron neutrinos when using local estimates of the average neutrino energy. This robustly drives the electron fraction of the polar ejecta to values $Y_e \gtrsim 0.25$, with an increase of $\Delta Y_e \sim 0.05-0.1$ with respect to results using a global estimate of the average neutrino energy. The low-density, bound regions of the remnant also see an increase in their electron fraction. Such a change in the average electron fraction of the polar ejecta could have important consequences for the observable electromagnetic counterpart of the merger due to r-process nucleosynthesis in the ejecta. In the absence of neutron-rich ejecta in the polar regions, the opacity of the ejecta along the line of sight of an observer viewing the merger face-on could be significantly reduced. This makes it possible to observe electromagnetic transients peaking in the optical when the merger is observed face-on, particularly if high-$Y_e$ disk winds continue to be ejected by the post-merger remnant over timescales significantly longer than the duration of our simulation~\cite{Just2014,Perego2014,Metzger2014,Fernandez:2014,Fernandez:2014b}. Although we believe that the new methods presented here provide a more accurate representation of the merger than our previous results (Paper I), the strong dependence of the polar electron fraction on the method used to estimate the average neutrino energy may offer us a first view of the limits of current gray neutrino transport schemes. After all, even our improved estimate of the neutrino spectrum remains fairly rudimentary. The difference between the composition of the ejecta in the two neutrino transport schemes is only slightly smaller than the difference observed in Paper I between a leakage scheme ignoring neutrino absorption and the neutrino transport schemes. It may thus be useful to obtain better predictions for the neutrino energy spectrum in the polar regions. We should also note that the moment formalism used here to approximate the neutrino distribution function is notoriously problematic in regions in which radiation beams emitted from different directions cross paths. This is obviously the case in large parts of the polar regions, where most of the neutrino-matter interactions that drive up the electron fraction of the wind take place. Even if the neutrino-matter interactions are reasonably well approximated by the current scheme in a volume-averaged sense, the exact impact of the moment formalism on the neutrino emission and properties of the outflows remain an important question for future studies of binary neutron star mergers. | 16 | 7 | 1607.07450 |
1607 | 1607.06937_arXiv.txt | We report the optical spectroscopy of four young radio sources which are observed with the Lijiang 2.4m telescope. The Eddington ratios of these sources are similar with those of narrow-line Seyfert 1 galaxies (NLS1s). Their Fe {\sc ii} emission is strong while [O {\sc iii}] strength is weak. These results confirm the NLS1 features of young radio sources, except that the width of broad H$\beta$ of young radio sources is larger than that of NLS1s. We thus suggest that the young radio sources are the high black hole mass counterparts of steep-spectrum radio-loud NLS1s. In addition, the broad H$\beta$ component of \astrobj{4C 12.50} is the blue wing of the narrow component, but not from the broad line region. | % \label{sect:intro} The young radio sources, including compact steep spectrum (CSS) and gigahertz peaked spectrum (GPS) radio sources, are believed to represent the earliest stages in the evolution of the powerful radio galaxies \citep{1998PASP..110..493O}. Some of them are found to be associated with galaxies mergers \citep{1986Natur.321..750G} or ultraluminous infrared galaxies (ULIRGs) \citep{2011MNRAS.410.1527H, 2012MNRAS.422.1453N}. Recent observations also manifest that some radio-loud narrow-line Seyfert 1 galaxies/quasars (NLS1s) have the radio properties consistent with CSS radio sources \citep{2010AJ....139.2612G, 2014MNRAS.441..172C, 2015arXiv151005584L}. NLS1s are classified based on their narrow Balmer lines (with the full width at half maximum, FWHM $< 2000~km~s^{-1}$), small ratio between [O {\sc iii}]$\lambda$5007 and H$\beta$ ([O {\sc iii}]$\lambda$5007/H$\beta$ $<$ 3) and strong emission of Fe {\sc ii} complexes (Fe {\sc ii} $\lambda$4570/H$\beta$ $>$ 0.5, \citealt{1985ApJ...297..166O, 2001A&A...372..730V}). These features are explained as a result of relatively small mass of the central black hole with high accretion rate (\citealt{2002ApJ...565...78B, 2012AJ....143...83X} , but also see \citealt{2008MNRAS.386L..15D}). Thus NLS1s are suggested to be during the early stage of the accretion activities. The accretion rates of young radio sources are found to be relatively high, with the average value similar with NLS1s \citep{2009MNRAS.398.1905W, 2016ApJ...818..185F}. Thus the young radio sources can also stand during the early stage of accretion activities. Moreover, there is another similar feature between the young radio sources and NLS1s. That is the blue wing of narrow [O {\sc iii}] \citep{2005MNRAS.364..187B, 2008MNRAS.387..639H, 2009MNRAS.398.1905W, 2009MNRAS.400..589H}. This feature is always explained as the outflow originated from the disk wind or galactic wind. Some results indicate that the strength of the blueshift is related to the Eddington ratio \citep{2008ApJ...680..926K}, while other explanations refer to the jet - ISM interaction \citep{2008MNRAS.387..639H}. Although the common radio properties between NLS1s and young radio sources have been discussed frequently in the literature, the common optical properties between them are less discussed. In this paper, we obtain the optical spectra of four young radio sources, and explore their emission lines properties. Throughout this paper, the luminosity is calculated using a $\Lambda$CDM cosmology model with h=0.71, $\Omega_{m}$=0.27, $\Omega_{\Lambda}$=0.73. | \label{sect:discussion} We estimate the black hole mass of three type 1 sources. The calculations for H$\beta$ follow the relation of single epoch reverberation mapping in \citet{2006ApJ...641..689V}, \begin{equation} log M_{BH} = 2~log(\frac{FWHM (H\beta)}{1000~km~s^{-1}}) + 0.5~log(\frac{\lambda L_{\lambda}(5100)}{10^{44}~erg~s^{-1}}) + 6.91 . \end{equation} For Mg {\sc ii}, we use the relation in \citet{2009ApJ...707.1334W}, \begin{equation} log M_{BH} = 1.51~log (\frac{FWHM (Mg~{\sc II})}{1000~km~s^{-1}}) + 0.5~log(\frac{\lambda L_{\lambda}(3000)}{10^{44}~erg~s^{-1}}) + 7.13 . \end{equation} Then we estimate the Eddington ratio $L_{bol}/L_{Edd} = 9 L_{5100}/(1.3\times 10^{38} M_{BH})$ \citep{2000ApJ...533..631K}. The results are listed in Table \ref{tab2}. \begin{table*} \caption[]{The black hole mass and Eddington ratio. The $\lambda$5100 luminosity of \astrobj{GB6 J0140+4024} is calculated from $\lambda$3000 with the spectral index -1.65. The black hole mass and Eddington ratio of \astrobj{4C 12.50} are taken from \citet{2006ApJ...638..745D} and \citet{2011MNRAS.410.1527H}. \label{tab2}} \vspace{-1mm}\footnotesize \begin{center}\doublerulesep 0.1pt \tabcolsep 0.5pt \begin{tabular}{cccccc} \hline\noalign{\smallskip} Source Name & redshift & $L_{5100}$ ($erg~s^{-1}$) & Log$M_{BH}$ & $L_{bol}/L_{Edd}$ \\ \hline\noalign{\smallskip} \astrobj{GB6 J0140+4024} & 1.62 & $2.58\times 10^{45}$ & 8.88 & 0.23 \\ \astrobj{TXS 0942+355} & 0.208 & $6.28\times 10^{43}$ & 7.66 & 0.10 \\ \astrobj{IRAS 11119+3257} & 0.189 & $1.51\times 10^{45}$ & 7.63 & 2.45 \\ \astrobj{4C 12.50} & 0.122 & --- & 7.82 & 0.27 \\ \noalign{\smallskip}\hline \end{tabular} \end{center} \end{table*} The estimated average values of Eddiongton ratios for NLS1s are from 0.15 \citep{2008MNRAS.390..752B} to 0.79 \citep{2012AJ....143...83X}, while the broad line Seyfert 1 galaxies (BLS1) have the average value about 0.16 \citep{2012AJ....143...83X}. The Eddington ratios of four young radio sources are generally larger than BLS1 and distributes in the range of NLS1s. The three type 1 sources all show strong Fe {\sc ii} emissions, and [O {\sc iii}] emission is weak for \astrobj{TXS 0942+355} and \astrobj{IRAS 11119+3257}. These features confirm that the emission lines properties of young radio sources are similar with NLS1s, except that the line width of young radio sources is broader than that of NLS1s. The estimated black hole mass is also larger than the average value of NLS1s \citep{2012AJ....143...83X}. Meanwhile, the radio powers of the compact steep-spectrum NLS1s are at the low end of that of young radio sources \citep{2010AJ....139.2612G, 2014MNRAS.441..172C, 2015arXiv151005584L}, but locate in the range of the low-luminosity compact radio sources ($P_{1.4GHz} < 10^{26}~W~Hz^{-1}$, \citealt{2010MNRAS.408.2261K}). Therefore, we suggest the young radio sources are the high mass counterparts of the steep-spectrum radio-loud NLS1s. The blue wing of [O {\sc iii}] is a prominent feature in NLS1s \citep{2005MNRAS.364..187B, 2008ApJ...680..926K}. Among the four sources, the blue wings of [O {\sc iii}] are presented in \astrobj{IRAS 11119+3257} and \astrobj{4C 12.50}. Both sources are ULIRGs, and with relatively high Eddiongton ratio, which links the outflow mechanism to high star formation rate or high accretion rate. | 16 | 7 | 1607.06937 |
1607 | 1607.04874_arXiv.txt | We investigate the cosmological implications of modified gravities induced by the quantum fluctuations of the gravitational metric. If the metric can be decomposed as the sum of the classical and of a fluctuating part, of quantum origin, then the corresponding Einstein quantum gravity generates at the classical level modified gravity models with a nonminimal coupling between geometry and matter. As a first step in our study, after assuming that the expectation value of the quantum correction can be generally expressed in terms of an arbitrary second order tensor constructed from the metric and from the thermodynamic quantities characterizing the matter content of the Universe, we derive the (classical) gravitational field equations in their general form. We analyze in detail the cosmological models obtained by assuming that the quantum correction tensor is given by the coupling of a scalar field and of a scalar function to the metric tensor, and by a term proportional to the matter energy-momentum tensor. For each considered model we obtain the gravitational field equations, and the generalized Friedmann equations for the case of a flat homogeneous and isotropic geometry. In some of these models the divergence of the matter energy-momentum tensor is non-zero, indicating a process of matter creation, which corresponds to an irreversible energy flow from the gravitational field to the matter fluid, and which is direct consequence of the nonminimal curvature-matter coupling. The cosmological evolution equations of these modified gravity models induced by the quantum fluctuations of the metric are investigated in detail by using both analytical and numerical methods, and it is shown that a large variety of cosmological models can be constructed, which, depending on the numerical values of the model parameters, can exhibit both accelerating and decelerating behaviors. | Modified gravity theories may provide an attractive alternative to the standard explanations of the present day observations that have shaken the well-established foundations of the theoretical physics. Astronomical observations have confirmed that our Universe does not according to standard general relativity, as derived from the Hilbert-Einstein action, $S=\int{\left(-R/2\kappa ^2+L_m\right)\sqrt{-g}d^4x}$, where $R$ is the Ricci scalar, $\kappa $ is the gravitational coupling constant, and $L_m$ is the matter Lagrangian, respectively. Extremely successful on the Solar System scale, somehow unexpectedly, general relativity faces, on a fundamental theoretical level, two important challenges, the dark energy and the dark matter problems, respectively. Several high precision astronomical observations, with the initial goal of improving the numerical values of the basic cosmological parameters by using the properties of the distant Type Ia Supernovae, have provided the result that the Universe underwent recently a transition to an accelerating, de Sitter type phase \cite{1n,2n,3n,4n,acc}. The necessity of explaining the late time acceleration lead to the formulation of a new paradigm in theoretical physics and cosmology, which postulates that the explanation of the late time acceleration is the existence of a mysterious component of the Universe, called dark energy (DE), which can describe the late time dynamics of the Universe \cite{PeRa03,Pa03}, and can explain all the observed features of the recent (and future) cosmological evolution. However, in order to close the matter-energy balance of the Universe, a second, and equally mysterious component, called Dark Matter, is required. Dark Matter, assumed to be a non-baryonic and non-relativistic (cold) component of the Universe, is necessary for explaining the dynamics of the hydrogen clouds rotating around galaxies, and having flat rotation curves, as well as the virial mass discrepancy in clusters of galaxies \cite{dm1,dm2}. The direct detection/observation of the dark matter is extremely difficult due to the fact that it interacts only gravitationally with the baryonic matter. After many decades of intensive observational and experimental efforts there is no direct evidence on the particle nature of the dark matter. One of the best theoretical descriptions that fits almost perfectly the observational data is based on the simplest theoretical extension of general relativity, which includes in the gravitational field equations the cosmological constant $\Lambda $ \cite{Wein,Wein1}. Based on this theoretical formalism the basic paradigm of modern cosmology has been formulated as the $\Lambda $CDM-$\Lambda $ Cold Dark Matter Model, in which the dark energy is nothing but the simple constant introduced almost one hundred years ago by Einstein. Even that the $\Lambda $CDM model fits the data well, it raises some fundamental theoretical questions about the possibility of explaining it. There is no theoretical explanation for the physical/geometrical nature of $\Lambda $, and, moreover, general relativity cannot give any hints on why it is so small, and why it is so fine tuned \cite{Wein,Wein1}. Therefore, the possibility that dark energy can be explained as an intrinsic property of a {\it generalized gravity theory}, going beyond general relativity, and its Hilbert-Einstein gravitational action, cannot be rejected a priori. In this context a large number of modified gravity models, all trying to extend and generalize the standard Einsteinian theory of gravity, have been proposed. Historically, in going beyond Einstein gravity, the first, and most natural step, was to extend {\it the geometric part of the Hilbert-Einstein action}. One of the first attempts in this direction is represented by the $f(R)$ gravity theory, in which the gravitational action is generalized to be an arbitrary function of the Ricci scalar, so that $S = \frac{1}{2\kappa ^2}\int {f(R) \sqrt{-g}d^4 x} +\int {L_m \sqrt{-g}d^4 x}$ \cite{Bu701,re4,Bu702,Bu703,re1,Bu704,Fel}. However, this as well as several other modifications of the Hilbert-Einstein action focussed only on the geometric part of the gravitational action, by explicitly postulating that the matter Lagrangian plays {\it a subordinate and passive role} only, as compared to the geometry \cite{Mat}. From a technical point of view such an approach implies a minimal coupling between matter and geometry. But a fundamental theoretical principle, forbidding an arbitrary coupling between matter and geometry, has not been formulated yet, and perhaps it may simply not exist. On the other hand, if general matter-geometry couplings are introduced, a large number of theoretical gravitational models, with extremely interesting physical and cosmological properties, can be easily constructed. The first of the modified gravity theory with arbitrary geometry-matter coupling was the $f\left( R,L_{m}\right) $ modified gravity theory \cite{fL1,fL2,fL3,fL4}, in which the total gravitational action takes the form $S = \frac{1}{2\kappa ^2}\int {f\left(R,L_m\right) \sqrt{-g}d^4 x} $. In this kind of theories matter is essentially indistinguishable from geometry, and plays an active role in generating the geometrical properties of the space-time. A different geometry-matter coupling is introduced in the $f(R,T)$ \cite{fT1,fT2} gravity theory, where matter and geometry are coupled via $T$, the trace of the energy-momentum tensor. The gravitational action of the $f(R,T)$ theory is given by $S=\int{\left[f(R,T)/2\kappa ^2+L_m\right]\sqrt{-g}d^4x}$. A recent review of the generalized $f\left(R,L_{m}\right)$ and $f(R,T)$ type gravitational theories with non-minimal curvature-matter couplings can be found in \cite{Revn}. Several other gravitational theories involving geometry-matter couplings have also been proposed, and extensively studied, like, for example, the Weyl-Cartan-Weitzenb\"{o}ck (WCW) gravity theory \cite{WCW}, hybrid metric-Palatini $% f(R,\mathcal{R})$ gravity \cite{HM1,HM2,Revn1}, where $\mathcal{R}$ is the Ricci scalar formed from a connection independent of the metric, $% f\left(R,T,R_{\mu \nu }T^{\mu \nu }\right)$ gravity theory, where $R_{\mu \nu } $ is the Ricci tensor, and $T_{\mu \nu }$ the matter energy-momentum tensor, respectively \cite{Har4,Odin}, or $f(\tilde{T},\mathcal{T})$ gravity \cite{HT}, in which a coupling between the torsion scalar $\tilde{T}$, essentially a geometric quantity, and the trace $T$ of the matter energy-momentum tensor is introduced. Gravitational models with higher derivative matter fields were investigated in \cite{HDM}. One of the interesting (and intriguing) properties of the gravitational theories with geometry-matter coupling is the non-conservation of the matter energy-momentum tensor, whose four-divergence is usually different of zero, $\nabla _{\mu}T^{\mu \nu}\neq 0$. This property can be interpreted, from a thermodynamic point of view, by using the formalism of open thermodynamic systems \cite{fT2}. Hence one can assume that the generalized conservation equations in these gravitational theories describe {\it irreversible matter creation processes}. Thus the non-conservation of the energy-momentum tensor describes an irreversible energy flow from the gravitational field to the newly created matter constituents, with the second law of thermodynamics requiring that space-time transforms into matter. In \cite{fT2} the equivalent particle number creation rates, the creation pressure and the entropy production rates were obtained for both $f\left(R,L_m \right)$ and $f(R,T)$ gravity theories. The temperature evolution laws of the newly created particles was also obtained, and studied. Due to the non-conservation of the energy-momentum tensor, which is a direct consequence of the geometry--matter coupling, during the cosmological evolution of the Universe a large amount of comoving entropy could be also produced. The prediction of the production of particles from the cosmological vacuum is one of the remarkable results of the quantum field theory in curved space-times \cite{P1,Z1,P2,Full,P3}. Particles creation processes are supposed to play a fundamental role in the quantum field theoretical approaches to gravity, where they naturally appear. It is a standard result of quantum field theory in curved spacetimes that quanta of the minimally-coupled scalar field are created in the expanding Friedmann-Robertson-Walker Universe \cite{P3}. Therefore, the presence of particle creation processes in both quantum theories of gravity and modified gravity theories with geometry-matter coupling may suggest that a deep connection between these two, apparently very different physical theories, may exist. And, interestingly enough, such a connection has been found in \cite{re11}, where it was pointed out that by using a nonperturbative approach for the quantization of the metric, proposed in \cite{re8,re9,re10}, a particular type of $f(R,T)$ gravity, with Lagrangian given by $L=\left[(1-\alpha )R/2\kappa ^2+\left(L_m-\alpha T/2\right)\right]\sqrt{-g}$, where $\alpha $ is a constant, naturally emerges as a result of the quantum fluctuations of the metric. This result suggests that an equivalent microscopic quantum description of the matter creation processes in $f(R,T)$ or $f\left(R,L_m\right)$ gravity is possible, and such a description could shed some light on the physical mechanisms leading to particle generation via gravity and matter geometry coupling. Such mechanisms do indeed exist, and they can be understood, at least qualitatively, in the framework of some quantum/semiclassical gravity models. It is the goal of the present paper to further investigate the cosmological implications of modified gravities induced by the quantum fluctuations of the gravitational metric, as initiated in \cite{re11,re8,re9,re10}. As a starting point we assume that a general quantum metric can be decomposed as the sum of the classical and of a fluctuating part, the latter being of quantum (or stochastic) origin. If such a decomposition is possible, the corresponding Einstein quantum gravity generates at the classical level modified gravity models with a nonminimal interaction between geometry and matter, as previously considered in \cite{fL1, fT1,Har4}. After assuming that the expectation value of the quantum correction can be generally expressed in terms of an arbitrary second order tensor $K_{\mu \nu}$, which can be constructed from the metric and from the thermodynamic quantities characterizing the matter content of the Universe, we derive from the first order quantum gravitational action the (classical) gravitational field equations in their general form. We analyze in detail the cosmological models obtained from the quantum fluctuations of the metric in tqo cases. First we assume that the quantum correction tensor $K_{\mu \nu}$ is given by the coupling of a scalar field and of a scalar function to the metric tensor, respectively. As a second case we consider that $K_{\mu \nu}$ is given by a term proportional to the matter energy-momentum tensor. The first choice gives a particular version of the $f(R,T)$ gravity model \cite{fT1}, while the second choice corresponds to specific case of the modified gravity theory of the form $f\left(R,T,R_{\mu \nu}T^{\mu \nu},T_{\mu \nu}T^{\mu \nu}\right)$ \cite{Har4}. For each considered model we obtain the gravitational field equations, and the generalized Friedmann equations for the case of a flat homogeneous and isotropic geometry. In some of these models the divergence of the matter energy-momentum tensor is non-zero, indicating a process of matter creation. From a physical point of view a non-zero divergence of the energy-momentum tensor can be interpreted as corresponding to an irreversible energy flow from the gravitational field to the matter fluid. Such an irreversible thermodynamic process is the direct consequence of the nonminimal curvature-matter coupling, induced in the present case by the quantum fluctuations of the metric \cite{fT2, creat, Pavon}. The cosmological evolution equations of these modified gravity models induced by the quantum fluctuations of the metric are investigated in detail by using both analytical and numerical methods. As a result of this analysis we show that a large variety of cosmological models can be constructed. Depending on the numerical values of the model parameters, these cosmological models can exhibit both late time accelerating, or decelerating behaviors. The present paper is organized as follows. In Section~\ref{sect2} we derive the general set of field equations induced by the quantum fluctuations of the metric. The relation between this approach and the standard semi-classical formulation of quantum gravity is also briefly discussed. In Section~\ref{sect3} we investigate in detail the cosmological implications of the quantum fluctuations induced modified gravity models with the fluctuation tensor proportional to the metric. Two cases are considered, in which the fluctuation couples to the metric via a scalar field, and a scalar function, respectively. Modified gravity models induced by quantum metric fluctuations proportional to the energy-momentum tensor are investigate in Section~\ref{sect4}. Finally, we discuss and conclude our results in Section~\ref{sect5}. The details of the derivation of the gravitational field equations for an arbitrary metric fluctuation tensor and for a fluctuation tensor proportional to the matter energy-momentum tensor are presented in Appendices~\ref{App1} and \ref{App2}, respectively. In the present paper we use a system of units with $c=1$. | \label{sect5} In the present paper we have considered the cosmological properties of some classes of modified gravity models that are obtained from a first order correction of the quantum metric, as proposed in \cite{re8,re11}. By assuming that the quantum metric can be decomposed into two components, and by substituting the fluctuating part by its (classical) average value $K_{\mu \nu}$, a large class of modified gravity models can be obtained. As a first step in our study we have derived the general Einstein equations corresponding to an arbitrary $K_{\mu \nu}$. An important property of this class of models is the non-conservation of the matter energy-momentum tensor, which can be related to the physical processes of particle creation due to the quantum effects in the curved space-time. By assuming that $K_{\mu \nu}\propto g_{\mu \nu}$, a particular class of the modified $f(R,T)$ gravity theory is obtained. We have investigated in detail the cosmological properties of these models, by assuming that the coupling coefficient between the metric and the average value of the quantum fluctuation tensor is a scalar field with a non-vanishing self-interaction potential, and a simple scalar function. The scalar field self-interaction potential was assumed to be of Higgs type \cite{Aad}, which plays a fundamental role in elementary particle physics as describing the generation of mass in the quantum world. We have investigated two cosmological models, in which the scalar field is in the minimum of the Higgs potential, and the case of the "complete" Higgs potential. In both approaches in the large time limit the Universe enters an accelerating phase, with the accelerating de Sitter solution acting as an attractor for these cosmological models. However, in the case of the "complete" Higgs potential the redshift evolution of the deceleration parameter $q$ indicates at low redshifts an extremely complex, oscillating behavior. Such a dynamics, as well as the corresponding cosmological evolution may play a significant role in the inflationary/post inflationary reheating phase in the history of the Universe. By assuming that the coupling between the metric and the quantum fluctuations is given by a scalar function, two distinct cases of cosmological models can be obtained. By imposing the conservativity of the energy-momentum tensor for a Universe filled with matter obeying a linear barotropic equation of state, a decelarating cosmological model is obtained. Thus a model could be useful to describe the evolution of the high density Universe at large redshifts. An alternative model with matter creation can be also constructed, by imposing a specific equation of evolution for the time evolution of the coupling function $\alpha $. This model leads to an approximately de Sitter type expansion, with the deceleration parameter tending to minus one in the large time limit, A second modified gravity model can be obtained by assuming that the average value of the quantum fluctuation tensor is proportional to the matter energy-momentum tensor, $K_{\mu \nu}\propto T_{\mu \nu}$. This choice leads to an extension of the $f\left(R,T,R_{\mu \nu}T^{\mu \nu}\right)$ gravity theory \cite{Har4,Odin}, with the quantum corrected action including an extra term $T_{\mu \nu}T^{\mu \nu}$. Hence by considering the effects of the quantum fluctuations of the metric proportional to the matter energy-momentum tensor a particular case of a general $f\left(R,T,R_{\mu \nu}T^{\mu \nu},T_{\mu \nu}T^{\mu \nu}\right)$ modified gravity theory is obtained. The numerical analysis of the cosmological evolution equations for this model show that, depending on the numerical values of the coupling constant $\alpha $, for a Universe filled with dust matter a large variety of cosmological behaviors can be obtained at low redshifts, with the deceleration parameter varying between a constant (approximately) zero value on a large redshift range, and a de Sitter phase reached at $z=0$. Quantum gravity represents the greatest challenge present day theoretical physics faces. Since no exact solutions for this problem are known, resorting to some approximate methods for studying quantum effects in gravity seems to be the best way to follow. A promising path may be represented by the inclusion of some tensor fluctuating terms in the metric, whose quantum mechanical origin can be well understood. Interestingly enough, such an approach leads to classical gravity models involving geometry-matter coupling, as well as non-conservative matter energy-momentum tensors, and, consequently, to particle creation processes. Hence even the study of the gravitational models with first order quantum corrections can lead to a better understanding of the physical foundations of the modified gravity models with geometry-matter coupling. In the present paper we have investigated some of the cosmological implications of these models, and we have developed some basic tools that could be used to further investigate the quantum effects in gravity, | 16 | 7 | 1607.04874 |
1607 | 1607.06089_arXiv.txt | We report the detection of morphology--dependent stellar age in massive quenched galaxies (QGs) at z$\sim$1.2. The sense of the dependence is that compact QGs are 0.5--2 Gyr older than normal-sized ones. The evidence comes from three different age indicators, D$_n4000$, H$_{\delta}$ and fits to spectral synthesis models, applied to their stacked optical spectra. All age indicators consistently show that the stellar populations of compact QGs are older than their normally--sized counterparts. We detect weak [\ion{O}{2}] emission in a fraction of QGs, and the strength of the line, when present, is similar between the two samples; however, compact galaxies exhibit significantly lower frequency of [\ion{O}{2}] emission than normal ones. A fraction of both samples are individually detected in 7--Ms Chandra X--ray images (luminosities$\sim10^{40}$--$10^{41}$ erg/sec). 7--Ms stacks of non-detected galaxies show similarly low luminosities in the soft band only, consistent with a hot gas origin for the X--ray emission. While both [\ion{O}{2}] emitters and non-emitters are also X--ray sources among normal galaxies, no compact galaxy with [\ion{O}{2}] emission is an X--ray source, arguing against an AGN powering the line in compact galaxies. We interpret the [\ion{O}{2}] properties as further evidence that compact galaxies are older and further along into the process of quenching star--formation and suppressing gas accretion. Finally, we argue that the older age of compact QGs is evidence of progenitor bias: compact QGs simply reflect the smaller sizes of galaxies at their earlier quenching epoch, with stellar density most likely having nothing directly to do with cessation of star--formation. | \label{Introduction} The formation and evolution of massive early-type galaxies remains poorly understood, despite much recent progress. Constraints from the local Universe indicate that their stellar ages are very old ($>$10 Gyr), indicating that they formed the bulk of their stellar masses at z$>$2 \citep{Bower1992, Renzini1993, vanDokkumEllis2003, Heavens2004, Renzini2006}, subsequently quenching star-formation, and remaining quenched, until the present. Constraints on stellar abundance ratios (high $\alpha$/Fe) indicate their star-formation took place on short timescales \citep{Thomas2005, Thomas2010, Renzini2006}. Additionally, it has been observed that galaxy morphology and star-formation properties are highly correlated, such that this quenched nature in massive galaxies appears coincident with morphological transformation to ellipsoidal stellar structure \citep[e.g.][]{Strateva2001,Kauffmann2003b,Franx2008}. Despite efforts to study this transition from star-forming galaxy to quenched ellipsoid, we have gained very little insight into both the transformational quenching process, as well as the mechanisms preventing further star-formation for the majority of the Universe's history. Of particular importance to this effort are constraints from observing the progenitors of these massive early-type galaxies at z$>$1, shortly after their transformation from star-forming galaxy. Recently quenched galaxies (QGs) begin to appear in large numbers at z$\sim$2 \citep{Cimatti2008, vanDokkum2008, Cassata2011,Cassata2013}, and have been observed even out to z$\sim$3-4 \citep[][]{Fontana2009, Guo2012, Gobat2012, Muzzin2013, Stefanon2013, Straatman2014}. The properties of these quenched high-redshift galaxies provide significant insight into both the formation process of the galaxies during the star-formation phase, as well as the quenching mechanisms causing their transformation. The most striking feature of these recently quenched galaxies at high-redshift is their stellar structure; while already having built up a similar amount of stellar mass as their z$\sim$0 counterparts, they are remarkably compact in stellar density \citep{Daddi2005, Trujillo2006, Bundy2006, Zirm2007, Toft2007, vanDokkum2008, Cimatti2008, vanderwel2008, Bezanson2009, Saracco2009, Damjanov2009, Williams2010}. The overwhelming majority of QGs ($>$80\%) at z$>$1.5 have stellar densities higher than the lower 1$\sigma$ of passive (early-type) galaxies at z$\sim$0 at the same stellar mass \citep{Cassata2013}. Additionally, they are much smaller than the majority of massive star-forming galaxies at the same epoch \citep{vanderWel2014}, and in fact, one of the strongest predictors of quiescence among high-redshift galaxies is centrally concentrated light \citep[i.e. a measure of compactness;][]{Franx2008, Bell2012, Omand2014, Teimoorinia2016, Whitaker2016}. It appears, therefore, that compactness and quenched nature at high-redshift are inextricably linked. However, the physical reason for this correlation is also poorly understood. Does the existence of the compact QGs at high-redshift imply something very fundamental about quenching? In other words, does some physical process associated with stellar compactness predispose galaxies to quench? Alternatively, are the earliest galaxies to form in the Universe and complete their evolution simply the densest because the Universe was denser at early times \citep[e.g.][]{LillyCarollo2016}), or, because of some highly dissipative gaseous process that could take place predominantly at high redshift \citep[e.g.][]{Dekel2009, Johansson2012,DekelBurkert2014,Zolotov2015, Ceverino2015}. There are some physically motivated reasons to believe that the former may be true. First, high stellar density implies a previous epoch of high surface density of star-formation, which would mean a higher energy input into the interstellar medium (ISM) of compact galaxies, than might be present in larger, extended galaxies. \citet{Hopkins2010} have made this argument, based on the observation that there appears to be a maximum stellar density of {\it any} structure in the Universe ($\Sigma \sim$ 10$^{11}$ M$_{\odot}$ kpc$^{-2}$). This limit in stellar density exists despite covering 8 orders of magnitude in stellar mass, from star clusters within galaxies, to the z$>$2 compact QGs. This empirical limit argues for some stellar feedback process, such as massive stellar winds, that truncate the star-formation and prevent further growth beyond this density limit. Studies of objects with high surface density of star-formation, where such extreme stellar feedback might be expected, have in fact found evidence of feedback in the form of very fast ($\sim$1000 km/s) galactic-scale outflows from extremely compact star-forming regions that approach the Eddington limit \citep{DiamondStanic2012, Sell2014}. The second, alternative scenario exists, that this connection is a very simple consequence of the size-evolution of star-forming galaxies, whose radii (at fixed mass) are observed to decrease with increasing redshift \citep[e.g.][]{vanderWel2014}. The most massive galaxies in the early Universe have formed the earliest in cosmic time, and therefore evolve to the end of their star-forming lifetime earliest. In this scenario, the density of the parent halo of the quenched population at any redshift reflects the density of the Universe at its formation epoch \cite[e.g.][]{Mo1998}, and therefore will progressively increase in size (and therefore stellar density) over time. Such a scenario, known as progenitor bias \citep[as described by][]{LillyCarollo2016}, would also contribute to the increasing size evolution of QGs over cosmic time \citep{Valentinuzzi2010a, Valentinuzzi2010b, Poggianti2013, Carollo2013, LillyCarollo2016} \citep[see also][]{Bezanson2009}. In this scenario, the significance of compactness is irrelevant for the quenching, and rather, galaxies quench when they've reached sufficient mass that they no longer support star-formation. The quenching mechanism then may be related to halo mass, or some other mass related mechanism to cut of gas supply for future star-formation \citep[e.g.][]{Peng2010, BirnboimDekel2003, DekelBirnboim2006}. Distinguishing between these scenarios is highly important to understanding the evolution of early-type galaxies. Each of these two scenarios have empirical predictions for the properties of QGs. In the stellar-density regulated star-formation scenario, galaxies with high enough surface density of star-formation will quench, and produce remnants with high stellar density. Thus at any given epoch, the most recently quenched objects should also be the densest \citep[e.g.][]{Whitaker2012}. There is no explicit prediction for a trend of stellar density with stellar age or mass \citep[e.g.][]{Hopkins2010}. However, progenitor bias explicitly predicts that stellar age, and stellar density be correlated, with the densest objects also being the oldest at any given epoch \citep[in the absence of size growth via merging;][]{LillyCarollo2016}. In this paper, we seek to distinguish between these two scenarios, and in the process, gain insight into why galaxies quench their star-formation early in cosmic time. In Section 2 we present the data used in this study. in Section 3 we present our results, and in Section 4 we discuss these results in the context of quenching and the formation of the QGs. Throughout this paper we assume a cosmology with $\Omega_{\Lambda}$=0.7, $\Omega_{M}$=0.3, and H$_{o}$ = 70 km/s/Mpc. \begin{figure*} [!t] \begin{center} \includegraphics[scale=0.48]{fig1.eps} \caption{The mass-size relation of z$\sim$1.2 QGs in this study. In red are QG defined as compact according to \citet{Cassata2013}, as being below the lower 1-$\sigma$ of the z$\sim$0 early-type galaxy mass-size relation \citep[orange dashed line;][]{Shen2003}. The z$\sim$0 mean early-type galaxy mass-size relation is the blue dot-dashed line \citep{Shen2003}. Blue galaxies are considered normally-sized QG (relative to early-type galaxies at z$\sim$0). Galaxies with triangles designate those in which [\ion{O}{2}]$\lambda$3727 emission was detected. Squares designate those galaxies with X-ray detections. Bottom panel: the mass distributions of each sample are roughly equivalent.} \label{masssize} \end{center} \end{figure*} | There are three main results in this study: 1) in the redshift range that we have considered, $1\le z\le 1.4$, massive compact QGs have stellar populations that are, on average, older than normally sized ones; 2) the frequency of [\ion{O}{2}] detection and X-ray detection are significantly lower among the compact galaxies; and 3) the X-ray properties generally disfavor the presence of strong AGN in both samples of recently QGs (low luminosities $\approx 10^{40}-10^{41}$ erg/sec from both the few X-ray detections, and average stacked emission). While X--ray detected normal QGs often also have [\ion{O}{2}] emission, not a single compact galaxy with [\ion{O}{2}] emission is individually detected in Chandra images. This strongly argues against AGN as the power source of the [\ion{O}{2}] emission in compact galaxies, and favors instead either warm gas, stellar remnants and/or residual star formation, minor merging with a gas rich companion, or LINER emission, possibly powered by stellar sources \citep[e.g.][]{YanBlanton2012, Singh2013}. These mechanisms may also be active in the normal galaxies. Although it is unclear what are the sources of [\ion{O}{2}] emission, it is clear that they are less active among the compact sample. Taken all together, these lines of evidence paint a picture in which compact galaxies formed and evolved earlier than normal ones and, consequently, quenched star-formation earlier. \begin{table*}[!!th] \begin{center} \caption{ Stacked 7Ms X-ray Fluxes \label{table1}} \begin{tabular}{llllll} & & & \\ \tableline Flux [ergs/s/cm$^{2}$]\footnote{Energy conversion factors evaluated using the Chandra PIMMS tool (http://cxc.harvard.edu/toolkit/pimms.jsp)} & Galaxy Sample & Soft band & Hard band & \\ \tableline & & & & \\ & Compact & 3.46$\pm$1.37x10$^{-18}$&-0.30$\pm$1.73x10$^{-17}$ \\ & Normal & 5.99$\pm$1.53x10$^{-18}$&2.48$\pm$1.68x10$^{-17}$ \\ \tableline \tableline Luminosity [ergs/s]\footnote{assuming the average redshift of each sample}& Galaxy Sample & Soft band & Hard band & \\ \tableline & & & \\ & Compact & 2.96$\pm$1.17x10$^{40}$&-0.26$\pm$1.48x10$^{41}$ \\ & Normal & 4.23$\pm$1.08x10$^{40}$&1.75$\pm$1.19x10$^{41}$ \\ \tableline \tableline \end{tabular} \label{xstack} \end{center} \end{table*} \subsection{Age constraints: evidence for progenitor bias} The evidence of the age difference between normal and compact QGs we have presented here comes from two independent age diagnostics, the $D_{n(4000)}$ and the H$\delta_{A}$, as well as stellar population synthesis modeling. The $D_{n(4000)}$ is larger in compact galaxies than normal ones, corresponding to an age difference of $\sim$0.3 Gyr, using the calibrations by \citet{Kauffmann2003}. The difference in H$\delta_{A}$ implies a larger age differential ($\sim$2.5 Gyr) but as discussed in Section 3, may be somewhat overestimated due to continuum absorption in old, metal-rich stars. Although these age conversions are model dependent, the evidence for an age {\it difference} between samples is independent of the adopted stellar libraries and assumed metallicity for the range of values found here. Stellar population synthesis modeling with pPXF provide consistent results with the Lick Indices. All age diagnostics considered here imply age differentials in the same direction, i.e. more compact passive galaxies are older. The age differential between normal and compact QGs that we discuss here is at redshift $z\sim 1.2$. However, the result is in qualitative and quantitative agreement with the observations by \citet{Belli2015} at $z\sim 2$ that at a given redshift, the largest galaxies (in radius, although related to stellar density) are among the youngest, suggesting that the property is a general feature of passive, massive galaxies at high redshift. Additionally, \citet{Saracco2009} have arrived at the same result via the opposite analysis from us; by separating the oldest QGs during the epoch 1$<$z$<$2 from the youngest ($\delta$age$\sim$1.5-2Gyr) they find the youngest to reside on the local z$\sim$0 early-type galaxy mass-size relation (like our normal QG sample), whereas the old QGs are denser, with radius a factor of 2.5-3 smaller than local early-type galaxies. Although, they note that their QG samples differ in mass with the younger sample being less massive \citep[see also][]{Thomas2010, Fagioli2016}. The fact that we observe the same trend with essentially an identically mass-matched sample indicates that mass is not the primary factor related to the age differential in the population, and rather the stellar density or size may be the primary factor \citep[][find a similar result]{Saracco2011}. In a complementary study, \citet{Fagioli2016} find that at lower redshifts than our sample (0.2$<$z$<$0.8) this trend of age and compactness in QGs persists in the stellar mass range explored here (however, \citet{Trujillo2011} do not find evidence for such a trend to z$\sim$0). At high redshift, compact galaxies dominate the population of QGs at the high mass end \citep[][]{Cimatti2008, vanDokkum2010, Cassata2011, Cassata2013, vanderWel2014}. This has prompted investigations of scenarios where quenching is more efficient in galaxies with high stellar density because of increased stellar feedback \citep[e.g.][]{Hopkins2010}. However, a causal relationship between high stellar density (compactness) and likelihood of quenching does not directly predict an age--density correlation; rather, at any given epoch, the densest galaxies should be the most likely to quench, but not necessarily the oldest \citep[see e.g.][]{Whitaker2012, Yano2016}. The presence of a relationship between age and stellar density where by denser galaxies are older is precisely the prediction of the progenitor bias scenario \citep{LopezSanjuan2012, Carollo2013, Poggianti2013, Belli2015, Keating2015, Wellons2015, Wellons2016, LillyCarollo2016}. That is, due to the observed size-evolution of star-forming galaxies \citep[e.g.][]{vanderWel2014}, the density of a galaxy reflects the density of the Universe when the galaxy formed (assuming very little, or average, structural disruption) and therefore older galaxies should be denser. Thus, at face value, our results support that galaxies which form earlier, and completed their evolution earlier, were simply denser than larger galaxies that form and evolve later, without the density necessarily having anything directly to do with the cessation of their star formation activity. Star-formation may be then affected by some other quenching agent \citep[e.g. halo or mass quenching;][]{BirnboimDekel2003,DekelBirnboim2006, Peng2010}. This idea is extensively discussed in \citep{LillyCarollo2016} who show by means of a simple toy model that this scenario will naturally explain the correlations between galaxy structure and star-formation properties, without the need of a stellar density-related quenching mechanism \citep[see also][]{Abramson2016}. \subsection{Energy sources: quenching agents in QGs} An independent, but complementary, piece of information comes from the X--ray and [\ion{O}{2}] properties of our two samples. In themselves the [\ion{O}{2}] emission line and the X--ray data do not provide any firm indication as to the causes of quenching. The average X--ray luminosity for the majority of our sample does not show evidence of any powerful AGN, but one could have been present prior to reaching the current, very low level of star formation activity. There is evidence that the sources of ionizing radiation or warm gas that are still present in the two samples are different at the time of observation; namely the compact galaxies show a significantly lower detection rate of [\ion{O}{2}] and of X-ray emission than the normal QGs. This difference is fully consistent with a scenario where the quenching occurred earlier in the compact sample and have therefore had a longer time to fade. Larger QGs from the normal sample are more likely to exhibit emission from energizing sources simply because on average, quenching in larger galaxies was initiated more recently. Alternatively, the presence of [\ion{O}{2}] emission may be evidence of rejuvenated SF due to gas from minor merging \citep[e.g.][]{Treu2002}. In general, our data do not provide any conclusive constraints on the quenching mechanisms that truncated the star-formation in these galaxies, nor if the quenching mechanisms differ with stellar density. It remains an important goal of galaxy evolution to understand the quenching mechanisms in massive galaxies. Over the last several years, efforts have been made to identify star-forming progenitors of soon-to-be QGs at z$>$2. These efforts have relied on the fortuitous observation that the first galaxies to quench are compact, making it relatively easy to identify their immediate star-forming progenitors among compact star-forming galaxies \citep[][]{Williams2014, Barro2013, Patel2013, Stefanon2013, Nelson2014, vanDokkum2015}. It is interesting to note that various studies have come to very different conclusions about the nature of feedback in compact star-forming galaxies. \citet{Barro2013} have claimed that the AGN fraction in compact SFGs may be up to 50\%, suggesting that the presence of AGN may truncate the star-formation on short timescales. However, \citet{Spilker2016} have discovered that a subset of this sample, despite being on the star-forming main sequence, have much lower molecular gas masses than normal main-sequence galaxies, suggesting that star-formation will be quenched on short timescales due to simple gas exhaustion if the influx of gas has been suppressed. Additionally, \citet{Williams2015} found no evidence for AGN among compact star-forming galaxies, but instead detected both faster outflow velocities in the ISM, and extreme, redshifted Lyman-$\alpha$ emission among compact star-forming progenitors. \citep[Similar Ly$\alpha$ signatures were identified among quenching galaxies by][]{Taniguchi2015}. They interpreted these observational signatures as related to the compact galaxies having enhanced feedback in the ISM, due to higher surface density of star-formation than their more extended counterparts \citep[see also][]{Alexandroff2015,DiamondStanic2012,Sell2014,HeckmanBorthakur2016}, plausibly leading to the truncation of future star-formation. From our data it is clear, however, that in the time since quenching was initiated in these QG samples, major energizing sources have already dissipated. Future detailed studies of galaxies closer to their quenching epoch may identify the physical processes that shut down star-formation at high-redshift. \begin{figure*} [!th] \begin{center} \includegraphics[scale=0.35]{fig8a.eps} \includegraphics[scale=0.35]{fig8b.eps} \caption{Left panel: 7 Ms hard-band detected QGs, including photometric QG without optical spectroscopy (gray points) compared to the local QG relation for total X--ray luminosity (0.3-10 keV) from LMXB \citep[solid black line;][]{Gilfanov2004}. We have additionally k-corrected the local QG relationship to z$\sim$1.2 (dashed line) assuming a spectral slope as described in the text. Right panel: The same relation for the 7 Ms soft-band detected QGs. The X--ray luminosity from the z$\sim$1.2 QG sample is larger than that emitted purely by LMXB in local QGs at a given stellar mass, plausibly suggesting additional sources of X-ray emission. } \label{xmass} \end{center} \end{figure*} \subsection{X-ray emission from passive galaxies: AGN, binaries, or hot gas?} An important feature of both samples is that some of the QGs have X-ray emission detected in the deep Chandra 7 Ms data of GOODS-South. AGN have long been suspected to be the key agent behind the quenching of star formation \citep[e.g.][]{Granato2004}, as well as prevention of future star formation in galaxies \citep[e.g.][]{CiottiOstriker1997, Fabian2012}. Thus an obvious question is whether or not the X-ray properties of our QGs are consistent with this idea. Only a minority of the galaxies are individually detected (see Table 1); stacking the images at the position of the the non-detected galaxies, however, yields measurable flux in the soft band, but not in the hard one. While the range of hardness ratio spanned by our galaxies is the same as narrow--line AGN and normal galaxies \citep[e.g.][]{Szokoly2004, Hasinger2008}, the distribution of hardness ratio and X-ray luminosity shown by our galaxies qualitatively looks very different from the distribution of the parent population of X-ray detections \citep[see gray points in Figure \ref{xray};][]{Luo2016}. In number, these X-ray sources are dominated by AGN with luminosities $>$ 10$^{42}$ erg/s \citep[see also][for the distribution among more powerful obscured AGN]{Wilkes2013}. The distribution exhibits large scatter in hardness ratio and X-ray luminosity, most likely a reflection of the diversity of obscuration, orientation, spectral slope, power and selection effects of the AGN population. Our X--ray detected sample have significantly lower luminosities than the majority of AGNs at comparable redshifts, suggesting a different origin for the X--ray emission (with the exception of the brightest QG, whose high X--ray luminosity and extreme hardness ratio suggest it may host an obscured AGN). Figure 7 also shows a plot of the X-ray luminosity versus the \ion{O}{2}\ luminosity (right panel). We do not find a correlation between X-ray and \ion{O}{2}\ luminosities, which, if present, may have suggested an AGN origin for both (however we note the sample has limited dynamic range in \ion{O}{2}\ luminosity). These arguments, taken together with the fact that we observe stacked X-ray emission in the soft band but not the hard one (i.e. on average our galaxies are soft X-ray sources) suggests that the X-ray emission is unlikely powered by AGN. Alternatively, the X--ray luminosity of our QGs could be primarily powered by emission from hot gas (e.g. coronal gas that formed as a combination of outflows, stellar winds, or gravitational heating), and/or low--mass X--ray binaries (LMXB). To investigate the contribution from hot gas and LMXB, we plot the relationship of both the hard and soft X--ray luminosity with stellar mass for our X--ray detected samples in Figure \ref{xmass}. At z$\sim$0, it is well established that X--ray luminosity from LMXB scales linearly with stellar mass \citep[e.g.][]{Gilfanov2004, KimFabbiano2004,Colbert2004}. We find that our spectroscopic sample (colored points) does not show any correlation between X--ray luminosity and stellar mass, as would be expected if the X--ray luminosity was primarily from LMXBs. Since our X--ray detected sample is small, we additionally plot photometrically selected QGs from the parent sample selected in \citet{Cassata2013} at 1$<$z$<$1.5 that are X--ray detected in the {\it Chandra} 7 Ms imaging (i.e. those that were excluded from this study due to lack of optical spectroscopy). We limit the photometric QG sample to this narrow redshift range to mitigate the uncertain effects due to comparing the luminosities without the k-correction, which becomes increasingly large as the redshift range within the sample increases. The X--ray luminosities we observe in our QG sample (both hard and soft) are well in excess of that seen from LMXB in local QGs \citep[solid line in Figure \ref{xmass};][]{Gilfanov2004}. The local scaling relation appears an order of magnitude lower than our data, despite the luminosity being the total sum of photons with energies spanning both the hard and soft bands \citep[0.3-10 keV;][]{Gilfanov2004}. We k-corrected the local QG relation to the observed-frame of the QG sample (with mean redshift z$\sim$1.2) assuming a conservative (steep) but typical power law spectral shape defined by a photon index of $\Gamma \sim$1.8, typical for low-redshift LMXB \citep{Lehmer2007, Boroson2011}. The k-correction increases the expected luminosity at z$\sim$1.2, however, is still well below our observations. Observationally it is unknown if the X--ray luminosity from LMXB evolves with redshift; although theoretical scaling relations indicate the possibility that LMXB luminosity may increase with redshift, at a given stellar mass, based on metallicity and star-formation history evolution in the Universe \citep{Fragos2013}. We are unable to place further constraints on the contribution of LMXB. We conclude by simply noting the following: 1) our QGs are an order of magnitude more luminous in the X--ray than the LMXB contribution in local QGs of the same mass, however, 2) their observed X--ray luminosities are comparable to the total emission of local massive QGs, which includes the luminosity emitted from hot gas \citep[e.g.][]{KimFabbiano2013, Goulding2016}. Locally, it is known that the contribution from $\sim$1 keV gas dominates the soft-band emission \citep{Matsumoto1997, Sivakoff2004, Gilfanov2004}, also consistent with the luminosity seen on average in the stacked z$\sim$1.2 QGs. This leaves open the likely possibility that some X--ray emission comes from hot gas in these QGs, heated either by gravitational or feedback processes. However, the dominant X--ray source in these z$\sim$1.2 QGs remains unknown. An in--depth study of the X--ray emission of our sample and the comparison with local counterparts is beyond the scope of this paper, which is devoted to the galaxies' stellar age. We simply note that the possibility that we are observing hot gas emission from individual quenched galaxies at $z\sim 1.2$, an observation which is very important to understanding the evolution of the interstellar medium and circum-galactic medium in early-type galaxies. In a forthcoming paper, we plan to accurately quantify selection effects and further investigate the X-ray properties of QGs. If confirmed, it will provide a powerful diagnostic of the feedback that took place during the star--formation phase \citep[e.g.][]{vandeVoort2016} and that is preventing the galaxies from forming stars again. We also note that even if the observed X--ray emission does indeed turn out to be from hot gas, this still does not rule out AGN as the cause of quenching, since the AGN activity could have been quenched together with the star-formation, for example, as a result of the cessation of gas accretion into the galaxies. The question of whether or not AGN is responsible for star-formation quenching in galaxies thus remains an open one. However, if hot gas is indeed a major component of the X--ray emission, this suggests that the quenching may be driven more by heating the gas than expelling it. | 16 | 7 | 1607.06089 |
1607 | 1607.03453_arXiv.txt | {Following an approach of Matarrese and Pietroni, we derive the functional renormalization group (RG) flow of the effective action of cosmological large-scale structures. Perturbative solutions of this RG flow equation are shown to be consistent with standard cosmological perturbation theory. Non-perturbative approximate solutions can be obtained by truncating the a priori infinite set of possible effective actions to a finite subspace. Using for the truncated effective action a form dictated by dissipative fluid dynamics, we derive RG flow equations for the scale dependence of the effective viscosity and sound velocity of non-interacting dark matter, and we solve them numerically. Physically, the effective viscosity and sound velocity account for the interactions of long-wavelength fluctuations with the spectrum of smaller-scale perturbations. We find that the RG flow exhibits an attractor behaviour in the IR that significantly reduces the dependence of the effective viscosity and sound velocity on the input values at the UV scale. This allows for a self-contained computation of matter and velocity power spectra for which the sensitivity to UV modes is under control. } \begin{document} | \label{intro} Understanding how density fluctuations in cold dark matter evolve under the influence of gravity is at the basis of analyzing data on large-scale structures and of constraining cosmological models from them. Technically, this task amounts to solving the collisionless Vlasov-Boltzmann equation for appropriate classes of initial conditions \cite{bernardeau0}. With recent progress in both computational power and coding techniques, $N$-body simulations describe by now structure evolution down to small scales where baryonic effects become increasingly important \cite{Kuhlen:2012ft, Vogelsberger:2014dza, Schneider:2015yka}. Even for pure cold dark matter simulations, perturbative techniques cannot be expected to apply on scales where non-linearities become strong and the process of virialization starts to dominate. Nevertheless, $N$-body simulations remain CPU-intensive, while suitable analytical techniques can help to scan the relevant range of initial conditions, as well as to explore non-standard cosmological models. Also, there is the general desire to understand the result of complex simulations in simple analytical terms for certain limiting cases. Last but not least, future large-scale structure surveys, including EUCLID~\cite{Amendola:2012ys}, LSST~\cite{Abell:2009aa}, DESI~\cite{Levi:2013gra} and eBOSS~\cite{Dawson:2015wdb}, will probe scales within the weakly non-linear regime with increasing precision. For these and other reasons, much effort has gone recently into complementing $N$-body simulations with advanced analytical techniques. Cosmological perturbation theory is based on writing the Vlasov-Boltzmann equation as a hierarchy of equations for the momentum moments of the phase-space distribution \cite{bernardeau0}. If truncated at the lowest moments, the resulting equations of motion are of ideal fluid dynamic form. A linearized perturbative treatment of this `cosmological fluid' works well for the small fluctuations at sufficiently large scales for redshift $z=0$. In addition, the leading non-linear corrections computed within this framework capture the dominant effects on the density power spectrum at scales $k \lesssim 0.075\, h/$Mpc \cite{Scoccimarro:1995if, bernardeau0, bernardeau, Carlson:2009it, Audren:2011ne, CrSc1, Taruya:2012ut}. Moreover, unequal-time correlators exhibit a damping due to the stochastic background of large-scale bulk flows. This effect can be rather accurately described within an ideal fluid dynamical framework by resumming certain classes of perturbative contributions related to long-wavelength perturbations \cite{CrSc2, Bernardeau:2011vy}. While most of this effect cancels out in equal-time correlators \cite{Jain:1995kx, Peloso:2013zw, Kehagias:2013yd, nonlinear}, its residual effect is important for the non-linear broadening and shift of the baryon acoustic peak \cite{Crocce:2007dt, Smith:2007gi, Sherwin:2012nh, Baldauf:2015xfa, Senatore:2014via}. This influence of very long-wavelength modes on the scale of baryon acoustic oscillations (BAO) can be computed systematically by a suitable reformulation and resummation of perturbation theory \cite{Baldauf:2015xfa, Zaldarriaga:2015jrj, Blas:2015qsi, Blas:2016sfa}. However, beyond the linear approximation, fluctuations around the BAO scale are also modified by interactions with the UV part of the spectrum of density fluctuations. Here, the standard cosmological perturbation theory fails, in the sense that higher-order `loop' corrections to the propagation of long-wavelength fluctuations become increasingly more sensitive to the UV part of the spectrum, see e.g.~\cite{Blas1}. Several properties are relevant for a discussion of this failure: First, the density contrast grows large in the UV. Second, at UV scales where the so-called single-stream approximation fails, higher moments of the phase-space distribution that are not accounted for in standard perturbation theory become important. Third, $N$-body simulations~\cite{Bagla:1996zb,nishimichi,Pueblas:2008uv} as well as general considerations about the decoupling of virialized substructures \cite{Peebles:1980} and the form of non-perturbative response functions \cite{Garny:2015oya}, indicate that the impact of UV modes on the BAO range is overestimated by standard perturbation theory. In other words, UV modes decouple more efficiently than expected perturbatively. To account for interactions with the UV part of density fluctuations, a consistent cosmological perturbation theory needs to be based on effective dynamics, which is applicable only above some length scale and in which non-perturbative parameters absorb the effect of small-scale perturbations that are `integrated out'~\cite{Pietroni:2011iz,Manzotti:2014loa}. Since the lowest moments of the phase-space distribution that are dropped in the single-stream approximation define the stress tensor of an imperfect fluid, it is natural to absorb the integrated out UV physics in the viscous transport coefficients that parametrize the non-ideal stress tensor. This strategy is adopted for instance in recent works on the effective field theory of large-scale structure where the viscous coefficients are not predictable within the effective theory, but are fixed by comparison of calculated correlations functions with either $N$-body simulations or observations \cite{Baumann:2010tm, Carrasco:2012cv, Porto:2013qua, Foreman:2015lca, Baldauf:2015aha, Assassi:2015jqa, Abolhasani:2015mra}. For a discussion of alternative approaches leading to the same strategy, see \cite{Fuhrer:2015cia} and references therein. The main point of the present paper is that the scale dependence of the viscous coefficients is calculable. Our main result is the formulation of a renormalization group for the coefficients that enter the effective dynamics of large-scale cosmological perturbations, and an explicit calculation of the renormalization-group trajectories of the effective viscosity and sound velocity derived from it. These coefficients account for momentum transfer with UV modes that are integrated out in the effective dynamics. Technically, we start from work of Matarrese and Pietroni~\cite{Max1} who formulated renormalization-group (RG) techniques for the mildly non-linear stages of large-scale structure evolution. For other applications of RG techniques to the problem of large scale structure, see Refs.~\cite{mcdonald,time}. Based on the central elements of their proposal, which we recall in section~\ref{sec2}, we derive in section~\ref{sec3} how the effective action for long-wavelength cosmological perturbations flows with the coarse-graining scale. In section~\ref{sec4}, we demonstrate that, if solved perturbatively, this exact functional RG flow equation for the effective action is consistent with results from standard cosmological perturbation theory. In section~\ref{sec5}, we then employ the physical intuition that the effective action at mildly non-linear scales describes the dynamics of an imperfect fluid to explore a truncated non-perturbative solution of the functional renormalization group. In this way, the dependence on the coarse-graining scale of the non-perturbative parameters entering the effective dynamics of large-scale structure can be described explicitly by a set of coupled ordinary differential equations. In section~\ref{sec6}, we solve for this RG flow numerically and we discuss physical implications. Furthermore, we compare the resulting power spectra of the density contrast and velocity scalar to those measured in $N$-body simulations. We close with a short discussion and outlook. | In summary, we have developed a renormalization-group approach that determines the effect of initial-state fluctuations on the effective action and equations of motion of the cosmological large-scale structure. The formalism is based on a functional renormalization-group equation for the one-particle irreducible effective action of a stochastic field theory. This formalism generates solutions of the classical field equations for stochastic initial conditions. Initial-state fluctuations are added gradually by lowering an infrared regulator scale. We discuss how RG flow equations for propagators and other correlation functions are obtained from the flow equation for the effective action and how they can be solved perturbatively or, within truncations, non-perturbatively. More concretely, we determine the RG trajectories of the parameters corresponding to the effective viscosity and sound velocity. These quantities receive a contribution from initial state fluctuations that is largely independent of their value at the microscopic UV scale. Correspondingly, the final value in the IR is also largely independent of the microscopic value as long as the latter is not too large. Using this RG improved effective theory, we derive perturbative two-loop expressions for different two-point correlation functions of the large-scale theory, and study them numerically. Our results for the density and velocity power spectra show reasonable agreement with results from $N$-body simulations for wavenumbers below about $0.2\, h/$Mpc. For higher momentum scales, one expects that additional effects, that have not been taken into account so far, become important. In particular, the initial values of effective viscosity and pressure for the renormalization-group flow will receive contributions coming from the higher moments of the distribution function at the microscopic scale. While the attractor behaviour of the RG flow that we have discussed in section~\ref{sec5} shows that the impact of the initial conditions is actually small, they are expected to become relevant at a certain level of precision and for larger wavevectors. The impact of the stress tensor of deep UV modes on the BAO range has been quantitatively analyzed in Ref.\,\cite{Pueblas:2008uv} based on $N$-body data. For the density power spectrum at $z=0$ its impact was found to be at the percent level for modes with wavenumber $k\simeq 0.2\, h/$Mpc, and growing rapidly for larger values of $k$. A contribution of this size is consistent with the level of accuracy that we obtain for the density power spectrum. On the one hand, this result is reassuring: the effective RG treatment improves over standard perturbation theory, and yields an error estimate that is consistent with the expected size of physical effects that have been neglected. On the other hand, it is straightforward to include these effects in future extensions of the RG framework (see below). It is also straightforward to extend the RG framework to theories for which dark matter has non-trivial material properties on the microscopic level \cite{Kunz:2016yqy, Kopp:2016mhm}. In particular, a possible viscosity or sound velocity due to fundamental (self-)interactions can be taken into account in a straightforward way in the form of additional contributions to the initial condition for the RG flow at the microscopic scale. We conclude with a few more general remarks: \begin{enumerate} \item The exact functional renormalization-group equation \eqref{eq:WetterichEqn} for the effective action $\Gamma_k$ specifies how the effective theory of structure formation changes with the coarse-graining scale $k$. In principle, for a given microscopic action one can follow the RG evolution and determine the RG trajectories of the parameters that describe the effective dynamics on macroscopic scales. The analysis of the RG flow provides a way to quantify the sensitivity of the macroscopic evolution of the LSS to properties in the UV. For instance, as illustrated in section~\ref{sec6}, the RG-flow \eqref{eq:WetterichEqn} can converge so rapidly to a (partial) IR fixed point that certain parameters of the effective theory are practically insensitive to the precise initial value at the UV scale. Other macroscopic parameters could be more sensitive to the microscopic physics. \item For the present paper, we have specialized from the onset to a simplified description with a relatively small set of fields. However, the derivation of the flow equation for $\Gamma_k$ given in section~\ref{sec3} generalizes to more complete dynamical descriptions. For instance, for calculations of LSS at scales $> 0.2 \, \text{h/Mpc}$, where virialization becomes gradually more important, it is interesting to account also for additional effects, such as those induced by fluid vorticity or by higher moments of the distribution function. To this end, one can formulate, for a more comprehensive field content, an action that encodes the initial state fluctuations and that obeys an exact RG flow of the form \eqref{eq:WetterichEqn}. \item The RG flow equation~\eqref{eq:WetterichEqn} provides also a tool for testing the self-consistency of any proposed effective dynamics. More precisely, any specific effective action \eqref{eq:operatorExp} is a truncation of the full dynamics at the coarse-graining scale $k$ to a finite-dimensional subspace. The exact RG flow \eqref{eq:WetterichEqn} will generate also terms that lie outside this finite dimensional subspace. A truncation is a good approximation as long as these terms are unimportant for the further evolution. We believe that it is interesting and possible to develop techniques that employ these properties of the exact equation \eqref{eq:WetterichEqn} in order to specify the range of applicability of a specific effective action, and to assign theoretical uncertainties to the scale-dependence of the effective parameters. \end{enumerate} \subsubsection* | 16 | 7 | 1607.03453 |
1607 | 1607.08788_arXiv.txt | { We refine the mass and environment dependent spherical collapse model of chameleon $f(R)$ gravity by calibrating a phenomenological correction inspired by the parameterized post-Friedmann framework against high-resolution $N$-body simulations. We employ our method to predict the corresponding modified halo mass function, and provide fitting formulas to calculate the enhancement of the $f(R)$ halo abundance with respect to that of General Relativity (GR) within a precision of $\lesssim 5\%$ from the results obtained in the simulations. Similar accuracy can be achieved for the full $f(R)$ mass function on the condition that the modeling of the reference GR abundance of halos is accurate at the percent level. We use our fits to forecast constraints on the additional scalar degree of freedom of the theory, finding that upper bounds competitive with current Solar System tests are within reach of cluster number count analyses from ongoing and upcoming surveys at much larger scales. Importantly, the flexibility of our method allows also for this to be applied to other scalar-tensor theories characterized by a mass and environment dependent spherical collapse. } \begin{document} | \label{sec:intro} The abundance of galaxy clusters depends on the growth rate of cosmic structures as well as on the expansion history of the universe. This makes it a powerful probe of cosmology as a function of redshift, and particularly suited to investigate the nature of dark energy and deviations from General Relativity (GR)~\cite{Albrecht:2006um,Rapetti:2009ri}. Current and upcoming galaxy cluster surveys, such as the Dark Energy Survey (DES)~\cite{Abbott:2005bi}, the extended Roentgen Survey with an Imaging Telescope Array (eROSITA)~\cite{Merloni:2012uf}, the South Pole Telescope Third-Generation survey (SPT-3G)~\cite{Benson:2014qhw}, the Large Synoptic Survey Telescope (LSST)~\cite{Abell:2009aa} and {\it Euclid}~\cite{Laureijs:2011gra}, will detect an unprecedented number of these objects covering two orders of magnitude in mass ($M \sim 10^{13.5} - 10^{15.5} \, \Msunh$) for redshifts $z \lesssim 2$, with accurate calibration of the mass-observable relations down to a few percent. In order to take full advantage of this wealth of data, numerical and theoretical predictions of the mass distribution of virialized structures (also known as halo mass function) must reach a similar level of precision. Extensive effort has gone into modeling and calibrating this fully nonlinear observable in the standard cosmological constant plus Cold Dark Matter ($\Lambda$CDM) paradigm (e.g.~\cite{Maggiore:2009rv,Maggiore:2009rw,Corasaniti:2011dr,Corasaniti:2010zt,Sheth:1999mn,Jenkins:2000bv,Tinker:2008ff,Tinker:2010my,Crocce:2009mg,Manera:2009ak,Warren:2005ey,Reed:2012ih,Lukic:2007fc,Watson:2012mt,Despali:2015yla,Bocquet:2015pva}), and some work in this direction has been carried out for alternative dark energy models and gravity theories (e.g.~\cite{Barreira:2013eea,Barreira:2013xea,Barreira:2014kra,Bhattacharya:2010wy,Cui:2012is,Brax:2012sy,Schmidt:2008tn,Li:2011uw,Lombriser:2013wta,Lombriser:2013eza,Kopp:2013lea,Achitouv:2015yha}). In this paper, we focus on the class of scalar-tensor theories known as $f(R)$ gravity (for reviews see e.g.~\cite{Sotiriou:2008rp,DeFelice:2010aj}), where the standard Einstein-Hilbert action is replaced by a general nonlinear function of the Ricci scalar $R$. The function $f(R)$ can be adjusted to mimic the $\Lambda$CDM expansion history, which in turn limits deviations from the standard model only to the growth of structure on both linear and nonlinear scales due to the fifth force mediated by the new scalar degree of freedom, known as \textit{scalaron}~\cite{STAROBINSKY198099,Oyaizu:2008sr,Li:2011vk,Llinares:2013qbh,Puchwein:2013lza,Pogosian:2007sw}. Constraints from local experiments~\cite{Will:2005va} are only consistent with functional forms that display the so-called \textit{chameleon} screening mechanism~\cite{Khoury:2003rn}. This ensures that force modifications are suppressed and GR is recovered for structures with deep potential wells, as the Solar System and the Galaxy~\cite{Hu:2007nk,Brax:2008hh}. However, the same coupling between the scalaron and the standard matter fields responsible for the chameleon mechanism may lead to catastrophic particle production in the early universe prior to Big Bang Nucleosynthesis (BBN), which can only be alleviated through fine tuning of the scalaron initial conditions~\cite{Erickcek:2013oma,Erickcek:2013dea}. In addition, the scalaron amplitude has been strongly constrained on small scales and late times using unscreened local dwarf galaxies, with allowed values in the range $|f_{R0}| \lesssim 10^{-7}$ at 95.4\% confidence level~\cite{Jain:2012tn,Vikram:2013uba}. It is also worth noting that this relatively recent technique would still greatly benefit from further investigation on various relevant astrophysical systematic uncertainties. All in all, these results further support the observation that $f(R)$ theories are unlikely candidates for a fundamental theory of gravity. Nevertheless, they can still be regarded as effective theories at low redshifts and on cosmological scales, with measurable deviations from GR predictions of the large scale structure. The first studies designed to test $f(R)$ gravity with cluster number counts constrained the allowed region of parameter space to $|f_{R0}| \lesssim 10^{-4}$ at 95.4\% confidence level~\cite{Schmidt:2009am,Lombriser:2010mp}. More recently, from the abundance of X-ray selected massive galaxy clusters and utilizing the conservative halo mass function (HMF) predictions of Schmidt et al.~(2009)~\cite{Schmidt:2008tn}, Cataneo et al.~(2015)~\cite{Cataneo:2014kaa} improved this upper bound by an order of magnitude. Upon accurate modeling of the nonlinear chameleon mechanism, and employing the same cluster abundance data, weak lensing mass calibration and cluster analysis~\cite{Mantz:2014paa} this constraint could be further reduced by about a factor of two. An even more interesting prospect comes from including lower mass objects ($M \sim 10^{13.5} \, \Msunh$) at low redshift ($z\sim0.1$) along with an improved mass calibration down to 5\%, which could further strengthen the upper limit to $|f_{R0}| \lesssim 10^{-6}$. Thus, cluster count constraints have the potential to be competitive with those set by astrophysical and local tests of gravity but on much larger scales~\cite{Lombriser:2011zw,Joyce:2014kja}. To this end, we present a phenomenological modification of the spherical collapse model of Lombriser et al.~(2013)~\cite{Lombriser:2013wta}, which we calibrate against high-resolution $N$-body simulations to predict the relative abundance of halos in $f(R)$ gravity with respect to GR within a $5\%$ precision (see~\cite{Li:2011uw,Kopp:2013lea,Achitouv:2015yha} for alternative approaches; for recent applications of the theoretical mass function presented in~\cite{Kopp:2013lea,Achitouv:2015yha} see~\cite{Peirone:2016wca,Liu:2016xes}). This is the first in a series of two papers dedicated to accurately modeling, robustly analyzing and tightly constraining chameleon $f(R)$ gravity from the abundance of massive clusters. While here we develop an accurate model of the $f(R)$ mass function, observational constraints will be presented in the second paper of the series. In Sec.~\ref{sec:theory} we review the main aspects of $f(R)$ gravity including the chameleon screening. Sec.~\ref{sec:scfR} summarizes the spherical collapse approach of Lombriser et al.~(2013)~\cite{Lombriser:2013wta} and introduces our new parametrization to correct for residual inaccuracies in that model. The dark matter only cosmological simulations that we use to calibrate the new model are described in Sec.~\ref{sec:simulations}, and we present our halo mass function predictions in Sec.~\ref{sec:hmf}. We conclude in Sec.~\ref{sec:conclusions} with an outlook on possible extensions and applications of our results. | \label{sec:conclusions} The abundance of galaxy clusters is sensitive to the growth of the large scale structure, and as such can effectively test departures from GR on cosmological scales. Upcoming and future cluster surveys will provide exquisite data, requiring accurate percent level theoretical predictions to realize the full potential of these measurements. In this work we have presented a novel semi-analytical approach that combines the advantages of the spherical collapse model of Lombriser et al.~(2013)~\cite{Lombriser:2013wta} with the information available in fully nonlinear cosmological simulations. Taking GR as a baseline theory of gravity, we have calibrated mass function ratios in the context of $f(R)$ gravity and obtained fitting functions for our additional parameters able to predict these ratios within a 5\% precision for the ranges $10^{13.5} \leqslant M_{300 \rm m} (\Msunh)^{-1} \leqslant 10^{15.5} $, $10^{-6} \leqslant |f_{R0}| \leqslant 10^{-4}$ and $0 \leqslant z \leqslant 0.5$. This corresponds to about a $50\%$ improvement on the purely spherical collapse results of~\cite{Lombriser:2013wta}. A similar level of accuracy can be achieved for the full $f(R)$ mass function on the condition that the modeling of the reference GR halo abundance is accurate at the percent level. Although in Eqs.~\eqref{eq:HMFzfits}-\eqref{eq:thetafits} we provide fits only for halo masses defined by mean matter densities of $\bar{\rho}=300\rhomb$, our relations can be readily refitted using other mass definitions (e.g. $\bar{\rho}=500\bar{\rho}_{\rm cr}$) bearing in mind the resolution limitations imposed by Eq.~\eqref{eq:nhalomin}. Our method can also be straightforwardly applied to calibrate theoretical mass functions of other scalar-tensor theories characterized by a mass and environment dependent spherical collapse threshold. This is for example the case of the dilaton and symmetron models investigated in Brax et al.~(2012)~\cite{Brax:2012sd}. Note also that baryonic physics is likely to currently be irrelevant for the HMF ratios~\cite{Cataneo:2014kaa,Puchwein:2013lza}, and that any departures from DM-only predictions due to baryons could be included through a baseline GR mass function calibrated against hydrodynamical simulations (see e.g.~\cite{Bocquet:2015pva}). Analogous considerations might hold as well for the impact of massive neutrinos on the $f(R)$ to GR halo mass function ratio. It would be interesting to test the performance of our method on cosmological simulations incorporating massive neutrinos in both GR and $f(R)$ (see e.g. Baldi et al.~(2014)~\cite{Baldi:2013iza}). If the accuracy of our predictions remains unchanged when allowing a varying effective sum of the neutrino masses, then it would be sufficient to implement the prescription of Castorina et al.~(2014)~\cite{Castorina:2013wga} on the baseline GR mass function. Finally, in addition to Poisson noise it will be necessary to account for the uncertainty due to sample variance in order to unbiasedly constrain the low mass end of the HMF with forthcoming cluster number count data~\cite{Hu:2002we}. For the specific cosmological models of interest, this will require the calculation of the linear bias parameter, which in itself depends on the spherical collapse threshold~\cite{Sheth:1999mn,Tinker:2010my}. For $f(R)$ gravity, we should be able to use our effective linearly extrapolated overdensity (see Eq.~\eqref{eq:dceff}) to evaluate the linear bias and hence the sample variance contribution (Cataneo et al., in preparation) needed for a series of upcoming key cosmological analyses from ongoing and future cluster surveys. | 16 | 7 | 1607.08788 |
1607 | 1607.03386_arXiv.txt | We present new MUSE observations of a galaxy group probed by a background quasar. The quasar sightline passes between multiple $z=0.28$ galaxies, whilst showing at the same redshift low ionised metal line species, including Ca~{\sc ii}, Mg~{\sc i}, Mg~{\sc ii} and Fe~{\sc ii}. Based on the galaxy redshifts measured from the MUSE data, we estimate the galaxies to be part of a small galaxy group with a halo mass of $\approx6\times10^{12}$~M$_{\odot}$. We use the MUSE data to reveal the two dimensional dynamical properties of the gas and stars in the group galaxies, and relate these to the absorber kinematics. With these data we consider a number of scenarios for the nature of the gas probed by the sightline absorbers: a co-rotating gas halo associated with a single galaxy within the group; outflowing material from a single group member powered by recent star-formation; and cool dense gas associated with an intra-group medium. We find that the dynamics, galaxy impact parameters, star-formation rates, and the absorber strength suggest the cool gas can not be clearly associated with any single galaxy within the group. Instead we find that the observations are consistent with a superposition of cool gas clouds originating with the observed galaxies as they fall into the group potential, and are now likely in the process of forming the intra-group medium. | The interplay between galaxies and the gaseous media that surround them is a crucial element in the study of galaxy formation, providing insights into both the infall of material onto galaxies and the extent and impact of galactic outflows. Observing the baryonic material surrounding galaxies is primarily performed (at least beyond modest redshifts) using absorption line studies, either in the line of sight to background QSOs (quasi-stellar object) or galaxies. Much work has been performed on both the small scales, linking absorption systems to individual galaxies \citep[e.g.][]{2011ApJ...733..105K,2012ApJ...760L...7K,2012MNRAS.426..801B,2013Sci...341...50B,2013ApJ...777...59T,2013ApJ...763..148S,2013MNRAS.436.2650P,2013ApJ...776..115N,2015ApJ...812...83N,2016ApJ...818..171N,2014ApJ...792....8W,2015MNRAS.446.3178F}; and on statistical scales: analysing the large scale distribution of gas around galaxies \citep[e.g.][]{adelberger03,adelberger05,2006MNRAS.367.1251R,2006MNRAS.367.1261M,2009ApJ...701.1219C,2010MNRAS.402.2520S,2011MNRAS.414...28C,2012ApJ...750...67R,2012ApJ...751...94R,2014MNRAS.437.2017T,2014MNRAS.442.2094T,2014MNRAS.445..794T,2015MNRAS.450.2067T,2016MNRAS.460..590F,2016arXiv161009144B}. An important, but poorly constrained, element in connecting absorbers with individual galaxies however, is the galaxies' place in the large scale structure of the Universe. Many studies focus on so-called field galaxies with little or no quantification of the nearby galaxy population. Moreover, much of the star forming galaxy population resides in groups and this element is rarely considered in the analysis of absorption line systems, especially at high redshift. Absorbers are often treated as part of a distinct galaxy halo, whereas in reality the distribution of gas around galaxies is more complex and it may not be clear or reasonable to constrain where the halo of one galaxy finishes and the next begins \citep[e.g.][]{1994Natur.372..530Y,2002ASPC..254...72M}. In terms of absorption of background light by gas within galaxy groups, \citet{1994ApJ...427..696M} suggested that a significant fraction of H~{\sc i} absorption systems (with column densities above of $\sim10^{13}$~cm$^{-2}$) are the result of gaseous pressure-confined tidal debris, finding comparable space densities between low-column density H~{\sc i} absorbers and galaxy groups in the local Universe. Such absorption studies are complemented by H~{\sc i} emission maps of nearby galaxies and small groups \citep[e.g.][]{1999AJ....117..811H}, which show significant clumps of H~{\sc i} gas apparently tidally or ram-pressure stripped from the surrounding galaxies. Further work has been performed analysing the place of H~{\sc i} gas (over a range of column densities) within the large-scale-structure, with \citet{2002ApJ...565..720P}, for example, looking at the relationship between Ly{$\alpha$} absorbers and voids and super-clusters. Through a clustering analysis, they reported that Ly$\alpha$ absorbers (with column densities of $10^{13.2}$ to $10^{15.4}$~cm$^{-2}$) cluster with galaxies, but more weakly than galaxies cluster with each other. As such, \citet{2002ApJ...565..720P} conclude that such Ly$\alpha$ absorption systems are primarily associated with the large-scale structures of galaxies, i.e. filaments and groups, rather than individual galaxies \citep[see also][]{2014MNRAS.437.2017T}. Continuing such considerations of the relationship between absorbers and the large scale structure, \citet{2010MNRAS.402.1273C} found evidence that galaxy groups trace Mpc-scale structures that appear in both H~{\sc i} gas and galaxies based on the study of a triplet of closely separated QSO sightlines. Additionally, \citet{2012MNRAS.425..245T} found peaks in the Ly{$\alpha$} absorbers distribution around the edges of large voids: i.e. tracing the filamentary structures enclosing such voids (a topic further developed in \citealt{2016MNRAS.455.2662T}). Consistent with the column densities of H~{\sc i} emission maps discussed above, \citet{2000ApJ...543L...9C} presented the detection of a Lyman limit system (LLS) apparently associated with a pair of galaxies at $z=0.167$. Significantly, the reported LLS corresponded in velocity space to an O~{\sc vi} absorption system, which the authors claimed as tracing warm gas within the intra-group medium of this small galaxy group. Given that this detection of O~{\sc vi} absorption fits with the gas temperatures expected for the intra-group medium of typical galaxy groups (i.e. $T\sim10^5$~K), much interest has since been placed in tracing the potentially huge reservoirs of gas within galaxy groups using such absorption systems. Indeed, \citet{2006ApJ...643..680P} find O~{\sc vi} absorption arises in a diverse set of galactic environments including the halos of individual galaxies, galaxy groups, filamentary-like structures, and also regions devoid of luminous galaxies, albeit based on a relatively small sample of galaxies and absorbers. Similarly, \citet{2009ApJ...701.1219C} find that tidal debris in groups/galaxy pairs may be principally responsible for observed O~{\sc vi} absorbers (although this is somewhat in contrast to other findings, e.g. \citealt{2011Sci...334..948T}). \citet{2014ApJ...791..128S} presented a study of absorber measurements in quasar sightlines, observed with HST/COS, passing through 14 galaxy groups. Their far-UV spectra produced 14 warm ($T\geq10^5$~K) absorption systems consisting of broad Ly$\alpha$ and broad O~{\sc vi} absorption, which were each found to be coincident with spiral-rich groups or cluster outskirts. Their analysis tentatively favours the absorbers to be associated with the intra-group/cluster medium rather than the halos of individual galaxies, albeit with some systematic uncertainties. Supporting this finding, \citet{2015MNRAS.449.3263J} analyse the effect of environment on their analyses of O~{\sc vi}, finding that galaxies with nearby neighbours exhibit a modest increase in O~{\sc vi} covering fraction compared to isolated galaxies suggest that environmental effects play a role in distributing heavy elements beyond the enriched gaseous haloes of individual galaxies, whilst \citet{2016MNRAS.458..733P} identify O~{\sc vi} and broad Ly~$\alpha$ absorbers associated with $z=0.4$ galaxy group. However, although there has been much focus on O~{\sc vi} as a tracer of intra-group gas, in reality the intra-group medium is likely a complex mix of phases incorporating cool gas, potentially in clumpy structures or clouds, as well as the warmer gas phases \citep{2016ApJ...830...87S}. Mg~{\sc ii} as a doublet observed at optical wavelengths represents an excellent tracer of cool gas and as such it is an important and complementary diagnostic to add to the O~{\sc vi} studies discussed above. Only a small number of Mg~{\sc ii}-galaxy group associations have been reported in detail however: \citet{2006MNRAS.368..341W} detected Mg~{\sc ii} absorption coincident with a galaxy group at $z=0.666$; \citet{2010MNRAS.406..445K} reported a strong Mg~{\sc ii} absorber ($W_r(2796)=1.8$~\AA) coincident with a galaxy group with 5 observed members at $z=0.313$; and \citet{2013MNRAS.432.1444G} investigated a large equivalent width (EW) Mg~{\sc ii} absorber detected in close proximity to a galaxy group at $z\sim0.5$. The latter study concluded that the absorption system is most likely associated with the intra-group medium, with the authors favouring a scenario whereby the absorption originates in tidal debris. To give a sense of the number of potential Mg~{\sc ii} absorber-galaxy group associations, we note that from a sample of $\sim200$ Mg~{\sc ii} absorbers and nearby galaxies, \citet{2013ApJ...776..114N} found that $\approx40$ appeared to be associated with group galaxies or were ambiguous in what galaxy they may be associated with. However, due to the complexity of disentangling the relationship between absorbers found in a group environment and the group galaxies, most studies have taken the approach of either removing the group absorbers from their analysis \citep[e.g.][]{2013ApJ...779...87C,2013ApJ...776..115N,2015ApJ...812...83N} or simply taking the group galaxy with the smallest impact parameter \citep[e.g.][]{2016ApJ...833...39S} as the absorber `host'. Here, we present serendipitous MUSE observations of a galaxy group observed at $z=0.2825$ in the foreground of a bright $z=1.71$ quasar and coincident with significant metal line absorption features. MUSE represents an excellent probe of galaxy-group structures around QSO sightlines, providing as it does a blind survey over an area of $1'\times1'$ \citep[e.g.][]{2016MNRAS.462.1978F,2017MNRAS.464.2053P}. Our observations are nominally part of the QSO Sightline and Galaxy Evolution (QSAGE) survey, which aims to provide deep galaxy data around HST/STIS observed QSO sightlines, primarily using the WFC3 G141 grism (HST Program 14594, PI: R. Bielby). In Section~\ref{sec:observations} of this paper, we present our MUSE observations and give an overview of the available ancillary galaxy data and the sightline spectral data. In Section~\ref{sec:analysis}, we present an analysis of the galaxy and group properties alongside an analysis of the sightline absorbers. In Section~\ref{sec:discussion}, we discuss the relationship between the group galaxies and the absorbers and the implications for the nature of the group environment. We present our conclusions in Section~\ref{sec:conclusions}. We assume a Planck 2013 cosmology \citep{2014A&A...571A..16P} throughout. All quoted distances are in proper coordinates unless stated otherwise. | 16 | 7 | 1607.03386 |
|
1607 | 1607.00385_arXiv.txt | We present radiation transfer models of rotating young stellar objects (YSOs) with hotspots in their atmospheres, inner disk warps and other 3-D effects in the nearby circumstellar environment. Our models are based on the geometry expected from the magneto-accretion theory, where material moving inward in the disk flows along magnetic field lines to the star and creates stellar hotspots upon impact. Due to rotation of the star and magnetosphere, the disk is variably illuminated. We compare our model light curves to data from the Spitzer YSOVAR project \citep[e.g.][]{morales11, cody14} to determine if these processes can explain the variability observed at optical and mid-infrared wavelengths in young stars. We focus on those variables exhibiting ``dipper" behavior that may be periodic, quasi-periodic, or aperiodic. We find that the stellar hotspot size and temperature affects the optical and near- infrared light curves, while the shape and vertical extent of the inner disk warp affects the mid-IR light curve variations. Clumpy disk distributions with non-uniform fractal density structure produce more stochastic light curves. We conclude that the magneto-accretion theory is consistent with certain aspects of the multi-wavelength photometric variability exhibited by low-mass YSOs. More detailed modeling of individual sources can be used to better determine the stellar hotspot and inner disk geometries of particular sources. | Multi-wavelength studies of the variability of young stellar objects (YSOs) probe the combined stellar and circumstellar properties of newly forming stars along with angular momentum driven phenomena such as stellar rotation and binary orbital motion. Optical and near-infrared data are sensitive to the stellar photosphere (hot and cool spots), and other energetically ``hot" regions (accretion columns, chromospheres), as well as scattering from the circumstellar material. Observations at mid-IR and longer wavelengths offer a new perspective as they are sensitive to variability associated with ``warm" or ``cool" regions --- the disks and envelopes of YSOs. Figure \ref{f:image} illustrates that different wavelengths dominate different regions by showing a 3-color plot of one of our models of a spotted star surrounded by a warped accretion disk. The optical variability of accreting YSOs has been successfully interpreted in the context of the magnetospheric accretion model. In this model, the inner disk is truncated, and material flows from the disk to the star along stellar magnetic field lines \citep{ghosh78,konigl91}. As the free-falling material reaches the star, the kinetic energy is dissipated in shocks at the stellar surface \citep{konigl91}. The stellar magnetic field is often inferred to not be aligned with the rotation axis based on line emission modeling \citep{donati11}, resulting in photometric modulation as the shock columns move in and out of view \citep{mahdavi98, gregory11}. Strong H$\alpha$ (and other) line emission and blue excesses are produced by the inflowing gas and shock columns \citep{hartmann94,gullbring98,muzerolle01}. The lightcurves of accreting YSOs show variations on a variety of timescales and with a variety of color-magnitude effects \citep{Herbst94}. Timescales on the order of a few hours track material in free-fall from the inner disk to the stellar surface. The time for an inner disk asymmetry to transit the stellar surface is $\sim 0.3$ days on average. The stellar rotational modulation is typically $\sim1-8$ days \citep{rebull04}. Disk accretion rates and magnetospheric structure changes occur on timescales of days to weeks to years. The color variability ranges from essentially colorless amplitude variability, indicating achromatic or ``black" processes, to large color variability, indicating substantial changes in accretion or extinction. \citet{morales11} and \citet{cody14} presented results of multi-wavelength photometric monitoring of the Orion Nebula Cluster (ONC) and NGC 2264, as part of the young stellar object variability (YSOVAR) project that also includes many smaller clusters, as summarized in \citet{rebull14}. Among thousands of YSOs, 70\% of those with mid-IR excess are variable at levels typically 0.1 to 0.2 magnitudes but some have amplitudes as high as 0.5 mag. The YSOs observed exhibit many different behaviors, but can be grouped into a few main categories based on light curve morphology: periodic/quasi-periodic, dippers (both periodic and irregular), bursters/accretors (almost always irregular), stochastic variables, and stars showing either brightening or fading trends covering the full duration of the time series. The periodic light curves can have relatively symmetric and regular flux variations, but there is also a sub-class of the periodic sources with asymmetric lightcurves that show pronounced ``dips." Other light curves exhibit quasi-periodic behavior, with additional upward or downward trends in brightness that render them not detected as significantly periodic under Fourier analysis though with semi-ordered and repeated variations diagnosed using the ``Q" statistic of \citet{cody14}. Like the periodic sources, the quasi-periodic objects may be roughly symmetric in their brightness variations, or with pronounced ``dips." Such ``dipper'' sources may be periodic with regular dips in brightness, quasi-periodic as described above, or irregular with dips occurring much more stochastically relative to a defined stable flux level. An obvious physical interpretation for this category is variable extinction, but we also propose an alternate model based on variable illumination. Another YSOVAR category is the inverse of the dippers, the ``bursters", that are characterized by flux bursts and excess brightness peaks on various time scales, with mostly constant flux otherwise, and irregular repetition. A popular interpretation for this category is variable accretion. Some light curves are neither periodic nor quasi-periodic but exhibit large and/or small, stochastic, brightness variations over a few days, possibly due to a combination of extinction and accretion events \citep{cody14,stauffer15}. The ``trender" category is likely dominated by processes occurring outside the magnetospheric region, where the dynamical time scales are longer than the few days to week long variations that typify the other categories. In this paper we present models intended to apply only to the various forms of periodic and quasi-periodic lightcurves, especially those of the ``dipper" variety. Periodicity naturally arises from the rotation of the star and Keplerian rotation within the disk. We illustrate how variations in accretion properties and inner disk geometry affect the brightness, including wavelength-dependent effects, which can be used to infer the physical processes responsible for the observed variations due to stochastic accretion. In section 2 we describe the star-hotspot-accretion disk geometry we adopt for our radiation transfer models. Section 3 presents the photometric and polarimetric variability from our models. In Section 4 we compare our models to observations and we summarize our findings in Section 5. | Using our models we can predict percentages of stars in each of the variability categories. We assume either a stable (ordered dipole behavior) or unstable disk (chaotic magnetic field, clumpy disk) for all of the stars, and treat these two types of disks as separate cases. The numbers we use in the following paragraphs are estimates from extensive modeling of a grid of inclination angles (e.g., variability is present at 65\arcdeg but dipping at 70\arcdeg). We will first assume that all of the accretion disks are stable, and that above $i=77^\circ$, the wavelengths we are observing would be extincted and the stars will be too faint to detect. We will therefore normalize the models over the range $0^\circ \le i \le 77^\circ$. Models with $0^\circ \le i \le 20^\circ$ will show little variability, since a high latitude hotspot will be visible throughout the entire rotation period. Thus we estimate that about 8\% of stars will show no variability. For inclinations $20^\circ \le i \le 67^\circ$ we expect to see some sort of periodic variation, so 71\% of stars with stable magnetospheric accretion should show this form of variability. Dippers are likely to show up for $67^\circ \le i \le 77^\circ$, which is 21\% of the stars. Now we consider the statistics if all of the accretion disks are unstable and therefore have 3-D variations in their disk structure rather than one or two warps. For $0^\circ \le i \le 50^\circ$ we predict that there will be no variation since with an unstable accretion disk there are not strong hotspots or a pattern of variation from the accretion disk except at high angles of inclination. This means that about 47\% would be non-variable. For $50^\circ \le i \le 77^\circ$ there will be aperiodic variations, which is about 53\% of the stars. Morales Calderon et al. (2011) report that about 70\% of the stars observed were variable. Using this percentage we can try to match our predictions with the observational data. In order to get around 30\% of sources that are non-variable we can estimate that about 50\% of the disks must be stable and 50\% are unstable, giving us 28\% that are not variable. Next we can apply this same 50\% to the rest of the categories to come up with some predictions. Table~4 summarizes this statistical analysis of our models. It is important to note that these results are only for Class II objects and do not include Class I objects that are heavily embedded, or spotted weak-lined T-Tauri stars, which are usually categorized as periodic or non-variables. These statistics include all the main sources of variability, since for Class II objects variability seems dominated by disk-related effects rather than the underlying cool spot rotational modulation, which is undoubtedly there, but not included in our models. In an optical study, \citet{cody14} find that only 3\% of a disk-selected sample showed purely periodic behavior due to spots. \begin{table}[ht] \centering \caption{Occurrence of YSOVAR classes in our models} % \begin{tabular}{c c c c} \hline \hline & Stable Disks & Unstable Disks & Total \\ [0.5ex] % \hline % Non-Variable & 8\% & 47\% & 28\% \\ % Periodic & 71\% & - & 36\% \\ Periodic Dippers & 21\% & - & 10\% \\ Wild & - & 53\% & 26\% \\ [1ex] \hline % \end{tabular} \label{table:predic} % \end{table} \vspace{20pt} | 16 | 7 | 1607.00385 |
1607 | 1607.02663_arXiv.txt | {By modeling the broadband spectral energy distributions (SEDs) of a typical flat spectrum radio quasar (FSRQ, 3C 279) and two GeV narrow-line Seyfert 1 galaxies (NLS1s, PMN J0948+0022 and 1H 0323+342) in different flux stages with the one-zone leptonic models, we find a universal correlation between their Doppler factors ($\delta$) and peak luminosities ($L_{\rm c}$) of external Compton scattering bumps. Compiling a combined sample of FSRQs and GeV NLS1s, it is found that both FSRQs and GeV NLS1s in different stages and in different sources well follow the same $\delta$--$L_{\rm c}$ correlation. This indicates that the variations of observed luminosities may be essentially due to the Doppler boosting effect. And the universal $\delta$--$L_{\rm c}$ relation between FSRQs and GeV NLS1s in different stages may be further evidence that the particle acceleration and radiation mechanisms for the two kinds of sources are similar. In addition, by replacing $L_{\rm c}$ with the observed luminosity in the \emph{Fermi}/LAT band ($L_{\rm LAT}$), this correlation holds, and it may serve as an empirical indicator of $\delta$. We estimate the $\delta$ values with $L_{\rm LAT}$ for 484 FSRQs in the \emph{Fermi}/LAT Catalog and they range from 3 to 41, with a median of 16, which are statistically consistent with the values derived by other methods. | % Flat spectrum radio quasars (FSRQs) and BL Lacertae objects (BL Lacs) are referred to as blazars, whose broadband spectral energy distributions (SEDs) are thought to be dominated by the jet emission. FSRQs are different from BL Lacs for having significant emission lines. It was proposed that many radio-loud (RL) narrow-line Seyfert 1 galaxies (NLS1s) display blazar characteristics and maybe also host relativistic jets\footnote{Gu et al. (2015) studied the compact radio structures of 14 NLS1s with Very Long Baseline Array observations at 5 GHz and reported that 50\% of the sources show a compact core only and the remaining 50\% exhibit a core-jet structure.} (Zhou et al. 2003; Yuan et al. 2008); this has been confirmed by the detection of $\gamma$-ray emission from NLS1s by \emph{Fermi}/LAT (Abdo et al. 2009; D'Ammando et al. 2012) and the observations of Kiloparsec--parsec scale radio structures (Doi et al. 2012; Gu et al. 2015), especially the observation of apparent superluminal velocity in the jet of SBS 0846+513. The broadband SEDs of GeV NLS1s are similar to blazars (Abdo et al. 2009; Zhang et al. 2013b; Paliya et al. 2013; Sun et al. 2014, 2015; Kreikenbohm et al. 2016; Paliya \& Stalin 2016) and their $\gamma$-ray emission is also dominated by external Compton scattering (EC) process by photons from their broad-line regions (BLRs) (e.g., Sikora et al. 1994; Ghisellini et al. 2009; Zhang et al. 2014; Chen et al. 2012; Liao et al. 2014; Sun et al. 2015). GeV NLS1s have significant emission lines, and sometimes a blue bump of the disk thermal radiation is observed in their SEDs. Therefore, the circumnuclear environment in NLS1s is analogous to that in FSRQs. Recently, on the basis of one-zone leptonic jet models, Sun et al. (2015) reported that the jet property of GeV NLS1s, including their jet power, radiation efficiency, and magnetization parameter, is indeed a bridge between FSRQs and BL Lacs, but more analogous to FSRQs than BL Lacs. Further more, Zhang et al. (2015) suggested a BL Lac--NLS1--FSRQ sequence with the increase of their BLR luminosity and Eddington ratio, which may correspond to the change of the accretion disk structure and the transformation of the dominant mechanism for jet launching. The luminosities of blazars are thought to be boosted since the emitting regions move with relativistic velocity and small viewing angle ($\theta$). Recently, Richards \& Lister (2015) reported that the jets of RL NLS1s are aligned at moderately small angles to the line of sight, which is similar to blazars. The active galactic nuclei (AGNs) that are not detected with the \emph{Fermi}/LAT have, on average, lower Doppler factors than those that are detected with the \emph{Fermi}/LAT (Lister et al. 2015). The measurements of these parameter values are very crucial for understanding the physics of jets (e.g., Nokhrina et al. 2015), for examining the unified models of AGNs (e.g., Hovatta et al. 2009; Savolainen et al. 2010), and even for investigating the intrinsic radiation physics of blazars with gamma-ray bursts (Wu et al. 2011; Wang et al. 2011; Nemmen et al. 2012; Zhang et al. 2013a). Some approaches have been proposed to estimate the Doppler factor values of AGNs. With the Very Long Baseline Interferometry (VLBI) measurements of the core angular size and radio flux, Ghisellini et al. (1993) estimated the $\delta$ values of $\sim$100 AGNs by comparing the observed X-ray fluxes to that predicted by the Self-Synchrotron-Compton scattering model. The derived $\delta$ values with this method usually have large uncertainty since it needs simultaneous X-ray and VLBI observations and strongly depends on the turnover values (L\"{a}hteenm\"{a}ki \& Valtaoja 1999). Jorstad et al. (2005) also used the VLBI observation data to derive the Doppler factors by comparing the flux decline timescale ($\tau_{\rm obs}$) with the light-travel time ($\tau_{\rm int}$) across the emitting region ($\tau_{\rm obs}\sim\tau_{\rm int}\delta$). Another more popular way to estimate the Doppler factors is to obtain the variability brightness temperatures of sources using total flux density variations (L\"{a}hteenm\"{a}ki \& Valtaoja 1999; Hovatta et al. 2009), which is boosted by $\delta^3$ compared with the intrinsic brightness temperature of the source. Although the superluminal motions in many of sources were resolved by VLBI observations (Homan et al. 2001; Kellermann et al. 2004; Jorstad et al. 2005; Piner et al. 2007; Lister et al. 2013), a quantitative assessment of the beaming parameters (i.e., the bulk Lorentz factor $\Gamma$ or the velocity of emitting region and the viewing angle $\theta$) is still lacking. With the transparency condition, one may also estimate the lower limit of $\delta$ (e.g., Fan et al. 2014). Theoretically, by modeling the observed broadband SEDs, the Doppler factors can also be constrained (Zhang et al. 2012, 2014, 2015; Sun et al. 2015; Yan et al. 2014; Kang et al. 2014). It is interesting that a tentative correlation between $\delta$ and the peak luminosities of the EC bumps ($L_{\rm c}$) is found for FSRQ 3C 279 (Zhang et al. 2013c) with four SEDs and two GeV NLS1s (Sun et al. 2015). In this paper, we firstly further test this correlation in individual sources with 14 SEDs of 3C 279. And then we compile a sample of FSRQs and GeV NLS1s to study the correlation between $L_{\rm c}$ and $\delta$ for different sources and also investigate whether this correlation can be used to estimate $\delta$ with the observed luminosity. The analogous observations in both FSRQs and GeV NLS1s also motivate us to explore whether they share the same $\delta$--$L_{\rm c}$ relation and the physics of this correlation. Model and SED fitting and the $\delta$--$L_{\rm c}$ correlation in different stages for 3C 279 are presented in \S 2. The $\delta$--$L_{\rm c}$ correlation in different sources is described in \S 3. The possible physical implications of this correlation are discussed in \S 4. Using this relation to derive the Doppler factors of FSRQs in \emph{Fermi}/LAT Third Source Catalog (3FGL) and comparing the results with others are reported in \S 5. Summary and conclusions are given in \S 6. | By modeling the broadband SEDs of a typical FSRQ 3C 279 and two GeV NLS1s PMN J0948+0022 and 1H 0323+342 in different stages with a one-zone leptonic model, we found a correlation between the Doppler factor ($\delta$) and EC peak luminosity ($L_{\rm c}$). We then compiled a sample of 30 FSRQs and 5 GeV NLS1 galaxies and found that the $\delta$--$L_{\rm c}$ correlation holds well. The main results are summarized as follows: \begin{itemize} \item $L_{\rm c}$ is strongly correlated with $\delta$ for both FSRQs and GeV NLS1s, and the two kinds of AGNs form a clear sequence in the $\delta$--$L_{\rm c}$ plane, which may imply a unified picture of the particle acceleration and cooling mechanisms in the comoving frame for the two kinds of sources. Therefore, the observed differences of $L_{\rm c}$ in different stages and different sources may be essentially due to their different Doppler factors. \item Replacing $L_{\rm c}$ with the observed luminosity in the \emph{Fermi}/LAT band ($L_{\rm LAT}$), this correlation holds. The linear fitting result of the $\delta$--$L_{\rm LAT}$ relation is well consistent with the $\delta$--$L_{\rm c}$ relation within the errors. $L_{\rm LAT}$ may serve as an empirical indicator of $\delta$. \item We estimated the $\delta_{\rm LAT}$ values with $L_{\rm LAT}$ for 484 FSRQs in 3FGL and they range from 3 to 41, with a median of 16. The derived values of $\delta$ are statistically consistent with the values calculated by other methods. \end{itemize} | 16 | 7 | 1607.02663 |
1607 | 1607.00499_arXiv.txt | This work presents the AYT2 line list: a comprehensive list of 114 million $^{1}$H$_2$$^{32}$S vibration-rotation transitions computed using an empirically-adjusted potential energy surface and an {\it ab initio} dipole moment surface. The line list gives complete coverage up to 11000 \cm\ (wavelengths longer than 0.91 $\mu$m) for temperatures up to 2000 K. Room temperature spectra can be simulated up to 20000 \cm\ (0.5 $\mu$m) but the predictions at visible wavelengths are less reliable. AYT2 is made available in electronic form as supplementary data to this article and at \url{www.exomol.com}. | The investigation of the sulphur chemistry in space is a subject of the active research \citep{01RuKixx.SO,04WaCaCe.H2S,06ViLoFe.H2S,09ZaMaFr.H2S,11AlMaMa.H2S,13ReSeBa.H2S}. In particular \citet{13ReSeBa.H2S} studied the atmospheric composition and the spectra of earth-like exoplanets with sulphur compounds such as hydrogen sulphide (H$_{2}$S) and sulphur dioxide (SO$_{2}$) using a one-dimensional photochemistry model and associated radiative transfer model to investigate sulphur chemistry in atmospheres ranging from reducing to oxidising. \citet{06ViLoFe.H2S} used thermochemical equilibrium and kinetic calculations to model sulphur chemistry in giant planets, brown dwarfs, and extrasolar giant planets, and found that H$_2$S is the dominant S-bearing gas throughout substellar atmospheres and approximately represents the atmospheric sulphur inventory. Therefore, observations of H$_{2}$S in these objects should provide a good estimate of their atmospheric sulphur content. H$_2$S has been, however, ruled out as a potential biosignature in atmospheres of exoplanets according to a biomass-based model study by \citet{13SeBaHu.exoplanet}. H$_2$S has long been known in the interstellar medium \citep{72ThWiKu.H2S} and is important in star-forming \citep{04WaCaCe.H2S,15NeGoGe.H2S} and circumstellar \citep{ 93OmLuMo.H2S} regions. \citet{11AlMaMa.H2S} detected H$_{2}$S for the first time in galaxy M82, where they studied the chemical complexity towards the central parts of the starburst galaxy, and investigated the role of certain molecules as tracers of the physical processes in the galaxy circumnuclear region. \citet{01RuKixx.SO} found evidence for SO$_{2}$, SO and H$_{2}$S sulphide in Io's exosphere. For Venus, the H$_{2}$S composition of the atmosphere at altitudes below 100 km was studied by \citet{85ZaMoxx.H2S} and \citet{06BeMoTa.H2S}. Determination of the abundances of gases such as CO, SO$_{2}$, OCS, S$_{2}$ or H$_{2}$S near the surface is important to constrain the oxidation state of the lower atmosphere and surface, and determine the stability of various minerals. Also, measurements at higher altitudes of, for example, SO$_{3}$, SO or elemental sulphur, are needed to better understand the sulphur cycle and the chemistry at work below the cloud base. Conversely a recent search for H$_2$S in volcanic emissions on Mars failed to detect any \citep{15KhViMu}. H$_2$S is known to be present in comets \citep{02BiBoCr.comets} being first detected by \citet{91BoCoCr.comets}. On Earth naturally occurring H$_2$S is associated with volcanic activity \citep{12HoHoLa.H2S}. Gaseous H$_2$S is also detected in a number of other situations including emissions from waste water \citep{12LlEsMa.H2S} and as a by-product of industrial processes \citep{13SzMoGu.H2S}. Known experimental absorption spectra of H$_2$S molecule cover the region from the microwave up to the visible (0.6 $\mu$m). Observations include transitions belonging to 59 vibrational bands associated with different 14 polyads \citep{12PoLaVo.H2S}, where the polyad number is defined as $n=v_{1}+v_{2}/2+v_{3}$, where $v_{i}$ are standard normal-mode vibrational quantum numbers. The rotational band has received attention from many experimentalists \citep{53BuJRGo.H2S,66HuDyxx.H2S,68CuKeGa.H2S,69MiLeHa.H2S,71Huxxxx.H2S,72HeCoDe.H2S,83FlCaJo.H2S,85BuFeMe.H2S,94YaKlxx.H2S,95BeYaWi.H2S,jt558,14CaPuGa.H2S}. The first bending vibrational band ($\nu_2$) at 1183~cm$^{-1}$ was studied by \citet{82LaEdGi.H2S}, \citet{83Stxxxxa.H2S} and \citet{96UlMaKo.H2S}. The two fundamental stretching; symmetric ($\nu_1$) and asymmetric ($\nu_3$) lying at 2615 and 2626 cm$^{-1}$, respectively, are not isolated but overlapped with strong Coriolis and Fermi resonance interactions. The first triad region ($2\nu_2,\nu_1,$ and $\nu_3$) was studied by \citet{81GiEdxx.H2S}, the second triad region ($3\nu_2,\nu_1+\nu_2,$ and $\nu_2+\nu_3$) was studied by \citet{69SnEdxx.H2S} and \citet{96UlOnKo.H2S}, while \citet{98BrCrCr.H2S} studied these two triad regions simultaneously. The 4500~--~5600 cm$^{-1}$ spectral region was investigated by \citet{97BrCrCr.H2S}. \citet{04BrNaPoa.H2S} and \citet{05UlLiBe.H2S} recorded and analysed the transitions in the region 5700~--~6600~\cm. In the 7300~--~7900~\cm\ region, more than 1550 transitions up to $J=14$ were recorded and analysed by \citet{04UlLiBe.H2S}. The absorption spectrum in the region 8400~--~8900~\cm\ was recorded by \citet{04BrNaPob.H2S}. \citet{94ByNaSm.H2S} recorded and analysed spectra between 2000 to 11 147 \cm. A number of shorter wavelength regions have been studied, namely, 9540~--~10~000~\cm\ by \citet{03DiNaHu.H2S}, 10~780~--~11~330~\cm\ by \citet{01NaCaxxb.H2S}, 11~930~--~12~300~\cm\ by \citet{94GrRaSt.H2S} and \citet{95FlGrRa.H2S}, 12~270~--~12~670~\cm\ by \citet{97VaBiCa.H2S}, near 13~200 \cm\ by \citet{99CaFlxx.H2S}, 14~100~--~14~400~\cm\ by \citet{98FlVaCa.H2S}, and 16~180~--~16~440~\cm\ by \citet{01NaCaxxa.H2S}. This situation is summarised in Fig.~\ref{summary-for-experimenatl-work-on-H2S}. All this work has been performed using cool samples, that is below 300 K. The highest recorded value of the rotational quantum number $J$ is 22 in the rotational band region, and the highest predicted value is 27 in the same region. Altogether around 10~000 ro-vibrational energy levels are known from these experiments. The spectroscopic data for the H$_{2}$S molecule has been used to populate various spectroscopic databases. Table~\ref{databases} summarises the contents of the HITRAN-2012 \citep{jt557}, GEISA \citep{jt504,jt636}, W@DIS \citep{12PoLaVo.H2S}, CDMS \citep{01MuThRo.db,05MuScSt.db}, and JPL \citep{98PiPoCo} databases. All these databases contain data resulting from fitted effective Hamiltonians, apart from W@DIS which contains only measured transitions without intensities. The 2012 release of HITRAN \citep{jt557} updated HITRAN 2008 \citep{jt453} using extra rotational data rotational band of H$_{2}$S spectrum from \citet{jt558} and the data published in IAO LMS Spectra [spectra.iao.ru]. While absorption spectra of H$_2$S at elevated temperature have recently been recorded in the ultra-violet by \citet{15GrFaCl.H2S}, we are unaware of any high resolution experimental studies of H$_2$S infrared spectra at higher than ambient temperature. H$_2$S cross-sections have been measured up to $T = 50$~C as part of PNNL database \citep{PNNL}. \begin{table*} \begin{center} \caption{Summary for the H$_{2}$S spectral data available in different databases.} \label{databases} \begin{tabular}{c c r r r r} \hline \hline Database & isotopologue &\multicolumn{1}{c}{Bands} &\multicolumn{1}{c}{Transitions} &\multicolumn{1}{c}{Wavenumber$_{min}$}& \multicolumn{1}{r}{Wavenumber$_{\rm max}$}\\ & & \multicolumn{1}{c}{\#} & \multicolumn{1}{c}{\#} & (\cm) & (\cm) \\ \hline HITRAN 2012 & H$_{2}$$^{32}$S & 49 & 36561 & 2 & 11330 \\ & H$_{2}$$^{33}$S & 19 & 6322 & 5 & 11072 \\ & H$_{2}$$^{34}$S & 24 & 11352 & 5 & 11227 \\ GEISA & H$_{2}$$^{32}$S & 14 & 12330 & 2 & 4257 \\ & H$_{2}$$^{33}$S & 8 & 3564 & 5 & 4098 \\ & H$_{2}$$^{34}$S & 8 & 4894 & 5 & 4171 \\ W@DIS & H$_{2}$$^{32}$S & 59 & 34148 & 1 & 16437 \\ CDMS & H$_{2}$$^{32}$S & 1 & 1501 & 1 & 554 \\ & H$_{2}$$^{33}$S & 1 & 4759 & 1 & 402 \\ & H$_{2}$$^{34}$S & 1 & 990 & 1 & 444 \\ JPL & H$_{2}$$^{32}$S & 1 & 1525 & 1 & 333 \\ \hline \hline \end{tabular} \end{center} \end{table*} \begin{figure} \centering {\leavevmode \epsfxsize=9.0cm \epsfbox{summary.eps}} \caption{Summary of experimental work on the H$_{2}$S absorption spectrum overlayed with our theoretical cross-sections. All spectra were recorded at room temperature. The polyad number of each band is also given.} \label{summary-for-experimenatl-work-on-H2S} \end{figure} A number of theoretical studies have considered H$_{2}$S. Its ro-vibrational spectrum was calculated by \citet{89SeCaZi.H2S}, \citet{jt266} and \citet{04TyReSc.H2S}. In the work of \citet{89SeCaZi.H2S}, the room temperature absorption ro-vibrational spectrum of H$_{2}$S was calculated variationally for the pure rotational and $\nu_{2}$, $2\nu_{2}$, $\nu_{1}$, and $\nu_{3}$ transitions from $J = 0$ to 13, where the full account of the anharmonicity effects and ro-vibration couplings were considered. \citet{89SeCaZi.H2S} calculated the vibrational band origins of the fundamental transitions with accuracy better than 10~\cm, and the ro-vibrational transitions to within a few tenths of a \cm~for low $J$'s and up to a few \cm~for high $J$'s; their work was also extended to deuterated isotopologues \citep{jt99}. The anomalies in the spectral intensity for this molecule were obtained qualitatively. \citet{jt266} calculated the vibrational band origins of H$_{2}$$^{32}$S with accuracy of 29~\cm~up to 14~300~\cm, and the rotational transitions of the ground vibrational state for $J = 17$ with deviations from the experimental values from 2 to 10~\cm. \citet{04TyReSc.H2S} used a spectroscopically-determined potential energy surface to compute the spectrum in the interval 0~--~8000~\cm, for $J$ up to 18 and with and intensity cut-off $\leq~10^{-27}$~cm$^{-1}$/(molecule$\times$cm$^{-2}$); they reported calculated transitions to be better than 0.01~\cm~for the line positions and 1-3\% in the intensities for the strong and medium lines and to $\sim$10\% for the weak lines for room temperature conditions. Very recently \citet{15CaLexx.H2S} explored the use of ladder operators to study the vibrational spectrum of this system. Remote detection of H$_{2}$S relies on well-characterised laboratory spectra. At higher temperatures the resulting list of transitions becomes very extensive and is best calculated using a robust theoretical model \citep{jt511}. The ExoMol project \citep{jt528} aims to provide molecular line lists for exoplanet and other atmospheres with a particular emphasis on hot species. In this work we present a comprehensive, hot linelist of vibration-rotation transitions of $^{1}$H$_2$$^{32}$S. This line list should be appropriate for temperatures up to 2000~K. The methodology used, which is discussed in the following section, closely follows that used to generate comprehensive hot line lists for triatomic species such as water \citep{jt378,jtpoz} and, recently, SO$_2$ \citep{jt635}. Section 3 presents our line list computations. Results and comparisons are given in section 4. Section 5 gives our conclusions. | A new hot line list for H$_2$S, called ATY2, has been computed containing 114 million transitions. The line list is divided into an energy file and a transitions file. This is done using the new ExoMol format \citep{jt631}. The full line list can be downloaded from the CDS, via \url{ftp://cdsarc.u-strasbg.fr/pub/cats/J/MNRAS/xxx/yy}, or \url{http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/MNRAS//xxx/yy}, as well as the ExoMol website, \url{www.exomol.com}. The line lists and partition function together with auxiliary data including the potential parameters and dipole moment functions, as well as the absorption spectrum given in cross section format \citep{jt542}, can all be obtained also from \url{www.exomol.com} as part of the extended ExoMol database \citep{jt631}. | 16 | 7 | 1607.00499 |
1607 | 1607.07046_arXiv.txt | In this paper, we scrutinize the effect of spectral index distribution on estimating the AGN (active galactic nucleus) radio luminosity function (RLF) by a Monte Carlo method. We find that the traditional bivariate RLF estimators can cause bias in varying degree. The bias is especially pronounced for the flat-spectrum radio sources whose spectral index distribution is more scattered. We believe that the bias is caused because the $K$-corrections complicate the truncation boundary on the $L-z$ plane of the sample, but the traditional bivariate RLF estimators have difficulty in dealing with this boundary condition properly. We suggest that the spectral index distribution should be incorporated into the RLF analysis process to obtain a robust estimation. This drives the need for a trivariate function of the form $\Phi(\alpha,z,L)$ which we show provides an accurate basis for measuring the RLF. | The luminosity function (LF), which provides a census of the galaxy and active galactic nucleus (AGN) populations over cosmic time, has been an important and also common tool for understanding the evolution of galaxies and ANGs \citep[][]{2008ApJ...682..874K}. With the abundance of multi-wavelength observed data for AGNs, AGN LFs have been estimated for various wave bands, such as the optical LF \citep[OLF, e.g.,][]{2000MNRAS.317.1014B}, the X-ray LF \citep[XLF, e.g.,][]{2000A&A...353...25M}, the $\gamma$-ray LF \citep[GLF, e.g.,][]{2012ApJ...751..108A,2013MNRAS.431..997Z} and the radio LF \citep[RLF, e.g.,][]{2001MNRAS.322..536W,2016ApJ...820...65Y,1990MNRAS.247...19D}. In this work, we will focus on the AGN RLFs. In an actual survey, only a very limited number of objects in the universe can be observed. How many sources entering the sample depends on the survey depth and selection function. Thus the estimation of LFs is inevitably based on a truncated sample of objects. Another difficulty is brought about by $K$-correction. It not only affects the accurate determination of intrinsic luminosity of individual sources, but also complicates the process of translating flux selection limits into luminosity selection limits, even for a single band selected survey \citep{2007ApJ...661..703S,2016arXiv160407493L}. The truncation boundary on the $L-z$ plane of a real sample is often a complicated region, but not a regular curve. Therefore, $K$-correction can affect the estimation of LFs by making it difficult to define the truncation boundary. For example, \citet{2004MNRAS.351..541I} found that a wide range of $K$-corrections being applied across different galaxy types can bias the shape of the global LF. For the radio AGNs, their spectra are frequently characterized as a simple power-law, $S\propto\nu^{-\alpha}$, and the $K$-correction has a simple form of $K(z)=(1+z)^{1-\alpha}$. Then the $K$-corrections of radio sources can be represented by their spectral properties (include spectral curvature, spectral index $\alpha$ and its distribution). Several Authors have discussed the potential problem of spectral curvature and its effect on obtaining reliable K-corrections for distant flat-spectrum radio sources \citep{1985MNRAS.217..601P,1996Natur.384..439S,2005A&A...434..133W} and also for steep-spectrum sources \citep{2011MNRAS.416.1900R}. Particularly, \citet{2000MNRAS.319..121J} highlighted the effect of spectral curvature that removes the evidence for the rapid decline in number density at high redshift suggested by \citet{1996Natur.384..439S}, and suggested that curvature would need to be incorporated in a full analysis of the RLF. Nevertheless, a recent study of \citet{2012MNRAS.422.2274C} pointed out that the effect of curvature only becomes important at higher frequencies ($\nu > 5 GHz$) and it can be avoided by measuring the spectra at lower frequencies. In the same paper, \citet{2000MNRAS.319..121J} also highlighted the importance of a distribution in spectral index in the parametric modeling of RLF \citep[also see][]{2001MNRAS.327..907J}. In this work, we analyze the bias caused by traditional bivariate RLF estimators based on Monte Carlo simulations, and further prove the necessity of incorporating the spectral index distribution into the RLF calculation. Throughout the paper, we adopt a Lambda Cold Dark Matter cosmology with the parameters $\Omega_{m}$ = 0.27, $\Omega_{\Lambda}$ = 0.73, and $H_{0}$ = 71 km s$^{-1}$ Mpc$^{-1}$. | 16 | 7 | 1607.07046 |
|
1607 | 1607.07877_arXiv.txt | In chameleon gravity, there exists a light scalar field that couples to the trace of the stress-energy tensor in such a way that its mass depends on the ambient matter density, and the field is screened in local, high-density environments. Recently it was shown that, for the runaway potentials commonly considered in chameleon theories, the field's coupling to matter and the hierarchy of scales between Standard Model particles and the energy scale of such potentials result in catastrophic effects in the early Universe when these particles become nonrelativistic. Perturbations with trans-Planckian energies are excited, and the theory suffers a breakdown in calculability at the relatively low temperatures of Big Bang Nucleosynthesis. We consider a chameleon field in a quartic potential and show that the scale-free nature of this potential allows the chameleon to avoid many of the problems encountered by runaway potentials. Following inflation, the chameleon field oscillates around the minimum of its effective potential, and rapid changes in its effective mass excite perturbations via quantum particle production. The quartic model, however, only generates high-energy perturbations at comparably high temperatures and is able remain a well-behaved effective field theory at nucleosynthesis. | \label{sec:Intro} Many explanations for the current accelerated expansion of the Universe posit the existence of a new light scalar field. These scalar fields are usually coupled to matter and so can mediate long-range forces, often of gravitational strength. Such scalars are not only cosmologically motivated, but also pervasive in high-energy physics and string theory. However, stringent experimental bounds imply tight constraints on any new fifth forces mediated by scalar fields. These constraints require the scalar's coupling to matter to be tuned to unnaturally small values in order to avoid detection. Another approach is to employ a screening mechanism, which suppresses effects of the field locally, allowing consistency with successful tests of general relativity. One of the few known screening mechanisms capable of reconciling the predictions of scalar-tensor gravitational theories and experimental constraints is the chameleon mechanism \cite{KhouryShort, KhouryLong}. In chameleon gravity theories, the scalar field's potential function and its coupling to the stress-energy tensor combine into an effective potential whose minimum is dependent on the matter density of its environment. Consequently, the effective mass of the chameleon field is also dependent on the environment, increasing enough in regions of high density to suppress the field's ability to mediate a long-range force. Because of this ability to hide within its environment, the chameleon can couple to matter with gravitational strength and still evade experimental detection in laboratory and Solar System tests of gravity. The vast majority of cosmological investigations of chameleon gravity have considered potentials of the runaway form, such as the exponential \mbox{$V(\phi)=M^4\exp[(M / \phi)^n]$} and power-law $V(\phi)=M^{4+n}\phi^{-n}$ potentials. In order to evade Solar System tests of gravity, $M$ has to be set to a value of $\sim\!\!10^{-3}$ eV, which is the energy scale of dark energy \cite{KhouryLong}. This coincident energy scale gave the chameleon a lot of attention early on as a possible explanation for cosmic acceleration. However, it was shown in Ref. \cite{NoGo} that the chameleon field cannot account for the accelerated expansion of the Universe without including a constant term in its potential. Nevertheless, light scalar fields arise in many theories that consider physics beyond the Standard Model (SM), and the chameleon mechanism remains one of the most-studied approaches to screening the unwanted forces mediated by these fields. Many laboratory experiments have been conducted to search for forces mediated by chameleon fields. Experiments that use atom \cite{atom1} and neutron \cite{neutron1,neutron2} interferometry and those that use $\mu$m-sized test masses \cite{microshperes} have already placed constraints on chameleon theories. Additional experiments have been proposed: one aims to measure the interactions between parallel plates to search for new forces \cite{pp1} and another suggests using atom interferometry between parallel plates of different densities to detect density-dependent chameleon forces \cite{pp3}. Laboratory searches for chameleon particles converted from photons in the presence of a magnetic field via the Primakov effect have placed constraints on the chameleon-photon coupling \cite{ADMX,Chase}. The CERN Axion Solar Telescope searched for chameleons created in the Sun by this effect \cite{sun2} and is currently conducting more sensitive searches \cite{sun4} to detect solar chameleons via their radiation pressure \cite{sun1}. Chameleon theories have also been constrained by their effects on the pulsation rate of Cepheids \cite{cepheids} and comparisons of x-ray and weak-lensing profiles of galaxy clusters \cite{cluster1,cluster2}. There have also been efforts to constrain the parameters of chameleon models by their effects on the cosmic microwave background \cite{CMBbeta,CMBDM}, though these analyses focus specifically on potentials of the power-law form. Given the tremendous experimental effort under way to detect or constrain chameleons, it is troubling that the most widely studied chameleon models have been shown to suffer a breakdown in calculability in the early Universe due to the discrepancy between the chameleon mass scale and that of the SM particles \cite{AdrienneShort, AdrienneLong}. We aim to identify a chameleon potential that can avoid the computational breakdown suffered by runaway models. We analyze a class of potential not often considered in chameleon theories: the quartic potential, $V(\phi)=\kappa\phi^4/4!$. Prevalent in high-energy theories, the quartic potential is also viable as a chameleon model because the self-interaction of this potential is sufficient to ensure that the field will be adequately screened in high-density environments \cite{Phi4}. The scale-free property of the quartic model is potentially beneficial as it can avoid the hierarchy of energy scales that arises due to the low-energy scale of the runaway potentials, and we investigate whether it is able to remain well-behaved in the early Universe.\footnote{Another proposed way to avoid the detrimental effect of the kicks is to include DBI-inspired corrections to the chameleon's Lagrangian that weaken the chameleon's coupling to matter at high energies. This modification effectively introduces a second screening mechanism analogous to a Vainshtein screening in which derivative interactions weaken the effect of the kicks \cite{SwiftKick}.} In runaway chameleon models, the field rolls to some value far from the minimum of its effective potential after inflation and remains stuck there during the radiation-dominated era due to Hubble friction. The chameleon's coupling to the trace of the stress-energy tensor makes it sensitive to the energy density, $\rho$, and pressure, $P$, of the radiation bath through the quantity $\Sigma\equiv(\rho-3P)/\rho$. While the Universe is radiation dominated, $\Sigma$ is nearly zero and the chameleon is light enough that Hubble friction is able to prevent it from rolling toward its potential minimum. However, as the temperature of the radiation bath cools, particle species in thermal equilibrium become nonrelativistic and $\Sigma$ momentarily becomes nonzero. The chameleon then gains mass, is able the overcome Hubble friction, and is seemingly ``kicked'' toward the minimum of its effective potential \cite{Brax}. Originally, the kicks were seen as an auspicious way to bring runaway chameleons to their potential minimum prior to Big Bang Nucleosynthesis (BBN).\footnote{A consequence of the chameleon's coupling to matter is that any variation in the chameleon field can be recast as a variation in particle masses in the Jordan frame. As we know particle masses differed very little between BBN and the present day, this constrains the chameleon to be at or near the minimum of its potential prior to the onset of BBN \cite{Brax}.} However, they impart such a high velocity to the field that the chameleon rebounds off the other side of its effective potential back to field values further from the potential minimum than where it was stuck when the kick began \cite{EM}. However, Ref. \cite{EM} also showed that the inclusion of a coupling between the chameleon and the electromagnetic field offers a solution. The chameleon's coupling to a primordial magnetic field allows the chameleon to overcome Hubble friction and begin oscillating about its potential minimum prior to the kicks. For a sufficiently rapidly oscillating field, the kicks then have little effect on the chameleon's evolution. These kicks further jeopardized chameleon theories by throwing into question their validity as a classical field theory \cite{AdrienneShort, AdrienneLong}. The effective potential in runaway models is minimized when $\phi\sim\!M$, and at field values $\phi\lesssim\!M$, the extremely steep slope of the bare potential leads to rapid changes in the chameleon's effective mass for small field displacements. Thus, the GeV-scale velocity with which the chameleon approaches its meV-scale minimum after the kicks causes nonadiabatic changes in the mass that excite extremely energetic fluctuations and lead to the quantum production of particles \cite{AdrienneShort, AdrienneLong}. Quantum corrections due to particle production then invalidate the classical treatment of the chameleon field and the particles' trans-Planckian energies cast doubt on the chameleon's viability as an Effective Field Theory (EFT) at the energy scale of BBN. We will show that the quartic potential is able to avoid these problems due to its scale-free nature. In the early Universe, a chameleon field in a quartic potential oscillates rapidly with a large amplitude far beyond the minimum of its effective potential. In the classical treatment, the chameleon would continue this behavior until the end of radiation domination and still be oscillating far outside its minimum at the onset of BBN. A quantum treatment of the chameleon's motion shows that these oscillations will create particles, albeit with much less energy than those created during the rebounds off runaway potentials. The same quantum effects that were catastrophic to previous chameleon models will cause the field to lose energy and bring the quartic chameleon to its potential minimum prior to the onset of the kicks. Consequently, for the quartic chameleon, these kicks do not have as significant an influence on the field's evolution as in models with runaway potentials. Depending on the value of $\kappa$, the rate at which the field loses energy can vary significantly. For large values of $\kappa$, the field can lose all of its initial energy to particle production within the first oscillation and fall to its minimum. However when $\kappa$ is closer to unity only a small percentage of the energy is lost during each oscillation, but the total effect accumulates over many oscillations to introduce a decay factor to the amplitude that still allows the field to reach its potential minimum before the kicks. We begin with a brief review of chameleon gravity and then analyze the evolution of a classical chameleon field in a quartic potential in Section II. Then, in Section III, we consider the effects of quantum particle production on the field and investigate how the energy lost to this process is affected by the choice of $\kappa$. Section IV explores how the kicks affect the field, and we follow up with concluding remarks in Section V. Throughout this paper we will use $\Mpl=(8 \pi G)^{-1/2}$ and $c=\hbar =1$. | \label{sec:end} Since the chameleon model was first proposed as an alternative to dark energy \cite{KhouryLong, KhouryShort}, its cosmological impacts have been studied extensively. While it has been shown that chameleon theories cannot account for the expansion of the Universe without the addition of a constant term to its potential \cite{NoGo}, the field's sensitive dependence on its environment gives it remarkable properties that are of great interest. However, for most chameleon models, the same matter coupling that gives it its unique phenomenology leads these theories into trouble in the early Universe. The meV mass scale of runaway potentials is at odds with the GeV mass scale of SM particles, which accelerate the chameleon field to very high velocities when they become nonrelativistic. The hierarchy between these two energy scales leads to the quantum production of particles that radically alters the field's evolution. Without very weak couplings or highly tuned initial conditions these chameleon models cannot be trusted as effective field theories at the time of BBN \cite{AdrienneShort, AdrienneLong}. In this paper, we have considered the quartic chameleon potential, which is not often studied in theories of chameleon gravity. A significant feature of this model is the fact that there is no mass scale in the potential: the chameleon's self-interaction is enough to ensure adequate screening. We have shown that this scale-free property of the potential allows the quartic chameleon to avoid the catastrophic effects of the small energy scales within runaway models. After inflation, the quartic chameleon oscillates in its potential well. The amplitude of its oscillations are damped due to Hubble friction. In the classical treatment, the minimum of the field's effective potential decreases faster than the oscillation amplitude throughout radiation domination. Consequently, the field cannot reach its potential minimum before BBN, though the oscillation amplitude is always sufficiently small that the variation of the field from this minimum does not imply an unacceptable variation from known particle masses. \begin{figure} \centering\includegraphics[width=3.4in]{m2min2Plot.eps} \caption{The numerical evaluation of the ratio in Eq. (\ref{eq:mmindot}), which is significantly greater than 1 throughout the kicks. It becomes less than 1 when $T\simeq3.9\times10^{-4}$GeV.} \label{m2Sig2} \end{figure} The rapid oscillations of the chameleon field cause changes in its effective mass that excite perturbations and lead to particle production. The effects of quantum particle production ensure that the field does reach its potential minimum while the Universe is radiation dominated. For large values ($\gtrsim 10$) of the self-interaction constant $\kappa$, the fractional loss of energy to these particles can be large, in which case the field loses all its energy in the course of a single oscillation. For smaller $\kappa$, the energy lost to particles constitutes only a small fraction of the field's energy. This much slower energy loss accumulates over multiple oscillations and introduces an additional decay term to the oscillation amplitude which allows the field to catch its minimum after many oscillations. At this point, the field will adiabatically track its potential minimum. It will track this minimum until the very tail end of the kicks, when the Boltzmann suppression of $\Sigma$ decreases the value of $\phimin$ faster than the field can follow. The value of the field at this point is sufficiently small that any deviation of particle masses implied by the deviation of the field from its potential minimum are entirely negligible. The energy of the modes that are excited in the quartic model are on order of the temperature: highly energetic modes are only excited at high temperatures. This is an important contrast to runaway models, which experience extremely energetic fluctuations at relatively low temperatures and can no longer be treated as EFTs during BBN. While quantum corrections lead to extremely energetic fluctuations in the field and a breakdown in calculability for runaway models, quantum corrections to the quartic potential are comparatively small and are, in fact, necessary to ensure the field can reach its minimum. Once it reaches the minimum of its effective potential the chameleon can then adiabatically track this minimum even throughout the kicks. Thus, the quartic chameleon's scale-free nature means is not susceptible to problems arising from a hierarchy of scales and can remain a well-behaved effective field theory throughout the evolution of the early Universe. | 16 | 7 | 1607.07877 |
1607 | 1607.02513_arXiv.txt | We present 1.3~mm observations of the Sun-like star $\tau$ Ceti with the Atacama Large Millimeter/submillimeter Array (ALMA) that probe angular scales of $\sim1\arcsec$ (4 AU). This first interferometric image of the $\tau$ Ceti system, which hosts both a debris disk and possible multiplanet system, shows emission from a nearly face-on belt of cold dust with a position angle of $90\degr$ surrounding an unresolved central source at the stellar position. To characterize this emission structure, we fit parametric models to the millimeter visibilities. The resulting best-fit model yields an inner belt edge of $6.2^{+9.8}_{-4.6}$~AU, consistent with inferences from lower resolution, far-infrared \emph{Herschel} observations. While the limited data at sufficiently short baselines preclude us from placing stronger constraints on the belt properties and its relation to the proposed five planet system, the observations do provide a strong lower limit on the fractional width of the belt, $\Delta R/R > 0.75$ with $99\%$ confidence. This fractional width is more similar to broad disks such as HD 107146 than narrow belts such as the Kuiper Belt and Fomalhaut. The unresolved central source has a higher flux density than the predicted flux of the stellar photosphere at 1.3~mm. Given previous measurements of an excess by a factor of $\sim2$ at 8.7~mm, this emission is likely due to a hot stellar chromosphere. | \label{sec:intro} The 5.8 Gyr-old \citep{mam08} main-sequence G8.5V star $\tau$ Ceti is the second closest \cite[3.65~pc,][]{vanL07} Solar-type star reported to harbor both a tentative planetary system and a debris disk \cite[after $\epsilon$ Eridani, e.g.][]{gre98,hat00}. The $\tau$ Ceti debris disk was first identified as an infrared excess by IRAS \citep{aum85} and confirmed by ISO \citep{hab01}. \cite{gre04} marginally resolved 850 $\mu$m emission from the system with the James Clerk Maxwell Telescope (JCMT)/SCUBA, revealing a massive (1.2 $M_\oplus$) disk extending to 55 AU from the star. Recent \emph{Herschel} observations at 70, 160, and 250~$\mu$m resolve the disk well and are best fit by a broad dust belt with an inner edge between $1-10$ AU and an outer edge at $\sim55$~AU \citep{law14}. Due to its proximity and similarity to our Sun in age and spectral type, $\tau$ Ceti has been the object of numerous searches for planets using the radial velocity technique \cite[e.g.][]{pepe11}, most of which have proved unsuccessful. Using extensive modeling and Bayesian analysis of radial velocity data from the High Accuracy Radial Velocity Planet Searcher (HARPS) spectrograph \citep{may03,pepe11}, the Anglo-Australian Planet Search (AAPS) on the Anglo Australian Telescope (AAT), and the High Resolution Echelle Spectrograph (HIRES) on the Keck telescope \citep{vogt94}, \cite{tuo13} report evidence for a tightly-packed five planet system. This purported planetary system consists of five super-Earths with masses of $4.0-13.2$~$M_\oplus$ (for orbits co-planar with the disk), semi-major axes ranging over $0.105-1.35$ AU, and small eccentricities, $e\sim0-0.2$. The veracity of these planet candidates, however, remains controversial. \cite{tuo13} acknowledge that the detected signals could also result from a combination of instrumental bias and stellar activity, although no further evidence is given to support these alternative interpretations. Also of note is the sub-Solar metallicity of $\tau$ Ceti, [Fe/H] $= -0.55\pm0.05$ dex \citep{pav12}, which makes it an interesting target for exoplanet searches due to the observed higher frequency of low-mass planets around low-metallicity stars \citep{jen13}. We present interferometric observations of the $\tau$ Ceti system at 1.3~mm using the Atacama Large Millimeter/submillimeter Array (ALMA). Millimeter imaging of this debris disk opens a unique window on the location and morphology of the underlying population of dust-producing planetesimals orbiting the star. While these large, kilometer-sized bodies cannot be detected directly, millimeter observations probe emission from the large dust grains produced through collisions that are not rapidly redistributed by stellar radiation and winds \citep{wya06}. These new ALMA observations provide limits on the disk location and width, which bear on the proposed planetary system within the disk. In Section~\ref{sec:obs}, we present the ALMA observations of the $\tau$ Ceti system. In Section~\ref{sec:results}, we describe the analysis technique and disk model results. In Section~\ref{sec:disc}, we discuss the significance of the best-fit model parameters for the dust belt inner edge, width, proposed planetary system, and the origin of a bright, unresolved central emission source. | \label{sec:conclusions} We observed the $\tau$ Ceti debris disk with ALMA at 1.3~mm with baselines that probe $1\arcsec$ (4~AU) scales. These are the first observations of this nearby system with a millimeter interferometer and reveal somewhat patchy emission from a dust disk surrounding an unresolved central emission peak. In order to characterize these two emission components, we fit simple parametric models directly to the visibility data within an MCMC framework. Our best-fit model yields an inner belt edge of $6.2^{+9.8}_{-4.6}$ AU, consistent with the analysis of previous far-infrared \emph{Herschel} observations. Given the relatively low sensitivity at short baselines in the ALMA observations, we are unable to place a tighter constraint on the inner edge and its position relative to the proposed five planet system. These data, however, provide a strong lower limit on the fractional width of the belt, $\Delta R/R > 0.75$ with $99\%$ confidence. This result implies that the $\tau$ Ceti debris disk is broad, much wider than the classical Kuiper Belt in our Solar System and more comparable to the HD 107146 debris disk \citep{ric15}. The bright central peak at the stellar position has a flux density of $F_\text{1.3mm}=0.69^{+0.02}_{-0.05}$~mJy, somewhat higher than the predicted flux of the stellar photosphere at 1.3~mm. At longer centimeter wavelengths, this excess is more significant, increasing to $\sim2\times$ the photospheric prediction in VLA observations at 8.7~mm \citep{vill14}. The spectral index between these two measurements is $1.74\pm0.15$, shallower than the expectation for an optically thick photosphere. Given the high brightness temperatures at both 1.3 and 8.7~mm, this excess emission is likely due to a hot stellar chromosphere. Similar spectra have been observed for other nearby Sun-like stars, e.g. $\alpha$ Cen A/B and $\epsilon$ Eridani. These first ALMA observations of the $\tau$ Ceti system allow us to probe the structure of the debris disk with higher resolution than previous work. However, higher sensitivity observations at shorter baselines are still needed to constrain the location of the inner edge of the dust belt more precisely. If the disk extends in towards the star, within the orbit of the outermost proposed planet, this provides strong evidence against the posited five planet system. However, if the disk inner edge is located well outside the proposed planetary system, an additional massive planet on a wide orbit may be required to clear out the central hole in the belt. Additional observations with the ACA could provide the necessary sensitivity to determine the position of the inner disk edge and its implications for an interior planetary system. | 16 | 7 | 1607.02513 |
1607 | 1607.05934_arXiv.txt | Magnetars are neutron stars endowed with surface magnetic fields of the order of $10^{14}-10^{15}$~G, and with presumably much stronger fields in their interior. As a result of Landau quantization of electron motion, the neutron-drip transition in the crust of a magnetar is shifted to either higher or lower densities depending on the magnetic field strength. The impact of nuclear uncertainties is explored considering the recent series of Brussels-Montreal microscopic nuclear mass models. All these models are based on the Hartree-Fock-Bogoliubov method with generalized Skyrme functionals. They differ in their predictions for the symmetry energy coefficient at saturation, and for the stiffness of the neutron-matter equation of state. For comparison, we have also considered the very accurate but more phenomenological model of Duflo and Zuker. Although the equilibrium composition of the crust of a magnetar and the onset of neutron emission are found to be model dependent, the quantum oscillations of the threshold density are essentially universal. | At the end point of stellar evolution, neutron stars are not only the most compact stars in the universe, but also the strongest magnets~\cite{hae07}. In particular, magnetic fields of the order of $10^{14}-10^{15}$~G have been measured at the surface of soft gamma-ray repeaters and anomalous x-ray pulsars~\cite{mcgill14}, thus dubbed \emph{magnetars}. On the other hand, numerical simulations have shown that the internal magnetic field could be even stronger, up to about $10^{18}$~G~\cite{deba15}. The outermost layer of a neutron star is thought to consist of a solid crust, whose atoms are fully ionized by the gravitational pressure (for a review, see Ref.~\cite{lrr}). With increasing depth, nuclei become progressively more neutron rich by capturing electrons until at some point, neutrons start to drip out of nuclei. The presence of a neutron liquid in the crust of a magnetar is expected to leave its imprint on various observed astrophysical phenomena like sudden spin-ups~\cite{dib08,alpar14} and spin-downs~\cite{archi13,mus14,duncan13,kantor14} (generally referred to as ``glitches'' and ``anti-glitches'' respectively), quasiperiodic oscillations detected in the giant flares from soft gamma-ray repeaters~\cite{and09,chamel2013,pass14} and cooling~\cite{agui09}. We have recently studied the effects of a strong magnetic field on the neutron-drip transition in the crust of a magnetar~\cite{chamel2012,chamel2015b}. We have shown that the neutron-drip density and pressure increase almost linearly with the magnetic field strength in the strongly quantizing regime. In the weakly quantizing regime, the variations of the neutron-drip density with magnetic field strength exhibit typical quantum oscillations. The neutron-drip transition in a magnetar, as compared to unmagnetized neutron stars, can thus be shifted to either higher or lower densities depending on the magnetic field strength. In this paper, we explore the role of nuclear uncertainties on the neutron-drip transition considering different nuclear mass models. | We have pursued our study of the effects of Landau quantization on the onset of neutron emission by nuclei in the crust of a magnetar~\cite{chamel2015b}. In particular, we have explored the impact of nuclear uncertainties considering the recent series of Brussels-Montreal Hartree-Fock-Bogoliubov nuclear mass models, from HFB-22 to HFB-27$^*$~\cite{gcp13a,gcp13b}. These models yield equally good fits to essentially all experimental masses, with a root-mean-square deviation of about $0.5-0.6$ MeV. However, they lead to different predictions of nuclear-matter properties thus reflecting the current lack of knowledge of the symmetry energy, and the stiffness of the neutron-matter equation of state. For comparison, we also considered the more accurate but also more phenomenological model of Duflo and Zuker~\cite{dz95}. Although the equilibrium nucleus at the bottom the outer crust, and the onset of neutron emission are found to be model dependent, the oscillations of the threshold density as a function of the magnetic field strength are almost universal. The role of the magnetic field on the neutron-drip density in magnetar crusts could thus be a priori inferred from the value of the neutron-drip density in unmagnetized neutron star crusts. On the other hand, the change of nuclear masses due to the strong magnetic field~\cite{bas2015}, which we have not taken into account, may undermine this universal behaviour. \ack This work was financially supported by Fonds de la Recherche Scientifique - FNRS (Belgium), Wallonie-Bruxelles-International (Belgium), the Bulgarian Academy of Sciences, the Bulgarian National Science Fund under contract No. DFNI-T02/19, and the European Cooperation in Science and Technology (COST) Action MP1304 ``NewCompStar''. | 16 | 7 | 1607.05934 |
1607 | 1607.08006_arXiv.txt | The IceCube Collaboration has previously discovered a high-energy astrophysical neutrino flux using neutrino events with interaction vertices contained within the instrumented volume of the IceCube detector. We present a complementary measurement using charged current muon neutrino events where the interaction vertex can be outside this volume. As a consequence of the large muon range the effective area is significantly larger but the field of view is restricted to the Northern Hemisphere. IceCube data from 2009 through 2015 have been analyzed using a likelihood approach based on the reconstructed muon energy and zenith angle. At the highest neutrino energies between $194\,\mathrm{TeV}$ and $7.8\,\mathrm{PeV}$ a significant astrophysical contribution is observed, excluding a purely atmospheric origin of these events at $5.6\,\sigma$ significance. The data are well described by an isotropic, unbroken power law flux with a normalization at $100\,\mathrm{TeV}$ neutrino energy of $\left(0.90^{+0.30}_{-0.27}\right) \times 10^{-18}\,\mathrm{GeV^{-1}\,cm^{-2}\,s^{-1}\,sr^{-1}}$ and a hard spectral index of $\gamma = 2.13 \pm 0.13$. The observed spectrum is harder in comparison to previous IceCube analyses with lower energy thresholds which may indicate a break in the astrophysical neutrino spectrum of unknown origin. The highest energy event observed has a reconstructed muon energy of $(4.5\pm1.2)\,\unit{PeV}$ which implies a probability of less than $0.005\%$ for this event to be of atmospheric origin. Analyzing the arrival directions of all events with reconstructed muon energies above $200\,\mathrm{TeV}$ no correlation with known $\gamma$-ray sources was found. Using the high statistics of atmospheric neutrinos we report the currently best constraints on a prompt atmospheric muon neutrino flux originating from charmed meson decays which is below $1.06$ in units of the flux normalization of the model in \cite{Prompt:ERS}. | \label{sec:introduction} The detection of high-energy cosmic neutrinos as cosmic messengers has been an important goal of astroparticle physics. Being stable, electrically neutral particles, high-energy neutrinos are able to propagate almost undisturbed through the universe from their production sites to Earth keeping their directional and energy information. Hence, they constitute excellent cosmic messenger particles, particularly at the highest energies. They arise from weak decays of hadrons, mostly pions and kaons, which are expected to be produced by hadronic interactions of cosmic-rays in the surrounding matter of the cosmic-ray accelerator. Their observation will help to elucidate the unknown sources of high-energy cosmic-rays \citep{Gaisser:Halzen:Stanev,Learned:Mannheim,Becker:2007sv}. Already in the 1960s the observation of high-energy neutrinos was discussed by \cite{Greisen:1960, Markov:1960vja,Reines:1960}, shortly after the discovery of the neutrino by \cite{Reines:1956}. The proposed method was the detection of up-going muons as a signature of a charged-current (CC) muon neutrino interaction below the detector. Soon it was realized that the expected astrophysical fluxes are small and cubic-kilometer sized detectors would be needed to accomplish the goal, see e.g. \cite{Roberts:1992re}. The construction of large Cherenkov detectors by instrumenting optically transparent natural media, i.e. deep oceans, lakes and glaciers with photo-sensors \citep{Belolaptikov:1997ry,Andres:1999hm,antares:2011nsa} proved to be a key concept. The largest instrument to date is the IceCube Neutrino Observatory at the geographic South Pole, \cite{IceCube:FirstYear}. Main backgrounds to the search for astrophysical neutrinos are high-energy atmospheric neutrinos and muons produced by cosmic-ray interactions in the Earth's atmosphere. In 2013, a diffuse all-flavor flux of high-energy astrophysical neutrinos was discovered \citep{Aartsen:2013jdh,IceCube:HESE3Years}. The analysis selected events due to high-energy neutrinos which interact within the detector by using its outer layers as a veto. This strategy enables a full-sky sensitivity for all neutrino flavors. The veto not only rejects atmospheric muons entering the detector from the outside extremely efficiently, but also atmospheric neutrinos from above the detector which are produced together with muons. \begin{figure*} \centering \includegraphics[width=1.\textwidth]{signature_conventional_prompt_powerlaw.pdf} \caption{Distribution of the expected neutrino energy (left) and zenith angle (right) for the data selection of this analysis. Shown are the distributions of conventional atmospheric neutrinos \citep{Conv:Honda2007}, prompt atmospheric neutrinos \citep{Prompt:ERS} where both are corrected for the cosmic-ray spectrum in \cite{Gaisser:2012zz} and a benchmark astrophysical signal $10^{-18}\,\mathrm{GeV^{-1}\,cm^{-2}\,sr^{-1}\,s^{-1}} (E_\nu/100\,\mathrm{TeV})^{-2}$. \label{fig:astro_signature}} \end{figure*} In this analysis we focus on up-going muons which arise from charged-current interactions of muon neutrinos both inside and outside the detector. By allowing neutrinos to interact outside the instrumented volume a larger effective area is achieved. However, at the same time it is necessary to restrict the analysis to the Northern hemisphere where the Earth filters atmospheric muons efficiently. Furthermore, the analysis is mainly sensitive to a muon neutrino flux because of the large muon range. Nevertheless this strategy will impose further constraints on possible models \citep{PhysRevD.87.063011, PhysRevD.88.047301, JCAP.2013.01.028, PhysRevLett.111.041103, JCAP.2012.06.030, PhysRevD.88.043009, PhysRevD.88.081302, GonzalezGarcia201439, PhysRevD.89.083003, Tamborra:2014:Starforming, MuraseBlazers:2014, Bechtol:2015, Senno:2015} that have been proposed to explain the observed astrophysical neutrino flux. This analysis is based on a high-purity and high-statistics selection of about $350 000$ well-reconstructed up-going muon events from six years of IceCube operation, improving the statistics compared to previous analyses \citep{IceCube:IC59NuMuDiffuse,IceCube:IC79NuMuDiffuse} by almost an order of magnitude. Even when individual astrophysical neutrino sources cannot be identified because they are too weak, their cumulative flux can be measured as a diffuse flux. The signature of an astrophysical neutrino signal with respect to the background of atmospheric neutrinos is illustrated in Fig. \ref{fig:astro_signature}. Astrophysical neutrinos from cosmic accelerators are generically expected to have a hard energy spectrum as originally predicted by Fermi: $dN_{\nu}/dE \simeq \phi_0 \cdot E^{-2}$. However, the spectral index depends in detail on the source properties and the acceleration mechanism \citep{Bell201356, Kashti:2005qa, Klein:2012ug}. Recent IceCube analyses \citep{Aartsen:2015zva:hese, IceCube:MESE2Years, Aartsen:2015knd, Aartsen:2015zva:casc} yielded a softer spectrum with a spectral index between 2.5 and 2.7. The energy spectrum of the atmospheric neutrino background is about one power steeper than the primary cosmic-ray spectrum ($dN_{CR}/dE \propto E^{-2.7..3.1} $), with the exception of prompt neutrinos from heavy meson decays, which follow the primary spectrum more closely. The astrophysical signal appears as an excess above energies of about $100\,\mathrm{TeV}$. As shown the zenith distribution differs for signal and backgrounds which themselves depend on the energy. At the highest energies the Earth becomes increasingly opaque to neutrinos and the signal is dominated by events near the horizon. The identification of an astrophysical signal is based on a two-dimensional likelihood fit in zenith and energy. It follows the methods of the previous analyses \citep{IceCube:IC59NuMuDiffuse,IceCube:IC79NuMuDiffuse} which are improved with respect to the treatment of systematic uncertainties. The data selection is described in Sec. \ref{sec:data_sample}. The method is described in Sec. \ref{sec:analysis_method}. The results of the analysis with respect to the astrophysical signal are presented in Sec. \ref{sec:results}, where we discuss the fit results, tests of alternative hypotheses and investigations on the most energetic event \citep{schoenen2015detection}. In Sec. \ref{sec:aniso} we present investigations on the directions of recorded events and the attempt to correlate these directions with astrophysical objects. In Sec. \ref{sec:prompt} we discuss implications of this analysis for the expected flux of high-energy prompt atmospheric neutrinos from the decay of charmed mesons and obtain the currently most constraining exclusion limit. | In this paper we have presented the result of analyzing 6 years of up-going muon data measured with the IceCube neutrino telescope. We measure an astrophysical flux of $ \Phi_{\nu+\overline{\nu}} = \left(0.90^{+0.30}_{-0.27}\right) 10^{-18}\,\mathrm{GeV^{-1}\,cm^{-2}\,sr^{-1}\,s^{-1}} \cdot (E_\nu/100\,\mathrm{TeV})^{-(2.13 \pm 0.13)} $ with statistical significance of $5.6$ standard deviations with respect to only being of atmospheric origin. With this result we have further established the observation of an astrophysical neutrino signal \citep{Aartsen:2013jdh,IceCube:HESE3Years, IceCube:IC79NuMuDiffuse} in a second, largely independent detection channel. The detection channel used here is of great interest because of the good directional reconstruction of detected muons and a large signal efficiency with an estimated number of about $500$ astrophysical neutrinos included in this data sample. The data include an exceptionally high-energy muon with $(2.6 \pm 0.3)$\,PeV deposited energy, which is the highest energy lepton that has been reported to date. A parametric unfolding of neutrino energies shows that the spectrum extends to about $10$\,PeV in neutrino energy with no significant spectral break or cut-off. The measured hard spectral index of $\gamma = 2.13 \pm 0.13$ is in tension with complementary measurements of IceCube, which have a lower energy threshold by about one order of magnitude and are predominantly sensitive to the Southern hemisphere. However, the consistency of the observed fluxes at high energies may be interpreted as indication of a spectral break or additional astrophysical component at lower energy to which this analysis is not sensitive. For the highest-energy events no correlation with known high-energy gamma-ray sources or other astrophysical objects could be identified. By splitting the data in right ascension, we find no significant correlation with the orientation of the galactic plane and conclude that the dominant fraction of the flux is largely all-sky and isotropic. The present analysis is also sensitive to a flux of prompt neutrinos which are expected from the decay of heavy mesons in the atmosphere. We find no indications for such a signal. However, because the prompt flux is subdominant to the astrophysical and conventional atmospheric neutrino flux, the exclusion depends on the assumed astrophysical model parameters. Variations of the astrophysical flux uncertainties lead to a conservative exclusion limit of approximately at the level of the mean expected flux normalization from \cite{Prompt:ERS}. For the first time, it is possible to constrain such a flux in this range of theoretical predictions. However, recent perturbative QCD calculations from \cite{Prompt:GMS}, \cite{Prompt:BERSS} and \citep{Prompt:GRSST} predict lower prompt neutrino fluxes which are not yet constrained by the upper limit. | 16 | 7 | 1607.08006 |
1607 | 1607.03778_arXiv.txt | We report on the search for gamma-ray emission from 20 magnetars using 6 years of \fermi\, Large Area Telescope (LAT) observations. No significant evidence for gamma-ray emission from any of the currently-known magnetars is found. We derived the most stringent upper limits to date on the 0.1--10\, GeV emission of Galactic magnetars, which are estimated between $\sim10^{-12}-10^{-11}$\ergscm2 . Gamma-ray pulsations were searched for the four magnetars having reliable ephemerides over the observing period, but none were detected. On the other hand, we also studied the gamma-ray morphology and spectra of seven Supernova Remnants associated or adjacent to the magnetars. | \label{intro} The magnetars (comprising the Anomalous X-ray Pulsars and Soft Gamma Repeaters; AXPs and SGRs) are a small group of X-ray pulsars (about twenty objects), with spin periods between 0.3 and 12 s. Their bright X-ray emission ($L_{X}\sim10^{33}-10^{35}$\ergs ) is marginally explained by the {commonly accepted emission models for isolated pulsars} or by accretion from a companion star. Their inferred magnetic fields assuming the spin-down dipolar loss formula (B = 3.2 $\times$10$^{19} \sqrt{P\dot{P}}$ G, $P$ is the spin period and $\dot{P}$ is the first derivative of spin period) appear to be in general as high as $B\sim10^{14}-10^{15}$\,G. Due to these high magnetic fields, the emission of magnetars is thought to be powered by the decay and the instability of their strong fields (Duncan \& Thompson 1992; Thompson \& Duncan 1993, Thompson, Lyutikov, \& Kulkarni 2002). The powerful X-ray output is usually well modelled by thermal emission from the neutron star hot surface (about 0.2--0.6\,keV), reprocessed in a twisted magnetosphere through resonant cyclotron scattering, a process favored only under these extreme magnetic conditions. On top of their persistent X-ray emission, magnetars emit very peculiar flares and outbursts on several timescales (from fractions of seconds to years) releasing a large amount of energy ($10^{40}-10^{46}$ erg). These flares are probably caused by {large-scale} rearrangements of the surface/magnetospheric field, either accompanied or triggered by displacements of the neutron-star crust (somehow analogous to quakes on the stellar surface). See Mereghetti (2008), Rea \& Esposito (2011), and Olausen \& Kaspi (2014) for recent reviews. Thanks to \emph{INTEGRAL}, \emph{Suzaku} and \emph{NuSTAR} we now have good spectra for magnetars in the hard X-ray energy range (Kuiper et al. 2004; Enoto et al. 2010; An et al. 2015). The gamma-ray emission (0.1--300\,GeV band) from 13 magnetars has been searched using the first 17 months of {\em Fermi}-LAT data (Abdo et al. 2010a). This search resulted in upper limits and no pulsations were found either. Taking advantage of about 5 more years of data, in this paper we re-analyze the {\em Fermi}--LAT observations for the previous objects, as well as for the 7 magnetars that were discovered in the meantime. Furthermore, the better response of the instrument at lower energies brought about by Pass 8 (Atwood et al. 2013) allows us to produce a more stringent search. In Sec.\,\ref{obs} we report on the {\em Fermi}-LAT analysis procedure, we summarize our results in Sec.\,\ref{results} and \ref{timing}, and provide a discussion of our findings in Sec.\,\ref{discussion}. \begin{table*}{} \centering \scriptsize \label{table1} \caption{{\em Fermi}-LAT upper limits on magnetars as obtained from the likelihood analysis. Fluxes in the different energy ranges are in units of 10$^{-11}$ erg~cm$^{-2}$s$^{-1}$.} \begin{tabular}{lcccccc} \\ \hline\hline % \\ Source & TS&TS &0.1$-$1 GeV & 1$-$10 GeV & 0.1$-$10 GeV & {6-years internal list} \\ & &(Abdo et al. 2010a) & $\Gamma$ (1.5) & $\Gamma$ (3.5) & $\Gamma$ (2.5) &srcs within 3$^{\circ}$ \\ \\ \hline \\ \\ \underline{SGR\,0418$+$5729} & 0.0 & 2.3 & $<${0.18} & $<${0.05 } &$<${0.15} & 3 \\ \underline{4U\,0142$+$614} & 0.0 &3.6 & $<${0.43} & $<${0.07} &$<${0.29} & 3 \\ Swift\,J1822.3$-$1606 & 0.0 & & $<$0.88 & $<$0.13 &$<$0.43 & 19 \\ Swift\,J1834.9$-$0846$^\blacktriangledown$$^\star$ & 0.0 & & $<$1.06 & $<$ 0.99& $<$ 0.45 & 19 \\ \underline{1E\,1048.1$-$5937} & 0.0 & 0.0 & $<${0.51} & $<${0.53} &$<${1.10} & 35 \\ XTE\,J1810$-$197 & 0.0 & 13.1 & $<$1.00 & $<$0.89 &$<$3.65 & 17 \\ \underline{PSR\,J1622$-$4950} & {0.8} & & $<${1.98} & $<${0.63} &$<${1.00} & 26 \\ 1E\,1841$-$045 $^\star$ & 1.0 &7.5 & $<$1.05 & $<$0.78 & $<$ 2.02 & 23 \\ 3XMM\,J185246.6$+$003317$^\star$ & 2.5 & & $<$1.32 & $<$0.49 &$<$1.88 & 23 \\ \underline{1E\,2259$+$586}$^\star$ & {2.8} & 15.6 & $<${0.21}& $<${0.12}& $<${0.13} & 13 \\ SGR\,1806$-$20 & 2.9 & 2.8 & $<$2.38 & $<$0.50 &$<$0.67 & 15 \\ SGR\,1935$+$2154 & 3.5 & & $<$0.11 & $<$0.22 &$<$0.17 & 6 \\ \underline{SGR\,0501$+$4516}$^\star$ & {3.8} & 16.3 & $<${0.51} & $<${0.08} &$<${0.40} & 4 \\ \underline{1E\,1547.0$-$5408} & {4.2} & 36.2 & $<${1.40} & $<${0.90} &$<${2.84} & 14 \\ SGR\,1900$+$14$^\star$ & 4.5 &0.0 & $<$ 0.63 & $<$0.49 &$<$ 2.07 & 13 \\ \underline{CXOU\,J164710.2$-$455216} &{6.0}& & $<${0.96} & $<${0.45} &$<${0.89}& 27 \\ \underline{1RXS\,J170849.0$-$400910} & {6.6} & 32.1 & $<${2.14} & $<${0.96} &$<${2.51} & 15 \\ SGR\,1833$-$0832 $^\blacktriangledown$$^\star$ & 18.9 && $<$2.75 & $<$1.29 &$<$0.87 & 19 \\ \underline{SGR\,1627$-$41}$^\star$ & {23.8} & 36.0 & $<${2.87} & $<${1.75} & $<${3.34} & 29 \\ CXOU\,J171405.7$-$381031$\S$ & 51.0$\dag$ & & $<$0.27 & $<$0.53 &$<$0.80 & 11 \\ \tablecomments {Properties of the magnetars studied in this work, ordered by the measured TS values derived from the binned analysis. {The magnetars underlined have zenith angle cut at 90$\degree$ (see section \ref{obs} for detail).} The GeV upper limits are reported at {95\% confidence level} (see Section \ref{obs} for details). $\blacktriangledown$: SGR\,1833$-$0832 and Swift\,J1834.9$-$0846 are included in the same spatial model. $\S$: the upper limits of CXOU\,J171405.7$-$381031 are estimated with photon index fixed at 1.71. $\star$: magnetars with associated or close-by SNRs, the TS value are the residual after modelling the {extended} SNRs. $\dag$:{The high TS value may come from the possible gamma-ray emitting SNR CTB 37B. See section \ref{CXOU_subsection} for more detail.}} \end{tabular} \end{table*} | \label{discussion} In this paper, we searched for gamma-ray emission from magnetars using 6 years of {\em Fermi}-LAT data, which greatly extended the results of Abdo et al. (2010a). With this larger data set, and with a total of 20 magnetars searched, we can confirm that no gamma-ray emission is detected for any of the sources studied. Twelve out of twenty magnetars are singled out by a \textit{gtlike} analysis as having TS values below 25 (Table~1). CXOU J171405.7$-$381031 is positionally {coincident} with a point source. However, we could not further distinguish if the emission originates from the magnetar or from SNR CTB 37B. For all the studied magnetars we derived the deepest upper limits to date in the 0.1--300 GeV energy range. Gamma-ray pulsations were searched for 4 magnetars having reliable ephemerides covering most of the {\em Fermi}-LAT data span, but no periodicity was significantly observed. By applying an outer-gap model to magnetars, Cheng \& Zhang (2001) and Zhang \& Cheng (2002) predicted detectable gamma-ray emission from SGR 1900+14 and five others AXPs within one year of {\em Fermi}-LAT observations. These predicted flux levels are much beyond the currently imposed upper limits. Admittedly, these models are based on a number of assumptions and parameters (see Vigan\`o et al. (2015a) for a detailed discussion). However, it would be unrealistic to suppose that these parameters (e.g., the inclination angle) should be such to secure non-detections in all cases. This point was analyzed by Tong et al. (2010) when considering the reasons for the early non-detection of 4U 0142+61. For this particular case, even doubling the distance and using an outer gap location at 10 stellar radii or beyond the flux would have been larger than the upper limit found (see Tong et al. 2010 for details). Takata et al. 2013 proposed a possible scenario for GeV gamma-ray emission from magnetars, in which the magnetic energy released from crust cracking could be carried into outer magnetosphere by Alfv$\acute{e}$n waves. In this case, the predicted gamma-ray flux for 1E\,2259$+$586 is higher than the upper limits we derived. Vigan\`o et al. (2015b) proposed an outer-gap model for gamma-ray pulsars which follows the particle dynamics to consistently compute the emission of synchro-curvature radiation. They have applied this model to both the phase-averaged and phase-resolved spectra of the brightest pulsars. By fitting the model to all observed {\em Fermi}-LAT pulsar spectral data, they found a strong correlation between the accelerating electric field and the magnetic field at the light cylinder (B$_{LC}$, Vigan\`o et al. 2015c). If the correlations they found hold up to the magnetar regime (where at the light cylinder (LC), B$_{LC}$ $\sim$ 10$^{-1}$-10$^{-2}$ G, P $\sim$ 2-10 s), the accelerating electric field in the gap would be just too weak to provide particles energetic enough as to emit gamma rays. Their predictions that magnetars should not emit gamma rays via synchro-curvature radiation is consistent with the observational results of this paper. Analyzing the gamma-ray emission of the regions around magnetars, we have significantly detected 7 SNRs, four of which are believed to be associated to the spatially coincident magnetar. We have now studied them using Pass 8 data, updating their morphology and spectral parameters with respect to the 3FGL catalog (Acero et al. 2015) and Acero et al. 2016a. The gamma-ray morphologies and luminosities of these magnetars are in line with what is expected for remnants associated with normal radio pulsars at the same age (see also Martin et al. 2014 for an X-ray comparison). | 16 | 7 | 1607.03778 |
1607 | 1607.04644_arXiv.txt | {Automated arc detection methods are needed to scan the ongoing and next-generation wide-field imaging surveys, which are expected to contain thousands of strong lensing systems. Arc finders are also required for a quantitative comparison between predictions and observations of arc abundance. Several algorithms have been proposed to this end, but machine learning methods have remained as a relatively unexplored step in the arc finding process.} {In this work we introduce a new arc finder based on pattern recognition, which uses a set of morphological measurements that are derived from the Mediatrix filamentation method as entries to an artificial neural network (ANN). We show a full example of the application of the arc finder, first training and validating the ANN on simulated arcs and then applying the code on four Hubble Space Telescope (HST) images of strong lensing systems.} {The simulated arcs use simple prescriptions for the lens and the source, while mimicking HST observational conditions. We also consider a sample of objects from HST images with no arcs in the training of the ANN classification. We use the training and validation process to determine a suitable set of ANN configurations, including the combination of inputs from the Mediatrix method, so as to maximize the completeness while keeping the false positives low.} {In the simulations the method was able to achieve a completeness of about $90\%$ with respect to the arcs that are input into the ANN after a preselection. However, this completeness drops to $\sim 70\%$ on the HST images. The false detections are on the order of $3\%$ of the objects detected in these images.} {The combination of Mediatrix measurements with an % ANN is a promising tool for the pattern-recognition phase of arc finding. More realistic simulations and a larger set of real systems are needed for a better training and assessment of the efficiency of the method.} | } Strong lensing provides a useful tool to uncover the mass distribution in galaxies \citep[e.g.,][]{0004-637X-575-1-87,2002MNRAS.337L...6T,2006ApJ...649..599K} and galaxy clusters \citep[e.g.,][]{1989ApJ...337..621K,1998MNRAS.294..734A,2007MNRAS.376..180N,2010AdAst2010E...9Z,2010ApJ...715L.160C,2010ApJ...723.1678C}. Gravitational arcs have also been used to constrain the background cosmological model \citep[e.g.,][]{1998A&A...330....1B,1999A&A...341..653C,2002A&A...387..788G,2002MNRAS.337L...6T,2001PThPh.106..917Y,2004MPLA...19.1083M, 2005MNRAS.362.1301M,2010Sci...329..924J,Magana2015,0004-637X-806-2-185,2016A&A...587A..80C}. More recently, arcs and Einstein rings, in combination with kinematic information for the lenses, have been employed for testing modified gravity \citep[e.g.,][]{Schwab2010,Enander2013,2016arXiv160203385P}. Furthermore, strong lenses are also being exploited as cosmic telescopes, enabling spectroscopic and spatially resolved studies of high-redshift sources, such as dwarf galaxies at distant redshifts \citep{Marshal2007}, star-forming galaxies \citep{Stark2008}, quasar accretion disks \citep{2008ApJ...673...34P}, and faint Lyman-alpha blobs \citep{2015arXiv151205655C}. The many applications of gravitational arcs in astrophysics and cosmology have spurred the search for these objects in both space-based and ground-based observations. This includes searches in Hubble Space Telescope (HST) mosaics, such as the {\it Hubble Deep Field} \citep[HDF;][]{1996ApJ...467L..73H}, {\it HST Medium Deep Survey} \citep{1999AJ....117.2010R}, {\it Great Observatories Origins Deep Survey} \citep[GOODS;][]{2004ApJ...600L.155F}, {\it Extended Groth Strip} \citep[EGS;][]{Marshall2009robot}, {\it HST Cosmic Evolution survey} \citep[COSMOS;][]{2008ApJS..176...19F,2008MNRAS.389.1311J} and in targeted observations of galaxies \citep{2006ApJ...638..703B,2012ApJ...744...41B} and clusters \citep{2005Smith,2005ApJ...627...32S,2010MNRAS.406.1318H,2016ApJ...817...85X}. Investigations from the ground include follow-ups of clusters \citep{1999A&AS..136..117L, 2003ApJ...584..691Z,2008AJ....135..664H,2010A&A...513A...8K,2013MNRAS.432...73F} and galaxies \citep{2006MNRAS.369.1521W}, and searches in wide-field surveys, such as the {\it Red-Sequence Cluster Survey} \citep[RCS;][]{2003ApJ...593...48G,2012ApJ...744..156B}, {\it Sloan Digital Sky Survey} \citep[SDSS;][]{2007Estrada,2009MNRAS.392..104B, 2010ApJ...724L.137K, 2011RAA....11.1185W,2012ApJ...744..156B}, {\it Deep Lens Survey} \citep[DLS;][]{Kubo2008}, {\it Canada-France-Hawaii Telescope (CFHT) Legacy Survey} \citep[CFHTLS;][]{2007A&A...461..813C, 2012More,Maturi2014,RINGFINDER,SPACEWARPSII,ParaficzCFHTLS}, {\it CFHT Stripe 82 Survey} (CS82; Caminha, More et al., in prep.), and Dark Energy Survey\footnote{\texttt{http://www.darkenergysurvey.org}} \citep[DES;][]{DES01} Science Verification data \citep{DESSL1Nord}. As of now, the largest homogeneous samples of gravitational arcs have on the order of a hundred systems. These numbers will increase by one order of magnitude with the close completion of the Kilo Degree Survey\footnote{\texttt{http://kids.strw.leidenuniv.nl/}} \citep[KiDS;][]{KiDSr2} and DES \citep{DESnonDE}, which will cover, respectively, 1000 and 5000 square degrees with sub-arcsecond seeing. Comparable numbers are expected from the ongoing Hyper Suprime-Cam\footnote{\texttt{http://www.naoj.org/Projects/HSC/surveyplan.html}} (HSC) and the forthcoming Javalambre Physics of the Accelerating Universe Astrophysical Survey \citep[J-PAS;][]{JPASredBook} projects. These numbers are expected to increase even further in the near future, with the operation of the Large Synoptic Survey Telescope \citep[LSST;][]{LSST20} and Euclid\footnote{\texttt{http://www.euclid-ec.org/}} \citep{EuclidSciBook}, which are both expected to detect $\mathcal{O}\left(10^5\right)$ systems with arcs \citep{Collet2015}. The vast majority of the current samples of arc systems involve a visual search and classification. This is true for the targeted surveys and also for the wide-field imaging surveys, where either the full footprint or cutouts around potential lenses (e.g., luminous red-galaxies, galaxy clusters) are visually inspected. This manual procedure is still possible for current surveys with good image quality, which cover at most few hundred square degrees. However, it will become prohibitive for DES and KiDS, and especially for LSST and Euclid. Therefore, the development of automated arc finding methods is absolutely needed for the scrutiny of these surveys in the quest for gravitational arcs. Regardless of the size of the survey, automated arc detection is important for an objective and reproducible definition of arc samples, which often includes the determination of arc properties. This is of course critical for arc statistics \citep[see, e.g.,][]{2013SSRv..177...31M,2016ApJ...817...85X} and for any comparison of real and simulated data \citep[e.g.,][]{2005Horesh,2011MNRAS.418...54H} and among different data sets \citep[e.g.,][]{2010MNRAS.406.1318H}. Motivated by these needs, several automated methods to find gravitational arcs have been proposed in recent years. Most of these methods focus on pattern recognition, i.e., on identifying shapes that look like gravitation arcs, in particular, thin and elongated structures \citep[e.g.,][]{2004Lenzen,2005Horesh, 2006Alard,2007Seidel, 2012More}; in some cases, these methods require a degree of curvature \citep[e.g.,][]{2007Estrada,Kubo2008}. \citet{Maturi2014} combine this approach with a multicolor selection of the sources. \citet{Marshall2009robot} use lens inverse modeling to find strong lenses, i.e., assuming that a given object in an image is a consequence of lensing and determining whether the lensing solution is favored by the data. More recently, new arc finders have been proposed that subtract the lens candidate (usually early-type galaxies) light distribution, either using two bands, as in \citet{RINGFINDER}, or by modeling the lens in a single band, as in \citet{JosephPCA2014} and \citet{Brault2015}. The residuals are then investigated, using their shapes \citep{JosephPCA2014}, by color selection \citep{RINGFINDER}, or with inverse modeling \citep{Brault2015}. The inverse modeling approach is particularly interesting as it uses the physics of lensing to find candidates, however its applicability has been restricted to galaxy-scale lenses; this is because the automated modeling is much more tractable in this case owing both to the simplicity of the lens model and the identification of the images. The lens subtraction approach is a necessity for galaxy-scale lenses, especially when observed from the ground, as the arcs can be embedded in the galaxy's light. On the other hand it is less critical for arcs on cluster scales, which span larger angular sizes than the galaxies and the PSF. Our main interest here is to look for arcs in galaxy groups and clusters that can be found even with a single band survey (as in CS82). Therefore, in this paper we focus only on the development of a pattern-recognition based arc finder. Most arc finders in this category use sets of measurements of the objects, such as ellipticity, length, $L$, width $W$, and {\it arcness,} to determine whether they are arc candidates or not. They usually employ hard (i.e., fixed and mutually independent) cuts, whose values may be arbitrarily assigned or tuned using data or simulations. However, given the diversity of arc properties (shapes, sizes, $S/N$ ratios, etc.) and their physical origin, different cuts could perform better in different regions of the multidimensional space of arc parameters. For example, arcs may be very elongated and not necessarily curved for galaxy cluster lenses, while arcs are not as drastically elongated for galaxy lenses but exhibit a clear curvature. Therefore, a flexible criterium based on a combination of parameters may be more efficient than applying hard cuts. This is a typical situation in which machine learning methods can be extremely helpful. A suitably trained algorithm can then classify the objects into arcs or not, given a set of input values for the object features. Such training can be carried out either on real data (on objects previously known to be arcs) or using simulations, by feeding the algorithm with a large set of arc and nonarc samples. This process is characteristic of supervised learning methods, of which the most well known is the back-propagation artificial neural network \citep[ANN;][]{williams1986learning,rumelhart1988learning}. The choice of the set of input parameters is as important as the choice of the classification method and its configurations. In this work we adopt measurements derived from the Mediatrix filamentation method \citep{2012Bom,BomMediatrix}, a novel iterative technique that decomposes elongated objects into segments along their intensity ridge line. This method provides several morphological parameters that are well suited to characterize arcs, including the length along the ridge line, $L$, the width $W$, and, most notably, estimates of the object center of curvature and its significance. Therefore, the purpose of this work is to construct an ANN gravitational arc finder based on the Mediatrix filamentation method, or {\it ANN Mediatrix Arcfinder} (AMA) for short. We use a sample of simulated gravitational arcs and a sample of nonarcs from HST images to train and validate the ANN. This sample is used to pin down a few configurations among the many possible choices involved in the ANN detection process: the types of images used for the training, the selection of inputs given to the ANN, the number of neurons, and the final threshold for classification. As an illustration of the application of the method to real data, we consider four galaxy cluster images from HST and run the AMA on them, comparing the results with the training and validation. The paper is organized as follows. Next section provides an overview of the AMA algorithm: the processes to detect and select the objects, the measurements carried out on them, in particular the Mediatrix filamentation, and the ANN used in the arc identification. Section \ref{train_and_validation} describes the training and validation process, the samples used, and the tuning of the ANN for arc detection. Section \ref{aplication} shows an example of application of the method on real data. Final discussions and concluding remarks are presented in section \ref{discusao}. | \label{discusao} The purpose of the paper is to present the AMA% \footnote{The method is implemented in the python language and the source code is available upon request to the authors. The training sample is also available upon request.} and provide a simple example of application, first illustrating the training and validation process on simulated arcs and then the application to a real, albeit small, data set. The major novel aspects of the present work are the use of the Mediatrix method in step $3$ of the arc finding process and the use of simulations. The simulated arcs are used not only to train, but also to find a good set of ANN configurations for step $4$. There is room for many improvements in the processes described in this paper, most notably in the use of more realistic simulations and increasing the number of systems in the application to real data, but also in other aspects arc detection. On the object identification % and segmentation side, betterments can be implemented in order to detect faint sources, avoid the breaking of large arcs, improve on the deblending with nearby objects, and to find and segment sources in high background regions, which are known issues for arc identification. Although \texttt{SExtractor} does not deal with all these issues in an optimal way for arc detection, it has been used in many arc finders for their object identification to some degree \citep[e.g.,][]{2005Horesh,2007Estrada,Kubo2008,Marshall2009robot,JosephPCA2014, Maturi2014}. A better performance than running \texttt{SExtractor} in a straight way, as in the current work, has been obtained by carrying out multiple runs of this code with different thresholds \citep{2005Horesh}, or by using \texttt{SExtractor} alone for pixel thresholding \citep{Kubo2008}. Other arc finders use different approaches for object detection and segmentation, which are specifically oriented toward identifying arcs \citep[e.g.,][]{2004Lenzen,2006Alard,2007Seidel,2012More,2016ApJ...817...85X}. A key issue for detection and segmentation for arcs is that these objects are often embedded in the haloes of bright galaxies (especially for galaxy-scale arcs and radial arcs) or blended with foreground galaxies (especially for arcs in clusters). One approach that has been implemented to address this issue is to fit and subtract the light profile of galaxies, as in \citet{2005ApJ...627...32S,Brault2015}. Several % codes have been proposed in the literature to this end \citep[e.g.,][]{PyMorph2010,GALPHAT2011,GALAPAGOS2012}, often running \texttt{galfit}~\citep{galfit32010} recursively to fit each galaxy by a combination of elliptical brightness distributions with \citet{sersic} profiles. Advanced versions of \texttt{SExtractor} also fit and subtract all identified objects in a field \citep[e.g.,][Moraes et al., in prep.]{SEMF2012ApJ...757...83D,SEMF2015A&A...578A..79D}. Other schemes to subtract objects from images, which could be useful for arc finding have also been proposed \citep[e.g.,][]{JimenezTeja2012}. A different approach has been carried out by \citet{2016ApJ...817...85X}, who propose a new detection and segmentation scheme, working in intensity difference space, which performs well in bright halos without the need of subtraction. Regarding the preselection, for applications to wide-field surveys this phase must also include the removal of image artifacts such as satellite tracks, star spikes, and regions with a large amount of noise or with a steeply varying background. In the current example we remove noisy regions by simply cutting off objects that are close to the CCD borders. However the AMA already includes a proper handling of survey masks, which are produced to avoid bright star halos and spikes, satellites, and other image features. The approach of \citet{2016ApJ...817...85X} is also well suited to remove diffraction spikes without the need to use masks and may be useful for less bright nonmasked stars whose spikes could contaminate the arc detection. For object measurement we propose the use of the Mediatrix filamentation method, and several parameters derived from it, as it was designed for elongated and curved objects. Most arc finders end up using less parameters and simpler measurement schemes to characterize the arc candidates, such as $L$ and $L/W$, and only in a few cases include estimates of the curvature \citep{2007Estrada,Kubo2008}. However, other sets of inputs in addition to the Mediatrix inputs could be given to the ANN, such as higher order moments of the brightness distribution, including the arcness \citep{Kubo2008}. We argue that by using a machine learning algorithm for the final candidate selection (in this case a back-propagation ANN) one may achieve a better efficiency in finding arcs than using hard cuts on a few variables. By working on a multidimensional parameter space, it is possible to deal with correlations among the variables and empirically obtain combinations that represent gravitational arcs. For example, arcs tend to be more curved and smaller at galaxy scales than in massive clusters, such that a single cut in $L/W$ or arcness would not be optimal for finding arcs at both scales. Artificial neural networks were first used by \citet{2007Estrada} to search for arcs. In their case the simulated arcs are simply sections of a circle with a surface brightness profile that is uniform along the tangential direction and is convolved with a Gaussian with FWHM similar to the typical seeing of the images. The ANN is trained using a hundred such simulated arcs covering a range of sizes and brightnesses, which are added to SDSS images. The objects are also identified with \texttt{SExtractor} and the preselection is also carried out using an estimate of the object's elongation. Finally, four inputs are given to the ANN, based on a fit of the object by a circle and on a determination of the object's length (using its furthest pixels). \citet{2007Estrada} study the efficiency for recovering the simulated arcs both for a visual inspection and for the automated process as a function of peak surface brightness and $L/W$. A maximum efficiency of $40\%$ (with respect to the number of simulated arcs) is achieved in the automated search. The present work can be seen as an improvement on \citet{2007Estrada} in the sense that we use more realistic simulated arcs and a wider set of input measurements % well suited to characterize the arcs, in addition to tuning the ANN configurations to improve the completeness. Of course the key to a good performance in a learning algorithm is the realism of the training sample. Many improvements can be incorporated into the simple \texttt{AddArcs} simulations described in this work, such as considering more realistic lenses \citep[e.g.,][]{2011MNRAS.418...54H,2016ApJ...817...85X} and sources \citep{Kubo2008,Marshall2009robot, 2011MNRAS.418...54H}. Moreover, in addition to having a realistic sample of isolated arcs, those have to be added to real images, for example, to address the issue of blending with other sources and of embedding in the halo of bright galaxies in cluster cores. Other works have used simulations to test arc finders, define their parameters, or train their methods, and in some cases determine the selection functions \citep[e.g.,][]{2005Horesh,Kubo2008,Marshall2009robot, 2011MNRAS.418...54H, RINGFINDER,JosephPCA2014,Brault2015,2016ApJ...817...85X}. Another possibility is to use the growing number of strong lensing systems detected in wide-field surveys and HST images to perform the training on real data sets. For example, over 600 candidate systems have been detected in recent studies using CFHTLS data \citep{Maturi2014,RINGFINDER,Brault2015,SPACEWARPSII,ParaficzCFHTLS}, which could be used to train and better characterize the AMA and other arc finders. By training and validating the ANN with more realistic simulated arcs or with real data, we expect to reach a better agreement for $c$ in comparison to applications to other data sets (and therefore to achieve a higher completeness), which is different from what we found when applying our trained ANN to the HST data. In general, the several arc finders proposed in the literature carry out an end-to-end approach from the science image to a list of arc candidates, implementing all four steps that we refer to in this paper. However, they have their own solutions for each step with different degrees of sophistication and specificity for finding arcs. For example, in this work we focus on the third and fourth steps, respectively, by using a set of object measurement parameters that are well suited to arcs and a trained ANN, while most methods use simple cuts on a few parameters for the final classification. On the other hand, we use a generic object segmentation code that is not optimal for arcs. For most methods, these four steps could be performed interchangeably. Therefore, if they are presented in a modular way, we would be able to test the performance of each step independently. Moreover, new arc finders could be created combining the solutions for each step from those available that work better in specific situations. Several possible concatenations of the arc finder modules could also be compared using their end-to-end performance. Given that about a dozen finders have already been proposed, and that we are on the verge of applying them on large data sets, it would be very useful to carry out such a comparison. This is one of the avenues we will pursue for future work \citep[see,][for preliminary results with a few arc finders]{2015mgm..conf.2088D}. After a decade of progress in the development of gravitational arc finders, several codes are ready for exploring the new generation of wide-field surveys in the quest for gravitational arcs. However more progress is still needed for fully automated runs so as to produce samples that can be readily exploited for their applications. Besides improving the efficiency in some situations, the most important is to limit the false positives to a level low enough to be corrected for and certainly less than the number of real systems, as thousands to hundreds of thousands strong lenses are expected in the forthcoming data. Different strategies have been proposed and implemented to address these issues. Combining aspects of these solutions, which are more suited to each step of arc detection, seems a natural way to proceed. We believe that using neural networks or other machine learning methods may provide an important contribution to the task of selecting more complete and pure samples of gravitational arcs for a broad range of deflector scales and backgrounds. | 16 | 7 | 1607.04644 |
1607 | 1607.06194_arXiv.txt | We analyze the SDSS DR12 quasar catalogue to test the large-scale smoothness in the quasar distribution. We quantify the degree of inhomogeneity in the quasar distribution using information theory based measures and find that the degree of inhomogeneity diminishes with increasing length scales which finally reach a plateau at $\sim 250 \hmpc$. The residual inhomogeneity at the plateau is consistent with that expected for a Poisson point process. Our results indicate that the quasar distribution is homogeneous beyond length scales of $250 \hmpc$. | The assumption that the Universe is statistically homogeneous and isotropic on sufficiently large scales is fundamental to modern cosmology. Our current understanding of the Universe comes from the vast amount of cosmological observations which are hard to interpret without relying on this assumption. Therefore it is important to verify this assumptions using various observations. There are a multitude of evidences favouring isotropy such as the isotropy of the CMBR \citep{penzias,smoot,fixsen}, isotropy in angular distributions of radio sources \citep{wilson,blake}, isotropy in the X-ray background \citep{peeb93,wu,scharf}, isotropy of Gamma-ray bursts \citep{meegan,briggs}, isotropy in the distribution of galaxies \citep{marinoni,alonso}, isotropy in the distribution of supernovae \citep{gupta,lin} and isotropy in the distribution of neutral hydrogen \citep{hazra}. But the local isotropy around us alone is not sufficient to assure large-scale statistical homogeneity. One requires to combine the local isotropy with the Copernican principle to infer the large-scale statistical homogeneity of the Universe. The Copernican principle states that we do not occupy a special location in the Universe which itself requires validation. One can infer the large-scale statistical homogeneity from local isotropy only when it is assured around each and every point in the Universe. So it is not straightforward to infer large-scale statistical homogeneity of the Universe from the local isotropy. A large number of studies \citep{martinez94,borgani95,guzzo97,cappi,bharad99,pan2000,yadav,hogg,prakash,scrim,nadathur,pandey15,pandey16} find that the galaxy distribution behaves like a fractal on small scales but on large-scale the Universe is homogeneous. Most of these studies claim to have found a transition to homogeneity on scales $70-150 \, h^{-1} \rm {Mpc}$. Contrary to these claims a number of studies \citep{pietronero,coleman92,amen,joyce,labini07, labini09, labini11} reported multi-fractal behaviour on different length scales without any transition to homogeneity out to the scale of the survey. The results from these studies clearly indicates that there is no clear consensus in this issue yet. There would be a major paradigm shift in cosmology if the assumption of cosmic homogeneity is ruled out with high statistical significance by multiple data sets. The most important implication of inhomogeneities comes from the averaging problem in General Relativity through their effect on the large scale dynamics known as backreaction mechanism. The backreaction mechanism is known to cause a global cosmic acceleration without any additional dark energy component \citep{buchert97, schwarz, kolb06, buchert08,ellis}. The Sloan Digital Sky Survey (SDSS) \citep{york} is the largest and finest galaxy redshift survey todate. The quasars are the brightest class of objects known as Active Galactic Nuclei (AGN). The high luminosities of quasars allow them to be detected out to larger distances. The SDSS DR12 quasar catalogue provides us an unique opportunity to test the assumption of cosmic homogeneity on the largest accessible scale due to its enormous volume coverage. The presence of large quasar groups (LQG) in the quasar distribution is known for quite some time. \citet{clowes} identified a huge LQG with characteristic size $\sim 500 \hmpc$ at $z\sim 1.3$ in the DR7 quasar catalogue and claimed that this structure is incompatible with large-scale homogeneity indicating possible violation of the cosmological principle. A number of subsequent studies \citep{nadathur,marinello} pointed out some flaws in interpreting LQGs as structures. However if such structures really exist then they owe an explanation. In the present study we test the large-scale homogeneity in the SDSS DR12 catalogue using information theory based methods \citep{pandey13,pandey16}. We do not address the LQGs separately but the presence of any such structures in the quasar distribution are clearly expected to boost the signal of inhomogeneity up to noticeably larger length scales. Throughout our work, we have used the flat $\Lambda$CDM cosmology with $\Omega_m = 0.3, \; \Omega_{\Lambda} = 0.7 \mbox{ and } h = 1$. A brief outline of the paper follows. In section 2 we describe the method of analysis followed by a description of the data in section 3. We present the results and conclusions in section 4. \begin{figure*} \resizebox{7cm}{!}{\rotatebox{0}{\includegraphics{qsomag.eps}}}% \resizebox{9cm}{!}{\rotatebox{0}{\includegraphics{qsoden.eps}}}\\ \caption{The left panel shows the SDSS quasars in the redshift-absolute magnitude plane. The upper region and the lower region in this panel represent the selected and discarded quasars respectively. The right panel shows the comoving number density of quasars as a function of radial distance $r$. We compute the density in shells of uniform thickness $ 30.73 \hmpc$ in the radial direction.} \label{fig:qsoden} \end{figure*} \begin{figure*} \resizebox{7cm}{!}{\rotatebox{0}{\includegraphics{qsoshn_SH.eps}}}% \resizebox{7cm}{!}{\rotatebox{0}{\includegraphics{dqso_DH.eps}}}\\ \resizebox{7cm}{!}{\rotatebox{0}{\includegraphics{qsoshn_DH.eps}}}% \resizebox{7cm}{!}{\rotatebox{0}{\includegraphics{dqso_SH.eps}}}\\ \resizebox{7cm}{!}{\rotatebox{0}{\includegraphics{qsokld.eps}}}% \resizebox{7cm}{!}{\rotatebox{0}{\includegraphics{qsomr.eps}}}\\ \caption{The top left and middle left panels show $1-\frac{H_{d}}{(H_{d})_{max}}$ and $(H_{d})_{max}-H_{d}$ as a function of length scales respectively. The slopes of the respective quantities in the quasar sample are shown in the top right and middle right panels. The bottom left panel shows the KL divergence as a function of length scales and the bottom right panel shows the number of voxels available at each length scales. The errorbars for the quasar sample are obtained from $30$ bootstrap samples and the errorbars for the Poisson samples are obtained from $30$ Monte Carlo realizations of the same.} \label{fig:qsores} \end{figure*} | We present our results in \autoref{fig:qsores}. In top left and middle left panel of the \autoref{fig:qsores} we show respectively the variation of $1-\frac{H_{d}}{(H_{d})_{max}}$ and $(H_{d})_{max}-H_{d}$ with increasing grid sizes for the quasar sample and its random mock counterparts. Both the quasar sample and the random samples have a mean inter-particle separation of $91 \hmpc$ and they are hardly distinguishable below this length scale but as the grid size increases the differences become evident. In both of the plots we find that the quasar sample has a higher information content than the Poisson samples throughout the entire length scale ranges due to the gravitational clustering. However the differences diminish with increasing length scales and at $\sim 250-300 \hmpc$ the results for the Poisson samples lies within $1-\sigma$ errorbars of the same for the quasar sample. The residual inhomogeneities beyond $250 \hmpc$ are consistent with what one would expect for a homogeneous Poisson point process. We show the rates of change of $1-\frac{H_{d}}{(H_{d})_{max}}$ and $(H_{d})_{max}-H_{d}$ in the quasar sample with increasing length scales in the top right and middle right panels of \autoref{fig:qsores} respectively. We find that the rates of change for both of the measures in the quasar sample diminish nearly to zero with tiny errorbars at $\sim 250 \hmpc$. In the bottom left panel of \autoref{fig:qsores} we show the KL divergence measure as a function of length scales in the quasar sample. The KL divergence also indicates that the inhomogeneities diminish with increasing length scales finally reaching a plateau at $\sim 250-300 \hmpc$. However it is worth mentioning here that though the entropy is sensitive to the higher order moments of a distribution it may not capture the signatures of the full hierarchy of correlation functions. A Minkowski Functional analysis of SDSS LRGs by \citet{wiegand} find significant deviations from the $\Lambda$CDM mock catalogues on scales of $500 \hmpc$. It has been suggested that the quasars inhabit dark matter halos of constant mass $\sim 2 \times 10^{12} h^{-1} M_{\odot}$ from redshifts $z \sim 2.5$ (the peak of quasar activity) to $z\sim 0$ and their large scale linear bias evolves from $b=3$ at $z\sim 2.2$ to $b=1.38$ at $z\sim 0.5$ \citep{shen, ross, geach}. Our quasar sample extends from $z=2.2$ to $z=3.2$ for which we expect a large scale linear bias of $\gtrsim 3$. As the quasars inhabit rarer high density peaks one would expect the quasar sample to be homogeneous on even larger scale than the SDSS main galaxy sample and the LRG sample. Our result is consistent with our earlier studies on the SDSS main galaxy sample \citep{pandey15} and the LRG sample \citep{pandey16} for which we find a transition scale to homogeneity at $\sim 150 \hmpc$. It may be noted here that there are $414$ independent voxels (bottom right panel of \autoref{fig:qsores}) at grid size of $250 \hmpc$ and each voxel is expected to host $\sim 21$ quasars provided the distribution is homogeneous beyond this length scale. This number is certainly very small due to the small number density of the quasar sample. Further we can not have access to spatial hypersurface of constant time. So any analysis of homogeneity on large scales would unavoidably incorporate some signatures of the time evolution. Despite these difficulties it is interesting to note the degree of homogeneity in the quasar sample beyond a length scale of $250 \hmpc$. Last but not the least we prepare $4$ different quasar samples with different set of values for the coefficients $(a,b,c,d)$ used to define the limiting magnitudes at different redshifts. Irrespective of our choice we find the same transition scale to homogeneity in each of these quasar samples. We finally conclude that the SDSS quasar distribution is homogeneous beyond $250 \hmpc$, for the information theoretic measures employed in this paper. | 16 | 7 | 1607.06194 |
1607 | 1607.05532_arXiv.txt | We study the evolution of galactic magnetic fields using 3D smoothed particle magnetohydrodynamics (SPMHD) simulations of galaxies with an imposed spiral potential. We consider the appearance of reversals of the field, and amplification of the field. We find magnetic field reversals occur when the velocity jump across the spiral shock is above $\approx$20~km~s$^{-1}$, occurring where the velocity change is highest, typically at the inner Lindblad resonance (ILR) in our models. Reversals also occur at corotation, where the direction of the velocity field reverses in the co-rotating frame of a spiral arm. They occur earlier with a stronger amplitude spiral potential, and later or not at all with weaker or no spiral arms. The presence of a reversal at a radii of around 4--6 kpc in our fiducial model is consistent with a reversal identified in the Milky Way, though we caution that alternative Galaxy models could give a similar reversal. We find that relatively high resolution, a few million particles in SPMHD, is required to produce consistent behaviour of the magnetic field. Amplification of the magnetic field occurs in the models, and while some may be genuinely attributable to differential rotation or spiral arms, some may be a numerical artefact. We check our results using {\sc athena}, finding reversals but less amplification of the field, suggesting that some of the amplification of the field with SPMHD is numerical. | Understanding the magnitude and morphology of galactic magnetic fields is a long-standing problem in galactic astronomy. The morphology of galactic magnetic fields is difficult to observe, and poorly understood theoretically. One difficulty is measuring the direction of the magnetic field. Faraday rotation has indicated that the Milky Way contains reversals of the magnetic field (where the magnetic field vector reverses direction), but as yet we do not know whether reversals occur in other galaxies. Another problem is understanding the location of `magnetic spiral arms', where the ordered component of the magnetic field is strongest. In some spiral galaxies, these tend to be aligned with the optical spiral arms as expected (e.g. M51, \citealt{Fletcher2011}) but in other galaxies the ordered component is strongest in the inter-arm regions \citep{Beck2007a}. Many models and simulations use dynamo theory to try and interpret these phenomena, but there is no consensus about the origin of reversals, the differences in morphologies between galaxies, or the physical cause or timescale of magnetic field growth in galaxies. Most observations indicate there is one reversal of the magnetic field in the Milky Way, in the region of the Sagittarius spiral arm \citep{Frick2001,Nota2010,VanEck2011,Beck2011}. \citet{Han2006} suggest there may be numerous reversals within the galactic disc, although most surveys rule out any reversals in the outer Galaxy \citep{Brown2001,VanEck2011}. Reversals have been observed in a number of simulations, with different types of numerical codes, with and without an explicit turbulent dynamo. Many explanations of reversals involve a turbulent dynamo, and for example, \citet{Chamandy2013} investigate reversals in different types of spiral galaxies starting from a primordial magnetic field strength. Alternatively reversals could be induced by vertical oscillations of the gas which produce a vertical dynamo \citep{Ferriere2000,Gressel2013}. However reversals are also seen in simulations without explicit dynamo terms (e.g. \citealt{Pakmor2013}). Although the reversals seem to readily occur, there is little explanation of what causes the reversal or what properties the reversals depend on. A further question is how the magnetic field has been amplified to the values seen in low-redshift galaxies, from primordial field strengths. Again, dynamo theory is often invoked to explain amplification of the magnetic field (\citealt{Parker1971}, see also recent review by \citealt{Brandenburg2015}). However, a number of processes may contribute to the amplification of the field including turbulence (as exemplified by standard turbulent dynamo theory), stellar feedback and differential rotation \citep{Gaensler2005}. Recent simulations of isolated galaxies indicate that supernovae feedback can drive a small scale dynamo \citep{Becka2012,Rieder2016}. Differential rotation also amplifies magnetic fields (see simulations by \citealt{Kotarba2009}) although there are doubts that differential rotation alone can be responsible for the present day magnetic field strengths \citep{Brandenburg2015}. In addition to explaining current field strengths, turbulent dynamos have also been explicitly included as an extra term in MHD calculations to reproduce ordered fields (or magnetic spiral arms) between the optical spiral arms \citep{Chamandy2013,Moss2015}. There are numerous challenges for studying galactic magnetic fields numerically, including the difficulty of modelling magnetic fields in the smoothed particle hydrodynamics (SPH) method, as well as achieving high enough resolution to effectively model the interstellar medium (ISM). In our previous work modelling spiral galaxies \citep{DP2008}, we used an SPMHD code whereby the magnetic fields were represented by Euler potentials. The magnetic fields were found to smooth out structure in the disc, and the spiral arms were also seen to induce some measure of disorder in the magnetic field, more prominent in the inter-arm regions. The Euler potentials method had the advantage that the magnetic divergence is zero by construction. However the morphology of the field was limited --- in particular, winding up of the field after multiple rotations of the disc cannot be captured with the Euler potentials. An alternative is to use divergence cleaning methods in SPH to try and limit the value of the magnetic divergence, but these have so far only been applied to galaxy simulations employing relatively low resolutions (e.g. \citealt{Pakmor2013}) or which study only small regions of the ISM. \citet*{Tricco2016} showed that it is possible to model a turbulent dynamo in SPMHD using 3D simulations of a small-scale dynamo in a turbulent, periodic box (representing the ISM). Recent grid-based simulations have been used to investigate the role of magnetic fields on gravitational and hydrodynamic instabilities in spiral arms. While magnetic fields were found to have only a small effect, they can limit instabilities seen in purely hydrodynamic models \citep{Lee2014,Kim2015}, and suppress gravitational instabilities in cases with low shear \citep{KOS2002}. A number of grid-based simulations have also self consistently included a galactic dynamo by modelling supernova feedback \citep{Hanasz2009,Kulpa2011}, but again the simulations were relatively low resolution. In the present paper, we model isolated spiral galaxies with an imposed spiral potential. Rather than studying the amplification of a primordial field, we start our calculations with magnetic field strengths closer to present day values. Since we observe reversals in our simulations, we investigate their dependence on the galactic potential in order to understand what is required for magnetic field reversals in galaxies. To test the robustness of our results, we examine the evolution of the magnetic divergence, perform resolution tests, vary the strength of the cleaning prescription for the magnetic field, and compare with a grid-based code. | We have performed SPMHD calculations of spiral galaxies with a fixed spiral potential. Our calculations adopt simple physics, assuming an isothermal ISM, and do not include self gravity, star formation or stellar feedback, but we vary the form of the galactic and spiral potential. We observe a clear dependence of the morphology of the magnetic field on the form of the spiral potential, in particular the presence of magnetic field reversals. We find that reversals of the magnetic field tend to be associated with large velocity changes across the spiral shock. We find reversals only occur if the velocity difference is $\gtrsim20$ km s$^{-1}$. The locations of the reversal in the disc are shown to be close to the location of the maximum change in velocity. This is where the magnetic field experiences the maximum distortion by the velocity field, and where the consequent straightening of the field leads to a reversal. Simulations with only weak, or no spiral arms do not show reversals, or at least not until times later than we examine. We tested our findings with the grid code {\sc Athena}, and although not completely comparable, we saw reversals occur at similar locations to with SPMHD, and with similar dependence on spiral arm strength. The location of reversals, and large velocity changes, tend to be coincident with the ILR, as noted previously by \citet{Linden1998}. However we do not rule out reversals due to large changes in velocity not coincident with resonances, e.g. at the centre of the disc, or due a perturber or collision at larger radii. Furthermore, although we have only examined the case of a fixed spiral potential, we would expect our idea that large velocity changes induce reversals to also be true in simulations with transient spiral arms, if there is similarly a large velocity gradient. In addition to the reversals typically seen near the ILR, we also see a reversal near corotation in our model MHDN4High$\Omega$, the only model where corotation is within the simulation domain. At corotation, even though there is no spiral shock (and no change in velocity), the velocity field reverses moving from inside to outside corotation, and so a reversal is not surprising. We also made a simple comparison of the location of our reversal with the Milky Way. Both our fiducial simulation and the observations of the Milky Way show a reversal between the Sun and the centre of the Galaxy, but we caution that the dynamics of the Milky Way are not well known, and we do not know how well our simple galaxy model resembles the real Galaxy. We also examined amplification of the field. We see significant amplification of the field, by a factor of 10 or more, particularly in the centre of the disc. Theoretical predictions indicate a linear increase in the field, whereas we see a superlinear increase in the centre. Furthermore our results with {\sc Athena} show only a small (typically $\lesssim$ 10 \%) increase in the field strength. Thus some of the amplification of the field appears to be numerical with SPMHD. Neither field amplification nor field reversals were seen in our simulations with Euler potentials \citep{DP2008} where the field evolution is limited. Other vector potential methods may avoid the problem of over amplification of the field, whilst also allowing phenomena such as reversals, but these methods may present other problems compared to the divergence cleaning method presented here \citep{Price2010}. Comparisons of our results with and without magnetic fields suggest that magnetic fields have only a minor effect on the disc structure, merely smoothing out substructure in the disc similar to an extra pressure term. This is a similar conclusion to our previous calculations with Euler potentials, and broadly similar to other work \citep{Lee2014,Kim2015}. However we note that our simulations are simplified, and we do not consider for example the formation or collapse of molecular clouds by self gravity. Similar also to our work with Euler potentials, the field tends to be slightly more ordered, and stronger in the spiral arms, and more random and weaker in the inter-arm regions. We have also performed a resolution study and examined the effect of the divergence cleaning method. We find that resolution is important in these simulations, in particular we conclude that 1 million particles in a global simulation is not sufficient to obtain reliable results, but that our simulations with 4 and 8 million particles are much more consistent. We checked how well our divergence cleaning method works, and find that typically $h \vert \nabla\cdot\mathbf{B}\vert/\vert\mathbf{B}\vert$ is $\sim 0.01$ for our calculations, thus they are not affected by erroneously high values of $\nabla\cdot\mathbf{B}$. Furthermore our results are independent of the strength of the damping of the divergence. We do see edge effects start to develop at later times, but these can be diminished with stronger divergence cleaning. \appendix | 16 | 7 | 1607.05532 |
1607 | 1607.02915_arXiv.txt | {The recently published Yemeni observing report about SN 1006 from al-Yam\={a}n{\={\i}} clearly gives AD 1006 Apr $17 \pm 2$ (mid-Rajab 396h) as first observation date. Since this is $\sim 1.5$ weeks earlier than the otherwise earliest reports (Apr 28 or 30) as discussed so far, we were motivated to investigate an early sighting in more depth. We searched for additional evidences from other areas like East Asia and Europe. We found that the date given by al-Yam\={a}n{\={\i}} is fully consistent with other evidence, including: (a) SN 1006 {\em rose several times half an hour after sunset} (al-Yam\={a}n{\={\i}}), which is correct for the location of \d{S}an$^{c}$\={a}' in Yemen for the time around Apr 17, but it would not be correct for late Apr or early May; (b) the date ({\em 3rd year, 3rd lunar month, 28th day wuzi}, Ichidai Yoki) for an observation of a guest star in Japan is inconsistent (there is no day {\em wuzi} in that lunar month), but may be dated to Apr 16 by reading {\em wuwu} date rather than a {\em wuzi} date; (c) there is observational evidence that SN 1006 was observed in East Asia early or mid April; for the second half of April, a bad weather (early monsoon) period is not unlikely -- there is a lack of night reports; (d) the observer in St. Gallen reported to have seen SN 1006 {\em for three months}, which must have ended at the very latest on AD 1006 Jul 10, given his northern location, so that his observations probably started in April. We conclude that the correctly reported details give quite high confidence in the fully self-consistent report of al-Yam\={a}n{\={\i}}, so that the early discovery date should be considered seriously. } | Historic observations of supernovae (SN) are essential to understand SNe, neutron stars, and SN remnants (SNR): Historic reports can in principle deliver the date of the explosion (hence, the age of the SNR and, if existing, the neutron star) together with a light curve (hence, possibly the SN type), sometimes the colour and its evolution, and the position of the SN, which is needed to identify the SNR and, if existing, the neutron star. Such historic observations have been used very successfully for SNe 1006 (from East Asia, Europe, and Arabia), 1054 (from East Asia and Arabia), 1181 (only from East Asia), and SNe 1572 and 1604 (from East Asia and Europe), plus a few more SNe from the 1st millenium AD (see Stephenson \& Green 2002, henceforth SG02, and references therein). While the Arabic report of SN 1054 merely confirms a bright new star in Gemini/Taurus around AD 1054, the reports of SN 1006 present a lot of detailed information (Goldstein 1965), which were used to identify the SNR (Gardner \& Milne 1965). The transient celestial object of AD 1006 was listed as comet in Pingre (1783). Humboldt (1851) lists it as {\em new star}, based on the St. Gallen chronicle, dated incorrectly to AD 1012, and placed incorrectly in Aries. Sch\"onfeld (1891) corrects the date to AD 1006 (consistent date shift in St. Gallen chronicle) and the location to Scorpius (previous misreading of the Syriac {\em $^{c}$aqrab\={a}} for Scorpius as {\em 'emr\={a}} for Aries in Bar Hebraeus); he already used the chronicle of Bar Hebraeus and its source, namely the annals of Ibn al-Ath\={\i}r.\footnote{In the report by Lynn (1891), an English summary of the paper by Sch\"onfeld (1891), which was written in German, it is said that Sch\"onfeld (1891) would have corrected the date from AD 1006 to 1012, but the opposite is true.} Convincing evidence for SN 1006 was presented first by Goldstein (1965) based on the Arabic reports: $^{c}$Al\={\i} ibn Ri\d{d}w\={a}n (lived from AD 988 or 998 until 1061 in Cairo, Egypt; indeed, it was considered seriously before that he observed SN 1006 at an age of only eight years, the calculations could have been done later) reported the ecliptic longitude (15th degree of Scorpius), strong scintillation ({\em it twinkled very much ... large ... round in shape}), the size and/or brightness ({\em 2.5 or 3 times as large as Venus ... the intensity of its light was a little more than the quarter of that of moonlight}), the duration of the observations (some four months until conjunction with the Sun), that it did not move relative to the stars ({\em It remained where it was and it moved daily with its zodiacal sign}) and that he observed it as an eyewitness during {\em the beginning of my studies ... all I have mentioned is my own personal experience}; he listed the (calculated) positions of the planets as well as those of the Sun, the Moon, and its ascending node, from which Goldstein (1965) deduced the date of his observations to be the evening of AD 1006 Apr 30, the earliest certain observation accepted in SG02. $^{c}$Al\={\i} ibn Ri\d{d}w\={a}n also mentioned that {\em Sun and Moon met in the 15th degree of Taurus} when the transient object first appeared; only from that statement we can conclude that it appeared close to the conjunction of moon and sun (new moon on AD 1006 Apr 30 at 9:08h UT), so that SN 1006 was probably sighted (by him) on the evening of Apr 30. Goldstein (1965) presented another Arabic report of SN 1006 from Ibn al-Jawz{\={\i} (a historian, who lived AD 1116-1201 in Baghdad, Iraq) and, based on Ibn al-Jawz{\={\i}, also by Ibn al-Ath\={\i}r (a historian, who was born AD 1160 in Jazirat Ibn Umar, now Cizre in Turkey, and who died in 1233 in Mosul, Iraq), both about a very bright new star -- as well as a report from Morocco\footnote{By the author Ab\={u} l-\d{H}asan $^{c}$Al\={\i} b. $^{c}$Abdall\={a}h b. Ab\={\i} Zar$^{c}$ al-F\={a}s\={\i} (short: Ibn Ab\={\i} Zar$^{c}$, died in or after AD 1326) in the book entitled {\em Al-An\={\i}s al-mu\d{t}rib bi-rau\d{d}at al-qir\d{t}\={a}s f\={\i} akhb\={a}r mul\={u}k al-maghrib wa-t\={a}r{\={\i}}kh mad\={\i}nat F\={a}s}, the mentioned town of {\em F\={a}s} is now called Fes in Morocco, an edition of the work appeared in 1972 in Rabat, Morocco.} mentioning a {\em great star} [{\em najm}] {\em among the comets} and a {\em nayzak} ({\em spectacle} or {\em guest star} or {\em transient celestial object}, see Kunitzsch 1995). While the reports of Ibn al-Jawz{\={\i} (and Ibn al-Ath\={\i}r) give the date of first appearance as Friday, the 1st day of the Muslim month of Sha$^{c}$b\={a}n of the Muslim year 396h\footnote{The Islamic year 396 hijra (396h) started 396 lunar years after the start of the lunar year in which the Hijra took place, i.e. the emigration of the Islamic Prophet Mu\d{h}ammad from Mecca to Medina, known as Hijra; this era, i.e. the year 1h started in the evening of AD 622 Jul 16 (1st day from evening of 16th to evening of 17th July) according to most scholars -- but it may have been one day earlier (evening of 15th to evening of 16th July), see, e.g., de Blois (2000); according to Gautschy (2011, 2014), see www.gautschy.ch/$\sim$rita/archast/mond/Babylonerste.txt, new moon was on AD 622 Jul 14 (Julian calendar), so that the crescent new moon was not visible in the evening of AD 622 Jul 14, hardly visible in the evening of AD 622 Jul 15, but it was well visible in Mecca, Saudi Arabia, in the evening of AD 622 Jul 16 (Neugebauer 1929). While the date of the 1st day of year 1h was mistakenly given one day too early in footnote 1 in Rada \& Neuh\"auser (2015), all other dates there are correct. The uncertainty in the exact date of the start of the Hijra era affects only the conversion between the calculated Muslim calendar and some other calendar (like the Julian or Gregorian calender), and this uncertainty amounts to only one day.} (and visibility until Dh\={u} al-Qa$^{\rm c}$dah, i.e. roughly three months), the Moroccan report mentions that {\em it began to appear in the beginning (i.e. 1st) of Sha$^{c}$b\={a}n} and that {\em it lasted for a period of six months} (Goldstein 1965), i.e. from early May 1006 (Sha$^{c}$b\={a}n 396h) until October; however, SN 1006 was in conjunction with the Sun in October, so that they could have observed it until at most the heliacal setting in the middle of September; it is more likely that this {\em period of six months} were meant as rough statement (like about half a year) than that they observed heliacal rising in November. The Islamic date of 1 Sha$^{c}$b\={a}n 396h corresponds to AD 1006 May 2/3, evening to evening (e.g. Goldstein 1965), but only in the calculated Islamic calendar, while the real start of a month was fixed by observations (not by a calculated calendar); Muslim dates run from one evening to the next evening, a month starts with the evening of the first sighting of the crescent of a new moon. It was confirmed in Rada \& Neuh\"auser (2015) that the conversion of 1 Sha$^{c}$b\={a}n 396h to AD 1006 May 2/3 is correct when considering the first observation of the crescent (and also regarding the given week-day). In general, unless more information is available, the conversion from the calculated Islamic calendar to the Julian or Gregorian calender has an uncertainty of some 2 days due to (a) uncertainty in the start of the Hijra era (one day), (b) uncertainty as to which months and years in history had one extra day (in addition to 354 days in 12 lunar months -- given that a synodic month is not exactly 29.5 days), and (c) uncertainty as to when the new crescent moon was sighted first (e.g. Spuler \& Mayr 1961, Spuler 1963, Neuh\"auser \& Kunitzsch 2014). Goldstein (1965) also gives the Arabic texts for the four Arabic reports presented. Goldstein (1965) also gives an English translation of a Syriac report of SN 1006 by Bar Hebraeus (born AD 1226 in Malatya in Turkey, died 1286 in Maragha, now Iran), where it was specified that the new star was observed in the zodiacal sign of Scorpius. Cook (1999) presented another Arabic report of SN 1006 from Ya\d{h}y\={a} ibn Sa$^{c}$\={\i}d al-An\d{t}\={a}k\={\i}, Patriarch of Antioch (now Antakya, Turkey), who extended the chronicle of Eutychius of Alexandria (Egypt) for the time since circa AD 939 and died in AD 1066, according to which the new star was seen for four months since {\em Saturday, 2nd day in Sha$^{c}$b\={a}n} of the year 396h (AD 1006 May 3/4). Most recently, Neuh\"auser, Ehrig-Eggert, and Kunitzsch (2016) presented another original report about SN 1006, namely written by Ibn S{\={\i}}n\={a} (Avicenna). From the ecliptic longitude of the SN as given by $^{\rm c}$Al\={\i} ibn Ri\d{d}w\={a}n (and an error bar from $^{\rm c}$Al\={\i} ibn Ri\d{d}w\={a}n's assumed measurement precision) together with the declination limit from a St. Gallen observation of this SN (Goldstein 1965) and the Chinese right ascension range (from the Chinese {\em lunar lodge}), it was then possible to constrain the location of the SN and to identify the SNR (Gardner \& Milne 1965). SN 1006 is believed to have taken place on AD 1006 Apr 30 or earlier (see Rada \& Neuh\"auser 2015) in the constellation of Lupus with the following basic parameters: \begin{itemize} \item distance being $2.18 \pm 0.08$ kpc from the proper motion of ejecta in mas/yr and the shock velocity of filaments in km/s (Winkler et al. 2003),\footnote{The distance determination by Jiang \& Zhao (2007), $\sim 1.56$ kpc, is highly uncertain: it was obtained by interpreting a presumable observation of SN 1006 in AD 1016 as an effect of re-brightening of parts of the SNR, while SG02 argued that this late observation date is a mistake in a historical document.} \item extinction being A$_{\rm V} = 0.32 \pm 0.03$ mag (Schaefer 1996 from various techniques) or A$_{\rm V} = 0.31 \pm 0.10$ mag (Winkler et al. 2003) from the reddening of the very blue sub-dwarf (Schweizer \& Middleditch 1980) showing strong, broad absorption lines due to the SNR, so that it is located in the background at 1.05-2.1 kpc (Burleigh et al. 2000) to 1.5-3.3 kpc (Schweizer \& Middleditch 1980), \item peak apparent brightness being $-7.5 \pm 0.4$ mag from distance for a SN type Ia (Winkler et al. 2003), and \item the apparently young SNR G327.6+14.6 was identified as its remnant (Gardner \& Milne 1965, Reynolds et al. 1994). \end{itemize} While Damon et al. (1995) and Firestone (2014) claim that a $^{14}$C signal from SN 1006 was observed (in AD 1009), Menjo et al. (2005) argued that the $^{14}$C amplitude around AD 1009 may be consistent with typical Schwabe cycle modulation. (A $^{14}$C detection three years after the SN might be possible due to the carbon cycle, which takes a few years.) | It would be quite surprising if the observer, who is the original source for the reports of the Yemeni authors, really has observed and detected SN 1006 already several days before the other Arabic observers (e.g. Apr 30 by $^{c}$Al\={\i} ibn Ri\d{d}w\={a}n), and all East Asian observers, so that this early date was rejected by SG02 as artificial (based only on the text by Ibn al-Dayba$^{c}$). We would now like to discuss several arguments which can be interpreted in favour of an early observation, at least not excluding an early observation: \subsection{Yemen: {\em half an hour after sunset}} The text of al-Yam\={a}n{\={\i}} says: \begin{quotation} In the night of mid-Rajab (or: 15th of Rajab), in the year 396h, a star appeared from the east at half an hour after sunset. It was four times as large as Venus. \end{quotation} The position of SN 1006 indeed did rise {\em half an hour after sunset} at the location of \d{S}an$^{c}$\={a}' in the middle of April 1006, while this statement would not be true at the end of April or early May. We cannot exclude that al-Yam\={a}n{\={\i}} (and Ibn al-Dayba$^{c}$ or their source) calculated much later that SN 1006 rose half an hour after sunset at the given night of mid-Rajab 396h (AD 1006 Apr 17), which may not be impossible (as $^{c}$Al\={\i} ibn Ri\d{d}w\={a}n had probably also calculated the positions of Sun, Moon, and planets as given in his report several decades after SN 1006, Goldstein 1965, SG02), but this would appear more doubtful. There are otherwise no obviously calculated facts in the report. The report by our most original and earliest Yemeni source, al-Yam\={a}n{\={\i}}, is self-consistent regarding the early date and rising time. We should therefore consider it seriously. \subsection{Southern location of \d{S}an$^{c}$\={a}'} Since \d{S}an$^{c}$\={a}' is at 2400 m sea level and since it has a clear horizon towards the south,\footnote{The Jabal al-Nabi Shu$^{c}$ayb, the highest mountain in Arabia with 3666 m, is due west-south-west from \d{S}an$^{c}$\={a}', while SN 1006 was rising in the evening in the south-east.} an observation from here earlier than all other known observations may not appear impossible: going from such a large height to roughly sea level, where most of the other, later observers were located (e.g. Cairo, Japan, China) can change the atmospheric extinction for object low on the horizon by some 4 mag (Schaefer 1993); the limit for serendipitous discovery of a new star on the sky by naked-eye is some 0 to 2 mag according to Clark \& Stephenson (1977) and Strom (1994). This consideration does not exclude that other observers at low altitude (e.g. in China or Japan, see Sect. 3.4) would have observed SN 1006 in mid April, as they could have observed later in the night, when SN 1006 was higher above the horizon. The professional Chinese and Japanese astronomers have observed all night. The Yemeni observers, though, may have observed mainly around and shortly after sunset, close to the time of the last prayers. \subsection{Rising of SN 1006 before the Moon on Apr 17 only from Yemen} It is well possible that the Yemeni (and other Arabic) observers checked for the full moon on the evenings of Apr 15, 16, and 17 (full moon on Apr 16 at 9:35h UT), in order to know the date in their lunar month relative to the full moon. While on AD 1006 Apr 16, the Moon was above the horizon earlier than the location of SN 1006 as seen from Yemen, the Moon rose half an hour after SN 1006 on Apr 17 and even later on Apr 18, so that an observation of SN 1006 may appear more probable on Apr 17 from Yemen. At the other relevant observing sites (Morocco, Iraq, Japan, and in particular also in Cairo, Egypt, and Kaifeng, China), the Moon was rising before SN 1006 even on Apr 17, e.g. 20 minutes before the SN 1006 as seen from Kaifeng, today's name of the capital of China at that time. \begin{table*} \caption{Times of apparent rising of SN 1006, sunset, and moonrise for \d{S}an$^{c}$\={a}', Yemen and Kaifeng, China (times given in UT). We also list the separation ({\em sep}) between SN 1006 and the moon for the dates given, as seen from \d{S}an$^{c}$\={a}' (left) and Kaifeng (right) -- always given for two hours after local moonrise, so that both the SN and the moon were visible. Bold face times indicate cases, where SN 1006 rose before the Moon (e.g. on Apr 17 in \d{S}an$^{c}$\={a}', but not in Kaifeng.} \begin{tabular}{lllll|llll} \hline 1006 & \multicolumn{4}{c}{for \d{S}an$^{c}$\={a}', Yemen} & \multicolumn{4}{c}{for Kaifeng, China} \\ Apr & SN rise & sunset & moonrise & sep SN/Moon & SN rise & sunset& moonrise & sep SN/Moon \\ \hline 15 & 15:54 & 15:16 & 14:34 & $26.2^{\circ}$ & 12:34 & 10:57 & 10:05 & $27.6^{\circ}$ \\ 16 & 15:50 & 15:17 & 15:27 & $20.4^{\circ}$ & 12:30 & 10:57 & 11:05 & $20.7^{\circ}$ \\ 17 & {\bf 15:46} & 15:17 & {\bf 16:20} & $22^{\circ}$ & 12:26 & 10:58 & 12:06 & $20.8^{\circ}$ \\ 18 & {\bf 15:42} & 15:17 & {\bf 17:17} & $30.1^{\circ}$ & {\bf 12:22} & 12:59 & {\bf 13:07} & $23.3^{\circ}$ \\ 19 & {\bf 15:38} & 15:17 & {\bf 18:14} & $41.4^{\circ}$ & {\bf 12:18} & 12:59 & {\bf 14:07} & $38.9^{\circ}$ \\ 20 & {\bf 15:34} & 15:17 & {\bf 19:10} & $54.4^{\circ}$ & {\bf 12:14} & 11:00 & {\bf 15:07} & $51.9^{\circ}$ \\ \hline \end{tabular} \end{table*} This consideration does not exclude that other observers further north (e.g. in China or Japan) would have observed SN 1006 in mid April close to full moon, because the moon was sufficiently well separated from SN 1006, so that SN 1006 may have been observable -- depending on its brightness. The Yemeni observers had a particulary good reason (full moon) to observe in mid April. \subsection{Possible observation in Japan on Apr 16 or 28} In the medieval Japanese chronicle Ichidai Yoki, an independant and original source based on {\em Abe Yoshimasa, teacher in astronomy}, we can read: \begin{quotation} [AD 1006 Apr 28:] ... in 3rd year, 3rd lunar month, 28th day wuzi [25], a guest star entered Qi, \end{quotation} (Stephenson et al. 1977). The date given as {\em 3rd year, 3rd lunar month, 28th day} corresponds to AD 1006 Apr 28. East Asian reports often specify the date in addition with the day count in the sexagenary system of numbering days continuously from 1 to 60. However, as discussed in SG02, there is no day called {\em wuzi} (25) in the third lunar month of that year; the name and number of the day AD 1006 Apr 28 is {\em gengwu} (7); the relevant characters ({\em wu} and {\em zi} compared to {\em geng} and {\em wu}) are very similar, so that already Kanda (1935) suggested that this guest star was indeed observed on the 28th day of that lunar month,\footnote{Note that the Chinese and Japanese started the day-count in each (lunar) month with what we call {\em new moon}, i.e. conjunction of moon and sun, as confirmed by the fact that all of the dates of solar eclipses from (at least) AD 700 to 1200 are dated to {\em the first day of the month}, see listing in Xu et al. (2000), while the Arabs started the lunar month day count with the first sighting of the crescent (Quran, Sura 2, 189).} i.e. already on AD 1006 Apr 28 (and that, later on, a scribe made a mistake with the date number). However, as pointed out by SG02, two separate scribal errors need to be assumed. The {\em 28th day} of that lunar month (i.e. 1006 Apr) would definitely be a few days before the Chinese ({\em 2nd day of the fourth lunar month}, i.e. May 2/3) and Arabian observations, the latter being {\em at the beginning of Sha$^{c}$b\={a}n}, i.e. beginning of May, and 1006 Apr 30 for $^{c}$Al\={\i} ibn Ri\d{d}w\={a}n. It is very well possible that the latter character in the sexagenary date alone was mistranscribed. If Kanda (1935) is correct regarding his emendation of the second character in the sexagenary date (from {\em zi} to {\em wu}), but the first character ({\em wu}) is left to stand as it is in the received text, the date becomes a {\em wuwu} day (55), i.e. AD 1006 Apr 16. {\em Wu} might easily be mistranscribed as {\em zi}, as both characters contain two horizontal and one vertical strokes; the two characters are differentiated by an initial curved stroke in {\em wu} and hooks at the end of the first and third stroke of {\em zi}. See Fig. 1. \begin{figure} \begin{center} {\includegraphics[angle=0,width=5cm]{chinese-sn1006.jpg}} \end{center} \caption{{\bf Chinese characters:} If the Chinese characters have unintentionally changed due to a mistake made by a copying scribe from those shown in the left ({\em wuwu}, i.e. day 55) to those shown in the right ({\em wuzi}), the original date of the observation could have been AD 1006 Apr 16.} \end{figure} The Apr 16 date is then not the 28th day of the lunar month, as also specified in the source. However, such dates often consist only of the year of the emperor, the number of the lunar month, and the date in the sexagenary system, leaving out the day within the lunar month (see, e.g., Sect. 3.5). It is possible that the original source did not contain such a lunar day, but that it was amended later, or that 16 was mistranscribed as 28. If we posit a {\em wuwu} day (Apr 16), then the Ichidai Yoki record corresponds very closely to that of the Yemeni observers. However, just as the sexagenary date casts doubt on the Apr 28 date, the lunar date casts doubt on the possible Apr 16 date. The dating of the event in the Ichidai Yoki remains uncertain. There are two other reports from Japan on this SN: Meigetsuki (13th century) and Dainihonshi (completed AD 1715) both list a report for May 1, see e.g. SG02, but it is obvious that the latter depends on the former. The report discussed here from Ichidai Yoki is an independent medieval chronicle of unknown date. According to SG02, the observation of SN 1006 within a few days at different places in Arabia (30 Apr to May 2 only) as well as in China and Japan (about 1 May) may provide evidence against an earlier observation elsewhere, we have to see that there were 8 to 6 days from the earliest previously accepted first observation (28 Apr in Japan or 30 Apr in Cairo) to the latest reported first observation (May 6 in China: {\em 3rd year, 4th lunar month, day wuyin [=AD 1006 May 6]. A Zhoubo star appeared ...} from SG02 from Wenxiao Tongkao), which is quite a long time for such a bright SN, in particular also for an observation around the new moon. A possibly relatively long time between the first observations in the different countries (it may have been even more days between the first detections in different countries) can also be seen as evidence for unstable weather at least in some of those places. E.g., Rada \& Neuh\"auser (2015) provided evidence for bad weather on AD 1006 May 1 in Antiochia, now Turkey, and maybe Mosul, Iraq: The reports from there mention explicitly {\em Saturday, the 2nd of Sha$^{c}$b\={a}n} and {\em Friday, the 1st of Sha$^{c}$b\={a}n}, respectively, so that both started the month of Sha$^{c}$b\={a}n on the evening of (our) Thursday, May 2, even though the crescent new moon would have been well visible at those sites on the evening of May 1. The non-detection of the crescent on May 1 may indicate bad weather. \subsection{Possible early observation on Apr 3 in the SE in China} Reports of an even earlier sighting run as follows: \\ (i) Wenxian Tongkao: {\em Jingde reign period, third year, third lunar month (day) yisi [42] (=AD 1006 Apr 3). A guest star [ke xing] appeared (chu) in the south-east direction}, \\ (ii) Songshi Annals: {\em Jingde reign period, third year, third lunar month (day) yisi [42] (=AD 1006 Apr 3). A guest star [ke xing] appeared at the south-east}, and \\ (iii) Songshi Astronomical treatise: {\em Jingde reign period, third year, third lunar month (day) yisi [42] (=AD 1006 Apr 3). A guest star [ke xing] appeared at the south-east}, \\ the complete citations from SG02 with their additions in round brackets and our additions in square brackets. This guest star may have been another object, e.g., a comet (SG02). There are, however, also a few arguments in favour of a possible interpretation of this guest star as SN 1006: \\ (a) The guest star was seen in the {\em south-east} like SN 1006. \\ (b) The more general word for {\em guest star} [ke xing] was used and not a more specific word for {\em broom star} or {\em tailed star} [hui xing] or {\em fuzzy star} [xing bo], which would have indicated a comet. \\ (c) There are no additional Chinese or other records available on any additional comet or other object in or around early April. \\ (d) An important political meeting on April 17 is reported, which could have been a consequence of the very bright magnitude of the new star: {\em On the jiwei day (Apr 17), admonitory ministers were summoned to court and asked to speak openly about what should and should not be done} (Song shi 7.130), but the reason for the meeting is not specified. That the information from China about this SN is sparse, in particular for April and May, may be due to the difficult interpretation at that time: A solar eclipse was expected for 1006 May 30, which would have to be interpreted in a more negative sense for the emperor; the new star became a Zhoubo star, for which the historic Chinese texts offered both positive or negative interpretations (SG02). Once it became obvious that the solar eclipse did not take place at the capital, it was not necessary any more to consider a negative interpretation (for both eclipse and the new star), so that one could opt for the positive interpretation of the new Zhoubo star, which of course met the approval of the emperor. This is fully consistent with the Chinese texts dated to May 30, see SG02, and it could possibly explain why Chinese sources are unusually quiet about the bright new star in its first few weeks. \subsection{No other East Asian observations Apr 17-28/30 (and Apr 4-15)} If the Japanese have observed SN 1006 on Apr 16 (and maybe the Chinese already on Apr 3), then again later since May 1, it would be surprising that there are no reports left from the professional astronomers in China and Japan about any observations inbetween, i.e. from Apr 17 to the end of April (or even Apr 4 to 15). Are there any East Asian observations known for the intermediate periods from Apr 17 to the end of April (or even from Apr 4 to 15)? Are there any East Asia night reports before 1006 Apr 16, where no guest star is mentioned~? There is only one Chinese observation known for April 1006, namely for AD 1006 Apr 14 reporting: \begin{quotation} Emperor Zhenzong of Song, 3rd year of the Jingde reign period, 3rd month, day bingchen (53). In the north a scarlet vapour extending across the sky (and a white vapour penetrated the Moon), \end{quotation} citing from Xu et al. (2000), a slightly different translation in Keimatsu (1975), both from Songshi 60.1308, without the text in brackets also in Yau et al. (1995). This is a probable aurora according to the criteria given in Neuh\"auser \& Neuh\"auser (2015), namely northern directions, aurora-typical colour, and night-time (implicitly given with the moon).\footnote{ In the list of candidate aurorae in Hayakawa et al. (2015), this event is listed twice, once with {\em 1006 Apr 14 R[ed] V[apour] n[orth] [in] Kaifeng [moon phase] 0.46} (near full moon) and once with {\em 1006 Apr 14 W[hite] V[apour] near the moon [in] Kaifeng [moon phase] 0.46} (near full moon); one of the two texts is from the astronomical treatise ({\em Tianwen zhi}) to the {\em Songshi}, the other from its treatise on general omenology ({\em Wuxing zhi}). In the same list of candidate aurorae in Hayakawa et al. (2015), there is an additional entry: {\em 1006 May [without day] Y[ellow] V[apour] near the moon [in] Kaifeng}, also from Songshi; the translation of this entry is: {\em On this date yellow vapour like a pillar penetrated the moon}. The date for this event is uncertain; however, as it is given as the {\em guimao} (40) day in the fourth month, when in fact there was no {\em guimao} day in the fourth month; a {\em guimao} day did occur at the beginning of the fifth month (1006 May 31) and at the beginning of the third month (1006 Apr 1) -- both, however, so close to new moon that the text ({\em penetrated the moon}) does not fit to the given sexagenary date. There is another instance of the same phrase ({\em Yellow vapour like a pillar penetrated the moon}) dated to the 4th month of the 3rd year of the Tianxi reign period (AD 1019) given without the {\em guimao} date, though there is a guimao date in that month; it is possible that the event somehow got transposed to the wrong reign period; Hayakawa et al. (2015) list {\em Y[ellow] V[apour]} for 1019 May 8, but by mistake omitted here {\em near the moon}, which is clearly given in the original Chinese; they give 0.04 as moon phase (new moon May 7), so that again the text ({\em near the moon}) is not consistent with the moon phase for this date; though the entry does not actually specify which day within the lunar month the event occured on, the {\em guimao} day (40) in that lunar month was AD 1019 May 23, i.e. close to full moon (May 21/22), when a lunar halo display would be possible. In any case, this event is not an aurora, but more likely a lunar halo pillar. } What is reported as {\em a white vapour} penetrating {\em the Moon} may well be some halo effect around the Moon, which is well possible two days before full moon; for a discussion of the aurora sighting around full moon, see Chapman et al. (2015). There are indeed no additional East Asian observations known for the remaining time of AD 1006 Apr 17 until the end of April. There is evidence for the fact that relevant Chinese documents are missing: \begin{quotation} ... on the 2nd day of the 4th lunar month the Zhoubo star was seen. The official astronomer reported it immediately. The [Song] Shilu [for] the Xiangfu reign period, 9th year, 4th lunar month, [day] gengchen should be consulted for further details, \end{quotation} quoting SG02 from Xu Zizhi Tongjian Changbian -- however, the mentioned Song Shilu is lost (SG02). Furthermore, there is also evidence that Zhou Keming, a prominent astronomer in China, was on a mission in April 1006, so that he may not have been able to consult documents (for the interpretation) on SN 1006 early. In the Biography of Zhou Keming (AD 954-1017), we can read: \begin{quotation} During the 3rd year of the Jingde reign period, a large star appeared in the sky at the west of Di. No-one could determine (its significance) ... At the time, (Zhou) Keming was away on a mission to Lingnan, On his return, he urgently requested to reply ... He said: "I have checked the (astrological manuals) Tianwen Lu and the Jingzhou Zhan ... the star is known by the name Zhoubo, which is yellow in colour and really brilliant in its light. The country where it is visible will prosper greatly ..." The Emperor approved and acceded to his request. He then promoted him to the post of Librarian and Escort of the Crown Prince, \end{quotation} cited from SG02 with their additions in brackets. We can see that Zhou Keming was on a mission to Lingnan (southern China, Goldstein and Ho Peng Yoke 1965), while the guest star first appeared, that no one present could (or was allowed to) interpret its astrological meaning, and that -- upon his return -- he checked the old documents about the astrological meaning of the bright guest star, identified it as Zhoubo star, and reported his interpretation to the Emperor. Other historical documents specify that the Emperor was informed on AD 1006 May 30: \begin{quotation} [AD 1006 May 30:] The Director of the Astronomical Bureau reported that previously, on the 2nd day of the 4th lunar month [May 1], during the initial watch of the night, a large star had been seen. Its colour was yellow ... According to the star manuals, there are four categories of auspicious stars. One of them is called Zhoubo; its colour is yellow and it is really brilliant; it presages great prosperity to the state over which it appears ... The officials congratulated the Emperor, \end{quotation} from SG02 from Song Huiyao Jigao. The astronomer, who is informing the Emperor here, is also called {\em superintendent astronomer} in the same document. Furthermore, we can read: \begin{quotation} Jingde reign period, 3rd year [AD 1006-1007], there was a large star seen in the sky ... Zhou Keming, the chief official of the Spring Academy reported that according to the Tianwen Lu and the Jingzhou Zhan, the star was a Zhoubo, \end{quotation} from SG02 from Shaofu Yitang Qinghua, also quoted in Goldstein and Ho Peng Yoke (1965), who point out that the {\em Tianwen Lu} and {\em Jingzhou Zhan} are lost. That the information from China about this SN is sparse, in particular for April and May, may be due to the difficult interpretation at that time, i.e. until after the expected solar eclipse at the end of May, as mentioned in Sect. 3.6. \subsection{Possibly bad weather (monsoon) in East Asia} The East Asian monsoon affects large parts of China, Korea, and Japan; the onset of the summer monsoon with pre-monsoonal rain over South China is typically in early May, but can also start a few days or weeks earlier; the summer monsoon with many rainy phases starts in the South China Sea and then moves northward to Japan (June) and Korea (July). There are no East Asian night-time observations known at all between AD 1006 Apr 17 and the end of April. For AD 1006 Apr 14, we have evidence for a halo display in the south as seen from China (see above: {\em white vapour penetrated the Moon}). The lack of reports for the time AD 1006 Apr 17 until the end of April may be due to either bad weather or the fact that the reports were lost. \subsection{SN 1006 observed in St. Gallen for 3 months} A monk from St. Gallen, Switzerland, reported for AD 1006: \begin{quotation} Nova stella apparuit insolitae magnitudinis, aspectu fulgurans, et oculos verberans, non sine terrore. Quae mirum in modum aliquando contractior, aliquando diffusior, et iam extinguebatur interdum. Visa est autem per tres menses in intimis finibus austri, ultra omnia signa quae videntur in coelo, \end{quotation} cited after Pertz (1826) from the Annales Sangallenses maiores (covering AD 709-1056); its second part (AD 919-1056) was written by different authors, Hepidannus being one of them, a St. Gallen monk, who lived in the 2nd half of the 11th century and died AD 1088, i.e. not necessarily an eyewitness of SN 1006 himself. The above text was translated as follows: \begin{quotation} [AD] 1006. A new star of unusual size appeared, it was glittering (fulgurans) in appearance and dazzling (verberans) the eyes, causing alarm. In a wonderful manner it was contracted, sometimes spread out, and moreover sometimes extinguished. It was seen, nevertheless, for three months in the inmost limits of the south, beyond all the constellations which are seen in the sky, \end{quotation} citing Stephenson et al. (1977) and SG02 with their brackets and additions; for what they translate as {\em constellations}, the Latin has {\em signa}, which can mean {\em signs} or {\em zodiacal signs}. \begin{figure*} \begin{center} {\includegraphics[angle=270,width=17cm]{sn-1006-gallen.pdf}} \end{center} \caption{{\bf Visibility of SN 1006 from St. Gallen:} The altitude (in degrees) of (the position of) SN 1006 is plotted versus the azimuth (in degrees) for St. Gallen. The black line shows the mountain top as seen towards the south from St. Gallen monastery (700 m high) according to Stephenson et al. (1977) with Mount S\"antis as highest peak at 2503 meter. The additional curve shows the path of the position of SN 1006 as seen from St. Gallen. Since Jul 18, the position of SN 1006 would be seen after sunset only as plotted (in green) to the right of the (rightmost, green) line, i.e. not visible any more above the horizon or the mountains; since Jul 10, the position of SN 1006 is seen after sunset only as plotted (in pink or green) to the right of the (pink) line; since Jun 22, the position of SN 1006 is seen after sunset only as plotted (in blue or pink or green) to the right of the (blue) line (i.e. visible for all azimuths $\ge 4^{\circ}$); and since Apr 25, the position of SN 1006 is seen after sunset as plotted (in red or blue or pink or green), i.e. it was above the mountain (except of course behind the mountain top at azimuth $\sim 6-10^{\circ}$ for a brief period). If SN 1006 was seen in St. Gallen {\em for three months}, i.e. at least for 2.5 months, and if it was visible last around Jul 10 (for about one minute after sunset) or earlier, then the first observation should have been in April. Given that the observer described the star to be {\em sometimes extinguished}, he must have been at an altitude such that the star, within the night, was sometimes seen above the mountain and sometimes being briefly occulted by the mountain top, so that the star was seen between $\sim 4^{\circ}$ to $\sim 5^{\circ}$ above a perfect flat horizon (given the height of Mount S\"antis). Therefore, the altitude of the observer was somewhere between 700 m and $\sim 1100$ m. } \end{figure*} The relevant part about the length of the observation of SN 1006 is {\em Visa est autem per tres menses}, i.e. {\em per tres menses}, which clearly means {\em for three months} or {\em throughout three months}, e.g. from the beginning of some month 1 (not neccessarily a calendar month) until the end of month 3. With {\em for three months}, the author of this part of the St. Gallen annales did not necessarily mean {\em three full months}, he may have rounded down or up (i.e. 2.5 to 3.5 months). The wording clearly does not mean {\em in three (different, subsequent calendar) months}, which could then have meant that it was observed first at the end of month 1 (e.g. May) and last at the beginning of month 3 (e.g. July). It is noteworthy to mention that the St. Gallen chronicle does not mention the duration of visibility just in passing, but explicitely ({\em nevertheless, for three months}), in spite of the difficult conditions (high mountains, strong extinction). The southern horizon as seen from St. Gallen has high mountains with Mount S\"antis straight towards the south being the highest one with 2503 m, located 20 km south of St. Gallen. The monastery is at an elevation of some 700 m, but the monks may have observed from a slightly higher point nearby; e.g. somewhat closer to the mountain. The highest point in today's St. Gallen is 1074 m. The summit of Mount S\"antis as seen from either the monastery or the higher point is only $\sim 4^{\circ}$ to $\sim 5^{\circ}$ above a perfectly flat mountain-less horizon. While the monk may in principle have observed from an even higher point, since the location itself is not specified in the text, the range in degrees given above, i.e. only about one degree fron $\sim 4^{\circ}$ to $\sim 5^{\circ}$ above horizon, must indeed be as small as given: the text specifies that {\em it} [SN 1006] ... {\em moreover sometimes extinguished.} The observer tells us that SN 1006 was sometimes seen and sometimes {\em extinguished}, which is well possible given the mountain range, where higher parts sometimes block the star light. Hence, this statement limits the range in altitude and, therefore, also the range in the height of the observing location, whereever it was (even if outside the monastery). See Fig. 2. If the star observed in St. Gallen was indeed SN 1006, then SN 1006 ($\delta _{1006} = -37^{\circ} 34^{\prime}$) was only up to $5^{\circ}$ above a perfectly flat (mountain-less) horizon at its location ($47^{\circ}25^{\prime}$ north). However, the horizon was furthermore limited by mountains (SG02): At an eastern azimuth, the true horizon due to mountains barely allowed celestial observations below $4^{\circ}$ above the perfect flat horizon, while at a western azimuth of $\ge 10^{\circ}$, celestial objects $\le 3^{\circ}$ above the perfect flat horizon were visible (Fig. 2). (If the observer went to a place higher up than the monastery, then SN 1006 could be seen a bit better and maybe a bit longer, but one criterion of seeing it for the last time only in the very last minute after sunset is already very hard.) We can now estimate the time of the year when SN 1006 was visible from St. Gallen above the mountains, see Fig. 2. Let us first estimate the last observing date: Given that there were no day-time observations reported for SN 1006 (except the report from Morocco: {\em its first appearance was before sunset}), we can assume that SN 1006 was visible only after local sunset in St. Gallen. Due to its very low altitude and strong atmospheric extinction as seen from St. Gallen, a day-time observaion of SN 1006 from St. Gallen is much less likely than from any other place, where SN 1006 was seen. For SN 1006 being $3^{\circ}$ above the perfect flat horizon (but less than 1 degree above the mountains), it was last visible from St. Gallen on AD 1006 July 10 at an extinction corrected apparent magnitude of about $-1 \pm 1$ mag; and on June 22 with $\sim -2 \pm 1$ mag for about one minute after sunset for $5^{\circ}$ above the perfectly flat horizon (but less than one degree above the mountain); see Fig. 2; we have neglected refraction here, which would amount to less than $1^{\circ}$. If it was last visible on or around July 10 (or earlier), when was it first sighted? As mentioned above, with {\em for three months}, the author(s) of the St. Gallen annales means at least 2.5 months, namely until July 10 (or earlier), see above. Then, he would have started to have seen SN 1006 on or around April 25 or earlier. On and around April 25, SN 1006 would have been $\ge 3^{\circ}$ above a perfectly flat St. Gallen horizon (and about one degree above the mountains) for quite some time after sunset. There is no regular weather pattern (like a monsoon in East Asia) in St. Gallen and central Europe as a whole in late April or early May (except maybe that the weather changes a lot in Europe in April), it could have been clear on many evenings. That not many other observers have noticed SN 1006 in Europe can be due to its extreme southerly declination, so that only very experienced and educated scholars (like monks) would detect it. Also the Annales Beneventani (southern Italy, $6^{\circ}$ south of St. Gallen) report about a new bright star in 1006 and use the wording {\em per tres menses} ({\em for three months}): \\ AD 1006: {\em Clarissima stella effulsit, et siccitas magna per tres menses fuit}, \\ which we translate as follows: \\ {\em A very brilliant star shone, and a large drought happened for three months}, (also given in SG02).\footnote{Additional European sightings do not mention the date or length of the observations (SG02).} \\ It is quite likely that the two items reported, a new bright star and a three-month drought, are meant to be connected. Given that this observer located $6^{\circ}$ south of St. Gallen and that he does not have high mountains towards the south, he should have been able to observe SN 1006 for longer than in St. Gallen -- and indeed, while the St. Gallen report ({\em for three months}) can mean at least 2.5 months, the Beneventani report ({\em for three months} can mean up to 3.5 months, both may be rounded, they are not inconsistent with slightly different time spans. A few more European annals mention a {\em cometes} for AD 1006, namely Li\'ege and Lobbes, Belgium, also Venice, Italy, as well as Metz and Mousson, France (SG02). Some of them are further north than St. Gallen, but probably just report what they heard from St. Gallen. That they use the Latin word {\em cometes} (usually translated as {\em comet}) should not worry us, because at that time it meant {\em transient celestial object}, like the Arabic {\em nayzak}. The annals from St. Gallen and Benevento, though, do not use the word {\em cometes} for 1006 indicating that the observers there noticed that this transient celestial object was different from what we today call a {\em comet} -- indeed, a new/very brilliant star. | 16 | 7 | 1607.02915 |
1607 | 1607.06916_arXiv.txt | s{We present a systematic study of modified gravity (MG) models containing a single scalar field non-minimally coupled to the metric. Despite a large parameter space, exploiting the effective field theory of dark energy (EFT of DE) formulation and imposing simple physical constraints such as stability conditions and (sub-)luminal propagation of perturbations, we arrive at a number of generic predictions about the large scale structures.} The goal of this work, in collaboration with F. Piazza, C. Marinoni and L. Hui \cite{Perenon:2015sla}, is to study the predictability of MG theories aiming at challenging the standard $\Lambda$CDM explanation of cosmic acceleration. We use the EFT of DE \cite{EFTOr} for it has established a common formalism to describe the widest set of MG theories, those adding a single extra scalar degree of freedom to the Einstein-Hilbert action. \vv Such a unifying description enables MG theories to be parametrized in a common framework in terms of structural functions of time, describing how matter perturbations evolve in the universe. The requirements needed to fully describe an EFT model can be reduced to two constants and three functions of time $ \left\{\omo,\; w, \; \mu(t), \; \mu_3(t),\; \epsilon_4(t) \right\}$. The three functions are non-minimal couplings, once ``turned on" they enable the description theories in the Horndeski class: $\mu$ is the Brans-Dicke (\emph{BD}) non-minimal coupling, adding $\mu_3$ models the cubic galileon (\emph{H3}) term and $\epsilon_4$ encodes the 4 and 5 Horndeski (\emph{H45}) Lagrangians. In parallel, the EFT of DE allows one to independently set the background expansion, $i.e$ the Hubble rate. This reduces to fixing the two constants, the present fractional matter density of non-relativistic matter in the perfect fluid approximation $\omo$ and the background DE equation of state parameter $w$ \cite{pheno}. They are set by the latest constraints \cite{planck} to reproduce a flat $\Lambda$CDM background, thus respectively $\sim 0.3$ and $-1$. \vv Another asset of the EFT of DE is to provide a clear and common means of assessing if theories are pathological or not, $i.e$ whether they suffer from ghost or gradient instabilities. The main purposes of this work is to show that despite large degrees of freedom, definite features common to all healthy ---stable and with no superluminal propagation of scalar and tensor modes--- EFT models arise within the vast Horndeski class. We show this by expanding the non-minimal coupling functions in power series up to second order in the reduced matter density, $ x $, used as time variable. An overall $(1-x)$ pre-factor ensures the vanishing of the couplings at early times, hence recovering general relativity. We randomly generate the coefficients of the expansions until we obtain $10^4$ healthy \emph{BD}, \emph{H3} and \emph{H45} theories. \vv At the linear level and under certain conditions, non-standard gravitational scenarios result in time dependent modifications of the Newton's constant $\geff$ and of the gravitational slip parameter $\gsp$. The requirement of a healthy theory leads to bounded evolutions and generic features in these quantities. For instance, as is it shown in fig.~\ref{fig}, the curves of $\geff/\gn$ display a characteristic {\it S-shape}, an alternation, stronger gravity at early times, weaker gravity at intermediate redshifs and stronger gravity again today. Two competing effects are at play: the matter density entering the perturbations of an EFT model is not the same as that of the background $x$ and being generally smaller it lowers the gravitational coupling; the contribution of the extra scalar field is attractive (healthy spin-0 field) which enhances the gravitational coupling. The effective Newton constant acts as a source term in the growth of matter perturbations. An observable of the large scale structures such as the growth function $\fs$ can be effectively used to constrain MG. As it is shown in fig.~\ref{fig}, its sensitivity to $\geff$ translates into having almost all models displaying lower growth than $\Lambda$CDM at $0.5\lesssim z\lesssim 1$ and always predicting stronger growth for $z>1.5$. The \emph{H3} and \emph{H45} cases, exhibiting more freedom, can introduce a few models deviating strongly from the $\Lambda$CDM reference however they are statistically insignificant. Complementing this analysis with $\gsp$ highlights more definite features and enhances the discriminating power. For \emph{BD} theories, one can show analytically that $\gsp \leqslant 1$ always holds. When going to \emph{H3} and \emph{H45}, a distinctive shape is still seen, $\eta$ is always smaller than unity at any redshifts except in the window $0.5\lesssim z\lesssim 1$. \vv We are now exploring, Perenon $et$ $al.$ in prep., the impact on the universality of our results when changing the background DE equation of state and changing the asymptotic behaviour of the non-minimal couplings (early dark energy). We are also investigating the link between self-acceleration and the production of weaker gravity. \begin{figure}[!] \begin{center}\vskip-2mm \includegraphics[scale=0.42]{PLOT_BD_geffgn.pdf} \hskip-4mm \includegraphics[scale=0.42]{PLOT_H3_geffgn.pdf} \hskip-4mm \includegraphics[scale=0.42]{PLOT_H45_geffgn.pdf} \vskip-8mm \includegraphics[scale=0.42]{PLOT_BD_gsp.pdf} \hskip-4mm \includegraphics[scale=0.42]{PLOT_H3_gsp.pdf} \hskip-4mm \includegraphics[scale=0.42]{PLOT_H45_gsp.pdf} \vskip-4mm \includegraphics[scale=0.42]{PLOT_BD_fsig8.pdf} \hskip-4mm \includegraphics[scale=0.42]{PLOT_H3_fsig8.pdf} \hskip-4mm \includegraphics[scale=0.42]{PLOT_H45_fsig8.pdf}\vskip-2mm \caption{The behaviour of $\geff/\gn$ (first row), $\gsp$ (second row) as a function of the reduced matter density $x$ and the redshift evolution of $\fs$ (third row) are shown for $10^4$ randomly generated viable EFT models (\emph{BD}, \emph{H3}, \emph{H45}). The dotted vertical lines identify, from left to right, the cosmic epochs $z=0.5$, $z=1$ and $z=2$. The thick red line represents the $\Lambda$CDM prediction. The density of curves passing through a region is shown by the levels of blue.}\vskip-2mm\vskip-2mm \label{fig} \end{center} \end{figure}\vskip-4mm\vskip-4mm | 16 | 7 | 1607.06916 |
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1607 | 1607.07910_arXiv.txt | We investigate the expected cosmological constraints from a combination of weak lensing and large-scale galaxy clustering using realistic redshift distributions. Introducing a systematic bias in the weak lensing redshift distributions (of 0.05 in redshift) produces a $>2\sigma$ bias in the recovered matter power spectrum amplitude and dark energy equation of state, for preliminary Stage III surveys. We demonstrate that these cosmological errors can be largely removed by marginalising over unknown biases in the assumed weak lensing redshift distributions, if we assume high quality redshift information for the galaxy clustering sample. Furthermore the cosmological constraining power is mostly retained despite removing much of the information on the weak lensing redshift distribution biases. We show that this comes from complementary degeneracy directions between cosmic shear and the combination of galaxy clustering with cross-correlation between shear and galaxy number density. Finally we examine how the self-calibration performs when the assumed distributions differ from the true distributions by more than a simple uniform bias. We find that the effectiveness of this self-calibration method will depend on the details of a given experiment and the nature of the uncertainties on the estimated redshift distributions. | Cosmic shear is potentially the most powerful tool available to cosmologists today. As an unbiased probe of the mass distribution, it offers powerful constraints on the mean density of the Universe and the clustering of dark matter. It is also expected to shed new light on the late-time accelerated expansion of the Universe and thus measure the dark energy equation of state and test General Relativity on the largest scales. \begin{figure} \centering \includegraphics[width=9cm]{results_plots/letter/observable_cls_svbinning_weighted.png} \caption{The components of the fiducial datavector used in this Letter. Shown are angular power spectra of cosmic shear (purple solid), galaxy clustering (red dotted) and the shear-density cross-correlation (green dot dashed). Each panel corresponds to a unique redshift bin pairing. In the panels where it is not visible, the $\delta _g \delta _g$ spectrum is below the range shown. We note that all values shown are positive, apart from $C^{1,3}_{\gamma \delta _g}$ (upper right), which becomes negative and is below the lowest point on this scale for $\ell<900 $.} \label{fig:cls} \end{figure} A three decade programme aiming to extract unprecedented constraints on our cosmological model from cosmic shear is now midway to completion. It began soon after the first detection in 2000 \citep{bacon00,vanWaerbeke00,wittman00,kaiser_detection} using $\sim$10000 galaxies and is expected to culminate in catalogues of more than a billion galaxies by the end of the coming decade (Stage IV, \citet{albrecht06}). Logarithmically, we are halfway there, with ongoing analyses of the preliminary Stage III datasets, containing $\sim$10 million galaxies (\citealt{dessv2pt15,kids16}, see also \citealt{heymans13,jee16}). The increase in the number of galaxies with reliable shape measurements has allowed tighter cosmology constraints, but also requires better control of systematic biases. In this Letter we focus on a potential Achilles' heel of galaxy imaging surveys for cosmology: the use of photometric redshifts to estimate distances to galaxies. Tomographic cosmic shear analyses bring a number of benefits \citep{hu99}, but place stringent requirements on our knowledge of galaxy redshift distributions. \citet{amara07,abdalla08,jouvel09} and \citet{ishak06} present detailed studies of the requirements for spectroscopic follow up of Stage IV cosmology surveys, while \citet{ma06,huterer06, bernstein09} offer numerical forecasts of cosmological impact from photometric redshift (\pz) biases. Many others (e.g. \citealt{bordoloi12,cunha_huterer12}) present detailed studies of specific \pz systematics, albeit with less focus on the ultimate cosmological impact. Tightening systematics requirements have sparked interest in spatial cross-correlations between photometric and spectroscopic galaxies within the survey volume, as a method for externally calibrating photometric redshifts \citep{newman08,menard13,dePutter14,rahman15,choi15,scottez16}. \begin{figure} \centering \includegraphics[width=1\columnwidth]{results_plots/letter/mock_sv_pz_last.png} \caption{ Redshift distributions considered in this paper. The upper panels show the shear catalogue redshift distributions used in this work, taken from DES SV \citep{bonnettsv15}: \blockfont{skynet} (solid purple; fiducial), \blockfont{skynet} with a bias of 0.05 (dashed green) and \blockfont{bpz} (dotted blue) without the shift of 0.05 in redshift used in \citet{bonnettsv15}. The lower panel displays the the fiducial galaxy clustering catalogue in three bins taken from~\citet{clampitt16} which uses~\citet{rozo15} (DES SV redMaGiC). } \label{fig:nofzs} \end{figure} \begin{figure*} \centering (a) \includegraphics[width=1\columnwidth]{results_plots/letter/sig8omm_constraints_skynet-skynet_biased.png} (b) \includegraphics[width=1\columnwidth]{results_plots/letter/w0S8_constraints_skynet-skynet_biased.png} \caption{ (a) Forecast constraints on the matter density and clustering amplitude in $\Lambda$CDM and (b) dark energy equation of state in $w$CDM for various assumptions about photometric redshifts (for reference constraints from Planck 2015 temperature and low frequency polarisation data alone are shown by the dot-dashed red contours \citep{planck2015cosmo}). The colours in each panel indicate three \pz scenarios. In green are the result of using the \blockfont{skynet} $n(z)$ in the theory calculation data and fixing $\delta z=0$. We show this as an unrepresentative ideal case, where the photometric estimates provide a perfect representation of the true galaxy distribution. Overlain are the results of using the \blockfont{skynet} $n(z)$, biased downwards by 0.05 in redshift under the (erroneous) assumption of no bias (blue dotted) and varying three additional $\delta z^i$ nuisance parameters marginalised with a wide Gaussian prior of width $\Delta \delta z =0.1$(purple solid). The true input cosmology is shown by the black cross. } \label{fig:biases} \end{figure*} Given the limited amount of spectroscopic information available, several authors have speculated about the possibility of calibrating cosmic shear redshift distributions from the imaging survey itself. \citet{huterer06} show that cosmic shear alone affords a limited capacity for self-calibration. \citet{schneider06} and \citet{sun15} investigate the photometric redshift calibration information available from Stage IV galaxy clustering. \citet{zhan06} explore the constraining power on $w_0$ using a similar technique and combine with cosmic shear constraints. \citet{zhang10} point out that shear-density cross-correlations (cross-correlation between shear and galaxy counts, also referred to as tangential shear or galaxy-galaxy lensing) can help to constrain \pz error, when combined with galaxy clustering. All the studies mentioned in the previous paragraph make a crucial assumption, which is unlikely to be realised in practice: that the galaxies used for cosmic shear have a systematics-correctable galaxy clustering signal. In practice regions of the sky with better (worse) seeing conditions are likely to contain a higher (lower) number density of galaxies usable for cosmic shear \citep[e.g. see Appendix C of][for more discussion]{choi15}. Therefore there will be a large spurious clustering signal from the galaxies selected for a shear catalogue, rendering standard galaxy clustering analyses useless. Thus in practice we will usually have a different galaxy sample selection for the weak lensing sample and the galaxy clustering sample. This is standard in galaxy-galaxy lensing analyses and was done in the first combined analyses of cosmic shear and large scale structure on data (\citealt{nicola16}, who also combined with the CMB), and was considered for Stage IV surveys with much tighter priors in the forecasts of \citet{cosmolike16}. This means one has twice as many redshift distributions to understand as in a shear-only analysis. However, this also offers an opportunity: we can choose to use a galaxy clustering sample with much better understood photometric redshift properties, which can in turn help to calibrate the redshift distribution of the weak lensing sample. In this Letter we explore the potential for simultaneously constraining \pz error and cosmology using cosmic shear, galaxy clustering and shear-density cross-correlations. Unlike previous studies using this data combination, we consider a scenario in which the redshift distribution of the shear catalogue and galaxy clustering catalogues differ significantly. We assume the galaxy clustering sample is highly homogeneous and dominated by luminous red galaxies, which tend to yield high quality \pz. This Letter is structured as follows. Section 2 outlines the setup of our analysis with a description of the simulated data vectors, redshift distributions and the photometric uncertainties considered. In Section 3 we investigate the power of cosmic shear, galaxy clustering and shear-density cross-correlations to internally constrain \pz biases. Finally a series of robustness tests are presented to explore the limits of this effect. We adopt a fiducial flat $\Lambda$CDM cosmology with $\sigma_8=0.82$, $\Omega_{\rm m} = 0.32$, $h=0.67$, $w_0 =-1$, $\Omega_{\rm b} = 0.049$. | We have investigated the potential for current cosmological galaxy imaging surveys to self-calibrate photometric redshift distribution uncertainties, for the first time considering the realistic case in which the weak lensing sample is different from the galaxy clustering sample, and has substantial calibration uncertainties. % We focus on a preliminary Stage III dataset with $\sim 40$M galaxies, in which the galaxy clustering sample has well-understood photometric redshifts ($\Delta \delta z=0.01$). We find that the combination of cosmic shear, galaxy clustering and shear-density cross-correlations is much more robust to errors and uncertainties in the redshift distribution calibration than cosmic shear alone. Specifically, the uncertainty on the clustering amplitude parameter $S_8 \equiv \sigma_8 (\Omega_{\rm m}/0.31)^{0.5}$ is increased by only $40\%$ on marginalising over three free independent bias parameters with a prior width $\Delta \delta z=0.1$, relative to the case $\Delta \delta z=0$. This contrasts with more than a factor of two degradation for cosmic shear alone. We illustrate that this is because cosmic shear constrains a different degenerate combination of cosmology and photometric redshift calibration parameters than galaxy clustering and shear-density cross-correlations. We find that the combination of all three two-point functions can correct even a substantial bias (of 0.05) in the $n(z)$ to accurately recover the input cosmology. This result is robust to allowing for a basic stochastic bias model, and strengthened by using less conservative cuts on the scales used in the galaxy clustering analysis. The self-calibration result disappears if the redshift distributions are less well understood for the galaxy clustering sample ($\Delta \delta z=0.05$). Using an alternative redshift distribution estimate (\blockfont{bpz}) we demonstrate that this result may change if the deviation of the redshift distribution from the truth is not fully captured by a uniform translation, if the only redshift uncertainty considered is a uniform translation in redshift. The validity of our findings should be verified on a case-by-case basis for specific realisations of the photo-z error This investigation advances on most previous numerical forecasts in implementing MCMC sampling rather than Fisher analyses, and assumes a low-density galaxy clustering sample with relatively well-known redshifts. We do, however, assume Gaussian covariance matrices, which tend to underestimate the uncertainties for cosmic shear and could thus make our forecasts over-optimistic. Investigation of non-Gaussian covariances is beyond the scope of this Letter. We have also assumed the Limber approximation is correct for the range of scales used, and ignored redshift-space distortions. The results suggest that self-calibration may be a practical solution for current cosmological surveys, assuming reliable \pz estimates can be obtained for the galaxy clustering catalogue, if the weak lensing redshift distributions cannot be easily calibrated through a different route. | 16 | 7 | 1607.07910 |
1607 | 1607.07317_arXiv.txt | { Models of galaxy formation in a cosmological framework need to be tested against observational constraints, such as the average stellar density profiles (and their dispersion) as a function of fundamental galaxy properties (e.g. the total stellar mass). Simulation models predict that the torques produced by stellar bars efficiently redistribute the stellar and gaseous material inside the disk, pushing it outwards or inwards depending on whether it is beyond or inside the bar corotation resonance radius. Bars themselves are expected to evolve, getting longer and narrower as they trap particles from the disk and slow down their rotation speed. } {We use 3.6~$\mu$m photometry from the Spitzer Survey of Stellar Structure in Galaxies (S$^{4}$G) to trace the stellar distribution in nearby disk galaxies ($z\approx0$) with total stellar masses $10^{8.5}\lesssim M_{\ast}/M_{\odot}\lesssim10^{11}$ and mid-IR Hubble types $-3 \le T \le 10$. We characterize the stellar density profiles ($\Sigma_{\ast}$), the stellar contribution to the rotation curves ($V_{3.6 \mu \rm m}$), and the $m=2$ Fourier amplitudes ($A_2$) as a function of $M_{\ast}$ and $T$. We also describe the typical shapes and strengths of stellar bars in the S$^4$G sample and link their properties to the total stellar mass and morphology of their host galaxy. } { For 1154 S$^4$G galaxies with disk inclinations lower than $65^{\circ}$, we perform a Fourier decomposition and rescale their images to a common frame determined by the size in physical units, by their disk scalelength, and for 748 barred galaxies by both the length and orientation of their bars. We stack the resized density profiles and images to obtain statistically representative average stellar disks and bars in bins of $M_{\ast}$ and $T$. Based on the radial force profiles of individual galaxies we calculate the mean stellar contribution to the circular velocity. We also calculate average $A_{2}$ profiles, where the radius is normalized to $R_{25.5}$. Furthermore, we infer the gravitational potentials from the synthetic bars to obtain the tangential-to-radial force ratio ($Q_{\rm T}$) and $A_2$ profiles in the different bins. We also apply ellipse fitting to quantitatively characterize the shape of the bar stacks. } { For $M_{\ast} \ge 10^{9}M_{\odot}$, we find a significant difference in the stellar density profiles of barred and non-barred systems: (i) disks in barred galaxies show larger scalelengths ($h_{\rm R}$) and fainter extrapolated central surface brightnesses ($\Sigma_{\circ}$), (ii) the mean surface brightness profiles ($\Sigma_{\ast}$) of barred and non-barred galaxies intersect each other slightly beyond the mean bar length, most likely at the bar corotation, and (iii) the central mass concentration of barred galaxies is higher (by almost a factor~2 when $T\le5$) than in their non-barred counterparts. The averaged $\Sigma_{\ast}$ profiles follow an exponential slope down to at least $\sim 10 M_{\odot} \rm pc^{-2}$, which is the typical depth beyond which the sample coverage in the radial direction starts to drop. Central mass concentrations in massive systems ($\ge 10^{10}M_{\odot}$) are substantially larger than in fainter galaxies, and their prominence scales with $T$. This segregation also manifests in the inner slope of the mean stellar component of the circular velocity: lenticular (S0) galaxies present the most sharply rising $V_{3.6 \mu \rm m}$. Based on the analysis of bar stacks, we show that early- and intermediate-type spirals ($0 \le T < 5$) have intrinsically narrower bars than later types and S0s, whose bars are oval-shaped. We show a clear agreement between galaxy family and quantitative estimates of bar strength. In early- and intermediate-type spirals, $A_{2}$ is larger within and beyond the typical bar region among barred galaxies than in the non-barred subsample. Strongly barred systems also tend to have larger $A_{2}$ amplitudes at all radii than their weakly barred counterparts. } { Using near-IR wavelengths (S$^{4}$G 3.6~$\mu$m), we provide observational constraints that galaxy formation models can be checked against. In particular, we calculate the mean stellar density profiles, and the disk(+bulge) component of the rotation curve (and their dispersion) in bins of $M_{\ast}$ and $T$. We find evidence for bar-induced secular evolution of disk galaxies in terms of disk spreading and enhanced central mass concentration. We also obtain average bars (2-D), and we show that bars hosted by early-type galaxies are more centrally concentrated and have larger density amplitudes than their late-type counterparts. } | In the lambda cold dark matter ($\Lambda$CDM) model, the seeds for dark matter halos arise from quantum fluctuations amplified by cosmic inflation. The halos gain angular momentum from cosmological torques as they grow. Galaxies are formed from the cooling and condensation of gas in their centres \citep[][]{1978MNRAS.183..341W,1980MNRAS.193..189F}, and the baryons inherit the angular momemtum from their host halos. In the early universe galaxies suffered recurrent mergers, which became less frequent with time. At $z\sim 2-2.5$ galaxies were still thick, gas-rich, and clumpy, and actively formed stars \citep[e.g.][]{2007ApJ...670..237B,2008ApJ...687...59G,2009ApJ...703..785D,2011ApJ...739...45F}. \citet[][]{2013ApJ...771L..35V} found that in galaxies with present-day stellar masses similar to that of the Milky Way (MW, log$_{10} (M_{\ast}/M_{\odot}) \approx 10.7$) most of the star formation had already taken place before $z \sim 1$, at which time these systems typically show an almost fully assembled backbone with a quiescent bulge\footnote{Throughout this paper, and unless stated otherwise, the term `bulge' refers to the excess of flux above the disk in the central regions of galaxies, independent of the stellar structures that emit this light.} and a slowly star-forming disk. This is consistent with studies showing a lack of strong evolution in the stellar mass-size relation in disk galaxies over the last 8 billion years \citep[e.g.][]{1998ApJ...500...75L,1999ApJ...519..563S,2005ApJ...635..959B}. In that period the bar fraction \citep[e.g.][]{2004ApJ...615L.105J,2004ApJ...612..191E} or the locii of outer rings relative to bars \citep[e.g.][]{2012A&A...540A.103P} have been found to be constant. However, \citet[][]{2008ApJ...675.1141S} found a constant bar fraction up to $z\sim 0.84$ only for the most massive spirals ($M_{\ast}\gtrsim10^{11}M_{\odot}$), whereas for fainter and bluer systems it declined substantially beyond $z\sim 0.3$. Recent work by \citet[][]{2015MNRAS.451....2S} concludes that disk galaxies have experienced a substantial peripheral growth since $z\sim1$ \citep[based on measurements of the Petrosian radius;][]{1976ApJ...209L...1P}. Disks in spiral galaxies exhibit luminosity profiles that tend to decay exponentially in the radial direction \citep[][]{1970ApJ...160..811F}. Exponential or quasi-exponential stellar disks have been produced in a $\Lambda$CDM framework in simulation models \citep[e.g.][]{2003ApJ...597...21A,2004ApJ...606...32R,2009MNRAS.396..121D}. Recent observations in optical \citep[e.g.][]{2005ApJ...626L..81E,2006A&A...454..759P,2008AJ....135...20E} and infrared wavelengths \citep[e.g.][]{2013ApJ...771...59M,2014MNRAS.441.1992L,2014ApJ...782...64K} show that disks can have breaks in their radial surface brightness profiles. Specifically, disks are classified as Type I if they are purely exponential, and Type II and Type III respectively if they become steeper or shallower after the break. Roughly two-thirds of the local galaxies have a bar \citep[e.g.][]{1991rc3..book.....D,2000ApJ...529...93K,2002MNRAS.336.1281W,2004ApJ...607..103L}. Stellar bars are known to conduct angular momentum of the baryonic and dark matter components throughout the disks of spiral galaxies \citep[e.g.][]{1985MNRAS.213..451W, 2002MNRAS.330...35A,2007ApJ...659.1176M}. Namely, angular momentum is emitted from the material in the surroundings of the inner Lindblad resonance of the bar, and absorbed by the material near the resonances associated with the spheroidal components (dark matter halo and bulge, when present) and with the outer disk. According to different dynamical models \citep[e.g.][]{1991MNRAS.250..161L,1998ApJ...493L...5D,2002MNRAS.330...35A,2003MNRAS.341.1179A,2006ApJ...637..214M}, while bars exchange angular momentum they become narrower, longer, stronger, and they slow down their rotation speed. Early analytical studies by \citet[][]{1980A&A....81..198C} proved that the orbits making up the bar do not extend beyond the corotation resonance radius ($r_{\rm cr}$). The bar slow-down predicted by N-body simulations is not easy to reconcile with bar pattern speed estimates from observations. \citet[][]{2005ApJ...631L.129R,2008MNRAS.388.1803R} showed that bars in early-type galaxies have typical ratios $r_{\rm cr}/r_{\rm bar}\le1.4$ (known as fast bars), while later-types have bars which can be both fast and slow. Recent measurements using the Tremaine-Weinberg method \citep[][]{1984MNRAS.209..729T} indicate fast bars for all morphological types \citep[][]{2015A&A...576A.102A}. Models by \citet[][]{1992MNRAS.259..345A} showed that $r_{\rm cr}/r_{\rm bar}$ determines the shape of the offset dust lanes. Among fast bars, the curvature of dust lanes was found to inversely scale with the strength of the bar in the recent work by \citet[][]{2015MNRAS.450.2670S}, while slow bars had similar values of the mean curvature for all bar strength values. This confirmed the theoretical prediction in \citet[][]{1992MNRAS.259..345A} \citep[see also][]{2002MNRAS.337..808K,2009ApJ...706L.256C}. Possible observational evidence for the bar evolution was provided in \citet[][]{2007MNRAS.381..401L}, \citet[][]{2007ApJ...670L..97E}, \citet[][]{2011MNRAS.415.3308G}, and \citet[][hereafter DG2016]{2016A&A...587A.160D}, by finding a dependence between different proxies of the bar strength and estimates of their length, and by studying the evolution of the bar parameters in the Hubble sequence. The bar evolution is also manifested in the buckling instability in the vertical direction that gives birth to boxy/peanut (B/P) bulges \citep[][]{1981A&A....96..164C,1990A&A...233...82C,1991Natur.352..411R,2002MNRAS.330...35A,2004ApJ...604L..93D,2004ApJ...613L..29M,2006ApJ...645..209D}. The so-called barlenses \citep[][]{2011MNRAS.418.1452L} are thought to be boxy/peanut bulges seen in the face-on view \citep[][]{2007MNRAS.381..401L,2014MNRAS.444L..80L,2015MNRAS.454.3843A}. For a review of the properties of B/P bulges, the reader is referred to \citet[][]{2016ASSL..418...77L} and \citet[][]{2016ASSL..418..391A}. Bars participate in the redistribution of stars and gas inside the disk \citep[][and references therein]{2013seg..book..305A} by pushing them outwards (inwards) beyond (within) the corotation radius (CR). Simulation models predict that this process increases the disk size \citep[e.g.][]{1971ApJ...168..343H} and causes secular evolution of bulges \citep[e.g.][]{1992MNRAS.259..328A,1992MNRAS.258...82W,1993A&A...268...65F} after the bar-funneled cold gas is transformed into stars, and also as a result of old stars being driven inward by the bar torques \citep[][]{2004A&A...423..849G}. Actually, signatures of disk-like central components in the kinematics of observed barred galaxies have been found \citep[e.g.][]{2009A&A...495..775P,2015MNRAS.451..936S}. In addition, an enhanced star formation in the central parts of barred galaxies has been detected \citep[e.g.][]{2011MNRAS.416.2182E,2012ApJS..198....4O,2015AJ....149....1Z,2015A&A...584A..88F}. Based on studies of the bar fraction in the Galaxy Zoo \citep[][]{2011MNRAS.411.2026M} and measurements of bar sizes, \citet{2013ApJ...779..162C} provided empirical evidence for bar-driven secular evolution within the central kpc of disk galaxies. Furthermore, bars are known to be responsible for the formation of resonance rings \citep[e.g.][]{1981ApJ...247...77S,1986ApJS...61..609B,2000A&A...362..465R}. Whether stellar bars drive spiral density waves has been a matter of debate \citep[e.g.][]{1979ApJ...233..539K,2003MNRAS.342....1S,2004AJ....128..183B,2005AJ....130..506B,2009AJ....137.4487B,2009MNRAS.397.1756D}. Recent work by \citet{2010ApJ...715L..56S} supports such causality. Using SDSS-DR2 photometry, observational evidence for the bar-induced disk secular evolution was provided by \citet[][]{2013MNRAS.432L..56S}, who found that barred galaxies with stellar masses $M_{\ast}>10^{10}M_{\odot}$ at redshifts $0.02 \le z \le 0.07$ typically have fainter extrapolated central surface brightness and larger disk scalelengths than their non-barred counterparts. Based on IFU observations of three prominent galactic bulges and full spectral fitting methods, recent work by \citet[][]{2015MNRAS.446.2837S} found that at least $50\%$ of the stars in those bulges formed at $z\sim4$. They also detected a younger component ($\sim 1-8$ Gyr). Two bulge population families were found in the models of MW-type galaxies in a cosmological framework by \citet[][]{2013ApJ...763...26O}, the first forming during an early starburst-collapse (old stars) and the second during the phase driven by processes such as disk instabilities and/or mergers (young stars). In order to obtain the observational constraints needed to check galaxy formation models, \citet[][]{2009MNRAS.396..121D} proposed measuring the average deprojected surface brightness profiles as a function of the primary galaxy parameters, such as stellar mass, colour, and size. Mid-IR rest-frame wavelengths are well suited to this as they are very sensitive to the old stellar populations and are barely affected by dust absorption. For this reason, the Spitzer Survey of Stellar Structure in Galaxies \citep[S$^{4}$G;][]{2010PASP..122.1397S}, which includes 2352 nearby galaxies, is an ideal sample to carry out such a study at $z\approx0$. Using S$^{4}$G 3.6~$\mu$m imaging, DG2016 provided a first-order estimate of the halo-to-stellar mass ratio ($M_{\rm halo}/M_{\ast}$) from comparisons of the stellar component of the circular velocity to kinematic H{\sc\,i} data compilation from \citet[][]{2009AJ....138.1938C,2011MNRAS.414.2005C} and HyperLEDA\footnote{We acknowledge the usage of the database (http://leda.univ-lyon1.fr).} \citep{2003A&A...412...45P}. \citet[][]{2016A&A...587A.160D} found a good agreement in the slope of the $M_{\rm halo}/M_{\ast}$-$M_{\ast}$ relation with the best-fit model at $z\approx0$ in $\Lambda$CDM cosmological simulations \citep[e.g.][]{2010ApJ...710..903M}. Based on various bar measurements, DG2016 carried out a statistical analysis of the properties of bars and their fraction at $z=0$ as seen in the S$^{4}$G sample, providing possible evidence for the growth of galactic bars within a Hubble time. In the current paper, the characterization of galactic bars is done with an inverted approach: first we stack images of individual galaxies to obtain average bar density distributions as a function of stellar mass, revised Hubble type and galaxy family; and then study the properties of these stacked bar images. This paper is organized as follows. In Sect.~\ref{data1} we present the S$^{4}$G data and the criteria used to define the samples. In Sect.~\ref{scaling-bars} we describe the methodology for resizing and stacking the galaxy images and 1-D profiles. In Sect.~\ref{disk-char} we study the luminosity profiles of the stellar disks based on the stacks and we study the mean $A_2$ profiles and stellar component of the circular velocity. In Sect.~\ref{bars-char} we analyse the shape and strength of the average bars. In Sect.~\ref{bar-hubble} we characterize the bars in the Hubble sequence. In Sect.~\ref{disk-assemb} and Sect.~\ref{disk-hubble} we discuss on the assembly and secular evolution of disk galaxies; we emphasize the role of bars in this process, whose effect is demonstrated based on the properties of the disk stacks. In Sect.~\ref{summarysection} we summarize the main results of this paper. | 16 | 7 | 1607.07317 |
|
1607 | 1607.05312_arXiv.txt | We present the spatially resolved star formation history (SFH) of the Carina dwarf spheroidal galaxy, obtained from deep, wide-field g,r imaging and a metallicity distribution from the literature. Our photometry covers $\sim2$\,deg$^2$, reaching up to $\sim10$ times the half-light radius of Carina with a completeness higher than $50\%$ at $g\sim24.5$, more than one magnitude fainter than the oldest turnoff. This is the first time a combination of depth and coverage of this quality has been used to derive the SFH of Carina, enabling us to trace its different populations with unprecedented accuracy. We find that Carina's SFH consists of two episodes well separated by a star formation temporal gap. These episodes occurred at old ($>10$\,Gyr) and intermediate ($2$--$8$\,Gyr) ages. Our measurements show that the old episode comprises the majority of the population, accounting for $54\pm5\%$ of the stellar mass within $1.3$ times the King tidal radius, while the total stellar mass derived for Carina is $1.60\pm0.09\times 10^{6} \, M_{\rm{\odot}}$, and the stellar mass-to-light ratio $1.8\pm0.2$. The SFH derived is consistent with no recent star formation which hints that the observed blue plume is due to blue stragglers. We conclude that the SFH of Carina evolved independently of the tidal field of the Milky Way, since the frequency and duration of its star formation events do not correlate with its orbital parameters. This result is supported by the age/metallicity relation observed in Carina, and the gradients calculated indicating that outer regions are older and more metal poor. | Dwarf galaxies are crucial for understanding galaxy assembly and evolution. They are some of the oldest systems in the Universe, and inhabit the most numerous type of dark matter halos in the framework of a Lambda-CDM Universe \citep[e.g.,][]{kauffmann93a}. These systems gave origin to larger galaxies like the Milky Way in the early Universe via hierarchical merging \citep[e.g.,][]{unavane96}. The dwarf galaxies in the Local Group are particularly interesting, since their proximity allows us to resolve them into individual stars. Thus, it is not surprising that these galaxies have been studied in more detail than any other galaxies \citep{tolstoy09,mcconnachie12}. Carina is a Local Group dwarf spheroidal (dSph) galaxy located at about $104$\,kpc from the Sun \citep{karczmarek15}, with a half-light radius of $250\pm39\,$pc \citep{irwin95a}, a dynamical mass within the half-light radius of $M_{\rm dyn}(<r_{\rm half})=3.4\pm1.4\times10^{6}\,$M$_{\odot}$ \citep{walker09a}, and an absolute magnitude of M$_{V}=-9.3$ \citep{mateo98a}. This galaxy is especially important as a constraint on Galactic evolution since it is one of the few Local Group galaxies showing an episodic star formation history (SFH) \citep[e.g.,][]{weisz14b}. Furthermore, Carina is the only galaxy where the episodic star formation history translates into two clearly distinct sub-giant branches \citep[see for example][] {bono10,deBoer14}. These star formation episodes may be either related to interactions with the Milky Way or with internal evolution of its gas and stars. The SFH of local systems like Carina have been studied mainly through the analysis of their color-magnitude diagrams (CMDs). In recent years this has been done using the synthetic CMD method \citep[e.g.,][]{deBoer14,weisz14b}. This technique consists of deriving the history of a stellar system by creating different combinations of synthetic single stellar populations and comparing their properties to the ones of the stars observed. Carina was first thought to be a purely intermediate-age population galaxy, but then RR-Lyraes were discovered in this system, indicating that an old ($>10$\,Gyr) stellar population was also present in this galaxy \citep{saha86}. Multiple main-sequence turnoffs confirmed that Carina had an episodic SFH \citep{smecker-hane96,hurley-keller98}, which means that there are clearly distinguishable episodes of star formation activity separated by episodes where practically no stars were formed. Another key feature of Carina's CMD is its narrow red giant branch (RGB), which was at first interpreted as the result of a low metallicity spread \citep[see][and references therein] {rizzi03}. In that work, the authors measured a color spread of the RGB of $\sigma_{\rm{V-I}}=0.021\pm0.005$ and derived a metallicity of [Fe/H]=$-1.91$ with a spread of $0.12$\,dex, in agreement with early spectroscopic studies of upper RGB stars in Carina \citep[e.g.,][]{armandroff91}. More recent spectroscopic observations have detected a much larger metallicity spread in this galaxy \citep{helmi06a,koch06}. The latter study measured a mean metallicity of [Fe/H]$\sim-1.4$ and a spread of $0.92$\,dex. More recently, \citet{deBoer14} used Koch's metallicity distribution function (MDF) along with CMD information from their deep photometry and derived a self-consistent, complex SFH for Carina, indicating a strong age/metallicity degeneracy. This degeneracy implies that some properties of a given stellar population, like its MDF or its RGB color and width, are equivalent to the ones of a population that is older and more metal poor. Confirmation or refutation of these results would shed important light on the origin of Carina's photometric and spectroscopic features. From all these previous works there is a general agreement that Carina has a well separated, episodic SFH consisting of at least two episodes producing old an intermediate-age populations. In addition, several studies \citep{hurley-keller98,mateo98b,monelli03} claimed a third episode in Carina, consisting of young ($<1$\,Gyr) stars. However, the exact age and duration of all these potential episodes are still uncertain. Another important feature found in Carina is the metallicity and age radial gradients of its stellar populations. In Carina's outer regions, the relative prevalence of older and more metal-poor stars increases \citep{munoz06b,battaglia12,mcMonigal14,deBoer14}. One of these contributions, \citet{munoz06b} presented several observations that were indicative of tidal effects on Carina, a plausible explanation for these population gradients. These pieces of evidence include a component in the density profile that extends well beyond the nominal King limiting radius, a distribution of outer stars that lie preferentially along the major axis and a velocity dispersion profile that rises well past the limiting radius. These results were then supported by the work of \citet{munoz08a} and \citet{battaglia12}. It is worth mentioning that the proper motion study of Carina by \citet{piatek03} results in an orbit that is consistent with the tidal scenario proposed by \citet{munoz08a}. In their study, \citet{piatek03} claimed that this dSph is currently at apocenter, has an orbital period close to $2$\,Gyr and its last close passage from the Milky Way occurred $\sim0.7\,$Gyr ago. Tidal effects are interesting in the context of the work by \citet{piatek03} and \citet{pasetto11} who tested if the star formation episodes of Carina could be explained as the result of close encounters with the Milky Way. These encounters can promote star formation for example by removing angular momentum from the gas and driving it to the central regions \citep[e.g.,][]{larson02}, or by compression produced by tidal gravitational shocks \citep[e.g.,][]{pasetto11}. They placed constraints on Carina's orbit, but have achieved limited success in explaining Carina's SFH as a result of tidal shocks or ram pressure. Another proposed origin for the properties of Carina's SFH is internal evolution. For example, gas depletion or radial migration \citep{el-badry15} might produce the positive age radial gradient (average stellar population age increasing with radius) and the negative metallicity gradient (average metal content decreasing with radius). Additionally, gas heating \citep[see for example][]{revaz09} might explain the temporal gap in star formation. In summary, we do not have at present a complete scenario explaining the evolution of Carina that is consistent with its SFH, chemical enrichment, orbital information and gas dynamics. In this work we use deep/wide photometry along with public metallicities to derive the SFH of Carina. By making use of the Talos routine \citep{deBoer12} we take into consideration all the information in the CMD (and not just some key fingerprints) along with the MDF to derive the SFH in a consistent way. The high quality of the photometry along with the spatial extent of the observations (two square degrees), enable us to make three independent measures of the Carina's SFH within different concentric regions. In this way, we can quantify the dependence of the SFH on the distance from the galaxy center by looking at radial gradients. In Section~2 we present the spectroscopic and photometric data. Section~3 describes the method for deriving the SFH along with the input files used. Section~4 presents the main results for the SFH which are discussed in Section~5. | In this work we have presented the spatially resolved star formation history (measured in the inner, middle, and outer region) of the Carina dwarf spheroidal galaxy using deep, wide-field $g$, $r$ imaging and the metallicity distribution function from the data of \citet{koch06}. This is the first time a combination of depth and coverage of this quality is used to derive the SFH of Carina, enabling us to trace its different populations with unprecedented accuracy. The main results of this work can be summarized as followed: 1. The SFH of Carina shows a majority of old metal poor stars accounting for $54\pm4\%$ of the stellar mass and an intermediate-age population with increasing metallicity with time. The fraction of old stars could be larger considering the possibility that tidal influence from the Milky Way might have removed preferentially old stars from Carina, given that these stars are distributed in a more extended region. 2. Both episodes are separated by a period of no star formation. This temporal gap started $\sim10\,$Gyr ago and stopped $\sim8\,$Gyr ago. 3. Carina displays a positive age radial gradient and a negative metallicity radial gradient. The inner region is dominated by the intermediate-age population, the middle region is composed of a similar fraction of old and intermediate-age stars, and the outer region is dominated by the old metal poor population, with an almost negligible component of intermediate-age stars. 4. Results are consistent with a total stellar mass of $1.60\pm0.09\times 10^{6} \, M_{\rm{\odot}}$ within $1.3\times r_{\rm{tidal}}$, and $1.45\pm0.12\times 10^{6} \, M_{\rm{\odot}}$ within the nominal $r_{\rm{tidal}}$. The latter value is larger than previous values derived for Carina. For example, \citet{deBoer14} derived a value of $1.07\pm0.08\times 10^{6} \, M_{\rm{\odot}}$ within the nominal $r_{\rm{tidal}}$. We attribute the difference in the masses derived to the increased depth of our photometry that enabled us to detect more faint old main-sequence stars. 5. The star formation is consistent with no young ($<2$\,Gyr) stars. We calculated the fraction of blue plume stars in Carina and the fraction of blue stragglers in $4$ other local dwarf spheroidal galaxies. These fractions were calculated using the same mass range (in terms of the mass of the main-sequence turnoff) for both cases, in order to make a meaningful comparison between both fractions. Given that blue straggler fraction in the local dwarf galaxies is constant and that the blue straggler fraction of Carina is completely consistent with this value, we concluded that the blue plume in this galaxy is consistent with being composed of blue stragglers. 6. The spatially resolved SFH is consistent with being dominated by internal evolution as opposed to tidal influence from the Milky Way. Our hypothesis is that Carina formed the majority of its stars in a first episode of star formation, then the star formation ceased for a couple of gigayears due to an internal process (e.g., gas heating) and finally, after gas cooling, re-accreted the gas producing the second episode of star formation. | 16 | 7 | 1607.05312 |
1607 | 1607.07245_arXiv.txt | We present observations of three active sites of star formation in the Taurus Molecular Cloud complex taken at 323 and 608\,MHz (90 and 50\,cm, respectively) with the Giant Metrewave Radio Telescope (GMRT). Three pointings were observed as part of a pathfinder project, targeted at the young stellar objects (YSOs) L1551~IRS~5, T~Tau and DG~Tau (the results for these target sources were presented in a previous paper). In this paper, we search for other YSOs and present a survey comprising of all three fields; a by-product of the large instantaneous field of view of the GMRT. The resolution of the survey is of order 10\,arcsec and the best rms noise at the centre of each pointing is of order $100\,\umu$Jy\,beam$^{-1}$ at 323\,MHz and $50\,\umu$Jy\,beam$^{-1}$ at 608\,MHz. We present a catalogue of 1815 and 687 field sources detected above $5\,\sigma_{\rm rms}$ at 323 and 608\,MHz, respectively. A total of 440 sources were detected at both frequencies, corresponding to a total unique source count of 2062 sources. We compare the results with previous surveys and showcase a sample of extended extragalactic objects. Although no further YSOs were detected in addition to the target YSOs based on our source finding criteria, these data can be useful for targeted manual searches, studies of radio galaxies or to assist in the calibration of future observations with the Low Frequency Array (LOFAR) towards these regions. | \label{intro} The Taurus Molecular Cloud (TMC, see Fig.~\ref{fig:taurus}) is one of the best studied star forming regions due to its relatively nearby rich population of low mass young stellar objects \citep[YSOs;][]{2008hsf1.book..405K}. The TMC is located at a mean distance of 140\,pc. However, recent observations with the Very Long Baseline Array have been used to refine distance measurements and determine the three-dimensional structure of the complex \citep[which has been shown to have a depth of 30\,pc;][]{2005ApJ...619L.179L, 2009ApJ...698..242T}. The TMC is not as densely populated as other star forming regions, such as $\rho$~Ophiuchus and Orion, which allows for the study of individual low-mass protostellar systems analogous to the Sun. The low stellar density also minimises the mutual influence of outflows, jets or gravitational effects on star formation and the lack of more luminous stars (there are no O stars and only very few B and A stars) limits the effects of strong stellar winds and ionising UV radiation \citep{2007A&A...468..353G}. For these reasons, the TMC allows tests of stellar evolution models and provides the standard initial mass function for a nearby young association with stellar ages of $\sim1$--10\,Myr \citep[e.g.][]{1995ApJS..101..117K, 2004ApJ...617.1216L}. \begin{figure*} \includegraphics[width=0.7\textwidth]{Taurus.png} \caption[The Taurus Molecular Cloud]{Overview of the TMC. Greyscale is the high resolution CO extinction map from \citet{2005PASJ...57S...1D}. Circles correspond to the GMRT observed fields (the FWHM of the GMRT primary beam is $\approx85$\,arcmin at 323\,MHz and $\approx44$\,arcmin at 608\,MHz). Axes are J2000.0 Galactic coordinates.} \label{fig:taurus} \end{figure*} The TMC has been extensively surveyed in most bands of the electromagnetic spectrum, from X-ray, infrared and optical \citep[see e.g.][]{2007prpl.conf..329G} to sub-millimeter \citep{2005ApJ...631.1134A} wavelengths, providing an unsurpassed database for the nearest major star forming cloud complex. There have not, however, been many large-scale surveys at radio (centimetre) wavelengths directed at the TMC. \citet{2015ApJ...801...91D} aimed to rectify this deficit by conducting a multi-epoch radio study of the TMC at 4.5 and 7.5\,GHz as part of the Gould's Belt Very Large Array Survey (GBS-VLA) to systematically characterise the centimetre-wave properties of the YSO population. This region was also covered by all sky surveys such as the NRAO VLA Sky Survey \citep[NVSS,][]{1998AJ....115.1693C} at 1.4\,GHz and the Low Frequency Array \citep[LOFAR,][]{2013A&A...556A...2V} Multifrequency Snapshot Sky Survey \citep[MSSS,][]{2015A&A...582A.123H} at 30--160\,MHz, albeit with relatively poor angular resolution and sensitivity. There are no published surveys targeted at YSOs in the TMC at frequencies $<1$\,GHz. Furthermore, the study of YSOs at radio wavelengths has previously been largely confined to frequencies $\nu>1$\,GHz. This is due to the past sensitivity limitations of radio telescopes, the radio weakness of YSOs (flux densities of order $\sim1$\,mJy at centimetre wavelengths), the fact that they are typically detected via thermal bremsstrahlung (free--free) radiation and thus their spectra rise with frequency. We conducted a pathfinder project with the Giant Metrewave Radio Telescope (GMRT) to extend the study of young stars to very low radio frequencies. We observed the well-studied TMC members L1551~IRS~5, T~Tau and DG~Tau at frequencies $\nu=323$ and 608\,MHz (90 and 50\,cm, respectively), the results of which were presented in \citet{2016MNRAS.459.1248A}. A natural by-product of these GMRT observations was a large instantaneous field of view within which to search for additional objects, in particular other YSOs. Although the TMC is not as densely populated with young stars as other star forming regions, there are several other known pre-main-sequence objects located within the GMRT field of view of each of the target fields. For example, in the southern region of the TMC, most pre-main sequence stars are in and around the Lynds~1551 (L1551) dark cloud. In addition to L1551~IRS~5, the protostar L1551~NE and a few deeply embedded T~Tauri stars (HL~Tau, XZ~Tau and HH~30~IRS) form a close group of pre-main sequence stars \citep{2008hsf1.book..405K}. The entirety of the L1551 cloud fits within the 44\,arcmin half-power point of the GMRT primary beam at 608\,MHz (see Fig.~\ref{fig:taurus}) and therefore the GMRT can be a potentially useful survey instrument for star forming regions at very long wavelengths due to the extent of its field of view. The importance of investigating the radio emission from young stars at very long wavelengths is twofold. First, observations of the long wavelength turnover in the free-free spectrum can constrain physical properties of the ionised plasma from these systems such as gas mass and electron density \citep[see e.g.][]{2016MNRAS.459.1248A}. Second, the spectra of non-thermal emission processes \citep[such as those observed from e.g.][]{2010Sci...330.1209C, 2014ApJ...792L..18A} typically rise at longer wavelengths making them easier to detect. It is important to note however, that non-thermal coronal emission, which has been detected from a large fraction of YSOs as part of the GBS-VLA \citep[e.g.][]{2013ApJ...775...63D, 2015ApJ...801...91D} can exhibit a broad range of spectral indices and may therefore be a poor discriminant between thermal and non-thermal emission. The new generation of radio interferometers such as LOFAR and upgrades to existing facilities such as the GMRT, will allow access to a previously unexplored wavelength regime for young stars. GMRT data can also be used to assist in the calibration of observations with LOFAR towards these regions as the radio sky at such low frequencies is largely unknown in these directions. In this paper we present a full catalogue of sources detected within the GMRT field of view of each target field, including a detailed description of the survey methodology and data products. In Section~\ref{sec:obs} we provide details of the observations and data reduction. In Section~\ref{sec:cat_creation} we discuss the method used for source fitting, measuring spectral indices, catalogue creation and a sample table of the final catalogue. We also showcase a sample of extended sources. In Section~\ref{sec:dis} we discuss the YSO detections and notable non-detections. We compare between the 323 and 608\,MHz images and with previous surveys. In Section~\ref{sec:dp} we list the resulting data products and summarise the survey in Section~\ref{sec:summary}. \vspace{-10pt} | \label{sec:dis} The results presented in \citet{2016MNRAS.459.1248A} focussed on the study of emission detected from the three target YSOs at the phase centres of each field (L1551~IRS~5, T~Tau and DG~Tau). In this work, the same set of observations has been used to make a catalogue of these regions towards the TMC at 323 and 608\,MHz. Although the TMC is not as densely populated as other star forming regions \citep[e.g. $\rho$~Ophiuchus, Orion;][]{2015ApJ...801...91D}, there are several other pre-main-sequence objects with known radio emission located within the observed fields. However, no additional protostellar objects were detected apart from the target YSOs. \subsection{YSO detections} We cross-referenced the GMRT-TAU catalogue with the GBS-VLA to search for additional known radio emitting YSOs within our fields. Catalogue sources GMRT-TAU~J043134.03+180803.9 and GMRT-TAU~J042159.51+193206.3 are identified as the target sources L1551~IRS~5 (GBS-VLA~J043134.16+180804.6) and T~Tau (GBS-VLA~J042159.43+193205.7), respectively. No additional known YSOs from the GBS-VLA were detected within the source finding criteria described in Section~\ref{sec:sf} (i.e. at the $5\,\sigma_{\rm rms}$ level). We also searched for sources within our catalogue which would have thermal spectral indices when compared with objects within the GBS-VLA. We find two thermal sources within the L1551~IRS~5 field other than L1551~IRS~5 itself: GMRT-TAU~J043150.37+182052.5 (GBS-VLA~J043150.44+182052.6) and GMRT-TAU~J043229.39+181359.8 (GBS-VLA~J043229.46+181400.2). GMRT-TAU~J043229.39+181359.8 was classified as a Class~II transition disk candidate (Tau~L1551~13) by \citet{2009ApJS..184...18G} through application of mid-infrared colour-based methods to \textit{Spitzer} data. This object has also been detected with 2MASS, XEST, NVSS and the GBS-VLA. Based on its 4.5--7.5\,GHz properties, \citet{2015ApJ...801...91D} suggested this object (GBS-VLA~J043229.46+181400.2) is instead an extragalactic background source based on its low variability ($<50$\,per~cent), negative spectral index ($\alpha_{\rm 4.5\,GHz}^{\rm 7.5\,GHz}=-0.25\pm0.71$) and lack of proper motion. Its radio flux density is also large at centimetre wavelengths ($\sim60$\,mJy) compared with most other YSOs in the TMC ($\sim1$\,mJy). Combining the GMRT 323 and 608\,MHz data, NVSS data at 1.4\,GHz and the GBS-VLA data at 4.5 and 7.5\,GHz, we find a spectral index $\alpha_{\rm 323\,MHz}^{\rm 7.5\,GHz}=0.31\pm0.03$ for GMRT-TAU~J043229.39+181359.8. This spectral index is consistent with thermal, free--free radiation which is typically detected within star forming regions from YSO outflows. However, the relatively high flux density compared to other YSOs within the TMC and the lack of proper motion \citep[members exhibit proper motions of $\sim20$\,mas\,yr$^{-1}$, see e.g.][]{2009ApJ...698..242T} still suggests that this object is extragalactic and the thermal spectrum may be related to star formation activity. We follow \citet{2015ApJ...801...91D} and do not consider this object to be a Taurus YSO for the rest of this paper. The only thermal source we detect within the T~Tau field when compared with the GBS-VLA is T~Tau itself. We detect two sources with apparent thermal spectral indices within the DG~Tau field: GMRT-TAU~J042452.49+264203.2 (GBS-VLA~J042452.48+264204.5) which has been classified as extragalactic by the Sloan Digital Sky Survey \citep[SDSS][]{2011yCat.2306....0A}, and GMRT-TAU~J042929.48+263151.6 (GBS-VLA~J042929.49+263152.8) which has been classified as a field star by \citet{2015ApJ...801...91D} but as a known galaxy by \citet{2011ApJS..196....4R} and SDSS. The GBS-VLA and SDSS properties of GMRT-TAU~J042929.48+263151.6 suggest that it is extragalactic. We therefore do not detect any YSOs in addition to the three target sources based on a cross-reference with the GBS-VLA. It is possible that targeted manual searches may reveal detections. \subsection{Notable YSO non-detections} Although no additional YSOs with known radio emission were detected within our fields, we emphasise that the work presented here is a by-product of a pathfinder project to detect the specific YSO target sources at the phase centre of each pointing, which was successful. We discuss some cases of notable YSO non-detections in the following paragraphs. Variability may partially account for the non-detections of additional YSOs within our single epoch observations. Non-thermal radio emission associated with active coronae (which we may expect to detect at such low frequencies) often exhibits high levels of variability \citep[>50\,per~cent;][]{1999ARA&A..37..363F, 2015ApJ...801...91D}. As mentioned earlier, the GBS-VLA results further showed that this emission mechanism can exhibit a broad range of spectral indices. The fact that only the YSOs at the phase centres were detected might suggest that either there are very few YSOs with active coronae in the TMC or that the non-thermal YSOs do not have very negative spectral indices as was shown by \citet{2015ApJ...801...91D}. However, the low angular resolution of our survey ($\sim10$\,arcsec) may also prevent detections of this non-thermal coronal emission which arises on much smaller scales close to the source. Longer on-source time may yield larger YSO detection rates within these fields. To test this, we conducted a follow-up to the current pathfinder project with a blind survey of the crowded star forming region, NGC~1333 in the Perseus Molecular Cloud with the GMRT at 610\,MHz. These observations, which will be presented in a forthcoming paper, have a much longer integration time ($\sim58$\,hr) in order to take full advantage of the large field of view and showcase the potential survey capability of star forming regions with the GMRT. Furthermore, the capabilities of the GMRT are currently being enhanced with wider frequency coverage (almost seamless from 130 to 1500\,MHz), larger bandwidth (ten-fold increase from 32 to 400\,MHz) and more sensitive receivers. This upgraded GMRT (uGMRT), which has been designated as a SKA pathfinder, will be the most suited instrument for large-scale surveys of star forming regions at low radio frequencies. \subsubsection{L1551~IRS~5 field} \label{sec:L1551} The entire L1551 cloud itself fits within the FWHM of the GMRT primary beam at 608\,MHz \citep[and therefore also at 323\,MHz, see overview of L1551 region in H$\alpha$ and \rm{[S~\textsc{ii}]} in Fig.~1 of][]{1999AJ....118..972D}. Other well-studied YSO systems within this complex include L1551~NE, XZ~Tau, HL~Tau, LkH$\alpha$~358 and the driving source of HH~30 (V1213~Tau), the positions of which are shown in Fig.~\ref{fig:L1551region}. It can be seen that this region suffers from residual calibration artefacts due to a bright radio galaxy (northern lobe: GMRT-TAU~J043144.24+181041.8 with $S_{\rm 608\,MHz, int}=34.74\pm1.96$\,mJy and southern lobe: GMRT-Tau~J043143.78+181024.48 with $S_{\rm 608\,MHz, int}=60.00\pm3.14$\,mJy) which results in increased noise levels in this area of the map. \begin{figure} \includegraphics[width=\columnwidth]{L1551_region.png} \caption{The L1551 region at 608\,MHz. Greyscale ranges from $-0.1$ to 2\,mJy and contours are $-3$, 3, 5, 8, 11, 14, 17 and $20\times\sigma_{\rm rms}$ where $\sigma_{\rm rms}=54\,\umu$Jy\,beam$^{-1}$ (Table~\ref{tab:srclist}) is measured in a local patch of sky near L1551~IRS~5. We note that $\sigma_{\rm rms}$ will increase in the areas close to the bright radio galaxy dominating the map. Axes are J2000.0 coordinates and the \textsc{clean} restoring beam is shown as a filled ellipse in the bottom left corner (see Table~\ref{tab:srclist} for dimensions). Crosses denote locations of known YSOs. See Section~\ref{sec:L1551} for details.} \label{fig:L1551region} \end{figure} L1551~NE has been previously studied at centimetre wavelengths \citep[e.g.][]{2002AJ....124.1045R, 2012MNRAS.423.1089A}. Based on the spectral index and 4.5\,GHz flux density values from the GBS-VLA, L1551~NE (GBS-VLA~J043144.49+180831.6) has a predicted flux density of $0.32\pm0.45$\,mJy at 608\,MHz. Based on a local rms noise value of $93\,\umu$Jy\,beam$^{-1}$ (measured with the \textsc{AIPS} task \textsc{imean}) in the 608\,MHz (catalogue) map, we might expect to detect this object at $3\,\sigma_{\rm rms}$, however we do not and place a $3\,\sigma_{\rm rms}$ upper limit of $279\,\umu$Jy\,beam$^{-1}$ at 608\,MHz. There is a source of emission immediately to the west (peak emission located $\approx16$\,arcsec west) of the L1551~NE position, catalogued as GMRT-TAU~J043143.29+180831.5 with $S_{\rm 323\,MHz, int}=4.44\pm0.40$\,mJy and $S_{\rm 608\,MHz, int}=1.10\pm0.15$\,mJy. It is possible that this non-thermal emission is associated with the L1551~NE outflow similar to the non-thermal emission seen from the DG~Tau outflow \citep[][see Section~\ref{sec:dgtau} below]{2014ApJ...792L..18A}, however further observations are required to confirm this. XZ~Tau and HL~Tau are separated by $\sim25$\,arcsec, are both drivers of impressive outflows and are known radio emitters \citep[e.g.][]{1994ApJ...427L.103R, 2009ApJ...693L..86C}. The GBS-VLA calculates a spectral index of $\alpha_{\rm 4.5\,GHz}^{\rm 7.5\,GHz}=-0.85\pm0.36$ for XZ~Tau (GBS-VLA~J043140.09+181356.7), which is indicative of non-thermal (gyro)synchrotron radiation associated with magnetic activity. XZ~Tau therefore has a predicted flux density of $2.19\pm1.72$\,mJy at 608\,MHz based on the GBS-VLA data. This should be easily detectable within our survey, however from Fig.~\ref{fig:L1551region} we do not detect XZ~Tau at $3\,\sigma_{\rm rms}$ (where $58\,\umu$Jy\,beam$^{-1}$ is the local rms noise measured with \textsc{imean}). We therefore place an upper limit of $174\,\umu$Jy\,beam$^{-1}$ for XZ~Tau at 608\,MHz. We place the same upper limit on HL~Tau (GBS-VLA~J043138.42+181357.3) which has a predicted 608\,MHz flux density of $0.03\pm0.03$\,mJy based on GBS-VLA data. LkH$\alpha$~358 and V1213~Tau also drive outflows but were not detected by the GBS-VLA and have not previously been detected at centimetre wavelengths \citep[to the authors' knowledge, although LkH$\alpha$~358 was recently imaged with the Atacama Large Millimeter/submillimeter Array at 2.9\,mm, see][]{2015ApJ...808L...3A}, consistent with the non-detections in this work. We provide a $3\,\sigma_{\rm rms}$ upper limit of $174\,\umu$Jy\,beam$^{-1}$ for LkH$\alpha$~358 and V1213~Tau at 608\,MHz. \subsubsection{DG~Tau field} \label{sec:dgtau} DG~Tau itself is not detected within the GMRT-TAU catalogue due to the fact that it is only detected at $3\,\sigma_{\rm rms}$ and the source fitting criteria requires a peak flux of $5\,\sigma_{\rm rms}$. We do however, detect the radio emission suggested to be a bow shock associated with the DG~Tau outflow \citep{2014ApJ...792L..18A} using the source finding criteria. The emission is detected as source GMRT-TAU~J042704.08+260559.3 with $S_{\rm 323\,MHz, int}=1.32\pm0.22$\,mJy, $S_{\rm 608\,MHz, int}=1.39\pm0.14$\,mJy and $\alpha_{\rm GMRT}=0.09\pm0.31$. These are not in agreement or within the errors of the results presented in \citet{2014ApJ...792L..18A} which could be due to the different methods of source fitting. \textsc{PyBDSM} fit this emission using a single Gaussian at 608\,MHz, despite its complicated structure, which may have given it an artificially high integrated flux density at this frequency. The peak flux densities of $S_{\rm 323\,MHz, peak}=0.90\pm0.15$\,mJy\,beam$^{-1}$, $S_{\rm 608\,MHz, peak}=0.50\pm0.12$\,mJy\,beam$^{-1}$ are in better agreement with the measurements and negative spectral index in \citet{2014ApJ...792L..18A}. Located $\approx55''$ to the southwest of DG~Tau is DG~Tau~B, the driving source of the asymmetrical optical jet HH~$159$ \citep{1983ApJ...274L..83M} and not thought to be related to DG~Tau except by projected proximity \citep{1986ApJ...311L..23J}. DG~Tau~B has previously been detected at radio wavelengths \citep[e.g.][]{2012RMxAA..48..243R, 2012MNRAS.420.3334S, 2013MNRAS.436L..64A} and, based on the GBS-VLA data (GBS-VLA~J042702.56+260530.4), has a predicted 608\,MHz flux density of $0.12\pm0.08$\,mJy. With a local rms noise of $80\,\umu$Jy\,beam$^{-1}$, we place a $3\,\sigma_{\rm rms}$ upper limit of $240\,\umu$Jy\,beam$^{-1}$ for DG~Tau~B at 608\,MHz. \subsection{Comparison between 323 and 608\,MHz images} Fig.~\ref{fig:cat} shows the distribution of peak flux density, integrated flux density and the \textsc{Resid\_Isl\_rms} flux (local rms noise) across the survey at 323 and 608\,MHz. The residual island rms flux is the rms of the flux left in an island after the modelled source(s) on the island have been subtracted, and is thus a combination of a measure of the local noise and the quality of the fit. \begin{figure*} \subfloat[323\,MHz Peak Flux Density]{\includegraphics[width=0.5 \textwidth]{fullcat_peak_flux_freq1_greyscale2.pdf} \label{fig:325peak}} \subfloat[608\,MHz Peak Flux Density]{\includegraphics[width=0.5 \textwidth]{fullcat_peak_flux_freq2_greyscale2.pdf} \label{fig:610peak}} \\ \subfloat[323\,MHz Integrated Flux Density]{\includegraphics[width=0.5 \textwidth]{fullcat_total_flux_freq1_greyscale2.pdf} \label{fig:325int}} \subfloat[608\,MHz Integrated Flux Density]{\includegraphics[width=0.5 \textwidth]{fullcat_total_flux_freq2_greyscale2.pdf} \label{fig:610int}} \\ \subfloat[Uncertainty in 323\,MHz Peak Flux Density]{\includegraphics[width=0.5 \textwidth]{fullcat_err_freq1_greyscale2.pdf} \label{fig:325uncertainty}} \subfloat[Uncertainty in 608\,MHz Peak Flux Density]{\includegraphics[width=0.5 \textwidth]{fullcat_err_freq2_greyscale2.pdf} \label{fig:610uncertainty}} \caption[Flux density distributions.]{Flux density distribution across the catalogue sources at 323 and 608\,MHz. Each row of panels present the peak flux density (a and b), integrated flux density (c and d) and \textsc{PyBDSM} peak flux density uncertainty (e and f) at 323 and 608\,MHz. The continuous line in (e) and (f) indicates the cumulative distribution function.} \label{fig:cat} \end{figure*} Fig.~\ref{fig:gmrt_spindx} shows the distribution of $\alpha_{\rm GMRT}$ which has a median of $-1.10$ and a standard deviation of $0.67$, consistent with a majority of sources in the survey being Active Galactic Nuclei (AGN) emitting synchrotron radiation \citep[see e.g.][]{1970ranp.book.....P}. Fig.~\ref{fig:gmrt_spindx_dist} shows the variation of spectral index with the log of the integrated source flux. As expected, fainter sources with lower signal to noise ratios result in a larger spread of spectral index values. A notable observational bias can be seen at the lowest fluxes, where faint non-thermal sources detected at 323\,MHz are unlikely to be detected at 608\,MHz, thus leaving an apparent gap in the bottom left of the plot. The dashed line plotted shows the range of 323\,MHz fluxes and spectral index values corresponding to the $3\,\sigma_{\rm rms}$ sensitivity at 608\,MHz of $\sim200\,\umu$Jy\,beam$^{-1}$. No detections are expected below this line. \begin{figure*} \subfloat[323 to 608 MHz Spectral Index]{\includegraphics[width=0.5 \textwidth]{fullcat_alpha_gmrt_greyscale2.pdf} \label{fig:gmrt_spindx}} \subfloat[Distribution of 323 to 608 MHz Spectral Index to Flux Density]{\includegraphics[width=0.5 \textwidth]{fullcat_alpha_gmrt_evo_greyscale2.pdf} \label{fig:gmrt_spindx_dist}} \caption[Spectral index distributions.]{(a) The 323 to 608\,MHz GMRT spectral index distribution with a median value of $\alpha_{\rm GMRT}=-1.10$ and a standard deviation of $0.67$. (b) Plot of $\alpha_{\rm GMRT}$ vs. the integrated flux density at 323\,MHz. The dashed line indicates the $3\,\sigma_{\rm rms}$, $200\,\umu$Jy\,beam$^{-1}$ sensitivity threshold at which the higher frequency (608\,MHz) would not be able to detect (non-thermal) emission.} \end{figure*} In Fig.~\ref{fig:spix} we present spectral index maps of three extended radio galaxies detected in our survey made using the \textsc{spix} operation within the AIPS task \textsc{comb}. We set \textsc{aparm(9)} and \textsc{(10)} to be the local $5\,\sigma_{\rm rms}$ cutoff for each frequency to eliminate low signal-to-noise (SNR) emission which can create artificially steep spectral indices. In general, the bright regions coincide with flatter indices ($\alpha_{\rm GMRT}\approx-1$) and the more diffuse regions exhibit steeper indices ($\alpha_{\rm GMRT}\approx-2$). This steeping can be interpreted in terms of spectral ageing of electrons as they move from acceleration sites in the bright radio knots to regions further along the jet, by which time synchrotron cooling has reduced the population of high energy electrons relative to the lower energies \citep{1962SvA.....6..317K}. Therefore the steepening trend observed is realistic, however we caution that the spectral indices themselves may not be due to the relatively high errors involved in a two point fit of measurements close in frequency. \begin{figure*} \subfloat[GMRT-TAU J043051.62+180819.1 and GMRT-TAU J043053.60+180752.4]{\includegraphics[width=0.29\textwidth]{L1551_field_source1_spix.png}} \qquad \subfloat[GMRT-TAU J043140.25+180325.8 and GMRT-TAU J043141.23+180248.3]{\includegraphics[width=0.26\textwidth]{L1551_field_source2_spix.png}} \qquad \subfloat[GMRT-TAU J042729.21+255046.0, GMRT-TAU J042730.16+255121.0 and GMRT-TAU J042730.32+255013.2]{\includegraphics[width=0.35\textwidth]{DGTau_field_source_spix.png}} \caption{A sample of spectral index maps for extended objects in the GMRT-TAU fields. (a) and (b) are located within the L1551~IRS~5 field, and (c) is located within the DG~Tau field. Greyscale (colour in the online journal) represents the range in spectral index. Contours are plotted as $-3$, 3, 5, 10, 15, 20, 25, 30, 40, 50, 60, 70, 80, 90 and $100\times\sigma_{\rm rms}$, where $\sigma_{\rm rms}$ is the local rms noise in the 323\,MHz image and equals (a) 158, (b) 148 and (c) $142\,\umu$Jy\,beam$^{-1}$. Axes are J2000.0 coordinates and the \textsc{clean} restoring beam is shown as a filled ellipse in the bottom left corner of each map. Fluxes were clipped below $5\,\sigma_{\rm rms}$ at both frequencies during spectral index map construction to eliminate artificially steep indices caused by low SNR emission, however we show $3\,\sigma_{\rm rms}$ contours to remain consistent with the source fitting procedure (Section~\ref{sec:sf}) and Fig.~\ref{fig:maps}.} \label{fig:spix} \end{figure*} \subsection{Comparison with NVSS} Fig.~\ref{fig:nvss_spindx} shows the distribution of $\alpha_{\rm NVSS}$ which has a median of $-0.80$ and a standard deviation of $0.36$, again consistent with synchrotron radiation from AGN. Note that these plots present a combination of two and three frequency calculations of the spectral index, depending on whether the source was detected at both GMRT frequencies or not. Fig.~\ref{fig:nvss_spindx_dist} again shows a gap in the bottom left of the plot due to observational bias, this time exaggerated due to the larger spread in frequencies. The dashed line plotted shows the range of 323\,MHz fluxes and spectral index values corresponding to the $3\,\sigma_{\rm rms}$ sensitivity at 1.4\,GHz of $\sim1.5$\,mJy\,beam$^{-1}$. No detections are expected below this line. \begin{figure*} \subfloat[GMRT to NVSS Spectral index]{\includegraphics[width=0.5 \textwidth]{fullcat_alpha_nvss_greyscale2.pdf} \label{fig:nvss_spindx}} \subfloat[Distribution of GMRT to NVSS Spectral Index to Flux Density]{\includegraphics[width=0.5 \textwidth]{fullcat_alpha_nvss_evo_greyscale2.pdf} \label{fig:nvss_spindx_dist}} \caption[Spectral index distributions.]{ (a) The GMRT to NVSS (1.4\,GHz) spectral index distribution which has a median value of $\alpha_{\rm NVSS}=-0.80$ with a standard deviation of $0.36$. (b) Plot of $\alpha_{\rm NVSS}$ vs. the integrated flux density at 323\,MHz. The dashed line indicates the $3\,\sigma_{\rm rms}$, 1.5\,mJy\,beam$^{-1}$ sensitivity threshold at which the higher frequency (1.4\,GHz VLA) would not be able to detect (non-thermal) emission.} \end{figure*} As described in Section~\ref{sec:obs}, a small systematic offset appeared to be present between the 323 and 608\,MHz maps of the L1551~IRS~5 and DG~Tau fields. The positions of the fitted bright, compact sources in each field at 608\,MHz are in general consistent with those from the NVSS to within an average of 0.5\,arcsec, making them more reliable than the positions at 323\,MHz. However, due to a lack of low frequency radio surveys at higher resolution for this patch of the sky (NVSS has a spatial resolution of approximately 45\,arcsec) the absolute coordinates may have a residual uncertainty of around 0.5\,arcsec. This should not have a significant impact as the GMRT spatial resolution is of order 5--10\,arcsec and we do not apply further correction as it is uncertain as to which survey the errors come from \citep[see e.g.][]{2007MNRAS.376.1251G}. \vspace{-10pt} \label{sec:summary} In this paper we have described a 323 and 608\,MHz (90 and 50\,cm, respectively) survey with the Giant Metrewave Radio Telescope of three regions towards the Taurus Molecular Cloud, specifically towards the target pre-main-sequence stars L1551~IRS~5, T~Tau and DG~Tau. This survey is a natural by-product of the large instantaneous field of view of the GMRT. Although we did not detect other YSOs in addition to the targets based on our source finding criteria, we provide a catalogue of field sources which can be useful for targeted manual searches, studies of radio galaxies or to assist in the calibration of future observations with LOFAR towards these regions. We emphasise that the work presented here is part of a pathfinder project to detect the specific YSO target sources at the phase centre of each pointing, which was successful \citep{2016MNRAS.459.1248A}. Longer on-source time may yield larger YSO detection rates within these fields and we have conducted a follow-up survey of the crowded star forming region NGC~1333 in the Perseus Molecular Cloud with a much longer integration time at 610\,MHz (Ainsworth et~al., in preparation). These observations will showcase the potential survey capability of star forming regions with the GMRT. The resolution of the survey is of order 10\,arcsec and the best rms noise at the centre of each pointing is of order $100\,\umu$Jy\,beam$^{-1}$ at 323\,MHz and $50\,\umu$Jy\,beam$^{-1}$ at 608\,MHz. The final data products comprise 12 images, the final source catalogue, an additional catalogue containing the results corresponding to all \textsc{PyBDSM} parameters and the scripts used to generate the catalogue and validation plots. These data will be made available on the project website\footnotemark[3] and on the VizieR online database. The final catalogue contains 1815 sources at 323\,MHz and 687 sources at 608\,MHz. A total of 440 sources were detected at both frequencies which yields a total unique source count of 2062. The catalogue has been cross-referenced with existing radio, infrared and X-ray surveys conducted towards this region and compared with other GMRT surveys and the NVSS. Notable YSO non-detections have been discussed and a sample of extended radio galaxies has been shown. The capabilities of the GMRT are currently being enhanced with wider frequency coverage, larger bandwidth and more sensitive receivers. This uGMRT has been designated as a SKA pathfinder and will be the most suited instrument for large-scale surveys of star forming regions at very low radio frequencies. \vspace{-10pt} | 16 | 7 | 1607.07245 |
1607 | 1607.05289_arXiv.txt | We use spectra from the ALFALFA, GASS and COLD GASS surveys to quantify variations in the mean atomic and molecular gas mass fractions throughout the \ms\ plane and along the main sequence (MS) of star-forming galaxies. Although galaxies well below the MS tend to be undetected in the Arecibo and IRAM observations, reliable mean atomic and molecular gas fractions can be obtained through a spectral stacking technique. We find that the position of galaxies in the \ms\ plane can be explained mostly by their global cold gas reservoirs as observed in the HI line, with in addition systematic variations in the molecular-to-atomic ratio and star formation efficiency. When looking at galaxies within $\pm0.4$ dex of the MS, we find that as stellar mass increases, both atomic and molecular gas mass fractions decrease, stellar bulges become more prominent, and the mean stellar ages increase. Both star formation efficiency and molecular-to-atomic ratios vary little for massive main sequence galaxies, indicating that the flattening of the MS is due to the global decrease of the cold gas reservoirs of galaxies rather than to bottlenecks in the process of converting cold atomic gas to stars. | \label{intro} The current galaxy evolution framework gives central stage to the cycling of gas in and out of galaxies, and the efficiency of the star formation process out of the gas that cools and settles into galactic discs. These elements are responsible for regulating star formation in galaxies, much more so than, for example, merger-driven starbursts. Simple numerical and analytical models that focus on the gas reservoir of galaxies, and how it is replenished by inflows and depleted by star formation and outflows \citep[e.g.][]{frenkwhite91,bouche10,dave11,dave12,krumholz12,lilly13}, are well supported by observations of the redshift evolution of the molecular gas contents of galaxies \citep{tacconi10,tacconi13,magdis12a,saintonge13,genzel15}. These studies show that the redshift evolution of the specific star formation rate (SSFR) can be explained simply by the changes in the molecular gas contents of galaxies, and by the mean star formation efficiency that increases slightly with redshift. Other studies have looked at variations in the gas contents and star formation efficiency of galaxies as a function of SSFR, or distance from the star formation main sequence. At all redshifts up to $z\sim2$, it has been established that galaxies well above the main sequence have star formation efficiencies enhanced by an order of magnitude \citep[e.g.][]{gao04,genzel10,daddi10}, and that almost all of these most extreme systems are major mergers \citep{sanders96,veilleux02,engel10}. Further studies have shown that star formation efficiency is not a bimodal quantity, but rather varies smoothly as a function of SSFR, or distance from the MS \citep{COLDGASS2}. The molecular gas contents of galaxies also varies with distance from the MS, indicating that the high SSFRs of starbursting galaxies are caused both by enhanced molecular gas mass fractions and increased star formation efficiencies, while the reverse effect is observed in bulge-dominated, below-MS galaxies \citep{saintonge12,genzel15}. In this paper, we push these investigations further in two main respects. Firstly, we focus on a local sample, which gives us access not only to the molecular gas, but also to the cold atomic gas contents of galaxies through HI 21cm observations. By studying both atomic and molecular gas fractions, we can identify if the conversion of atomic to molecular gas is in any region of the SFR-\mstar\ plane a bottleneck in the star formation process. Secondly, instead of restricting our study to gas-rich main sequence galaxies, we take advantage of the size and completeness of the GASS sample and make use of a spectral stacking technique to measure mean molecular and atomic gas fractions in the entire SFR-\mstar\ plane, including galaxies with SFRs more than two orders of magnitude lower than main sequence galaxies of the same mass. All rest-frame and derived quantities assume a \citet{chabrier03} IMF, and a cosmology with $H_0=70$\kms\ Mpc$^{-1}$, $\Omega_m=0.3$ and $\Omega_{\Lambda}=0.7$. | \label{summary} The objective of this paper was to provide a first quantitative measure of the mean atomic and molecular gas contents of galaxies across the entire \ms\ plane for galaxies with \mstar$>10^{10}$\msun. This is made possible by the GASS / COLD GASS surveys, which have obtained HI 21cm and CO(1-0) line flux measurements for a large, representative SDSS-selected sample. Even though individual galaxies are not all detected in the HI and/or CO line, we can derive accurate mean gas fractions over the entire \ms\ plane using a spectral stacking technique, owing to the homogeneous observing strategy of the surveys and their well-defined selection functions out of the larger SDSS parent sample. Building on these strengths, we set out with the following objectives: (1) to track systematically how gas contents, star formation efficiency and the molecular-to-atomic ratio vary across the \ms\ plane, (2) to understand how, in an average sense, the position of galaxies in the plane are related to their gas contents, (3) to identify regions of the plane where galaxies are out of equilibrium and in the process of quenching, and (4) to follow specifically the gas properties of the objects that form the main sequence of star forming galaxies. To address questions (1)-(3) we separated the \ms\ plane in 14 bins defined by lines of constant SSFR and \mstar. Stacking the HI and CO spectra gave us mean atomic and molecular gas fractions in each of these bins (see Figure \ref{gasMS}). We find that \fhi\ can explain to first order the SSFRs of different galaxy populations, although note that \fhi\ has an additional dependence on \mstar\ (eq. \ref{fHIeq}) and cannot alone account for the very high SSFRs of above main sequence galaxies. On the other hand, \fgas\ depends tightly on SSFR with no residual \mstar\ dependence because the molecular-to-atomic ratio is a function of stellar mass, with the most massive galaxies at fixed SSFR having a higher $R_{mol}$. Our quantitative analysis of the variations of the molecular gas mass fraction in the \ms\ plane (eq. \ref{fH2eq}) reveals the sub-linear slope of the relation between \fgas\ and SSFR; this implies that star formation efficiency must also be varying with SSFR to account for the full range of values \citep[a result previously reported in a number of studies, including ][]{saintonge12,genzel15}. We therefore conclude that, on average, {\it the position of galaxies in the \ms\ plane are determined by (1) how much cold gas is present as traced by HI, (2) how much of that cold gas is in the molecular phase and available for star formation, and (3) how efficient is the process of converting that molecular gas into stars.} We also identify a region of the \ms\ plane ($\log M_{\ast}/M_{\odot}>10.8$ and $-10.4<\log{\rm SSFR}<-9.6$) that is populated by galaxies which are within $\sim1$ Gyr of quenching. Galaxies in this ``danger zone" have on average very high molecular-to-atomic mass ratios, $R_{mol}>0.7$, more than twice as large as the mean for main-sequence galaxies. In the absence of extended HI envelopes or other sources of accretion to replenish the gas reservoirs, these galaxies will cease actively forming stars and migrate to the red cloud. Galaxies in the ``danger zone" are disc-like but with important bulge components ($<\log \mu_{\ast}>=8.9$), while also having young stellar populations in their central regions ($<D_n(4000)>=1.38$). This is an unusual combination: throughout the rest of the \ms\ plane, there is a strong correlation between bulge-like morphology and older stellar ages (see Fig. \ref{morphMS}). Indeed, a control sample to the galaxies in the ``danger zone" matched on \mstar\ and \must\ has a distribution of values of $D_n(4000)$ with a Kolmogorov-Smirnoff probability $<0.001\%$ of being extracted from the same parent sample, and a mean value of $<D_n(4000)>=1.85\pm0.05$. Our interpretation is that the bulge-dominated galaxies in the ``danger zone" benefit from a mechanism that efficiently drives gas towards their central, high-density regions, explaining the large values of $R_{mol}$ and the high levels of star formation in the central bulge region driving down the $D_n(4000)$ index. We observe that strong stellar bars are a common feature of galaxies in the ``danger zone", and are indeed an example of a mechanism that can trigger inward radial gas motions and therefore increased central gas concentrations and star formation rates \citep[e.g.][]{sakamoto99,sheth05,masters12}. We also looked specifically at the properties of galaxies along the main sequence, which importantly we define here not as a linear relation between $\log M_{\ast}$ and $\log {\rm SFR}$, but by following the ridge traced by star-forming galaxies in the \ms\ plane. We take main sequence galaxies to be those with $| \log ({\rm SFR/SFR_{MS}}) |<0.4$ dex (see Fig. \ref{figsample} and Eq. \ref{MSeq}). As shown in Figure \ref{HIH2}, the mean molecular and atomic gas fractions of galaxies decline steadily with \mstar\ along the main sequence. In contrast, both the molecular gas depletion timescale, \tdep, and the molecular-to-atomic ratio, $R_{mol}$ are very near constant along the main sequence. This implies that the reason for the flattening of the MS at \mstar$>10^{10}$\msun\ is the gradual decrease in the total cold gas mass fraction of star-forming galaxies, and not because of a reduction of the conversion of atomic into molecular gas or of the efficiency of star formation. Other quantities are well known to vary systematically across the \ms\ plane, including measures of stellar population properties and different morphological indicators such as stellar mass surface density, concentration index, and S\'ersic index. As these quantities appear mostly constant at fixed SSFR \citep[see Fig. \ref{morphMS} and e.g.][]{wuyts11}, the flattening of the MS also means they vary systematically as star-forming galaxies grow more massive. Our analysis of the properties of massive galaxies along the main sequence thus reveals that multiple transformations are occurring: galaxies grow central bulges, the mean age of their stellar populations increases, and their entire cold gas reservoir (atomic and molecular) decreases. A significant outstanding question in galaxy evolution studies concerns the mechanisms responsible for quenching, meant as the transition of galaxies from the star forming to the passive population. A vast number of mechanisms have been explored to explain quenching. These include processes that actively remove gas from galaxies such as feedback from AGN and star formation or stripping \citep[e.g.][]{gunngott72,cicone14,forster14}, to more passive mechanisms that simply prevent galaxies from accreting fresh gas \citep[e.g.][]{keres05,dekel06,peng15}. Such mechanisms are usually invoked to explain the relatively quick transition of galaxies form the main sequence to the passive population, but our results here suggest that quenching has already started happening for massive galaxies while on the main sequence. The mechanism responsible for this must be able to account for the simultaneous reduction of the gas fractions of massive main sequence galaxies, ageing of their stellar population and growth of their central bulges. | 16 | 7 | 1607.05289 |
1607 | 1607.05847_arXiv.txt | \vspace{-0.1in} The interaction of the shock waves of young supernovae (SNe) with the circumstellar medium (CSM) is the main driving power of the observed thermal and non-thermal emission that is detected predominantly in radio and X-rays. The shocked CSM is heated to high temperatures ($\sim 10^7$~K), while the shock-accelerated electrons ({\it primaries}) emit synchrotron radiation that typically peaks in the radio bands. Together with electrons, protons of the CSM are also accelerated at the shock to multi-TeV/PeV energies (for a review see \cite{blasi2013}). Inelastic proton-proton (\pp) collisions between the accelerated and non-relativistic protons of the shocked CSM lead to the production of pions ($\pi^{+},\pi^{-},\pi^{0}$), which further decay into lighter particles including $e^{-} e^{+}$ pairs ({\it secondaries}), neutrinos and $\gamma$-rays. Due to the small cross section for \pp collisions ($\sigma_{pp}\simeq 3\times 10^{-26}$~cm$^2$), high CSM densities are required for a non-negligible contribution of secondary particles to the observed emission. Such conditions may be realized whenever a shock plunges through a dense stellar wind. As SNe occur in CSM with a wide range of densities, spreading over at least six orders of magnitude, they provide a promising testing ground for cosmic-ray (CR) acceleration theories through the detection of multi-messenger signatures produced by the secondary particles. Although the smoking gun for CR acceleration in interaction-powered SNe would be the detection of high-energy neutrinos, a firm association of the IceCube events with one (or more) astrophysical candidate classes of sources is still lacking \citep[e.g.][]{aartsen2015a, aartsen2015b}. Yet, the radio synchrotron emission from secondary $e^{-}$ and the $\gamma$-ray emission from $\pi^0$ decays may be observable under certain conditions. Our preliminary semi-analytic estimates of the former are encouraging, as we have predicted observable features in the radio-mm spectra and light curves \citep{petropoulou2016radio}. Here, we present the first results of a theoretical investigation of the $\gamma$-ray emission in the energy band of \emph{Fermi}-LAT. For this purpose, we have been developing an one-zone model for the time-dependent calculation of the multi-wavelength (MW) non-thermal emission in SNe with dense CSM, which makes testable predictions in the \emph{Fermi}-LAT band. Ultimately, the model predictions can be used to probe CR acceleration in young SNe with dense CSM and, in particular, to indirectly estimate microphysical parameters, such as the (primary) electron-to-proton ratio $K_{ep}$ and the CR acceleration efficiency $\epsilon_p$. | \vspace{-0.1in} We presented the temporal evolution of the non-thermal MW emission from interaction-powered SNe (e.g. Type IIn) while focusing on the $\gamma$-ray emission produced by \pp collisions between the shock-accelerated protons and the non-relativistic protons of the CSM. We have shown that for sufficiently dense CSM ($n_{csm} > 10^{11}$ cm$^{-3}$) or for a nearby SN Type IIn ($d_L<10$Mpc), the $\gamma$-ray ($>100$MeV) emission would be detectable by \emph{Fermi}-LAT. The non-detection of $\gamma$-rays from extragalactic interaction-powered SNe by the \emph{Fermi}-LAT so far, can be thus used to constrain the CR acceleration efficiency. We have explicitly shown that the $\gamma$-ray emission ($\sim 100$~GeV) would be attenuated at early times (i.e. $<10$~d) due to internal $\gamma \gamma$ absorption. The optical SN emission, which typically peaks several ($\gtrsim 10$ d) days after the shock break out, may also attenuate the $\sim$100 GeV emission. We plan to include this in our numerical code by modelling the SN optical radiation as an external photon field of variable photon number density. Another significant process for producing $\gamma$-rays in such dense environments is the bremsstrahlung radiation of relativistic electrons. Although the results presented here do not include relativistic bremsstrahlung radiation, we plan to implement it in the code both as an energy loss process for electrons and a photon production process. We have also verified that the effect of $p\gamma$ interactions upon the synchrotron photons is negligible. Yet, $p\gamma$ interactions with the X-ray bremsstrahlung photons produced by the hot shocked plasma may be important for creating additional secondaries, especially at early times when the X-ray luminosity is high and the source is compact. \vspace{-0.2in} \small % | 16 | 7 | 1607.05847 |
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1607 | 1607.02460_arXiv.txt | \noindent All measurements are comparisons. The only physically accessible degrees of freedom (DOFs) are dimensionless ratios. The objective description of the universe as a whole thus predicts only how these ratios change collectively as one of them is changed. Here we develop a description for classical Bianchi IX cosmology implementing these relational principles. The objective evolution decouples from the volume and its expansion degree of freedom. We use the relational description to investigate both vacuum dominated and quiescent Bianchi IX cosmologies. In the vacuum dominated case the relational dynamical system predicts an infinite amount of change of the relational DOFs, in accordance with the well known chaotic behaviour of Bianchi IX. In the quiescent case the relational dynamical system evolves uniquely though the point where the decoupled scale DOFs predict the big bang/crunch. This is a non-trivial prediction of the relational description; the big bang/crunch is not the end of physics - it is instead a regular point of the relational evolution. Describing our solutions as spacetimes that satisfy Einstein's equations, we find that the relational dynamical system predicts two singular solutions of GR that are connected at the hypersurface of the singularity such that relational DOFs are continuous and the orientation of the spatial frame is inverted. | The description of the universe as a whole can not depend on external units of length or duration since all physical clocks and rods are part of the universe itself. The universe differs fundamentally from its subsystems in this aspect. The problem with GR's distance defining spacetime geometry is thrown into focus when considering cosmology. In a laboratory experiment an observer can easily justify the separation between the measuring apparatus (the external clock and rod) and the objects being measured. Cosmologists however are part of the universe and can not separate themselves form the studied system. Cosmological measurements are intrinsic, as rods and clocks are constructed from the dynamical objects in the universe. Units are constructed intrinsically using physical reference structures that define GR's notion of geometry. Hence, all dimensional quantities are intrinsic ratios. This leads us to concentrate on the dynamics of \textit{relational variables,} \textit{i.e.} dimensionless ratios and their relative infinitesimal variations\footnote{\textit{E.g.} the scale factor is \textit{not} a relational variable, while Misner's anisotropy parameters $\beta_+$, $\beta_-$~\cite{MTW} and the ratio of their variations $d\beta_+/d\beta_-$ are relational variables.} \cite{Tutorial}. Remarkably, the dynamics can be expressed entirely in terms of relational variables, which turn out to evolve autonomously and predict all \textit{intrinsic observables} of GR. In this letter we study spatially homogeneous cosmology (Bianchi IX) with a massless scalar field. This model is believed to correctly describe the near-singularity behaviour of full GR due to the BKL conjecture and Wheeler's insight that ``matter doesn't matter'' except a stiff component (such as a massless scalar field). In fact, the theorems of \cite{AR}, show that a dense set of inhomogeneous GR solutions obey the BKL conjecture, \textit{i.e.} spatial points decouple in the approach to the singularity and evolve as independent Bianchi IX systems. Moreover, a massless scalar field is compatible with Standard Model physics\footnote{Interestingly, an RG improved gravitational action, as obtained in the functional renormalization group setting, also offers a mechanism to achieve this quiescent behavior~\cite{FrankGiulio}.} (\textit{e.g.} the Goldstone mode of the Higgs field). We thus distinguish two cases:\\ 1. In absence of a massless scalar the approach to the singularity is given by the vacuum Bianchi IX evolution, in which the dynamics never actually reaches the singularity. It rather goes through an infinite amount of change, with infinitely many billiard-like `bounces' against steep triangular potential walls, alternating with intervals of free geodesic evolution (Kasner epochs). This fact was observed by Misner~\cite{Misner_Mixmaster}, and its consequences for the status of the singularity was discussed in~\cite{MTW}. This has an important consequence in the relational framework, where physical clocks necessarily possess internal relational DOFs which register time. An infinitesimal clock can not be treated as the idealized worldline of a point with its proper time, but has to be viewed as the infinitesimal limit of a sequence of ever smaller time-recording systems with internal structure \cite{Ohanian}. It has been noted \cite{Ohanian} that the change of the internal relational DOFs of an infinitesimal limit clock will be subject to the {\textit same} tidal effects as measured by its large counterparts. It follows that infinitesimal clocks register an unbounded lapse of time (\textit{i.e.} change of internal relational DOFs), when the gravitational field experiences an infinite number of Kasner epochs. It follows that the infinitesimal clocks, unlike their pointlike idealizations \textit{i.e.} proper time, will not reach the big bang/crunch in a finite time. \\ 2. In the presence of a massless scalar one experiences ``quiescent" behaviour~\cite{Tutorial}. The potential becomes irrelevant for the dynamics and the equations of motion asymptote into a geodesic evolution. Matter clocks will measure a finite amount of change between the singularity and any other point. It remains to investigate this case, because it is the one in which the singularity is reached in finite relational time and we have to establish what happens to the relational DOFs there. | We showed that the relational description of quiescent Bianchi IX universes evolves through the big bang, which is not a singularity of the equations on shape space. We will therefore call it a `Janus point', because it is the point of qualitative time-symmetry of each solution~\cite{ArrowPaper,EntropyPaper,JanusPointAnswerToZeh}. The Janus-point data, \textit{i.e.} the specification of $(\alpha,\gamma,\omega,\sigma)$ at $\beta=\frac{\pi}{2}$, determines a unique curve on the two hemispheres of shape space (with a single intersection with the equator) that can be effectively described as two quiescent Bianchi IX spacetimes glued together at the big bang, where a change of orientation occurs. The prediction of a classical change of orientation of the spatial manifold at the big bang could have profound implications for discrete symmetries in particle physics, particularly regarding matter/antimatter asymmetry. In vacuum dominated Bianchi IX (\textit{i.e.} in absence of massless scalars) on the other hand, one finds that physical time (measured by the change of shape of any finite clock) will go on forever in the relational description of vacuum-dominated Bianchi IX cosmology. These curves do not terminate in shape space. As such an intrinsic observer will never encounter any singularity. We established that big bang/crunch singularities in homogeneous cosmologies with compact topology are spacetime artifacts, that do not have any physical (\textit{i.e.} completely relational) meaning. This has important consequences for a dense set of GR solutions (those described in~\cite{AR}) whose behaviour near the cosmological singularity is completely described by an independent Bianchi IX universe at each point. Moreover, the BKL conjecture posits that this behaviour is generic in GR. Finally, this result allows discussing the typicality of universes in terms of their Janus point data, and the spontaneous emergence of an arrow of time (along the lines of \cite{ArrowPaper,EntropyPaper,JanusPointAnswerToZeh}). | 16 | 7 | 1607.02460 |
1607 | 1607.06075_arXiv.txt | We study the correlation of galaxy structural properties with their location relative to the SFR-$M_{*}$ correlation, also known as the star formation ``main sequence'' (SFMS), in the CANDELS and GAMA surveys and in a semi-analytic model (SAM) of galaxy formation. We first study the distribution of median S{\'e}rsic index, effective radius, star formation rate (SFR) density and stellar mass density in the SFR-$M_{*}$ plane. We then define a redshift dependent main sequence and examine the medians of these quantities as a function of distance from this main sequence, both above (higher SFRs) and below (lower SFRs). Finally, we examine the distributions of distance from the main sequence in bins of these quantities. We find strong correlations between all of these galaxy structural properties and the distance from the SFMS, such that as we move from galaxies above the SFMS to those below it, we see a nearly monotonic trend towards higher median S{\'e}rsic index, smaller radius, lower SFR density, and higher stellar density. In the semi-analytic model, bulge growth is driven by mergers and disk instabilities, and is accompanied by the growth of a supermassive black hole which can regulate or quench star formation via Active Galactic Nucleus (AGN) feedback. We find that our model qualitatively reproduces the trends described above, supporting a picture in which black holes and bulges co-evolve, and AGN feedback plays a critical role in moving galaxies off of the SFMS. | Out to $z\sim3$, galaxies can be split into star-forming and quiescent populations based on the bimodality observed in their colors and derived star formation rates \citep{Baldry2004, Bell2004b, Brinchmann2004, Kauffmann2003, Strateva2001, Brammer2011, Ilbert2013}. When focusing specifically on the galaxies classified as star-forming, a strong correlation is observed between the star formation rate and stellar mass of galaxies at a fixed redshift (the SFR-$\rm{M_{*}}$ correlation) \citep{Noeske2007, Daddi2007, Elbaz2007, Rodighiero2011}. This correlation is also sometimes referred to as the ``star-forming main sequence'' (SFMS). This stands in contrast to the less rigidly defined quiescent population, for which there is no such strong correlation. The SFR-$\rm{M_{*}}$ correlation can be defined by a (redshift-dependent) normalization and slope, with a straight line in log-log space providing a reasonable fit, although there is evidence that the slope of the main sequence may flatten above a mass of $\sim10^{10}M_{\odot}$ \citep{Whitaker2012, Whitaker2014}. It is still unclear whether this flattening is simply due to the fact that more of the stellar mass in high mass galaxies is likely to be in a non-star-forming bulge component, as suggested by \citet{Abramson2014} or \citet{Tacchella2015}, or whether there is something else going on. It has also been suggested that the presence of non star-forming bulges in star-forming galaxies may increase the scatter in the SFR-$\rm{M_{*}}$ relation around the main sequence \citep{Whitaker2015}. In any case, many studies have examined the SFR-$\rm{M_{*}}$ correlation and found that it holds over at least four orders of magnitude in mass and exists out to $z\sim6$ (see \citet{Speagle2014} and references therein, as well as \citet{Salmon2015}). The value of the slope in the SFR-$\rm{M_{*}}$ plane is measured to be $\sim1$ \citep{Rodighiero2011} and the relationship has an intrinsic $1-\sigma$ scatter of only $\sim0.2-0.4$ dex \citep{Whitaker2012, Kurczynski2016}. In general, SFMS galaxies at high redshift have much higher SFRs than galaxies on the main sequence today \citep{Sobral2014}, and the evolution of the normalization of the SFMS appears to be independent of galaxy environment \citep{Peng2010}. The small scatter of the SFR-$\rm{M_{*}}$ correlation leads us to believe that galaxy evolution is dominated by relatively steady star formation histories, rather than being highly stochastic and bursty. This places constraints on the duty cycle of processes such as galaxy mergers or disk instabilities, which may trigger starburst and quenching events that drive galaxies above or below the main sequence. Furthermore, observations show that since $z\sim 2$ there has been a build-up of quiescent galaxies, while the mass density of galaxies on the SFMS has remained relatively constant, implying that galaxies are being moved \emph{off} of the SFMS into the quiescent population, and remaining there permenantly or at least over rather long timescales \citep{Bell2004b, Borch2006, Bell2007, Faber2007}. As the processes which move galaxies off of the main sequence are often associated with morphological change, it is interesting to examine the correlation between distance from the SFR-$\rm{M_{*}}$ relation, or some other measure of quiescence, and galaxy structural properties. \citet[][B15]{Brennan2015} defined a redshift dependent SFMS by which to judge galaxies in order to divide them into star-forming and quiescent populations. We split the sSFR-S{\'e}rsic index plane into four quadrants in star-formation activity and morphology: star-forming disk-dominated galaxies, star-forming spheroid-dominated galaxies, quiescent disk-dominated galaxies, and quiescent spheroid-dominated galaxies. After dividing galaxies up, we examined the evolution of the fraction of galaxies in each of these categories with redshift. In order to constrain which processes were responsible for moving galaxies between these different categories, we did the same analysis on a sample of model galaxies generated from the ``Santa Cruz'' semi-analytic model described in \citet{Somerville2008} with updates as described in \citet{Somerville2012} and \citet{Porter2014}. In addition to prescriptions for the main physical processes believed to be important for shaping galaxy properties (described below), the model includes bulge formation due to mergers and disk instabilities, and concurrent growth of supermassive black holes and AGN feedback, allowing us to predict how model galaxies evolve in the SFR-S{\'e}rsic index plane. The SAM is a useful tool for studying the evolution of large populations of galaxies, as it can generate large cosmologically representative samples with modest computational resources, allowing us to efficiently test the effects of various physical processes. In B15, we found that our prescriptions for quenching and morphological transformation were able to transform galaxies in a manner in qualitative agreement with the observations as long as bulge growth due to disk instabilities was included. Bulge growth due to mergers and disk instabilities and subsequent AGN feedback produced roughly the right fraction of galaxies in each of our four subpopulations. Models in which bulge growth occured only due to mergers did not produce as many spheroid-dominated galaxies as seen in observations. Our goal in this paper is to study the structural properties of model galaxies \emph{continuously} across and off the main sequence, rather than using the main sequence to sort our galaxies into bins based on their SFRs and morphologies as in B15 and Pandya et al. (in prep.). The latter explicitly examines galaxies with intermediate star-formation and structural properties. We learned in B15 that our model could broadly produce the right fractions of different types of galaxies and the evolution of these fractions, and now we will examine more closely if it can produce both ``typical'' main sequence galaxies, as well as match how the structural properties of galaxies change as they move farther from the main sequence. In this way, we hope to continue to build our understanding of the physical processes which drive the correlation between star formation, quenching, and galaxy structural properties. Many observational studies have examined the structure of galaxies | In this work, we have investigated the correlation of galaxy structural properties with their location in the plane of star formation rate and stellar mass. We studied structural properties such as morphology as represented by S{\'e}rsic index, radial size, and mean stellar surface density as a continuous function of a galaxy's distance from the mean star forming main sequence at its observation time. We carried out a parallel analysis on the GAMA survey of nearby galaxies, the CANDELS survey which can measure galaxy structural properties to $z\sim 3$, and a semi-analytic model that tracks the evolution of galaxy properties within a cosmological framework. We focus on the population of galaxies with stellar mass >$10^{10}M_{\odot}$, for which these surveys are highly complete and the measurement of structural properties is robust. Our main findings are as follows: \begin{itemize} \item Within $\pm 0.5$ dex of the SFMS, we find a weak dependence of galaxy structural properties on the distance from the MS. Below the main sequence, we see a rapidly steepening dependence such that galaxies with larger negative MS residuals had higher median S{\'e}rsic index, smaller size, and higher stellar surface density. These trends are seen in both nearby galaxies (GAMA) and out to $z\sim 2.5$ (CANDELS), and are qualitatively very similar in the theoretical models. \item Our observational results are very similar overall to the results of an earlier study by \citet[][W11]{Wuyts2011}. One difference between our results and those of W11 is that we do not find a significant population of galaxies with high S{\'e}rsic index ($n\sim 3.5$--4) in the extreme starburst region above the SFMS. Similarly, we do not see as large a population of galaxies with small radii above the SFMS. We suspect that these galaxies are removed from our sample due to our requirement of being well fit by a single component S{\'e}rsic profile. \item The good qualitative agreement between our model results and the observations suggests a plausible causal explanation for the observed correlations; namely, that central spheroids and black holes grow together, and black holes play a major role in quenching star formation in galaxies. \item Quantitatively, our models disagree with the observations in some important respects. Our models do not produce as large a quiescent population at high redshift ($z>1.5$) as seen in the observations (as already noted by B15), and the SFR for the model quiescent galaxies are lower than those of observed quiescent galaxies. This suggests the need to refine our modeling of AGN feedback. Moreover, the S{\'e}rsic indices of galaxies below the SFMS are systematically lower (more disk-like) in the models, while on and below the SFMS, especially at low redshift, the sizes of our galaxies are too large. As a result, there is not as large a separation between the sizes for the star forming and quiescent populations in the models as what is seen in the observations. This suggests that we also need to refine our determination of galaxy sizes in the model. \end{itemize} | 16 | 7 | 1607.06075 |
1607 | 1607.06769_arXiv.txt | We present an analysis of a deep (1$\sigma$=13\,$\mu$Jy) cosmological 1.2-mm continuum map based on ASPECS, the ALMA Spectroscopic Survey in the Hubble Ultra Deep Field. In the 1\,arcmin$^2$ covered by ASPECS we detect nine sources at $>3.5\sigma$ significance at 1.2-mm. Our ALMA--selected sample has a median redshift of $z=1.6\pm0.4$, with only one galaxy detected at z$>$2 within the survey area. This value is significantly lower than that found in millimeter samples selected at a higher flux density cut-off and similar frequencies. Most galaxies have specific star formation rates similar to that of main sequence galaxies at the same epoch, and we find median values of stellar mass and star formation rates of $4.0\times10^{10}\ M_\sun$ and $\sim40~M_\sun$ yr$^{-1}$, respectively. Using the dust emission as a tracer for the ISM mass, we derive depletion times that are typically longer than 300\,Myr, and we find molecular gas fractions ranging from $\sim$0.1 to 1.0. As noted by previous studies, these values are lower than using CO--based ISM estimates by a factor $\sim$2. The 1\,mm number counts (corrected for fidelity and completeness) are in agreement with previous studies that were typically restricted to brighter sources. With our individual detections only, we recover $55\pm4\%$ of the extragalactic background light (EBL) at 1.2\,mm measured by the {\it Planck} satellite, and we recover $80\pm7\%$ of this EBL if we include the bright end of the number counts and additional detections from stacking. The stacked contribution is dominated by galaxies at $z\sim1-2$, with stellar masses of (1--3)$\times$10$^{10}$\,M$_\odot$. For the first time, we are able to characterize the population of galaxies that dominate the EBL at 1.2\,mm. | One of the most fundamental discoveries with regard to the cosmic evolution of galaxies has been the determination that a substantial fraction of the integrated Extragalactic Background Light (EBL) arises at infrared-to-millimeter wavelengths: the Cosmic Infrared Background (CIB). Quantitative observations of the CIB began with the Cosmic Background Explorer (COBE). At a low angular resolution ($0.7\deg$), COBE provided the first large-scale measurement of the spectral energy distribution (SED) of the EBL from the far-infrared to the (sub)millimeter \citep{puget96,fixsen98}. The CIB consists of the combined flux of all extragalactic sources, and contains much information about the history and formation of galaxies, and of the large scale structure of the Universe. The observation that the cosmic density of star-formation was an order of magnitude higher at cosmological redshifts, $z\sim2-4$ \citep[e.g.,][]{madau96, lilly96}, opened the possibility that most of the CIB arose from dust re-processed UV-light from distant galaxies. These studies used the Lyman dropout technique to identify normal galaxies at high-redshift, being mostly insensitive to dust obscured star formation. Later, sensitive maps obtained with submillimeter/millimeter bolometer arrays were thus able to directly detect and identify luminous dusty star forming galaxies (DSFGs), which were soon found to contribute a fraction to the EBL at these wavelengths \citep[e.g.,][]{smail97}. Since then, a number of groups have conducted (sub)millimeter surveys of the sky, currently yielding up to hundreds of sources in contiguous areas of the sky \citep[e.g.,][]{hughes98, barger98, eales00, bertoldi00, scott02, cowie02, voss06, bertoldi07,scott08,greve08, weiss09,austermann10,vieira10,aretxaga11,hatsukade11,scott12,mocanu13}. These blank field bolometer (sub)millimeter surveys discovered a population of luminous DSFGs at high redshift that were not accounted for in optical studies. These galaxies -- also called ``submillimeter galaxies'' (SMGs) due to the region of the electromagnetic spectrum in which they were first discovered -- have been characterised as massive starburst galaxies with typical stellar and molecular gas masses of $\sim10^{11}\ M_\sun$, typically located at $z=1-3$ \citep[e.g.,][]{chapman05} with a tail out to $z\sim6$ \citep{weiss13, riechers13}, and most likely driven by relatively bright mergers \citep{engel10}. As such, these galaxies are found to be gas/dust rich, with gas fractions typically exceeding 0.2 \citep[e.g.][]{daddi10a, tacconi10, magdis12, tacconi13,bothwell13}. Despite their large SFRs implied by the large IR luminosities ($>10^{12.0-12.5}\ L_\sun$) and significant abundance at high-redshift, these galaxies (e.g. $S_{\rm 1.2mm}>2-3$ mJy) were found to contribute only a minor fraction of the EBL at submillimeter wavelengths \citep{barger99,eales99,smail02,coppin06,knudsen08,weiss09,scott12, chen13}. Hence, questions about the properties of the population of galaxies that dominate this EBL remain. \begin{figure*}[ht] \centering \includegraphics[scale=0.55]{UDF_1mm_continuum_image_v4.ps}\hspace{4mm} \includegraphics[scale=0.55]{UDF_1mm_sensitivity_image_v4.ps} \caption{({\it Left:}) ALMA 1.2-mm signal-to-noise continuum mosaic map obtained in the HUDF. Black and white contours show positive and negative emission, respectively. Contours are shown at $\pm2,3,4,5,8,12,20$ and $40\sigma$, with $\sigma=12.7\mu$Jy beam$^{-1}$ at the field center. The boxes show the position of the sources detected with our extraction procedure at $S/N>3.5$. The synthesized beam ($1''\times2''$) is shown in the lower left. ({\it Right:}) ALMA 1.2-mm observations primary beam (PB) pattern to represent the sensitivity obtained across the covered HUDF region. PB levels are shown by the black/white contours at levels 0.3, 0.5, 0.7 and 0.9 of the maximum. Both the signal-to-noise and PB maps are shown down to PB$=0.2$.\label{fig_map1mm}} \end{figure*} \begin{figure*}[ht] \centering \includegraphics[scale=0.55]{UDF_3mm_continuum_image_v4.ps}\hspace{4mm} \includegraphics[scale=0.55]{UDF_3mm_sensitivity_image_v4.ps} \caption{({\it Left:}) ALMA 3-mm signal-to-noise continuum mosaic map obtained in the HUDF. Black and white contours show the positive and negative signal, respectively. Contours are shown at $\pm2,3,4,5,8,12,20$ and $40\sigma$, with $\sigma=3.8\mu$Jy beam$^{-1}$ at the field center. The boxes show the position of the sources detected in the 1.2-mm map, with our extraction procedure at $S/N>3.5$. The synthesized beam ($2''\times3''$) is shown in the lower left. ({\it Right:}) ALMA 3-mm observations primary beam (PB) pattern. PB levels are shown by the black/white contours at levels 0.3, 0.5, 0.7 and 0.9. Both the signal-to-noise and PB maps are shown down to PB$=0.2$.\label{fig_map3mm}} \end{figure*} To locate and characterise the population of faint DSFGs that make up most of the EBL at (sub)millimeter wavelengths, we must overcome several observational limitations. First, the poor resolution of (sub)millimeter bolometer maps taken with single-dish telescopes, typically with beam sizes between $10-30''$, makes the identification of an optical counterpart difficult and thus limits the characterisation of submillimeter sources. In addition, this affects the number counts, since the brightest sources are seen to split into multiple components in high-resolution (sub)millimeter images \citep{younger07, wang11, smolcic12, hodge13, karim13, miettinen15}. Secondly, the sensitivity of single dish bolometer maps, typically down to $0.5-1.0$ mJy, along with confusion at the faint levels limits our view to the most luminous sources. An important approach to reach fainter galaxies has been the use of gravitational lensing enabled by massive galaxy clusters \citep[e.g.,][]{smail97, smail02, sheth04,knudsen08,noble12,johansson12, chen13}. However, these surveys suffer severely from cosmic variance, due to the small areas covered in the source plane, source confusion, and the need for accurate lens models and magnification maps. A parallel approach has been to perform stacking of the submillimeter emission using pre-selected samples of optical/infrared galaxies. This approach has successfully resolved significant amounts of the EBL at (sub)millimeter wavelengths, reaching down to sources with $S_{\rm 1.2mm}>0.1$ mJy \citep{webb04,knudsen05,greve10,decarli14}. The major limitation of this approach is that it yields average properties over a population of galaxies that must be assumed to have similar (sub)millimeter properties. The advent of the Atacama Millimeter/submillimeter Array (ALMA) is opening up a new window for the study of the faint DSFG population. Its significantly higher angular resolution compared to single-dish telescopes ($<3''$), and the unparalleled sensitivity allow us to reach flux density levels in (sub)millimeter continuum maps even deeper than those achieved by studies of galaxy cluster fields or based on stacking analysis. Several recent studies have individually pinpointed (sub)millimeter sources down to 0.1 mJy in the 1-mm band \citep{hatsukade13,ono14,carniani15,oteo15,hatsukade16, dunlop16}. Some of these surveys have used clever approaches by taking advantage of archival data \citep{ono14,carniani15,fujimoto16}, including ALMA calibration fields \citep{oteo15}. Recently, \citet{fujimoto16} were able to reach down to a flux limit of 15$\mu$Jy at 1.2-mm, providing the deepest measurements of the number counts to date, and allowing them to resolve most of the CIB into individual sources. Despite the substantial progress, the current studies are still affected significantly by cosmic variance and are not ``blank-field'' in nature (as some of them target overdense fields). Most importantly, the lack of sufficiently deep complementary data have only permitted the characterisation of a handful of sources \citep{hatsukade15,fujimoto16, yamaguchi16}. Using ALMA in Cycle 2, we have conducted a deep ALMA Spectroscopic Survey (ASPECS) of a region of the {\it Hubble} Ultra Deep Field (UDF), covering the full 3-mm and 1-mm bands. In this paper, we present the {\it deepest} millimeter continuum images obtained to date in a contiguous 1 arcmin$^2$ area. This is the Paper~II in the ASPECS series. A full description of the survey and spectral line search is presented in Paper~I \citep{walter16}. Measurements of the CO luminosity function and cosmic density of molecular gas are shown in Paper~III \citep{decarli16a}. A detailed analysis of the CO brightest objects is presented in Paper~IV \citep{decarli16b}. A search for {\sc [CII]} line emission is shown in Paper~V \citep{aravena16b}. This paper is organised as follows: in \S \ref{sec_obs}, we summarise the ALMA observations and multi-wavelength ancillary data available. Here, we also present the obtained ALMA continuum maps at 1.2-mm and 3-mm. In \S \ref{sec_results}, we present the detected sources and compute the fidelity and completeness of our extraction procedures in the 1.2-mm map. In \S \ref{sec_counts}, we derive the number counts at 1.2-mm. In \S \ref{sec_prop}, we characterise the multi-wavelength properties of the individually detected sources, including their typical stellar masses, SFRs and redshifts, and discuss whether our sources are starbursts or more quiescent star forming galaxies. In \S \ref{sec_stack}, we conduct a stacking analysis to determine the average properties of the faintest population of galaxies, not detected individually by our survey. In \S \ref{sec_ism}, we investigate the ISM properties of the individually detected sources based on measurements of the ISM masses from the 1.2-mm fluxes. We estimate their gas masses, depletion timescales and fractions. In \S \ref{sec_cib}, we determine the contribution of both our individually-detected and stacked sample to measure the fraction of the EBL at 1.2-mm resolved by our observations. We discuss the properties of the galaxies that dominate the CIB. Finally, in \S \ref{sec_concl}, we summarise the main results of this paper. Throughout the paper, we assume a standard $\Lambda$CDM cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_\Lambda=0.7$ and $\Omega_{\rm M}=0.3$. | Using ALMA in cycle-2, we have conducted a millimeter spectroscopic survey by scanning the full 3-mm and 1.2-mm bands over a region in the {\it Hubble} UDF. The collapsed cubes constitute the deepest continuum images ever obtained over an 1 arcmin$^2$ contiguous area of the sky. The main results of our continuum measurements can be summarised as follows: \begin{itemize} \item We detect nine sources with significances $>3.5\sigma$ at 1.2-mm and only one source at 3-mm. From these detections, we measure the 1.2-mm number counts over the flux density range $S_{\rm 1.2mm}=0.036-0.57$ mJy. Our number counts are similar to previous measurements, with differences within a factor of $\sim2$. \item We measure the properties of the individually detected galaxies at S/N$>3.5$. We find that there is a large spread in stellar masses and SFRs, with median values of $4\times10^{10}\ M_\sun$ and $\sim40~M_\sun$ yr$^{-1}$, much lower than found in brighter SMGs. We find that these faint DSFGs are systematically located at lower redshifts than millimeter-selected SMGs, with a median redshift of $z=1.7$. All galaxies are consistent with being close to the main sequence at their respective redshift. \item We use stacking analysis to estimate the average emission from samples of galaxies selected by redshift, stellar mass and SFRs. We only find detections in samples selected in the redshift range $1<z<2$, as well as in the stellar mass ranges log$(M_*/M_\sun)=9.5-10.0$ and log$(M_*/M_\sun)=10.0-10.5$, with typical SFRs of $3-10\ M_\sun$ yr$^{-1}$ . This suggests that the rest of the emission, not individually detected in our survey, comes from galaxies less massive, with lower SFRs, but at a similar redshift than the detected sources. \item We use the 1.2-mm flux as a proxy for the ISM masses in our individually detected galaxies. We find that most of our sources are located in the star-forming trend occupied by main-sequence galaxies and local spirals, implying relatively large gas time depletion timescales, typically above 300 Myr, and a large spread in the molecular gas fractions ranging from 0.1 to 1.0. We compare these results to ISM mass estimates using CO as a tracer in \citet[][; Paper~IV]{decarli16b}. \item Our individual detections alone are able to resolve $55\pm4\%$ of the EBL at 242 GHz measured by the {\it Planck} satellite. By adding up the integrated intensity from our number counts, to the contribution from the bright end of the number counts -- mostly composed by SMGs -- and the contribution of faint galaxies detected using stacking, we are able to resolve between 77--84\% of the CIB at 242 GHz. The typical properties of the population that makes up most of the EBL at these frequencies corresponds to that of the galaxies described in this work. \end{itemize} | 16 | 7 | 1607.06769 |
1607 | 1607.03699_arXiv.txt | {The current generation of X-ray instruments is progressively revealing more and more details about the complex magnetic field topology and the geometry of the accretion flows in highly magnetized accretion powered pulsars. } % {We took advantage of the large collecting area and good timing capabilities of the EPIC cameras on-board \xmm\ to investigate the accretion geometry onto the magnetized neutron star hosted in the high mass X-ray binary \exo\ during the rise of a source Type-I outburst in 2014.} % {We carried out a timing and spectral analysis of the \xmm\ observation as function of the neutron star spin phase. We used a phenomenological spectral continuum model comprising the required fluorescence emission lines. Two neutral absorption components are present: one covering fully the source and one only partially. The same analysis was also carried out on two \suzaku\ observations of the source performed during outbursts in 2007 and 2012, to search for possible spectral variations at different luminosities. } % {The \xmm\ data caught the source at an X-ray luminosity of $2\times10^{36}$\,erg\,s$^{-1}$ and revealed the presence of a narrow dip-like feature in its pulse profile that was never reported before. The width of this feature corresponds to about one hundredth of the neutron star spin period. From the results of the phase-resolved spectral analysis we suggest that this feature can be ascribed to the self-obscuration of the accretion stream passing in front of the observer line of sight. We inferred from the \suzaku\ observation carried out in 2007 that the self-obscuration of the accretion stream might produce a significantly wider feature in the neutron star pulsed profile at higher luminosities ($\gtrsim$$2\times10^{37}$\,erg\,s$^{-1}$).} % {This discovery allowed us to derive additional constraints on the physical properties of the accretion flow in this object at relatively small distances from the neutron star surface. The presence of such a narrow dip-like feature in the pulse profile is so far unique among all known high mass X-ray binaries hosting strongly magnetized neutron stars.} % | \label{sec:intro} \exo\ is a prototypical high mass Be X-ray binary (BeXRB), comprising a neutron star (NS) and a Be companion. X-ray outbursts are generally produced when the orbit of the pulsar intercepts the companion's decretion disk. Be/X-ray binaries typically show two types of outbursts: (1) giant outbursts (type II), which lasts tens of days and are characterized by high luminosities and high spin-up rates (i.e., a significant increase in pulse frequency), and (2) normal outbursts (type I), which are characterized by lower luminosities and less pronounced spin-up rates (if any). Type I outbursts are known to occur (almost) regularly at each periastron passage in some systems \citep{stella1986, bildsten1997}. During these events, the material lost by the Be star is first focused toward the NS as a consequence of its strong gravitational field, and then funneled by its intense magnetic field ($B\sim10^{12-13}$\,G) down to the magnetic poles, where one or more accretion columns are formed. The bulk of the continuum X-ray emission from BeXRBs is produced within the accretion columns due to the Compton scattering of seed thermal photons from the hot spot on the NS surface or by bremsstrahlung processes occurring along the column \citep[][and references therein]{becker2007}. The spectral energy distribution of these sources is expected to show a remarkable dependence on the spin phase due to the changes in the viewing angle of the observer and the angular dependence of the Compton scattering cross section in a strong magnetic field \citep{meszaros1985a,meszaros1988b}. Iron fluorescence lines corresponding to different ionization levels of these heavy ions are commonly observed in BeXRBs and ascribed to the illumination of the accreting material at different distances from the NS by the intense X-ray radiation. The broad band spectra of BeXRBs are usually described by different phenomenological models, the most widely used one being an absorbed power-law modified at high energy by an exponential cut-off. Depending on the statistics of the data and the energy coverage, it proved necessary in several cases to complement these relatively simple spectral models with the additions of broad Gaussian components \citep[see, e.g.,][]{klochkov2007, suchy2008} and/or partial covering absorbers \citep{Naik2011,Naik2012,Naik2014}. Furthermore, the cyclotron resonant scattering of electrons in the high magnetic field of the NS is known to produce characteristic absorption lines that have been observed in several of these systems and included in the spectral fits by using either Gaussian or Lorentzian profiles \citep[see, e.g.,][for a recent review]{walter15}. Cyclotron Resonant Scattering Features (CRSFs) can be used to infer the NS surface magnetic field strength as the centroid energy of the fundamental line is at $E_{\rm cyc} \simeq 11.6\,B_{12}\times (1+z)^{-1}$\,keV, where $B_{12}$ is the NS magnetic field strength in units of $10^{12}$\,G and $z$ the redshift of the scattering medium. \exo hosts a 42\,s pulsar discovered with \textsl{EXOSAT} during a giant type-II outburst in 1985 \citep{parmar1989b}. The compact object orbits a B0\,Ve star \citep[][and references therein]{coe1988} every 46 days \citep{wilson2005,wilson2008}. The estimated distance to the source is 7.1\,kpc \citep{wilson2002}. Type-I outbursts reaching peak fluxes of about 100\,mCrab (15-50\,keV) have been regularly detected from \exo at virtually all periastron passages since 1991 \citep[the outburst peak usually occurs about $\sim7$\,d after the periastron passage; see, e.g.,][]{wilson2005}. In the period spanning from 1992 to 1994, the type-I outbursts have been brighter than average and the NS showed a remarkable spin-up. From 1994 to 2002, the outbursts showed somewhat lower peak luminosities (a factor of few) and the pulsar displayed a clear spin-down trend. A new re-brightening period followed until June 2006, when \exo underwent its second observed giant type-II outburst. After a number of binary orbits characterized by a higher persistent luminosity than average and type-I outburst achieving a peak flux of 200-300\,mCrab (15-50\,keV), the source returned back to its normal behavior. During type-I and type-II outbursts, the broad-band spectrum of \exo\ can be reasonably well described by using an absorbed ($N_\mathrm{H}\simeq10^{22}\,\mathrm{cm}^{-2}$) power-law with a high-energy exponential roll-over. Contrasting results have been published concerning the presence of possible CRSFs in the X-ray emission from the source. \citet{reig1999} reported on the possible detection of such a feature with a centroid energy of $\sim36$\,keV, while \citet{wilson2008} found evidence of a CRSF at $\sim11$\,keV. \citet{klochkov2007} showed, however, that the data used by \citet{wilson2008} could also be reasonably well characterized without the cyclotron line by including in the fit a broad Gaussian emission feature at $\sim15$\,keV \citep[as observed in other high mass X-ray binaries; see, e.g., the discussion in][and references therein]{ferrigno2009}. \citet{klochkov2008} also reported on the detection of a CRSF at $\sim63$\,keV in spin-resolved spectra extracted during the peak of the 2006 giant outburst. This could correspond to a higher harmonics of the previously suggested cyclotron line at $\sim36$\,keV. The pulse profile of \exo is known to be strongly dependent on the X-ray luminosity. At the peak of the outbursts its characteristic shape is usually interpreted in terms of a fan beam-like emission, while it become more reminiscent of what is expected in the case of a pencil beam emission during the decay of the outburst \citep{parmar1989b,klochkov2008}. A similar interpretation was also suggested by the detailed study carried out by \citet{sasaki2010} with a pulse decomposition method. In this paper, we report on the first \xmm\ observation of \exo performed during the rise of a type-I outburst in 2014. The large collecting area and good timing resolution of the EPIC-pn camera on-board \xmm\ allowed us to extract the source pulse profiles with more than 100 phase bins, revealing a peculiar sharp and deep feature never detected before (Sect.~\ref{sec:timing}). Following our spectral results (Sect.~\ref{sec:spectral}), we suggest that this feature is caused by the obscuration effect of the accretion stream passing in front of the observer line of sight to the source (Sect.~\ref{sec:phase}). The implications of our results are discussed in Sect.~\ref{sec:discussion} and summarized in Sect.~\ref{sec:conclusion}. | \label{sec:conclusion} The discovery of a sharp and narrow dip-like feature in the soft X-ray pulse profile of \exo\ during a typical type-I X-ray outburst allowed us to have a novel insight on the physical properties of the accretion flow in this object. The presence of such feature is so far unique among all known high mass X-ray binaries hosting strongly magnetized stars (at the best of our knowledge). Further investigations on other similar systems with X-ray instruments endowed with a good timing resolution and large effective areas at $\sim$10\,keV could complement the existing legacy data, typically available only at higher luminosities and harder X-rays, and allow us to probe in more details the accretion geometry and magnetic field topology of these systems. A model comprising a phenomenological Comptonization continuum and a combination of homogeneous and inhomogeneous absorbers is shown to provide a reasonably good fit to the \xmm\ data of \exo,\ as well as to the broader energy coverage spectra of the source provided by the \suzaku observations. Exploiting the suitability of such model to fit the X-ray spectra of other BeXRBs will permit to possibly achieve a more homogeneous description of the high-energy emission from these sources. % | 16 | 7 | 1607.03699 |
1607 | 1607.03520_arXiv.txt | We have imaged GM~Aurigae with the Hubble Space Telescope (HST), detected its disk in scattered light at 1400~\AA\ and 1650~\AA, and compared these with observations at 3300~\AA, 5550~\AA, 1.1$\mu$m, and 1.6$\mu$m. The scattered light increases at shorter wavelengths. The radial surface brightness profile at 3300~\AA\ shows no evidence of the 24~AU radius cavity that has been previously observed in sub-mm observations. Comparison with dust grain opacity models indicates the surface of the entire disk is populated with sub-$\mu$m grains. We have compiled an SED from 0.1~$\mu$m to 1~mm, and used it to constrain a model of the star~+~disk system that includes the sub-mm cavity using the Monte Carlo Radiative Transfer (MCRT) code by Barbara Whitney. The best-fit model image indicates that the cavity should be detectable in the F330W bandpass if the cavity has been cleared of both large and small dust grains, but we do not detect it. The lack of an observed cavity can be explained by the presence of sub-$\mu$m grains interior to the sub-mm cavity wall. We suggest one explanation for this which could be due to a planet of mass $<$~9~M$_J$ interior to 24~AU. A unique cylindrical structure is detected in the FUV data from the Advanced Camera for Surveys/Solar Blind Channel (ACS/SBC). It is aligned along the system semi-minor axis, but does not resemble an accretion-driven jet. The structure is limb-brightened and extends 190~AU~$\pm$~35~AU above the disk midplane. The inner radius of the limb-brightening is 40~$\pm$~10~AU, just beyond the sub-millimeter cavity wall. | Transitional disks are protoplanetary disks that are in the process of evolving from a gas-rich primordial disk to a gas-poor debris disk. During this transition period, material in the disk within the first few tens of AU from the star clears to form an optically thin gap or cavity. Consistent with the presence of accretion, these objects retain an inner disk within a few AU of the star \citep{Hartigan_90}. If the cleared cavities were devoid of material, these systems would be unable to replenish their inner disk and accretion would cease relatively quickly. Many of these objects continue to accrete material while having cavities tens of AU in radii, as evidenced by mid-IR/sub-mm/mm observations \citep{Rice_06,Lubow_06,Salyk_13}. One possible explanation for such cavity generation is a filtration mechanism that allows only small dust grains, which would be undetected in the sub-mm/mm data, to migrate inward through the gap entrained with the still-accreting gas from the outer disk \citep{Quillen_04}. \citet{Cieza_12} compared several processes driving disk evoution including grain growth, the effect of planets and photoevaporation mechanisms. Starting with the hypothesis that small grains in a cavity are part of the process of grain growth, they found that grain growth can account for $>40$\% of transition disks around K- and M-type stars, though the process can have complicating factors like fragmentation or replenishment \citep{Dullemond_05}. The dust grains that make up the interstellar medium (ISM) are estimated to include a range in size from 0.005~$\mu$m to 1~$\mu$m \citep{Mathis_77}. \citet{Rice_06} suggest that filtration of small ISM-like grains into the cavity devoid of larger grains could be an indication that a planet with a mass of 1-6~M$_J$ resides within the cavity. Simulations by \citet{Zhu_12} and \citet{de_Juan_Ovelar_13} produce similar results, and they also suggest the presence of one or more gas giant planets as a disk clearing mechanism. FUV-driven neutral atomic or molecular disk winds have also been considered as a possible mechanism for disk dispersal. At large radii, giant planets have difficulty clearing material on a sufficiently short time scale to be consistent with estimates of the disk-dispersal time for T~Tauri stars \citep[$\approx 10^5$ yr;][]{Simon_95,Wolk_96}. In rich cluster environments, \citet{Johnstone_98} showed that FUV radiation from nearby stars can dominate the photoevaporation rate. This, however, does not explain the rapid dispersal of disks around young stars with no nearby source of the FUV flux. \citet{Gorti_09} have considered flows driven by FUV radiation from the central star and argued that mass-loss rates of the order of 10$^{-8}$~M$_{\sun}$~yr$^{-1}$ can be obtained at large radius (>100 AU). Mapping the dust grain size distribution throughout the disk can indicate which mechanism is responsible for clearing material from the disk. Dust grains scatter light most efficiently at wavelengths comparable to their size, so determining whether sub-$\mu$m dust grains are present inside a disk cavity requires observations to be in sub-$\mu$m bandpasses. Disks with large cavities provide test-beds for multi-wavelength observations probing the surface of the disk to analyze the light scattering properties and size distribution of the dust grains. This method for determining grain properties and distributions has been employed successfully at sub-mm and longer wavelengths \citep{Perez_12, Banzatti_11}, but relies on scattered light rather than thermal emission. High contrast far-ultraviolet (FUV) and optical images have the added benefit of a smaller inner working angle (IWA) than at longer wavelengths with the same instrument. For disks with cavities tens of AU in diameter, such observations allow us to map the spatial distribution of gas and small (sub-$\mu$m) grain reflection nebulosity present at the disk surface, both interior and exterior to the cavity wall. If small-grain dust exists interior to the large grain cavity wall, we would not expect transitional disks with relatively high accretion rates to have an optically thin cavity at short (FUV and optical) wavelengths. The disk associated with the classical T~Tau star GM~Aur \citep[K5.5 $\pm$ 1.0, B-V=1.12,][]{Espaillat_10} has been studied extensively for two decades with the Hubble Space Telescope (HST). The distance to GM~Aur is 136$_{-29}^{+50}$ pc \citep{Bertout_06}, and the inclination from pole-on for its disk is 55\degree \citep{Calvet_05,Hughes_09,Andrews_11}. The disk has been detected in scattered optical and NIR light \citep{Stapelfeldt_95,Schneider_03}. The disk became categorized as transitional when millimeter and sub-millimeter observations detected a cleared inner cavity extending from the star to between 20~AU and~28~AU \citep{Calvet_05,Hughes_09,Grafe_11,Andrews_11}; in this paper we assume the sub-mm cavity is 24~AU in radius. GM~Aur is also accreting material at a rate of $\approx$ 10$^{-8}$ M$_{\sun}$ yr$^{-1}$ \citep{Gullbring_98,White_01,Ingleby_15}, suggesting that material is migrating inward from the outer disk, through the sub-mm cavity, and accreting onto the star. This can be directly tested with FUV and short wavelength optical observations. We have acquired the requisite datasets with HST's Advanced Camera for Surveys (ACS) using the Solar-Blind Channel (SBC), as well as archival data from both the High Resolution Channel (HRC) and the 2nd generation Wide Field Planetary Camera (WFPC2). In this paper HST data of GM~Aur, along with data from a wide variety of other instruments, are used to create a Spectral Energy Distribution (SED), all of which we then use to model the GM~Aur star+disk system. We probe material within the sub-mm cavity region, to test the hypothesis that small grain dust exists within the 24~AU cavity of GM~Aur's transitional disk. A description of the data used and the data reduction are described in \S~2. Our analysis and results are found in \S~3, with a discussion in \S~4. | The disk of GM~Aur has now been detected in scattered light from 0.1450~$\mu$m to 1.6~$\mu$m, more than a decade in wavelength coverage. This large wavelength lever allows us to explore the surface dust grain opacity as a function of wavelength, in a manner similar to \citet{Pinte_08} but with the addition of FUV wavelength coverage. In contrast to \citet{Schneider_03} who inferred a dust grain composition similar to Model 1 in \citet{Wood_02}, which is composed of large grains (up to 1~mm), we find that the wavelength dependence of F$_{disk}$/F$_{star}$ is consistent with ISM-like grains. This indicates that little grain growth has occurred in the upper layers of the disk. ISM-like grains are also consistent with the radial surface brightness profile of the disk at 3300 \AA . The overall color of the dust disk surface is blue ($m_{F330W}-m_{F555W}~\approx~-0.2$), and the presence of ISM-like grains is undoubtedly a factor in disk detection in direct imaging using WFPC2 \citep{Stapelfeldt_95}. \subsection{Cavity Non-Detection at Short Wavelengths} The cavity detected in mm, sub-mm, and mid-IR wavelengths \citep{Calvet_05,Hughes_09,Grafe_11, Andrews_11} is large enough in radius to be detected in the FUV and optical F330W bandpasses. The consistency in detections of the cavity at mid-IR and longer wavelengths makes it safe to assume there is a cavity in the large grain dust. However, we report a non-detection of the cavity exterior to 15 AU at optical and FUV wavelengths. Our analysis provides no evidence for a break in the radial surface brightness profile that would indicate a depletion in the surface density of the grains or signal the presence of a change in disk composition. This can only be the case if the small grains at the surface of the disk persist inside the radius of the sub-millimeter cavity wall. Our non-detection of the cavity indicates that there is a mechanism in place actively filtering the material within the cavity, allowing only sub-$\mu$m grains to migrate inward beyond the cavity wall. (This is not precluded by the SED fit, and does not conflict with a statement by \citet{Calvet_05} on the outer boundary of the optically thin region in the IR and/or redward. The inner region is optically thick in the far-UV.) This conclusion is supported by the analysis of ro-vibrational CO lines in \citet{Salyk_11}. Their models suggest that an inner disk of CO gas extends out to a radius of 0.2 AU of GM~Aur, and that it is likely being replenished via gas migration through the cavity from the outer disk. This is consistent with Monte Carlo Radiative Transfer and hydrodynamical model predictions in \citet{Paardekooper_06a, Rice_06,Zhu_12,Dong_12a} \&\ \citet{de_Juan_Ovelar_13}, which suggest giant planet formation within the disk cavity as a likely culprit in this scenario. There are other mechanisms which can also preferentially decrease the density of micron-sized grains. Grain growth is one possibility, although it could be expected to vary smoothly with radius \citep{Cieza_12}. \citet{Dullemond_05} concluded that if grain growth were responsible for cavity development, small grains must be replenished, possibly by aggregate fragmentation via high-speed collisions. \citet{Owen_11} suggested that X-ray photo-evaporation could explain a large fraction ($\gsim 50$\%) of transitional disks. However, \citet{Alexander_06b} suggested that GM~Aur in particular had too high an accretion rate to have its cavity produced by photo-evaporation, and that it was rather caused by another mechanism such as planet formation or grain growth/coagulation. The detection of small dust grains within the millimetric cavity seems to rule out photoevaporation and dust coagulation processes as the main origin for the mm feature, as in both cases small dust particles would not be expected inside the hole. Our result here confirms the predictions of \citet{Calvet_05} and \citet{Espaillat_10}, using an independent data set. \subsection{Limits on Giant Planets in the Disk of GM~Aur} Grain filtration occurs when larger grains are restricted to the outer disk and small grains, which are more tightly coupled to the gas, can freely penetrate the cavity \citep{Rice_06}. Grain filtration is also a predicted consequence of the presence of one or more giant planets. Based on the models presented in \citet{Zhu_12}, a 3~$M_J$ object would clear a gap in the 30~$\mu$m grains on a timescale of 10$^5$ yrs, and create a noticeable depletion ($\approx$ 3 orders of magnitude) in the gas density within the cavity. The models presented in \citet{de_Juan_Ovelar_13} produce similar results to those presented in \citet{Zhu_12}. They find that a 1 $M_J$ planet does deplete 1~$\mu$m sized grains at the surface of the disk that could be detected in optical ($0.65~\mu$m) observations \citep[Fig.~3]{de_Juan_Ovelar_13}, given high enough spatial resolution. They find that a 9~M$_J$ planet depletes the 1~$\mu$m grains at the disk surface by a factor of 1000 \citep[Fig.~7]{de_Juan_Ovelar_13} at the planet location; a 15 $M_J$ planet eliminates dust grains of all sizes entirely from the disk surface at the location of the planet. Their results suggest that a planet with a mass $>~9~M_J$ at $\approx$~20 AU would create a depletion in sub-$\mu$m surface grains that would be detectable in our PSF-subtracted 1400\AA, 1650\AA, and 3300\AA\ data. However, a depletion is not detected in our UV or optical observations, and its absence places an upper limit on the mass of a planet in the disk of GM~Aur to a mass of $<~9~M_J$, likely even lower, though the model grid is too course to be more specific. Furthermore, the models by \citet{Zhu_12} and \citet{de_Juan_Ovelar_13} also place a $1~M_J$ lower limit on planet mass due to the detection of a cavity in mm and sub-mm data \citep{Calvet_05,Hughes_09,Grafe_11,Andrews_11}. \subsection{A Possible Molecular Outflow} \label{subsec-possmolecularoutflow} Molecular outflows launched from the inner disk are expected in very young star $\plus$ disk systems \citep{Ercolano_09}, but their persistence in older systems is less well explored. Extended H$_2$ emission has been detected around a host of young objects, for example RU~Lupi \citep[different signals in different apertures,][]{Herczeg_05a}, plus T~Tau \citep{Walter_03,Saucedo_03} and DG~Tau \citep{Schneider_13}. Spectroscopically, it is clear that the H$_2$ is coming from an emitting region that is not from the star but rather from the inner part of the disk, thus it is extended spatially. For DG~Tau, the UV H$_2$ emission appears as a limb brightened ``bubble" with a length of about 0$\farcs$3= 42~AU located toward the approaching lobe of the outflow. A very similar morphology is observed in the near-IR \citep{Agra-Amboage_14}. The H$_2$ emission lines in several systems are also red/blue-shifted, which is expected in the case of molecular outflows \citep{Herczeg_05,Herczeg_06,France_12}; \citet{Herczeg_06} discuss the spectral evidence for other potential sources of H$_2$ emission as well. \citet{France_12} found no clear evidence for extended UV fluorescent H$_2$ emission lines for GM Aur in their Cosmic Origins Spectrograph (COS) data around 1450 \AA . The COS line profiles they present do not show obvious extended wings, though there may be a hint of marginal additional flux at -30 to -40 km~s$^{-1}$ in their Fig.~3. Otherwise, they are centered at the systemic velocity and appear compatible with emission from the inner disk (at radius 0.5~AU). Their analysis was designed to map the distribution of H$_2$ in the disk, assuming Keplerian motion of the fluorescent gas. As discussed in \S \ref{FUV_structure}, we detect an extended signal aligned along the disk's semi-minor axis in the F140LP and F165LP data that we do not detect in any other bandpass. This structure projects beyond the boundary where we expect to observe reflection and scattered light from the disk surface. The protruding region of the feature is limb-brightened and its geometry appears cylindrical with a radius of 40 $\pm$ 10~AU. A possible interpretation is that it is a gaseous photoevaporative wind of H$_2$, with some of the smallest sub-$\mu$m dust grains embedded within, being driven from the disk surface just beyond the location of the 24~AU sub-mm cavity. We only detect the feature in the observations where numerous H$_2$ transitions exist (in the FUV), indicating its composition could be primarily molecular hydrogen. The smallest dust grains are the most tightly coupled to the gas, so without detection in the F330W data, any dust grains embedded within the gas must be either small in radius or low in abundance. The structure, being aligned with the system minor axis, may offer an explanation for the 5\arcsec\ distant ``polar lobes'' seen in the NICMOS data \citep{Schneider_03}, which they interpreted as likely due to shocked line emission from a molecular outflow. We exercise some caution in this interpretation, subject to confirmation of this feature, which would benefit from more detailed modeling and further FUV spectroscopy specifically designed to cover the candidate outflow region. The dominant excitation mechanisms in classical T~Tauri stars are the high-energy X-ray and EUV photons that are produced during accretion. These high-energy photons photo-dissociate molecular gas and excite the atomic components that remain, so they are easily detected in the optical forbidden emission lines \ion{O}{1}, \ion{S}{2} etc. \citep{Schneider_13}. Strong UV \ion{C}{4}, \ion{He}{2}, \ion{Si}{4} and \ion{N}{5} lines have been detected towards GM~Aur \citep{Ardila_13}. Furthermore, \citet[their figure 2]{Schneider_13} show the bandpass of the F140LP and F165LP filters superposed on the spectrum of the active T~Tauri star DG Tau. \ion{C}{4}, \ion{He}{2}, and \ion{Si}{4} emission may also contribute to the flux in the F140LP filter. The \ion{C}{4} lines are the strongest features longward of 1300~\AA\ but as noted in \citet{France_12}, the integrated H$_2$ flux is twice that in the \ion{C}{4} doublet. However, we have not detected any extended atomic or ionic emission lines in the broadband optical WFPC2 imagery or UV/optical STIS spectra, and conclude that it is not an atomic or ionic jet that is responsible for the extended emission detected in the FUV ACS/SBC data. For clarification, we follow the terminology convention of \citet{Klaassen_13,Ray_07} \&\ \citet{Reipurth_Bally_01}, and define a jet as high velocity gas (>~100~km~s$^{-1}$) launched from the inner 0.1 AU of the disk, collimated by the magnetic field. The geometry of the outflow is inconsistent with typical collimated atomic jets. The outflow appears to launch from the inner portions of the outer disk rather than the central star, and smoothly extends 190~AU $\pm$ 35~AU until it is no longer detectable above the background. There is also no detection of accompanying H I in the outflow, so it is likely that the material does not dissociate, which precludes X-ray and EUV as the excitation mechanism. H$_2$ is far more likely this far out from the star rather than hot (up to $10^5$~K) gas. The hot transition region lines (\ion{C}{4}, \ion{Si}{4}) seem to be emitted from the stellar surface, while the molecular lines are circumstellar. These are further arguments against the possibility that the source of the emission is hot. It is possible that the H$_2$ is being excited by FUV photons at Ly$\alpha$. If it is fluorescently excited by Lyman $\alpha$ from the vicinity of the star, we should expect the radial brightness of the structure to decline as $\approx r^{-2}$. However, we find that the radial profile of the structure follows an $r^{-0.7 \pm 0.1}$ dependance. One possibility this implies is that some fraction of the fluorescent excitation could be shocked gas \citep{Herczeg_02}. GM~Aur may possibly present the first direct detection of the dominant mechanism responsible for clearing gas from the outer disk. If this phenomenon were representative of other transitional disks at 1-2 Myr, then it would place tight constraints on the time frame over which gas giants may form. How common this could be among transitional disks and over what portion of a disk's lifetime this continues are points that can only be addressed with observations of additional transitional disks in the FUV. \subsection{Summary of Results} $\bullet$ \ We confirm previous results for disk orientation and inclination, but at shorter wavelengths than those in the literature. Additionally, we have combined multiple datasets from the literature and unpublished data to create a detailed SED that spans nearly 4 orders of magnitude (from 0.14 $\mu$m to 860 $\mu$m). $\bullet$ \ We resolve the disk down to a radius corresponding to 15 AU (the distance to GM~Aur assumed to be 140 pc) in optical and FUV wavelengths. We do not detect a change in the radial surface brightness profile at or near the location of the sub-mm cavity wall. We conclude that small grain dust and gas exist within the cavity, which is consistent with models that describe dust filtration via planet-disk interaction \citep{Zhu_12,Dong_12a,de_Juan_Ovelar_13}. $\bullet$ \ Comparing the surface brightness of the disk imaged at the multiple wavelengths discussed here and reported in \citet{Schneider_03} with grain models (see \citet{Whitney_92,Kim_94,Cotera_01,Wood_02}), we conclude that the surface of the disk is populated by small grains. $\bullet$ \ The FUV observations detect a signal, undetected at longer wavelengths, that extends along the disk semi-minor axis. One possible explanation we put forth is that it is a FUV photoevaporative disk wind composed of H$_2$ and small grain dust. However, radial velocity measurements along with additional FUV long-slit spectral data are needed in the future, in order to test this hypothesis. \subsection{Implications for the Future} As mentioned in \S \ref{subsec-possmolecularoutflow}, small scale H$_2$ molecular outflows have been detected in T~Tauri stars, especially in the near-IR \citep{Beck_08,Beck_12}, but also in the FUV \citep[DG Tau;][]{Schneider_13}. They could be something other than thermally driven winds. For example, at $r=40$~AU around a $1.2~M_{\odot}$ star the escape velocity would be $v_{esc}\sim 7$~km~s$^{-1}$, which would require a gas temperature $T_{gas}>5000$~K, high for the molecular gas component. GM~Aur allows possibly the first direct imaging detection of such a molecular outflow from a T-Tauri star in the FUV. If confirmed in future observations of GM~Aur, its presence would have far-reaching implications. Observations with ALMA could determine the abundance \&\ chemical composition of GM~Aur's disk, as well as provide high precision radial velocity data for the gas in the extended region \citep{Bruderer_14,Klaassen_13,Mathews_13}. In the past, a topic of much speculation has been over what time frame and by what mechanism gas is cleared from the outer disk. H$_2$ is the dominant gas species in disks. Therefore, if this outflow feature is composed primarily of H$_2$, and found in future observations to be present at an early age for a significant number of star + disk systems, it will constrain the time frame over which gas giants may form. ALMA data could also have an impact on how we approximate gas to dust ratios in Monte Carlo radiative transfer and hydrodynamical models of transitional disk systems \citep{Bruderer_14}. When available, high-contrast imagery in FUV and short-wavelength optical bandpasses enhances our ability to determine whether small-grain dust exists within the cavities of transitional disks. Transitional disks with larger cavities than GM~Aur would constrain the dust opacity and particle size distribution, as well as place limits on the ice content at the dust disk surface. Higher S/N data, such as might be provided by the next generation of UV instrumentation, are needed to probe ice chemistry in the outer dust disk surface. They could also more stringently constrain the abundance of pure ice grains from the most abundant ice species. Converting flux data into mass loss rates - which can then be quantitatively compared with predicted photo-evaporation and photo-dissociation rates - requires velocity data at the location of the molecular H$_2$ outflow. The loss rate of the dominant gas species in the disk is directly related to the time allowed for gas giant formation. It is important for transitional disk investigators to fully explore the feasibility of obtaining high-contrast FUV imagery on these objects while the opportunity exists, as there is no approved successor to HST with a FUV imaging capability. | 16 | 7 | 1607.03520 |
1607 | 1607.07465_arXiv.txt | Turbulence profoundly affects particle transport and plasma heating in many astrophysical plasma environments, from galaxy clusters to the solar corona and solar wind to Earth's magnetosphere. Both fluid and kinetic simulations of plasma turbulence ubiquitously generate coherent structures, in the form of current sheets, at small scales, and the locations of these current sheets appear to be associated with enhanced rates of dissipation of the turbulent energy. Therefore, illuminating the origin and nature of these current sheets is critical to identifying the dominant physical mechanisms of dissipation, a primary aim at the forefront of plasma turbulence research. Here we present evidence from nonlinear gyrokinetic simulations that strong nonlinear interactions between counterpropagating \Alfven waves, or strong \Alfven wave collisions, are a natural mechanism for the generation of current sheets in plasma turbulence. Furthermore, we conceptually explain this current sheet development in terms of the nonlinear dynamics of \Alfven wave collisions, showing that these current sheets arise through constructive interference among the initial \Alfven waves and nonlinearly generated modes. The properties of current sheets generated by a strong \Alfven wave collisions are compared to published observations of current sheets in the Earth's magnetosheath and the solar wind, and the nature of these current sheets leads to the expectation that Landau damping of the constituent \Alfven waves plays a dominant role in the damping of turbulently generated current sheets. | The ubiquitous presence of turbulence impacts the evolution of many space and astrophysical plasma environments, mediating the transport of energy from violent events or instabilities at large scales down to the small scales at which the energy is ultimately converted to heat of the protons, electrons, and minor ions. It is widely believed that plasma turbulence plays an important role in heating the solar corona to millions of degrees Kelvin, accelerating the solar wind to hundreds of kilometers per second, regulating star formation, transporting heat in galaxy clusters, and affecting the injection of particles and energy into the Earth's magnetosphere. At the forefront of plasma turbulence research is the effort to identify the physical mechanisms by which the turbulent fluctuations are damped and their energy converted to plasma heat or some other energization of particles. In contrast to the intermittent filaments of vorticity that arise in hydrodynamic turbulence \citep{She:1990}, intermittent current sheets are found to develop in plasma turbulence \citep{Matthaeus:1980,Meneguzzi:1981}. Recent work investigating the statistics of these coherent structures, self-consistently generated by the plasma turbulence, has demonstrated that the dissipation of turbulent energy is largely concentrated in these current sheets \citep{Uritsky:2010,Osman:2011,Zhdankin:2013}. Since current sheets are associated with enhanced dissipation, illuminating their origin and nature is critical to identifying the dominant physical mechanisms of dissipation in plasma turbulence. How current sheets develop in plasma turbulence is a longstanding question in the study of space and astrophysical plasmas \citep{Parker:1972,Pouquet:1978,Priest:1985,vanBallegooijen:1985,Antiochos:1987,Zweibel:1987,Biskamp:1989,Longcope:1994,Cowley:1997,Spangler:1999,Biskamp:2000,Merrifield:2005,Greco:2008}. Early work focused on the study of solar coronal loops, asking if the continuous motion of line-tied footpoints in ideal MHD would lead an initially smooth magnetic field to develop a tangential discontinuity \citep{Parker:1972}, necessarily supported by a sheet of finite current, according to Maxwell's equations. It was argued that continuous footpoint motion cannot generate a discontinuous magnetic field \citep{vanBallegooijen:1985,Antiochos:1987,Zweibel:1987}, but later shown that current layers of finite but arbitrarily small thickness were realizable through continuous footpoint motion \citep{Longcope:1994,Cowley:1997}. Of course, in the more complete kinetic plasma description, current layers generally have structure at both characteristic ion and electron length scales; kinetic simulations of plasma turbulence indeed observe the development of current sheets of finite thickness \citep{Wan:2012,Karimabadi:2013,TenBarge:2013a}. Recent spacecraft measurements of current sheets in the near-Earth solar wind have lead to fundamental questions about their origin and their influence on plasma heating. Do the measured current sheets represent advected flux tube boundaries \citep{Borovsky:2008,Borovsky:2010}, or are they generated dynamically by the turbulence itself \citep{Boldyrev:2011,Zhdankin:2012}? In the past few years, vigorous activity has focused on the spatial localization of plasma heating by the dissipation of turbulence in current sheets through statistical analyses of solar wind observations \citep{Osman:2011,Borovsky:2011,Osman:2012a,Perri:2012a,Wang:2013,Wu:2013,Osman:2014b} and numerical simulations \citep{Wan:2012,Karimabadi:2013,TenBarge:2013a,Wu:2013,Zhdankin:2013}. To understand the origin of coherent structures, we must investigate how the turbulent nonlinear interactions govern their development \citep{Howes:2015b}. In plasma turbulence, the \Alfven wave represents the fundamental response of the plasma to an applied perturbation. Early research on incompressible MHD turbulence in the 1960s \citep{Iroshnikov:1963,Kraichnan:1965} emphasized the wave-like nature of turbulent plasma motions, suggesting that nonlinear interactions between counterpropagating \Alfven waves---or \emph{\Alfven wave collisions}---mediate the turbulent cascade of energy from large to small scales. The \Alfven wave remains central to modern theories of MHD turbulence that provide explanations for the anisotropic nature of the turbulent cascade \citep{Goldreich:1995} and the dynamic alignment of velocity and magnetic field fluctuations \citep{Boldyrev:2006}. In this Letter, we demonstrate that the generation of current sheets in plasma turbulence is a natural consequence of strong \Alfven wave collisions. Furthermore, we present a first-principles explanation for this current sheet development in terms of the nonlinear dynamics, showing that the current sheet can be accurately reconstructed from a linear superposition of the interacting \Alfven waves and a surprisingly small number of nonlinearly generated modes. The properties of the resulting current sheet are compared to previously published observations of current sheets in turbulence in the Earth's magnetosheath \citep{Retino:2007,Sundkvist:2007} and the solar wind \citep{Perri:2012a}. Finally, we discuss implications for the damping of current sheets generated by strong \Alfven wave collisions. | The nonlinear dynamics of strong \Alfven wave collisions provides a natural explanation for the ubiquitous development of current sheets in plasma turbulence. The discovery that current sheets arise transiently through constructive interference among the primary waves and nonlinearly generated modes provides valuable insight into the physical mechanisms by which the turbulent fluctuations are damped. Because the dominant constructively interfering modes are $k_z=+k_\parallel$ ($k_z=-k_\parallel$) \Alfven waves that propagate in the $+\zhat$ ($-\zhat$) direction and $k_z=0$ modes nonlinearly generated by the interactions between these counterpropagating waves, collisionless damping of the constituent \Alfven waves via the Landau resonance with protons and electrons is expected to play an important role in the dissipation of the current sheets, as previously suggested \citep{TenBarge:2013a}. Of course, other mechanisms exist that can produce current sheets, such as particular flow and magnetic field geometries, like the Orszag-Tang vortex. A final question that remains to be answered is whether such alternative mechanisms play any role in the development of current sheets in solar wind turbulence, or are strong \Alfven wave collisions sufficient to account for all current sheets observed in the turbulent solar wind? This work was supported by NSF grant PHY-10033446, NSF CAREER Award AGS-1054061, and NASA grant NNX10AC91G. This work used the Extreme Science and Engineering Discovery Environment (XSEDE), which is supported by National Science Foundation grant number ACI-1053575. | 16 | 7 | 1607.07465 |
1607 | 1607.01844_arXiv.txt | \par We develop a new Monte-Carlo-based method to convert the SDSS (Sloan Digital Sky Survey) $u$-band magnitude to the SCUSS (South Galactic Cap of $u$-band Sky Survey) $u$-band magnitude. Due to more accuracy of SCUSS $u$-band measurements, the converted $u$-band magnitude becomes more accurate comparing with the original SDSS $u$-band magnitude, in particular at the faint end. The average $u$ (both SDSS and SCUSS) magnitude error of numerous main-sequence stars with $0.2<g-r<0.8$ increase as $g$-band magnitude becomes fainter. When $g=19.5$, the average magnitude error of SDSS $u$ is 0.11. When $g=20.5$, the average SDSS $u$ error is up to 0.22. However, at this magnitude, the average magnitude error of SCUSS $u$ is just half as much as that of SDSS $u$. The SDSS $u$-band magnitudes of main-sequence stars with $0.2<g-r<0.8$ and $18.5<g<20.5$ are converted, therefore the maximum average error of converted $u$-band magnitudes is 0.11. The potential application of this conversion is to derive more accurate photometric metallicity calibration from SDSS observation, especially for those distant stars. Thus, we can explore stellar metallicity distributions either in the Galactic halo or some stream stars. | \par It is an increasing perception that the Galactic halo system comprises at least two spatially overlapping components with different kinematics, metallicity and spatial distribution \citep{Carollo07, Carollo10, An13, An15}. Chemical abundance is the direct observational ingredient in investigating the dual nature of the Galactic halo. Since the chemical abundance of stars have strong effect on the emergent flux, especially at blue end, the natural endeavor is to recover the metal information from large photometric surveys such as SDSS \citep[Sloan Digital Sky Survey;][]{York00}. The advantage of photometric metallicity estimate is that the metallicity information of large numbers of stars can be obtained. \par Based on the SDSS $ugriz$ photometry, \cite{Ivezic08} used polynomial-fitting method from spectroscopic calibration of de-reddened $u-g$ and $g-r$ colors to derive the photometric metallicity \citep[see also][]{Peng12}. However, due to the relatively large error of SDSS $u$-band magnitude, only the metallicities [Fe/H], of stars brighter than $g=19.5$ are obtained. Combining the more accurate SCUSS \citep{Zhou16} $u$-band photometry, SDSS $g$ and $r$ photometry, \cite{Gu15} developed a three-order polynomial photometric metallicity estimator, in which $u$-band magnitude can be used to faint magnitude of $g=21$. However, both estimator developed by \cite{Ivezic08} and \cite{Gu15} based on polynomial-fitting have their intrinsic drawback that they can not be extended to metal-poor end. In order to solve this problem, \cite{Gu16} (hereafter denoted as Paper I) devised a Monte-Carlo method to estimate stellar metallicity distribution function (MDF) which appears particularly good at both metal-rich and metal-poor ends. The natural forward step is to combine the SCUSS $u$, SDSS $g, r$ photometry with the method introduced in Paper I to investigate the MDF of the Galactic halo stars. But only those stars in South Galactic cap are surveyed by SCUSS. How can we estimate the photometric metallicity distribution of faint stars (deep in Galactic halo) in both South and North hemisphere? This paper provides a new method to achieve this goal. Due to the fact that SCUSS $u$ is more accurate than SDSS $u$, we convert SDSS $u$ to SCUSS $u$ using a Monte-Carlo method, through which we make the converted $u$ magnitude becomes as accurate as SCUSS $u$ magnitude. \par We organize this paper as follows. In Section 2, we take a brief overview of the SDSS and SCUSS. The technical details for converting SDSS $u$ to SCUSS $u$ are presented in section 3. Section 4 evaluate the effectiveness of this conversion. The discussion of the potential application of the conversion is given in Section 5. | \par As we all know, $u$-band measurements is very important to derive the photometric metallicity and therefore to construct a precise MDF. Because of the relatively shallow survey limit ($u\sim22$) and the relatively large error in the SDSS $u$-band near the faint end, the application of the photometric metallicity estimates is greatly restricted in the range of $g<19.5$, an insufficient depth to explore the distant halo and substructures. However, the SCUSS $u$ is 1.5 mag deeper than SDSS $u$, and its error is smaller than SDSS $u$ error on the whole. The potential application of the conversion from SDSS $u$ to SCUSS $u$ is very important to derive relative accurate photometric metallicities of distant stars. In Paper I, we developed a new method to estimate the photometric metallicity distribution of large number of stars. Compared with other photometric calibration methods, this method in Paper I effectively reduces the error induced by the method itself, and therefore enables a more reliable determination of the photometric MDF. However, another error source still matters: the error of SDSS $u$-band magnitude. This error behavior limits the application of the method in the range of $g<19.5$ in Paper I. This range is same as that of Ivezi\'c et al.'s (2008) photometric metallicity estimator. The more accurate SCUSS $u$-band measurements guarantee the accuracy of the stellar distribution in $u-g$ versus $g-r$ panel, and it extends the application of method in Paper I to even fainter stars. Thus, the photometric MDF of distant stars such as halo stars or some stream stars can be estimated. However, only the stars in South Galactic Cap are surveyed by SCUSS which have relatively more accurate $u$ band magnitude, how to derive the photometric metallicity of stars in the North Galactic hemisphere? The conversion from SDSS $u$ to SCUSS $u$ statistically diminish the error of $u$-band magnitude, which make it possible to estimate the photometric MDF of stars in the whole sky. In this study, we have done the conversion for stars in $18.5<g<20.5$. The conversion combined with the method introduced in Paper I enable us to estimate the photometric metallicity distribution function for stars at least in the range of $g<20.5$, which is 1 mag deeper than that of spectroscopically-surveyed stars. So we can study the chemical structure of the Galactic halo more detailed. Besides the application described above, the more accurate $u$ band magnitude from the conversion, can be applied to address other scientific issues. | 16 | 7 | 1607.01844 |
1607 | 1607.03593_arXiv.txt | The moment of inertia, the spin-induced quadrupole moment, and the tidal Love number of neutron-star and quark-star models are related through some relations which depend only mildly on the stellar equation of state. These ``I-Love-Q'' relations have important implications for astrophysics and gravitational-wave astronomy. An interesting problem is whether similar relations hold for other compact objects and how they approach the black-hole limit. To answer these questions, here we investigate the deformation properties of a large class of thin-shell gravastars, which are exotic compact objects that do not possess an event horizon nor a spacetime singularity. Working in a small-spin and small-tidal field expansion, we calculate the moment of inertia, the quadrupole moment, and the (quadrupolar electric) tidal Love number of gravastars with a polytropic thin shell. The I-Love-Q relations of a thin-shell gravastar are drastically different from those of an ordinary neutron star. The Love number and quadrupole moment are negative for less compact models and the I-Love-Q relations continuously approach the black-hole limit. We consider a variety of polytropic equations of state for the matter shell, and find no universality in the I-Love-Q relations. However, we cannot deny the possibility that, similarly to the neutron-star case, an approximate universality might emerge for a limited class of equations of state. Finally, we discuss how a measurement of the tidal deformability from the gravitational-wave detection of a compact-binary inspiral can be used to constrain exotic compact objects like gravastars. | Self-gravitating astrophysical objects are roughly spherically symmetric. If they rotate and/or are in binary systems, however, their shape deviates from spherical symmetry due to centrifugal and/or tidal forces. The response to such deformations provides us with important information on the matter distribution inside the object. Deviations from spherical symmetry may be characterized by the multipole moments of the gravitational fields outside the body. The properties of a nonspherical, self-gravitating distribution of matter are then encoded in its multipole moments. For instance, the spin-induced quadrupole moment encodes information about the deformability of a rotating body, which is related to the stiffness of the matter distribution (see, e.g., Refs.~\cite{Laarakkers,PoissonWill}). The multipole moments of a neutron star depend in similar ways on its internal structure. Thus, measurements of the first multipole moments can be used to constrain the equation of state of the neutron-star core, which is still unknown at nuclear and the super-nuclear densities. Another interesting application of the multipole moments of compact objects is related to tests of the black-hole no-hair theorems. Astrophysical black holes are fully characterized by only their mass $M$ and angular momentum $J$, all higher multipole moments being uniquely determined only by these two quantities~\cite{hage1,hage2,hage3,hage4}. Thus, independent measurements of at least three multipole moments can be used to distinguish black holes from exotic compact objects or to provide null-hypothesis tests of the Kerr metric and of general relativity (see, e.g., Ref.~\cite{Berti:2015itd} for a review). In this context, the lowest-order multipole moments, namely the mass $M$, angular momentum $J$, and quadrupole moment $Q$, would play a prime role. It has been recently pointed out that there exist various approximate relations among the lowest-order multipole moments of compact stars which are only mildly dependent on the equation of state. By considering a reasonable set of equations of state for the neutron-star interior, approximately universal relations between the moment of inertia and the lowest-order multipole moments have been found (e.g., see Refs.~\cite{lattimer,urbanec}). More recently, it has been found that, for slightly deformed neutron stars, the moment of inertia $I$, the tidal quadrupole deformability $\lambda$ (the tidal Love number), and the spin-induced quadrupole moment $Q$, also satisfy some universal relations (dubbed as ``I-Love-Q'') that are nearly independent of the equation of state to within a few percent level~\cite{yagi, yagi2}. Here, the tidal Love number $\lambda$ is defined as the ratio of the tidally-induced quadrupole moment of a nonspinning star to the strength of the perturbing quadrupolar tidal field. These nearly universal relations will be helpful to extract physically important quantities from observational data if they exist within the required accuracy. For example, the I-Love-Q relations can break degeneracy in the models used for X-ray and gravitational-wave observations of neutron stars. For these reasons, nearly-universal relations of compact objects have attracted much attention in the last few years. Although the I-Love-Q relations were originally found~\cite{yagi,yagi2} for slightly deformed, isolated, nonmagnetized compact stars, they have been extended to more general cases, namely dynamical configurations~\cite{vitor1}, rapid rotation~\cite{doneva,Pappas:2013naa,Chakrabarti:2013tca,Yagi:2014bxa}, nonbarotropic~\cite{Martinon:2014uua} and anisotropic~\cite{yagi3,yagi4,yagi5} fluids, strong magnetic fields~\cite{haskell}, and also deviations from general relativity~\cite{yagi,yagi2,Sham:2013cya,Pani:2014jra,Doneva:2014faa}. Possible explanations for the emergence of this approximate universality are given in Refs.~\cite{yagi7,Sham:2014kea}. Similar universal relations have been also found among higher-order multipole moments induced by tidal effects for nonrotating neutron stars~\cite{yagi6}. The inclusion of the rotational Love numbers in the I-Love-Q relations has been also argued in Ref.~\cite{pani3}. These universal relations among higher-order multipole moments of a neutron star are reminiscent of the black-hole no-hair theorems, namely they show that only a handful of low-order multipole moments can (approximately) characterize the gravitational field of a compact object. To investigate this problem and to elucidate an origin of the I-Love-Q relation, it is interesting to understand how the I-Love-Q relations behave in the black-hole limit. Since the ratio of the mass $M$ to the radius $R$ (i.e., the compactness) of any standard spherical perfect-fluid star has the upper limit $M/R<4/9$ (in geometrized units, see, e.g., Refs.~\cite{Buchdahl:1959zz,hartle0}), the compactness of a perfect fluid star is disconnected from that of a black hole, $M/R=1/2$. To examine the black-hole limit of the I-Love-Q relations, therefore, we have to consider somewhat peculiar compact objects that can sustain higher compactness. Yagi and Yunes~\cite{yagi3,yagi4,yagi5} have studied this problem using fluid stars with anisotropic pressure. In the models they employed~\cite{bow}, stars with the maximally anisotropic pressure can approach the black-hole limit continuously, i.e. their compactness can be as high as $M/R\sim1/2$ in the nonrotating case. They confirmed that the universality of the I-Love-Q relations hold for stars with weakly anisotropic pressure and that the tidal Love number and the quadrupole moment of anisotropic fluid stars continuously approach their corresponding black hole's values as the compactness increases. On the other hand, they observed that, for strongly anisotropic pressure models, the behavior of the Love number and quadrupole moment are quite different from those of standard neutron stars and quark stars. These peculiar models with strongly anisotropic pressure have negative Love number and quadrupole moment, in contrast with the neutron-star case. A negative value of the spin-induced quadrupole moment basically indicates that the matter distribution of the object is prolate. Indeed, Yagi and Yunes observed that the rotating stars with strongly anisotropic pressure become prolate despite of the centrifugal force, due to a strongly anisotropic pressure~\cite{yagi5}. Similar results have been also obtained for another peculiar compact star model, namely for thin-shell gravastars. The original gravastar model has been proposed by Mazur and Mottola~\cite{gravastar} as an alternative to the final state of the stellar evolution of very massive stars, whose collapse would form a black hole in the standard scenario. The thin-shell gravastar model is a simplified version of the original gravastar and is composed of a vacuum core with a positive cosmological constant (de Sitter core) surrounded by an infinitesimally thin shell, which is required to match the interior core to the Schwarzschild exterior metric~\cite{vis}. Although this scenario lacks a precise formation mechanism, the de Sitter core is assumed to appear through a quantum phase transition in the vicinity of the would-be event horizon during the gravitational collapse of very massive objects. Gravastars are a hypothetical exotic compact objects which can be as compact as black holes but are free from the theoretical problems associated with an event horizon and a spacetime singularity~\cite{gravastar}. The possibility of testing the gravastar scenario with electromagnetic and gravitational-wave observations has been recently argued in Refs.~\cite{BN,cr,Cardoso:2016rao,Giudice:2016zpa}. Some models of thin-shell gravastars are stable against small radial disturbances~\cite{vis} although nonlinear instability might occur for those models whose compactness is larger than $1/3$ (and therefore possess a light ring)~\cite{jk,vitor3} and for highly-spinning compact models due to the ergoregion instability~\cite{Cardoso:2007az}. Recently, one of us investigated the I-Love-Q relations for a particular model of thin-shell gravastar in which the energy density of the thin shell vanishes~\cite{pani2}. Similarly to the case of fluid stars with maximum anisotropic pressure, in this model, $I$, $\lambda$ and $Q$ smoothly connect to their black-hole values as the compactness increases. Furthermore, low-compactness gravastars have negative Love number and quadrupole moment and become prolate shaped when rotating. At the same time, by assuming that the equation of state for the thin shell is given by the same form as that given by equilibrium sequences of spherical solutions with fixed values of the gravitational mass, two of us~\cite{uchi2} have studied rotational effects on a thin-shell gravastar in the absence of an external tidal field. Similarly to the case of gravastars with zero thin-shell energy, we obtained prolate shaped rotating gravastars and showed that some solutions may have the same quadrupole moment of a black hole with same mass and spin. This implies a confusion problem, namely these particular solutions cannot be distinguished from a black hole through independent measurements of their mass, spin and quadrupole moment. Given the plethora of applications related to the multipole moments and the tidal deformability of relativistic compact objects, the motivation for the present study is manifold. On the one hand, we extend our previous work~\cite{pani2,uchi2} to investigate the universality and the black-hole limit of the I-Love-Q relations for various thin-shell gravastars with a generic polytropic equation of state for the shell. We use the Hartle-Thorne formalism~\cite{hartle,thorne}, assuming the tidal and rotational effects on the gravastar are sufficiently small, which is a standard assumption to study the tidal Love numbers and the spin-induced multipole moment for compact objects~\cite{chandra,urbanec,thorne2,flanagan,hinderer}. Since the thin-shell gravastar is composed of two distinct spacetimes, namely a de Sitter core and a Schwarzschild exterior in the spherically symmetric case, we need to take account of the junction conditions of spacetime at the thin-shell position, where the two spacetimes are matched together~\cite{israel,ba}. On the other hand, our novel results for the tidal deformability of thin-shell gravastars allow us to investigate the extend to which gravitational-wave observations~\cite{GW150914,GW151226} can constrain gravastar models. Tidal effects enter the two-body inspiral gravitational waveforms at high post-Newtonian order~\cite{flanagan,hinderer} (cf., e.g., Ref.~\cite{Buonanno:2014aza} for a review). Because the tidal Love numbers of static~\cite{bin,Gurlebeck:2015xpa} (and, presumably, also rotating~\cite{pani3a,pani3}) black holes are identically zero, any gravitational-wave measurement of a nonvanishing tidal deformability would imply that one of the two objects is not a black hole. Conversely, gravitational-wave observations of compact-binary inspirals~\cite{GW150914,GW151226} may be used to put upper bounds on the tidal Love number of the two bodies, thus constraining exotic alternatives. The plan of this paper is the following. In Sec.~II, we briefly give the formulation for constructing models of distorted thin-shell gravastars whose small deformation is caused by the centrifugal and the tidal forces, respectively. In Sec.~III, we present numerical results for stationary and static axisymmetric models. We first construct spherically symmetric, thin-shell gravastars and examine their radial stability, which depends on the equation of state for the thin-shell matter. Next, we present the results for the I-Love-Q relations of thin-shell gravastars with polytropic equation of state and investigate their black-hole limit. Finally, in Sec.~IV we discuss how a measurement of the tidal deformability from the gravitational-wave detection of a compact-binary inspiral can be used to constrain models of thin-shell gravastars. We conclude in Sec.~V. In this study, we use geometrized units in which $G=v_c=1$, where $G$ and $v_c$ are the gravitational constant and the speed of light, respectively. Note that, as mentioned later, we employ the symbol ``$c$'' to denote the compactness of the object, i.e., $c:=M/R$ with $M$ and $R$ being the mass and radius of the object with spherical symmetry. | We have studied the rotational and tidal quadrupole deformations of a thin-shell gravastar, which are characterized by the tidal Love number and the rotational quadrupole moment, respectively. We worked in a perturbative regime in which the tidal and the rotational effects are described by small deviations from spherical symmetry. We considered a thin shell made of a polytropic fluid and studied in detail the cases when the polytropic indices are $n=1$ and $n=3$. We found that the I-Love-Q relations of a thin-shell gravastar are drastically different from those of an ordinary compact star like a neutron star. The Love number and quadrupole moment are negative for less compact models and the I-Love-Q relations continuously approach the black-hole limit. The appearance of negative values of the Love number and quadrupole moment means that the elongation direction of the matter distribution on the meridional cross section is turned by $90$~degrees relative to the case of an ordinary compact star. The reason for the appearance of this counterintuitive deformation is not entirely clear, but seems to be related to the strongly anisotropic stress of the thin shell in the horizontal and vertical directions, and to the peculiar equation of state of the gravastar's interior. Similar counterintuitive properties have been found also for others infinitesimally-thin shells deformed by rotational or tidal effects~\cite{dlc,pf,pf2,pani2,uchi2}. We considered a wide range of polytropic equations of state for the thin-shell matter. Within this range, there is no universality in the I-Love-Q relations, unlike the case of neutron stars and quark stars for which such relations depend only mildly on the stellar equation of state, at the level of a few percent. Although (some degree of) approximate universality might be restored by restricting to a subclass of equations of state, it is nevertheless interesting that thin-shell gravastars provide an example of very compact objects for which the universality of the I-Love-Q relations is manifestly broken. We also evaluated the stability of the unperturbed spherical gravastars with a polytropic perfect-fluid thin shell against small radial perturbations. It is found that less compact models are stable and that the stability changes at the maximum mass models of equilibrium sequences characterized by a single equation of state. The stability analysis implies that the gravastars become unstable as they approach the black-hole limit. Therefore, we inevitably had to use unstable models to investigate the I-Love-Q relations very close to the black-hole limit. The instability region shrinks to zero as $k$ increases for a fixed value of $n$. The $k\to\infty$ limit corresponds to the case of a thin shell with vanishing energy density studied in Ref.~\cite{pani2}. Despite some differences in our analysis (we corrected the definition of the moment of inertia used in Ref.~\cite{pani2} and kept the properties of the thin-shell matter fixed along different sequences of solutions) our results are qualitatively similar to those presented in Ref.~\cite{pani2} and extend the latter to more realistic configurations. In this study, we focus on the case of a perfect-fluid thin shell. However, a thin shell with anisotropic pressure might be more reasonable. For example, if the phase transition replacing the would-be horizon~\cite{gravastar} is due to a scalar field, the latter would naturally give rise to anisotropic stresses. Thus, it would be interesting to investigate in detail the case of gravastars with an anisotropic thin shell. Likewise, it would be interesting to extend our computation to the case of finite-thickness shells, although that would likely require a numerical integration of the field equations within the shell. A detailed analysis of the observational constraints on the tidal deformability of gravastars coming from gravitational-wave measurements goes beyond the scope of this work. Nonetheless, our results suggest that near-future constraints on the tidal Love numbers of compact objects from inspiral gravitational waveforms can place very stringent lower limits on the compactness of the two objects, ruling out several gravastar models and giving further support to the fact that events like GW150914 and GW151226 are coalescences of a pair of black holes. So far, studies of rotating models of gravastars are restricted to the case of slow rotation. On the other hand, it is likely that a substantial fraction of black holes are rapidly spinning, and this is surely the case for black holes formed in the coalescences recently detected in the gravitational-wave band by the LIGO interferometer~\cite{GW150914,GW151226}. Highly-spinning gravastars might be unstable against the ergoregion instability~\cite{Cardoso:2007az}, and a detailed studied is therefore necessary to assess their viability as exotic compact objects and black-hole mimickers. These investigations remain as future work. | 16 | 7 | 1607.03593 |
1607 | 1607.01769_arXiv.txt | Upcoming cosmic microwave background (CMB) experiments will measure temperature fluctuations on small angular scales with unprecedented precision. Small-scale CMB fluctuations are a mixture of late-time effects: gravitational lensing, Doppler shifting of CMB photons by moving electrons (the kSZ effect), and residual foregrounds. We propose a new statistic which separates the kSZ signal from the others, and also allows the kSZ signal to be decomposed in redshift bins. The decomposition extends to high redshift, and does not require external datasets such as galaxy surveys. In particular, the high-redshift signal from patchy reionization can be cleanly isolated, enabling future CMB experiments to make high-significance and qualitatively new measurements of the reionization era. | On large angular scales ($l \lsim 2000$), anisotropy in the cosmic microwave background is mainly sourced by fluctuations at redshift $z \approx 1100$. On smaller angular scales ($l \gsim 2000$), this ``primary'' anisotropy is exponentially suppressed, and CMB fluctuations are mainly a mixture of several ``secondary'' or late-time effects. Among secondary effects with the same blackbody spectrum as the primary CMB, the largest are gravitational lensing, and the kinematic Sunyaev-Zel'dovich (kSZ) effect. The kSZ effect refers to Doppler shifting of CMB photons as they scatter on radially moving inhomogeneities in free electron density \cite{Ostriker:1986fua,Sunyaev:1980vz,Sunyaev:1972eq}. The kSZ anisotropy can be roughly decomposed into a ``late-time'' contribution from redshifts $z \lsim 3$, when inhomogeneities are large due to gravitational growth of structure, and a ``reionization'' contribution from redshift $z \sim 7$, when the ionization fraction is expected to be inhomogeneous during ``patchy'' reionization \cite{Battaglia:2012im,McQuinn:2005ce,Park:2013mv, Alvarez:2015xzu, Zahn:2011vp}. \begin{figure}[t] \centerline{\includegraphics[width=7cm]{figs/cltt_v3.pdf}} \caption{Fiducial model for the CMB temperature power spectrum $C_l^{TT}$ used throughout this paper, split into primary, lensing, late-time kSZ, and reionization kSZ contributions.} \label{fig:cltt} \end{figure} In Fig.~\ref{fig:cltt} we compare contributions to the temperature power spectrum $C_l^{TT}$ from weak lensing of the CMB, late-time kSZ, and reionization kSZ, in a fiducial model to be described shortly. Individually, these three contributions are very interesting. Gravitational lensing depends on cosmological parameters such as neutrino mass~\cite{Lewis:2006fu}, late-time kSZ probes the distribution of electrons in dark matter halos as well as the large-scale velocity field, and reionization kSZ may provide the first observational window on patchy reionization, which will shed light on the formation of first stars and other sources of ionizing photons. Although the total power spectrum will soon be measured very precisely at high $l$ \cite{Calabrese:2014gwa}, it is unclear how well these signals can be disentangled, since all three components have large astrophysical modelling uncertainties, and the two kSZ contributions are essentially degenerate at the power spectrum level. In this paper we will propose a higher-order statistic which isolates the kSZ signal, and moreover gives information about its source redshift dependence, allowing the late-time and reionization kSZ to be separated. This will complement measurements from future 21cm experiments \cite{Furlanetto:2015apc,Morales:2009gs}. We describe the intuitive idea here, with a more formal description in the next section. We first recall that to a good approximation, the kSZ power spectrum may be written as an integral \cite{Ma:2001xr}: \be C_l^{\rm kSZ} = \int dz \, Q(z) \, \big\langle v_r(z)^2 \big\rangle \, P_e\!\left( \frac{l}{\chi(z)}, z \right) \label{eq:clksz_integral} \ee where $P_e(k,z)$ is the free electron power spectrum, $\langle v_r^2 \rangle = \langle v^2 \rangle/3$ is the mean squared radial velocity, and $\chi(z)$ is the comoving distance to redshift $z$. The radial weight function $Q(z)$ is given by \be Q(z) = T_{\rm CMB}^2 \, \frac{H(z)}{\chi(z)^2} \left( \frac{d\bar\tau}{dz} \right)^2 e^{-2\bar\tau(z)} \ee where $d\bar\tau/dz$ is the optical depth per unit redshift. Throughout the paper we will use the following notation frequently. Let $\bar K$ be the sky-averaged small-scale power spectrum in a fixed high-$l$ band (say $3000 \le l \le 5000$). For each direction $\n$ on the sky, let $K(\n)$ be the locally measured small-scale power spectrum near sky location $\n$. The precise definitions of $\bar K$ and $K(\n)$ will be given in the next section. It is intuitively clear that $K(\n)$ will be better approximated by using the actual realization of $v_r^2(\n,z)$ along the line of sight in direction $\n$ in the integral in Eq.~(\ref{eq:clksz_integral}), rather than the cosmic average $\langle v_r^2 \rangle$. This leads to anisotropy in $K(\n)$ on large angular scales. To estimate the level of anisotropy, suppose we divide the line of sight into segments of size 50 Mpc (the coherence length of the velocity field), and roughly model the radial velocity $v_r$ as an independent Gaussian random number in every segment. Since the line of sight is $10^4$ Mpc in length, $K(\n)$ can be roughly modelled as the sum of squares of $N=200$ independent Gaussians. This suggests that fluctuations in $K(\n)$ between different lines of sight are of fractional size $\sqrt{2/N} \approx 0.1$. Rephrasing, if we measure the CMB in two regions of sky separated by more than $\sim$1 degree, so that the lines of sight sample independent realizations of the velocity field, the kSZ power spectra will differ by $\approx 10\%$. This is a large non-Gaussian effect which is not present for lensing and other secondaries, allowing statistical separation of the kSZ signal. In fact we can go further by considering $C_L^{KK}$, the angular power spectrum of $K(\n)$. Suppose we write $C_L^{KK}$ as a sum of contributions from multiple source redshift bins. In the next section we will show (Fig.~\ref{fig:clkk_model}) that the contribution from redshift $z$ has a broad peak at wavenumber $L_* \sim k_* \chi(z)$, where $k_* \approx 0.01$ $h$ Mpc$^{-1}$. Thus the shape of $C_L^{KK}$ is source redshift dependent. In the general case where $C_L^{KK}$ is a sum over redshift bins, we can ``deconvolve'' the observed $C_L^{KK}$ to infer the contribution from each bin, thus separating the late-time and reionization kSZ signals. The main advantage of this method (compared to an analysis based on $C_l^{TT}$) is its robustness: we can make statements about reionization which do not depend on precise modelling of the other contributions. | In this section, we will present signal-to-noise forecasts. The first type of forecast we will consider is ``single-bin detection'': total signal-to-noise of the kSZ $C_L^{KK}$ summed over all source redshifts, marginalized over an arbitrary constant $\delta C_L^{KK}$ as previously described. In our fiducial model, a single-bin detection would get 86\% of its signal-to-noise from reionization and 14\% from the late-time kSZ. Therefore a single-bin detection of $C_L^{KK}$ at the expected level would be strong evidence for patchy reionization in the CMB. The next observational milestone would be a ``two-bin detection'', in which we fit for the overall amplitude of the high-$z$ signal ($z \ge 4$), marginalized over an arbitrary multiple of the low-$z$ contribution ($z \le 4$). A two-bin detection would measure the amplitude of the patchy reionization signal without any assumptions on the low-$z$ amplitude. Finally, we consider a ``three-bin detection'': detection significance of the $z \ge 5$ contribution, marginalized over independent redshift bins with $0 \le z \le 2.5$ and $2.5 \le z \le 5$. A three-bin detection would establish the bimodal redshift dependence of the kSZ sources, with peaks at late time and during reionization, and little or no power in between, which would also provide a powerful check on systematics. The precise definitions are as follows: given $N$ contributions $(\delta C_L^{KK})_1$, $\cdots$, $(\delta C_L^{KK})_N$ to $C_L^{KK}$ such that $C_L^{KK} = \sum_i A_i (\delta C_L^{KK})_i $ (for example $N$ redshift bins), we forecast signal-to-noise on the amplitudes $A_i$ by computing the $N$-by-$N$ Fisher matrix \be F_{ij} = \frac{f_{\rm sky}}{2} \sum_{L=L_{\rm min}}^{L_{\rm max}} (2L+1) \frac{(\delta C_L^{KK})_i \, (\delta C_L^{KK})_j}{(C_L^{KK})_{\rm tot}^2} \label{eq:fisher} \ee The signal-to-noise of $(\delta C_L^{KK})_i$, marginalizing over signals $j \ne i$, is given by $S/N = (F^{-1}_{ii})^{-1/2}$. The significance of our ``$N$-bin detection'', where $N=1,2,3$, is defined to be the signal-to-noise of the highest redshift bin, taking $(C_L^{KK})_{\rm tot}$ to be the sum of contributions from the lower redshift bins plus reconstruction noise, and marginalized over the other redshift bins plus a contribution of the form $\delta C_L^{KK}=\mbox{constant}$. The maximum multipole $L_{\rm max}$ in Eq.~(\ref{eq:fisher}) will depend in practice on the extent to which secondary contributions to $C_L^{KK}$ can be modelled as $L$ increases. As a fiducial value, we have used $L_{\rm max}=300$ here. In Fig.~\ref{fig:forecast} we show forecasted signal-to-noise as functions of instrumental noise level and beam size. These results include improvements from a generalization of the Fisher matrix in Eq.~(\ref{eq:fisher}) in which multiple $K$-fields are defined corresponding to bins in CMB wavenumber $l$. However, we find that the only case where this significantly improves signal-to-noise is the three-band detection with $\theta_{\rm FWHM} \lsim 2'$. It is seen that the signal-to-noise is a steep function of noise level, favoring a deep small-field observing strategy. However, for surveys smaller than a few hundred square degrees, there is a signal-to-noise penalty beyond the simple $f_{\rm sky}$ scaling in Fig.~\ref{fig:forecast}, since $C_L^{KK}$ cannot be measured on super-survey scales $L \le L_{\rm min} = (2\pi/\theta_{\rm surv})$. This signal-to-noise penalty ranges from 5--10\% for a 1000 deg$^2$ survey, and 25--50\% for a 100 deg$^2$ survey, depending on the noise level and forecast chosen. Including this penalty, some example surveys which achieve $N$-bin detections are as follows. A $3\sigma$ one-bin detection can be achieved by a survey with area $A=500$ deg$^2$, noise $\Delta_T = 4$ $\mu$K-arcmin, and beam $\theta_{\rm FWHM}=1.4$ arcmin. Likewise two-bin and three-bin $3\sigma$ detections can be achieved by surveys with $(A,\Delta_T,\theta_{\rm FWHM}) = (900, 3, 1)$ and $(2400, 2, 1)$ respectively. An ambitious future survey with 1 $\mu$K-arcmin noise, 1 arcmin beam, and $f_{\rm sky}=0.5$ can achieve 1-bin, 2-bin, and 3-bin detections with significance 279$\sigma$, 44$\sigma$, and 16$\sigma$. \begin{figure} \centerline{\includegraphics[width=7cm]{figs/sn_forecast_v4.pdf}} \caption{Forecasted $S/N$ for 1-bin, 2-bin, and 3-bin detections as defined in the text, for varying noise level, and beam size $\theta_{\rm FWHM}=1'$ (blue/upper curves) or $\theta_{\rm FWHM}=3'$ (red/lower curves).} \label{fig:forecast} \end{figure} In this paper, we have identified a new non-Gaussian signal in the CMB which is a distinctive observational signature of the kSZ effect. It should soon be detectable, and an exciting milestone will be a ``clean'' detection of patchy reionization, with minimal assumptions on modelling of other CMB secondaries. Future experiments such as CMB-S4 will have sufficient signal-to-noise to measure the signal with more granularity and constrain the redshift and wavenumber dependence of the kSZ sources, opening up a qualitatively new observational window on the epoch of reionization. \vskip 0.2cm {\em Acknowledgements.} We thank Joel Meyers and Alex van Engelen for discussions and initial collaboration. We are also grateful to Nick Battaglia, Emmanuel Schaan and David Spergel for discussion. KMS was supported by an NSERC Discovery Grant and an Ontario Early Researcher Award. SF was supported by the Miller Institute at the University of California, Berkeley. Some computations were done on the Scinet cluster at the University of Toronto. Research at Perimeter Institute is supported by the Government of Canada through Industry Canada and by the Province of Ontario through the Ministry of Research \& Innovation. | 16 | 7 | 1607.01769 |
1607 | 1607.03907_arXiv.txt | I study the location of the $\gamma$-ray emission in blazar jets by creating a Compton-scattering approximation valid for all anisotropic radiation fields in the Thomson through Klein-Nishina regimes, which is highly accurate and can speed up numerical calculations by up to a factor $\sim10$. I apply this approximation to synchrotron self-Compton, and external Compton-scattering of photons from the accretion disk, broad-line region (BLR), and dust torus. I use a stratified BLR model and include detailed Compton-scattering calculations of a spherical and flattened BLR. I create two dust torus models, one where the torus is an annulus, and one where it is an extended disk. I present detailed calculations of the photoabsorption optical depth using my detailed BLR and dust torus models, including the full angle dependence. I apply these calculations to the emission from a relativistically moving blob traveling through these radiation fields. The ratio of $\gamma$-ray to optical flux produces a predictable pattern that could help locate the $\gamma$-ray emission region. I show that the bright flare from 3C 454.3 in 2010 November detected by the Fermi Large Area Telescope is unlikely to originate from a single blob inside the BLR since it moves outside the BLR in a time shorter than the flare duration, although emission by multiple blobs inside the BLR is possible; and $\gamma$-rays are unlikely to originate from outside the BLR from scattering of photons from an extended dust torus, since then the cooling timescale would be too long to explain the observed short variability. | \label{intro} Blazars are active galactic nuclei (AGN) with jets of nonthermal plasma moving at relativistic speeds oriented close to our line of sight. They produce nonthermal radiation across the electro-magnetic spectrum, from radio to $\g$ rays, the lower portion of which is almost certainly produced by nonthermal synchrotron. Blazars are classified based on the strength or weakness of broad lines in their optical spectra as FSRQs or BL Lacs, respectively \citep[e.g.][]{marcha96,landt04}, and on the location of their synchrotron peak \citep[$\nu_{pk}$;][]{abdo10_sed} in $\nu F_{\nu}$ spectral energy distributions (SEDs) as low-synchrotron peaked (LSP; $\nu_{pk}<10^{14}\ \Hz$), intermediate-synchrotron peaked (ISP; $10^{14}\ \Hz < \nu_{pk} < 10^{15}\ \Hz$), or high-synchrotron peaked (HSP; $10^{15}\ \Hz\ < \nu_{pk}$). In the 0.1--300 GeV energy range, blazars dominate the sky in terms of number of associated sources \citep{acero15_3fgl}. Hadronic models for $\g$-ray production in these sources often face difficulties due to requiring excessive jet powers, especially for FSRQs and LSPs \citep{boett13,zdziarski15,petro15_3c279,petro16}, although there are some examples, especially HSPs, where jet powers are more reasonable \citep[e.g.,][]{cerruti15}. Inverse Compton scattering of soft photons by relativistic nonthermal electrons in the jet is usually invoked as the most likely mechanism for $\g$-ray production in FSRQs and some LSPs. Possible seed photons for Compton scattering include the synchrotron photons produced by the same population of electrons that produces the $\g$ rays \citep[known as synchrotron self-Compton, or SSC;][]{bloom96}; or seed photons from another portion of the jet \citep[e.g.,][]{ghisellini05,macdonald15,sikora16}; or by seed photons external to the jet entirely, for instance, from the accretion disk \citep{dermer92,dermer93}, broad line region \citep[BLR;][]{sikora94}, or dust torus \citep{kataoka99,blazejowski00} in the object; or from the cosmic microwave background \citep[CMB;][]{boett08,yan12_1101,meyer15,sanchez15,zacharias16}. The dominant seed photon source depends critically on the location in the jet of the emitting region. In order of increasing distance from the black hole, the dominant external seed photon source could be the accretion disk, the BLR, the dust torus, and/or the CMB. The dominant seed photon source, and the location of the primary emitting region, is an important topic in the understanding of blazar jets that has not yet been resolved. I endeavor to make progress in answering this question by providing detailed calculations of Compton scattering of the relevant external radiation fields. Calculations of Compton scattering of various external radiation fields (external Compton or EC hereafter) using the full Compton cross section, valid in the Thomson and extreme Klein-Nishina regimes, including anisotropic external radiation fields, have been explored by many authors \citep[e.g.,][]{boett97,boett00,dermer09,hutter11,hunger16}. These calculations can be quite numerically intense, especially if calculations are repeated numerous times in, for example, fitting routines \citep{finke08_SSC,mank10,mank11,yan13,cerruti13_fit} or multi-zone models \citep[e.g.,][]{jamil12,joshi14}. Often, $\delta$-function approximations, valid in the Thomson regime, are used to approximate Compton scattering processes \citep[e.g.,][]{dermer92,dermer93}. In these approximations, {\em all} of the scattering photons are assumed to have the same energy, that of the mean scattered photon energy. A $\delta$-function approximation valid at all energies, in the Thomson through extreme Klein-Nishina regimes, was developed by \citet{moderski05} for isotropic radiation fields, which I made use of recently to compute theoretical power spectral densities and Fourier-frequency dependent time lags of blazar light curves \citep{finke15}. In the present manuscript, I generalize the $\delta$-function approximation of \citet{moderski05} to anisotropic radiation fields (Section \ref{deltafcnsection}) and apply it to EC for several external isotropic and anisotropic external radiation fields (Section \ref{ECsection}), and, for completeness, to SSC (Section \ref{SSCsection}), all in the context of relativistic jets. Researchers may find this $\delta$-function approximation useful for Compton-scattering calculations in other astrophysical contexts besides blazars, such as microquasars \citep[e.g.,][]{gupta06a,dubus08,dubus10,zdziarski14}, colliding winds of massive stars \citep[e.g.,][]{reimer06_wind}, or gamma-ray bursts \citep[e.g.,][]{lu15}. The speed and accuracy of the approximations are explored in Section \ref{numericalsection}. The $\delta$-function approximations are used to derive the beaming pattern for the scattering of various external radiation fields by relativistic jets in Section \ref{beampatternsection}. Another process necessary to the Compton-scattering model for blazar jet emission is $\g\g$ absorption. The interaction of $\g$ rays with soft photons from the accretion disk, BLR, and dust torus can limit the escape of $\g$ rays. This process is explored in Section \ref{absorbsection}. As the $\g$-ray emitting region moves at relativistic speeds, it travels through regions where various external radiation fields dominate the Compton scattering process. In Section \ref{signaturesection} I look at the effect this would have on the ratio of $\g$-ray to optical flux that one would expect as a function of time. This can give critical clues to the location of the $\g$-ray emitting region in blazars. In particular, I apply the calculations presented here to the optical and $\g$-ray light curves of the giant flare in 3C 454.3 in 2010 November. I conclude with a summary in Section \ref{discussion}. In Appendix \ref{BLRmodel}, I present a simple model for determining the luminosities and radii of line emission in a stratified BLR based on a composite quasar spectrum from the Sloan Digital Sky Survey (SDSS). | \label{discussion} The location of the $\g$-ray emitting region is blazar jets is poorly known; modeling EC jet emission is necessary to determine its origin, but doing this correctly requires the full angle-dependent Compton cross-section. Consequently, I have created a novel $\delta$-function approximation for Compton scattering that is fully angle dependent, valid in both the Thomson and extreme Klein-Nishina regimes. It assumes that all of the scattered photons have the same energy of the mean scattered photon. It allows one to eliminate an integral in Compton scattering calculations, so that in applications for scattering in blazar jets, it is numerically faster by up to approximately a factor of 10. It is likely to have applications to other high-energy astrophysical situations where accurate Compton scattering calculations are important. I have used this approximation to derive formulae for Compton scattering of external radiation fields of interest to blazars; namely, external radiation originating from the accretion disk, BLR, and dust torus. The Compton-scattered accretion disk formula is essentially the same as the one from \citet{dermer09}, only with the new approximation. Formulae for Compton-scattering of BLR photons are provided for two geometries, considering the line emission as originating from spherical shell and from a flat annulus. I have also provided a method for estimating the relative radii and luminosities of the various broad lines that make up the BLR (Appendix \ref{BLRmodel}). I also provide formulae for computing the scattering of dust torus photons using two dust torus models, one as a thin annulus, and one as an extended torus made up of many dust clouds. The latter is similar to the one described by \citet{sikora13}, although I take into account the full angle dependence in the Compton-scattering calculation. I derive beaming patterns for scattering of spherical shell external radiation fields or annulus external radiation fields in the Thomson regime, the latter of which could represent either a broad emission line or a dust torus. I show that if the shell or annulus has a radius $R_{\rm re}$, then for distances $r\ll R_{\rm re}$, the beaming pattern resembles that of an isotropic radiation field; and for distances $r\gg R_{\rm re}$, it resembles the beaming pattern of a point source behind the jet. The beaming pattern for the spherical shell and annulus seed radiation field geometries are similar enough to each other that beaming pattern studies are unlikely to distinguish between them. Using my models for radiation from the BLR and dust torus, I have computed the $\g\g$ absorption optical depth as a function of $\g$-ray energy and distance from the black hole. I believe these are the most detailed of this sort of calculation that has yet been performed. This calculation shows that photons below $\approx20\ \GeV$ in the source frame are always able to escape unabsorbed. For $\g$ rays emitted $r \ga 10R({\rm Ly}\alpha)$, $\tau_{\g\g}<1$ everywhere, at least for the fiducial parameters I used. For emission of $\g$ rays in between these radii, the escape of $\g$ rays is dependent on the details of the BLR structure, which I have realistically computed. The flattened, ring geometry lowers the opacity a bit, and the threshold is a bit higher, but the conclusions on the escape of $\g$ rays are not changed drastically. Finally, I explore the ratio of the Compton-scattered flux to synchrotron flux for a blob moving out in a jet. This ratio, and how it changes with time, can be a clue to finding the location of the $\g$-ray emitting region. For flares lasting longer than $\sim$\ a day, the flaring region rapidly moves out of the BLR (for fiducial values $\G=\delta_D=40$), so that scattering of BLR photons by a single blob cannot explain these flares, although emission by many separate blobs within the BLR would be possible. Scattering of photons from an extended dust torus at $\sim10\ \pc$ from the black hole is not able to cool photons fast enough to explain the rapid flares seen from FSRQs. If $\g$-ray flares originate from EC at these distances, another source of seed photons is needed to rapidly cool the electrons producing these flares. A possible alternative photon source is other regions of the jet: such as a sheath surrounding the emitting region \citep{ghisellini05,macdonald15,sikora16} or scattering of photons from a standing shock or other moving regions of the jet \citep{marscher14}. The location of the $\g$-ray emitting region in blazars and the seed photon source is still unclear. | 16 | 7 | 1607.03907 |
1607 | 1607.08538_arXiv.txt | I start by providing an updated summary of the penalized pixel-fitting (\ppxf) method, which is used to extract the stellar and gas kinematics, as well as the stellar population of galaxies, via full spectrum fitting. I then focus on the problem of extracting the kinematic when the velocity dispersion $\sigma$ is smaller than the velocity sampling $\Delta V$, which is generally, by design, close to the instrumental dispersion $\sigma_{\rm inst}$. The standard approach consists of convolving templates with a discretized kernel, while fitting for its parameters. This is obviously very inaccurate when $\sigma\la\Delta V/2$, due to undersampling. Oversampling can prevent this, but it has drawbacks. Here I present a more accurate and efficient alternative. It avoids the evaluation of the under-sampled kernel, and instead directly computes its well-sampled analytic Fourier transform, for use with the convolution theorem. A simple analytic transform exists when the kernel is described by the popular Gauss-Hermite parametrization (which includes the Gaussian as special case) for the line-of-sight velocity distribution. I describe how this idea was implemented in a significant upgrade to the publicly available \ppxf\ software. The key advantage of the new approach is that it provides accurate velocities regardless of $\sigma$. This is important e.g. for spectroscopic surveys targeting galaxies with $\sigma\ll\sigma_{\rm inst}$, for galaxy redshift determinations, or for measuring line-of-sight velocities of individual stars. The proposed method could also be used to fix Gaussian convolution algorithms used in today's popular software packages. | The kinematics of the stars and gas in galaxies is a key ingredient in our understanding of how they form and evolve. Nowadays, this information is routinely extracted from integral-field spectroscopic (IFS) data, which provide a three dimensional view of galaxies. The technology is now available on all major telescopes and has become the standard way of obtaining spectra for galaxy evolution studies \citep[see][for a review]{Cappellari2016}. The IFS data provide a ``fossil record'' of galaxy formation. Specifically, the IFS kinematics of the stars allow one to distinguish a galaxy assembly dominated by gas poor merging from a growth driven by gas accretion and star formation \citep[e.g.][]{Emsellem2011,Cappellari2013p20,Naab2014}. Dynamical models based on IFS kinematics allow one to infer galaxy mass distributions to study scaling relations \citep[e.g.][]{Cappellari2013p15,Scott2015}, the stellar and dark matter content \citep[e.g.][]{Cappellari2012,Cappellari2015dm}, or measure black hole masses \citep[e.g.][]{Krajnovic2009,Seth2014,Walsh2016}. The gas content and kinematics tell us about the role of gas accretion in galaxy assembly \citep[e.g.][]{Sarzi2006,Davis2011b,Barrera-Ballesteros2015} or the mechanism that regulates star formation in galaxies \citep[e.g.][]{Alatalo2011,Cheung2016,Ho2016}. For more than a decade, this information has been extracted from IFS data by surveys targeting one galaxy at a time (SAURON \citealt{deZeeuw2002}; DiskMass \citealt{Bershady2010}; ATLAS$^{\rm 3D}$ \citealt{Cappellari2011a}; CALIFA \citealt{Sanchez2012}). But the observational panorama is undergoing a revolution, with the arrival of large multiplexed IFS surveys targeting 10--20 galaxies at a time. Thousands of galaxies have already been observed in this manner (MaNGA \citealt{Bundy2015}; SAMI \citealt{Bryant2015}). These ongoing IFS surveys depend critically on the full spectrum fitting technique to deliver their science. The method is used to extract stellar and gas kinematics, as well as stellar population from the spectral data cubes. But, given the large number of objects and the variety of galaxy morphological types, these surveys are pushing the existing techniques to their limits. This paper is motivated by the existence of these ongoing large surveys and mostly arises from initial experiences with the analysis of the MaNGA data. In fact, one new characteristic of both the MaNGA and SAMI survey is that they are observing large numbers of galaxies (especially spirals) with stellar velocity dispersion $\sigma$ well below the instrumental dispersion $\sigma_{\rm inst}$ of the spectrographs. This situation was until now not very common, as the observers generally tried to target galaxies with $\sigma_{\rm inst}\la\sigma$. However this restriction is actually not necessary to obtain useful kinematic information. In fact, although it is true that $\sigma$ becomes intrinsically difficult to measure reliably when $\sigma\ll\sigma_{\rm inst}$, the kinematics of low-$\sigma$ galaxies is dominated by the stellar velocity $V$ \citep[e.g.][]{Cappellari2016}, which remains a well defined observable. I realized that in this observational regime, all the spectrum fitting approaches, which are nearly universally adopted to extract stellar and gas kinematics, suffer from limitations and can be significantly improved. Here in \autoref{sec:concepts} I describe general concepts about kinematic extraction, in \autoref{sec:ppxf} I give an updated summary of the \ppxf\ method, in \autoref{sec:low_dispersion} I discuss problems of the current approach and propose a clean solution, I summarize my paper in \autoref{sec:summary}. | \label{sec:summary} In the first part of the paper, I provided an overview, or tutorial, of general concepts useful to understand and properly interpret the extraction of kinematics from galaxy spectra. I tried to clarify in particular the questions I received more often, over more than a decade, from users of my publicly available \ppxf\ software. Then I gave an updated overview of the \ppxf\ method. I concentrated especially on the description of features of the method which were included after the publication of the original paper in 2004, some of which had never been properly explained and precisely documented in the literature. Subsequently, I focused on the problem of extracting kinematics via full spectrum fitting, when the velocity dispersion is smaller than the spectral sampling, which is generally chosen to be the same as the instrumental dispersion. I illustrated the obvious but dramatic problems that arise when one completely ignores the issue, as well as the limitations of the previous solution, which consist of oversampling the spectra. Finally I provided a clean solution to the long-standing under-sampling issue, which consist of using the analytic Fourier transform of the LOSVD in conjunction with the convolution theorem. This approach completely removes the need for oversampling and makes the full spectrum fitting method suitable for measuring reliable kinematics at any velocity dispersion. This is especially crucial for the mean velocity, which now becomes a well-determined quantity even when the dispersion becomes negligible, and consequently impossible to reliably recover from real data. The approach described in this paper was implemented in a significant upgrade to the publicly available \ppxf\ code, and is already being used as part of the MaNGA Data Analysis Pipeline (Wesfall et al.\ in preparation). Further tests on real IFS data will be published elsewhere. The proposed solution appears quite natural, however, perhaps surprisingly, it is currently not being used by any of the popular software packages. Given the simplicity of our approach for accurate convolutions, we argue it should become standard practice. | 16 | 7 | 1607.08538 |
1607 | 1607.03136_arXiv.txt | The likelihood function for cosmological parameters, given by e.g. weak lensing shear measurements, depends on contributions to the covariance induced by the nonlinear evolution of the cosmic web. As nonlinear clustering to date has only been described by numerical $N$-body simulations in a reliable and sufficiently precise way, the necessary computational costs for estimating those covariances at different points in parameter space are tremendous. In this work we describe the change of the matter covariance and of the weak lensing covariance matrix as a function of cosmological parameters by constructing a suitable basis, where we model the contribution to the covariance from nonlinear structure formation using Eulerian perturbation theory at third order. We show that our formalism is capable of dealing with large matrices and reproduces expected degeneracies and scaling with cosmological parameters in a reliable way. Comparing our analytical results to numerical simulations we find that the method describes the variation of the covariance matrix found in the SUNGLASS weak lensing simulation pipeline within the errors at one-loop and tree-level for the spectrum and the trispectrum, respectively, for multipoles up to $\ell\leq 1300$. We show that it is possible to optimize the sampling of parameter space where numerical simulations should be carried out by minimising interpolation errors and propose a corresponding method to distribute points in parameter space in an economical way. | \label{sec:1} Measurements of cosmological parameters and investigations into the properties of gravity on large scales are the focus of a number of upcoming surveys of the cosmic large-scale structure. These investigations require probing how the expansion dynamics of the Universe and the gravitational model affect the growth rate of structures, as well as understanding the relation between redshift and distance. A tool combining both these sources of cosmological information is weak gravitational lensing \citep[e.g.][]{Kaiser1998, Bacon2000, Kaiser2000, Maoli2000, Mellier2000, Bartelmann2001, Kilbinger2003} which, as a line-of-sight integrated quantity of the Newtonian tidal shear field, probes both structure growth and the evolution of the background cosmology by measuring a correlation in the shapes of galaxies. The estimation of cosmological parameters based on large-scale structure observations requires a precise knowledge of the covariance matrix, which describes the cosmic variance, the statistical dependence of the modes of the cosmic matter distribution, and the noise inherent in the surveys. Due to mode coupling in non-linear structure formation the covariance matrix is non-diagonal, acquires large amplitudes on small scales, and renders the statistical properties of the cosmic matter distribution non-Gaussian; In this respect cosmological large-scale structure observations differ significantly from observations of primary covariance matrix-fluctuations \citep{Komatsu2011}, where the assumption of Gaussian statistics is very good. The scaling of non-linear structure growth with cosmological parameters is necessarily non-linear, which is immediately apparent in perturbative approaches \citep{Bernardeau1994, Bernardeau1994a, Taruya2002}. Each order of perturbation theory is characterised by a different dependence on cosmology, and in assembling a perturbation series these dependences are mixed by superposition. Mode coupling in non-linear structure formation generates off-diagonal entries in the covariance matrix \citep[e.g.][etc.]{Scoccimarro1999a, Cooray2001, Takada2007, Takada2009, Sato2009, Kayo2012} and therefore reduces the information content of the density field \citep[e.g.][]{Hu2003, Takada2007,Sato2009, Sato2011}. On the other hand, fluctuations in the cosmic density field are strongly amplified by non-linear structure formation, which allows measurements on small scales which are otherwise inaccessible due to the sparsity of galaxies. Future experiments such as the Euclid mission\footnote{http://www.euclid-ec.org} \citep{Laureijs2011} will use the weak gravitational lensing effect to probe the cosmic web on scales deep in the non-linear regime \citep[e.g.][]{Benjamin2007, Laureijs2011, VanWaerbeke2013, Kitching2014}. In fact, Euclid's anticipated weak lensing signal, with a significance of close to $1000\sigma$, is largely generated by non-linear scales. Because non-linear structure formation cannot yet be fully described by analytical methods, estimates of the covariance matrix require simulations of cosmic structure formation. Due to the large volume of future surveys and the necessity to observe at non-linear scales, cosmological simulations require both large volumes and high resolutions. In addition, a large suite of statistically equivalent simulations is required to estimate covariance matrices using ensemble-averaging. This estimation needs to be undertaken throughout the anticipated parameter space, because non-linear structure formation depends strongly on the choice of cosmological parameters. Standard spatially flat dark energy cosmologies typically have six parameters. Thus, even a rather coarse sampling of parameter space would require a tremendous number of $N$-body ray-tracing simulations \citep{Fosalba2008, Hilbert2009} or other techniques such as line-of-sight integrations \citep{Kiessling2011}, which are, up to now, the only robust method to determine the mode coupling and induced higher order cumulants to the desired accuracy. The computational load to produce large suites of simulations at Gpc-scales, while retaining resolution at sub-Mpc-scales, quickly becomes prohibitive. As a consequence, it is inevitable, that variations in the covariance matrices in parameter space are being investigated \citep{Eifler2009} and a way to interpolate between these points must now be developed. In this context a number of questions arise: $(i)$ How strong are the variations of the covariance matrix with varying cosmological parameters? $(ii)$ To which cosmological parameters is the covariance matrix most sensitive? $(iii)$ Is there a way of predicting variations and in which directions in parameter space the strongest variations are encountered? $(iv)$ Is it possible to decompose changes to the shape, size and orientation of the covariance matrix in a geometrically clear way? $(v)$ What would be sensible choices of cosmological parameters for simulations in order to cover the relevant parameter space economically? $(vi)$ Is there a natural way to interpolate between covariance matrices from numerical simulations? Recently \citet{Schafer2016} introduced a method to interpolate between Fisher matrices at different points in parameter space We now intend to apply this formalism to the variation of the covariance matrix of the matter and convergence power spectrum estimators. This should be possible because both Fisher-matrices and covariance matrices share positive-definiteness as a common property, which is required by our formalism. The focus will be on the power spectrum of the weak lensing convergence, as it is directly linked to observables provided by Euclid. Non-linear structure formation on small scales generates a non-Gaussian contribution to the covariance matrix, where we employ Eulerian perturbation theory at tree-level to predict the trispectrum as the lowest-order non-Gaussian contribution \citep{Scoccimarro1999a}. We consider perturbation theory as an easily manageable tool for predicting non-linear corrections to the covariance matrix and do not imply that it describes all non-linearities accurately, but we will check its validity against numerical simulations. The fiducial cosmological model is a spatially flat $\Lambda$CDM model with base parameters $\Omega_\text{m} = 0.3$, $\Omega_\Lambda = 1-\Omega_\text{m}$, $h = 0.7$, $n_\text{s} =1$, $w_0=-1$ and $w_\text{a} = 0$. Moreover, we will use the sum convention throughout this paper, thus implying summation over repeated indices. After a brief review of the lensing observables we will review the covariance matrix theory for the matter and convergence spectrum in Sect. \ref{sec:2}. In Sect. \ref{sec:3} the Lie basis is constructed and applied to the covariance matrix in Sect. \ref{sec:4}, where we also compare the theoretical prediction with simulations. In Sect. \ref{sec:5} we summarize. | \label{sec:5} In this paper we investigated variations of the covariance matrix of cosmic large-scale structure observations, where non-linear structure formation processes generate non-Gaussian and non-diagonal contributions. We described the variation of the covariance matrix with a change of the cosmological model by constructing a basis, and considered as specific examples the matter density and weak lensing convergence power spectra. We worked with an analytical model for non-linear structure formation based on Eulerian perturbation theory and derived non-Gaussian contributions to the covariance matrices by evaluating the spectrum and the trispectrum in third order. This analytical model was juxtaposed with the results from numerical simulations, which showed that the fundamental scaling of the analytical model with the parameters $\Omega_\text{m}$ and $\sigma_8$ was reproduced correctly. The covariance matrix of estimates of spectra depends on cosmological parameters, both in the linear and non-linear regime. We investigated the scaling of the covariance matrix with parameters from a $w$CDM-cosmology. By constructing a basis for the transformation which relates covariance matrices at different points in parameter space to each other we were able to predict the magnitude and degeneracies rather well, and were able to identify directions in parameter space associated with large changes in the covariance. Our formalism was able to represent variations in the covariance matrix for a wide region of the parameter space. In fact, it could describe variations much larger than that allowed by current experiments like Planck. Furthermore, the formalism also captured degeneracy lines, i.e. parameter combinations along which the covariance matrices effectively remain constant, which we showed to have clear physical explanations. The identification of directions in parameter space in which the largest variations of the covariance matrix occur allows for an economical sampling with numerical simulations; This is feasible because our formalism effectively provides a metric which determines the distance in different directions in parameter space where the variation of the covariance matrix would be larger than a predefined threshold. Apart from predicting variations, our formalism is also well suited for inter- and extrapolation of covariance matrices which are ultimately determined from a large set of numerical simulations at discrete, specifically chosen, parameter points. | 16 | 7 | 1607.03136 |
1607 | 1607.06239_arXiv.txt | We report on \textit{Fermi}/Large Area Telescope observations of the accreting black hole low-mass X-ray binary V404\,Cygni during its outburst in June--July 2015. Detailed analyses reveal a possible excess of $\gamma$-ray emission on 26 June 2015, with a very soft spectrum above $100$\,MeV, at a position consistent with the direction of V404 Cyg (within the $95\%$ confidence region and a chance probability of $4 \times 10^{-4}$). This emission cannot be associated with any previously-known \textit{Fermi} source. Its temporal coincidence with the brightest radio and hard X-ray flare in the lightcurve of V404\,Cyg, at the end of the main active phase of its outburst, strengthens the association with V404\,Cyg. If the $\gamma$-ray emission is associated with V404 Cyg, the simultaneous detection of $511\,$keV annihilation emission by INTEGRAL requires that the high-energy $\gamma$ rays originate away from the corona, possibly in a Blandford-Znajek jet. The data give support to models involving a magnetically-arrested disk where a bright $\gamma$-ray jet can re-form after the occurrence of a major transient ejection seen in the radio. | \begin{figure} \includegraphics[scale=1]{V404Cyg_Lightcurve_LAT511kev_small.pdf} \vspace{-10pt} \caption{$\gamma$-ray flux and $95\%$ upper limits ${>}100\,$MeV and corresponding TS at the position of V404$\,$Cyg in 12-h bins, shifted by $6$ hours in time. The upper panels display the optical (extracted from the AAVSO database), hard X-ray \textit{Swift}/BAT \citep{2013ApJS..209...14K}, INTEGRAL/ISGRI count rate in the $80$--$150$\,keV band and the photon flux in the annihilation line \citep{2016Natur.531..341S} LCs. \label{fig:lc_12h}} \end{figure} The microquasars consist of an accreting black hole or neutron star in a binary system with transient or persistent relativistic jets \citep{1999ARA&A..37..409M}. They display a wide range of behaviour at all wavelengths \citep[e.g][]{2006csxs.book..381F}, but they have rarely been detected at high-energy $\gamma$ rays, despite the high-energy particles produced in their jets \citep{2002Sci...298..196C}. Relativistic particles in the jet could emit $\gamma$ rays, either by Compton up-scattering the low-energy photons from the accretion disc/stellar field \citep[e.g.][]{2002A&A...388L..25G, 2006A&A...447..263B} or by pion decays from inelastic collisions between jet particles and stellar wind protons \citep{2003A&A...410L...1R}. Gamma-ray photons could also be produced at the shock regions where the jets encounter the interstellar medium \citep{2011A&A...528A..89B} or within the jet itself \citep{1999MNRAS.302..253A}. Despite their recurrent outbursts, only the microquasar Cyg X$-$3 \citep[and perhaps Cyg X$-$1,][]{2013ApJ...775...98B, 2013MNRAS.434.2380M} behaves like a clearly identified transient high-energy emitter \citep{2009Sci...326.1512F, 2009Natur.462..620T}. The latter's $\gamma$-ray flares are strongly correlated to the radio emission originating from the relativistic jets \citep{2012MNRAS.421.2947C}. With a donor star of mass ${>}10\,\mbox{M}_{\sun}$, they are consistent with being high-mass X-ray binaries. The low-mass X-ray binary (LMXB) V404$\,$Cygni (also known as GS$\,$2023$+$338) underwent an exceptional outburst phase in June 2015. V404$\,$Cyg is a nearby system \citep[$2.39 \pm 0.14\,$kpc, ][]{2009ApJ...706L.230M} harbouring a ${\sim}9\,\mbox{M}_{\sun}$ black hole and a ${\sim}1\,\mbox{M}_{\sun}$ companion star \citep{2010ApJ...716.1105K, 1992ApJ...401L..97W} with a $6.5$-day orbital period \citep{1992Natur.355..614C}. After a $26$ year long quiescent period \citep[{e.g.}][]{2016ApJ...821..103R}, renewed activity was detected with the \textit{Swift}/BAT \citep{2015GCN..17929...1B} and \textit{Fermi}/GBM \citep[][up to $300\,$keV]{2015GCN..17932...1Y, 2016arXiv160100911J} on June, 15 (MJD 57188). This triggered a worldwide multi-wavelength monitoring campaign from radio \citep[][]{2015ATel.7658....1M} to hard X-rays \citep[][]{2015A&A...581L...9R} until V404$\,$Cyg faded towards its quiescent accretion level in August 2015 \citep{2015ATel.7959....1S}. The source has since undergone a fainter re-brightening, in December 2015 \citep[{e.g.}][]{2015ATel.8455....1B}. We took advantage of the intense activity and monitoring of V404$\,$Cyg to probe its high-energy emission with the Large Area Telescope \citep[LAT,][]{2009ApJ...697.1071A} on board \textit{Fermi}. The $\gamma$-ray analysis leading to the marginal detection of a flare is described in \S\ref{sec:analysis}. Its origin is discussed in \S\ref{sec:discussion}. | \label{sec:discussion} We have found transient excess $\gamma$-ray emission at a position consistent with that of V404\,Cyg, which we could not associate with a previously-known {\it Fermi}/LAT source. Formally, the TS of this excess is not sufficiently high to claim a detection, once the number of trials associated with constructing the LC is taken into account. The probability of measuring a TS (distributed as a $\chi^2_2$) value $\geq 15.3$ over the 10 days of the outburst, given 320, $6$\,h-long, independent trials is ${\sim}2\%$. The integrated PSF of the LAT for the soft excess seen toward V404\,Cyg, in the $0.1$--$2$\,GeV band, has a $1\fdg5$ HWHM. By comparing to $12$-hour-long all-sky maps averaged before and after the time interval of the X-ray outburst, we estimate the chance probability of having a transient or fluctuation as bright as the excess at ${\sim}2\%$. So, the ${\sim}4 \times 10^{-4}$ probability of having a random, point-source, $\gamma$-ray excess at the time of the X-ray outburst, toward V404\,Cyg supports the detection of an associated $\gamma$-ray flare. Its temporal coincidence with the brightest daily averaged \textit{Swift}/BAT flare, associated with a marked change in the multi-wavelength properties of the source, provides a compelling reason to posit a detection of transient $\gamma$-ray emission from V404\,Cyg. \subsection{Multi-wavelength behaviour} \label{sec:mw_picture} V404\,Cyg showed strong flaring activity from radio to hard X-rays over a period of about 10 days following the initial detection of the outburst on MJD 57188. The brightest flares occurred at the end of this period, close in time to the detection of the $\gamma$-ray excess we report. As shown by the AMI $13.9\,$GHz LC (Fig.~\ref{fig:lc_mw}), a giant radio flare with a peak above $3.4\,$Jy occurred ${\sim} 6\,$h before the $\gamma$-ray flare peak. Our $15\,$GHz OVRO data (Fig.~\ref{fig:lc_mw}), as well as VLA and VLBA (J.~Miller-Jones, priv. com.) and RATAN \citep{2015ATel.7716....1T} observations conducted around MJD 57198 and 57200 do not indicate the presence of other major radio flares. At hard X-rays, the brightest flare reaches about $57\,$Crab at $20$--$40\,$keV and occurs ${\sim}14\,$h after the radio flare (i.e.\ starting on MJD 57199.5). This is preceded by another hard X-ray flare around MJD 57199.2 that is coincident with the \textit{Fermi}/LAT $\gamma$-ray excess. Although its $20$--$40\,$keV intensity is lower than the MJD 57199.5 flare, its $80$--$150\,$keV flux is higher (and corresponds to the maximum at ${\sim}29\,$Crab over the whole outburst period, see Fig.~\ref{fig:lc_mw}). The activity on MJD 57198--57200 in radio and hard X-rays is also accompanied by the third (and last) detection of $e^-e^+$ pair annihilation \citep{2016Natur.531..341S}. The annihilation flux LC follows the $100$--$200\,$keV flux over these two days, including the dip at MJD 57199.4. The X-ray spectral fits suggest annihilation occurs in two zones with, respectively, $kT {\sim} 2\,\rm keV$ and ${\sim} 500\,\rm keV$. \citet{2016Natur.531..341S} performed a LAT study during the largest positron flare (MJD 57199.616--57200.261) and derived a $8 \times 10^{-7}\,$ph$\,$cm$^{-2}\,$s$^{-1}$ upper limit that is consistent with our upper limit of $8.7 \times 10^{-7}\,$ph$\,$cm$^{-2}\,$s$^{-1}$ on the same time interval, above $100$\,MeV. The activity and emission levels decrease markedly at all wavelengths after MJD 57200. The $\gamma$-ray excess thus appears related to the last spur of activity of the source, preceding the brightest enhancement of $511$\,keV emission, before it started to fade. \subsection{Gamma-ray emission from the jet?} The association of the $\gamma$-ray excess with major radio activity and marked transition to a softer X-ray state (Fig.~\ref{fig:lc_mw}) are reminiscent of Cyg X$-$3, where $\gamma$-ray emission is detected when the radio flux is above ${\sim} 0.2\,$Jy and when X-ray emission is soft -- but not too much so \citep{2012MNRAS.421.2947C}. The $\gamma$-ray luminosity from V404\,Cyg is ${\sim} 2\times 10^{35} (d/{\rm 2.4\,kpc})^2\,\rm erg\,s^{-1}$, a factor ${\sim} 5$ fainter than the typical ${>}100$\,MeV luminosity from Cyg X$-$3. This is a minor fraction of the overall luminosity output, which reached ${\sim} 2\times 10^{38}\,\rm erg\,s^{-1}$ \citep{2015A&A...581L...9R}. Gamma-ray emission may thus be much fainter in V404\,Cyg than in Cyg X$-$3, either intrinsically or because of the boosting depending on the Lorentz factor and orientation of the emission in the relativistic jet \citep{2005Ap&SS.297..109G}. % Radio VLBI imaging of the jet should be able to narrow down such a possibility. In both cases, it is likely that we have detected only the brightest flaring episode during the outburst. \begin{figure} \includegraphics[scale=1]{V404Cyg_TSmap_0p1-2GeV_spectralcolor_truenames.pdf} \vspace{-10pt} \caption{$12$-h residual TS map on MJD~$57199.25$ ($0.1$--$2\,$GeV, $0\fdg1$\,pixel$^{-1}$). % 3FGL sources are marked as black crosses. $68\%$ and $95\%$ confidence regions on the best fit position are represented. \label{fig:tsmap}} \end{figure} V404\,Cyg appeared to stay in a hard or intermediate state during this outburst \citep{2015A&A...581L...9R, 2016arXiv160100911J, 2016arXiv160103234R, 2016Natur.529...54K}. The tentative high-energy $\gamma$-ray detections of the microquasar Cyg X-1 are also associated with this spectral state \citep[and/or at the transition to the soft state,][]{2013ApJ...766...83S}. \citet{2013ApJ...775...98B} found 21 daily measurements for which $\rm{TS}>9$ at the position of the source, mostly when it was in a hard or intermediate state. \citet{2013MNRAS.434.2380M} obtained a total $\mathrm{TS} {\sim} 15$ by integrating all the hard state data and no detection in the soft state. The $\gamma$-ray emission from Cyg X$-$1 is not a simple extrapolation of the ${>}100\,$keV power-law emission detected in X-rays, associated with non-thermal emission from the corona and/or the jet \citep{2015ApJ...807...17R}. Models suggest that the high-energy $\gamma$-ray emission is located elsewhere, most likely associated with non-thermal emission from the jet \citep{2013MNRAS.434.2380M}. The association in V404\,Cyg of high-energy $\gamma$-ray emission contemporaneous with annihilating $e^+e^-$ pairs further supports this picture. Radiative models of coronal plasmas show significant $511\,$keV line emission when the pair energy input is primarily non-thermal and when the corona is very dense i.e. a large opacity to pair production on ${>}100\rm\,MeV$ photons \citep[e.g.][]{1987MNRAS.227..403S}. The high-energy $\gamma$-ray emission must thus originate away from the corona. An intriguing possibility is that the ejection during the radio flare is associated with the disruption of a jet extracting energy from the black hole through the Blandford-Znajek process \citep{1977MNRAS.179..433B}. Numerical simulations indicate this process to be efficient when magnetic flux dragged in by a thick accretion flow piles up to create a `magnetically-arrested disk' (MAD) close to the black hole \citep{2011MNRAS.418L..79T, 2012MNRAS.423.3083M}. \citet{2016ApJ...819...95O} show that such a configuration leads to high-energy $\gamma$-ray emission in the hard state, due to Compton upscattering of jet or disk synchrotron photons. In contrast, the $\gamma$-ray emission is very weak when the built-up magnetic flux is insufficient to suppress standard accretion. In this scenario, a transient ejection occurs when the MAD configuration is disrupted (reconnects) due to incoming magnetic flux of opposite polarity \citep{2014MNRAS.440.2185D}. \citet{2016ApJ...819...95O} indicate that the re-formed jet is brighter in $\gamma$-rays than prior to the major ejection. Hence, the brief $\gamma$-ray emission from V404 Cyg following a major radio flare might be associated with the temporary re-formation of a powerful Blandford-Znajek jet. The present observations support the presence of $\gamma$-ray emission associated with ejections in microquasars, and for the first time from an accreting black hole with a low-mass donor companion. This emission is weak and variable, with a low statistical significance. Some perseverance will be required to clarify the multi-wavelength context in which it appears in this and other microquasars, thus fulfilling its potential as a diagnostic of the accretion-ejection process. | 16 | 7 | 1607.06239 |
1607 | 1607.01023_arXiv.txt | In this Letter we compare the abundance of member galaxies of a rich, nearby ($z=0.09$) galaxy cluster, Abell 2142, with that of halos of comparable virial mass extracted from sets of state-of-the-art numerical simulations, both collisionless at different resolutions and with the inclusion of baryonic physics in the form of cooling, star formation, and feedback by active galactic nuclei. We also use two semi-analytical models to account for the presence of orphan galaxies. The photometric and spectroscopic information, taken from the Sloan Digital Sky Survey Data Release 12 (SDSS DR12) database, allows us to estimate the stellar velocity dispersion of member galaxies of Abell 2142. This quantity is used as proxy for the total mass of secure cluster members and is properly compared with that of subhalos in simulations. We find that simulated halos have a statistically significant ($\gtrsim 7$ sigma confidence level) smaller amount of massive (circular velocity above $200\,{\rm km\, s^{-1}}$) subhalos, even before accounting for the possible incompleteness of observations. These results corroborate the findings from a recent strong lensing study of the Hubble Frontier Fields galaxy cluster MACS J0416 \citep{grillo2015} and suggest that the observed difference is already present at the level of dark matter (DM) subhalos and is not solved by introducing baryonic physics. A deeper understanding of this discrepancy between observations and simulations will provide valuable insights into the impact of the physical properties of DM particles and the effect of baryons on the formation and evolution of cosmological structures. | \label{sect: intro} In the hierarchical structure formation scenario, galaxies form and evolve in dark matter (DM hereafter) halos that merge with other halos to assemble larger systems. Because of this process, halos are composed of a diffuse matter component and a population of subhalos, whose motion and spatial distribution is determined by dynamical processes taking place after a halo has merged into another one. Dynamical friction makes subhalos sink toward the halo center, where strong tidal fields are very effective at stripping material from the external regions of subhalos \citep[see, e.g.][]{kravtsov2004,BK2008}. Several other processes can act and affect the DM and baryonic components, such as tidal heating, ram-pressure stripping and harassment \citep[e.g.][]{moran2007,biviano08,bruggen2008,delucia2012}. In \cite{grillo2015}, the subhalo distribution inferred from the strong lensing analysis of a massive galaxy cluster, MACS J0416, is compared with the predictions of N-body simulations. This comparison shows a significant lack of massive subhalos in simulations. The latter do not include baryonic physics and this could be one of the reasons for the disagreement with the observed subhalo population. In fact, the simulated subhalos are less concentrated than they would be if they had baryons, therefore they are more fragile against tidal stripping. On the other hand, that cluster has a total density profile characterized by an inner core. This would cause tidal fields to be weaker than in the case of a cuspy profile, such as the Navarro-Frenk-White profile \citep{navarro96,navarro97} found in simulated halos, leading to a larger massive subhalo population in MACS J0416 when compared with simulations. In this paper, we analyze the subhalo distribution of a massive, nearby cluster, by utilizing internal kinematics of cluster galaxies as a proxy of subhalo masses, as opposed to the strong lensing modeling techniques used in \cite{grillo2015}. We then compare the observed subhalo population with the predictions of numerical simulations. In particular, we study Abell 2142 (A2142 hereafter), a massive ($\rm M_{200,cr} = (1.25 \pm 0.13) \times 10^{15}\rm M_\odot$) cluster at $z \sim 0.09$ \citep[][M14 hereafter]{munari14}. The cluster was studied by several authors using different probes, namely X-ray \citep{markevitch00,akamatsu11,rossetti13}, the Sunyaev-Zel'dovich effect \citep{umetsu09}, weak lensing \citep{okabe08} and galaxy dynamics (\citealt[][O11 hereafter]{owers11}; M14), and although it possesses several subclumps in the galaxy distribution (O11), these do not appear to affect the dynamical equilibrium of the cluster significantly (M14). | \label{sect: conclusions} Our analysis indicates that current numerical simulations predict a significant smaller amount of massive (circular velocity above $200\,{\rm km\, s^{-1}}$) subhalos. This result is robust, as it holds even when we compare the predictions of simulations and the direct measurements of velocity values of cluster members, without addressing incompleteness issues. When accounting for the latter, the actual number of observed galaxies becomes larger, making the discrepancy even more significant. These results support the findings of a recent strong lensing study of the Hubble Frontier Fields galaxy cluster MACS J0416 at $z=0.4$ \citep{grillo2015}, suggesting that this discrepancy, which is already present in DM-only simulations, is not alleviated by the inclusion of baryonic physics. | 16 | 7 | 1607.01023 |
1607 | 1607.08334_arXiv.txt | { By revisiting the {\it Suzaku} and {\it XMM-Newton} data of the North Polar Spur, we discovered that the spectra are inconsistent with the traditional model consisting of pure thermal emission and neutral absorption. The most prominent discrepancies are the enhanced \ovii and \neix forbidden-to-resonance ratios, and a high \oviii Ly$\beta$ line relative to other Lyman series. A collisionally ionized absorption model can naturally explain both features, while a charge exchange component can only account for the former. By including the additional ionized absorption, the plasma in the North Polar Spur can be described by a single-phase CIE component with temperature of 0.25 keV, and nitrogen, oxygen, neon, magnesium, and iron abundances of $0.4-0.8$ solar. The abundance pattern of the North Polar Spur is well in line with those of the Galactic halo stars. The high nitrogen-to-oxygen ratio reported in previous studies can be migrated to the large transmission of the \oviii Ly$\alpha$ line. The ionized absorber is characterized by a balance temperature of $0.17-0.20$ keV and a column density of $3-5 \times 10^{19}$ cm$^{-2}$. Based on the derived abundances and absorption, we speculate that the North Polar Spur is a structure in the Galactic halo, so that the emission is mostly absorbed by Galactic ISM in the line of sight. } | The North Polar Spur (NPS hereafter) is a prominent structure emitting both in the soft X-ray and radio bands, with a projected distribution from the Galactic plane at $l \sim 20^{\circ}$ towards the north Galactic pole. Despite its vicinity, the origin of the NPS remains largely unclear. Early research suggested that the NPS is an old supernova remnant, or a front created by stellar wind from the Scorpio-Centaurus OB association, at a distance of several hundred pc from the Sun (Berkhuijsen et al. 1971; Egger \& Aschenbach 1995). Alternatively, the NPS can be explained as a shock front produced by an energetic event, such as starburst, in the Galactic center $\sim 15$ Myr ago (Sofue et al. 1977; Bland-Hawthorn \& Cohen 2003). Recent morphological studies further indicated a possible relation between the NPS and the {\it Fermi} $\gamma-$ray bubbles (e.g., Kataoka et al. 2013). In the ``Galactic center origin'' scenario, the distance to the NPS is expected to be several kpc. X-ray studies of the NPS hot plasma provide important information, including the plasma temperature, density, and metal abundances, which are essential to understand its origin. Using {\it XMM-Newton} observations of three regions in the NPS, Willingale et al. (2003, hereafter W03) identified a thermal component, with a temperature of $\approx 0.26$ keV and metal abundances of $\sim 0.5$ $Z_{\odot}$, associated with the enhanced NPS emission. They further deduced that the thermal energy contained in the NPS is consistent with the energy released by one or more supernovae events. Based on a {\it Suzaku} observation, Miller et al. (2008, hereafter M08) measured a slightly higher thermal temperature $\sim 0.30$ keV, and a quite enhanced nitrogen abundance, with nitrogen-to-oxygen abundance ratio $\approx 4.0$ times of the solar value. They proposed that additional enrichment from stellar evolution in the NPS vicinity is required to explain the observed abundance pattern. The above two X-ray studies are both based on an assumption that the NPS emission is purely from thermal plasma in collisional ionization equilibrium (CIE), affected by only neutral absorption in the line-of-sight. Recent research shows that some diffuse objects also emit non-thermal X-rays, such as charge exchange emission produced at the interface between hot and cold materials (e.g., Lisse et al. 1996, Katsuda et al. 2011, Gu et al. 2015). There are also cases in which a portion of the foreground absorber appears to be highly ionized, as reported in e.g., Yao \& Wang (2005) and Hagihara et al. (2010, 2011). Both the charge exchange component and ionized absorption can strongly affect the line emission, and hence deviate the resulting physical model. Indeed, the NPS is a potential target for charge exchange, because it might be surrounded by a shell of neutral gas (e.g., Heiles et al. 1980), providing substantial environment for ion-neutral charge exchange. Meanwhile, the NPS might be subject to ionized absorption, since in the ``Galactic center origin'' scenario, it is expected to be located behind a layer of hot interstellar medium (ISM) of the Galactic halo/bulge. Actually, the spectra presented in W03 and M08 already showed a hint for such additional components: the central line energies of unresolved \ovii and \neix triplets are shifted by a few 10 eV to longer wavelength, relative to the energies of \oviii and \nex (Lallement 2009). This might indicate either charge exchange enhancement of the forbidden lines, or possible absorption in the resonance ones. In this paper, we present a detailed spectral analysis by revisiting the high quality X-ray data of the NPS. The latest {\tt SPEX} version 3.01 is employed, as it includes a new charge exchange code (Gu et al. 2016) and a model for ionized absorption. \S2 gives a brief description of data reduction, and the data analysis and results are presented in \S3. We discuss the physical implication of the results in \S4 and summarize our work in \S5. Throughout the paper, the errors are given at 68\% confidence level. We adopt the proto-Solar abundance table of Lodders et al. (2009), and convert the previous abundance measurements to the new standard. | \begin{figure*}[!] \centering \resizebox{0.8\hsize}{!}{\hspace{-1cm}\includegraphics[angle=0]{merge_int_abu.eps}} \caption{(a) Integrated Earth-centered ISM column density profiles based on the models reported in Almy et al. (2000, black) and Miller \& Bregman (2013, red). Extra column density from the local hot bubble is included in the central 70 pc. The dashed line shows the observed value based on {\tt CIE} model with ionized absorption. (b) The [N/O] verseus [O/H] diagram. The best-fit NPS results obtained with {\tt CIE} and {\tt CIE} plus {\tt CX} models with the new absorptions are shown in black and grey error bars, respectively. The M08 result is plotted in green. The abundance patterns of the cold ISM based on local star and distant star absorption measurements by Knauth et al. (2006) are plotted in blue filled circles and blue empty circles, respectively. The metal-poor and metal-rich Galactic halo star data from Israelian et al. (2004) are marked by red filled diamonds and red empty diamonds, respectively. \vspace{0.5cm} } \label{CIE_fig} \end{figure*} By revisiting the {\it Suzaku} and {\it XMM-Newton} data of the NPS region, we discovered that the soft X-ray spectra can be well described by a single-phase thermal component, with a temperature of $0.23-0.25$ keV, absorbed by at least two species of foreground materials, in both neutral and ionized states. The key evidence for the ionized absorber is the unusually high \oviii Ly$\beta$ line relative to other Lyman series. Assuming a nearby turbulent-free plasma, the hot absorber exhibits a balance temperature of $0.17-0.20$ keV and column density of $\sim 3-5 \times 10^{19}$ cm$^{-2}$. A charge exchange component is marginally detected only with the {\it Suzaku} XIS data. The oxygen abundance of the NPS is then obtained to be $0.6-0.7$ $Z_{\odot}$, apparently higher than those reported in W03 and M08. The Fe/O ratio is consistent with the solar values within measurement uncertainties, while the N/O becomes slightly sub-solar. Next we shed light on the origin of the ionized absorber based on the derived properties. As shown in \S 3.2, the balance temperature of the absorber is $0.17-0.20$ keV, lower than the NPS plasma temperature $\approx 0.23-0.25$ keV on $> 90$\% confidence level. This means that it cannot be fully ascribed to the self-absorption of the NPS plasma. On the other hand, the local hot bubble alone cannot be the absorber either. The temperature of the local hot bubble ($\approx 0.1$ keV) is lower than the observed value, and assuming a line-of-sight scale of $40-90$ pc and density of $0.01$ cm$^{-3}$ (e.g., W03), the column density is estimated to be $1.2-2.7 \times 10^{18}$ cm$^{-2}$, accounting $<10$\% of the obtained value (\S3.2). The obtained properties of the ionized absorber are consistent with those reported in the absorption studies on Galactic compact object 4U 1820$-$303 (Hagihara et al. 2011), and extragalactic objects PKS 2155$-$304 (Hagihara et al. 2010), LMC X$-$3 (Yao et al. 2009), and Mrk 509 (Pinto et al. 2012). This leads us to the scenario that the Galactic ISM contributes significantly to the observed ionized absorption. The temperature $\approx 0.2$ keV appears to be self-consistent with that of the Galactic ISM included as a background component in \S2.2. It also agrees well with previous measurements of the ISM temperature in the Galactic halo (e.g, Smith et al. 2007) and Galactic bulge (e.g., Almy et al. 2000). To calculate the ISM absorption column density, we employ the 3-D ISM density models from Almy et al. (2000, in their Fig. 5) and Miller \& Bregman (2013, ``spherical-saturated'' model in their Table 2). The former is based on {\it ROSAT} 3/4 keV observation of the emission in the Galactic bulge region, while the latter is focused on the Galactic halo, and utilized line absorption measurement on background objects. The two models provide galactocentric ISM density profiles, which are then transformed into Earth-centered line-of-sight distance profiles by using Eqs.~1$-$3 of Miller \& Bregman (2013). The sky coordinate of the {\it Suzaku} pointing is used in the center transformation. By integrating the density over distance, we calculate the column density distance profiles and present them in Figure 5a. It shows that the two ISM models are roughly consistent with each other. Let us consider an extreme case, in which the ionized absorption is fully due to the Galactic ISM, the models predict that the part of NPS covered by the XIS would be at a distance of $\sim 6-7$ kpc. This is still well in line with the recent measurements of the NPS distance with radio data by Sun et al. (2014) and Sofue (2015). As described in \S3.2, the over-solar N/O abundance ratio reported in M08 can be migrated to the large opacity of \oviii Ly$\alpha$ line. In Figure 5b, the new results are plotted in a [N/O] verseus [O/H] diagram, and are compared with the abundances of Galactic halo stars measured in Israelian et al. (2004). The NPS values are consistent with the implied Galactic stellar evolution by lying in the gap between the ``metal-poor'' and ``metal-rich'' subsamples of stars. At the same time, the abundance patterns of Galactic cold ISM, based on {\it HST} and {\it FUSE} observations of \oi and \nni absorptions against stars (Knauth et al. 2006), are also plotted in the same diagram of Figure 5b. Despite of the large uncertainties, the NPS results appear to agree better with the abundance patterns of the distant ISM ($d > 500$ pc), than with those of the local ISM ($d < 500$ pc). This also supports the scenario that the NPS is a structure in the Galactic halo rather than in the solar neighborhood. | 16 | 7 | 1607.08334 |
1607 | 1607.06329_arXiv.txt | In this Letter, we study analytically the evolutions of the flat Friedmann-Lemaitre-Robertson-Walker (FLRW) universe and its linear perturbations in the framework of {\em the dressed metric approach} in loop quantum cosmology (LQC). Assuming that the evolution of the background is dominated by the kinetic energy of the inflaton at the quantum bounce, we find that both evolutions of the background and its perturbations are independent of the inflationary potentials during the pre-inflationary phase. During this period the effective potentials of the perturbations can be well approximated by a P\"oschl-Teller (PT) potential, from which we find analytically the mode functions and then calculate the corresponding Bogoliubov coefficients at the onset of the slow-roll inflation, valid for any inflationary model with a single scalar field. Imposing the Bunch-Davies (BD) vacuum in the contracting phase prior to the bounce when the modes are all inside the Hubble horizon, we show that particles are generically created due to the pre-inflation dynamics. Matching them to those obtained in the slow-roll inflationary phase, we investigate the effects of the pre-inflation dynamics on the scalar and tensor power spectra and find features that can be tested by current and forthcoming observations. In particular, to be consistent with the Planck 2015 data, we find that the universe must have expanded at least $141$ e-folds since the bounce. | The paradigm of cosmic inflation has achieved remarkable successes in solving several problems of the standard big bang cosmology and predicting the primordial perturbation spectra whose evolutions explain both the formation of the large scale structure of the universe and the small inhomogeneities in the cosmic microwave background (CMB) \cite{inflation}. Now they are matched to observations with unprecedented precisions \cite{WMAP, Planck2013, Planck2015}. However, such successes are contingent on the understanding of physics in much earlier epochs when energies were about the Planck scale. This leads to several conceptual issues. For example, to be consistent with observations, the universe must have expanded at least $60$ e-folds during its inflationary phase. {However, if the universe had expanded a little bit more than $70$ e-folds during inflation (as it is the case in a large class of inflationary models \cite{MRV}), then one can show that the wavelengths of all fluctuation modes which are currently inside the Hubble radius were smaller than the Planck length at the beginning of the period of inflation. This was referred to as the trans-Planckian issue in \cite{trans-planck}, and leads to the question about the validity of the assumption: {\em the matter fields are quantum in nature but the spacetime is still classical}, which are used at the beginning of inflation in order to make predictions \cite{inflation}. In addition, insisting on the use of general relativity (GR) to describe the inflationary process will inevitably lead to an initial singularity \cite{singularity}. Moreover, the inflation paradigm usually sets the BD vacuum state at the time when the wavelength of fluctuations were well within the Hubble horizon during the inflationary process. However, such treatment ignores the pre-inflationary dynamics which could lead to non-BD states at the onset of inflation, even when these modes were well inside the Hubble horizon during inflation. For more detail about the sensibility of the inflationary paradigm to Planckian physics, we refer the readers to \cite{trans-planck,DB}. } All the issues mentioned above are closely related to the fact that we are working in the regime where GR is known to break down. One believes that new physics in this regime - a quantum theory of gravity, will provide a complete description of inflation as well as its pre-inflationary dynamics. LQC is one of such theories that offers a framework to address these issues, in which the inflationary scenarios can be extended from the onset of the slow-roll inflation back to the Planck scale in a self-consistent way \cite{planck_extension,planck_extension_CQG,quadratic_loop}. Remarkably, the quantum geometry effects of LQC at the Planck scale provide a natural resolution of the big bang singularity (see \cite{bounce,Ashtekar2015CQG,BB16, Yang_alternative_2009} and references therein). In such a picture, the singularity is replaced by a quantum bounce, and the universe that starts at the bounce can eventually evolve to the desired slow-roll inflation \cite{AS10,bounce_inflation, bounce_inflation2, deformed_tensor, deformed_scalar,deformed,Starobinsky_loop,bounce_effects}. An important question now is whether the quantum bounce can leave any observational signatures to current/forth-coming observations, so LQC can be placed directly under experimental tests. The answer to this question is affirmative. { In fact, with some (reasonable) assumptions and choice of the initial conditions, the {\em deformed algebra approach} already leads to inconsistence with current observations \cite{deformed}.} Note that in general there are two main approaches to implement cosmological perturbations in the framework of LQC, the {\em dressed metric} and {\em deformed algebra approaches} \cite{bounce,Ashtekar2015CQG,BB16}. In both, the primordial perturbations have been intensively studied {\em numerically} \cite{planck_extension_CQG,quadratic_loop, bounce_effects, deformed, deformed_scalar, deformed_tensor, Starobinsky_loop}. One of our purposes of this Letter, in contrast to the previous numerical studies, is to present an {\em analytical} analysis of the effects of the quantum bounce and pre-inflation dynamics on the evolutions of both background and spectra of the scalar and tensor perturbations, in the framework of the dressed metric approach \citep{planck_extension,planck_extension_CQG,quadratic_loop}. It is expected that such an analysis will provide a more complete understanding of the problem and deeper insights. In the following, we will focus on the case that the kinetic energy of the inflaton dominates the evolutions at the bounce, because a potential dominated bounce is either not able to produce the desired slow-roll inflation \cite{Starobinsky_loop}, or leads to a large amount of e-folds of expansion. { This will wash out all the observational information about the pre-inflation dynamics and the resulting perturbations are the same as those given in GR \citep{bounce, Ashtekar2015CQG,BB16}. Assuming that the influence of the potential at the bounce is negligible, } our studies show that: { \begin{itemize} \item { During the pre-inflationary phase, the evolutions of the background and the scalar and tensor perturbations are independent of the inflationary potentials. Thus, the evolution of the background is the same for any chosen potential, and in this sense we say that it is {\em universal}. } \item { During this phase the potentials of the scalar and tensor perturbations can be well approximated by an effective PT potential, for which analytic solutions of the mode functions can be found. The Bogoliubov coefficients at the onset of the slow-roll inflation can thereby be calculated [cf. (\ref{main})], which are valid for any slow-roll inflationary model with a single scalar field. Assuming that the universe is in the BD vacuum in the contracting phase (the moments where $t \lesssim - t_s$ as shown in Fig. \ref{length}) we find that particle creations occur generically during the pre-inflation phase. } \item Oscillations always happen in the power spectra, and their phases for both scalar and tensor perturbations are the same, in contrast to other theories of quantum gravity \cite{trans-planck,Zhu1}. \item Fitting the power spectra to the Planck 2015 data \cite{Planck2015}, we find the lower bound for $N_{\text{tot}} \equiv \ln{(a_0/a_B)} >141$ (95\% C.L.), where $a_B$ and $a_0$ denote the expansion factor at the bounce and current time, respectively. Details of the calculations will be reported elsewhere \cite{bounce_uniform}. \end{itemize} | In this Letter, we {\em analytically} studied the evolutions of the background and the linear scalar and tensor perturbations of the FLRW universe in LQC within the framework of {\em the dressed metric approach} \citep{planck_extension,planck_extension_CQG,quadratic_loop}, and showed that, { if the pre-inflationary phase is dominated by the kinematic energy of the inflaton, the evolutions will be {\em independent of the slow-roll inflationary models} during this phase [cf. Fig.~\ref{scalar_factor} and Eqs.~(\ref{Bsolution}) and (\ref{pw})]. Imposing the BD vacuum in the contracting phase ($t \lesssim -t_s$ as shown in Fig. 2), } we obtained the Bogoliubov coefficients (\ref{main}) at the onset of the slow-roll inflation, which shows clearly that during the pre-inflationary phase, particles are generically created ($\left.\beta_k\right|_{t=t_i} \not= 0$), and the resulting power spectra are $k$-dependent. This is in contrast to GR (where the BD vacuum ($\left.\beta_k\right|_{t=t_i} = 0$) is usually imposed at the onset of the slow-roll inflation \cite{inflation}. {This provides a potential window to test LQC directly by the measurements of CMB and galaxy surveys \citep{S4-CMB}. In particular,} fitting the power spectra to the Planck 2015 temperature (TT+lowP) and polarization (TT,TE,EE+lowP) data, we found the lower bound for $N_{\text{tot}} \equiv \ln{(a_0/a_B)} >141$ (95\% C.L.). That is, to be consistent with current observations of CMB, the universe must have expanded at least 132 e-folds since the bounce. | 16 | 7 | 1607.06329 |
1607 | 1607.06059_arXiv.txt | In this work, we are trying to extent the existing photometric redshift regression models from modeling pure photometric data back to the spectra themselves. To that end, we developed a PCA that is capable of describing the input uncertainty (including missing values) in a dimensionality reduction framework. With this ``spectrum generator'' at hand, we are capable of treating the redshift regression problem in a fully Bayesian framework, returning a posterior distribution over the redshift. This approach allows therefore to approach the multimodal regression problem in an adequate fashion. In addition, input uncertainty on the magnitudes can be included quite naturally and lastly, the proposed algorithm allows in principle to make predictions outside the training values which makes it a fascinating opportunity for the detection of high-redshifted quasars. | The exploration of the past development of the universe has been mainly driven by the detection and investigation of highly-redshifted extragalactic sources, such as the quasi-stellar objects (QSO, \citealp{1993ARA&A..31..473A}). The study of the distribution of these objects over space and time allows do draw precise conclusions about how the universe has initially formed and developed since then \citep{2010ASSP...14..255A}. Additonally, photometric redshifts have been used in the studies of galaxy clusters \citep{2008MNRAS.387..969A} and in constraining the galaxy luminosity function \citep{1996AJ....112..929S}. Since the detection of the first quasars a significant time of research has been spent in estimating the redshift, caused by the expansion of the universe, to these ultra-luminous objects. While spectroscopic surveys are extremely precise in doing so, they are extremely time-intensive and can not be used to study a large fraction of the objects known to date. Instead photometric surveys are used to infer knowledge about the nature and the redshift of the quasars. Originally, this was done in a template-based way \citep{2000A&A...363..476B} and only rather recently the number of data-driven approaches has increased drastically (\citealp{2010MNRAS.406.1583W, 2011MNRAS.413.1395O, 2011MNRAS.418.2165L} and many more). In these works the main focus has been on the comparison of methodology instead of the introduction of new concepts and the community seemed to have agreed on, that the random forest is tailored for this task in terms of reproducability, precision and computational complexity. In our work, we want to present an algorithm that considers a number of problems in redshift regression that have been known to the community for a long time but have not been tackled and/or been ignored over the last decade. As it can be study in all plots showing the regressed redshift versus the actual redshift: the redshift regression problem is actually \emph{multimodal}. This means that a given color can generally be explained by \emph{more than one} redshift, cf. Fig.\,\ref{fig:multimodal}. In our work, we will also show that the RMS is an inadequate measure to estimate the accuracy of photometric redshift regression algorithms and present a more useful measure. This measure will be based on the posterior distribution over the redshift which should be rather considered in a multimodal problem. Another striking problem of the existing methodology is, that the uncertainty of the input data can not be considered so far, i.e. that uncertainties, or even more drastic missing values, of the colors can not be considered. Despite that, it has been claimed, that these regression algorithms can be used for predicting out-of-sample (i.e., higher redshifts than provided by the training) regression values. While in practice this might be possible, it is conceptually highly questionable whether this is the right concept. Instead we can provide with the model-based approach an alternative in the search for highly-redshifted extra-galactic objects, however, a lot more work has to be done in order to achieve this goal. \begin{figure} \centering \includegraphics[width=.5\textwidth]{multimodal.pdf} \caption{Example of a multimodal problem. It is evident that the Gaussian process is not capable of describing the multimodal data accordingly.\label{fig:multimodal}} \end{figure} We start with a description of the data used in \mbox{Section\,\ref{sec:data}}, followed by a description of the methodology in \mbox{Section\,\ref{sec:methodology}}. In the following part the results are presented which contains also a direct comparison with the random forest. Subsequently, we summarize in \mbox{Section\,\ref{sec:discussion}} the work and show potential prospects for continuing the presented work. | \label{sec:discussion} In this work, we presented an alternative way of modeling photometric redshifts. The idea was to combine the advantages of model-based approaches (full model control, Bayesian framework, multimodality) with the ones from the data-driven approaches (no explicit model-formulation). For this reason a probabilistic PCA was developed which can produce projections in the presence of missing and uncertain data. In a first step, we tried to utilize this projection to optimize for the coordinates of the PCA, but learned that if the coordinates are unconstrained already a 10-dimensional PCA can produce unphysical spectra which then lead to overfitting. Instead, we used the provided reconstructed spectra themselves as prototypes and therefore converted the continuous model into a discrete one. With these prototypes at hand, we were able to treat the regression problem in a fully Bayesian framework which allowed us to compute the posterior distribution over the redshift. We compared the distribution of our model with predictions made by the random forest and the Gaussian process. Even though, we emphasize that this comparison is fairly meaningless as the different approaches are tailored for different tasks. The standard regression models (RF, GP) are capable of describing a function-like behaviour (one input maps to a redshift), while it is long known that the photometric redshift regression problem is a multimodal one (one input can map to several redshifts). Therefore, also the measure of the $RMS$ and the $MAD$, as commonly done in the literature, are questionable measures of the quality of a regression algorithm. For this reason, we introduced two additional measures which reflect the likelihood of capturing the true underlying redshift and on the other hand a measure denoting how well we can describe similar colors with different redshifts. By design, our presented algorithm is capable of describing this behaviour, while the standard regression models are not, as they expect a one-to-one (or many-to-one) mapping. We would like to highlight, that the choice of the methodology is only of minor importance, as long as the wrong objective is measured. To treat this multimodality purely with data-driven methodology is quite complicated, even though, some approaches already exist (e.g., mixture density networks). Consequently, past publications have utilized a measured that simply not adequate for the task at hand. \subsection*{Prospects} The presented approach can be developed further in two aspects: the methodological one and the astronomical one. The biggest disadvantage of the current design is, that rather prototypes are extracted and used for prediction than a continuous PCA model. The problem of the continuous PCA model is its high flexibility which leads to severe overfitting. For this reason, it would be desirable to constrain the PCA in the lower dimensional space by, for example, estimating the density of the spectra. This density estimate could then be used as a prior weighting the frequency of the a PCA coordinate in the real spectra. In case that this density can be written as a mixture of Gaussians, it might be even possible to marginalize the posterior over the PCA coordinates analytically. This would make the model extremely fast and would omit the point of sampling or drawing prototypes. From an astronomical point of view, more work has to be done. So far we have just provided the functioning on a small toy dataset where the magnitudes were extracted with the provided filter curves and (known and noise-free) zero points were added. We chose this setting, as we wanted to have full control on the model and not to be distracted by errorneous and noisy calibrations. It is important to notice that a purely data-driven approach can deal with this quite naturally, while the presented algorithm depends heavily on the correctness of these calibrations. On the other hand, it is of course also possible to include a given uncertainty of the zero points into the model and also this can be cross-validated using a training set. In summary, a much more detailed understanding of how photometric measurements relate to the spectra is required. Another striking advantage of our model is, that we can include uncertainty of the photometric measurements directly into the model. This includes also \emph{missing} values which are a common struggle in astronomy due to the different coverage and depth of the surveys. This is contrast to nearly all data-driven algorithms which can at maximum handle the input uncertainty by sampling (which is hard for a missing value). A different prospect of our model would be to extent the PCA model towards the infrared. At the moment, our coverage above $1\mu m$ is very shallow and thus it would be desirable to retrieve near-infrared to mid-infrared spectra of low-redshifted quasars (as otherwise the rest-frame would be in the optical again). This would allow us to include also infrared data as then the coverage of the prototypes would reach into the near infrared. It is important to notice that it does not matter whether the infrared spectra are the same objects as the optical ones, it only has to be guaranteed that there is considerable overlap with the prototypes as they are now. \label{lastpage} | 16 | 7 | 1607.06059 |
1607 | 1607.01954_arXiv.txt | MICADO will equip the E-ELT with a first light capability for diffraction limited imaging at near-infrared wavelengths. The instrument's observing modes focus on various flavours of imaging, including astrometric, high contrast, and time resolved. There is also a single object spectroscopic mode optimised for wavelength coverage at moderately high resolution. This contribution provides an overview of the key functionality of the instrument, outlining the scientific rationale for its observing modes. The interface between MICADO and the adaptive optics system MAORY that feeds it is summarised. The design of the instrument is discussed, focussing on the optics and mechanisms inside the cryostat, together with a brief overview of the other key sub-systems. | \label{sec:intro} % MICADO is the Multi-AO Imaging Camera for Deep Observations, that is being designed and built by a consortium of institutes in Germany, France, Netherlands, Austria, and Italy. It will equip the European Extremely Large Telescope (E-ELT) with a first light capability for diffraction limited imaging at near-infrared wavelengths (0.8--2.4\,$\mu$m). The instrument is optimised to work with the laser guide star multi-conjugate adaptive optics (MCAO) module developed by the MAORY consortium\cite{dio16}. It also includes a jointly developed single-conjugate adaptive optics (SCAO) mode\cite{cle16} that uses just a single natural guide star. MICADO has to be able to function in a 'stand-alone' deployment configuration that involves operating the instrument with the SCAO wavefront sensor but without the full MCAO bench. It is required for the acceptance, integration and verification (AIV) phase of the MICADO project, and is also part of the programmatic risk mitigation strategy for the E-ELT. Following the approval of the E-ELT construction proposal in December~2014, the Agreement to design and build MICADO was signed in September~2015, and the Phase~B Kick-Off meeting was held shortly afterwards. The schedule for MICADO is tied to that of the E-ELT, for which first light is planned to be in 2024. In order to achieve this, and yet also to accommodate development of the interfaces to the observatory (and also to MAORY) that is expected to occur during the early stages of the project, the preliminary design phase (Phase~B) of MICADO has been scheduled to last for 3 years, with the associated design review expect to be held towards the endof 2018. Mid-way through that phase, there will be an internal review of the system and sub-system specifications and interfaces. While the design has evolved considerably since the end of Phase~A\cite{dav10} in 2009, the rationale for the instrument remains fundamentally the same. The key capabilities of MICADO exploit the most unique features of the E-ELT: sensitivity and resolution, precision astrometry, and wide wavelength coverage spectroscopy. The primary observing mode is imaging, with a focus on astrometry\cite{tri10,fri10}. To achieve the stability necessary to provide spatial resolution better than 10\,mas and astrometric precision below 50\,$\mu$as, the instrument is supported above the Nasmyth platform in a gravity invariant orientation, includes an optical path comprising entirely of fixed mirrors, uses a state-of-the-art atmospheric dispersion corrector, and has a dedicated astrometric calibration plan and data pipeline. The array of $3\times3$ 4k$^2$ detectors at the focal plane can image a small field of about 20$^{\prime\prime}$ with a fine 1.5\,mas pixel sampling that is especially useful in very crowded fields or at short wavelengths; as well as a large field more than 50$^{\prime\prime}$ across, with a coarser 4\,mas pixel scale that still fully samples the H- and K-band diffraction limit. In both cases, a wide selection of broad and narrow band filters are available. This mode will provide comparable sensitivity to the James Webb Space Telescope at 6 times better spatial resolution, and enable proper motions as small as 5\,km\,s$^{-1}$ to be measured at distances of up to 100\,kpc. While exoplanets are among the primary science drivers for the E-ELT, a dedicated camera for such studies will not be available for a number of years after first light. As such, a secondary role of MICADO is to bridge this gap, using coronagraphy to providing a high contrast imaging capability. This is enabled via a variety of focal plane and/or pupil plane masks, and is envisaged to make use of angular differential imaging techniques. Its novel feature will be the very small angular scales on which it will be possible to detect exoplanets. A slit spectroscopic mode is also included in MICADO, optimised for compact objects. It emphasizes simultaneous wavelength coverage at moderately high resolution: covering IJ and HK bands together at a spectral resolution of $R\sim15000$ for point sources (the resolution will be lower for extended objects that fill the width of the slit). Slit widths suitable for compact and extended objects will be provided. Time resolved imaging is enabled by defining suitable windows on the detector which enable the frame rate to be in the range 1--100\,Hz, and providing precise ($<$1\,ms time) stamping of the frames. | \label{sec:conc} MICADO is the first light imaging camera for the E-ELT. The consortium designing, building, and commissioning the instrument comprises institutes in Germany, France, the Netherlands, Austria, and Italy. The project entered Phase~B in September~2015 and the next milestone is consolidation of the basic design concept together with the system and sub-system specifications and interfaces, followed in 2018 by the preliminary design review. The instrument is optimised to operate with the multi-conjugate laser guide star adaptive optics system MAORY. The two consortia are also jointly developing a simple and robust single-conjugate natural guide star adaptive optics system. The observing modes offered by MICADO include: standard imaging, astrometric imaging, high contrast imaging, time resolved imaging, and slit spectroscopy. Observers will be assisted with dedicated preparation tools, data processing pipelines, and PSF reconstruction. | 16 | 7 | 1607.01954 |
1607 | 1607.06868_arXiv.txt | The origin of the observed steep rotation curves of blue compact dwarf galaxies (BCDs) remains largely unexplained by theoretical models of BCD formation. We therefore investigate the rotation curves in BCDs formed from mergers between gas-rich dwarf irregular galaxies based on the results of numerical simulations for BCD formation. The principal results are as follows. The dark matter of merging dwarf irregulars undergoes a central concentration so that the central density can become up to six times higher than those of the initial dwarf irregulars. However, the more compact dark matter halo alone can not reproduce the gradient differences observed between dwarf irregulars and BCDs. We provide further support that the central concentration of gas due to rapid gas-transfer to the central regions of dwarf--dwarf mergers is responsible for the observed difference in rotation curve gradients. The BCDs with central gas concentration formed from merging can thus show steeply rising rotation curves in their central regions. Such gas concentration is also responsible for central starbursts of BCDs and the high central surface brightness and is consistent with previous BCD studies. We discuss the relationship between rotational velocity gradient and surface brightness, the dependence of BCD rotation curves on star formation threshold density, progenitor initial profile, interaction type and merger mass ratio, as well as potential evolutionary links between dwarf irregulars, BCDs and compact dwarf irregulars. | Blue compact dwarf galaxies (BCDs) resemble the spectra of H{\sc ii} regions of spiral galaxies while being optically small, low luminosity (Thuan \& Martin 1981), and typically exhibit subsolar metallicities [Fe/H] $< -1$ (Drozdovsky \& Tikhonov 2000; Drozdovsky et al. 2001). Many studies have turned their focus to BCDs as they exhibit features similar to those of high-redshift galaxies (Papaderos 2006) and were originally thought to have been undergoing their first starbursts (e.g. Searle, Sargent \& Bagnuolo 1973; Aloisi, Tosi \& Greggio 1999; Thuan, Izotov \& Foltz 1999), making them potential insights into the early Universe. The more recent confirmation of an underlying population of old stars in most BCDs (e.g. James 1994; Cair{\'o}s et al. 2001; Amor\'in et al. 2007) did not compromise their potential as valuable low-redshift laboratories to study star formation and galaxies analogous to the early Universe due to their low metallicities (Thuan et al. 1995; Izotov \& Thuan 1999). The rotation curves of dwarf galaxies generally resemble solid body rotation in their central regions, but when compared to their disc scale length they rise as steeply as spirals (Swaters et al. 2009). Comparatively, BCDs exhibit central rotation curves significantly steeper than the solid body rotation profile implying a strong central concentration of mass (Lelli et al. 2012a,b; Lelli et al. 2014a, hereafter L14; Koleva et al. 2014, hereafter K14). Dark matter of BCDs is observed to have high central densities around 0.1 M$_{ \odot} $ pc$^{-3}$ (Meurer et al. 1996; Meurer, Stavley-Smith \& Killeen 1998, hereafter M98) and centrally peaking H{\sc i} maps (e.g. van Zee et al. 1998; Lelli et al. 2014b) have frequently been related to their elevated central starburst regions and small optical sizes (e.g. McQuinn et al. 2015). One of the unresolved problems in BCDs is their steep central rotational velocity profiles (e.g. M98; Lelli et al. 2012a,b; L14). Mass models of the BCDs NGC 2915 and NGC 1705 show rotation curves dominated by dark matter with central dark matter densities to up to 10 times higher than those of dwarf irregulars (Meurer et al. 1996; M98; Elson, de Blo \& Kraan-Korteweg 2010) but the definite cause of this concentration remains undecided. With it becoming more obvious that the concentration of not only dark matter is necessary for BCD formation but the concentration of gas also, several mechanisms have already been suggested. Torques from internal gas and star clumps (Elmegreen et al. 2012) and rotating triaxial dark matter haloes (Bekki \& Freeman 2002) are among the many proposed. Verbeke et al. (2014, hereafter V14) showed that in-falling gas clumps consistently increase central concentration of gas consequentially causing a steeper rotation curve, ignites starbursts extending over a few 100 Myr and increases central surface brightness in simulated dwarf galaxies. Another mechanism that has been gaining rapid support recently is dwarf--dwarf galaxy mergers (e.g. Noeske et al. 2001; {\"O}stlin et al. 2001; Bekki 2008, hereafter B08). Recent observational studies frequently find that gas distributions in BCDs resemble that of merger remnants (e.g. Ekta et al. 2008), the majority of BCDs show signs of tidal interaction or galaxy merging in their starburst component (e.g. Pustilnik et al. 2001) and Pak et al. 2016 observed an apparent in-progress merger with two blue star forming cores and exponential surface brightness. However, even though dwarf galaxy mergers efficiently transport large amounts of gas to the centre of the interaction (Bekki 2015, hereafter B15), these models have not yet discussed whether BCDs formed from dwarf--dwarf merging explain the observed steep rotation curves of BCDs (L14). The purpose of this paper is to investigate the rotational velocity profile of BCDs produced through merger interactions based on a set of new hydrodynamical simulations with star formation, chemical evolution, and dust evolution. A primary focus of this paper is on the origin and evolution of steep rotation curves of BCDs. We mainly investigate 20 models of BCDs formed from dwarf--dwarf merging with different star formation threshold densities, merger mass ratios, and merger orbits. The rotation curves of galaxies are a useful tool for investigating mass distributions of different components (e.g. stellar disc and dark matter halo) in galaxies and distinguishing between contributions from each of these to the mass distributions. Therefore, the present numerical study contributes greatly to the better understanding of the origin of the observed mass profiles of BCDs. The present study is complementary to previous simulations of BCD formation and evolution through dwarf--dwarf galaxy mergers (e.g. B08; B15) because the major focuses of these previous studies were mostly on structures, metallicities, and post-merger evolution of BCDs. The plan of the paper is as follows. First, we describe the details of the adopted dwarf--dwarf merger model of BCD formation. Then, we present some key results of the new simulations on the rotation curves of BCDs formed from dwarf--dwarf mergers and their dependence on the model parameters of dwarf--dwarf merging. In this section, we also discuss how the steep rotation curves of BCDs can be achieved during dwarf--dwarf merging. In \S 4, we discuss the importance of major merging, star formation threshold gas density, and initial dark matter mass profiles in the formation of the steep rotation curves of BCDs. We summarize our conclusions in \S 5. | We have investigated the rotation curves of BCDs formed from mergers between gas-rich dwarf irregular galaxies through numerical simulations. Using 20 different dwarf--dwarf merger models of BCD formation, we have investigated the correlation between rotation curve gradient and surface brightness, dark matter and gas concentration, the dependence of rotation curve on star formation threshold and merger mass ratio. The principal results are as follows: (1) The slowly rising rotation curves of initial dwarf irregulars can be transformed into steeply rising ones in the formation of BCDs. The rotation curve gradients of BCDs with strong central starbursts can be by 20 km s$^{-1}$ steeper than those of initial dwarfs. This result implies that dissipative dwarf--dwarf merging can be responsible for the observed steep rotation curves of BCDs. (2) Although the rotation curves in the outer parts ($>2$ kpc) of BCDs are dominated by dark matter, both baryonic and dark matter contribute equally to the rotation curves in the central regions ($R<1$ kpc). Dissipative merging process can enhance the central mass concentration of dark matter by a factor of 6, and consequently transform dark matter haloes with initially cored distributions into cuspy ones that are similar to the NFW profile. (3) Dark matter provides the base rotation curve but the increase in gradient between the progenitor and the BCD is due to the strong central concentration of gas. This result suggests that baryonic physics is a key for better understanding the origin of the rotation curves of BCDs. (4) BCD rotation curves are dependent on star formation threshold density and merger mass ratio. BCDs formed from major dwarf--dwarf merging can show steeper rotation curves that minor mergers. Simple tidal interaction without merging cannot create BCDs with steep rotation curves, minor mergers are more likely to create BCDs with off centre starbursts. (5) BCDs formed from dwarf--dwarf merging are likely to evolve into compact dwarf irregulars with a significant amount of residual star formation after strong starburst phases, because there is still a plenty of gas left in the merger remnants. The BCDs inevitably show both high surface brightness and steep rotation curves, because efficient gas infall into the central regions of merging can drive both central mass concentration and enhanced star formation. | 16 | 7 | 1607.06868 |
1607 | 1607.02350_arXiv.txt | The origin of Phobos and Deimos is still an open question. Currently, none of the three proposed scenarios for their origin (intact capture of two distinct outer solar system small bodies, co-accretion with Mars, and accretion within an impact-generated disk) is able to reconcile their orbital and physical properties. Here, we investigate the expected mineralogical composition and size of the grains from which the moons once accreted assuming they formed within an impact-generated accretion disk. A comparison of our results with the present day spectral properties of the moons allows us to conclude that their building blocks cannot originate from a magma phase, thus preventing their formation in the innermost part of the disk. Instead, gas-to-solid condensation of the building blocks in the outer part of an extended gaseous disk is found as a possible formation mechanism as it does allow reproducing both the spectral and physical properties of the moons. Such a scenario may finally reconcile their orbital and physical properties alleviating the need to invoke an unlikely capture scenario to explain their physical properties. | During the 70s and 80s, dynamicists have demonstrated that the present low eccentricity, low inclinations and prograde orbits of Phobos and Deimos are very unlikely to have been produced following capture \citep{bur78,pol79}, thus favoring a formation of the moons around Mars \citep{sze83,caz80,gol63}. Despite such early robust evidence against a capture scenario, the fact that the moons share similar physical properties (low albedo, red and featureless VNIR reflectance, low density) with outer main belt D-type asteroids has maintained the capture scenario alive \citep{fra12,fra14,paj13}. Whereas the present orbits of the moons are hardly compatible with a capture scenario, they correspond to the expected outcome of an in situ formation scenario either as the result of co-accretion or of a large impact. Co-accretion with Mars appears unlikely because Phobos and Deimos would consist of the same building block materials from which Mars once accreted. Those building blocks would most likely comprise water-poor chondritic meteorites (enstatite chondrites, ordinary chondrites) and/or achondrites (e.g., angrites), which are all suspected to have formed in the inner ($\leq$2.5 AU) solar system, namely interior to the snowline. This assumption is supported by the fact that the bulk composition of Mars can be well reproduced assuming ordinary chondrites (OCs), enstatie chondrites and/or angrites as the main building blocks \citep{san99, bur04,fit16}. Yet, OCs as well as the remaining candidate building blocks (enstatite chondrites, angrites) are spectrally incompatible with the martian moons, even if space weathering effects are taken into account (see panel b in Figure~\ref{fig1}). It thus appears from above that accretion from an impact-generated accretion disk remains as the only plausible mechanism at the origin of the martian moons. As a matter of fact, the large impact theory has received growing attention in recent years \citep{cra11,ros12,can14,cit15}. This hypothesis is attractive because it naturally explains the orbital parameters of the satellites as well as some features observed on Mars such as (i) its excess of prograde angular momentum possibly caused by a large impact \citep{cra11}, and (ii) the existence of a large population of oblique impact craters at its surface that may record the slow orbital decay of ancient moonlets formed from the impact-generated accretion disk \citep{sch82}. Along these lines, \citet{cit15} have recently shown that a large impact (impactor with 0.01--0.02 Mars masses) would generate a circum-Mars debris disk comprising $\sim$1--4\% of the impactor mass, thus containing enough mass to form both Phobos and Deimos. Although the impact scenario has become really attractive, it has not yet been demonstrated that it can explain the physical properties and spectral characteristics of the martian moons. Here, we investigate the mineralogical composition and texture of the dust that would have crystallized in an impact-generated accretion disk. Since there are no firm constraints regarding the thermodynamic properties of the disk, we perform our investigation for various thermodynamic conditions and impactor compositions. We show that under specific disk's pressure and temperature conditions, Phobos and Deimos' physical and orbital properties can be finally reconciled. | Here, we have opened the possibility that gas-to-solid condensation in the external part of an extended gaseous disk is a likely formation mechanism for the martian moons building blocks as it would lead to the formation of small ($\leq$ 2 microns) dust particles. Accretion from such tiny grains would naturally explain the similarity in spectral properties between D-type asteroids (or comets) and the martian moons as well as their low densities. It therefore appears that accretion in the external part of an impact-generated gaseous disk is a likely formation mechanism for the martian moons that would allow reconciling their orbital and physical properties. Future work should attempt modeling the mineralogy resulting from the gas-to-solid condensation sequence and verify that a space weathered version of the derived composition sieved to sub-micron sized grains is compatible with the spectral properties of the moons. Such investigation would greatly benefit form detailed numerical models that would constrain the thermodynamic properties of a circum-Mars impact generated disk as a function of radial distance. % Finally, note that our proposed scenario is not incompatible with the the presence of weak hydration features in the spectra of Phobos and Deimos \citep{giu11,fra12,fra14}. Water has been delivered to the surfaces of most if not all bodies of the inner solar system, including the Moon \citep{sun09}, Mercury \citep{law13} and Vesta \citep{scu15}. It has thus become clear over the recent years, that space weathering processes operating at the surfaces of atmosphere-less inner solar system bodies do not only comprise the impact of solar wind ions and micrometeorites, which tend to redden and darken the spectra of silicate-rich surfaces, but they also comprise contamination and mixing with foreign materials including water-rich ones \citep{pie14}. | 16 | 7 | 1607.02350 |
1607 | 1607.05600_arXiv.txt | {Rotating bodies in General Relativity produce frame dragging, also known as the {\it gravitomagnetic effect} in analogy with classical electromagnetism. In this work, we study the effect of magnetic field on the gravitomagnetic effect in neutron stars with poloidal geometry, which is produced as a result of its rotation. We show that the magnetic field has a non-negligible impact on frame dragging. The maximum effect of the magnetic field appears along the polar direction, where the frame-dragging frequency decreases with increase in magnetic field, and along the equatorial direction, where its magnitude increases. For intermediate angles, the effect of the magnetic field decreases, and goes through a minimum for a particular angular value at which magnetic field has no effect on gravitomagnetism. Beyond that particular angle gravitomagnetic effect increases with increasing magnetic field. We try to identify this `null region' for the case of magnetized neutron stars, both inside and outside, as a function of the magnetic field, and suggest a thought experiment to find the null region of a particular pulsar using the frame dragging effect. } | \label{secintro} Neutron stars are ideal testing grounds for studying strong gravitational effects as predicted by General Relativity. Frame-dragging \cite{hp} is one such important relativistic effect in compact stars. General relativity predicts that the curvature of spacetime is produced not only by the distribution of mass-energy but also by its motion. Similarly, in electromagnetism, the fields are produced not only by charge distributions but also by currents, i.e., as the moving charge produces the magnetic field. Considering this analogy, this effect is also called the {\it gravitomagnetic effect}. The rotation of a massive body distorts the spacetime \cite{lt}, making the spin of nearby test gyroscope precess \cite{schiff}. It has been pointed out in Ref. \cite{hp} that most of the merit for the discovery of frame-dragging belongs to Einstein but in the existing literature this is referred as Lense-Thirring (LT) effect. Frame-dragging is generally studied in the context of the accretion disc theory where the orbital plane precession frequency is calculated for a {\it test particle} orbiting in the equatorial plane relative to the asymptotic observer. In this article, we derive the frame-dragging precession frequency or LT precession frequency of a test gyroscope ({\it spinning test particle}) relative to the ``fixed stars'' (Copernican system), which was first derived by Schiff \cite{schiff} for the weak field regime. This effect has also been derived using the Kerr metric \cite{cm} without invoking the weak field approximation and it has been shown that the exact formulation reduces to the well-known LT frequency in weak gravity regime, derived in \cite{schiff}. This has been experimentally confirmed by the Gravity Probe B \cite{ev} in the field of the earth and the precession of its orbit has also been measured precisely by the LAGEOS satellite \cite{nat}. The prospect of exploring frame-dragging in the field of a rotating neutron star is attractive, as LT precession frequencies there can be larger than those caused by classical oblateness effects \cite{Morsink}. Further, as the frame dragging effect is also found to occur inside rotating neutron stars and the magnitude of this effect depends on the density profile, it is an interesting prospect to try to constrain the equation of state exploiting this effect \cite{Morsink,Kalogera}. The matter orbits in accretion disks can provide ways to probe the spacetime properties in the vicinity of neutron stars, including frame dragging effects. Quasi-periodic oscillations (QPOs) observed in X-ray Bursts and diagnostics of X-ray spectra have been suggested as promising tools to investigate the motion of matter around neutron stars \cite{Stella}. Recently, it has been predicted \cite{rp} that frame-dragging can be generated by the electromagnetic fields. Using the slow-rotation approximation proposed in \cite{hartle}, the frame dragging effect inside a slowly rotating neutron star can be estimated. In this approximation, to second order in rotation frequencies, the structure of a star changes only by quadrupole terms and the equilibrium equations describing the structure reduce to a set of ordinary differential equations. The frame-dragging frequency in this scheme is only a function of the radial distance from the centre of the star. The precession frequency at the centre is then given by the frequency of rotation of the neutron star, and drops off monotonically towards the surface. Recently some of the authors of this article derived the exact LT frequency of a test gyro \cite{chandra}. The frame dragging rate was found to be not only a function of the rotation frequency of the star but also of the other metric components. Unlike in the slow rotation limit, it was found that the precession rate depends both on the radial distance ($r$) and the colatitude ($\theta$) of the gyroscope. Thus, in rapidly rotating neutron stars, the precession rates are not found to be the same in the equatorial and polar directions. It is well known that neutron stars are endowed with strong magnetic fields. The surface magnetic field value for normal pulsars, estimated from their spin down rates, points to values around $10^{12} - 10^{13}$ G. There have been reports of discovery of high magnetic field neutron stars, both in isolated systems (XDINSs, RRATs) \cite{Olausen, Kaspi} as well as in binaries (SXTs) \cite{Ho}. The largest magnetic fields have been observed in Anomalous X-ray Pulsars (AXPs) and Soft Gamma-ray Repeaters (SGRs), commonly called magnetars. Various observations, including direct detection of cyclotron features in the spectra \cite{Ibrahim, Mereghetti} confirm that these objects possess surface magnetic fields as large as $10^{15} - 10^{16}$ G. However, magnetic fields in the interior of magnetars could be even larger. As it cannot be directly measured, the maximum interior magnetic field is usually estimated using the scalar virial theorem, which points towards a value of $10^{18}$ G. The presence of a strong magnetic field breaks of spherical symmetry of the star, resulting in a strong deformation of its shape from isotropy. The investigation of the effect of magnetic field on the frame-dragging rate in neutron stars is the aim of this paper. The scheme of the paper is as follows. In Sec. \ref{seceq} of this paper, we define the basic equations to describe the exact frame dragging rate in rotating neutron stars in presence of an electromagnetic field. Sec. \ref{secnum} explains the numerical scheme for solving the equations described in the previous section. The results obtained are discussed in Sec. \ref{secres} and Sec. \ref{secnull} is devoted to suggest a thought experiment to find the `null region' of a pulsar using LT precession. Finally Sec. \ref{seccon} closes the article with a summary and future outlook. | \label{secres} An EoS of neutron star matter is needed for this calculation. The recently observed 2M$_{\odot}$ neutron star puts a strong constraint on the EoS \cite{watts}. In an earlier calculation of the LT precession rate without magnetic field \cite{chandra}, we adopted three different EoSs based on the model of Akmal, Pandharipande and Ravenhall (APR), chiral model and density dependent (DD) relativistic mean field model. All three EoSs satisfy the two solar mass neutron star constraint. It was shown in \cite{chandra} that the behaviour of the LT-precession rate for different EoSs is qualitatively similar. Therefore in this study, we employ only the APR \cite{APR} EoS. The maximum masses as well as the corresponding radii of static neutron stars as well as those rotating at Kepler frequency were displayed already in \cite{chandra}, hence we do not repeat this study here. Our aim is to perform a study of the effect of magnetic field on the frame dragging rate $\Omega_{LT}$ for neutron stars rotating at sub-Kepler frequencies. We construct equilibrium configurations of neutron stars for constant values of magnetic moment $\cal{M}$ = $5 \times 10^{30}, 10^{31}, 5 \times 10^{31}, 8 \times 10^{31}, 10^{32}, 1.1 \times 10^{32}~ A.m^2$. For a given magnetic moment, we calculate the value of the frame-dragging precession rate inside ($r \leq r_s$) as well as outside the pulsar ($r>r_s$), where $r_s$ is the distance of the surface of the pulsar from its centre (at the equator $r_s=r_{eq}$; $r_{eq}$ is the equatorial coordinate radius). We should clarify that the LT frequency can be calculated using Eq. (\ref{eq:omega_lt_mod}) both inside and outside of a pulsar. As all the parameters ($A, B, N, N^{\phi}$) in Eq. (\ref{eq:omega_lt_mod}) are functions of $r$ and $\theta$, the equation is automatically modified by changing the values of $r$ and $\theta$. Inside the pulsar ($r \leq r_s$) all of the above mentioned parameters take the values corresponding to the solution of the Einstein equation with the energy-momentum tensor mentioned in Eq. (\ref{emtensor}). Outside the pulsar, ($r > r_s$) Eq. (\ref{emtensor}) reduces to the case with $\varepsilon=P=0$ in absence of matter and Eq.(\ref{eq:metric}) describes the electrovacuum solution of the Einstein equation. Thus, the frame-dragging formula (Eq. (\ref{eq:omega_lt_mod})) also gets modified automatically with the corresponding values of $A, B, N, N^{\phi}$ and the LT precession rates (for a given angle) are continuous from the centre to the asymptotically large distances as well as on the surface of the pulsar. One must note that the effect of magnetic field on the frame-dragging frequency is two fold:\\ (i) Firstly, there is the effect of electromagnetic field directly on the matter and vacuum solutions, given by Eq. (\ref{eq:omega_lt_mod}). \\ (ii) Secondly, magnetic field causes a deformation of the stellar surface. Thus it affects the radial distance at which the energy density falls to zero, and consequently the LT frequency also shifts at the radial distance at which the matter solution matches the electrovacuum solution. \subsection{Interior region ($r \leq r_s$)} In Panel (a) and Panel (b) of Fig. \ref{fig:in}, we display the frame dragging rate as a function of radial distance ($r/r_{eq}$) for different angular values from 0 (along polar direction) to 90 (equatorial direction) for two different constant magnetic moments. The results are qualitatively the same as obtained in \cite{chandra}: the frame dragging frequency has a smooth behaviour along the pole from the centre to the surface of the star but the behaviour is quite different along the equator, which is already been explained in \cite{chandra} and \cite{cc2}. Along the equatorial direction the frame dragging frequency goes through a ``dip'' or local minimum at a certain angle, labeled as ``critical angle". In this case, the critical angle occurs around $\sim 60^0$. For magnetic moments ${\cal{M}} \le 10^{31}$ A m$^2$, the results are indistinguishable from the non-magnetic case. It can also be seen from Panel (a) and Panel (b) of Fig.\ref{fig:in} that frame-dragging rate is quite large $(\sim70\%$ of the total spin of the pulsar) at the centre of the pulsar, where the mass-energy density is largest, and is decreasing in the radial outward direction. This result is compatible with the earlier result obtained in \cite{chandra}. Actually, the frame-dragging rates inside of the pulsars can vary in a wide range depending on their masses, spins and mass-energy densities \cite{weber}. \begin{figure}[tbp] \begin{center} \subfigure[at ${\cal{M}}=5 \times 10^{30}~ A.m^2$]{ \includegraphics[width=.4\textwidth,angle=270]{th5e30.eps}} \subfigure[${\cal{M}}=1.1 \times 10^{32}~ A.m^2$]{ \includegraphics[width=.4\textwidth,angle=270]{th1.1e32.eps}} \caption{Normalized frame-dragging rate $\Omega_{LT}/\Omega$ inside the pulsar as a function of radial distance $r$ in units of $r_{eq}$ along the pole ($\theta = 0^0$) to equator ($\theta = 90^0$ ) for two values of magnetic moment $5 \times 10^{30}$ and $1.1 \times 10^{32}~ A.m^2$.} \label{fig:in} \end{center} \end{figure} Panels (a) to (f) of Fig. \ref{diffmag} reveal that for different values of magnetic moment, depending on the angular direction, the frame dragging rate differs inside the pulsar. Interestingly, the frame dragging rates along the pole ($\theta = 0^0$) decrease with increasing the magnetic moments and the reverse effect is seen along the equator ($\theta = 90^0$ ), i.e., an increase in the frame dragging rates with increasing magnetic moments. Therefore, there should exist a `crossover angle' where the effect of the magnetic field on the frame-dragging will be `null'. Magnetic field of the pulsar does not affect the gravitomagnetic effect in this region, which we can call the `null region'. To find the null region we show the evolution of the frame-dragging for various magnetic moments (starting from $0-1.1\times 10^{32}$ A.m$^2$) from Panel (a) to Panel (f), i.e., from the pole to the equator for every $10^0$ interval. The solid red line represents the zero magnetic moment and black dashed line represents the highest magnetic moment considered. In Panel (a) and Panel (f), the deviation between these two lines is maximum which means that frame-dragging rate is affected maximally by the magnetic moment of the pulsar in these two regions. The evolution from Panel (a) to (c) shows that the distance between these two lines decreases if we proceed from the pole $(0^0)$ to an angle $\theta\sim 40^0$. Deviation is minimum ($\rightarrow 0$) around $\theta \sim 40^0$ (Panel (c)). This shows that magnetism has no effect on frame-dragging at this particular angle. If we proceed further from Panel (d) to Panel (f), we see that the deviation between those two lines (solid red and dashed black) again increases but in this case it is in the reverse direction. Thus the maximum effect of magnetic field is along the polar direction (reducing the frame-dragging frequency) and in the equatorial direction (increasing the frame-dragging frequency), while for intermediate angles its effect goes through a minimum, having practically no influence on the frame-dragging frequency around $40^0$. \begin{figure}[tbp] \begin{center} \subfigure[at $0^0$ (along the pole)]{ \includegraphics[width=2.3in,angle=270]{null0.eps}} \subfigure[at $20^0$ ]{ \includegraphics[width=2.3in,angle=270]{null20.eps}} \subfigure[at $40^0$ (along the null Region) ]{ \includegraphics[width=2.3in,angle=270]{null40.eps}} \subfigure[at $50^0$]{ \includegraphics[width=2.3in,angle=270]{null50.eps}} \subfigure[at $70^0$]{ \includegraphics[width=2.3in,angle=270]{null70.eps}} \subfigure[at $90^0$ (along the equator)]{ \includegraphics[width=2.3in, angle=270]{null90.eps}} \caption{\label{figlt}Plot of normalized frame-dragging rate $\Omega_{LT}/\Omega$ inside the pulsar as a function of radial distance $r$ in units of $r_{eq}$, from the pole ($\theta = 0^0$) to equator ($\theta = 90^0$ ), for different values of magnetic moment.} \label{diffmag} \end{center} \end{figure} \subsection{Exterior region ($r \geq r_s$)} The frame dragging rates are plotted for radial distances extending outside the pulsar ($r \geq r_s$) along polar (panel a), $\theta=50^0$ (panel b) and equatorial directions (panel c) in Fig. (\ref{out}), respectively. This shows that the null region also exists outside the pulsar, at angle of about $50^0$. This is expected as the effect of the magnetic field exists not only inside but also outside (in principle $r\rightarrow \infty$) the pulsar and the frame-dragging should therefore also extend upto $r\rightarrow \infty$. The change in the null angle inside and outside can be understood from the fact that, as the energy density vanishes at the surface of the pulsar, the frame-dragging frequency is no longer determined by the matter solution where $T^{\mu \nu}$ is given by Eq. (\ref{emtensor}) but by the electrovacuum solution with $\varepsilon=P=0$ in $T^{\mu \nu}$. The curves for frame-dragging frequencies converge asymptotically at large distances ($r >> r_s$). However, there is an interplay between the influence of the magnetic field on both the matter and vacuum solutions, as well as on the radial distance at which there is a change from the matter to vacuum solutions. \begin{figure}[tbp] \begin{center} \subfigure[at $0^0$ (along the pole)]{ \includegraphics[width=2.3in, angle=270]{ext_pol.eps}} \subfigure[at $50^0$ (along the null Region)]{ \includegraphics[width=2.3in, angle=270]{ext_theta50.eps}} \subfigure[at $90^0$ (along the equator)]{ \includegraphics[width=2.3in, angle=270]{ext_eq.eps}} \caption{Normalized frame dragging rate $\Omega_{LT}/\Omega$ vs radial distance $r/r_{eq}$ outside the neutron star along (a) polar direction, (b) null Region $(50^0)$ and (c) equatorial direction for different values of magnetic moment.} \label{out} \end{center} \end{figure} Comparing the energy density profiles along different angles for different magnetic moments, we find that for small values of magnetic moments (e.g. $10^{31}$ A.m$^2$ ) the energy density profiles remain the same along all angular directions (see Fig. \ref{fig:enprof_mm} (a)). This is expected as for small magnetic fields the stellar surface remains spherically symmetric. However for large magnetic moments (e.g. $10^{32}$ A.m$^2$ ), the energy density profiles vary along different angles (see Fig. \ref{fig:enprof_mm} (b)). This can be understood from the fact that the stellar surface is distorted from spherical symmetry due to the strong magnetic fields. For a poloidal configuration, the deformation of the stellar surface is oblate (polar flattening and equatorial bulging). This means the energy density goes to zero at smaller radial distances along the polar direction and at larger radial distances along the equatorial direction, defining the distorted stellar surface. It is found that the energy density profile (and hence the stellar surface) is least affected along $50^0$ (comparing Panel (a) and Panel (b) of Fig. \ref{fig:enprof_mm}). \begin{figure}[tbp] \begin{center} \subfigure[$10^{31}$ A.m$^2$]{ \includegraphics[width=2in, angle=270]{enprof_mm1e31.eps}} \subfigure[$10^{32}$ A.m$^2$]{ \includegraphics[width=2in, angle=270]{enprof_mm1e32.eps}} \caption{Energy density profile (in units of nuclear density $\rho_{nuc}$) along polar direction (red solid line), $40^0$ (magenta dashed line), at $50^0$ (green dotted line) and equatorial direction (brown dot-dashed line) for magnetic moment of (a) $10^{31}$ A.m$^2$ and (b) $10^{32}$ A.m$^2$.} \label{fig:enprof_mm} \end{center} \end{figure} \subsection{Role of the rotation frequency of the pulsar} \label{sec:rot} To determine the influence of the value of pulsar rotation on the position of the null region, we repeat the numerical calculations for $\Omega$ = 0, 100, 1000 $s^{-1}$. As the rotation frequency increases to large values. the surface of the star is also deformed to a non-spherical shape. However, we found that when the rotation frequency is zero, the LT frequency vanishes, no matter what the value of the magnetic field. This implies that there is no contribution of free precession due to the oblateness caused by the magnetic field, and that the LT effect arises purely due to the rotation of the spacetime. The LT frequency is non-zero only in presence of rotation, and its magnitude is affected by the presence of the magnetic field. However, we have also observed that the value of the rotation frequency of the pulsar has no influence on the value of the angle at which the magnetic field ceases to affect the frame dragging rate. For the interior region, the null region is found at $\sim 40^0$ while for the exterior region it is found at $\sim 50^0$, as obtained for $\Omega$ = 276.8 $s^{-1}$ for the PSR J0737-3039. Here, we should emphasize that any kind of magnetic field is not directly responsible for the frame-dragging. This effect solely depends on the rotation of the spacetime (there is also an exception where LT effect is not related to the rotating spacetime but some ``rotational sense'' was involved (\cite{cm}) there). Thus, if a magnetic field exists in a non-rotating spacetime, the spacetime will not be able to show the frame-dragging. If the spacetime starts to rotate, gravitomagnetism is induced in this spacetime and the magnetic field comes into play to show its effect on frame-dragging. In an analogy with electromagnetism, we can say that as a non-magnetic material is not affected by the magnetic field of any magnet, the non-rotating spacetime is also not affected by the magnetic field of that spacetime in the case of frame-dragging. It is evident form Eq.(\ref{eq:metric}) that $N^{\phi}$ vanishes for a non-rotating spacetime ($\Omega=0$) and therefore the gravitomagnetic effect vanishes ($\Omega_{LT}=0$) which is also evident from Eq.(\ref{eq:omega_lt}), though the magnetic field is non-zero (see Eq.(\ref{emtensor})). The {\it tussle between magnetic field and gravitomagnetic effect is only possible if they are both present in a particular spacetime}. If any one of them is absent, the tussle will also be absent. That is why it is meaningless to isolate the role of magnetic moment on LT effect from that due to rotation only, but the deviation of the LT effect due to the presence of the magnetic field can easily be deduced from the Panels (a)-(f) of Fig. \ref{diffmag} for various magnetic moment values at various angles. \subsection{Role of the magnetic field geometry} \label{sec:geometry} It is instructive to discuss the role played by the simplifying assumptions that went into the construction of the model for rotating magnetized neutron stars. It was outlined in Sec. \ref{sec:global} that the poloidal field geometry considered in this work was intuitive, but definitely not the most general case. Unlike the external field geometry, the internal magnetic field configuration cannot be determined by direct observations. It has been suggested that differential rotation in newborn neutron stars could create a strong toroidal magnetic component. The effect of purely toroidal fields on the structure of neutron stars has been elaborately studied in ~\cite{Kiuchi, Yasutake, Frieben}. Numerical simulations of purely poloidal and toroidal geometries have found them to be unstable \cite{Lander2011,Braithewaite06}. However some studies found rapid uniform rotation to suppress instability in both purely poloidal and toroidal configurations. Among mixed-field configurations, the twisted-torus geometry was found to be the most promising candidate for a stable magnetic field configuration \cite{Yoshida06, Glampedakis2012, Ciolfi09,Pili2014}. All these models consider a poloidal-dominated geometry, with toroidal-to-poloidal energy ratio restricted to less than 10$\%$. Recently, Ciolfi and Rezzolla \cite{CiolfiRezzolla} succeeded in producing toroidal dominated twisted-torus configurations, but their stability in nonlinear simulations is still to be established. Although the frame dragging rate depends only on the spacetime metric and not directly on the magnetic field geometry (see Sec. \ref{sec:ltprec}), the interplay of the frame dragging rate with the magnetic field profile should depend on the chosen magnetic field geometry. It has been found that while purely poloidal magnetic fields deform a neutron star shape and matter distribution to an oblate shape (equatorial radius larger than polar radius) as in Fig. \ref{fig:xzcut}, purely toroidal fields fields deform them to a prolate shape. On the other hand, rotation induces oblateness in the surface deformation. Thus for rotating magnetized neutron stars with a purely toroidal geometry, different combinations of the surface deformation are possible depending on the relative strength of the field and rotation rate \cite{Frieben}. Thus to summarize, the matter distribution of the star can vary with the magnetic field geometry in a very complicated way, and this would affect its interplay with frame-dragging. But this is beyond the scope of this work and we leave the general case for future study. However it must be outlined here that for both purely poloidal and toroidal cases, the stellar shape would be least distorted close to an angle $40-50^0$, and there should exist a null region following the same arguments. \label{seccon} In this work, we extended the recent calculation \cite{chandra} for exact LT precession rates inside and outside neutron stars to include effects of electromagnetic field. This should be particularly interesting for the case of high magnetic field pulsars and magnetars. We employed a well established numerical scheme \cite{BGSM,chatterjee,Bocquet} to obtain equilibrium models of fully relativistic magnetized neutron stars with poloidal geometry. We computed the frame dragging rates inside and outside the stellar surface for such models. We obtained qualitatively similar peculiar features for the frame dragging rates along polar and equatorial directions as previously discussed for non-magnetic calculations \cite{chandra}, with the effect of the electromagnetic field evident on the frame dragging rate inside as well as outside the star. This could be potentially interesting for high magnetic field pulsars to constrain the magnetic field from a measurement of the LT frequency. As an example for PSR J0737-3039, we found that the maximum effect of magnetic field appears along the polar direction (decreasing the frame-dragging frequency) and in the equatorial direction (increasing the frame-dragging frequency), while for intermediate angles its effect goes through a minimum, having practically no influence on the frame-dragging frequency around $40^0$ inside and $50^0$ outside the pulsar. We showed that this can be attributed to the twofold effect of the magnetic field, on the LT frequency as well as on the breaking of spherical symmetry of the neutron star. We also devised a thought-experiment to determine the null angle using a measurement of its LT frequency. Measurements of geodetic and frame dragging effects by dedicated missions such as LAGEOS and Gravity Probe B \cite{Overduin} have permitted us to impose constraints and better understand gravitomagnetic-magnetic aspects of theories of gravity. There is a need to expand the sensitivity of future gravitational tests to neutron stars to understand phenomena such as LT precession. We have discussed briefly the possibility of measuring frame dragging in neutron stars in the future. | 16 | 7 | 1607.05600 |
1607 | 1607.02216_arXiv.txt | We describe a directed search for continuous gravitational waves in data from the sixth initial LIGO science run. The target was the nearby globular cluster NGC~6544 at a distance of $\approx 2.7$ kpc. The search covered a broad band of frequencies along with first and second frequency derivatives for a fixed sky position. The search coherently integrated data from the two LIGO interferometers over a time span of 9.2 days using the matched-filtering $\mathcal{F}$-statistic. We found no gravitational-wave signals and set 95\% confidence upper limits as stringent as $6.0 \times 10^{-25}$ on intrinsic strain and $8.5 \times 10^{-6}$ on fiducial ellipticity. These values beat the indirect limits from energy conservation for stars with characteristic spindown ages older than 300 years and are within the range of theoretical predictions for possible neutron-star ellipticities. An important feature of this search was use of a barycentric resampling algorithm which substantially reduced computational cost; this method will be used extensively in searches of Advanced LIGO and Virgo detector data. | 16 | 7 | 1607.02216 |
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1607 | 1607.07928_arXiv.txt | { The global structure of solar corona observed in optical window is governed by the global magnetic field with different characteristics over solar activity cycle. Ludendorff flattening index becomes a popular measure of the global structure of solar corona as observed during eclipse. In this study, 15 digital images of solar corona from 1991 to 2016 were analyzed in order to construct the coronal flattening profiles as a function of radius. In most of the cases, the profile can be modeled with 2nd order polynomial function so that the radius with maximum flattening index ($\rmax$) can be determined. Along with this value, Ludendorff index ($a+b$) was also calculated. Both Ludendorff index and $\rmax$ show anti-correlation with monthly sunspot number, though the $\rmax$ values are more scattered. The variation of $\rmax$ can be regarded as the impact of changing coronal brightness profile over equator. | \label{sect:intro} Corona is the outer part of solar atmosphere with density of $\sim10^{15}$ m$^{-3}$ and temperature of millions of Kelvin. This layer has a total brightness of about $4\times10^{-6}$ times to the brightness of solar phototsphere \citep{hanaoka2012} such that observable only in short wavelength range (extreme ultraviolet and x-ray) or when the glaring photosphere is blocked by coronagraph or lunar disk during solar eclipse. Therefore, relatively rare occurrence of total solar eclipses provide opportunity to study solar corona in optical window (or white light) from the ground. In this window, the corona can be categorized into K (\emph{kontinuerlich}) corona and F (\emph{Fraunhofer}) corona with different properties. Continuum radiation of K corona, which dominates the inner part ($r<2$ \rsun), is caused by Thompson scattering of photosphere radiation by electron in the corona. In the outer part, interplanetary dust scatter Sun's radiation and create the F corona \citep{foukal2004}. Global structure of solar corona observed in optical window represents the electron distribution in this layer which is influenced by local and global magnetic field extending from photosphere to corona \citep[e.g.][]{sykora2003,pasachoff2009}. The change of coronal structure or shape over solar activity cycle is clearly observed. During minimum, there are few active regions and helmet streamers are relatively concentrated near the equator so that the corona tends to be flattened. Conversely, solar corona becomes more radially symmetric during maximum phase as the streamers are more evenly distributed over heliographic latitudes. Quantitative parameters were defined to describe global structure of solar corona, some of which are photometric or Ludendorff flattening index \citep{ludendorff1928}, geometric flattening index \citep{nikolsky1956}, angular extent of streamer-free polar regions \citep{loucif1989}, modified flattening index \citep[accounting magnetic tilt][]{gulyaev1997}, and latitudinal span of helmet streamers \citep{tlatov2010}. Among those parameters, Ludendorff index becomes the most popular measure and regularly obtained from every corona observation during solar eclipse \citep{pishkalo2011}. This index can be regarded as the flattening of solar corona at 2{\rsun} heliocentric distance whose values range from $\sim0$ during solar maximum to $\sim0.4$ during solar minimum. \cite{pishkalo2011} has already compiled Ludendorff index from 1851 to 2010 and demonstrated the correlation between this index and monthly sunspot number ($SSN$). However, a rather scattered data obscures this correlation. Similar problem occurred at the variation of flattening index as a function of solar activity phase. From particular eclipse event, some observers may get different flattening index. This difference arise from several influencing factors such as observational bias \citep{sykora1999}, diverse detector characteristics (e.g. film emulsion), exposure time, number of isophotal contours used to calculate flattening index, poorly-oriented image, to the different statistics implemented \citep{pishkalo2011}. These factors can be minimized or even eliminated by implementing homogeneous method of analysis that becomes a main focus of this study. The objectives of this study are to re-analyze publicly available coronal images from 1991 to 2016 solar eclipses and to construct radial profiles of flattening index. From each profile, $\rmax$ that represents the equatorial radius with maximum flattening index can be determined so that the variation of this value over solar cycle can be examined together with Ludendorff index. Data used and method applied in this study are explained in \autoref{sec:data}, while the result and discussion are presented in \autoref{sec:result}. The study is concluded in \autoref{sec:conclusion}. | \label{sec:conclusion} In this study, 15 white-light solar coronal images taken during solar eclipses that occurred in 1991 to 2016 have been analyzed using semi-autonomous method such that shape parameter of the corona can be determined. Flattening profiles as a function of radius for very images have been produced. In most cases (80\% sample), the flattening profile can be modeled using 2nd order polynomial function such that the radius of maximum flattening ($\rmax$) can be determined. At small heliocentric distances ($r\leq\rmax$), flattening index increases almost linearly and the Ludendorff index (flattening at $r=2$ \rsun) can be extrapolated. In agreement with previous studies, Ludendorff index anti-correlates with monthly sunspot number. Additionally, this study shows that $\rmax$ change over solar cycle at the same phase of flattening index variation. The change of $\rmax$ can be interpreted as the observational consequences of the change in equatorial brightness profile that is different to the brightness profile in polar direction. \normalem | 16 | 7 | 1607.07928 |
1607 | 1607.01329_arXiv.txt | There is evidence that coronal heating is highly intermittent, and flares are the high energy extreme. The properties of the heat pulses are difficult to constrain. Here hydrodynamic loop modeling shows that several large amplitude oscillations ($\sim 20$\% in density) are triggered in flare light curves if the duration of the heat pulse is shorter that the sound crossing time of the flaring loop. The reason is that the plasma has not enough time to reach pressure equilibrium during the heating and traveling pressure fronts develop. The period is a few minutes for typical solar coronal loops, dictated by the sound crossing time in the decay phase. The long period and large amplitude make these oscillations different from typical MHD waves. This diagnostic can be applied both to observations of solar and stellar flares and to future observations of non-flaring loops at high resolution. | \label{sec:intro} One fundamental question about coronal heating is whether the energy released in coronal loops is gradual or impulsive \citep{Klimchuk2006a,Parnell2012a,Reale2014a}. There is increasing evidence in the time series that in active regions the heating might be highly irregular \citep[e.g.,][]{Sakamoto2008a,Sakamoto2009a,Vekstein2009a,Terzo2011a,Viall2011a,Viall2012a,Ugarte-Urra2014a,Tajfirouze2016a,Tajfirouze2016b}. There are different colors about this question, for instance whether the heat pulses are frequent or not with respect to the plasma cooling time \citep[e.g.,][]{Warren2011a}. This is very difficult to constrain because coronal loops are likely structured into thin strands, where the heating is released through a storm of small-scale pulses. The single heating episodes can be hardly resolved up to date. In the framework of intermittent heating, an important issue is the duration of the heat pulses, because it links directly to the basic mechanisms of magnetic energy conversion, e.g., reconnection. For instance, recent work has suggested that the pulses are preferentially short \citep[$\leq 1$ min,][]{Testa2014a,Tajfirouze2016a}. Even in flares the duration of the pulses is difficult to diagnose because the efficient heat conduction thermalizes the whole flaring loop in few seconds, cancelling all heating signatures in the EUV and soft X-rays. Some diagnostics about the features of the heating come from the hard X-rays, which track the emission from non-thermal electron beams. However, it is debated whether the electron beams are entirely responsible for the flare heating or they co-exist with other mechanisms excited by fast magnetic reconnection, e.g., the dissipation of current sheets \citep[e.g.,][]{Battaglia2015a}. Here we propose a way to diagnose how long is the heating release in brightening coronal loops. Coronal loops can be described as closed magnetic tubes where the plasma is confined, moves and transports energy along the field lines. If the heat pulse is short, strong pressure waves are triggered inside the magnetic flux tube. Since the temperature is very uniform along the loop because of the efficient thermal conduction, the pressure waves \aftr{manifest as} steep density fronts that slosh up and down along the tube. These are purely hydrodynamic waves. The density fronts determine visible periodic fluctuations in the light curves, which may be detected. This kind of fluctuations has been detected in the solar non-flaring \citep{Harrison1987a,Wang2003a} and flaring \citep{Nakariakov2009a} corona and in stellar coronal flares \citep{Mitra-Kraev2005a,Welsh2006a,Pandey2009a,Lopez-Santiago2016a}, and generally interpreted in terms of MHD harmonic modes. In the following we investigate the details and conditions for the presence of these fluctuations through hydrodynamic loop modeling. | \label{sec:discus} We have shown that short heat pulses can excite large amplitude wavefronts of plasma confined in coronal loops. The critical time scale is the return sound crossing time (or the sound crossing time along the entire loop length) at the temperature peak. If the pulse duration is shorter than this time scale, then there is not enough time to equalize the pressure in the initial transient, and plasma sloshing is triggered back and forth between the apex and the footpoints. Since the efficient thermal conduction keeps the temperature very uniform along the loop, the pressure fronts are mostly density fronts, and determine strong fluctuations in the emission that can be detected in the light curves taken in the appropriate bands. We remark that we are modeling plasma flowing freely along the flux tube and that there is no direct interaction with the magnetic field, except for confinement and channelling. The excited waves are therefore purely hydrodynamic waves in a compressible plasma, different from low-order MHD modes, such as sausage or kink modes. The assumption of closed loops symmetric with respect to the apex makes the model evolution particularly clear and well-behaved. This scenario is \aftr{an acceptable simplification} because we might expect that magnetic reconnection triggers heat pulses deposited symmetrically at both loop footpoints. Also if the heat pulses are spread in the coronal part of the loop, the efficient thermal conduction would level out the temperature along the whole loop. Twin density fronts would anyway arise from the chromosphere at both footpoints and with a very small time difference, and they would hit against each other high in the loop determining the initial accumulation that triggers the sloshing. This might occur not exactly at the loop apex, so, if the evolution is not totally symmetric, we might expect more irregular quasi-periodic patterns. \aftr{In addition, loops that are not symmetric could have very different gravitational stratifications in each leg, leading to more irregular patterns.} The amplitude of the density waves is large and even larger in the plasma emission, because of the dependence on the square of the density. These are not standing waves nor acoustic harmonic oscillations inside the loop \citep{Selwa2005a}, and they have been customarily found in previous loop modeling \citep[e.g.,][]{Reale2012a,Bradshaw2013a}. We might expect to detect them easily in the light curves, whenever present. At variance from typical magnetoacoustic waves, their period is relatively large, minutes (or more for longer loops), and therefore easy to identify. They may be best detected if the heating is released almost all at once across a loop, to have a coherent evolution, as in proper flares. Also, the large amplitude makes them different from typical MHD waves. Another point is interesting to remark. A general decay time has been found for plasma flaring in single loops \citep{Serio1991a,Reale2007b,Reale2014a}: \begin{equation} \tau_d = 120 \frac{L_9}{\sqrt{T_7}} \label{eq:tau_d} \end{equation} This decay time scales exactly as the sound crossing time~\ref{eq:tau_s}, i.e., the period of the waves scales as the decay time of the flare. Since the decay is typically the longest part of a flare, the implication is that, whatever the flare duration, any flare light curve will contain a similar number of major oscillations (not many, typically around 5). In the end, we propose that periodic oscillations detected in the light curves of solar and stellar flares are often due to plasma sloshing as modelled in the present study and that their presence depends on the duration of the flare heating related to the flaring loop length (whereas the dependence on the temperature is relatively weak). Thus, this becomes a new way to identify pulsed heating and to constrain its duration. This does not seem to be so frequent in solar flares, probably because the length of the loops that brighten initially is often quite small. In spite of the smaller signal to noise, it is instead more frequent in stellar flares which can occur in giant magnetic channels \citep{Lopez-Santiago2016a}. We can extend this result to flares at any scale, and in particular to small scales (nanoflares). We could expect to detect oscillations in light curves from non-flaring coronal loops in active regions, as observed, for instance, by SDO/AIA. This is not typically the case \citep[e.g.][]{Sakamoto2008a,Sakamoto2009a,Viall2011a,Tajfirouze2016a,Tajfirouze2016b}. but short heat pulses may still be present in the framework of multi-stranded pulse-heated loops. \cite{Tajfirouze2016a} find that short pulses match better the observed light curves and \cite{Tajfirouze2016b} show that, even if there are strong oscillations in the single light curves, they are washed out when they are mixed up along the line of sight across a loop with a multitude of independently heated strands. We might hope to detect such oscillations even in non-flaring loops with appropriate resolution of next generation instruments. | 16 | 7 | 1607.01329 |
1607 | 1607.08208_arXiv.txt | Measurements of cosmic microwave background (CMB) anisotropies provide strong evidence for the existence of dark matter and dark energy. They can also test its composition, probing the energy density and particle mass of different dark-matter and dark-energy components. CMB data have already shown that ultra-light axions (ULAs) with mass in the range $10^{-32}~{\rm eV} \to 10^{-26}~{\rm eV}$ compose a fraction $\lsim 0.01$ of the cosmological critical density. Here, the sensitivity of a proposed CMB-Stage IV (CMB-S4) experiment (assuming a 1 arcmin beam and $\sim 1~\mu K{\rm-arcmin}$ noise levels over a sky fraction of 0.4) to the density of ULAs and other dark-sector components is assessed. CMB-S4 data should be $\sim 10$ times more sensitive to the ULA energy-density than {\sl Planck} data alone, across a wide range of ULA masses $10^{-32}\lsim m_{a}\lsim 10^{-23}~{\rm eV}$, and will probe axion decay constants of $f_{a}\approx 10^{16}~{\rm GeV}$, at the grand unified scale. CMB-S4 could improve the CMB lower bound on the ULA mass from $\sim 10^{-25}~{\rm eV}$ to $10^{-23}~{\rm eV}$, nearing the mass range probed by dwarf galaxy abundances and dark-matter halo density profiles. These improvements will allow for a multi-$\sigma$ detection of percent-level departures from CDM over a wide range of masses. Much of this improvement is driven by the effects of weak gravitational lensing on the CMB, which breaks degeneracies between ULAs and neutrinos. We also find that the addition of ULA parameters does not significantly degrade the sensitivity of the CMB to neutrino masses. These results were obtained using the \textsc{axionCAMB} code (a modification to the \textsc{CAMB} Boltzmann code), presented here for public use. | \label{intro} \begin{figure}[htbp!] \includegraphics[width=0.5\textwidth, trim = 5mm 0mm 10mm 10mm, clip]{Log_constraints_axions_s4.pdf} \caption{\textbf{Projected CMB-S4 sensitivity to the axion energy density as a function of axion mass, compared with Fisher-matrix {\sl Planck} sensitivity}: Vertical bars show $1\sigma$ errors at fixed neutrino mass $\Sigma m_\nu = 0.06~\mathrm{eV}$ while the shaded bars show the errors marginalizing over $\Sigma m_{\nu}$. We classify axions as DE-like if $m_a < 10^{-29}\,\mathrm{eV}$, `DM-like' if $m_a > 10^{-25} \mathrm{eV}$ and `fuzzy' DM for masses in between. In the `fuzzy' DM region, CMB-S4 will allow for percent-level sensitivity to the axion mass fraction, improving significantly on current constraints. For {\sl Planck} data alone, neutrino degeneracies significantly degrade sensitivity to axions, even at the 1$\sigma$ level. In contrast, CMB-S4 constraints remain robust to varying neutrino mass in the `fuzzy' region. The solid and dashed lines show the $2\sigma$ and $1\sigma$ exclusion limits, i.e. the lowest axion fraction that could be excluded at those masses. \label{fig:axions}} \end{figure} Identifying dark matter (DM) remains one of the outstanding cosmological challenges of the current age. While searches for direct or indirect evidence of a dark matter candidate continue \cite{Buckley:2013bha,Cushman:2013zza}, the effect of dark matter on cosmological observables provides a complementary approach to constraining the dark sector. In the face of increasingly strict experimental limits to Weakly Interacting Massive Particle (WIMP) DM, axions are re-emerging as a popular alternative (see Ref.~\cite{Marsh:2015xka} for an extensive review of axions). Cosmological axion production can proceed through decays of exotic particles (e.g. moduli) or topological defects, thermal production from the standard-model plasma, or coherent oscillation around a misaligned (from the vacuum state) initial value, known as vacuum realignment. If axions are also produced because of non-vanishing matter couplings, a relativistic population can be produced, contributing to the relativistic energy density in the early universe (parameterized by a generic parameter $N_\mathrm{eff}$, describing the number of \textit{relativistic degrees of freedom}). Constraints on these axion models were presented in Refs. \cite{Acharya:2010zx,Weinberg:2013kea,Iliesiu:2013rqa,Baumann:2016wac}. Vacuum realignment is the only axion production mechanism that occurs independent of assumptions about axion couplings or inflationary physics, and produces an extremely cold population of axions, in contrast with other mechanisms. Here, we consider only axions produced by vacuum realignment.\footnote{We do vary $N_\mathrm{eff}-3.046$, but without bias as to its physical nature.} Ultralight axions (ULAs) produced via vacuum realignment with masses in the range $10^{-33}~\mathrm{eV}\leq m_{a}\leq 10^{-20}~\mathrm{eV}$ are well motivated by string theory, and can contribute to either the dark matter or dark energy components of the Universe, depending on their masses \cite{Marsh:2015xka}. They are distinguishable from standard dark energy (DE) and cold dark matter (CDM) using cosmological observables such as the cosmic microwave background (CMB) temperature and polarization power spectra, the matter power spectrum (as probed using the correlations of galaxy positions and shapes) and the weak gravitational lensing of the CMB. Constraints on the allowed contribution of ULAs to the total DM component using these observables provide a test of the CDM scenario. A key goal of future cosmological experiments is to measure the sum of the neutrino masses, $\Sigma m_\nu$ (see Ref.~\cite{Lesgourgues:2012uu} for a review of neutrino cosmology). The current bound on $\Sigma m_\nu$ from ground-based oscillation experiments is $\Sigma m_\nu \gtrsim 0.06~\mathrm{eV}$ \cite{Otten:2008zz}. Current cosmological neutrino bounds indicate that $\Sigma m_\nu < 0.23~\mathrm{eV}$ at 95\% confidence, using data from Planck \cite{Ade:2015xua} and measurements of Baryon Acoustic Oscillations (BAO) from the Baryon Oscillation Spectroscopic Survey \cite[BOSS,][]{Beutler:2014yhv}. Forecasted constraints for neutrino masses are that $\sigma(\Sigma m_\nu)= 15~\mathrm{meV}$ for a fiducial model with $\Sigma m_\nu = 60~\mathrm{meV}$, for a CMB-S4-like experiment and BAO measurements from a `DESI-like' survey \cite{Abazajian:2013oma}, promising a 4$\sigma$ detection of neutrino mass \cite{Allison:2015qca}. Much of this improvement is driven by weak gravitational lensing of the CMB, in particular at high multipoles $\ell\gtrsim 1000,$ although the change in the lensing convergence power is of order 25\% even at low multipoles. The lensing deflection power-spectrum is determined from $4$-pt functions of CMB maps, extracting a factor of $\sim\sqrt{3}$ as much information from CMB experiments~\cite{Scott:2016fad}. The promise of CMB experiments in probing neutrino masses motivates us to wonder: will future CMB experiments offer dramatic improvements in sensitivity to axion parameters? Given the known similarity of ULA and massive neutrino imprints \cite{Marsh:2011bf} on cosmological observables at low mass ($m_{a}\lsim 10^{-29}~{\rm eV}$), how significant are ULA-neutrino degeneracies at CMB-S4 sensitivity levels and will they degrade our ability to do fundamental physics with the CMB? To answer these questions, we conduct a Fisher-matrix analysis to explore the sensitivity of future CMB experiments to ULA masses, densities, and $\Sigma m_{\nu}$. We find that CMB-S4 will allow a $2-5\sigma$ detection of axion mass fractions that agree with pure {\sl Planck} limits, covering an axion mass range of $10^{-32}~{\rm eV}\lsim m_{a}\lsim10^{-24}~{\rm eV}$. Near the top of this range, CMB-S4 will break the degeneracy of axions and CDM. Sensitivity persists (but tapers off) towards higher axion masses of $m_{a}\sim 10^{-23}~{\rm eV}$. CMB-S4 will push CMB tests of the ULA hypothesis towards the mass range probed by subtle observables, like the size of DM-halo cores and the number of missing Milky-Way satellites. In the ``dark-energy-like" (``DE-like" ULAs henceforth) ULA regime ($m_{a}\lsim 10^{-29}~{\rm eV}$) we find that the the ULA mass fraction is degraded by degeneracies with the sum of the neutrino masses, but that this degeneracy disappears at higher masses. We find also that future measurements of the Hubble constant could break this degeneracy. We denote ULAs in the mass range $10^{-29}~{\rm eV}\lsim m_{a}\lsim 10^{-25}~{\rm eV}$ as ``fuzzy DM", and those with $m_{a}\gsim 10^{-25}~{\rm eV}$ as ``dark-matter-like" (or DM-like). We find that measurements of the lensing-convergence power spectrum $C_{\ell}^{\kappa \kappa}$ drive much of the improvement in sensitivity; if lensing is omitted, the fractional error bar on the axion mass fraction degrades by a factor of $\sim 3-5$ in the `fuzzy' regime. Finally, we explore the dependence of our results on CMB-S4's experimental design parameters. We begin this paper by summarizing the physics and cosmology of ULAs and neutrinos in Section~\ref{sec:axion_cosmo}. In Section~\ref{observables}, we discuss the effects of ULAs and neutrinos on cosmological observables (e.g., the CMB's primary anisotropies and its lensing-deflection power spectrum), as well as the degeneracies between axions and cosmic neutrinos. Our assumptions about future data, forecasting techniques, and key science results are presented in Section \ref{sec:results}. We conclude in Section \ref{conclusion}. All power spectra presented here were computed using the \textsc{AxionCAMB} code, a modification to the CMB anisotropy code \textsc{CAMB} \cite{cambnotes}, which is described in Appendix \ref{code_appendix}, is publicly available, and was used to obtain the ULA constraints of Ref.~\cite{Hlozek:2014lca}.\footnote{The code may be downloaded from \url{http://github.com/dgrin1/axionCAMB}.} In Appendix \ref{nonlinear_appendix}, we discuss the computation of the nonlinear matter power-spectrum (relevant for understanding the effect on weak lensing on the CMB). | } We live in the age of precision cosmology. Future experiments like the proposed CMB-S4 will significantly improve constraints on the composition of the dark sector. We have shown in detail how this is achieved in the case of ultra-light axions, including degeneracies with dark radiation and massive neutrinos. CMB-S4 will move the wall of ignorance for the heaviest axion candidates from $m_a=10^{-26}~\mathrm{eV}$ to $m_a=10^{-24}~\mathrm{eV}$ (detection with an axion fraction of 20\% at $>3\sigma$). The lower limit on the dominant DM particle mass will be increased from $m_a=10^{-25}~\mathrm{eV}$ to $m_a=10^{-23}~\mathrm{eV}$ ($1\sigma$ constraints rule out large fractions). This begins to make contact with the much more systematic-laden upper bounds on the axion mass and fraction from high-$z$ galaxies and reionization: $\Omega_a/\Omega_d<0.5$ for $m_a=10^{-23}\text{ eV}$ and $m_a\gtrsim 10^{-22}\text{ eV}$ for the dominant component~\cite{Bozek:2014uqa,Schive:2015kza,Sarkar:2015dib}. This value approaches the mass range needed to explain dwarf galaxy cores and missing Milky Way satellites~(e.g. Refs.~\cite{Hu:2000ke,Marsh:2013ywa,Schive:2014dra,Marsh:2015wka}). Perhaps more impressively, the constraints on the axion energy density at intermediate mass could improve by an order of magnitude. CMB-S4 could detect an axion fraction as low as $0.02$ at $>13\sigma$ for an axion mass of $10^{-27}\, \mathrm{eV}$. Given the power of these future efforts, it will be possible to probe the degeneracies between ULAs and other potential DM components, such as massive neutrinos, and light species such as massless sterile neutrinos. Improved independent constraints on measurements of the expansion rate (through measurements of the Hubble constant, for example) will improve sensitivity to the lightest, DE-like axions, and reduce the degeneracy between these species and both $\Sigma m_\nu$ and $N_\mathrm{eff}$. Even when marginalizing over the neutrino mass, the error on the axion fraction for a mass of $m_a = 10^{-32}\, \mathrm{eV}$ improves by a factor of three with a prior on the expansion rate. As $\Omega_a\propto f_a^2$ the improved sensitivity to the axion energy density improve the axion decay constant which could be detected from $10^{17}\mathrm{ GeV}$ with {\sl Planck} to $10^{16}\mathrm{ GeV}$ with CMB-S4 (over the relevant range of ULA masses). The improved sensitivity to $f_a$ will begin to test the predictions of the string axiverse scenario~\cite{Arvanitaki:2009fg}. Axions are a well motivated dark matter candidate, and future CMB experiments suggest an exciting opportunity to explore the rich complexity of their parameter space, moving towards sub-percent level sensitivity to the axion energy density or a $10\sigma$ detection if current limits to $\Omega_{a}$ are saturated by the true axion density, all over for a wide range of masses. As a spectator field during the inflationary era, axions would also carry isocurvature perburbations (see Ref. \cite{Hlozek:2014lca} and references therein), leading to distinct imprints on CMB observables and providing a unique new lever arm on the inflationary energy scale, which is otherwise only accessible through measurements of primordial CMB B-mode polarization \cite{Marsh:2014qoa}. In future work, we will extend {\sl Planck} constraints and CMB-S4 forecasts to include the impact of isocurvature. Unraveling the mystery of dark matter is an important goal for cosmology in the coming decades. The axion represents the lowest mass DM-candidate, and a `CMB-S4-like' survey will help identify (or rule out) these models of DM. Constraints on the light, DE-like axions are improved by independent measurements of the expansion rate of the Universe, thereby probing our knowledge of the cosmological constant, quintessence, and cosmic acceleration in general. In this work, we have illustrated that future CMB experiments will shed new light on the nature or existence of the axion and usher axiverse cosmology into a new era. | 16 | 7 | 1607.08208 |
1607 | 1607.03095_arXiv.txt | Should we expect most habitable planets to share the Earth's marbled appearance? For a planetary surface to boast extensive areas of both land and water, a delicate balance must be struck between the volume of water it retains and the capacity of its perturbations. These two quantities may show substantial variability across the full spectrum of water-bearing worlds. This would suggest that, barring strong feedback effects, most surfaces are heavily dominated by either water or land. Why is the Earth so finely poised? To address this question we construct a simple model for the selection bias that would arise within an ensemble of surface conditions. Based on the Earth's ocean coverage of $71\%$, we find substantial evidence (Bayes factor $K \simeq 6$) supporting the hypothesis that anthropic selection effects are at work. Furthermore, due to the Earth's proximity to the waterworld limit, this model predicts that most habitable planets are dominated by oceans spanning over $90\%$ of their surface area ($95\%$ credible interval). This scenario, in which the Earth has a much greater land area than most habitable planets, is consistent with results from numerical simulations and could help explain the apparently low-mass transition in the mass-radius relation. | \begin{figure*} \includegraphics[width=\columnwidth]{fig8p1} \hspace{0.2cm} \includegraphics[width=\columnwidth]{fig7p2} \caption{The oceanic fine-tuning problem. \emph{Left:} The ocean coverage as a function of the surface water volume, normalised in terms of the capacity of surface perturbations. The solid line describes any solid surface with a Gaussian hypsometry. This will serve as our model for the statistical average across all habitable worlds. The dot-dashed line depicts the effect of adopting the shape of the Earth's elevation profile \citep{eakins2012hypsographic}, while the dotted line (barely distinguishable) corrects for the isostatic depression of the seabed. The thick dashed line shows the response for the elevation profile of Mars. The thin vertical dashed line demonstrates that the Earth's value of $\Searth \simeq 4$ is precariously close to the waterworld limit. \emph{Right:} Three different models of the habitable land area, expressed as a fraction of the total surface area. The dotted, dashed and solid lines correspond to values of $\alpha = \frac{1}{5}, 1, 5$ respectively, and the habitable area is defined by equation (\ref{eq:H}). These curves are generated using a Gaussian hypsometry, as depicted in the left hand panel.} \label{fig:landfrac} \end{figure*} The methane seas of Titan are the only exposed bodies of liquid known to exist beyond our planet. They differ markedly from the Earth's oceans, not only in terms of chemistry, but also in their modest expanse. As a result, Titan is a world whose surface remains heavily dominated by dry land. Remarkably, \citet{dermott1995tidal} were able to deduce this fundamental feature long before detailed surface observations became available. They argued against the presence of extended oceans on the basis that Titan's orbit would have been circularised by the dissipative motions of their tides. This left only one viable hypothesis: a surface where the liquid was confined to sparse, disconnected pockets. In due course, Cassini's radar was able to construct a high resolution map of the surface, by piercing the haze of Saturn's largest moon. These observations vindicated the theoretical predictions in spectacular fashion \citep{stofan2007lakes}. The liquid hydrocarbons on Titan appear to account for little more than one per cent of the total surface area. We are currently faced with the even more daunting task of characterising the surfaces of habitable exoplanets. But one subtlety which appears to have been overlooked is that the prediction of \citet{dermott1995tidal} could have been made \emph{even without the orbital data}. On a purely statistical basis, and in the absence of correlations, one expects the division of liquid and solid surface areas to be highly \emph{asymmetric}. This is because the volume of liquid need not match the capacity of perturbations in the solid. The two quantities often differ by several orders of magnitude. If it is the liquid which dominates, the solid surface becomes completely immersed. Enceladus and Europa offer exemplary cases of this phenomenon. Beneath each of their icy crusts, a single ocean completely envelops a solid core \citep{kivelson2000galileo, waite2009liquid}. If, on the other hand, the liquid's volume is subdominant, it settles into small disconnected regions, as was found to be the case on the surface of Titan. Does this trend of asymmetric surface partitions extend to habitable exoplanets? And if so, why do we observe the Earth's water and land areas to be so finely balanced, differing in extent by only a factor of two? These are the core questions we shall aim to address in this work. Simulations of terrestrial planet formation provide us with the first clues for solving these puzzles. \cite{raymond2007high} explored the viability of delivering water to habitable planets from icy planetesimals which originate beyond the snow line. The chaotic nature of this process ensures habitable planets garner a broad spectrum of water compositions. This variety reinforces our expectation that their surfaces tend to be dominated by either solid or liquid. However not all water will reside on a planet's surface. Some will remain locked in the mantle, while a further portion will be lost through the upper atmosphere. Indeed a number of processes can influence the depths of the oceans \citep{1985Schubert, mcgovern1989thermal, kasting1992oceans, holm1996hydrothermal, 2012Abbot, cowan2014water}. If sufficiently strong feedback mechanisms are at work, it may be possible to ensure that the depths of oceans match the amplitude of perturbations in the crust. In which case, we ought to expect many habitable planets to resemble the Earth's division of land and sea. However it remains unclear if any are strong enough to correct for variations in water volume of more than one order of magnitude. Alternatively, habitable planets display a broad distribution of surface conditions, and for the case of the Earth we just `got lucky' \citep{cowan2014water}. But, given that trillions of dice have been rolled, do we require any luck at all? Perhaps the dice were weighted in favour of a balanced surface. The earliest applications of anthropic selection were of a binary nature, in that they addressed the question of whether a particular set of conditions forbade our existence \citep{1974IAUS...63..291C, carter1983anthropic, barrow1986anthropic, 1987PhRvL..59.2607W}. Later, more refined studies invoked Bayesian statistics to deliver quantitative assessments of how our cosmic environment may be biassed by our existence \citep{1995MNRAS.274L..73E, 2004Garriga, tegmark2006dimensionless, 2007PeacockAnthropic, simpson2016habitable, 2015SimpsonAliens}. Until recently, the application of Bayesian anthropic reasoning was restricted to the cosmological realm. The hypothetical ensemble of cosmic conditions has a number of theoretical motivations, yet any experimental evidence lies tantalisingly beyond our grasp. No such limitations exist for the ensemble of habitable planets. \citet{2015SimpsonAliens} used a simple population model to argue that our planet is likely to be towards the large end of the spectrum,\footnote{See also the pedagogical animation by MinutePhysics: \\ \url{https://youtu.be/KRGca_Ya6OM}} inferring the radius $R$ of a given planet with intelligent life to be $R < 1.2 R_\oplus$ (95\% confidence bound). Empirical analyses by \citet{2015Rogers} and \citet{2016ChenKipping} appear to support these findings, with the latter study concluding that the Terran-Neptunian divide occurs at approximately $1.2 R_\oplus$. Whether its the multiverse, extra-terrestrial life, or even the longevity of our species \citep{gott1993implications, 2016SimpsonDoomsday}, putting this predictive framework to the test is rarely practical. Yet the characterisation of habitable exoplanets provides a remarkable opportunity to do just that. In this work we turn our attention to the selection effect involving a planet's ocean coverage. Our understanding of the development of life may be far from complete, but it is not so dire that we cannot \emph{drastically} improve on the implicit approximation that all habitable planets have an equal chance of hosting intelligent life. Should we consider planets with different land-ocean divides to have an equal chance of producing an intelligent species such as \emph{Homo Sapiens}? Few would doubt whether the Earth's surface configuration is better suited to supporting a diverse biosphere than \emph{Tatooine}. It is this small piece of knowledge that can be exploited to update our prior belief for the surface conditions among the ensemble of habitable planets. In \S \ref{sec:tuning} we explore the fine-tuning problem associated with the Earth's oceans, and review two approaches for tackling the problem: feedback processes and observational selection effects. In \S \ref{sec:fecund} we quantify the relative probability of observing a host planet based on its habitable area. The model we use for the ensemble of surface conditions is defined in \S \ref{sec:model}. Our main results are presented in \S \ref{sec:results}, before concluding with a discussion in \S \ref{sec:conclusions}. | \label{sec:conclusions} On a purely statistical basis, one na\"{\i}vely expects to find a highly asymmetric division of land and ocean surface areas. A natural explanation for the Earth's equitably partitioned surface is an evolutionary selection effect. We have highlighted two mechanisms which could be responsible for driving this selection effect. First of all, planets with highly asymmetric surfaces (desert worlds or waterworlds) are likely to produce intelligent land-based species at a much lower rate. Secondly, planets with larger habitable areas are capable of sustaining larger populations. Both of these factors imply that our host planet has a greater habitable area than most life-bearing worlds. We have exploited this model of planetary fecundity to draw two major conclusions. First of all, we find that the Earth's oceanic area provides substantial evidence in favour of the selection model. Secondly, in the context of this model, we find that most habitable planets have surfaces which are over $90\%$ water ($95\%$ credible interval). Our results are robust to a broad variety of modifications to the model. The only critical assumption is that there is a significant variance in the basin saturation among habitable worlds, specifically $\sigmaX/\Vnorm \gtrsim 0.5$. This appears likely given that there are many variable factors which contribute to a planet's surface water volume and basin capacity. The anticipated prevalence of waterworlds is driven by the fact that our home planet is close to the waterworld limit. Such proximity to a critical limit is precisely what one expects to find in the presence of a selection effect, provided only a smooth tail of the distribution lies below the critical limit. This reasoning was previously exploited by \citet{1987PhRvL..59.2607W} to successfully predict the value of the cosmological constant. If the Earth's basin saturation is biased low, this implies that (a) its water mass fraction is likely to be biased low and (b) its elevation amplitude is likely to be biased high (and as with the basin saturation, the magnitude of this bias will depend on the planet-to-planet variance of these quantities). Do these two scenarios appear feasible? The water mass fraction among habitable planets could be considerably higher than the Earth. For example, numerical simulations based on delivering water from planetary embryos found a median water mass fractions of approximately $1\%$ \citep{raymond2007high}, ten times higher than the terrestrial value. Extremely elevated water compositions have been associated with the inflation of planetary radii \citep{2016Thomas}. This scenario, in which the Earth is among the driest habitable planets, could help explain the appearance of a low-mass transition in the mass-radius relation of exoplanets \citep{2015Rogers, 2016ChenKipping}. If it transpires that the Earth is indeed unusually dry for a habitable planet, then one might wonder what the mechanism was. Does the Solar System have some distinguishing feature that was responsible? For example, perhaps the low eccentricities and inclinations of solar system planets are inefficient at promoting water delivery. Another possibility could be the influence of the Grand Tack model, where Jupiter underwent a reversal of its migration \citep{walsh2011low}. This has been found to yield a delivery of water that is approximately consistent with terrestrial levels \citep{2014Icar..239...74O}. However recent simulations of the Grand Tack scenario suggest that, if anything, this may enhance the delivery of water to terrestrial planets \citep{2016Matsumura}, rather than curtail it. Alternatively, a dry Earth may not necessarily have arisen from an identifiable macroscopic feature, it could simply be associated with the inherently stochastic nature of the water delivery process. It also appears feasible that the Earth has an unusually deep ocean basin. The gravitational potential associated with its surface fluctuations is much higher than any other body in the solar system. In turn this may suggest the Earth has unusually strong tectonic activity, and consequentially, an abnormally strong magnetic field. This exemplifies how selection effects can easily be transferred to correlated variables. Could the planet-to-planet variability in $\Vnorm$ be very small? Feedback mechanisms may have acted to regulate the depths of planetary oceans relative to the magnitude of their surface perturbations \citep{2012Abbot}. Earlier we denoted this possibility $\mathcal{H}_2$. Fortunately this hypothesis leads to a very different forecast for the surface conditions of Earth-like planets. If $\mathcal{H}_2$ is correct, we shall discover that a substantial proportion of habitable planets share the Earth's equitable water-land divide. This is in stark contrast to the prediction of our selection model, based on $\mathcal{H}_1$, where habitable planets are dominated by oceans. Other aspects of the Earth's surface that are susceptible to selection effects include the spatial configuration of land. For example, if a planet's land area were retained in a single contiguous piece, akin to Pangea, it may be that either a larger proportion of the land is rendered uninhabitable, or the ecological diversity is significantly suppressed. The Earth's land configuration may be optimised to ensure that the majority of the available area is habitable, thereby maximising its fecundity, as defined in (\ref{eq:bayeshabitable}). This work builds on \citet{2015SimpsonAliens} by providing a further demonstration of why the Earth is likely to appear as a statistical outlier, across a broad spectrum of physical properties, when compared to other life-bearing worlds. In general, if a planet's population size is correlated with any variable, then the mean value witnessed by individuals will always exceed the true mean. This is true for \emph{any} distribution of population sizes (see Appendix \ref{sec:proof}). Ordinarily a single random sample is not particularly helpful in informing us on the nature of a broad population. However this limitation only applies to \emph{fair} samples. A single \emph{biased} sample can be used to place a lower (or upper) bound on the entire population distribution. For example, if the only data point we had regarding human running speed was taken from an Olympic 100 metres final, then we can be confident that a subsequent fair sample, across the global population, would not be significantly faster. Provided the population variance is significant, than we can be confident in finding a substantial deviation between the fair sample and the biased sample. To give a further pedagogical example: imagine that you look at a kitchen worktop and notice some spilled coffee granules. One of those granules, selected at random, is found to lie within $0.1mm$ from the edge of the $600mm$ worktop. This proximity could of course be entirely coincidental. But it is much more likely that the bulk of the granules fell on the floor, and what you are seeing is merely the tail end of the distribution. The fine-tuning of the Earth's parameters is closely related to the proposition that various cosmological parameters correspond to those which optimise star formation \citep{tegmark2006dimensionless}. The key difference here is that many elements in the planetary ensemble are observable, and thus our predictions are experimentally falsifiable. Indeed, it may not be long before we begin to build a census of nearby habitable planets, and begin to develop an understanding of how the Earth compares to other habitable worlds \citep{catala2009plato}. If habitable planets systematically differ from the Earth in some way - such as the ocean coverage discussed in this work - this provides a hint as to the conditions which favoured the development of intelligent life. It would show that there is a bias between the inner sets of Figure \ref{fig:sets}. This bias tells us something about why we evolved on this particular lump of rock. It has been argued that the finely-tuned properties of our planet is indicative of the sparsity of life in the Universe - the so-called `Rare Earth hypothesis'. However this interpretation overlooks one of the key factors which control the selection effect: the number of observers produced by each planet. The conditions on an individual's home planet is heavily skewed in favour of those conditions which maximise the \emph{abundance} of life. As an analogy, consider the contiguous piece of dry land you live on. It \emph{is} extremely special, in the sense that it is one of the largest pieces of contiguous land on the Earth's surface. But at the same time, there are hundreds of thousands of smaller chunks of land scattered across the Earth's surface. The selection effect that takes place when studying the ground beneath your feet is not a fair one. Likewise, the rarity of the Earth's parameters need not reflect the sparsity of life in the cosmos. On the contrary, it may be driven precisely because we are a small piece within a vast ensemble. When physiologists seek a deeper understanding of our body's features, such as our eyes and ears, a great deal of progress can be made from laboratory experimentation. Yet the \emph{only} way to arrive at a comprehensive answer is by including a complementary analysis of our origins. This allows biological function to be placed in an evolutionary context. A similar statement can be made regarding the features of our planet. No matter how formidable our understanding of planet formation becomes, one can never hope to fully appreciate the Earth's features without addressing the issue of how we came into being upon it. | 16 | 7 | 1607.03095 |
1607 | 1607.01090_arXiv.txt | We have measured an annual parallax of the Mira variable R~Ursae~Majoris (R~UMa) with the VLBI exploration for Radio Astronomy (VERA). From the monitoring VLBI observations spanning about two years, we detected H$_2$O maser spots in the LSR velocities ranges from 37 to 42 km\,s$^{-1}$. We derived an annual parallax of 1.97$\pm$0.05\,mas, and it gives a corresponding distance of 508$\pm$13\,pc. The VLBI maps revealed 72 maser spots distributed in $\sim$110 au area around an expected stellar position. Circumstellar kinematics of the maser spots were also revealed by subtracting a systemic motion in the Hipparcos catalog from proper motions of each maser spots derived from our VLBI observations. Infrared photometry is also conducted to measure a $K$ band apparent magnitude, and we obtained a mean magnitude of $m_K$ = 1.19$\pm$0.02\,mag. Using the trigonometric distance, the $m_K$ is converted to a $K$ band absolute magnitude of $M_K = -$7.34$\pm$0.06\,mag. This result gives a much more accurate absolute magnitude of R~UMa than previously provided. We solved a zero-point of $M_K - \log P$ relation for the Galactic Mira variables and obtained a relation of $M_K = -$3.52 $\log P$ + (1.09 $\pm$ 0.14). Other long period variables including red supergiants, whose distances were determined from astrometric VLBI, were also compiled to explore the different sequences of $M_K - \log P$ relation. | Mira variables and Long Period Variables (LPVs) are low- to intermediate-mass ($1- 8 M_{\odot}$) asymptotic giant branch (AGB) stars that pulsate with a period range of 100 $-$ 1000 days. They are surrounded by large and extended dust and molecular shells, and in sources with mass-loss rate higher than $\sim 10^{-7} M_{\odot}$yr$^{-1}$, we sometimes find maser emissions of H$_2$O, SiO, or OH \citep{gai14, hab03}. Because of their high mass-loss rate, they are also important source to study chemical composition of the universe. Another characteristic of the sources is concerned to their periodic variation. The relation between K band apparent magnitude ($m_K$) and logarithm of pulsation period ($\log P$) of Mira variables is well known in the Large Magellanic Cloud (LMC) \citep{fea89,ita04-1}. If the LMC distance is given, the relation can be converted to a relation of absolute magnitude ($M_K$) and $\log P$, then it can be used as a distance indicator. Since there is a metallicity difference between LMC and our galaxy, it is also important to establish this $M_K-\log P$ relation using sources in our own galaxy. Since the LPVs are very bright in infrared, we can use them to probe a region where interstellar extinction is strong such as the direction of the Galactic Center and Galactic plane. However, a construction of the $M_K-\log P$ relation for the Galactic Mira variables has long been difficult because of large errors in $M_K$ due to distance uncertainty. Making use of a high performance of the VERA array \citep{kob03}, which is a Japanese VLBI project dedicated to the Galactic astrometry, we aim to construct the $M_K-\log P$ relation for Galactic LPVs. R Ursae Majoris (R~UMa) is an O-rich Mira variable \citep{kna00} with a pulsation period of 301.6 days (GCVS)\footnote[1]{General Catalog of Variable Stars\\http://heasarc.gsfc.nasa.gov/W3Browse/all/gcvs.html}. The H$_2$O masers associated with R~UMa also exhibit regular periodic variation with a phase lag of $\sim$0.3 with respect to the optical light curve \citep{shi08}. Although the Hipparcos satellite has measured the annual parallax of 2.37$\pm$1.06 mas \citep{van07}, its corresponding distance of 422$^{+341}_{-130}$ pc still has a large error and it brings an uncertainty to the absolute magnitude. In this study, we observe this source with VERA to obtain more accurate distance based on an precise astrometry of H$_2$O maser positions. A high resolution VLBI map gives an angular distribution of maser spots around the star. Then, a time series of multiple VLBI observations can add kinematic information to the masers. This helps us to understand a global picture of circumstellar medium on the schemes of time and space. In case the maser distribution is relatively isotropic like a Mira variable `T~Lep' in our previous study \citep{nak14}, a circumstellar kinematics can be derived from an analysis only using VLBI maps. The detailed procedure is given in section 3.4 in \citet{nak14}. On the other hand, if the number of the maser spots is small or the distribution is far from isotropic, the same procedure can not offer a reliable kinematic picture of the maser spots. In this study, we propose a new method to obtain a kinematics of the circumstellar masers in R~UMa. We solve the $M_K-\log P$ relation of the Galactic Mira variables based on the latest results from astrometric VLBI observations. Infrared photometry is also used to measure an apparent magnitude in $K$ band. | \subsection{Internal motion of circumstellar masers} \label{sec_internalmotion} If we can directly detect a central star with our VLBI method, it is easy to know internal motions of maser spots with respect to the central star. But it is difficult, and we have to introduce some reasonable assumption in order to estimate internal motions of the maser spots. For example, an isotropy of circumstellar kinematics was assumed in our previous study \citep{nak14}. In this section, we try to reveal the internal motions of the maser spots with a method using an astrometric measurement from Hipparcos satellite~\citep{per97}. Proper motions of each maser spot measured in our VLBI observations ({\boldmath $\mu$}$^{\mathrm{VERA}}$) involve various kinematics, such as the Galactic rotation, a systemic motion of the star, and their internal motions. A proper motion of R~UMa measured by Hipparcos ({\boldmath $\mu$}$^{\mathrm{HIP}}$) also include the same kinematics as {\boldmath $\mu$}$^{\mathrm{VERA}}$ except for internal motions of the maser spots. Therefore, a remainder of two measurements ({\boldmath $\mu$}$^{\mathrm{VERA}}$ $-$ {\boldmath $\mu$}$^{\mathrm{HIP}}$) should give internal motions of the maser spots on the rest frame fixed to the central star. In the revised Hipparcos catalog~\citep{van07}, the proper motion of R~UMa is reported to be {\boldmath $\mu$}$^{\mathrm{HIP}}$ $=(-40.51\pm0.79, -22.66\pm0.78)$ mas\,yr$^{-1}$. For maser spots detected more than two continuous epochs, we estimated their proper motions. Since the parallax of 1.97 mas determined in section~\ref{sec_parallax} is the same for all maser spots, we used the fixed parallax and re-fitted all the individual maser spots solving only for internal motion. By subtracting the {\boldmath $\mu$}$^{\mathrm{HIP}}$ from the proper motions of each maser spot, we obtained internal motions $\mu^{\mathrm{int}}_{\mathrm{x}}$ and $\mu^{\mathrm{int}}_{\mathrm{y}}$ for 38 out of all 72 maser spots, and presented them in table~\ref{maser_table} in unit of mas\,yr$^{-1}$. We also presented errors of the internal motions $\sigma^{\mathrm{int}}_{\mu\mathrm{x}}$ and $\sigma^{\mathrm{int}}_{\mu\mathrm{y}}$ on the table which are quadratic sums of two errors from Hipparcos catalog and our VLBI observation. Since measurement accuracy of proper motions in our VLBI observation is, on average, 0.2\,mas\,yr$^{-1}$, resultant errors are almost dominated by the error from Hipparcos measurement. A proper motion of R~UMa system can also be estimated from our VLBI observations. We averaged out all the proper motions of maser spots and obtain a motion of $-40.77\pm0.39$ mas\,yr$^{-1}$ and $-24.75\pm0.38$ mas\,yr$^{-1}$ in RA and DEC, respectively. In the RA, two proper motions from Hipparcos and VERA are consistent within their errors. In the DEC, however, there is a difference of $\sim$2 mas\,yr$^{-1}$ between two measurements. As shown in figure~\ref{fig_internalmotion}, we found unisotropic distribution of maser spots. We think the systemic motion derived from VERA possibly be biased due to this unisotropy. Therefore, in this section, we used the proper motion in Hipparcos catalog to inspect internal motions of the maser spots. In figure~\ref{fig_internalmotion}, internal motions of maser spots on sky plane are indicated with arrows. Proper motion of 1 mas\,yr$^{-1}$ corresponds to a transverse velocity of 2.41 km\,s$^{-1}$ at the source distance of 508 pc. An arrow at bottom right shows a proper motion of 3 mas\,yr$^{-1}$ ($=$ 7.22km\,s$^{-1}$). Average errors of internal motions are 0.90 mas\,yr$^{-1}$ and 0.88 mas\,yr$^{-1}$ in RA and DEC, respectively. Most $\mu^{\mathrm{int}}_{\mathrm{y}}$ show essentially larger velocities than their errors $\sigma^{\mathrm{int}}_{\mu\mathrm{y}}$, and therefore, we think our analysis gives an reliable picture of the internal motions especially along DEC axis. Maser spots at the southern area of the map show southward motions, and on the contrary, two spots at the northern area show northward motions. With regards to the RA axis, $\mu^{\mathrm{int}}_{\mathrm{x}}$ values are same as their errors, and this brings a large uncertainty to directions about the internal motions along RA axis. Nevertheless, we can conclude that there is no remarkable systemic motion along RA axis enough to be detected with this method. As a result, our analysis of internal motion revealed an outward motions with respect to the central region of the map. A root sum square of internal motions along two axes $( = \sqrt{ (\mu^{\mathrm{int}}_{\mathrm{x}})^2 + (\mu^{\mathrm{int}}_{\mathrm{y}})^2}\,)$ gives a transverse velocities on the skyplane. An average of the transverse velocities was obtained to be 6.61 km\,s$^{-1}$. In the next section, we will give a more detailed study to capture a global picture of the maser spots and the central star. Finally in this subsection, we mention a binarity of the R~UMa system. There are some AGB samples known as a binary stars, i.e., Mira AB system \citep{kar97}, R~Aqr \citep{wil81}, and so on. The R~UMa is one of possible samples of binary system. \citet{sah08} studied fluxes of some AGB stars at near-UV, far-UV, optical, and near-IR bands. They concluded that the far-UV excess likely results either directly from the presence of a hot binary companion or indirectly from a hot accretion disk around the companion. If the R~UMa forms a binary system and the binarity affects internal motions of maser spots, astrometric measurements of the maser spots include the binary motion. Then, a remainder {\boldmath $\mu$}$^{\mathrm{VERA}}$ $-$ {\boldmath $\mu$}$^{\mathrm{HIP}}$ still have some contribution from the binary motion. However, as presented above, we could not detect any systematic motion for all maser spots but detected expanding like motions. This can implies that there is no large binary effect that seriously disturb our consideration of internal motion. Therefore, we have not considered the binarity in this section. \subsection{Stellar position and 3D-picture of the maser spots} \label{sec_stellarposition} Based on the distribution and internal motions of the maser spots revealed in previous sections, we will discuss the position of the central star. Now we focus on the motions of the maser spots along DEC axis. In figure~\ref{fig_yvsvy}, we show a relation between $\mu^{\mathrm{int}}_{\mathrm{y}}$ and $Y$ of 38 maser spots. Since there seems to be a gradient of $\mu^{\mathrm{int}}_{\mathrm{y}}$, we fitted this data to a linear function. Then we obtained a relation of $\mu^{\mathrm{int}}_{\mathrm{y}} = 0.062 Y - 8.758$, which is presented with a solid line in figure~\ref{fig_yvsvy}. For $V_{\mathrm{y}} = 0$, this relation gives a $Y = 140.45$, and this $Y$ value can be considered as a possible $Y$ position of the central star. In figure~\ref{fig_internalmotion}, this $Y$ value is presented with a solid horizontal line with peripheral region (gray colored) indicating its error. Along the RA axis, it is difficult to find a gradient, we could not conduct the same analysis as the DEC axis. Independently, we tried to estimate a stellar position using the angular distribution of maser spots. A cross on the map indicates an estimated stellar position of $(X, Y) = (-12.1, 136.8)$ mas which was obtained by simply calculating medians of two ends in RA and DEC axes. Assuming this position, we calculated angular distances $\theta$ of the maser spots from the central star. Here, we introduce a simple uniform expanding shell model \citep{oln77,rei77} of \begin{eqnarray} \left(\frac{\theta}{\theta_\mathrm{m}}\right)^2 + \left(\frac{V_\mathrm{LSR}-V_\ast}{V_\mathrm{exp}}\right)^2 = 1, \nonumber \end{eqnarray} where $\theta$ is an angular distance from the central star, $\theta_\mathrm{m}$ is the shell radius, $V_\ast$ is the radial velocity of the star, and $V_\mathrm{exp}$ is an expansion velocity of the shell. Figure~\ref{fig_pv} shows a $\theta$ vs $(V_\mathrm{LSR}-V_\ast)$ diagram of the 38 maser spots. A shape of the model on this figure shows an ellipse. We assumed a $V_\ast$ of 40.49 km\,s$^{-1}$ which is obtained as a center of $V_\mathrm{LSR}$ values of maser spots. Since a distribution of the data on the $\theta - (V_\mathrm{LSR}-V_\ast)$ plane is too narrow to numerically solve suitable $V_\mathrm{exp}$ and $\theta_\mathrm{m}$, we assumed a $V_\mathrm{exp}$ to be 6.6 km\,s$^{-1}$ which is an average of transverse velocities of the maser spots obtained in previous section. Using all maser spots, we solved a $\theta_\mathrm{m}$ to be 85$\pm$2 mas, which is presented with solid line in figure~\ref{fig_pv}. For comparison, we also presented two models whose outer radius of 120 mas (dashed--line) and an inner radius of 65 mas (one--dotted chain line). Latter two models are not obtained from numerical fitting but we assumed fixed radii. In figure~\ref{fig_internalmotion}, the shell model with a radius of $\theta_\mathrm{m}=$ 85 mas is presented as a dotted circle whose center is assumed to be the cross symbol. \begin{figure} \begin{center} \includegraphics[width=80mm, angle=0]{fig07_rev01.eps} \end{center} \caption{ Velocity gradient of $V_{\mathrm{y}}$ along DEC ($Y$) axis. An intersection point of $V_{\mathrm{y}} = 0$ and the slope gives $Y = 140.45$ mas. } \label{fig_yvsvy} \end{figure} \begin{figure} \begin{center} \includegraphics[width=85mm, angle=0]{fig08.eps} \end{center} \caption{ Relation between an angular radius $\theta$ and $V_\mathrm{LSR}-V_\ast$ of the maser spots. A solid ellipse indicates an uniform shell model with a radius of 85 mas obtained from a numerical fitting. Dashed and one-dotted-chain lines indicate models with shell radii of 120 and 65 mas, respectively. Expansion velocities of $V_\mathrm{exp} =$ 6.6 km\,s$^{-1}$ are fixed for all models. } \label{fig_pv} \end{figure} By revealing accurate motion of circumstellar matter, properties of stellar wind will be inspected and it helps us to understand mass-loss process. However, the distributions of circumstellar matter sometimes show symmetry, and sometimes show asymmetry. A conventional method for determining the circumstellar motion based on the VLBI data, which assumes uniformity or symmetry for circumstellar matters, intrinsically has limitation to derive the real picture. Therefore, it is important to create different method which does not need any assumption about symmetry or uniformity for circumstellar matter. A capability of the new method, which combines VLBI and other independent astrometric data, was presented here. To acquire more accurate picture of internal motions using the same method, accurate systemic motions of stars are required. The Gaia satellite~ \footnote[1]{Gaia Mission, ESA; http://sci.esa.int/gaia/}, launched in 2013, is an ongoing astrometric program which expected to measure accurate proper motions for large amount of stars. For very bright stars, like nearby Mira variables, a Japanese satellite Nano--JASMINE~ \footnote[2]{Nano - JASMINE, NAOJ ;\\ http://www.jasmine-galaxy.org/index-en.html} will also be a powerful and promising telescope to determine their proper motions. In near future, more accurate proper motions of stars measured by new satellites will be tied up with VLBI measurements of maser spots, and circumstellar dynamics of many sources can be studied with the same method in this study. We expect a successful launch of Nano--JASMINE scheduled in near future. \subsection{R~UMa in the $M_K-$log$P$ Diagram} \label{sec_plr} For AGB stars in the LMC, it is known that there are several distinct sequences on the $m_K-$log$P$ diagram \citep{woo99, ita04-1}, where $m_K$ and $P$ indicate their $K$ band apparent magnitude and pulsation period. Sequences C and C' in their classification correspond to groups of variables pulsating in a fundamental tone and a first-overtone, respectively. Since distances of the Galactic sources can not be treated as identical like sources in the LMC, the same relation should be studied in $M_K-$log$P$ plane, where $M_K$ is a $K$ band absolute magnitude. In this section, we confirm a location of R~UMa on the $M_K-$log$P$ plane, and solve a zero-point of the Galactic $M_K-$log$P$ relation based on recent astrometric results. During the last decade, we have conducted astrometric observations for Mira and semiregular variables using VERA \citep{nak08}. Determination of $M_K = -8.33 \pm 0.10$ mag of U~Lyn by \citet{kam15} is the latest published result of Galactic Mira variable based on a parallax from our ongoing program. In this study, using our trigonometric distance of 508$\pm$13 pc and the infrared magnitude of $m_K = 1.19\pm0.02$ mag, we derived a $M_K = -7.34\pm0.06$\,mag for R~UMa. The error is determined as a root mean square of distance based error and apparent magnitude error. Adding this new source, we summarized the sources in table~\ref{parallax_table} in an increasing order of pulsation period together with all Galactic LPVs whose distances are determined with astrometric VLBI. There are samples of Mira, semiregular (SRa, SRb in table~\ref{parallax_table}), and red supergiants (SRc in table~\ref{parallax_table}). Semiregular variables show amplitude smaller than Miras, and among them, stars with better periodicities are referred to as SRa to distinguish from SRb with poorly defined periodicities. Species of the masers used in the parallax measurements are also shown. References of the parallaxes and apparent magnitudes $m_K$ are given in the footnote to table~\ref{parallax_table}. Using the distance and $m_K$, we derived absolute magnitudes $M_K$ of the sources. In estimation of $M_K$ errors, only the distance errors were considered. Now, we define a $M_K-$log$P$ relation in the form of $M_K = -3.52 \, \mathrm{log}P + \delta$, where we assume a fixed slope of $-3.52$ determined by \citet{ita04-1} using LPVs in the LMC. Using Miras and four semiregular variables (RW~Lep, S~Crt, RX~Boo, and W~Hya) in table~\ref{parallax_table}, we solved the constant $\delta$. Unweighted and weighted least squares fitting to the data gives $\delta$ of 1.09$\pm$0.14 and 1.45$\pm$0.07, respectively. In \citet{kam12}, various periods of a well-studied semiregular source RX~Boo are reported, and ratios of the periods are found to be around two. In order to include all Mira and semiregular variables in the fitting, the periods of semiregular variables are multiplied by two and used. Red supergiants are not included in the fitting. Since $M_K$ errors of a few sources are quite small compared to other source, it gives a $\delta$ discrepancy of 0.36 between two fittings. In figure~\ref{plr}, we presented all sources in table~\ref{parallax_table} on $M_K-$log$P$ plane with red squares. To distinguish the current result of R~UMa from published ones, R~UMa is presented with an open square. Two solid lines indicate the results from unweighted (upper) and weighted (lower) fitting, respectively. And, two dashed lines indicate relations derived by \citet{ita04-1} for sequence C' (first-overtone) and $C$ (fundamental tone), respectively. The LPVs in LMC reported by \citet{ita04-1} are presented with small dots in a shaded area of the figure. We can find that R~UMa falls on the sequence C, and it is likely that the star pulsates in a fundamental mode. To calibrate absolute magnitudes of the sources in LMC, we assumed a distance modulus of 18.49 \citep{van07} for LMC. Within a determination accuracy of $\delta$ in our study, we find a consistency of the relations for Mira variables between our Galaxy and the LMC. In addition to Miras and semiregular variables, now we focus on four sources (3 red supergiants and NML~Cyg) in table~\ref{parallax_table}. They are $\sim$3 magnitudes brighter than Mira and semiregular variables, and show longer pulsation periods of 822 ($\log P = 2.915$) to 1280 ($\log P = 3.107$) days. If we extrapolate C and C' sequences of \citet{ita04-1} to brighter and longer period region, we find that three sources fall near C' and the other falls near C sequence, that possibly indicating a difference of the pulsation mode. Since the number of the Galactic red super giants with parallactic distance is small, it is important to increase the sample number for better calibration of their $M_K-$log$P$ relation. The VLBI astrometry can be a useful method to provide their accurate distances. \begin{figure} \begin{center} \includegraphics[width=85mm, angle=0]{fig09_rev01.eps} \end{center} \caption{ Absolute magnitudes ($M_K$) $-$ $\log P$ diagram of the Galactic long period variables. Filled red squares indicate sources in table~\ref{parallax_table}, whose distances are derived from astrometric VLBI. Only the result of R~UMa is presented with an open square. Solid lines show fitting results of $M_K - \log P$ relations for Galactic Mira variables in table~\ref{parallax_table}. See section~\ref{sec_plr} for detail of the fitting procedure. In a shaded area, LPVs in the LMC are presented with small dots \citep{ita04-1}. Labels C and C' and corresponding dashed lines indicate sequences in \citet{ita04-1}. } \label{plr} \end{figure} \begin{table*}[!tb] \caption{Results from VLBI astrometry} \label{parallax_table} \begin{center} \begin{tabular}{lcccccccc} \hline Source & Type & Parallax & $P$ & $\mathrm{Log}P$ & $m_K$ & $M_K$ &Maser & Reference\footnotemark[$\dag$] \\ & & [mas] &[day]& & [mag] & [mag] & & (Parallax, $m_K$) \\ \hline \hline RW~Lep & SRa & 1.62$\pm$0.16 & 150 & 2.176 & 0.639 & $-8.31\pm0.22$ &H$_2$O& kam14, a\\ S~Crt & SRb & 2.33$\pm$0.13 & 155 & 2.190 & 0.786& $-7.38\pm0.12$ &H$_2$O& nak08, a\\ RX~Boo & SRb & 7.31$\pm$0.5 & 162 & 2.210 & $-$1.96 & $-7.64\pm0.15$ &H$_2$O& kam12, b\\ R~UMa & Mira & 1.97$\pm$0.05 & 302 & 2.480 & 1.19 & $-7.34\pm0.06$ &H$_2$O& $\cdots$ \\ W~Hya & SRa &10.18$\pm$2.36 & 361 & 2.558 & $-$3.16 & $-8.12\pm0.51$ &OH & vle03, c\\ S~CrB & Mira & 2.39$\pm$0.17 & 360 & 2.556 & 0.21 & $-7.90\pm0.15$ &OH & vle07, c\\ T~Lep & Mira & 3.06$\pm$0.04 & 368 & 2.566 & 0.12& $-7.45\pm0.03$ &H$_2$O& nak14, c\\ R~Aqr & Mira & 4.7$\pm$0.8 & 390 & 2.591 & $-$1.01& $-7.65\pm0.37$ &SiO & kam10, c\\ R~Aqr & Mira & 4.59$\pm$0.24 & 390 & 2.591 & $-$1.01& $-7.70\pm0.11$ &SiO & min14, c\\ RR~Aql & Mira & 1.58$\pm$0.40 & 396 & 2.598 & 0.46 & $-8.55\pm0.56$ &OH & vle07, c\\ U~Her & Mira & 3.76$\pm$0.27 & 406 & 2.609 & $-$0.27 & $-7.39\pm0.16$ &OH & vle07, c\\ SY~Scl & Mira & 0.75$\pm$0.03 & 411 & 2.614 & 2.55& $-8.07\pm0.09$ &H$_2$O& nyu11, b\\ R~Cas & Mira & 5.67$\pm$1.95 & 430 & 2.633 & $-$1.80 & $-8.03\pm0.78$ &OH & vle03, c\\ U~Lyn & Mira & 1.27$\pm$0.06 & 434 & 2.637 & 1.533 & $-7.95\pm0.10$ &H$_2$O& kam15, a\\ UX~Cyg & Mira & 0.54$\pm$0.06 & 565 & 2.752 & 1.40 & $-9.94\pm0.24$ &H$_2$O& kur05, a\\ S~Per & SRc & 0.413$\pm$0.017 & 822 & 2.915 & 1.33 & $-10.59\pm0.09$ &H$_2$O& asa10, b\\ PZ~Cas & SRc & 0.356$\pm$0.026 & 925 & 2.966 & 1.00 & $-11.24\pm0.16$ &H$_2$O& kus13, b\\ VY~CMa & SRc & 0.88$\pm$0.08 & 956 & 2.980 & $-$0.72 & $-11.00\pm0.20$ &H$_2$O& cho08, b\\ NML~Cyg & --- & 0.62$\pm$0.047 & 1280 & 3.107 & 0.791 & $-10.25\pm0.16$&H$_2$O & zha12, a\\ \hline \hline \multicolumn{9}{@{}l@{}}{\hbox to 0pt{\parbox{145mm}{\footnotesize \smallskip \par\noindent \footnotemark [$\dag$] References of the parallax are as follows : (kam14) \cite{kam14}, (nak08) \cite{nak08}, (kam12) \cite{kam12}, (vle03) \cite{vle03}, (vle07) \cite{vle07}, (nak14) \cite{nak14}, (kam10) \cite{kam10}, (min14) \cite{min14}, (nyu11) \cite{nyu11}, (kam15) \cite{kam15}, (kur05) \cite{kur05}, (asa10) \cite{asa10}, (kus13) \cite{kus13}, (cho08) \cite{cho08}, and (zha12) \cite{zha12}. References of the apparent magnitudes ($m_K$) are as follows : (a) The IRSA 2MASS All-Sky Point Source Catalog \citep{cut03}, (b) Catalogue of Stellar Photometry in Johnson's 11-color system \citep{duc02}, (c) Photometry by \citet{whi00}. }\hss}} \end{tabular} \end{center} \end{table*} \subsection{Conclusion} \label{sec_concl} We conducted astrometric VLBI observation of a Mira variable R~UMa. Positions of H$_2$O maser spots at 22\,GHz were measured using the VERA array. Obtained parallax of 1.97$\pm$0.05 mas gives a distance of 508$\pm$13 pc. Circumstellar kinematics of 38 maser spots and angular distribution of all 72 maser spots were revealed from our observations, then we gave a constraint on a stellar position. The $M_K - \log P$ relations for the Galactic LPVs, whose distances were measured from astrometric VLBI, were explored and we obtained the relations of $M_K = -3.52 \, \mathrm{log}P + (1.09\pm0.14)$ and $M_K = -3.52 \, \mathrm{log}P + (1.45\pm0.07)$ from unweighted- and weighted-least squares fittings, respectively. R~UMa was found to fall on C sequence that pulsate in a fundamental tone. Positions of red supergiants in the same $M_K - \log P$ plane also found to fall on C' and C sequences. From this we can infer the pulsation mode of the sources, and also this represents an capability of VLBI astrometry for a calibration of $M_K - \log P$ relation applicable to red supergiants. | 16 | 7 | 1607.01090 |
1607 | 1607.08178_arXiv.txt | The power spectrum of the X-ray fluctuations of accreting black holes often consists of two broad humps. We quantitatively investigate the hypothesis that the lower frequency hump originates from variability in a truncated thin accretion disc, propagating into a large scale-height inner hot flow which, in turn, itself is the origin of the higher frequency hump. We extend the propagating mass accretion rate fluctuations model \textsc{propfluc} to accommodate double hump power spectra in this way. Furthermore, we extend the model to predict the cross-spectrum between two energy bands in addition to their power spectra, allowing us to constrain the model using the observed time lags, which in the model result from both propagation of fluctuations from the disc to the hot flow, and inside the hot flow. We jointly fit soft and hard power spectrum, and the cross-spectrum between the two bands using this model for 5 \textit{Swift X-ray Telescope} observations of MAXI J1659-152. The new double hump model provides a better fit to the data than the old single hump model for most of our observations. The data show only a small phase lag associated with the low frequency hump. We demonstrate quantitatively that this is consistent with the model. We compare the truncation radius measured from our fits with that measured purely by spectral fitting and find agreement within a factor of two. This analysis encompasses the first joint fits of stellar-mass black hole cross-spectra and power spectra with a single self-consistent physical model. | \label{sec:int} Transient black hole X-ray binaries (BHBs) evolve in very characteristic ways during their outbursts (e.g. Belloni et al. 2005; Remillard \& McClintock 2006; Belloni 2010; Gilfanov 2010). A typical BHB outburst passes through a number of different states, each state being defined by particular spectral and timing properties of the source. At the beginning of the outburst, the source is in the low-hard state (LHS): it shows high aperiodic variability ($rms$ $\utilde{>}$ 30\%) and its energy spectrum is dominated by a hard power law component (photon index $\Gamma \approx$ 1.7). As the source luminosity increases, the source moves towards the high-soft state (HSS): the aperiodic variability drops off ($rms \approx$ 3\%), the power law softens ($\Gamma \approx$ 2.4), and the spectrum becomes dominated by a multi-colour blackbody component peaking in soft X-rays ($\approx$ 1 keV). At the end of the outburst, the source hardens again, turning back in the LHS. \\ Looking at the power spectrum of the source during the outburst, it is possible to identify several different components representing rapid variability on time scales between $\approx$ 0.01 and $\approx$ 100 s, which have different characteristics for each state. In particular, the LHS is usually characterized by the presence of a quasi periodic oscillation (QPO) superimposed on broad band continuum noise. During the evolution of the outburst, all the characteristic frequencies of the power spectral components correlate with hardness (e.g. Wijnands \& van der Klis 1998; Psaltis, Belloni \& van der Klis 1999; Homan et al. 2001). The initial transition between LHS and HSS usually takes place through intermediate states with spectral and timing properties in between those of LHS and HSS. For example, after the LHS, the source can enter the hard-intermediate state (HIMS) where its spectrum is characterized by the presence of both a disc and a power law component, the aperiodic variability decreases to $rms \approx$ 10-20\%, and the QPO superimposed on the broad band noise is still present.\\ The transition between LHS and HSS can be explained considering two different emitting regions in the accreting flow interacting with each other: an optically thick disc producing the blackbody emission (Shakura \& Sunyaev 1973), and an optically thin Comptonizing region producing the power law (Thorne \& Price 1975; Sunyaev \& Truemper 1979). The latter is often referred to as \textit{corona} (e.g. Melia \& Misra 1993; Svensson \& Zdziarski 1994; Churazov, Gilfanov \& Revnivtsev 2001) or \textit{flow} depending on whether the region is vertically or radially separated from the disc respectively. In particular, the \textit{truncated disc model} (e.g. Esin, McClintock \& Narayan 1997; Done, Gierli\'nski \& Kubota 2007) considers an optically thick geometrically thin accretion disc truncated at a certain radius $r_o$ and an optically thin geometrically thick hot flow extending from $r_o$ down to a radius equal or larger than the innermost stable circular orbit (ISCO). At the beginning of the outburst, the truncation radius is still relatively far from the black hole (BH) and the energy spectrum is dominated by the power law component. When the mass accretion rate increases, the truncation radius approaches the BH and the energy spectrum becomes dominated by the blackbody emission. Disc photons up-scatter in the hot flow cooling it down and, as a consequence, the power law softens. \\ Although the spectral properties of BHBs can be explained considering this two-regime accreting configuration (even though the precise way in which the disc and hot flow interact with each other is not clear), the origin of the fast variability is not fully understood, and a single model explaining both spectral and timing properties is a still matter of debate. The recently proposed model \textsc{propfluc} (Ingram \& Done 2011, 2012, hereafter ID11, ID12; Ingram \& van der Klis 2013, hereafter IK13) is based on the truncated disc model described above. Additionally, \textsc{propfluc} contains the ingredients of mass accretion rate fluctuations propagating through the hot flow, and precession of the entire hot flow caused by frame dragging close to the BH. Mass accretion rate fluctuations are generated at every radius of the hot flow and propagate towards the BH giving rise to a broad band noise component in the power spectrum (single hump power spectrum). The characteristic time scale of the noise is set by the viscous time scale in the hot flow (e.g. Lyubarskii 1997; Churazov, Gilfanov \& Revnivtsev 2001; Arevalo \& Uttley 2006). As a consequence of the propagation of the fluctuations, the time variability of the emission from every ring of the flow is correlated (with a time delay). Because the mass accretion rate fluctuations at larger radii, after propagating inward, modulate the amplitude of the fluctuations at smaller radii by multiplication, the process gives rise to the linear rms-flux relation observed in BHBs (Uttley \& McHardy 2001; Uttley, McHardy \& Vaughan 2005). Meanwhile, the Lense-Thirring (LT) precession of the entire hot flow (Stella \& Vietri 1998; Fragile et al. 2007; ID11) produces the QPO at a frequency depending on the mass distribution in the hot flow and on its radial dimension. \\ Rapisarda et al. (2014) (hereafter RIK14) presented the first application of \textsc{propfluc} to study the BH candidate MAXI J1543-564. They fitted selected power spectra of the rising phase of the 2011 outburst of the source with the single hump power spectrum calculated by \textsc{propfluc} and traced the evolution of the physical parameters in these observations.\\ The \textsc{propfluc} version used in RIK14 produces a single hump power spectrum originating from mass accretion rate fluctuations arising only in the hot flow. However, timing analysis of BHBs shows that their power spectrum in the LHS/HIMS is often characterized by a more complex structure than a single hump (e.g Belloni et al. 1997; Homan et al. 2001; Kalamkar et al. 2015a), requiring two or three broad Lorentzians to be fitted (low, mid, and high frequency Lorentzian). Additionally, BHBs often show time lags between different energy bands associated with this broad band variability (e.g., Miyamoto et al. 1988; Nowak, Wilms \& Dove 1999a). The delay between emission in different energy bands depends on the geometry of the accreting system and can be used to constrain different accretion models (e.g. Miyamoto \& Kitamoto 1989; B$\rm \ddot{o}$ttcher \& Liang 1999; Misra 2000; Nowak et al. 1999b; Kotov et al. 2001; Arevalo \& Uttley 2006).\\ Combining spectral and timing analysis it is possible to obtain clues about the origin of the different power spectral components. In particular, Wilkinson \& Uttley (2009), on the basis of measurements of variability amplitudes of X-ray spectral components, suggested that low frequency noise is the result of intrinsic variability generated in the disc and propagating through the flow. By its very nature, propagation also predicts time lags between soft and hard energy bands, but up to now these two aspects of the propagation hypothesis have never been jointly considered in a quantitative analysis. \\ As pointed out in IK13, with the model \textsc{propfluc} we can simultaneously predict these time lags, the variability amplitudes, and the coherence between energy bands by calculating power spectra at different energies, and cross-spectra between those energies. These predictions can then jointly be fitted to observed power and cross-spectra. The model can also be adapted to simulate extra disc variability and produce a two-hump power spectrum by considering mass accretion rate fluctuations generated both in the disc and in the flow, all propagating towards the BH. Fits to cross- and power spectra of BHBs in the LHS/HIMS characterized by a two-hump profile can then be attempted using observations spanning the low energy range where the disc emission is concentrated. \\ In this paper, we analyze data from MAXI J1659-152, a BH discovered in 2010 (Mangano et al. 2010; Negoro et al. 2010). During its 2010 outburst, MAXI J1659-152 followed the usual behavior observed in BHB outbursts (Mu{\~n}oz-Darias et al. 2011). Previous timing analysis of the source using the \textit{Rossi X-ray Timing Explorer} (RXTE; Jahoda et al. 1996) and \textit{Swift} (Gehrels et al. 2004) observations (Kalamkar et al. 2011; Kalamkar et al. 2015a), showed that its power spectra in the HIMS are characterized by several broad band components with characteristic frequencies between $\approx 0.001$ and $\approx 5$ Hz. We explore the hypothesis put forward by Kalamkar et al. (2015a) that some of this enhanced low frequency variability originates in the disc by performing joint fits of the power and cross-spectra of MAXI J1659-152 in the HIMS using \textit{Swift} XRT data in two different energy bands (0.5 - 2.0 keV and 2.0 - 10.0 keV). The \textit{Swift} XRT data allow us to study the source from the beginning of the outburst (RXTE started observing the source 3 days later) in an energy range where the disc emission is significant.\\ Sec. \ref{sec:mod} is dedicated to the description of the new two-hump version of \textsc{propfluc}, Sec. \ref{sec:obs} briefly describes how we reduced and analyzed the data, and in Sec. \ref{sec:res} and Sec. 5 we present and discuss the results of our fits, respectively. By a strictly quantitative analysis, we find, perhaps counterintuitively, that the small lag observed in the broadband noise between these two energy bands is entirely consistent with mass accretion rate fluctuations in the disc propagating to the hot flow. | We applied the double hump model \textsc{propfluc} to investigate the HIMS of MAXI J1659-152 using \textit{Swift} data. In the model, low frequency broad band components are interpreted as the result of mass accretion rate fluctuations arising in the disc and propagating towards the BH through the hot precessing flow. This double hump model was statistically preferred to a single hump model for most of the GTIs analyzed. In our analysis we detected only small phase lag associated with the low frequency variability, however model predictions are consistent with the data. Using both spectral and timing analysis we estimated recessing trend in truncation radius, and from that we inferred the mass accretion rate. Considering the truncation radius estimate from the \textsc{propfluc} fit and the maximum temperature in the disc (spectral fit parameter), we found a peak in the average increasing mass accretion rate trend that matches the variability properties of the accreting system (the amount of variability generated and the viscous frequency in the disc and in the flow). Considering the truncation radius estimate from spectral fit, would have lead to an almost constant mass accretion rate, in contrast to observations (the total counts increase almost linearly in our observation sample). Our analysis constitutes the first joint fitting of compact object cross-spectra and power spectra with a single self-consistent physical model. \\ \\ \textbf{\textit{ACKNOLEDGEMENTS}}\\ We thank the anonymous referee for his/her useful comments that greatly helped to improve the manuscript. S. Rapisarda, A. Ingram, and M. van der Klis acknowledge support from the Netherlands Organization for Scientific Research (NWO). M. Kalamkar acknowledges support by Marie Curie FP7-Reintegration-Grant under contract no. 2012-322259. This research has made use of the XRT Data Analysis Software (XRTDAS) developed under the responsibility of the ASI Science Data Center (ASDC), Italy. | 16 | 7 | 1607.08178 |
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