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1808 | 1808.08170_arXiv.txt | Using a sample of Nova-like stars from the Ritter and Kolb catalog, we examine the relationship between their Gaia determined absolute magnitude and the inclination of the binary system. Webbink et al. (1987) derived a relationship between these two variables that provides a good fit and allows differentiation between $\dot{M}$ (and possibly M$_{WD}$) as a function of inclination. We show that the spread in M$_V$, at a given $i$, is dominated by the mass transfer rate with only a small dependence on the white dwarf mass. The validated relation shows that present-day theoretical population studies of cataclysmic variables as well as model fits to observational data yield mass transfer rates and white dwarf masses consistent with the Gaia derived M$_V$ for the nova-like stars. | \label{sec:intro} There have been a number of previous works that attempted to discern the relationship between the absolute magnitude (M$_V$) of an accretion disk bearing binary star (i.e., a cataclysmic variable) and its system inclination. In principle, such a relationship is easy to imagine since seeing the accretion disk face-on will make the binary system substantially brighter than seeing only the thin edge of cooler material. From both first principles (Paczynski \& Schwarzenberg-Czerny, 1980; Mayo et al., 1980) and observational results mixed with theory (Warner, 1986; Webbink et al., 1987), astronomers have produced a number of empirical relationships for M$_V$ vs. $i$ as an aid in the quest to understand the disk structure and distances to such binaries. More recently, general values and equations, based on the above accepted paradigms pervade the literature and are casually used and excepted as correct for individual or classes of stars (e.g., Warner 1995, Patterson 2011; Ramsay 2017). The original relationships mentioned above were developed for specific cases; U Gem, UX UMa-like disks, recurrent novae, and nova remnants. Most of these types of cataclysmic variable generally have one thing in common - their accretion disk is believed to be (or was modeled as) optically thick with the disk dominating the light output in the visible part of the spectrum. However, these past studies used a mixture of system types, orbital periods, and techniques to formulate their relationships. In this paper, we revisit the connection between a bright disk system's absolute magnitude and its inclination. We use a model-independent approach based on a set of nova-like stars and new results from Gaia data release two (DR2). We assume that our stars M$_V$ is dominated by the light from the optically thick disk and that, for a constrained range in orbital period, their system properties are similar such that the disk inclination will dominate M$_V$; white dwarf mass, $q$, and mass transfer rate being second order effects. There is, of course, much observational evidence that disks of nova-like systems are optically thick (e.g. Bisol et al. 2012, Baptista et al. 1995a,b). Our assumptions are well founded based on theoretical results (e.g., Howell et al., 2001; Kolb \& Baraffe 2000) and the literature reviewed later on. We discuss our sample in the next section, review some of the original relationships and highlight their features in \S3. Finally we summarize our results, and as a spoiler, we find that the old, venerable relationships work pretty well. | We have used a sample of nova-like systems with orbital periods of 3-6 hr and containing accretion disks that dominate the light in the visible part of the spectrum. Taking the new Gaia DR2 parallax results, we determined the absolute magnitudes for these stars, and with system inclinations taken from the Kolb \& Ritter CV catalog and current references, produced the relationship between M$_V$ and $i$. Two previous M$_V$-$i$ relationships, based on disk models, initially applied to nova remnants, and still is use today, were examined and proved to be fairly representative of the data. Webbink et al., (1987) derived a relationship between these two variables that provides a good fit to the observations and allows differentiation between $\dot{M}$ values for a given $i$. We show that the spread in M$_V$, for a given inclination, is indeed dominated by the mass transfer rate (i.e., the disk luminosity being proportional to T$^{4}$; see eq. 5.20, Frank, King, \& Raine 1992) with a small, but perhaps measurable dependence on the white dwarf mass. Additionally, we confirm that modern theoretical population studies of cataclysmic variables as well as model fits to observational data for individual systems, as discussed in various literature articles, yield derived mass transfer rates consistent with the true M$_V$ for the nova-like stars. | 18 | 8 | 1808.08170 |
1808 | 1808.08988_arXiv.txt | Gaia Data Release 2 includes observational data for 14,099~pre-selected asteroids. From the sparsely sampled $G$ band photometry, we derive lower-limit lightcurve amplitudes for 11,665 main belt asteroids in order to provide constraints on the distribution of shapes in the asteroid main belt. Assuming a triaxial shape model for each asteroid, defined through the axial aspect ratios $a > b$ and $b=c$, we find an average $b/a=0.80\pm0.04$ for the ensemble, which is in agreement with previous results. By combining the Gaia data with asteroid properties from the literature, we investigate possible correlations of the aspect ratio with size, semi-major axis, geometric albedo, and intrinsic color. Based on our model simulations, we find that main belt asteroids greater than 50~km in diameter on average have higher $b/a$ aspect ratios (are rounder) than smaller asteroids. We furthermore find significant differences in the shape distribution of main belt asteroids as a function of the other properties that do not affect the average aspect ratios. We conclude that a more detailed investigation of shape distribution correlations requires a larger data sample than is provided in Gaia Data Release 2. | \label{sec:intro} Asteroid physical properties provide important constraints on the formation and evolution of our Solar System, including clues to the different phases of orbital instability and migration of the giant planets \citep[see, e.g.,][]{Morbidelli2015}. Collisions within the asteroid belt can play a major role altering the asteroid population. These clues are recorded in the shapes of individual asteroids, so that a global understanding of the shape distribution --- as informed by a large catalog --- provides insight into our Solar System's evolution. Shape information can be determined for individual asteroids from photometric data over a long period of time, with Doppler-delay radar imaging, with ground-based adaptive optics techniques, or {\it in situ} via spacecraft data. These last three techniques are limited in the number of targets that can be reasonably observed. As a result, only a small percentage of km-scale main belt asteroids have known shapes and spin-pole orientations. For instance, the DAMIT database\footnote{Database of Asteroid Models from Inversion Techniques: \url{http://astro.troja.mff.cuni.cz/projects/asteroids3D/web.php}} \citep{Durech2010} holds information for less than 1,000 asteroids as of writing this. Currently, the predominant method for determination of asteroid shape and spin properties is the inversion of dense photometric lightcurves, as initially developed by \citet{Kaasalainen2001a} and \citet{Kaasalainen2001b}. Without dense photometric data it is difficult to derive detailed shape and spin pole orientations for individual asteroids. With a sufficiently large dataset, however, it is possible to determine a statistical shape distribution, even if only a small number of data points are present for each individual asteroid. The advantage of a large untargeted dataset is that it provides an estimate for a population's shape distribution without being subject to observational biases in favor of elongated objects (i.e.,~higher lightcurve amplitudes) that are commonly present in targeted observations. Previous work on the shape distribution for main belt asteroids (MBAs) has been carried out by \citet{McNeill2016}, \citet{Cibulkova2016}, \citet{Nortunen2017}, and \citet{Cibulkova2018}. European Space Agency's Gaia astrometric space observatory \citep{Gaia2016} measures the positions, distances, proper motions, and other physical properties of more than one billion stars in our galaxy with unprecedented accuracy. In addition to its main mission, Gaia will also observe a significant fraction of the currently known asteroid population (${\sim}$350,000 Solar System small bodies) and discover previously unknown asteroids \citep{Spoto2018}. The sample of Gaia-observed asteroids provides a large uniform and magnitude-limited set of observations perfectly suited to independently investigate the distribution of asteroid shapes. | \label{sec:conclusions} We measure lower limit lightcurve amplitudes for 11,665 main belt asteroids in the Gaia Data Release 2 catalog. We derive the mean aspect ratio for main belt asteroids as $b/a=0.80\pm0.04$, which is in agreement with previous studies. We investigate trends in the shape distribution as a function of semi-major axis, intrinsic color, diameter, and geometric albedo, using data from the literature. Based on our model simulations, we find that main belt asteroids greater than 50~km in diameter have on average higher $b/a$ aspect ratios (are rounder) than smaller asteroids. We furthermore find significant differences in the derived lower limit amplitude distributions with respect to semi-major axis, intrinsic color ($a^*$), and geometric albedo. We predict that more detailed population and shape distribution studies will be possible with the availability of Gaia Data Release 3. | 18 | 8 | 1808.08988 |
1808 | 1808.09291.txt | In this work we study an alternative scheme for an Emergent Universe scenario in the context of a Jordan-Brans-Dicke theory, where the universe is initially in a truly static state supported by a scalar field located in a false vacuum. The model presents a classically stable past eternal static state which is broken when, by quantum tunneling, the scalar field decays into a state of true vacuum and the universe begins to evolve following the extended open inflationary scheme. | \label{Int} The standard cosmological model (SCM) \cite{weinberg,peebles,kolb} and the inflationary paradigm \cite{Guth1,Albrecht,Linde1,Linde2} are shown as a satisfactory description of our universe \cite{weinberg,peebles,kolb}. However, despite its great success, there are still important open questions to be answered. One of these questions is whether the universe had a definite origin, characterized by an initial singularity or if, on the contrary, it did not have a beginning, that is, it extends infinitely to the past. Theorems about spacetime singularities have been developed in the context of inflationary universes, proving that the universe necessarily has a beginning. In other words, according to these theorems, the existence of an initial singularity can not be avoided even if the inflationary period occurs, see Refs.~\cite{Borde:1993xh,Borde:1997pp, Borde:2001nh, Guth:1999rh,Vilenkin:2002ev}. In theses theorems it is demonstrated that null and time-like geodesics are generally incomplete in inflationary models, regardless of whether energy conditions are maintained, provided that the average expansion condition ($H> 0$) is maintained throughout of these geodesics directed towards the past, where $H$ is the Hubble parameter. The search for cosmological models without initial singularities has led to the development of the so-called Emergent Universes models (EU) \cite{Ellis:2002we,Ellis:2003qz, Mulryne:2005ef, Mukherjee:2005zt,Mukherjee:2006ds,Banerjee:2007qi,Nunes:2005ra,Lidsey:2006md}. In the EU scheme it is assumed that the universe emerged from a past eternal Einstein Static (ES) state to the inflationary phase and then evolves into a hot big bang era. % These models do not satisfy the geometrical assumptions of the theorems \cite{Borde:1993xh,Borde:1997pp, Borde:2001nh, Guth:1999rh,Vilenkin:2002ev} and they provide specific examples of nonsingular inflationary universes. % % Usually the EU models are developed by consider a universe dominated by a scalar field, which, during the past-eternal static regime, is rolling on the asymptotically flat part of the scalar potential (see Fig.~(\ref{fig:Potential-1})) with a constant velocity, providing the conditions for a static universe, see for example models \cite{Ellis:2002we, Ellis:2003qz, Mulryne:2005ef,Nunes:2005ra, Lidsey:2006md, delCampo:2007mp, delCampo:2009kp, delCampo:2010kf,delCampo:2011mq,Guendelman:2011zza,Guendelman:2011fr, Guendelman:2013dka, Guendelman:2014bva, Guendelman:2015uca, delCampo:2015yfa}. % Other possibility is to consider EU models in which the scale factor only asymptotically tends to a constant in the past \cite{Mukherjee:2005zt, Mukherjee:2006ds, Banerjee:2007sg,Debnath:2008nu, Paul:2008id, Beesham:2009zw,Debnath:2011qi, Mukerji:2011wq, Labrana:2013oca, Huang:2015zma}. % We can note that in these schemes of Emergent Universe are not all truly static during the static regime. At this respect a new scheme for an EU model was proposed in Ref.~\cite{Labrana:2011np}, where the universe is initially in a truly static state supported by a scalar field located in a false vacuum, also see Refs.~\cite{Labrana:2013kqa,Labrana:2014yta}. The universe begins to evolve when, by quantum tunneling, the scalar field decays into a state of true vacuum. % For simplicity, in this first approach to this new scheme of EU, the model was developed in the context of General Relativity (GR). % In particular in Ref.~\cite{Labrana:2011np} was concluded that this new mechanism for an Emergent Universe is plausible and could be an interesting alternative to the realization of the Emergent Universe scenario. However, as this first model was developed in the context of General Relativity, the past eternal static period, suffer from classical instabilities associated with the instability of Einstein's static universe. % \begin{figure}[t] \centering \includegraphics[scale=0.75]{Fig1.pdf} \caption{Schematic representation of a potential for a standard Emergent Universe scenario.} \label{fig:Potential-1} \end{figure} % The ES solution is unstable to homogeneous perturbations, as was early discussed by Eddington in Ref.~\cite{Eddington} and more recently studied in Refs.~\cite{Gibbons:1987jt, Gibbons:1988bm, Harrison:1967zz, Barrow:2003ni}. % The instability of the ES solution ensures that any perturbations, no matter how small, rapidly force the universe away from the static state, thereby aborting the EU scenario. This instability is possible to cure by going away from GR, for example, by consider a Jordan-Brans-Dicke (JBD) theory at the classical level, where it have been found that contrary to general relativity, a static universe could be classically stable, see Refs.~\cite{delCampo:2007mp, delCampo:2009kp, Huang:2014fia}. \begin{figure}[t] \centering \includegraphics[scale=0.75]{Fig2.pdf} \caption{Scalar field potential $U(\psi)$. Here $U_F=U\left(\psi_F\right)$ and $U_T=U\left(\psi_T\right)$.} \label{fig:Potential-2} \end{figure} In this work, we are interested in apply the scheme of Emergent Universe by Tunneling of Ref.~\cite{Labrana:2011np} to EU models which present stable past eternal static regimes. % In particular, we study this scheme in the context of a JBD theory, similar to the one studied in Refs.~\cite{delCampo:2007mp, delCampo:2009kp}, but where the static solution is supported by a scalar field located in a false vacuum. In this case we are going to show that, contrary to what happens in Ref.~\cite{Labrana:2011np}, the ES solution is classically stable. The Jordan-Brans-Dicke \cite{Jbd} theory is a class of models in which the effective gravitational coupling evolves with time. The strength of this coupling is determined by a scalar field, the so-called Brans-Dicke field, which tends to the value $G^{-1}$, the inverse of the Newton's constant. The origin of Brans-Dicke theory is found in the Mach's principle according to which the property of inertia of material bodies arises from their interactions with the matter distributed in the universe. In modern context, Brans-Dicke theory appears naturally in supergravity models, Kaluza-Klein theories and in all known effective string actions \cite{Freund:1982pg, Appelquist:1987nr, Fradkin:1984pq, Fradkin:1985ys, Callan:1985ia, Callan:1986jb, Green:1987sp}. In particular in this work we are going to consider that the universe is initially in a truly static state, which is supported by a scalar field $\psi$ located in a false vacuum ($\psi = \psi_F$), see Fig.~(\ref{fig:Potential-2}). The universe begins to evolve when, by quantum tunneling, the scalar field decays into a state of true vacuum. Then, a small bubble of a new phase of field value $\psi_W$ can be formed, and expands as it converts volume from high to low vacuum energy and feeds the liberated energy into the kinetic energy of the bubble wall. This process was first studied by Coleman \& De Luccia in \cite{Coleman:1977py, Coleman:1980aw} in the context of General Relativity. If the potential has a suitable form, inflation and reheating may occur inside the bubble as the field rolls from $\psi_W$ to the true minimum at $\psi_T$, in a similar way to what happens in models of Open Inflationary Universes and Extended Open Inflationary Universes, see for example \cite{linde, re8, delC1, delC2, Balart:2007je}, where the interior of the bubble is modeled by an open Friedmann-Robertson-Walker universe. %%%%%%%%%%%%%%%%%%%%%%%%%% \begin{figure}[t] \begin{center} \includegraphics[scale=0.65]{Fig3} \caption{Potential with a false and true vacuum.} \label{fig:Potential-3} \end{center} \end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The advantage of the EU by Tunneling scheme (and of the Emergent Universe in general), over the Eternal Inflation scheme is that it correspond to a realization of a singularity-free inflationary universe. % As was discussed in Refs.~\cite{Borde:1993xh,Borde:1997pp, Borde:2001nh, Guth:1999rh,Vilenkin:2002ev}, Eternal Inflation is usually future eternal but it is not past eternal, because in general space-time that allows for inflation to be future eternal, cannot be past null complete. On the other hand Emergent Universe are geodesically complete. %%%%%%%%%%% Notice that in the EU by Tunneling scheme, the metastable state which support the initial static universe could exist only a finite amount of time. Then, in this scheme of Emergent Universe, the principal point is not that the universe could have existed an infinite period of time, but that in theses models the universe is non-singular because the background where the bubble materializes is geodesically complete. This implies that we have to consider the problem of the initial conditions for a static universe. Respect to this point, there are very interesting possibilities discussed for example in the early works on EU \cite{Ellis:2003qz} and more recently in \cite{Labrana:2011np}. % One of these options is to explore the possibility of an Emergent Universe scenario within a string cosmology context \cite{Antoniadis}. Other possibility is that the initial Einstein Static universe is created from "nothing" \cite{Tryon,Vilenkin-cre}, see Refs.~\cite{Mithani} for explicit examples. % It is interesting to mention that the study of the Einstein Static solution as a preferred initial state for our universe have been considered in the past, where it has been proposed that entropy considerations favor the ES state as the initial state for our universe \cite{Gibbons:1987jt,Gibbons:1988bm}. In this paper we consider a simplified version of this scheme, with the focus on studying the process of creation and evolution of a bubble of true vacuum in the background of an ES universe in the context of a JBD theory. % This is motivated because we are mainly interested in the study of new ways of leaving the static period and begin the inflationary regime for Emergent Universe models which present a classically stable static state period. In particular, in this paper we consider a JBD theory where one of the matter content of the model is a scalar field (inflaton) with a potential similar to Fig.~(\ref{fig:Potential-3}) and we study the process of tunneling of the scalar field from the false vacuum $U_F$ to the true vacuum $U_T$ and the subsequent creation and evolution of a bubble of true vacuum in the background of an stable ES universe. % The simplified model studied here contains the essential elements of the scheme we want to present, so we postpone the detailed study of the inflationary period, which occurs after the tunneling, for future work. The paper is organized as follow. In Sect.~\ref{sec:estadoestatico} we study a Einstein static universe supported by a scalar field located in a false vacuum and its stability in the context of a JBD theory. In Sect. \ref{sec:tunel} we study the tunneling process of the scalar field from the false vacuum to the true vacuum and the subsequent creation of a bubble of true vacuum in the background of the Einstein static universe for a JBD theory. In Sect.~\ref{sec:evolucionbur} we study the evolution of the bubble after its materialization. In Sect.~\ref{sec:conclu} we summarize our results. | \label{sec:conclu} In this paper we study an alternative scheme for an Emergent Universe scenario called Emergent Universe by tunneling. In this scheme the universe is initially in a truly static state supported by a scalar field which is located in a false vacuum. The universe begins to evolve when, by quantum tunneling, the scalar field decays into a state of true vacuum. The EU by tunneling scheme was originally developed in Ref.~\cite{Labrana:2013kqa}, in the context of General Relativity, where it was concluded that this mechanism is feasible as an EU scheme. Nevertheless, this first model present the problem that the ES solution is classically instable. % The instability of the ES solution ensures that any perturbation, no matter how small, rapidly force the universe away from the static state, thereby aborting the EU scenario. The present work is the natural extension of the idea presented in Ref.~\cite{Labrana:2013kqa}, but where the problem of the classical instability of the static solution is solved by going away from General Relativity and consider a JBD theory. In particular, in this work we focus our study on the process of tunneling of a scalar field and the consequent creation and evolution of a bubble of true vacuum in the background of a classically stable Einstein Static universe. % % Our principal motivation is the study of new ways of leaving the static period and begin the inflationary regime in the context of Emergent Universe models. In the first part of the paper, Sect.~\ref{sec:estadoestatico}, we study an Einstein static universe supported by a scalar field located in a false vacuum and its stability in the context of a JBD theory. % Contrary to General Relativity, we found that this static solution could be stable against isotropic perturbations if some general conditions are satisfied, see Eqs.~(\ref{condgamma})-(\ref{vdosprima}). % This modification of the stability behavior has important consequences for the emergent universe by tunneling scenario, since it ameliorates the fine-tuning that arises from the fact that the ES model is an unstable saddle in GR and it improves the preliminary model studied in Ref.~\cite{Labrana:2011np}. % % In this study, for simplicity, we have not considered inhomogeneous or anisotropic perturbations. % At this respect, the stability of the ES solution under anisotropic, tensor and inhomogeneous scalar perturbations have been studied in the context of JBD theories in Refs.~\cite{delCampo:2009kp, Huang:2014fia}. It was found for theses JBD models that different from General Relativity \cite{Barrow:2003ni} and others modified theories of gravity as $f(R)$ \cite{Seahra:2009ft} or modified Gauss-Bonnet gravity \cite{Huang:2015kca}, that a static universe which is stable against homogeneous perturbations, could be also stable against anisotropic and inhomogeneous perturbations. We expect a similar behavior for our JBD model, were the static universe is supported by a scalar field located in a false vacuum. Then, we expect that in our case the inhomogeneous and anisotropic perturbations do not lead to additional instabilities. Nevertheless, we intend to return to these points in the near future by working an approach similar to that followed in Refs.~\cite{Barrow:2003ni,delCampo:2009kp,inhomogeneous, Huang:2014fia}. In Sect. \ref{sec:tunel} we study the tunneling process of the scalar field from the false vacuum to the true vacuum and the consequent creation of a bubble of true vacuum in the background of Einstein static universe for a JBD theory. In particular we determinate the nucleation rate of the true vacuum bubble using the approaches developed in Ref.~\cite{Holman:1989gh} and previous results obtained in Ref.~\cite{Labrana:2011np}. The classical evolution of the bubble after its nucleation is studied in Sect.~\ref{sec:evolucionbur} where we found that once the bubble has materialized in the background of an ES universe, it grows filling the background space. % This demonstrates the viability of our EU model, since there is the possibility of having an open inflationary universe inside the bubble. % % During this study we consider the gravitational back-reaction of the bubble by using the formalism developed in Ref.~\cite{Sakai:1992ud} applied to a JBD theory. At this respect we found a system of coupled differential equations, which we solved numerically. Three specific examples of these solutions were shown in Sect.~\ref{sec:evolucionbur} concerning to different background material contents. It is worth to note that once the bubble has materialized, from conditions \eqref{tresocho}, it follows that if one of the regions of spacetime separated by the wall is homogeneous, then the other region is, in general, inhomogeneous \cite{Sakai:1992ud}. Given that in our case the exterior of the bubble is a homogeneous universe, then the interior of the bubble will be, in general, inhomogeneous. However, since the degree of inhomogeneity depends on the difference in the energy density of the interior and the exterior of the bubble, it is possible in our case to decrease this inhomogeneity by adjusting the parameters of the static solution as was discussed in Ref.~\cite{Sakai:1992ud}. % Then in our model, it is possible to study the feasibility of having an open inflationary universe inside the bubble. % Nevertheless, given the similarities, we expect that the behavior inside the bubble of the EU by tunneling, will be similar to the models of single-field open and extended open inflation, as the ones studied in Refs.~\cite{linde, re8, delC1, delC2,Balart:2007je}. We expect to return to this point in the near future. | 18 | 8 | 1808.09291 |
1808 | 1808.06737_arXiv.txt | The turbulent burning of nuclei is a common phenomenon in the evolution of stars. Here we examine a challenging case: the merging of the neon and oxygen burning shells in a 23\,\Msun star. A previously unknown quasi-steady state is established by the interplay between mixing, turbulent transport, and nuclear burning. The resulting stellar structure has two burning shells {\em within a single convection zone}. We find that the new neon burning layer covers an extended region of the convection zone, with the burning peak occurring substantially below where the Damk\"ohler number first becomes equal to unity. These characteristics differ from those predicted by 1D stellar evolution models of similar ingestion events. We develop the mean-field turbulence equations that govern compositional evolution, and use them to interpret our data set. An important byproduct is a means to quantify sub-grid-scale effects intrinsic to the numerical hydrodynamic scheme. For implicit large eddy simulations, the analysis method is particularly powerful because it can reveal where and how simulated flows are modified by resolution, and provide straightforward physical interpretations of the effects of dissipation or induced transport. Focusing on the mean-field composition variance equations for our analysis, we recover a Kolmogorov rate of turbulent dissipation without it being imposed, in agreement with previous results which used the turbulent kinetic energy equation. | \label{s-intro} It is now feasible to simulate stellar convection in three-dimensions (3D), with realistic microphysics, multiple species of nuclei, and sufficient resolution in space and time to represent turbulent flow \citep{MeakinArnett2007,Mocak2011,Woodward2013}. Historical work on stellar convection \citep{bohmvitense1958} and 3D simulations of stellar atmospheres \citep{steinnordlund1998} have generally focused on flows having uniform composition, a case which is usually appropriate for the outer layers of stars. By contrast convection in stellar interiors is generally characterized by nuclear burning and nonuniform composition. Here we examine the interaction between turbulent convection, thermonuclear burning, and entrainment at boundaries. Simulations of convective shells, driven by nuclear burning, show entrainment of material from surrounding stable layers. Erosion at boundaries introduces inhomogeneities in composition, entropy, and buoyancy into the convective flow. This can be viewed as a multi-stage process of entrainment, transport, dispersion (or stirring) induced from the largest to smallest eddies, as well as diffusion, spanning the full spectrum of space-time scales of the flow \citep{Dimotakis2005}. The feedback of such mixing on nuclear burning, convection, and its impact on the evolution of the star remains largely unexplored: an ad hoc diffusion operator is almost universally used in stellar evolution. This paper begins to analyze these stages. We focus on the oxygen burning shell in a massive supernova progenitor. Oxygen burning and neon burning occur at sufficiently similar temperatures that these burning shells may interact \citep{wdaNe74,wdaO74}. Interaction was indeed found in the 3D simulations of \citet{meakinphd} and \citet{MeakinArnett2007}, but was not analyzed in detail there. Here we present a detailed account of the compositional mixing and modified nuclear burning, which occurs as the convective oxygen-burning shell merges with the (initially) stable overlying neon shell. We use 3D numerical hydrodynamics in the implicit large eddy simulation framework ILES \citep{ILES}, which means we solve the Euler equations with a non-oscillatory finite volume numerical fluid solver. In the current study we use the PPM method \citep{ColellaWoodward1984}. The dissipation in such solvers comes from solving the Riemann shock problem over each zone $\Delta_s$, giving a dissipation rate\footnote{The change in specific kinetic energy over the shock traversal time gives the dissipation rate. PPM approximates sub-grid structure as piece-wise parabolic, smoothing higher-order terms and decreasing information (complexity). Variances in velocity and in scalar variables dissipate/diffuse at this same rate at the sub-grid scale.} of $\sim v_s^3/\Delta_s$ as a shock of speed $v_s$ traverses a zone. In a turbulent cascade the mean damping is ${\cal D} \sim v_s^3/\Delta_s \sim v^3/\ell,$ which is determined by the rms velocity $v$ and dimension $\ell$ of the turbulent region. Use of such solvers introduces an implicit sub-grid model which corresponds to a Kolmogorov turbulent cascade, freeing computational resources to capture the large scales relevant to astrophysics. See \cite{ILES} for references and a more rigorous discussion. This is in contrast to the direct numerical simulation approach DNS \citep{Pope2000}, which solves the Navier-Stokes difference equations on the grid, all the way down to the dissipation scale. In this approach most of the computational effort is spent on these small scales, which are buried in the turbulent cascade. \cite{sytine2000} showed that both methods converge to the same result, but that ILES is more efficient for highly turbulent flows. We have confirmed that our simulations extend from the integral (large) scale down into the inertial range of the turbulent cascade (e.g. \citealt{Cristini2017}). In this paper, we extend our analysis by using an approach inspired by Reynolds-averaged Navier-Stokes (RANS) methods \citep{VialletMeakin2013,mmva14,amvclm15}. Our approach differs from traditional RANS \citep{Besnard1992,Chassaing2010} in two fundamental ways: (1) the fluctuations are taken from our simulations, and hence are dynamically constrained, and (2) we solve the ILES Euler equations. A more accurate acronym than ``RANS'' (which we have used previously) is needed for clarity; we choose ``Reynolds averaged ILES'' (RA-ILES), to distinguish our approach. {\em Unlike unconstrained RANS analysis, our RA-ILES equations are complete, exact to the accuracy of our grid, and require little added computational cost.} There is no closure issue\footnote{There is no explicit Navier-Stokes viscosity term to generate higher order moments in the conventional way \citep{tritton1988}; the implicit turbulent cascade gives closure.}. % Previous papers \citep{MeakinArnett2007,VialletMeakin2013,amvclm15,Cristini2017} have focused on the turbulent kinetic energy equation (TKE); here our analysis shifts to mean-field transport equations for the density of $^{16}$O and $^{20}$Ne, and their turbulent fluxes and variances, complemented by analysis of relevant timescales and nuclear burning processes. These transport equations are the ones which deal with changes in composition variance: i.e., mixing. The paper is organized as follows: In \S\ref{sect:simulation-model} we describe the initial conditions and the 3D stellar model that we investigate in this paper. In \S\ref{sect:rans-definitions} we develop the RA-ILES equations used for our analysis. In \S\ref{sect:timescales} we define several timescales which we use to characterize physical processes operating in the simulated flow. In \S\ref{sect:results} we present the results from our RA-ILES analysis of oxygen and neon entrainment, transport, dissipation, and burning. We use the composition related mean-field equations and provide a systematic description of each term in the budget equations with an emphasis placed on physical interpretation. We then look at effects of resolution on our results in \S\ref{sect:resolution}. Finally, we conclude with a summary and discussion in \S\ref{sect:summary}. \begin{figure*} \includegraphics[width=0.49\hsize]{Figures/fig1a_ob3dB_initial_model_rho_t_new-eps-converted-to.pdf} \includegraphics[width=0.49\hsize]{Figures/fig1b_ob3dB_initial_model_x_new-eps-converted-to.pdf} \includegraphics[width=0.49\hsize]{Figures/fig1c_ob3dB_tavg300_rho_t_new-eps-converted-to.pdf} \includegraphics[width=0.49\hsize]{Figures/fig1d_ob3dB_tavg300_mean_x_insf_new-eps-converted-to.pdf} \caption{{\texttt Top row:} ({\em left}) Initial 1D background temperature $T$ and density $\rho$; and ({\em right}) composition profiles (mass fraction). {\texttt Bottom row:} Radial profiles after 300 seconds of evolution at which point the model has obtained a quasi-steady character. Only the most energetically important nuclear species are shown. Vertical dashed lines mark boundaries of convection.} \label{fig:initial-model} \end{figure*} | \label{sect:summary} We have analyzed coupled turbulence and nuclear burning in a 3D hydrodynamic implicit large eddy simulation (ILES) of an oxygen burning shell in the core of a $23~\Msun$ supernova progenitor star. \subsection{Hydrodynamic simulation}\label{sect:hydrosummary} Our initial stellar structure encompassed a single oxygen-burning convection zone with a stable Ne-burning layer above. After an initial transient episode which led to a gentle merging of the O-burning convective shell with the Ne-burning shell, a quasi-steady layered-shell state was formed. Two layers of stable nuclear burning were set up in the joined velocity field of a single convection zone, with $^{16}$O dominating the nuclear energy production at the bottom of the convection zone, but $^{20}$Ne dominating in the upper layers. The quasi-static state showed a Ne-burning rate in the convection zone up to seven orders of magnitude higher than in the initial 1D model. This was due to a steady entrainment of $^{20}$Ne into the hotter, oxygen burning shell. Further, neon burning was found to occur over a wide region in the convection zone. This is in contrast to the situation found in 1D codes for similar convective-reactive cases, for example during proton ingestion into He burning regions in low-mass stars. As a consequence of limitations in their mixing algorithms, 1D codes generally show burning in a very thin layer, at a depth where the ratio of the burning timescale equals that of the convective turnover timescale, i.e. where the Damk\"ohler number equals one. In our simulation this defines only the top of a wider burning region, and the peak of the burning is well below this point (Fig.~\ref{fig:convective-reactive2}). \subsection{RA-ILES analysis} The key results presented in this paper are based on a Reynolds-Averaged ILES (RA-ILES) analysis of the composition-related mean fields found for our 3D ILES simulation data, and extend our previous investigations which used the turbulent kinetic energy (TKE) equation. We develop RA-ILES equations (\S\ref{sect:rans-definitions}) which describe the time evolution of mean composition, turbulent composition flux, and composition variance for each isotope present in the stellar plasma. The equations derived are exact and do not employ approximations to the extent that the continuum approximation is appropriate and flow features are resolved. This contrasts with use of closure relationships and truncations of the RANS equations to construct models of turbulence. {\em The RANS equations, when closed by 3D ILES numerical simulations (RA-ILES), are able to represent the full range of hydrodynamical behavior that is present in a stellar interior simulation. } Our detailed % RA-ILES analysis revealed, over the timescale of the simulation ($\sim$4 turnovers): \begin{itemize} \item A simple balance between turbulent transport and burning in the convection zone is maintained, whereby any amount of material entrained into the convection zone is counterbalanced by nuclear burning (\S\ref{sec:meancomp}, Fig.\ref{fig:composition-transport}). \item The turbulent composition flux % appears to be quasi-static due to a complex balance between turbulent production, nuclear burning, turbulent transport, pressure force coupling, and centrifugal forces (\S\ref{sec:compflux}, Fig.\ref{fig:flux-transport}). \item Composition variance provides information about composition fluctuation amplitudes, and the residual of the composition variance mean field equation then gives direct information about numerical dissipation timescales (\S\ref{results:variance}). \end{itemize} \subsection{Dissipation and ILES} The residual of the composition variance equation implies a match with the Kolmogorov damping timescale (\S\ref{results:variance}, Fig.\ref{fig:sigma-transport}). This independently confirms a similar result found using the turbulent kinetic energy equation \citep{MeakinArnett2007,ArnettMeakin2010}, and indicates a deep consistency in our ILES approach. In the simulation, the dissipation of both turbulent kinetic energy and composition variance occurs at the sub-grid level, as desired. This should be true of the entropy and temperature variances as well, because ILES joins onto the inertial range of the turbulent cascade. \subsection{Implications for Stellar Evolution Theory} As discussed above (\S\ref{sect:hydrosummary}) we have identified very substantial differences between 1D stellar evolution modelling and our 3D modelling. These findings indicate that if shell mergers and entrainment do happen in real stars, 1D evolutionary models that do not include these phenomena are likely to be substantially in error. Moreover, based on our finding of the way in which ingested fuel burns in the convective zone, it appears that even if 1D models do incorporate shell mergers and entrainment, current 1D theory would be unable to model them reliably. The differences identified here will need to be addressed to improve modeling of convective-reactive phases. If not addressed the errors from these events will accumulate through the subsequent evolutionary phases and will cause uncertainty in, for example, pre-supernova evolution and the 3D explosion models that rely on the 1D models for their initial states. Further 3D modeling -- and, just as importantly, detailed analysis such as that presented in this paper -- is vital to guide these improvements by illuminating the physical processes at play. | 18 | 8 | 1808.06737 |
1808 | 1808.04504_arXiv.txt | \noindent With the number of confirmed rocky exoplanets increasing steadily, their characterisation and the search for exoplanetary biospheres is becoming an increasingly urgent issue in astrobiology. To date, most efforts have concentrated on the study of exoplanetary atmospheres. Instead, we aim to investigate the possibility of characterising an exoplanet (in terms of habitability, geology, presence of life etc.) by studying material ejected from the surface during an impact event. For given parameters characterising the impact event, we estimate the escaping mass and assess its subsequent collisional evolution in a circumstellar orbit, assuming a Sun-like host star. We calculate the fractional luminosity of the dust as a function of time after the impact event and study its detectability with current and future instrumentation. We consider the possibility to constrain the dust \emph{composition}, giving information on the geology or the presence of a biosphere. As examples, we investigate whether calcite, silica or ejected microorganisms could be detected. For a 20\,km diameter impactor, we find that the dust mass escaping the exoplanet is roughly comparable to the zodiacal dust, depending on the exoplanet size. The collisional evolution is best modelled by considering two independent dust populations, a \textit{spalled} population consisting of non-melted ejecta evolving on timescales of millions of years, and dust \textit{recondensed} from melt or vapour evolving on much shorter timescales. While the presence of dust can potentially be inferred with current telescopes, studying its composition requires advanced instrumentation not yet available. The direct detection of biological matter turns out to be extremely challenging. Despite considerable difficulties (small dust masses, noise such as exozodiacal dust etc.), studying dusty material ejected from an exoplanetary surface might become an interesting complement to atmospheric studies in the future.\\ \textbf{Keywords:} Biosignatures -- Exoplanets -- Impacts -- Interplanetary dust -- Remote sensing | Although terrestrial planets orbiting solar-type main-sequence stars seem to be relatively common in the Galaxy \citep[e.g.][]{Fressin_etal_2013,Petigura_etal_2013}, it is at the moment completely unclear whether the phenomenon of life is widespread in the universe or unique to our home planet. Therefore, considerable effort has been undertaken to identify suitable signatures of bioactivity (biosignatures) on exoplanets, in parallel to the search for extraterrestrial life within the Solar System (e.g.\ on Mars or Europa). A promising approach to identify a biosphere is to use the influence of life on the composition of the atmosphere \citep[e.g.][]{Kaltenegger_etal_2010,Seager_Deming_2010,Rugheimer_etal_2013,Seager_2014}. For example, the detection of species out of chemical equilibrium has been argued to be indicative of a biosphere. This idea was first discussed by \citet{Lederberg_1965}, \citet{Lovelock_1965,Lovelock_1975} and \citet{Hitchcock_Lovelock_1967}. These authors argued that states out of thermodynamic equilibrium can be seen as a generalised signature of life. For instance, in Earth's atmosphere, oxygen and methane are severely out of redox equilibrium because of biological forcing \citep[e.g.][]{Sagan_etal_1993}. Without life, methane would indeed rapidly be removed from the atmosphere by reacting with oxygen. The atmospheric properties of exoplanets as small as Earth have already been constrained \citep{deWit_etal_2016,Southworth_etal_2017}. However, measurements of the atmospheric composition or mass are not yet feasible. In the future, investigating the atmospheres of Earth analogues will become possible by using telescope concepts similar to \textit{Darwin} or the \textit{Terrestrial Planet Finder Interferometer} (TPF-I). These instruments are designed to directly image Earth-sized exoplanets in the habitable zone. Besides the spectroscopic characterisation of the atmosphere, they can also be used to study the surface reflectance \citep[e.g.][]{Hegde_Kaltenegger_2013}, including surface reflectance biosignatures \citep[e.g.][]{DesMarais_etal_2002}, for example from vegetation and also microbial mats \citep[e.g.][]{Seager_etal_2005,Hegde_etal_2015,Schwieterman_etal_2015}. In general, the best strategy to avoid biosignature false positives \citep[e.g.][]{Rein_etal_2014} is to characterise the potentially inhabited exoplanetary system (including the host star) as accurately as possible. Only when seen in context can the significance of a biosignature detection be assessed. The study of methods yielding information in addition to atmospheric mass and composition, or general bulk properties (exoplanet radius, bulk density, orbital period), is clearly warranted. In particular, ways to probe the geology of exoplanets or detect non-atmospheric biosignatures would be valuable complements. In this work, we investigate the idea of characterising an exoplanet by observing dust generated during an impact event. The composition of the dust holds information on the geology of the impacted planet and might contain biosignatures if the planet is inhabited. Depending on the impact parameters, the escaping debris can have a much larger surface area than the planet itself. Impact ejecta most easily escape from low gravity bodies. Thus, the largest amounts of dust are expected for small planets (or moons), the atmospheres of which are the most difficult to study. Several studies have considered dust generated from collisions or impacts, covering different regimes in terms of sizes of both the impactor and the target. Giant collisions involving two bodies of planetary size are expected to occur frequently during the final stage of terrestrial planet formation lasting for approximately 100\,Myr \citep[e.g.][]{Kenyon_Bromley_2006,Kokubo_Genda_2010}. For example, \citet{Jackson_Wyatt_2012} modelled the evolution and detectability of dust originating from the Moon-forming collision. \citet{Genda_etal_2015} estimated the total amount of dust produced from giant impacts and compared to observations of warm debris disks (i.e.\ dust in the terrestrial region). \citet{Jackson_etal_2014} modelled observational signatures of giant impacts occurring at large orbital radii. \citet{Morlok_etal_2014} linked infrared observations of dust from collisions to laboratory spectra of terrestrial and martian rocks. Dust can also be produced from mutual collisions of smaller bodies. This is commonly observed in the form of dusty debris disks, objects akin to the asteroid belt or the Kuiper belt in the solar system \citep[e.g.][]{Wyatt_2008}. In general, debris disk dust is believed to originate from a so-called collisional cascade: asteroidal or cometary bodies\footnote{These objects can be seen as leftover planetesimals that were not incorporated into planets.} collide and produce smaller bodies, which further collide to produce even smaller fragments, resulting in copious amounts of micron-sized dust grains. Analysing the dust properties allows for characterisation of the parent bodies at the top of the collisional cascade \citep[e.g.][]{deVries_etal_2012}. Collisions among asteroids or comets in debris disks can also be responsible for observed dust clumps \citep{Wyatt_Dent_2002,Kenyon_etal_2014}. In this work, we are concerned with yet another impact regime: we consider impacts of asteroidal or cometary bodies, typically tens of kilometres in size, with planetary bodies. An example for such an event on Earth is the famous Chicxulub impact that caused the Cretaceous-Paleogene (K-Pg) extinction about 65\,Myr ago. In contrast to mutual collisions of asteroids or comets, such events generate dust originating from our objects of interest, namely exoplanets. In addition, such impacts are expected to occur over the entire lifetime of an exoplanetary system and are not restricted to the first 100\,Myr. In particular, they may occur once a planet has been extensively modified by geological activities or the presence of life, the signatures of which could be imprinted in the ejected dust. We emphasise that the present work does not attempt to develop a detailed model covering all the subtleties of the impact process and the dust evolution. Rather, the intent of the paper is to present the idea and make some quantitative estimates. We employ a simplified approach to get a general idea of the ejected dust masses, the timescales on which the dust evolves and what instruments would be needed to observe impact-generated dust. Our paper is organised as follows: section \ref{Sec:interesting_substances} is a general discussion of potentially interesting substances escaping during an impact event. Section \ref{Sec:Modelling} presents the modelling of the impact and the subsequent collisional evolution. In section \ref{Sec:Discussion} we discuss our results and estimate instrument capabilities needed to detect minerals or biological matter. Section \ref{Sec:Conclusions} gives a summary and the conclusions. | \label{Sec:Conclusions} We have calculated the mass escaping from an exoplanet during an impact event for a limited number of impact scenarios (exoplanet sizes, impactor types) using a simple model. We have also determined the fraction of escaping ejecta that is not considerably shock damaged and remains in the solid state. We then computed the collisional evolution of the debris with a simplified analytical model based on timescales of collisions (production of new, smaller grains) and removal by PR-drag and radiation pressure. For the relatively small dust masses considered here, PR-drag is an important removal process, in contrast to most known debris disks. We consider two dust populations that we assume to evolve independently of each other: the spalled population that consists of ejecta that remained in the solid state during the impact event, and the recondensed population that formed from melt droplets or vapour. The recondensed population is initially much brighter since it has a larger mass and smaller grain sizes than the spalled population. However, it is also removed faster, such that the spalled population becomes dominant typically after a million years. The fractional luminosity of the impact generated dust is roughly comparable to the fractional luminosity of the zodiacal dust. Such a fractional luminosity is potentially in the reach of the presently available LBTI. Future instruments such as TPF-like telescopes will be able to detect the presence of dust both in thermal emission and scattered light. By studying the composition of the dust, one would gain information on the impacted exoplanet, its geology or the presence of a biosphere. The escaping masses we derive can be used to assess the detectability of (bio)signatures present in the subsurface of an exoplanet. As examples, we considered three different substances. Calcite would likely indicate the presence of liquid water at some point in the history of the exoplanet, provided the dust indeed originated from a planetary body. Our estimates show that a detection of calcite would require advanced far-IR instruments not yet available. Glassy silica on the other hand would potentially be detectable with JWST, and should be within the reach of a TPF-like telescope. Since glassy silica is linked to violent impacts or collisions, the detection of substantial amounts would suggest that the observed dust indeed originated from an impact event involving a planetary body or large planetesimals. Finally, we considered the direct detection of ejected biological matter (microorganisms) in reflected light. While microorganisms have absorption features due to water of hydration, their detection within the ejected debris seems unfeasible. Indeed, a large fraction of any microorganisms would be destroyed during the impact. In addition, taking Earth as a reference, the density of microorganisms is too low to allow detection. Our calculations show that ejected dust masses are relatively small. We note however that the dust cross-section is in general larger than the planetary cross-section. Also, studying ejected dust could be an interesting complement to atmospheric studies. Looking at dust from impact events seems to be the only possibility to study the subsurface composition of exoplanets directly, apart from exoplanets evaporating very close to their host star \citep[e.g.][]{vanLieshou_etal_2014}. Dust from impact events is a valuable piece of information with which to achieve a complete characterisation of an exoplanet (including its geology), which is important when assessing its habitability or potential signatures of a biosphere. However, the detailed study of dust from impact events on terrestrial planets has to wait for future instruments with high sensitivity and the ability to observe material in the terrestrial region of a star. | 18 | 8 | 1808.04504 |
1808 | 1808.09426_arXiv.txt | The most heavily polluted white dwarfs often show excess infrared radiation from circumstellar dust disks, which are modeled as a result of tidal disruption of extrasolar minor planets. Interaction of dust, gas, and disintegrating objects can all contribute to the dynamical evolution of these dust disks. Here, we report on two infrared variable dusty white dwarfs, SDSS~J1228+1040 and G29-38. For SDSS~J1228+1040, compared to the first measurements in 2007, the IRAC [3.6] and [4.5] fluxes decreased by 20\% by 2014 to a level also seen in the recent 2018 observations. For G29-38, the infrared flux of the 10~$\mu$m silicate emission feature became 10\% stronger between 2004 and 2007, We explore several scenarios that could account for these changes, including tidal disruption events, perturbation from a companion, and runaway accretion. No satisfactory causes are found for the flux drop in SDSS~J1228+1040 due to the limited time coverage. Continuous tidal disruption of small planetesimals could increase the mass of small grains and concurrently change the strength of the 10~$\mu$m feature of G29-38. Dust disks around white dwarfs are actively evolving and we speculate that there could be different mechanisms responsible for the temporal changes of these disks. | \label{sec:intro} G29-38 was the first single white dwarf discovered to display excess infrared emission \citep{ZuckermanBecklin1987}, and follow-up studies have shown that the excess flux arises from a close-in hot dust disk \citep{Graham1990b}. The origin of such a dust disk remained as a mystery until the asteroid tidal disruption model was proposed \citep{DebesSigurdsson2002, Jura2003}. According to this model, the disks are remnants of minor planets that were perturbed into the tidal radius of the white dwarf and eventually became totally disrupted. The infrared excess is often modeled as a geometrically thin and optically thick disk within the tidal radius of the white dwarf \citep{Jura2003}. These compact hot dust disks (temperature $\sim$ 1000~K, size 0.01~au) around white dwarfs are morphologically different from debris disks around main sequence stars (temperature $\sim$ 100~K, size a few tens to hundreds au). There are more than 40 white dwarfs that show infrared excess emission consistent with the presence of dusty disks \citep{Farihi2016}. Some of the dusty white dwarfs also display calcium triplet emission from circumstellar gas that spatially coincide with the dust disk \citep[e.g.][]{Gaensicke2006, Gaensicke2008}. About 25--50\% white dwarfs are polluted -- they display elements heavier then helium in their spectra \citep{Zuckerman2003,Zuckerman2010,Koester2014a}. In many cases, continuous accretion onto the white dwarf from circumstellar material is needed due to the short settling times of heavy elements. The connection between atmospheric pollution and dust disks was first explored in \citet{vonHippel2007}. The most heavily polluted white dwarfs are accompanied by an infrared excess from a dust disk. Spectroscopic observations of these polluted atmospheres have opened up a new field of measuring chemical compositions of extrasolar planetary material \citep{JuraYoung2014, Harrison2018, Hollands2018}. Some polluted white dwarfs are dynamically active and they vary on short timescales. For example, the infrared flux of SDSS~J0959$-$0200 dropped by 35\% between two observations in 2010, and remained at the same level afterwards until at least 2014 \citep{XuJura2014}. The gas emission lines around WD~J1617+1620 disappeared within a few years \citep{Wilson2014}. Most gas disks show gradual variations over a few years \citep{Wilson2015,Manser2016a, Manser2016b, Redfield2017, Dennihy2018}. Recently, transits from an actively disintegrating asteroid were detected around WD~1145+017 \citep{Vanderburg2015} -- a white dwarf that is also heavily polluted, has an infrared excess from a dust disk, and displays absorption lines from circumstellar gas \citep{Xu2016}. The optical light curve of WD~1145+017 is changing on a daily basis, likely due to the vigorous nature of tidal disruption \citep[e.g.][]{Gaensicke2016,Gary2017}. The dynamical mechanism responsible for white dwarf pollution and tidal disruption is an area of active research \citep[e.g.][]{Veras2016}. The general consensus is that minor planets (i.e. asteroids, comets) and giant planets beyond a few au can survive the post main sequence evolution and orbit around white dwarfs \citep{NordhausSpiegel2013, Mustill2014}. Through different dynamical interactions, e.g. mean motion resonance, planet-planet scattering, secular resonance sweeping, and the Kozai-Lidov effect, the orbits of these minor planets are perturbed -- some are ejected from the system while others can enter into the white dwarf's tidal radius ($\sim$100R$_\mathrm{wd}$, \citealt{Debes2012a, Stephan2017, Mustill2018, Smallwood2018}). In addition, there is evidence for continuous accretion of small planetesimals \citep{Wyatt2014}. In this paper, we focus on two systems, SDSS~J1228+1040 and G29-38. Their basic parameters are listed in Table~\ref{tab:WDPar}. SDSS~J1228+1040 is the prototype of white dwarfs with circumstellar gas debris \citep{Gaensicke2006}. Its infrared excess was reported in \citet{Brinkworth2009}. Through 12-yr optical spectroscopic monitoring, \citet{Manser2016a} found a gradual variation of the gas emission lines and they proposed it as a result of precession of an asymmetric pattern under general relativity. SDSS~J1228+1040 is also heavily polluted and the composition of the accreting material resembles that of bulk Earth \citep{Gaensicke2012}. \begin{deluxetable}{lccc} \tablecaption{White Dwarf Parameters \label{tab:WDPar}} \tablewidth{0pt} \tablehead{ \colhead{} & \colhead{SDSS~J1228+1040} & \colhead{G29-38} } \startdata $T_\mathrm{eff}$ (K) & 23510 & 11240 \\ $log$ $g$ (cgs) & 8.16 & 8.00\\ Dom.\tablenotemark{a} & H & H \\ $M_\mathrm{wd}$ ($M_\odot$) & 0.70 & 0.59\\ $R_\mathrm{wd}$ ($R_\mathrm{\odot}$) & 0.012 & 0.013\\ $d$ (pc)\tablenotemark{b} & 127 &17.5\\ $V$ (mag) & 16.2 & 13.0\\ Ref & \citet{Tremblay2011}& \citet{Subasavage2017}\\ \enddata \tablenotetext{a}{Dominant element in the white dwarf's atmosphere.} \tablenotetext{b}{Distance is taken from {\it Gaia DR2} \citep{GaiaDr2}. } \end{deluxetable} G29-38 was the first white dwarf discovered to have an infrared excess and also the first discovered to display a 10~$\mu$m silicate emission feature \citep{Reach2005a, Reach2009}. The star also has a polluted atmosphere \citep{Koester1997} and recent HST/COS observations show that it is accreting from volatile-poor material that is similar in composition to the bulk Earth \citep{Xu2014}. G29-38 is among the first variable white dwarfs discovered \citep{ShulovKopatskaya1974,McGrawRobinson1975}. The newest addition to the wonders of G29-38 comes from the discovery of molecular hydrogen in its atmosphere, which provides an additional constraint of its stellar parameters \citep{Xu2013b}. It is among the hottest stellar environments with a molecular hydrogen detection. Here, we report infrared observations of the dust disks around SDSS~J1228+1040 and G29-38, demonstrating for the first time that these two disks are intrinsically variable. The rest of the paper is organized as follows. Observation and data reduction are presented in Section~\ref{sec:data}. In Section~\ref{sec:model}, we present some disk models that could explain the temporal variations of the infrared luminosity. Possible scenarios are explored in Section~\ref{sec:interp} and results are summarized in Section~\ref{sec:conclusion}. | } In this paper, we presented infrared variabilities of two white dwarfs with dust disks, SDSS~J1228+1040 and G29-38. For SDSS~J1228+1040, the IRAC [3.6] and [4.5] fluxes dropped by 20\% within 7 years and remained the same afterwards. The general behavior is very similar to the flux drop around another dusty white dwarf WD~J0959$-$0200 \citep{XuJura2014} with a further similarity being the appreciable amounts of circumstellar gas around both objects. The flux drop can be explained by either an increase of an inner disk radius, a decrease in the outer disk radius, or a change in the disk inclination, assuming the excess comes from an opaque dust disk. G29-38 appears to represent a different kind of infrared variability and the flux of the 10~$\mu$m feature has increased by 10~\% within 3 years while the 5--7~$\mu$m flux remained the same. We presented a two-component disk model to fit the infrared spectra, and concluded that the change is unlikely to be related to the photospheric pulsation of G29-38 with a static disk. We propose that the most likely cause is an increase in the mass of small grains in the optically thin component. To explain the observed infrared variability, we explored several scenarios, including tidal disruption events, external perturbation, and runaway accretion. Although continuous tidal disruptions of small planetesimals could explain the increased dust mass in G29-38, no satisfactory scenarios can explain the sudden drop of infrared flux for SDSS~J1228+1040 and SDSS~J0959$-$0200. Looking forward, a self-consistent radiative transfer disk model would be valuable in constraining white dwarf disk parameters. In the future, continued photometric monitoring in the infrared and high quality infrared spectroscopy from the {\it James Webb Space Telescope} will be crucial in constructing a complete picture of dust disks around white dwarfs. {\it Acknowledgements} The authors would like to thank the late UCLA Professor Michael Jura for his inspirations and support to study white dwarf disks. We thank S. Kleinman and A. Nitta for helpful discussions about white dwarf pulsation. A.\,B. acknowledges a Royal Society Dorothy Hodgkin Fellowship. J.\,O. acknowledges financial support from the ICM (Iniciativa Cient\'ifica Milenio) via the N\'ucleo Milenio de Formaci\'on Planetaria grant, from the Universidad de Valpara\'iso, and from Fondecyt (grant 1180395). D.\,V. gratefully acknowledges the support of the Science and Technology Facilities Council via an Ernest Rutherford Fellowship (grant ST/P003850/1). T. G. W. wishes to acknowledge funding from a STFC studentship. The research leading to these results has also received funding from the European Research Council under the European Union's Seventh Framework Programme (FP/2007-2013) / ERC Grant Agreement n. 320964 (WDTracer). This work is based in part on observations made with the {\it Spitzer Space Telescope}, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. This publication makes use of data products from the Wide-field Infrared Survey Explorer (WISE), which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics and Space Administration. This publication also makes use of data products from the Near-Earth Object Wide-field Infrared Survey Explorer (NEOWISE), which is a project of the Jet Propulsion Laboratory/California Institute of Technology. NEOWISE is funded by the National Aeronautics and Space Administration. \end{CJK} \software{IRAF \citep{IRAF1,IRAF2}, Mercury \citep{Chambers1999}, Matplotlib \citep{Matplotlib}} | 18 | 8 | 1808.09426 |
1808 | 1808.07111_arXiv.txt | We present four daytime thermal images of Europa taken with the Atacama Large Millimeter Array. Together, these images comprise the first spatially resolved thermal dataset with complete coverage of Europa's surface. The resulting brightness temperatures correspond to a frequency of 233 GHz (1.3 mm) and a typical linear resolution of roughly 200 km. At this resolution, the images capture spatially localized thermal variations on the scale of geologic and compositional units. We use a global thermal model of Europa to simulate the ALMA observations in order to investigate the thermal structure visible in the data. Comparisons between the data and model images suggest that the large-scale daytime thermal structure on Europa largely results from bolometric albedo variations across the surface. Using bolometric albedos extrapolated from \textit{Voyager} measurements, a homogenous model reproduces these patterns well, but localized discrepancies exist. These discrepancies can be largely explained by spatial inhomogeneity of the surface thermal properties. Thus, we use the four ALMA images to create maps of the surface thermal inertia \edit1{and emissivity at our ALMA wavelength}. From these maps, we identify a region of \edit1{either} particularly high thermal inertia \edit1{or low emissivity} near 90 degrees West and 23 degrees North, which appears anomalously cold in two of our images. | \label{sec:intro} Europa's icy surface is marked by fractured, ridged, and chaotic terrain suggestive of a history of geologic activity \citep[e.g.][]{Kattenhorn2009, Prockter2009, Collins2009}. Spectroscopic studies have revealed multiple compositions that reflect the influences of both endogenous geologic processes \citep[e.g.][]{McCord1998, Fischer2015} and exogenous radiolytic processing by Jovian magnetospheric particles \citep[e.g.][]{Carlson1999, Carlson2002}, but the balance of these influences in shaping surface properties is not well understood. Surface temperature measurements can provide an additional window onto both types of processes. Such measurements present perhaps the best means for identifying regions of active geologic activity. Indeed, active hotspots persist at both the ``tiger stripes" of Enceladus \citep{Spencer2006} and the volcanoes of Io \citep{Pearl1982,Spencer1990Io}. In addition, \textit{Cassini} thermal observations of Saturn's moons Mimas and Tethys have shown that temperature measurements can reveal details on the effects of magnetospheric particle bombardment on surface texture \citep{Howett2011Mimas, Howett2012Tethys}. Finally, thermal data can provide insight on diurnal sublimation cycles, impact gardening by micrometeorites, and sputtering from particle impacts, which also affect the surface compositions and morphologies and, thus, the surface thermophysical properties. To date, the only spatially resolved thermal data of Europa were collected by the \textit{Galileo} Photopolarimeter-Radiometer (PPR). These data provided the first brightness temperature maps of the surface and included both daytime and nighttime measurements \citep{Spencer1999}. Modeling efforts using the PPR data have found thermal inertia values \edit1{between 30 -- 140 $J/(m^2 \cdot K \cdot s^{1/2})$,} consistent with a particulate, uncompacted regolith texture unlike that of solid water-ice \citep{Spencer1999,Rathbun2010}. However, the \textit{Galileo} PPR only obtained limited coverage of the surface. Furthermore, from the end of the \textit{Galileo} mission until very recently, subsequent brightness temperature measurements of similar quality and spatial resolution have been impossible to achieve. Recently, however, the Atacama Large Millimeter Array (ALMA) has made the collection of spatially resolved, high-quality thermal datasets possible. Here, we present four ALMA thermal images that together cover the entire surface of Europa at a frequency of 233 GHz (1.3 mm) with a typical linear resolution of $\sim$200 km. Using a global thermal model of Europa, we fit the observations and investigate the nature of thermal structure visible across the surface. | \label{sec:conclusions} We obtained four ALMA thermal observations of Europa, which together cover the entire surface and reveal significant thermal structure. Using a globally homogenous, one-dimensional thermal model and a bolometric albedo map constructed from \textit{Voyager} measurements, we are able to reproduce much of this structure well, indicating that it is primarily a product of bolometric albedo variation across the surface and the passive absorption and re-emission of sunlight. However, despite the similarity of the data and model images, there are localized disagreements, which may indicate variability in the surface thermophysical properties. We examine the possibility that these discrepancies can be explained by local thermal inertia variations and construct a corresponding thermal inertia map, assuming a globally homogenous \edit1{millimeter} emissivity. The map suggests typical values of the surface thermal inertia ranging from 40 to \edit1{300} $J/(m^2 \cdot K \cdot s^{1/2})$, with the lowest thermal inertias on the sub-Jovian hemisphere and the highest between the leading and anti-Jovian hemispheres. \edit1{We also construct a complementary map of emissivity at our ALMA wavelength (1.3 mm), assuming a globally homogenous thermal inertia, which suggests emissivities of 0.67 -- 0.84.} We find little correlation \edit1{of thermal properites} with geology or composition and few noteworthy anomalies, with the exception of an elevated thermal inertia surrounding Pwyll crater and a region of \edit1{low emissivity} or extremely elevated thermal inertia near 90$\degree$W and 23$\degree$N on the leading hemisphere, in a region of relatively low-quality \textit{Galileo} imaging. This leading hemisphere location corresponds to the region spectroscopically determined to be the iciest (and potentially most crystalline) on the surface. However, it does not coincide with any unique geologic or morphological features, nor was it it covered by the \textit{Galileo} PPR. Thus, while we suggest that the area \edit1{is distinct in its thermal properties}, we can only speculate as to its origins. Future ALMA observations will provide measurements of each location on the surface at other times of day, which will allow for better constraint on the surface thermal properties and, thus, their potential influences. | 18 | 8 | 1808.07111 |
1808 | 1808.05114_arXiv.txt | We have analyzed the differences in positions of 9081 matched sources between the Gaia DR2 and VLBI catalogues. The median position uncertainty of matched sources in the VLBI catalogue is a factor of two larger than the median position uncertainty in the Gaia DR2. There are 9\% matched sources with statistically significant offsets between both catalogues. We found that reported positional errors should be re-scaled by a factor of 1.3 for VLBI and 1.06 for Gaia, and in addition, Gaia errors should be multiplied by the square root of chi square per degree of freedom in order to best fit the normalized position differences to the Rayleigh distribution. We have established that the major contributor to statistically significant position offsets is the presence of optical jets. Among the sources for which the jet direction was determined, the position offsets are parallel to the jet directions for 62\% of the outliers. Among the matched sources with significant proper motion, the fraction of objects with proper motion directions parallel to jets is a factor of 3 greater than on average. Such sources have systematically higher chi square per degree of freedom. We explain these proper motions as a manifestation of the source position jitter caused by flares that we have predicted earlier. Therefore, the assumption that quasars are fixed points and therefore, differential proper motions determined with respect to quasar photocenters can be regarded as absolute proper motions, should be treated with a great caution. | Since 1980s very long baseline interferometry (VLBI) has been the most accurate absolute astrometry technique. The accuracy of VLBI absolute positions can reach the 0.1~mas level. With few exceptions, VLBI is able to provide absolute positions of only active galactic nuclei (AGNs). In 2016, the Gaia Data Release~1 (DR1) \citep{r:gaia_dr1} ushered an emergence of the technique that rivals VLBI in accuracy. A quick analysis by \citet{r:gaia_icrf2} found that in general, the differences between common AGNs in VLBI and Gaia DR1 catalogues are close to their uncertainties, except for a 6\% of common objects. \citet{r:gaia_icrf2} claims that ``individual examination of a number of these cases shows that a likely explanation for the offset can often be found, for example in the form of a bright host galaxy or nearby star''. They conclude (page 13) that ``the overall agreement between the optical and radio positions is excellent''. We see it differently. If two independent observing campaigns produced small (negligible) differences, that also implies that the contribution of a new campaign is also small (negligible) with respect to what has been known before. Science does not emerge from agreements. It emerges from disagreements. Therefore, we focused our analysis on differences between VLBI and Gaia AGN positions. Our analysis of Gaia DR1 confirmed the existence of a population of sources with statistically significant VLBI/Gaia offsets \citep{r:gaia1}. We found that such factors as the failures in quality control in both VLBI and Gaia, blended nearby stars, or bright host galaxies can account at maximum for 1/3 of that population. This analysis, as well as recent works of others \citep{r:gaia_icrf2,r:mak17,r:Frouard18,r:Liu18a,r:Liu18b,r:Liu18c}, used arc lengths of VLBI/Gaia differences. Including the second dimension, the position angle of VLBI/Gaia offsets, resulted in a breakthrough. Though the distribution of the position angles counted from the declination axis turned out to be close to uniform, the distribution of the position angles with respect to the jet direction determined from analysis of VLBI images of matched sources revealed a strong anisotropy \citep{r:gaia2}: the offsets have a preferable direction along the jet, and at a smaller extent in the direction opposite to the jet. We interpret it as a manifestation of a presence of optical jets at scales finer than the Gaia point spread function (PSF), i.e., 100--300~mas. Known optical jets in AGNs resolved with Hubble Space Telescope are cospatial \citep{r:gabuzda2006,r:perlman2010, r:meyer18_m84}. Even in that case there will be position differences. It was emphasized in \citep{r:gaia3} that the response to an extended structure of a power detector used by Gaia and an interferometer that records voltage is fundamentally different. The Gaia positions correspond to the location of the optical centroid, while the VLBI positions are associated to the most compact and bright feature at the jet base. Therefore, the physical meaning of a VLBI/Gaia offset is a displacement of the optical centroid with respect to the jet base. In April 2018, the Gaia DR2 was published \citep{r:gaia_dr2}. It has 48\% more sources than Gaia DR1 and a significantly higher accuracy. \mbox{\citet{r:gaia_dr2_crf}} reported that in general, the VLBI/Gaia DR2 differences are small with some exceptions. They set out five reasons for discrepancies (page 10): 1)~real offsets between the centres of emission at optical and radio wavelengths; 2)~error in matching VLBI and Gaia objects; 3)~an extended galaxy around the quasar; 4)~double or lensed quasars; or 5)~simply statistical outliers. The presence of optical jets was not put in the list as a likely explanation. In \citet{r:gaia3} we examined the consequences of our interpretation of the VLBI/Gaia offsets due to the presence of optical jets. Among others, we made two predictions: 1)~``further improvement in the position accuracy of VLBI and Gaia will not result in a reconciliation of radio and optical positions, but will result in improvement of the accuracy of determination of these position differences'', 2)~``we predict a jitter in the Gaia centroid position estimates for radio-loud AGNs''. Since the Gaia DR2 accuracy is noticeably better than the Gaia DR1 accuracy, this motivated us to extend our previous analysis to the Gaia DR2 and check whether these predictions came true. To answer the question what is the most significant contributor to systematic position differences is the goal of this article. | Here we summarize the main results of our comparison of AGN positions and proper motions from the Gaia DR2 against the most complete catalogue of VLBI positions to date, the RFC. \begin{enumerate} \item The Gaia DR2 AGN position uncertainties of VLBI matched sources are a factor of two smaller than the VLBI position uncertainties. \Note{Gaia position catalogues are becoming the most precise astrometry catalogues at present.} \item We predicted in \citet{r:gaia3} that the improvement in accuracy of VLBI and/or Gaia will not reconcile the VLBI and Gaia positions, but will make these differences more significant. This prediction has come true. The fraction of outliers grew from 6 to 9\%, and the distribution of the position offset directions as a function of $\psi$ angle became sharper. \item We demonstrated that the main reason for the statistically significant VLBI/Gaia position offset is the presence of optical structure. Among the matched sources with the normalized arc lengths exceeding 4 that have measured jet directions, 52\%--62\%, i.e., {\it the majority}, have the position offsets parallel to the jet direction. Therefore, we conclude that the optical jet is the cause. Although this fraction may be somewhat lower for the entire population of matched AGNs, we got its firm lower limit: 27\%. Other reasons mentioned by \mbox{\citet{r:gaia_dr2_crf}} can explain only a small fraction of outliers. The presence of emission from a host galaxy within the Gaia point spread function may shift the centroid with respect to the nucleus if the galaxy central region structure is asymmetric or the AGN is dislodged with respect to the galaxy center of mass, \Note{we assume such a shift is independent on jet direction angle in the absence of evidence of such a dependency}. Table~\ref{t:hist_fit} provides the upper limit of the fraction of outliers which position offsets do not depend on $\psi$: 33\%. It does not seem likely that all of these offsets are caused by the contribution of host galaxies, because the fraction of AGNs with discernible host galaxies is much less. \item We found that scaling the Gaia position uncertainties by $\sqrt{\chi^2/\mbox{ndf}}$ eliminated the dependence of the fraction of outliers on $\chi^2/\mbox{ndf}$. Examining the subset of matches with dominating VLBI or Gaia errors allowed us to evaluate the scaling factors for the VLBI uncertainties, 1.30, and the Gaia position uncertainties: $1.06 \, \sqrt{\chi^2/\mbox{ndf}}$. Eliminating the observations within 0.5~rad of $\psi=0$ and $\psi=\pi$ and using re-scaled uncertainties, made the distribution of normalized VLBI/Gaia arc-lengths much closer to the Rayleigh distribution: compare Figures~\ref{f:norm_arc_all} and \ref{f:norm_arc_off}. \item The contribution of VLBI and/or Gaia systematic errors on estimates of the orientation angles of the Gaia DR2 catalogue with respect to the VLBI catalogue does not exceed 0.02~mas. \item We predicted in \citet{r:gaia3} that flares in AGNs would cause a jitter in their positions \Note{because an increase of flux in one of the components of an extended source will change the centroid position}. The analysis of Gaia proper motions provided us an indirect confirmation of this prediction: the sources with excessive Gaia residuals, i.e., large $\chi^2$/ndf, have proper motion directions predominately parallel to the jet directions. The median magnitude of statistically significant proper motions is larger than 1~mas/yr over a 1.16~year interval, which is significantly higher than $<0.05$~mas/yr over 5~years anticipated before the Gaia launch \citep{r:perryman14}. Although AGNs proper motions should not be interpreted as a bulk tangential motion, at the same time, these proper motions are not always artifacts of Gaia data analysis. The photo-centers of at least some quasars are not fixed points and the possibility of quasar proper motion should be taken into account in interpreting results of differential astrometry. \item We found that VLBI proper motions have a preferable direction along with the jet. Median VLBI proper motions of AGNs are a factor of 50 smaller than Gaia proper motions. \end{enumerate} We do not claim that we have solved the problem of establishing the nature of {\it all} outliers. The distribution in Figure~\ref{f:norm_arc_off} still deviates from Rayleigh and we still did not uncover the nature of the 1/3 statistically significant offsets, but we made a quite substantial progress. We anticipate that a study of VLBI/Gaia position offsets will become a power tool for probing properties of the accretion disk and the relativistic jet in the AGNs, in line with the work of \citet{r:gaia5}. | 18 | 8 | 1808.05114 |
1808 | 1808.07920_arXiv.txt | \label{intro} Hierarchical multi-body star systems (Evans 1968) form from different ways, such as from interaction/capture in a globular star cluster (van den Berk \etal 2007), from a massive primordial disk involving accretion processes and/or local disk instabilities (Lim and Takakuwa 2006; Marzari \etal 2009) or from a non-hierarchical star system by angular momentum and energy exchange via mutual gravitational interactions (Reipurth 2000). These systems can be basically classified into two groups; circumbinary and circumstellar systems. In circumbinary systems, one or more additional bodies move around a binary star and they are known as companions on P-type orbits (Dvorak 1986). A transiting circumbinary planet, PH1b, around \hbox{{KIC 4862625}} which consists of two binary pairs; the quadruple systems \hbox{{HD 98800}} (Furlan \etal 2007) and {SZ Her} (Lee \etal 2012) can be given as examples of such a hierarchy. On the other hand, the systems with companions orbiting one component of a binary pair are the other type of hierarchical systems (circumstellar or S-type configuration; Schwarz \etal 2011). The example of such a system can be found in Neuhäuser \etal (2007) and Chauvin \etal (2007). A hierarchical circumbinary system can be detected by observing the timings of the mid-eclipse times of the binary companion. The presence of an additional body causes a change in the relative distance of the eclipsing pair to the observer depending on the motion of the third body around the barycenter of the triple system. This binary wobble leads a periodic variation in conjunction times. As a result, the eclipses present lags or advances in the timings of minimum light (Irwin 1952). As known, the light-time effect is a geometrical feature and the third object produces a sinusoidal-like variation in the binary orbital. If the archival database is large and sufficient enough, this variation in eclipse timings provides an opportunity to understand the nature of the multi-body system (Pribulla \etal 2012). In this sense, space-based missions offer a unique opportunity for the discovery of companions orbiting eclipsing binaries. For example, Kepler provides continuous and highly homogeneous light curves over the time interval of four years. Thus, its photometric observations enable new discoveries of multiple star systems, such as triple, quadruple or even quintuple ones. Indeed, there are a large number of multiple star systems identified from the Kepler observations. Conroy \etal (2014) present a catalog, which includes precise minimum times and third body signals for 1279 close binaries in the latest Kepler Eclipsing Binary Catalog. They find 236 binaries having third body signals. Borkovits \etal (2015) report {\sl O-C} analysis of 26 compact hierarchical triple stars in the Kepler field. Borkovits \etal (2016) identify the existence of a third body in 222 of 2600 Kepler binaries. The quadruple system KIC 7177553 (Lehmann \etal 2016) consists of two eccentric binaries with a separation of 0.4 arcsec (167 au). The outer orbit's period is in the range of 1000-3000 yr. Another quadruple star system EPIC 220204960 contains two slightly eccentric binaries with orbital periods of 13.27 and 14.41 days (Rappaport \etal 2017). These binaries are in a quadruple system with an outer period of 1 yr and a physical separation of ≤ 30 au. An example for a quintuple star system is EPIC 212651213 and EPIC 212651234 (Rappaport \etal 2016). In this system, EPIC 212651213 hosts two eclipsing binaries with orbital periods of 5.1 and 13.1 days. EPIC 212651234 is a single star with a projected physical separation of about 0.013 pc to EPIC 212651213. It is also stated that EPIC 212651213 and EPIC 212651234 are gravitationally bound to each other. {DI Peg} ({HIP 116167}, {GSC 01175-00013}, {BD+14 5006}) was discovered by Morgenroth (1934) and identified to be an Algol type eclipsing binary (F4IV+ K4) by Rucinski (1967) and Lu (1992). From the photographic observations, Jensch (1934) determined the period of the system to be \hbox{$\sim 0_{\cdot}^\mathrm{d}711811$}. Rucinski (1967) analyzed the photoelectric observations of Kruszewski (1964) and derived the first orbital solutions. Based on the results, he suggested the existence of a third light which provided 24$\%$ contribution to the total light of the system. More photometric studies were performed by Chou and Kitamura (1968), Binnendijk (1973), Chaubey (1982), Lu (1992), and Yang \etal (2014). Gaposchkin (1953) detected a variation in the orbital period of the star. Ahnert (1974) and Vinkó (1992) proposed a possible light-time effect in the system and they gave periods of $\sim62.4$ and $\sim22.1$ yr. By using the spectroscopic solutions, Lu (1992) determined the system parameters as $a = 4.14(0.05)$ R$_\odot$, $V_0 = +43.8(2.0)$ \kms, $K_1$ = 185.2(2.4) \kms, $K_2 = 109.0(2.1)$ \kms, $T_0$ = HJD 48213.8851(0.0022) and $q_\mathrm{sp}$ = 0.59(0.01). Rafert (1982) derived a downward quadratic ephemeris with a cyclic variation in the {\sl O-C} diagram. Unlike this, Hanna and Amin (2013) obtained a cyclic modulation with the period of 55 years, superimposed on an upward parabolic variation. The long-term orbital period increase was found to be $dP/dt = 0.17$ s/century and interpreted as a mass transfer from the evolved secondary component to the primary one with the rate of $1.52 \times 10^{-8}$ M$_\odot$/yr. The cyclic variation was attributed to a low-mass third body with the mass of $M_3 \sim 0.2200 \pm 0.0006$ M$_\odot$. The parameters of the third body were given as $e_3 = 0.77(7)$ and $w_3 = 300^{\circ} \pm 10^{\circ}$. Recently, Yang \etal (2014) reproduced the photometric models by the help of new multi-color observations and previously published ones in literature. They determined the system parameters as $i = 89_{\cdot}^\mathrm{\circ}02\pm 0_{\cdot}^\mathrm{\circ}11$, $M_1 = 1.19(2)$ M$_\odot$, $M_2$ = 0.70(2) M$_\odot$, $L_1$ = 3.70(4) L$_\odot$, and $L_2$ = 0.53(2)~L$_\odot$. According to the results, they stated that the system had a low third light whose fill-out factor for the more massive component was $f_\mathrm{p} = 78.2(4)$. Their {\sl O-C} curve also indicated that the orbital period of DI\,Peg has changed in the complicated mode, such that the period of the star possibly showed two light-time orbits with the modulation periods of $P_3 \sim 54.6(5)$ yr and $P_4 \sim 23.0(6)$ yr, respectively. The masses of the inner and outer sub-stellar objects were given to be $M_{in}\sim 0.095$ M$_\odot$ and $M_{out}\sim 0.170$ M$_\odot$. On the basis of these results, Yang \etal (2014) suggested that the system has consisted of four objects. The aim of this study is to perform a detailed period analysis of {DI Peg} for the parameter determination of the additional bodies in the system by using the new and all available archival minimum times. For this purpose, the paper is organized as follows; the observations are presented in Section~\ref{observations}, the analysis is described in Section~\ref{analysis}, the results related to the analysis are discussed in Section~\ref{conclusion}. | 18 | 8 | 1808.07920 |
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1808 | 1808.07327_arXiv.txt | The Galactic centre supermassive black hole (SMBH), in sharp contrast with its complex environment, is characterized by only three classical parameters -- mass, spin, and electric charge. Its charge is poorly constrained. It is, however, usually assumed to be zero because of neutralization due to the presence of plasma. We revisit the question of the SMBH charge and put realistic limits on its value, timescales of charging and discharging, and observable consequences of the potential, small charge associated with the Galactic centre black hole. The electric charge due to classical arguments based on the mass difference between protons and electrons is $\lesssim 10^9\,{\rm C}$ and is of a transient nature on the viscous time-scale. However, the rotation of a black hole in magnetic field generates electric field due to the twisting of magnetic field lines. This electric field can be associated with induced charge, for which we estimate an upper limit of $\lesssim 10^{15}\,{\rm C}$. Moreover, this charge is most likely positive due to an expected alignment between the magnetic field and the black-hole spin. Even a small charge of this order significantly shifts the position of the innermost stable circular orbit (ISCO) of charged particles. In addition, we propose a novel observational test based on the presence of the bremsstrahlung surface brightness decrease, which is more sensitive for smaller unshielded electric charges than the black-hole shadow size. Based on this test, the current upper observational limit on the charge of Sgr~A* is $\lesssim 3\times 10^{8}\,{\rm C}$. | The observations of the Galactic centre across the electromagnetic spectrum, ranging from radio to gamma wavelengths, revealed the complex structure of the Nuclear Star Cluster (NSC) as well as that of the gaseous-dusty medium of the central parsec \citep{2005bhcm.book.....E,2007gsbh.book.....M,2010RvMP...82.3121G,2017FoPh...47..553E}. The presence of the concentrated, dark mass at the dynamical centre of the NSC was revealed by the near-infrared observations of stars using adaptive optics. The first proof for the compact dark single object in the Galactic centre came with the detection and the analysis of the first proper motion of stars orbiting Sgr~A* inside $1'' \sim 0.04\,\rm{pc}$, so-called S stars \citep{1996Natur.383..415E,1997MNRAS.284..576E,1998ApJ...509..678G, 2018arXiv180411014Z}. Based on these and follow-up observations \citep{2016ApJ...830...17B,2017ApJ...845...22P,2009ApJ...692.1075G,2017ApJ...837...30G,2018A&A...615L..15G}, the large mass of the dark object has been confirmed and a more precise value has been determined -- $\sim (4.15 \pm 0.13 \pm 0.57)\times 10^6\,M_{\odot}$ \citep{2017ApJ...845...22P}. If we associate this dark mass with a non-rotating black hole for simplicity, this yields a Schwarzschild radius of $R_{\rm Schw}=1.2 \times 10^{12}\,{\rm cm}(M_{\bullet}/4\times 10^6\,M_{\odot})$ and the expected mean density is, \begin{equation} \rho_{\bullet}=1.7 \times 10^{25} \left(\frac{M_{\bullet}}{4\times 10^6 \, M_{\odot}}\right) \left(\frac{R_{\rm Schw}}{3.9 \times 10^{-7}\,{\rm pc}} \right)^{-3}\,M_{\odot}{\rm pc^{-3}}\,. \label{eq_density_blackhole} \end{equation} In case of stellar orbits, the tightest constraint for the density of the dark mass comes from the monitoring of B-type star S2 with the pericentre distance of $r_{\rm P} \simeq 5.8 \times 10^{-4}\,{\rm pc}$ \citep{2017ApJ...845...22P,2002Natur.419..694S,2009ApJ...692.1075G,2017ApJ...837...30G,2018A&A...615L..15G} \begin{equation} \rho_{\rm S2}=5.2 \times 10^{15} \left(\frac{M_{\bullet}}{4.3 \times 10^6\,M_{\odot}}\right)\left( \frac{r_{\rm P}}{5.8 \times 10^{-4}\,{\rm pc}} \right)^{-3}\,M_{\odot}{\rm pc^{-3}}\,. \label{eq_density_blackholeS2} \end{equation} The most stringent density constraint was given by 3$\sigma$ VLBI source size of $\sim 37\mu {\rm as}$ \citep{2008Natur.455...78D,2018ApJ...859...60L}. When combined with the lower limit on the mass $M_{\rm SgrA*} \gtrsim 4 \times 10^5\,M_{\odot}$ based on the proper motion measurements \citep{2004ApJ...616..872R}, VLBI yields the lower limit of $\rho_{\rm SgrA*} \geq 9.3\times 10^{22}\,M_{\odot}{\rm pc^{-3}}$. This is about two orders of magnitude less than the density expected for a black hole of $\sim 4\times 10^6\,M_{\odot}$, see Eq.~\eqref{eq_density_blackhole}. The most plausible stable configuration that can explain such a large concentration of mass emerges within the framework of general relativity: a singularity surrounded by an event horizon -- a black hole, ruling out most of the alternatives \citep{2017FoPh...47..553E}. According to the uniqueness or the general relativistic ``no-hair" theorem \citep{1996bhut.book.....H}, any stationary black hole is fully characterized by only three classical and externally observable quantities: mass $M_{\bullet}$, angular momentum $J_{\bullet}$ (often the quantity $a_{\bullet}=J_{\bullet}/M_{\bullet}c$ is used which has a dimension of length), and the electric charge $Q_{\bullet}$\footnote{In case a magnetic monopole could exist, it could be the forth parameter.}. Thanks to the high-precision observations of stars in the Nuclear Star Cluster, including the innermost S cluster, the current value for the SMBH mass is $M_{\bullet}=(4.3 \pm 0.3) \times 10^6\,M_{\odot}$ \citep{2017FoPh...47..553E}, which is based on different methods, primarily the orbits of S stars \citep{2017ApJ...845...22P}, the Jeans modelling of the properties of the NSC \citep{2013ApJ...779L...6D}, and the general relativistic fits to the double-peaked X-ray flares that show signs of gravitational lensing \citep{2017MNRAS.472.4422K}. The constraints for the spin $J_{\bullet}$ were inferred indirectly based on the variable total and polarized NIR emission \citep{2006A&A...455....1E}. The spin can be determined based on the modelling of spin-dependent quantities, mainly the light curves of a hot spot or a jet base. In this way, \citet{2006A&A...460...15M} obtained constraints for the spin, which are rather weak and the spin parameter is $a_{\bullet} \gtrsim 0.4$, as well as the inclination, which is inferred based on the stable polarization angle of the flares and tends to be rather large $i \gtrsim 35^{\circ}$. The value of the spin parameter determined based on quasi-periodic oscillations for Sgr~A* reaches a unique value of $\approx 0.44$ \citep{2010MNRAS.403L..74K}, which is consistent with the value inferred from the fitting of the NIR flares. In general, the charge of the black hole $Q_{\bullet}$ is often set equal to zero due to the presence of plasma around astrophysical black holes. However, a black hole can acquire primordial charge because it was formed by a collapse of a charged (compact) star \citep{2003PhRvD..68h4004R}. It is not clear on which timescales such a charged black hole discharges or alternatively, can increase its charge. Also, from an astrophysical point of view, it is of a general interest if a charged black hole can be observationally distinguished from a non-charged case, clearly depending on the value of the charge. In addition, electric charge can be loaded or induced by black hole due to its rotation in external magnetic field within the mechanism similar to the Faraday unipolar generator. Such a mechanism is more relevant for supermassive black holes in the local Universe, since the primordial charge information is expected to be lost. The induction mechanism works in such a way that the rotation of a black hole generates electric potential between horizon and infinity which leads to the process of selective accretion of charged particles of plasma surrounding the black hole. In particular, a rotating black hole embedded in a uniform, aligned magnetic field will acquire an electric charge until an equilibrium value is reached $Q_{\bullet, W}=2B_{\rm 0} J_{\bullet}$, a so-called Wald charge \citep{1974PhRvD..10.1680W}, where $B_0$ is an asymptotic magnetic field strength. There is an evidence that significant and highly aligned magnetic field must be present in the Galactic center with equipartition strength of $10\,{\rm G}$ in the vicinity of the event horizon of the SMBH \citep{2012A&A...537A..52E,2013Natur.501..391E,2015llg..book..391M}. The twisting of magnetic field lines threading the horizon of rotating black hole produces an electric field which accelerates the charged particles along the magnetic field lines. Moreover, magnetic field plays the role of a catalyzing element that allows the extraction of rotational energy from rotating black hole through interaction of charged particles with an induced electric field in such processes as the Blandford-Znajek mechanism \citep{1977MNRAS.179..433B} and the magnetic Penrose process \citep{1985ApJ...290...12W}. Both of these processes that allow the energy extraction from rotating black holes require the presence of an induced electric field \citep{2018MNRAS.478L..89D}. Even a small charge associated with the black hole can have considerable effects on the electromagnetic processes in its vicinity, such as the bremsstrahlung emission and the motion of charged particles as we will show. The value of this small electric charge for black holes embedded in plasma will be necessarily temporary and fluctuating, mainly due to the attraction of oppositely charged particles and/or the variability of the magnetic field in which the black hole is immersed. Even for an extreme case of a charged black hole in vacuum, a spontaneous loss of charge would occur due to pair production with an exponential time-dependency \citep{1975CMaPh..44..245G}. In this paper, we revisit the question of a charge, mainly of an electric origin, associated with the Galactic centre SMBH. Previously, several theoretical studies have focused on the spacetime structure of charged black holes \citep{1991JPhy1...1.1005K,1991JMP....32..714K,2011PhRvD..83j4052P,2011PhRvD..84h4002K}. Here we are aiming at the connection between the current theoretical knowledge with a real astrophysical case -- Sgr~A* supermassive black hole, for which we gathered most constraints on its nearby plasma environment \citep{2017FoPh...47..553E} -- in order to put realistic constraints on electric charge of our nearest supermassive black hole. The study is structured as follows. In Section~\ref{section_prospects} we analyse the potential for charging given the plasma properties in the surroundings of Sgr~A*. Subsequently, we put constraints on the charge of the Galactic centre black hole in Section~\ref{limits_charge}, including different processes that can induce charge and change its value, namely accretion of charged matter and the induction mechanism based on the black hole rotation in the magnetic field. In Section~\ref{section_observable_effects}, we focus on possible observational consequences of the charged SMBH, specifically the effect of charge on the black hole shadow size, the bremsstrahlung brightness profile, and the position of the innermost stable orbits of charged particles. We summarize the charge constraints in Section~\ref{section_summary_discussion}, where we discuss additional effects of the black-hole rotation and a potential non-electric origin of the charge. Finally, we conclude with Section~\ref{section_conclusions}. | \label{section_conclusions} We performed analytical calculations to find out if the supermassive black hole at the Galactic centre can get charged and what the realistic values of its charge are. Based on the classical estimates and total amount of the charge in the sphere of influence of the black hole, we expect that the black hole can acquire a small, transient positive charge of $\lesssim 10^9\,{\rm C}$, which does not have an influence on the spacetime metric. Based on the general relativistic calculations, we further explore the induced charge based on the rotating black hole that is immersed in the external magnetic field. Such a configuration is in general expected in almost all galactic nuclei. If the black hole spin axis is approximately aligned with the external magnetic field, we again expect that the induced charge is positive, with the uppermost limit of $Q_{\bullet}\lesssim 10^{15}\,{\rm C}$. Although the spacetime metric is not influenced significantly by electric charge within the limits we found, even such a small charge can significantly influence the typical viscous timescales for protons and electrons as well as an infall of small charged particles (dust particles). Most importantly, for like charges of test particles and Sgr~A*, the ISCO shifts significantly in comparison with the no-charge case even for a small charge of the order of $10^{6}$--$10^{10}\,{\rm C}$, which effectively mimics the black hole spin of $a_{\bullet}\sim 0.6$. This effect should be taken into account in the future numerical calculations as well as the analysis of observation data. We also revisited observational tests of the presence of the charge for the Galactic centre black hole. The black shadow size, which was proposed previously, is only sensitive for large values of the charge, close to an extremal value, which are unrealistic as we showed. We propose a new test based on the observed surface brightness profile of the thermal bremsstrahlung inside the innermost $10^5$ Schwarzschild radii, which is the region that coincides with the S cluster. Within this range, a flattening and a decrease in the bremsstrahlung surface brightness is expected to occur due to the presence of the charged, unshielded black hole starting with about twenty orders of magnitude smaller values than the extremal case. Since the Chandra X-ray observations with the angular resolution of $0.5''$ did detect a weak indication of the drop in the brightness profile at $R_{\rm proj}\lesssim 0.4''$, it puts an observational upper limit on the charge $Q_{\rm SgrA*}\lesssim 3\times 10^{8}\,{\rm C}$. | 18 | 8 | 1808.07327 |
1808 | 1808.05052_arXiv.txt | We refine the gap size measurements of the disk surrounding the Herbig Ae star HD\,100546 in the N band. Our new mid-infrared interferometric (MIDI) data have been taken with the UT baselines and span the full range of orientations. The correlated fluxes show a wavy pattern in which the minima separation links to a geometrical structure in the disk. We fit each correlated flux measurement with a spline function, deriving the corresponding spatial scale, while assuming that the pattern arises interferometrically due to the bright emission from the inner disk and the opposing sides of the wall of the outer disk. We then fit an ellipse to the derived separations at their corresponding position angles, thereby using the observations to constrain the disk inclination to $i=$47\,$\pm$\,1$^{\circ }$ and the disk position angle to $PA=$135.0\,$\pm$\,2.5$^{\circ}$ East of North, both of which are consistent with the estimated values in previous studies. We also derive the radius of the ellipse to 15.7\,$\pm$\,0.8\,au. To confirm that the minima separations translate to a geometrical structure in the disk, we model the disk of HD\,100546 using a semi-analytical approach taking into account the temperature and optical depth gradients. Using this model, we simultaneously reproduce the level and the minima of the correlated fluxes and constrain the gap size of the disk for each observation. The values obtained for the projected gap size in different orientations are consistent with the separation found by the geometrical model. | \label{sec:intro} The protoplanetary disks surrounding Young Stellar Objects (YSOs) evolve by various physical mechanisms such as photoevaporation \citep[]{1994ApJ...428..654H,2006MNRAS.369..229A,2011ARA&A..49..195A}, accretion/ejection \citep[]{1974MNRAS.168..603L,2011ARA&A..49..195A,2014prpl.conf..475A} and dust growth and planet formation \citep[]{2007MNRAS.377.1324C,2014prpl.conf..497E,2014prpl.conf..339T}. These processes can occour at the same time with timescales of up to a few million years and can cause substantial evolution in the disks such as clearing disk cavities. In this case, the disks are named the transitional disks. The first transitional disk was recognized as a subclass of protoplanetary disks by \citet{1989AJ.....97.1451S} and \citet{1989AJ.....98.1409S} from near-infrared (NIR) ground-based photometry and IRAS (MIR) photometry. Detailed studies of these objects became feasible as progressively more sophisticated instruments, the Spitzer Space Telescops \citep{2004ApJS..154....1W} identified a new class of disks called pre-transitional disks around stars LKCa 15 and UX Tau A \citep{2007ApJ...670L.135E}. The spectral energy distribution (SED) of such disks, which provide indirect evidence of gaps have lack of MIR flux (5--20 $\mu$m ) and a significant excess at NIR (2--5 $\mu$m) and longer wavelengths compared to the full protoplanetary disks. The specified shapes of these SEDs demonstrates that pre-transitional disks are comprised of an inner disk and an outer disk separated by a gap. One of the proposed formation scenarios of the gapped disks are the disk-planet interaction while the geometrical structure of such gaps are still under debate. HD\,100546 \citep[110\,$\pm$\,0.6\,pc from the second Gaia Data Release,][]{2016A&A...595A...1G,2018arXiv180409365G} hosts a widely studied pre-transitional disk, which was first identified by \citet{2003A&A...401..577B} modeling the spectral energy distribution (SED). They concluded that HD\,100546 has a small inner disk of a few au and a more massive outer disk beyond $\sim$\,10\,au. They speculated about the presence of a planetary companion in the gap. The gas and dust-rich disk of HD\,100546 is one of the best laboratories to study forming planetary systems. Direct and indirect observations reveal that the disk hosts two embedded companion candidates \citep[e.g.][]{2013ApJ...766L...1Q, 2015ApJ...807...64Q,2003A&A...401..577B,2014ApJ...791..136B,2015ApJ...814L..27C}. The geometrical structures of the disk around HD\,100546 were widely studied in multi-wavelength observations. The results of the various studies are listed in Table\,1. Please note that in the literature the distance of 97\,pc is commonly used. In this work, we use the updated distance of 110\,pc, measured by GAIA. As a complementary study, in this paper, we present our mid-infrared interferometric data, consisting of 15 independent two-element baseline measurements obtained with the Very Large Telescope Interferometer (VLTI) and its instrument MIDI \citep{2003Msngr.112...13L}. With a large span of the projected position angles, these data provide a unique measurements of the gap geometry in the mid-infrared in unprecedented detail. It is unique since we have multi-epoch N-band data, which has not been done before. These data give us more information about the geometry of the disk than previous N-band observations on HD 100546. Constraining the precise size of the gap of HD\,100546 at various wavelengths is important to learn about the processes at work that form this gap. Moreover, HD\,100546 exhibits multiple signposts of planets and will be a prime target for all future instruments. | We have presented multi-epoch interferometric N-band data with a good uv-coverage obtained with VLTI/MIDI. The measured correlated fluxes are characterized by a wavy pattern, which varies with baseline length and orientation. However, some minima are less significant than others, which may or may not be related to the silicate feature that is present in the total N band fluxes \citep{2005A&A...437..189V}. Our multi-epoch data allows us to estimate the gap size of the disk of HD\,100546 in the N band. We show that the separation obtained by direct measurement of the separation of the minima of our new data in 2012 and 2013 are satisfactorily fitted by an ellipse. We derive an inclination of 47\,$\pm$\,1$^{\circ }$ and a PA of 135\,$\pm$\,2.54$^{\circ}$, which are consistent with the values estimated by previous studies. We also constrain gap size to 15.7\,$\pm$\,0.8\,au. We show by our model that the inner radius of the outer disk estimated for each observation is consistent with the separation obtained from the geometrical model. Please note that our projected gap size estimation takes into account the updated distance of HD\,100546 measured by GAIA, $\sim$\,13$\%$ larger than the distance of 97\,pc commonly used in literature. Taking into account this fact, our projected gap size estimation bears striking similarity to the multi-wavelengths values found by previous studies listed in Table\,1. However, one has to be careful in comparing the sizes, since the different wavelength regimes and observing methods trace potential different structures. Reconstructing images of the disk of HD\,100546 at L and N band with the upcoming second-generation VLTI instrument MATISSE \citep[the Multi AperTure mid-Infrared SpectroScopic Experiment,][]{2006SPIE.6268E..0ZL} will constitute a unique perspective to further assess the nature of the geometrical structure of HD\,100546 and will be a natural continuation for our study. MATISSE will recombine up to four telescopes. More detailed models including also our data obtained with the Auxiliary Telescopes (ATs) will be presented in a forthcoming paper. | 18 | 8 | 1808.05052 |
1808 | 1808.10856_arXiv.txt | We discuss prospects of identifying and characterizing black hole (BH) companions to normal stars on tight but detached orbits, using photometric data from the \textit{Transiting Exoplanet Survey Satellite} (\tess). We focus on the following two periodic signals from the visible stellar component: (i) in-eclipse brightening of the star due to gravitational microlensing by the BH (self-lensing), and (ii) a combination of ellipsoidal variations due to tidal distortion of the star and relativistic beaming due to its orbital motion (phase-curve variation). We evaluate the detectability of each signal in the light curves of stars in the \tess\ input catalog, based on a pre-launch noise model of \tess\ photometry as well as the actual light curves of spotted stars from the prime \kepler\ mission to gauge the potential impact of stellar activity arising from the tidally spun-up stellar components. We estimate that the self-lensing and phase-curve signals from BH companions, if exist, will be detectable in the light curves of effectively $\mathcal{O}(10^5)$ and $\mathcal{O}(10^6)$ low-mass stars, respectively, taking into account orbital inclination dependence of the signals. These numbers could be large enough to actually detect signals from BHs: simple population models predict some 10 and 100 detectable BHs among these ``searchable" stars, although the latter may be associated with a comparable number of false-positives due to stellar variabilities and additional vetting with radial velocity measurements would be essential. Thus the \tess\ data could serve as a resource to study nearby BHs with stellar companions on shorter-period orbits than will potentially be probed with \gaia. | Because the most massive stars end their lives as black holes (BHs), stellar mass BHs should exist ubiquitously: the stellar mass function suggests that about $10^8$ BHs exist in the Galaxy \citep[e.g.,][]{1994ApJ...423..659B}, and the nearest ones are expected to be within $\mathcal{O}(10)\,\mathrm{pc}$ \citep{1983bhwd.book.....S, 2003ApJ...596..437C}. Nevertheless, only a part of them have been probed via X-ray/radio emissions from interacting binaries with stellar companions (X-ray binaries) or from pulsars, which are presumably observable only in a fraction of the systems' lifetime and/or the parameter space of such binaries. The nearest known BH in an X-ray binary is about $1\,\mathrm{kpc}$ away \citep{2010ApJ...710.1127C, 2018arXiv180411349G}, and we likely miss many nearby BHs. A larger population of stellar mass BHs can be probed if we have a means to search for the quiescent systems with stellar companions on wider, \textit{detached} orbits. They do not only help completing the census of nearby compact objects, but the visible companions allow for reliable measurements of BH mass and kinematics of the system in the Galaxy, which are both direct probes of the mass loss and kick during the supernova (SN) explosion \citep[e.g.,][]{2017hsn..book.1499C}. If they are in tight orbits, we may also learn how the outcomes of binary interactions depend on the systems' property. In this aspect, elemental abundance of the stellar companion also helps to probe the signature of mass exchange. The information will be useful to understand the formation of compact objects and those in close binaries, such as X-ray binaries \citep[e.g.,][]{2006ARA&A..44...49R} and merging BH binaries % \citep[e.g.,][]{2016PhRvL.116f1102A, 2016PhRvX...6d1015A}. Detached BH companions of normal stars can be searched using the same techniques to identify unresolved binaries. The spectroscopic (i.e., radial velocity) search has been considered since 1960s \citep[e.g.,][]{1966SvA....10..251G}, and has recently identified massive, yet dark companions to stars both in a cluster \citep{2018MNRAS.475L..15G} and in the field \citep{2018arXiv180602751T}, whose minimum masses imply that they are BHs or massive neutron stars (NSs). The potential of \gaia\ astrometry has also been discussed extensively \citep{2017MNRAS.470.2611M, 2017ApJ...850L..13B, 2018ApJ...861...21Y, 2018MNRAS.481..930Y}, and hundreds or thousands of BHs may be found by the end of its five-year mission. The typical targets will be $\sim10\,M_\odot$ BHs in au-scale binaries. This paper focuses on all-sky photometry as another means to search for stars with BH/NS companions: we especially consider the potential of the {\it Transiting Exoplanet Survey Satellite} \citep[\tess,][]{2014SPIE.9143E..20R} to identify and characterize such binaries on tight orbits ($\lesssim0.3\,\mathrm{au}$). To achieve its main science goal to find transiting exoplanets around nearby stars, \tess\ will provide photometric light curves with sub-precent precision and at least 27-day long, for $>10^7$ stars in the almost entire sky \citep{2015ApJ...809...77S}. The number could be sufficiently large to find such rare binaries with compact objects, as shown below. We discuss two methods that have successfully identified white dwarf (WD) companions to normal stars in the photometric data from the \kepler\ mission \citep{2009Sci...325..709B}: (i) periodic brightening due to in-eclipse microlensing known as ``self-lensing" \citep{2014Sci...344..275K, 2018AJ....155..144K, 2019arXiv190707656M}, and (ii) phase-curve modulation due to ellipsoidal variations and Doppler beaming \citep[e.g.,][]{2007ApJ...670.1326Z, 2010A&A...521L..59M, 2015ApJ...815...26F}. Those BH--star binaries, if detected with \tess, would necessarily be nearby systems amenable to various follow-up observations, % with their short-period and repeating signals being ideal for detailed characterization. In particular, the self-lensing BHs, if detected, provide unambiguous evidence for their compact nature, which is otherwise difficult to confirm. They are also complementary to the BHs detectable with \gaia: they have shorter orbital periods and will be a more sensitive probe of the post-interaction systems that likely followed similar formation paths to X-ray binaries or merging BHs. In the following, we estimate how many \tess\ stars will be searchable for BH/NS companions with given masses and orbital periods (Section \ref{sec:searchable}). This information will then be combined with simple population models of BH--star binaries to estimate the number of detectable BHs around those searchable stars (Section \ref{sec:yields}). The estimated yields are checked against the known population of X-ray binaries in Section \ref{sec:xb}. We summarize and discuss future prospects in Section \ref{sec:summary}. | \label{sec:summary} We estimated that the self-lensing and phase-curve signals induced by BH companions on tight ($\lesssim0.3\,\mathrm{au}$) but detached orbits will be strong enough to be detectable in the \tess\ light curves of effectively $\sim 10^5$ and $\sim 10^6$ stars, respectively (Figure \ref{fig:ns}), taking into account inclination dependence of the signals. The detectability of these signals were evaluated using the stellar properties in the \tess\ input catalog \citep{2017arXiv170600495S}, \tess\ noise model in \citet{2015ApJ...809...77S}, and the light curves of spotted stars from the prime \kepler\ mission \citep{2014ApJS..211...24M} to gauge the impact of stellar activities. If combined with simple models for the population of detached BH--star binaries (Figure \ref{fig:bhpops}), these ``searchable" stars are expected to host $\sim 10$ and $\sim100$ detectable BH companions (Figures \ref{fig:yield} and \ref{fig:yield_ce}), although we cannot exclude the possibility that the latter may be associated with a comparable number of false-positives due to stellar activities. The most promising targets turned out to be BHs with masses of a few $10\,\msun$ and orbital periods of a few days. This large population of BHs, if identified, will reveal the BH mass function down to $\sim1\,\msun$ range, semi-major axis/eccentricity distribution of their orbits, positions and velocities of the systems in the Galaxy, and chemical compositions of the companion stars. These constraints will be valuable to probe mass ejection and natal kick during the BH formation, as well as the binary evolution process that may result in the observed close-in orbits. Since non-zero detections are expected from our models, even the null detection, if quantified, will provide critical information on the models of interacting binaries containing BHs. Although we have focused on systems with BHs, the \tess\ light curves are also sensitive to NS companions down to $\sim1\,\msun$ (Figure \ref{fig:ns}). We may also detect phase-curve signals from optical counterparts of X-ray binaries in a quiescent phase, which potentially allow for better characterization of both known and unknown X-ray binaries. Suppose that candidate BH companions are identified from the \tess\ photometry, what needs to be done next? The self-lensing signal, if identified, will provide the least ambiguous targets that are also best suited for further characterization: the precise orbital inclination and period from the light curve allow for the precise mass determination with radial velocity measurements. The stellar radius estimate using the \gaia\ parallax, if available, even allows for mass determination with light curves alone; because the pulse height constrains $\mbh/\rs^2$, the BH mass can be determined at least to the precision of $20\%$, or even better if spectroscopic effective temperature of the companion is available. The candidates identified with the phase-curve signals would require further vetting with follow-up spectroscopy to confirm their ``SB1" nature and to measure spectroscopic orbits. While the candidates identified from either method will most likely be bright enough for follow-up spectroscopy, the archival data from large spectroscopic surveys, such as APOGEE \citep{2018ApJS..235...42A}, RAVE \citep{2017AJ....153...75K}, LAMOST \citep{2012RAA....12.1197C}, and GALAH \citep{2018MNRAS.478.4513B}, will also play an essential role to complement the \tess\ photometry, both in terms of target vetting and dynamical/chemical characterization. Indeed, those archival data alone might even provide sufficient information to confirm or reject some of the candidates. Eventually, \gaia\ will also provide further information based on astrometric orbits and/or multi-epoch radial velocity measurements. There are also other classes of objects that can be searched with similar methods. This includes BH/NS companions of WDs, which were not considered in this paper because the eclipse probability is small, timescales of the detectable signals may be too short for the $30$-minute cadence photometry, and most of them are likely too faint for \tess. Nevertheless, they may still be good targets for all-sky photometric surveys from the ground \citep[cf.][]{2002A&A...394..489B}. If the \tess\ mission is extended, and/or for stars in the overlapping regions of observing sectors, BH/NS companions on longer-period orbits around evolved stars, as identified in \citet{2018arXiv180602751T} and \citet{2019ApJ...872L..20G}, will also be within reach. | 18 | 8 | 1808.10856 |
1808 | 1808.07057_arXiv.txt | {We examine the spatial distribution and mass segregation of dense molecular cloud cores in a number of nearby star forming regions (the region L1495 in Taurus, Aquila, Corona Australis, and W43) that span about four orders of magnitude in star formation activity. We use an approach based on the calculation of the minimum spanning tree, and for each region, we calculate the structure parameter \Q\ and the mass segregation ratio $\Lambda_{\rm MSR}$ measured for various numbers of the most massive cores. Our results indicate that the distribution of dense cores in young star forming regions is very substructured and that it is very likely that this substructure will be imprinted onto the nascent clusters that will emerge out of these clouds. With the exception of Taurus in which there is nearly no mass segregation, we observe mild-to-significant levels of mass segregation for the ensemble of the 6, 10, and 14 most massive cores in Aquila, Corona Australis, and W43, respectively. Our results suggest that the clouds' star formation activity are linked to their structure, as traced by their population of dense cores. We also find that the fraction of massive cores that are the most mass segregated in each region correlates with the surface density of star formation in the clouds. The Taurus region with low star-forming activity is associated with a highly hierarchical spatial distribution of the cores (low \Q\ value) and the cores show no sign of being mass segregated. On the other extreme, the mini-starburst region W43-MM1 has a higher \Q\ that is suggestive of a more centrally condensed structure and it possesses a higher fraction of massive cores that are segregated by mass. While some limited evolutionary effects might be present, we largely attribute the correlation between the star formation activity of the clouds and their structure to a dependence on the physical conditions that have been imprinted on them by the large scale environment at the time they started to assemble.} | The star formation process yields a large number of physical quantities and distribution functions that can be quantified by observations. Comparing these quantities and distributions between different star forming regions and/or with theoretical models and numerical simulations helps us gain insight into the relevance of various physical processes in different galactic environments. Some of the most studied quantities are the mass distribution of dense cores and of the stellar initial mass function in the early and emerging phases of the formation of stellar clusters (e.g., Johnstone \& Bally 2006; Dib 2014; Hony et al. 2015; Dib et al. 2017). The spatial distribution of dense cores in the early phases of star formation, of protostars in the phases of a cluster's buildup, and of stars in the (pseudo)gas-free phase of a young cluster, in conjunction with their masses and dynamics can also encapsulate critical information on how clusters assemble and form in different environments. The spatial distribution of stars in young and in evolved clusters has received a substantial amount of attention both in observations (McNamara \& Sekiguchi 1986; S\'{a}nchez \& Alfaro 2009; Gouliermis et al. 2014; Parker \& Alves de Oliveira 2017; Dib et al. 2018) and in numerical simulations of star forming regions (Schmeja \& Klessen 2006; Lomax et al. 2011; Parker \& Dale 2015; Gavagnin et al. 2017). However, quantifying the structure of star forming regions in the very early phases of star formation has remained elusive. This was principally due to the scarcity of observational data with the adequate spatial resolution to probe core masses in the stellar mass regime. With the advent of the {\it Hershel space observatory} (hereafter Herschel) and the {\it Atacama Large Millimeter Array} of radiotelescopes (ALMA), it is now possible to probe the spatial and mass distributions of dense structures in nearby star forming regions down to the mass regime of proto-brown dwarfs. The high sensitivity and spatial resolution of both Herschel and ALMA has provided so far unprecedented insight into quantities such as the dense core mass function (CMF) in the regions of Aquila, Taurus, Corona Australis, and W43 (Andr\'{e} et al. 2010; K\"{o}nyves et al. 2010,2015; Marsh et al. 2016; Bresnahan et al. 2018; Motte et al. 2018). The published CMFs of these regions suggest, at the very least, a striking difference between the low mass star forming regions such as Aquila and Taurus, and regions of massive star formation such as W43. When described by a power law of the form $dN/dM \propto M^{-\alpha}$, the derived value of $\alpha$ for W43 is $\approx 1.9$ which makes the CMF of W43 ostensibly shallower than in low mass star forming regions for which the derived values of $\alpha$ close to $2.3$ (Motte et al. 1998, but see Sadavoy et al. 2010), or even steeper as in the case of the California Molecular Cloud (Zhang et al. 2018). Dib et al. (2008a) showed that the slope of the CMF is steeper for cores that are defined using molecular species that trace higher densities of the gas and that are associated with an increasing degree of gravitational boundedness of the cores. However, in the case of the aforementioned star forming regions, the differences in the slopes of the CMF cannot be attributed to differences in the choice of the density tracers as all of the cores in those regions are observed in the continuum submillimeter emission. Within individual star forming regions, it has now been established that there are environmental differences between the populations of cores that are found on and off the filamentary structure of the clouds (Polychroni et al. 2013; Olmi et al. 2016; Kainulainen et al. 2017; Bresnahan et al. 2018). This can be attributed to effects of early mass segregation and/or to an extended phase of gas accretion by cores that reside inside the filaments (Dib et al. 2010) . \begin{figure*} \centering \includegraphics[width=0.46\textwidth]{f1.eps} \includegraphics[width=0.46\textwidth]{f2.eps}\\ \includegraphics[width=0.46\textwidth]{f3.eps} \includegraphics[width=0.46\textwidth]{f4.eps} \caption{Spatial distribution of dense cores in the four star forming regions considered in this work. The sizes of cores have been scaled (with an arbitrary formula) according to their masses (i.e., larger sizes relate to more massive cores) in order to visually highlight the location of the most massive cores.The minimum spanning tree (MST) for the ensemble of the 6 and 50 most massive cores in each region are displayed with the yellow and purple lines, respectively.} \label{fig1} \end{figure*} \begin{figure} \centering \includegraphics[width=0.9\columnwidth]{f5.eps} \vspace{0.5cm} \caption{The mass segregation ratios $\Lambda_{\rm MSR}$ (Allison et al. 2009, top row) and $\Gamma_{\rm MSR}$ (Olczak et al. 2011, bottom row) as a function of the number of most massive cores used in computing them, $n_{\rm MST}$. The values of $\Lambda_{\rm MSR}$ and $\Gamma_{\rm MSR}$ are calculated for the entire population of cores in the star forming regions (left column) and for the populations of bound cores (prestellar cores and cores with a protostar). A version of this figure that includes estimates of the uncertainties on $\Lambda_{\rm MSR}$ and $\Gamma_{\rm MSR}$ for each individual region is shown in Fig.~\ref{figappb1}.} \label{fig2} \end{figure} We complement the picture provided by the CMF of these different regions by analyzing the spatial distribution of their populations of dense cores. The structure and mass segregation of the cores in these regions lock important information on how star clusters form and evolve in different galactic environments. A number of studies have suggested a similarity between the structure of young clusters and of molecular clouds (Elmegreen et al. 2006; Gouliermis et al. 2014), while other works have argued that the similarity between the gas distribution and the distribution of stars can be quickly altered by the effects of rapid gas expulsion (Dib 2011; Dib et al. 2011;2013) or by gas tidal shocking of the young star clusters by surrounding gas clouds (Kruijssen et al. 2012). The question of whether the most massive stars in young clusters are mass segregated with respect to the total population of stars is still highly debated and has received a significant amount of attention both in observational studies (Hillenbrand \& Hartmann 1998; de Grijs et al. 2002; Gouliermis et al. 2004; Kerber \& Santiago 2006; Chen et al. 2007; Bontemps et al. 2010a; Er et al. 2013; Habibi et al. 2013; Lim et al. 2013; Elmegreen et al. 2014; Yu et al. 2017; Kuhn et al. 2017; Dib et al. 2018; Moser et al. 2019) and in theoretical/numerical works (Bonnell et al. 2003; Dib et al. 2007a; Dib 2007; Dib et al. 2008b; K\"{u}pper et al. 2011; Olczak et al. 2011; Maschberger \& Clarke 2011; Geller et al. 2013; Sills et al. 2018). The debate expands on whether the observed levels of mass segregation are primordial or due to dynamical interaction between stars in the clusters (Khalisi et al. 2007; Dib et al. 2018). In this work, we measure the spatial distribution of dense cores in the regions of Aquila, Taurus, Corona Australis, and W43 using the \Q\ parameter (Cartwright \& Whitworth 2004). We also quantify the levels of mass segregation in those regions using a measurement of the mass segregation ratio following the methods of Allison et al. (2009) and Olczak et al. (2011). We explore how the structure of these star forming regions and their levels of mass segregation correlate with their star formation activity. The paper is organized as follows: in \S.~\ref{datasets}, we briefly present the data sets that are used in this study. In \S.~\ref{methods}, we recall the basics of the methods used to quantify the structure parameter and mass segregation ratios. In \S.~\ref{results}, we present our results and in \S.~\ref{discussion}, we discuss them in light of previous work. In \S.~\ref{conclusions}, we conclude. | \label{discussion} A few studies on individual star forming regions by other group corroborate our findings. Alfaro \& Rom\'{a}n-Z\'{u}\~{n}iega (2018) analyzed the structure and mass segregation of dense cores in the quiescent Pipe nebula cloud which is characterized by a very low star formation activity (Lada et al. 2008). While in our case we chose to present an analysis of \Q\ and $\Lambda_{\rm MSR}$ (and $\Gamma_{\rm MSR}$) for the total and bound population of dense cores, Alfaro \& Rom\'{a}n-Z\'{u}\~{n}iega (2018) took a different approach and measured these parameters for cores selected in different bins of mass and gas density. For the Pipe nebula cloud, they found \Q\ values that are smaller for dense cores/peaks selected at either lower masses or lower gas volume densities of the order of $\approx 0.4$, similar to what we have measured for Taurus. They also found that the $\Lambda_{\rm MSR}$ values are relatively insensitive to the choice of the range of these physical parameters. Using data from the James Clerk Maxwell Telescope (JCMT) Gould Belt Survey, Kirk et al. (2016) applied both an MST based method (on the fluxes of cores rather than their masses) and the surface density-core mass method (see App.~\ref{appb}) and found evidence of primordial flux/mass segregation in the three subregions of Orion B (L1622, NGC 2023/2024, and NGC 2068/2071). Lane et al (2016) performed a very similar analysis to Kirk et al. for the entire Orion A cloud and reached similar conclusion. It is noteworthy that Kirk et al. (2016) reported very high values of \Q\ for the three sub-regions in Orion B, namely \Q\ $=1.18, 0.99$ and $0.91$ for L1622, NGC 2023/2024, and NGC 2068/2071, respectively. This is puzzling, given that in all three sub-regions, there is clear visual evidence of substructure. Parker (2018) re-analyzed the data of Kirk et al. (2016). He reported values of \Q\ that are $< 0.8$ ($\Q=0.72$, $0.65$, and $0.71$ for the sub-regions L1622, NGC 2068/2071, and NGC 20223/2024, respectively), and also variations in the values of $\Lambda_{\rm MSR}$ between those sub-regions with one sub-region (NGC 2023/2024) displaying strong levels of mass segregation ($\Lambda_{\rm MSR} \approx 28,$ for the $4$ most massive cores), another one (NGC 2068/2071) mild levels of mass segregation ($\Lambda_{\rm MSR} \approx 2$) and the third region, L1622, showing no mass-segregation. Parker 2018 argued (see Appendix A in his paper) that the large \Q\ values obtained by Kirk et al. (2016) are due to choices these authors made in how they normalize the quantities $\bar{\ell}_{\rm MST}$ and $\bar{s}$ (see discussion above in \S.~\ref{qparam}). We do not attempt to place the results of Alfaro \& Rom\'{a}n-Z\'{u}\~{n}iega (2018) and Parker (2018) on Fig.~\ref{fig3}. This is primarily due to the fact that the core selection for the Pipe and Orion B clouds were performed using different core extraction methods/algorithms (i.e., the CLUMPFIND and FellWalker algorithms, respectively) and a quantitative comparison using inhomogeneous data sets can lead to misleading conclusions. Nonetheless, the results of these two studies are broadly consistent with our findings in that the low star forming Pipe molecular cloud has a small \Q\ parameter, similar to the one we derived for Taurus, while the more intensely star forming regions in the Orion B cloud have a higher \Q\ values. Stutz \& Gould (2016) showed that dense cores in Orion A are well connected, both spatially and dynamically, to the integral shaped filament and the underlying gas structure which sets the gravitational potential of the region. Their approach is also another way of looking at the issue of whether there is a connection between the morphology of a star forming region, as traced by its population of dense cores, and the star formation activity in the region. The existence of substructure in star forming regions is a direct consequence of turbulent fragmentation. When turbulence is injected into the cloud on large physical scales (i.e., scales equal or larger than than cloud scale), a natural consequence of the turbulent cascade is the formation of a network of compressed, post-shock regions of different sizes (Dib et al. 2007b, Federrath et al. 2010; Burkhart et al. 2012; Padoan et al. 2014). Furthermore, as these substructures continue to fragment to smaller scales, forming dense cores, a non-zero level of mass-segregation can be expected, as more massive cores are statistically more likely to form in more massive substructures due to the availability of a larger mass reservoir (Padoan \& Nordlund 2002). In supersonic clouds that are magnetically subcritical such as Taurus \footnote{Hildebrand et al. 2009 and Chapman et al. 2011 report values of the mass-to magnetic flux ratio of $\mu=0.1\pm0.02$ and $0.32\pm0.02$, respectively, where $\mu$ is expressed in units of the critical value for collapse (Nakano \& Nakamura 1978). In W43, $\mu$ is highly supercritical and is found to be much in excess of unity across the entire region (Cortes et al. 2016); i.e., in the range $\approx 10-60$), and in Aquila, the available data about the magnetic field strength (Sofue \& Nakanishi 2017) and mean column density (K\"{o}nyves et al. (2015) suggest that $\mu$ falls in the range $5-7$ using the formula $\mu=7.6 \times 10^{-21} (N(H_{2})/B_{los}$ (Troland \& Crutcher 2008). There is no available information about the magnetic field strength in Corona Australis.}, star formation, which is mediated by ambipolar diffusion, proceeds at a slower pace due to effects of magnetic pressure which prevents substructure from merging efficiently (Nakamura \& Li 2005; Dib et al. 2007b). Thus, a higher level of substructure in the clouds could be indicative of either a young age for the region and/or of the existence of a strong magnetic support. With regard to the question of whether mass segregation is primordial (i.e., set at the dense core phase) or induced by the dynamical evolution of an initially sub-structured cluster, Dom\'{i}nguez et al. (2017) argued that a young cluster will quickly settle into a state where the most massive stars are mass segregated, regardless of the existence/absence of mass segregation within its different levels of substructure. On the other hand, Dib et al. (2010) argued that mild levels of mass segregation that can be generated by turbulent fragmentation could be significantly enhanced by gas accretion onto the cores. Dib et al. (2007a) proposed that another efficient channel for imprinting a significant level of primordial mass segregation is by the coalescence of cores. The process is more efficient in the densest regions of protocluster clouds where cores are more closely packed, and it may be the dominant mode of star formation in the centre of massive starburst clusters such as Arches and NGC 3603. | 18 | 8 | 1808.07057 |
1808 | 1808.00488_arXiv.txt | We explore the structure of galaxy cluster Abell 2029 and its surroundings based on intensive spectroscopy along with X-ray and weak lensing observations. The redshift survey includes 4376 galaxies (1215 spectroscopic cluster members) within 40\arcmin of the cluster center; the redshifts are included here. Two subsystems, A2033 and a Southern Infalling Group (SIG) appear in the infall region based on the spectroscopy as well as on the weak lensing and X-ray maps. The complete redshift survey of A2029 also identifies at least 12 foreground and background systems (10 are extended X-ray sources) in the A2029 field; we include a census of their properties. The X-ray luminosities ($L_{X}$) -- velocity dispersions ($\sigma_{cl}$) scaling relations for A2029, A2033, SIG, and the foreground/background systems are consistent with the known cluster scaling relations. The combined spectroscopy, weak lensing, and X-ray observations provide a robust measure of the masses of A2029, A2033, and SIG. The total mass of the infalling groups (A2033 and SIG) is $\sim 60\%$ of the M$_{200}$ of the primary cluster, A2029. Simple dynamical considerations suggest that A2029 will accrete these subsystems in next few Gyr. In agreement with simulations and with other clusters observed in a similar redshift range, the total mass in the A2029 infall region is comparable with the A2029 M$_{200}$ and will mostly be accreted in the long-term future. | Galaxy clusters grow hierarchically through accretion of generally lower mass systems. More massive clusters typically form later than lower massive systems \citep{Neto07, BoylanKolchin09, McBride09}. Numerical simulations suggest that the mass accretion rate is roughly proportional to the cluster mass: more massive clusters accrete more mass \citep{vandenBosch02, Fakhouri08, Fakhouri10, Giocoli12, deBoni16}. Observational estimates of the mass within the infall region of galaxy clusters enable measurement of mass accretion rates \citep{Diaferio97, Rines02, deBoni16}. The measurement of the mass infall rate is challenging because detailed observations covering the clusters outskirts (or infall regions) are required. Wide field-of-view redshift surveys, X-ray observations, and weak lensing offer complementary views of the infall region of an individual cluster. So far, there are relatively few systems where all of these observations are available. Wide field-of-view redshift surveys \citep{Geller99, Reisenegger00, Rines02} apply the caustic technique \citep{Diaferio97, Diaferio99} to estimate the mass in the infall region. \citet{Rines13} show that the typical mass in the infall region is comparable with $M_{200}$ ($= (4\pi/3)R_{200}^{3} 200 \rho_{crit}$). Identification of X-ray emitting groups in the infall region provides another probe of the future accretion by the cluster. The X-COP project \citep{Eckert17} surveys galaxy clusters with very deep {\it XMM} images to study infalling groups. \citet{Haines18} also conduct a systematic survey of X-ray groups in the infall region of 23 clusters at $z \sim 0.2$. They estimate that the galaxy clusters typically accrete $32\%$ of their mass by redshift 0 through the accretion of these surrounding X-ray groups. Weak gravitational lensing is another method for estimating the amount of mass in the outskirts of clusters (e.g. \citealp{Geller13, Umetsu17}). Unlike the X-ray mass estimates, weak lensing mass estimates are independent of the cluster dynamical state. The estimated mass in the infall region derived from weak lensing observations is consistent with caustic estimates from dense redshift surveys \citep{Geller13}. Combining these complementary probes strongly constrains the mass within the cluster and its infall regions. Each method of measuring the potentially infalling mass has limitations. For example, mass estimates based on the X-ray depend on the assumption of hydrostatic equilibrium. Dense spectroscopy is critical. Without redshifts, association between extended X-ray emission and the main cluster is ambiguous. Weak lensing mass estimates may be contaminated by the presence of foreground/background systems (e.g. \citealp{Hoekstra11, Geller13, Hwang14}). The redshift survey facilitates the separation of cluster members from these foreground/background structures. Here, we combine spectroscopy, X-ray, and weak lensing observations to study the future mass accretion by the nearby massive cluster A2029. A2029 is one of the most massive clusters at $z = 0.079$. A2029 is well studied cluster with {\it ROSAT}, {\it XMM}, {\it Suzaku} and {\it Chandra} observations (e.g. \citealp{Lewis02, Clarke04, Walker12, PaternoMahler13}). \citet{McCleary18} construct weak lensing map of A2029. \citet{Sohn17a} conduct a redshift survey of this cluster (see also \citealp{Tyler13}). They examine the statistical properties of the A2029 member galaxies including luminosity, stellar mass, and velocity dispersion functions. The complete redshift survey we discuss extends the survey of \citet{Sohn17a}. Based on complete spectroscopy, we investigate the core of A2029 and its infall region. We identify two relatively massive subsystems in the infall region and investigate their physical properties based on spectroscopy, X-ray and weak lensing maps. In this process, we refine the X-ray estimates of the subsystem masses. We probe the future dynamical evolution of the A2029 system based on the physical properties of the infalling groups. We also use the complete redshift survey to make a census of foreground/background systems. Construction of this census is critical to removing ambiguous contributions to the mass within the infall region. Including A2029 and the two infalling groups, we find a total of 13 extended X-ray sources. Their physical properties are consistent with the well-known scaling relation between X-ray luminosity and velocity dispersion. The combined analysis we discuss sets the stage for future large datasets including these complementary probes of the mass distribution in and around clusters of galaxies. eROSITA \citep{Merloni12}, Prime-Focus Spectrograph (PFS) on {\it Subaru} \citep{Takada14}, and Euclid \citep{Amendola18} will provide these observations for clusters with a wide range of masses and redshifts thus tracing the detailed evolution of these systems. We describe the redshift survey of A2029 in Section \ref{data}. We explain the identification of the cluster members using spectroscopic data in Section \ref{memsel}. In Section \ref{a2029struct}, we identify two groups within the infall region of A2029 along with foreground/background systems in the A2029 field. We summarize their aggregate properties by placing them on the well-known $L_{X} - \sigma_{cl}$ scaling relation. Finally, we discuss the past and future accretion history of A2029 (Section \ref{discussion}) based on the dynamical connection between A2029 and the massive infalling groups. Throughout this paper, we use the standard $\Lambda$CDM cosmology parameters: $H_{0} = 70 ~\kmsmpc$, $\Omega_{m} = 0.3$ and $\Omega_{\Lambda} = 0.7$. | We combine a dense redshift survey of the local massive cluster A2029 with X-ray and weak lensing maps to elucidate the future accretion story of this massive system. The total dataset for A2029 is unusually rich. The redshift survey is essentially complete within a wide field of $R_{cl} < 40\arcmin (= 3.5$ Mpc) around A2029. We refine analysis of the {\it ROSAT} images and of the weak lensing map to improve mass estimated for two massive subsystems, A2033 and SIG, within the A2029 infall region. The infalling groups, A2033 and SIG, appear in the weak lensing map and the X-ray image and the spectroscopic survey. Interestingly, the brightest galaxies in these subgroups are offset from the group centers (determined by X-ray or cluster members). The astrophysical implications of these offsets are unclear. The complete redshift survey facilitates the identification of foreground and background groups in the A2029 field. This identification is critical for removing spurious contributions to the mass within the infall region. We identify at least 12 foreground/background systems. Among these systems, 10 systems have {\it ROSAT} X-ray counterparts; a very bright X-ray group LOS7 lies $z = 0.223$. Oddly its position makes it appear to be a filamentary connection between A2033 and A2029. The redshift survey makes it clear that this apparent connection is merely a superposition. Taking these extended X-ray sources together with A2029, A2033, and SIG we demonstrate that they are all consistent with the well-known scaling relation between X-ray luminosity and velocity dispersion. We measure the mass of A2029 based on the three different mass proxies: caustics, weak lensing and X-ray luminosity (or temperature). The caustic mass based on the spectroscopic members is $M_{200} = (8.47 \pm 0.25) \times 10^{14} M_{\odot}$, agrees to within $1\sigma$ with the X-ray estimate. We also estimate the masses of infalling groups using velocity dispersions, weak lensing and X-ray luminosities. Within the much larger uncertainties, the estimates agree. They imply that the total mass in these two subsystems in $\sim 60\%$ of the mass of the main cluster. A simple two-body model traces the future accretion of the infalling groups. The model suggests that the infalling groups are obviously bound to A2029 and may be accreted by the primary cluster within $\sim 3$ Gyr. This accretion rate is larger than the average predicted by simulations. The infall region as a whole contains an amount of mass comparable with the A2029 M$_{200}$. The two massive subsystems contribute about $\sim 60\%$ of the mass in the infall region. Numerical simulations suggest that $90\%$ of the mass in the infall region will be accreted in the long-term future of the cluster. In the future a combination of eROSITA, PFS, and Euclid observations will make similar analyses possible for clusters across a broad range of cluster mass and over a wide redshift range. These combined spectroscopic, X-ray and weak lensing observations will enable construction of the full picture of the accretion story of clusters of galaxies. They will provide a strong test of the hierarchical structure formation picture. | 18 | 8 | 1808.00488 |
1808 | 1808.01501_arXiv.txt | The aim of this work is to contribute to the understanding of the stellar velocity distribution in the solar neighborhood (SN). We propose that the structures on the $U$--$V$ planes, known as the moving groups, can be mainly explained by the spiral arms perturbations. The applied model of the Galactic disk and spiral arms, with the parameters defined by observational data and with pattern speed $\Omega_p=$28.0\,km\,s$^{-1}$\,kpc$^{-1}$, is the same that allowed us to explain the origin of the Local Arm and the Sun's orbit trapped inside the corotation resonance (CR). We show that the $U$--$V$ picture of the SN consists of the main component, associated with the CR, and the inner and outer structures, which we could associate with the Hercules and Sirius streams, respectively. The Coma-Berenices and Hyades-Pleiades groups and the Sun itself belong to the main part. The substructures of Hercules are formed mainly by the nearby 8/1, 12/1, and even 6/1 inner Lindblad resonances, while Sirius is shaped by the bulk of overlapping outer Lindblad resonances, -8/1, -12/1, -16/1, which are stuck to the CR. This richness in resonances only exists near corotation, which should be of the spiral arms, not of the Galactic bar, whose stable corotation zone is far away from the Sun. The model's predictions of the velocity distribution match qualitatively and quantitatively the distribution provided by Gaia DR2. | \label{intro} In recent years, great efforts have been dedicated to give plausible explanations to the moving groups and the rough bimodality which are observed in the velocity distribution of the solar neighborhood (SN) \citep[][]{dehnen2000AJ, quillenMinchev2005AJ, antojaEtal2014AA}. However, these kinematic structures are not yet settled and it seems that, at present, we lack models which could explain the observations as complete as those presented by Gaia DR2 \citep{Gaia2018A}. The stellar velocity distribution near the Sun shows density clumps, which were first related to the disruption of open clusters (see \citealp{antojaEtal2010LNEA}, for a review). Nevertheless, the heterogeneity of the ages of these groups \citep{bensbyEtal2007ApJL,famaeyEtal2008AA} was in contradiction with this hypothesis, demanding a new scenario. Furthermore, a question, which has drawn attention, was the apparent bimodality in the $U$--$V$ distribution, in which one of these groups, Hercules, is separated from the main component of the velocity distribution in the SN. Some dynamical scenarios explain this separation considering it as an imprint of the Galactic bar perturbations \citep[][]{dehnen1999ApJL, dehnen2000AJ, bovy2010ApJ, antojaEtal2014AA, perezvillegasEtal2017ApJL}. Others study spiral arms perturbations \citep{desimoneEtal2004MNRAS,quillenMinchev2005AJ,antojaEtal2011MNRAS,quillenEtal2018arXiv} or a mix of bar+spirals \citep{antojaEtal2009ApJL}. Additionally, transient spiral structures and phase wrapping were recently proposed to explain $U$--$V$ distribution in the SN \citep{Hunt2018arXiv180602832H,antojaEtal2018arXiv}. Our explanation presented here is based on strong observational evidence, which include the Sagittarius-Carina and Perseus spiral arms, the Local Arm \citep[][hereafter Paper-I]{LepineEtal2017ApJ} and the proximity between the spiral corotation and the solar circle \citep{mishurovZenina1999AA}. We find that the complex structure of the $U$--$V$ plane in the SN can be represented roughly by three components: the main component, associated with the spiral corotation resonance (CR), Hercules, associated with inner Lindblad resonances (ILRs), and Sirius, associated with the bulk of overlapping outer Lindblad resonances (OLRs). Both ILRs and OLRs are high-order resonances, which naturally appear in the corotation neighborhood. \begin{figure*}[h] \begin{center} \epsfig{figure=fig1.pdf,width=1.0\textwidth,angle=0} \caption{Left: Dynamical map on the $U$--$V$ plane constructed at the initial position of the Sun, $R=8.0$\,kpc and $\varphi=90^\circ$ (blue cross). The main resonances are identified by the corresponding ratios. The dashed lines separate roughly the main component from the Hercules (V$<$-30\,km/s) and Sirius (V$>$5\,km/s) regions. Right: Heliocentric velocity distribution of the Gaia stars within 150\,pc from the Sun as a 2D histogram of the velocity with a bin-size of 1\,km\,s$^{-1}$ in both directions. The color scale indicates the percentage of stars per (km/s)$^2$. } \label{fig:fig1} \end{center} \end{figure*} | \label{discussion} We presented a novel approach to explain the distribution of stars on the $U$--$V$ plane of the SN, based on a spiral arms dynamical model. Our basic assumptions are: i) the proximity between the spiral corotation radius and the solar circle, and ii) the dynamical stability of the Local Arm and the Sun. The suggestion that both the Local Arm and the Sun evolve inside the spiral CR yields natural constraints on the magnitude of the Galaxy's spiral pattern speed. For the adopted rotation curve and solar position (Paper-I), we choose $\Omega_p = 28$\,km\,s$^{-1}$\,kpc$^{-1}$, which results in a corotation radius of $8.2$\,kpc. The comparison of the stellar velocity distribution provided by Gaia DR2 with the dynamical maps of the $U$--$V$ plane shows that it is related to the CR and nearby Lindblad resonances. It is important to state that, in general, resonances modify qualitatively the dynamics in their environment: they capture and trap stars inside the stable resonant zones, enhancing the density, and deplete regions close to saddle points and separatrices. Corotation is special among the resonances, being stronger and having a wider zone of influence surrounded by many high-order resonances. That is exactly what we observe on the $U$--$V$--plane (right) in Fig.\,\ref{fig:fig1}: the central region of enhanced density, known as main component, is the CR stable zone according to the dynamical map of the same plane (left). The notable moving groups, Coma-Berenices and Hyades-Pleiades, are inside the CR, together with the Sun, located at the origin of the plane. Immediately above the main component, for V$>$0, we can observe the Sirius group, whose origin is related to the overlapping high-order OLRs stuck to the CR. % Finally, the strong 4/1\,OLR is responsible for distant stars to visit the SN. In the region below the main component, at V$<$0, the nearby resonances are separated sufficiently on the $U$--$V$ plane to create several substructures well defined by Gaia DR2. The most prominent is the Hercules group related mainly to the 8/1 and 12/1\,ILRs. The overdensities produced by the inner resonances are explained by the fact that the objects, coming from the inner Galaxy, spend longer times at the outer edges of their orbits in the SN, since their velocities are smaller there. This kinematics produces the effect of a bimodal distribution on the $U$--$V$ plane. Despite the simplicity of our model, we are able to explain quantitatively the SN $U$--$V$ structure and its changes with different Galactic radii. The results presented here can be considered as an additional test, which confirms the robustness of our model. They are also another strong indicator that the spiral corotation radius lies near the solar circle, since only in the neighborhood of the CR this abundance of Lindblad resonances can be observed. The bar's CR has a minor role, since, due to its spacial orientation, the $L_4$--points of the bar are distant from the Sun, as shown in \cite{Michtchenkoetal2018AA}. | 18 | 8 | 1808.01501 |
1808 | 1808.04112_arXiv.txt | { We re--examine the case of anapole dark matter as an explanation for the DAMA annual modulation in light of the DAMA/LIBRA--phase2 results and improved upper limits from other DM searches. If the WIMP velocity distribution is assumed to be a Maxwellian, anapole dark matter is unable to provide an explanation of the DAMA modulation compatible with the other searches. Nevertheless, anapole dark matter provides a better fit to the DAMA--phase2 modulation data than an isoscalar spin--independent interaction, due to its magnetic coupling with sodium targets. A halo-independent analysis shows that explaining the DAMA modulation above 2 keVee in terms of anapole dark matter is basically impossible in face of the other null results, while the DAMA/LIBRA--phase2 modulation measurements below 2 keVee are marginally allowed. We conclude that in light of current measurements, anapole dark matter does not seem to be a viable explanation for the totality of the DAMA modulation.} \begin{document} | \label{sec:introduction} Weakly Interacting Massive Particles (WIMPs) provide one of the most popular explanations for the Dark Matter (DM) that is believed to make up 27\% of the total mass density of the Universe~\cite{planck} and more than 90\% of the halo of our Galaxy. The scattering rate of DM WIMPs in a terrestrial detector is expected to present a modulation with a period of one year due to the Earth revolution around the Sun~\cite{Drukier:1986tm}. For more than 15 years, the DAMA collaboration~\cite{dama_2008,dama_2010,dama_2013} has been measuring a yearly modulation effect in their sodium iodide target. The DAMA annual modulation is consistent with what is expected from DM WIMPs, and has a statistical significance of more than $9\sigma$. However, in the most popular WIMP scenarios used to explain the DAMA signal as due to DM WIMPs, the DAMA modulation appears incompatible with the results from many other DM experiments that have failed to observe any signal so far. This has prompted the need to extend the class of WIMP models. In particular, one of the few phenomenological scenarios that have been shown \cite{anapole_2014} to explain the DAMA effect in agreement with the constraints from other experiments is Anapole Dark Matter (ADM) \cite{anapole1,anapole2,anapole3,anapole4}, for WIMP masses $m_{\chi}\lsim$ 10 \GeV/c$^2$. Recently the DAMA collaboration has released first results from the upgraded DAMA/ LIBRA-phase2 experiment \cite{dama_2018}, increasing the significance of the effect to 12 $\sigma$. The two most important improvements compared to the previous phases are that now the exposure has almost doubled and the energy threshold has been lowered from 2 keV electron--equivalent (keVee) to 1 keVee. While for $m_{\chi}\lsim 10$~\GeV/c$^2$ the DAMA phase--1 data where only sensitive to WIMP--sodium scattering events, the new data below 2 keVee are in principle also sensitive to WIMP-iodine scattering, for WIMP speeds below the escape velocity in our Galaxy. This feature has worsened the goodness of fit of the DAMA data using a standard Spin-Independent interaction (SI) \cite{freese_2018,dama_2018_sogang}. In light of the DAMA/LIBRA--phase2 result, in the present paper we re--examine the ADM scenario. Moreover, compared to the analyses in \cite{anapole_2014}, we upgrade the constraints from other direct detection experiments. In this analysis we use results from CDEX~\cite{cdex}, CDMSlite~\cite{cdmslite_2017}, COUPP~\cite{coupp}, CRESST-II~\cite{cresst_II,cresst_II_ancillary}, DAMIC~\cite{damic}, DarkSide--50 ~\cite{ds50}, KIMS~\cite{kims_2014}, PANDAX-II~\cite{panda_2017}, PICASSO~\cite{picasso}, PICO-60~\cite{pico60_2015,pico60}, SuperCDMS~\cite{super_cdms_2017} and XENON1T~\cite{xenon_1t,xenon_2018}. The paper is organized as follows. In Section \ref{sec:model} we summarize the main features of the ADM scenario, providing the formulas for WIMP direct detection expected rates; our main results are in Section \ref{sec:analysis}, where we provide an updated assessment of ADM in light of the DAMA--phase2 data and of the latest constraints from other direct detection experiments, both assuming a Maxwellian WIMP velocity distribution and in a halo--independent approach. Section \ref{sec:conclusions} is devoted to our conclusions. In Appendix \ref{app:exp} we provide some details on how the experimental constraints on ADM have been obtained. | \label{sec:conclusions} We have re--examined the case of anapole dark matter as an explanation for the DAMA annual modulation in light of the DAMA/LIBRA--phase2 results and improved upper limits from other DM searches. For a Maxwellian WIMP velocity distribution, anapole dark matter is unable to provide an explanation of the DAMA modulation compatible with the other direct DM search results. Nevertheless, anapole dark matter provides a better fit to the DAMA--phase2 modulation data than a a standard isoscalar spin--independent interaction. This is due to the contribution from the magnetic moment of sodium, which reduces the hierarchy between the ADM WIMP--iodine and WIMP--sodium cross sections compared to the SI case. A halo-independent analysis shows that explaining the DAMA modulation above 2 keVee in terms of anapole dark matter is basically impossible in the face of the null results of XENON1T, PANDAX-II, and PICO(C${}_3$F$_{8}$). On the other hand, the DAMA/LIBRA--phase2 modulation measurements below 2 keVee lie near the border of the excluded region. We conclude that in light of current measurements, anapole dark matter does not seem to be a viable explanation for the totality of the DAMA modulation, not even in a halo--independent analysis, although the DAMA/LIBRA--phase2 modulation measurements below 2 keVee are marginally allowed. | 18 | 8 | 1808.04112 |
1808 | 1808.10475.txt | \label{sec-intro} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% In standard single field inflation, the linearised dynamics of the scalar curvature fluctuation ${\cal R} $ is described by a quadratic action \cite{Mukhanov:1990me} \be\label{scal-eq} S_{R}\,=\,\int d \eta\,d^3 x\,\frac{z_S^2}{2}\,\left[ {\cal R'}^2-(\vec \nabla {\cal R} )^2\right] \ee with: \be \label{scal-eq-ws} z_S\,\equiv\,a\,\frac{\dot \phi}{H}, \ee where $\phi$ is the homogeneous scalar field profile, and prime and dot indicate respectively derivatives along conformal and physical time. The combination $z_S$ %we identified earlier is known as the scalar pump field. If $z_S$ is an increasing function of time -- as in a slow-roll regime, where the Hubble parameter and $\dot \phi$ are approximately constant -- then inflation is in an {\em attractor phase}, ${\cal R}$ is conserved at superhorizon scales and its spectrum is almost scale invariant. However, if $z_S$ is rapidly decreasing for a brief interval, then the inflationary evolution is no longer an attractor: the would-be decaying mode of the curvature perturbation becomes dominant, and the power spectrum of modes leaving the horizon during this {\em non-attractor phase} can be enhanced by several orders of magnitude in a short time interval. This can occur for example in models where the scalar derivative rapidly decreases for a short period, as in inflection point ultra slow-roll and in constant roll inflationary systems (see e.g. \cite{Inoue:2001zt,Linde:2001ae,Kinney:2005vj,Martin:2012pe,Motohashi:2014ppa,Yi:2017mxs,Dimopoulos:2017ged,Pattison:2018bct}) % (during which $\dot \phi\sim 1/a^3$) or in the Starobinsky model with a rapid change in the potential slope \cite{Starobinsky:1992ts}. Interestingly, although the non-attractor phase of inflation lies well outside the slow-roll regime, a duality exists \cite{Wands:1998yp} which allows an analytical description of the statistical features of the enhanced spectrum of fluctuations. Recently these scenarios have received a renewed interest, since an amplification of scalar perturbations can lead to the production of primordial black holes from single field inflation (see e.g. \cite{Garcia-Bellido:2017fdg,Sasaki:2018dmp,Carr:2016drx} for general reviews and \cite{Garcia-Bellido:2017mdw,Germani:2017bcs,Motohashi:2017kbs,Ballesteros:2017fsr,Ezquiaga:2017fvi,Cicoli:2018asa,Ozsoy:2018flq,Biagetti:2018pjj} for a specific models). \smallskip Can we have a similar enhancement of primordial tensor modes during a phase of non-attractor single-field inflation? This question is phenomenologically interesting: while current and forthcoming constraints from CMB polarization can probe the amplitude of the primordial tensor spectrum at very large CMB scales (see e.g. the reviews \cite{Chongchitnan:2006pe,Kamionkowski:2015yta}), interferometers can probe a stochastic background of gravitational waves at much smaller scales (see the textbooks \cite{Maggiore:1900zz,Maggiore:2018sht}). Hence inflationary scenarios that enhance the spectrum of primordial tensor modes at interferometer scales make predictions that are easier to test with interferometers instead of CMB experiments. So far, two main approaches have been proposed. The first usually exploits instabilities for additional source fields during inflation: primordial gravity waves can be enhanced by coupling fields driving inflation with additional scalars \cite{Cook:2011hg,Senatore:2011sp,Carney:2012pk,Biagetti:2013kwa,Biagetti:2014asa,Goolsby-Cole:2017hod}, U(1) gauge vectors \cite{Sorbo:2011rz, Anber:2012du, Barnaby:2010vf , Barnaby:2012xt}, non-Abelian vector fields \cite{Maleknejad:2011jw,Dimastrogiovanni:2012ew,Adshead:2013qp,Adshead:2013nka,Obata:2014loa,Maleknejad:2016qjz,Dimastrogiovanni:2016fuu,Agrawal:2017awz,Adshead:2017hnc,Caldwell:2017chz,Agrawal:2018mrg}, or Standard Model fields \cite{Espinosa:2018eve}. The second approach implements space-time symmetry breaking during inflation. Ways to do so are scenarios of (super)solid inflation -- see e.g. \cite{Endlich:2012pz,Bartolo:2015qvr,Ricciardone:2016lym,Ricciardone:2017kre,Domenech:2017kno,Ballesteros:2016gwc,Cannone:2015rra,Lin:2015cqa,Cannone:2014uqa,Akhshik:2014bla} -- or massive gravity/bigravity models, \cite{Biagetti:2017viz,Dimastrogiovanni:2018uqy,Fujita:2018ehq}. See e.g. \cite{Bartolo:2016ami} for a more extensive survey of various models proposed so far, focussing on the detectability of inflationary tensor modes with LISA. The aim of this work is to present a new mechanism to enhance the spectrum of primordial tensor fluctuations in single field inflation, which can be used to enhance spin 2 fluctuations at arbitrary scales. It is based on the hypothesis that the single field inflationary dynamics encounters a brief non-attractor phase during its evolution: then the would-be decaying tensor mode grows at super-horizon scales, and enhance the tensor power spectrum. The advantage of working in single field inflation is that we do not have to deal with backreaction of additional fields that can interfere with the inflationary dynamics. We consider a set-up with non-minimal couplings between the inflaton field and the metric, in order to have an adjustable function of time in the quadratic action for tensor fluctuations, which we shall use %analogously to the scalar case %(see the previous discussion after eqs \eqref{scal-eq}, \eqref{scal-eq-ws}) to enhance the primordial tensor spectrum. We proceed as follows. \begin{itemize} \item In Section \ref{sec-enha} we study the second order action for primordial tensor fluctuations in single field inflation. We identify conditions for obtaining a large enhancement of the spectrum of primordial gravity waves, by exploiting a non-attractor phase for tensors that enhance the would-be decaying tensor mode. These conditions are analogous to the requirements discussed in various works, starting with \cite{Leach:2000yw,Leach:2001zf}, for enhancing scalar modes during non-attractor phases, and motivate our search for models of inflation with specific non-minimal couplings of tensors to the inflationary scalar field. \item In general, it is difficult to have analytic control of the dynamics of fluctuations in a non-attractor phase. In Section \ref{sec-dua} we identify a criterium, which we call {\it tensor duality}, that ensures identical behaviour, up to an overall factor, for the dynamics of perturbations in two different regimes of inflationary evolution. This is the generalization to the tensor case of the duality discussed by Wands \cite{Wands:1998yp} for the scalar sector. We determine the tensor dual of a phase of standard slow-roll inflation, which corresponds to a period of non-attractor inflation, with a scale invariant spectrum of tensor fluctuations amplified with respect to the standard case. \item Using tensor duality as a guide, in Section \ref{sec-modb} we build and analyse in detail a representative model of single field kinetically driven inflation, belonging to the G-inflation set-up of \cite{Kobayashi:2011nu}, which is able to amplify tensor modes. %during a non-attractor era. Our system is analogous to the Starobinsky model \cite{Starobinsky:1992ts}, where instead of having discontinuities in the potential, we have a discontinuity in the kinetic functions which causes a short non-attractor phase. Using tensor duality, we are then able to analytically investigate the dynamics of fluctuations during the non-attractor era, showing that the amplitude of the spectra of tensor (and scalar) fluctuations increases by several orders of magnitude with respect to a standard slow-roll regime. \item We conclude in Section \ref{sec-out} with a discussion of possible future directions to explore, and provide technical appendices for some of the results of the main text. \end{itemize} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 18 | 8 | 1808.10475 |
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1808 | 1808.04330_arXiv.txt | Following the detection of a $\sim$300 TeV neutrino potentially associated with the flaring blazar TXS~0506+056, an excess of neutrinos around its position in 2014-2015 was revealed by IceCube. However, its contemporaneous quiescence in $\gamma$-rays is challenging to interpret consistently. Meanwhile, the blazar PKS~0502+049, positioned within the neutrino localization uncertainties, was seen to be flaring in $\gamma$-rays. We show that dense, line-emitting gas clouds that interact with its jet and induce cosmic ray acceleration and hadronuclear interaction can plausibly explain the 2014-2015 neutrino flare. | 18 | 8 | 1808.04330 |
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1808 | 1808.03984_arXiv.txt | The high spatial resolution and precise astrometry and photometry of the {\it Gaia} mission should make it particularly apt at discovering and resolving transients occurring in, or near, the centres of galaxies. Indeed, some nuclear transients are reported by the {\it Gaia} Science Alerts (GSA) team, but not a single confirmed tidal disruption event (TDE) has been published. In order to explore the sensitivity of GSA, we performed an independent and systematic search for nuclear transients using {\it Gaia} observations. Our transient search is driven from an input galaxy catalogue (derived from the Sloan Digital Sky Survey Release 12). We present a candidate detection metric that is independent from the existing GSA methodology, to see if {\it Gaia} Alerts are biased against the discovery of nuclear transients, and in particular which steps may have an impact. Our technique does require significant manual vetting of candidates, making implementation in the GSA system impractical for daily operations, although it could be run weekly, which for month-to-year long transients would make a scientifically valuable addition. Our search yielded $\sim$480 nuclear transients, 5 of which were alerted and published by GSA. The list of (in some cases ongoing) transients includes candidates for events related to enhanced accretion on to a super-massive black hole and TDEs. An implementation of the detection methodology and criteria used in this paper as an extension of GSA could open up the possibility for {\it Gaia} to fulfil the role as a main tool to find transient nuclear activity as predicted in the literature. | The European Space Agency (ESA)--{\it Gaia} mission has been operational since mid-2014 and has provided accurate photometric, astrometric, and spectroscopic measurements for roughly a billion stars in the Milky Way (\citealt{2016A&A...595A...1G,2016A&A...595A...2G}). {\it Gaia}'s on-board detection algorithms are optimised for the detection of point-like sources, which implies that extended sources have to have an effective radius less than 0.6 arcsec in order to be detected (\citealt{2014sf2a.conf..421D,2015A&A...576A..74D}). Thus, the mission is also collecting data for a significant number of resolved extragalactic objects, such as small elliptical galaxies, and galaxies with compact bulges, or point-like sources such as high redshift quasars. The observing strategy (Nominal Scanning Law) of {\it Gaia} is optimised to deliver data for parallax measurement. As a result of this scanning law, most of the sources will be scanned more than about 70 times from different angles during five year mission. Each position on the sky is scanned, on average, once every 30 days (\citealt{2016A&A...595A...4L}). These repeat visits make {\it Gaia} an all-sky, multi-epoch photometric survey, that allows us to monitor variability with high precision as well as detect new transient sources (\citealt{2013RSPTA.37120239H,2017arXiv170203295E}). The Data Processing and Analysis Consortium (DPAC) {\it Gaia} data flow enables detections of transients within 24-48 hours of the observation (in the best case). However, due to different reasons the delay might be up to a few days (Hodgkin et al. in prep). From September 2014 onwards new transients from {\it Gaia} have been made publicly available after manual vetting of candidate transients detected by the {\it Gaia} Science Alerts (GSA) team. To this end, AlertPipe - dedicated software for data processing, transient searching, and candidate filtering was employed (Hodgkin et al.~in prep). Such a near-real-time survey is predicted to detect about 6000 low-redshift supernovae brighter than $G=19$ mag and 1300 microlensing events during the first five years of the mission (\citealt{2002MNRAS.331..649B,2003MNRAS.341..569B,2012Ap&SS.341..163A}). Accurate photometry and low--resolution spectroscopy should allow for a robust classification and reduce the rate of false positives. {\it Gaia} could therefore play an important role in transient detection. AlertPipe employs two different transient detection algorithms. Transient discovery is either based on the detection of a new source ({\it NewSource} detector; the event has to have 2 or more observations above flux threshold, equivalent to $G=19$), or a significant deviation in brightness of a known source compared to previous {\it Gaia} photometry ({\it OldSource} detector, either a source brightened by more than 1 magnitude and this rise is more than 3 sigma above the rms of the historic variations from all data available, or a source brightened by more than 0.15 magnitudes and this rise is more than 6 times the rms of the historic variations from all data available). The thresholds and other detection parameters of both detectors are tuneable but have been kept fixed over the period June 2016 -- June 2017 under consideration here. Optical variability occurring in the centres of galaxies can be associated with transients such as (superluminous) supernovae (SNe) and tidal disruption events (TDEs). On the other hand active galactic nuclei (AGNs) exhibit variability across the whole electromagnetic spectrum caused by activity in the accretion disk and jet (see e.g.~\citealt{2012ApJ...753..106M,2017MNRAS.470.4112G} for recent work showing transient and variability phenomena associated with the central super-massive black hole in galaxies). Core-Collapse Supernovae (CCSNe) originate from explosions of massive stars (masses $M>8 M_\odot$). CCSNe have been proposed as environmental metallicity probes (\citealt{2014MNRAS.440.1856D}). Together with Type Ia Supernovae (SNe Ia), they shape and influence galaxy structure and star formation (\citealt{2017ApJ...848...25M}). Being standard candles, SNe Ia can also be used as probes of distribution of dust in their host galaxies. Tracing it is particularly important in the very cores of galaxies, typically containing large amounts of obscuring dust, to test the relations between SN Ia observed brightness and distance from the core of their host galaxy as well as morphology of the host. Superluminous supernovae (SLSNe) are associated with deaths of the most massive stars, which means that they may have an impact on the chemical evolution and re-ionization of the Universe (\citealt{2010ApJ...724L..16P,2012Sci...337..927G}). The SLSN explosions are probably induced by different physical mechanisms than other, more common types of SNe (\citealt{2018MNRAS.475.1046I}). Tidal disruption events can be used to determine the presence and study the properties, such as the mass, of supermassive black holes (SMBHs) in quiescent galaxies (\citealt{1988Natur.333..523R}). TDE properties probe the stellar populations and dynamics in galactic nuclei, the physics of black hole accretion including the potential to detect relativistic effects near the SMBH, and the physics of jet formation and evolution (e.g. \citealt{2016MNRAS.461..371K,2018ApJ...852...72V}). In addition, because the rate of TDEs is temporarily massively enhanced in binary SMBH systems, TDEs might point us to galaxies that host compact binary SMBHs (\citealt{2011ApJ...738L...8W,2011ApJ...729...13C}). Finally, the volumetric TDE rate is a proxy for the mass of the black hole seeds that grow into SMBHs (\citealt{2016MNRAS.455..859S}). However, the inhomogeneous and small sample of the events currently available (about several tens\footnote{\tt http://tde.space}) probably prevents us from reaping the full potential of TDE studies. Most of the ground based surveys hunting for supernovae as well as spectroscopic follow-up observations had the tendency to avoid the central regions of host galaxies. This was largely due to various difficulties in the data processing and lower signal-to-noise of observed transients due to the core brightness. Although recent developments in difference image analysis techniques mitigate these issues (e.g. \citealt{2016ApJ...830...27Z}), the high spatial resolution afforded by {\it Gaia} and the lack of atmospheric seeing variations should also allow {\it Gaia} to resolve transients at closer angular separations to their host galaxy nuclei and enable discrimination between genuinely nuclear transients (e.g. TDEs) and near-nuclear events (e.g. circumnuclear SNe). Recently, a number of peculiar nuclear transients were discovered by various surveys (e.g. \citealt{2016NatAs...1E...2L,2017ApJ...844...46B,2017NatAs...1..865K,2017MNRAS.465L.114W}). Some of these transients were not discovered by AlertPipe, even though the sources had been detected by {\it Gaia}. Part of the goal of this paper is to investigate the reasons for this. The predicted number of detected SNe is around 1300 per year assuming a 19 magnitude minimum threshold for the brightness of the transient. About 15 per cent of these are predicted to occur in the host nuclei with offsets smaller than 1 arcsec. Moreover, $20\pm1$\footnote{The value of uncertainty from \cite{2016MNRAS.455..603B} was corrected (private communication). Poisson noise was not taken into account in the simulations by \cite{2016MNRAS.455..603B}.} TDEs should be discovered every year (\citealt{2016MNRAS.455..603B}). From mid-2016 to mid-2017 - when a stable version of AlertPipe was operating - GSA detected and published about 50 events preliminarily classified as nuclear transients\footnote{\tt http://gsaweb.ast.cam.ac.uk/alerts} (i.e. transients - likely SNe close to the host centre or AGN variability - observed within 0.5 arcsec from their host centre if the host is recognised using external catalogues) which is roughly less than 25 per cent of the expected number of supernovae and TDEs. We note that the predictions of \citet{2016MNRAS.455..603B} do not include events due to AGN variability, which are a significant contributor to the published {\it Gaia} Alerts, hence the missing fraction of nuclear transients is probably significantly larger than 75 per cent. In this study we performed a large-scale and systematic search for transient events in the nuclei of galaxies detectable by {\it Gaia} between mid-2016 and mid-2017. We started with objects classified as "galaxy" by the Sloan Digital Sky Survey Data Release 12 (SDSS DR12, \citealt{2015ApJS..219...12A}). We used a different method to search for transients than the AlertPipe daily search (\citealt{2012IAUS..285..425W,2013RSPTA.37120239H}, Hodgkin et al. in prep.). This paper is organized as follows. In Section 2 we introduce our data sample and present our transient-selection method. Next, we describe the newly found candidate nuclear transients in Section 3. In Section 4, we discuss our results and implications for {\it Gaia} Science Alerts. We conclude in Section 5. Throughout this paper we assume a flat $\Lambda$-Cold Dark-Matter ($\Lambda$CDM) concordance cosmological model of the Universe with parameters $\Omega_\Lambda = 0.7$, $\Omega_\mathrm{M} = 0.3$ and $H_0 = 70\mathrm{~km~s^{-1}~Mpc^{-1}}$, $h = 0.70$. | \cite{2016MNRAS.455..603B} predicted that {\it Gaia} Science Alerts will discover about 215 nuclear transients (from supernovae and TDEs) per year brighter than 19 mag and with an increase in magnitude of 0.3 mag or more. These sources would be discovered using both the {\it NewSource} and {\it OldSource} detectors. Our study presented here comprises about one third of the sky, and we report $\sim$160 ($\sim$480) candidates for transients brighter than 19 mag (20.5 mag). All these transients were discovered using historical data from the GSA DB (the same used by the {\it OldSource} detector). Our sample does not contain transients that would have been found by the {\it NewSource} detector. Our dedicated search for nuclear transients may be more sensitive than AlertPipe that is designed to discover all types of transients (like supernovae, cataclysmic variables, microlensing events, flare stars etc) with a low false-positive rate. Nevertheless, significant manual vetting of candidate nuclear transients has been necessary. \begin{figure} \includegraphics[width=\columnwidth]{gal_lb.eps} \caption{The map of all our candidate transients in Galactic coordinates, i.e. the 6k candidates (green dots), before our eye-balling reduced this to $\sim$480 (magenta diamonds). Most of the real transients were detected far from the Galactic plane. A significant number of false positives is located close to the Galactic disk (about 33 per cent). Hence, removing the area of the Galaxy disk and bulge would decrease the number of false positives and objects to vet without losing (many) real transients (about 50 from $\sim$480)} \label{fig:gallb} \end{figure} \begin{figure} \includegraphics[width=\columnwidth]{vonnskewhfilterlogsuccessrate.eps} \caption{The false-positive rate on the skewness vs. the reciprocal of von Neumann statistic plane. The percentage of rejected sources during the final vetting of the light curves and finding charts is presented. Only 8 per cent of selected objects on the skewness -- von Neumann parameter space were deemed to be transient events. There are regions on the skewness vs. the reciprocal of von Neumann statistic plane that are very efficient with finding nuclear transients candidates.} \label{fig:vns_success} \end{figure} Limiting the number of false positives seems to be a crucial requirement before one can consider including a new algorithm into AlertPipe. As using the von Neumann and skewness statistics needs a significant amount of time spent on eyeballing we explored the properties of the objects classified as false positives. About 33 per cent of these objects are (manually vetted) unresolved binary systems from which 74 per cent are located close to Milky Way's disc ($|b|<25^{o}$) where only 10 per cent of transient candidates were found (see Fig. \ref{fig:gallb}). Hence, removing this area of the sky increases the number of real transients. We also noticed that using more conservative cuts on the skewness -- the reciprocal of von Neumann plane might help here (see Fig. \ref{fig:vns_success}). Furthermore, one can also repeat this study every $\sim$2 weeks, and announce the candidates after vetting, given that most of these transients last months to years. In attempt to address the nature of the transient sources discovered in our search we investigated mid-IR Wide-field Infrared Survey Explorer (WISE) data. For 99.6 per cent of our sources we obtained a cross-match within 6 arcsec (approximately the FWHM of the $W1$ WISE data) of which 78 per cent has robust measurements in the three $W1, W2$, and $W3$ filters (i.e. the detection in all bands $W1, W2$, and $W3$ has a flux signal-to-noise ratio greater than 2). Comparing the WISE colour-colour diagram ($W2-W3$ vs.~$W1-W2$, see Fig.~\ref{fig:wise}) with that in \cite{2010AJ....140.1868W} we deduce that the majority of our sample of detected transients are from QSO-like objects. \begin{figure} \includegraphics[width=\columnwidth]{wise.eps} \caption{The WISE colour-colour diagram of our sample of transients discovered in the nuclei of galaxies. For 99.6 per cent of selected objects we found a corresponding source in the AllWise data base exists. We plotted the 78 per cent of objects with robust measurements in the $W1, W2$, and $W3$ filters (i.e. the source is detected in all bands $W1, W2$, and $W3$ with a flux signal-to-noise ratio greater than 2). The sample is dominated by QSO-like objects ($W1-W2>0.5$ for $\sim$74 per cent of hosts). A small fraction of our sources have WISE colours consistent with those of elliptical galaxies ($W2-W3<\sim1$). The sample also contains spirals and starburst galaxies ($W1-W2<\sim0.5$ and $W2-W3>\sim1$). The typical colours and location on the WISE colour-colour diagram for various types of objects can be found in \citet{2010AJ....140.1868W}.} \label{fig:wise} \end{figure} For a subsample of the transient sources discovered in our search with spectroscopic redshifts provided by SDSS we obtained absolute magnitude values at the light curve peak. The redshift range spans between 0 and 0.6. The histogram in Figure \ref{fig:absmag} shows this subsample separated according to the SDSS classification of the source before the occurrence of the transient events, into galaxies and quasars. The absolute magnitude range spans between -17 and -26 mag with the brightest transient occur in hosts associated with quasars where the absolute magnitude of -26 mag is not unusual for this redshift range (\citealt{2018A&A...613A..51P}). Figure \ref{fig:absmag} includes 50 spectroscopically classified galaxies, and 92 QSOs, supporting our finding from the WISE data that a significant fraction of the transients we detect come from QSO-like objects. However, about one third of our nuclear transients are associated with galaxies which are not classified as AGN by SDSS. Transients detected in these galaxies could possibly arise from circumnuclear supernovae, TDEs, or AGN switching on, after they were classified in SDSS. \begin{figure} \includegraphics[width=\columnwidth]{GNT_absmag_hist_hostsub.eps} \caption{The histogram of absolute magnitude values at the {\it Gaia} light curve peak for a subsample of the {\it Gaia} nuclear transient candidates where spectroscopic redshifts from SDSS were available. The subsample was separated into galaxies (orange dashed line) and quasars (blue line) using classification provided by SDSS. The magnitudes are host subtracted using the median value of the data from the first year of the mission (mid-2014 to mid-2015).} \label{fig:absmag} \end{figure} \subsection{Validation of {\it Gaia} Science Alerts} During the period of 1 year between July 2016 and June 2017 GSA discovered 48 transients in galaxy nuclei (i.e. transients observed within 0.5 arcsec from their host centre if the host is recognised using external catalogues). From this sample 22 events were detected by the {\it OldSource} detector. The rest of 26 transients detected by the {\it NewSource} detector could not be found by the method described here due to the lack of the historical measurements in the light curves. Sixteen events (from the sample detected in the {\it OldSource} detector) were discovered in the SDSS objects but only 5 of them are photometrically classified as galaxies. We re-discovered five transients in the centres of SDSS galaxies that were previously announced as transients by the GSA team (Gaia16ajq, Gaia17ays, Gaia17bib, Gaia17cff, Gaia17dko, see Tab. \ref{tab:alrt}). One source, Gaia16avf, that was alerted on was not re-discovered by our search. Another source, Gaia17arg, was discovered on the skewness -- von Neumann plane but removed from the final list as the second field of view was pointed on the Galactic plane during the peak. One transient in the centre of an SDSS galaxy was also not found (Gaia17bje). However, the host galaxy is fainter than 20 mag in SDSS $r$--band, hence it was not included in our search. Fifteen transients from our final sample were found by AlertPipe but then rejected through automated filtering and human visual inspection, and finally not published (see Tab. \ref{tab:unpub}). Further examination of the reasons for these rejections will be discussed in Hodgkin et al. (in prep.). \begin{table} \centering \caption{The {\it Gaia} Nuclear Transient (GNT) candidates detected by GSA AlertPipe but not published due to subsequent filtering.} \begin{tabular}{l l} \hline GNT ID & SDSS galaxy \\ \hline GNTJ003643.62$+$330622.42 & SDSSJ003643.61$+$330622.52 \\ GNTJ042910.72$-$052040.28 & SDSSJ042910.73$-$052040.25 \\ GNTJ073442.35$+$453623.24 & SDSSJ073442.36$+$453623.28 \\ GNTJ080115.97$+$110156.53 & SDSSJ080115.97$+$110156.52 \\ GNTJ081152.11$+$252521.39 & SDSSJ081152.11$+$252521.34 \\ GNTJ143701.50$+$264019.19 & SDSSJ143701.50$+$264019.18 \\ GNTJ150512.77$+$202240.70 & SDSSJ150512.78$+$202240.70 \\ GNTJ170356.27$+$231426.67 & SDSSJ170356.27$+$231426.62 \\ GNTJ171558.78$+$362323.05 & SDSSJ171558.79$+$362323.05 \\ GNTJ172027.48$+$103210.17 & SDSSJ172027.49$+$103210.17 \\ GNTJ210213.94$+$001327.17 & SDSSJ210213.93$+$001327.18 \\ GNTJ220801.33$+$304627.97 & SDSSJ220801.33$+$304628.04 \\ GNTJ232841.41$+$224847.96 & SDSSJ232841.40$+$224848.02 \\ GNTJ233520.51$+$280204.32 & SDSSJ233520.51$+$280204.25 \\ GNTJ233855.86$+$433916.86 & SDSSJ233855.86$+$433916.87 \\ \hline \end{tabular} \label{tab:unpub} \end{table} Figure \ref{fig:deltamAP} shows the comparison between the properties of transients found by the {\it OldSource} detector and by the search in this study. The {\it OldSource} detector and GSA filtering tends to only find the brighter transients with high amplitude whereas the transients detected on the skewness -- von Neumann parameter space are usually fainter and with lower amplitudes. A significant number of transients were not alerted on by the regular GSA system. There are various reasons for this situation. Sometimes multiple source IDs are assigned by the {\it Gaia} Initial Data Treatment to galaxy cores (see Subsection \ref{sec:addch} where we explain in detail how this works and how we corrected for this). We combine the magnitude measurement of different source IDs that actually belong to the same source, thereby recovering data points for the light curve of that object that increases the sensitivity of our detector. AlertPipe assumes that the {\it Gaia} Initial Data Treatment that matches sources detected during new observations with sources those previously detected on the basis of the first pass astrometric parameters of the objects works flawlessly. This helps with removing most of the close binary systems but because the centres of galaxies are not described well by a simple PSF profile more than one source might be assigned to a galaxy core. The presence of multiple entries for the same galaxy core in the GSA DB causes GSA to exclude these events from the Alerts stream. About 45 per cent of transients found in the study presented here were flagged during {\it Gaia} detection and cross-matching as confused with different sources within the GSA DB and thus discarded by the GSA. Another reason for a non-detection by the GSA system is that it currently requires a candidate transient to be detected at least once in each field of view within 40 days of each other, whereas {\it Gaia}'s scanning law implies that, like iPTF16fnl, the sky area of some transient sources is covered only by one of the two field of views in this 40-day window. For the sample presented in this paper, about 11 per cent of objects have a single detection within 40 days of the maximum of the {\it Gaia} light curve. The requirement of multiple detections mainly affects short transients as the second observation may happen when the source is back to quiescence and it is not the main cause for missing new transients by GSA (the {\it Gaia} scanning law was taken into account in simulations by \citealt{2016MNRAS.455..603B}). Moreover, 25 per cent of objects were detected more than once within the 40-day window, however the detection was each time in the same field of view meaning that these candidates are rejected by GSA. Because of these three reasons at least 56 per cent of the sources found by the method detailed in this paper could not be detected by AlertPipe. The fractions given above (45, 11, 25 per cent) cannot be directly added as a particular candidate transient might be rejected for multiple reasons. Several other AlertPipe thresholds set to reduce the number of false positives also reduce the number of nuclear transients such as the minimum difference between the magnitude of the latest photometric data point and the median from the previous detections. Similarly, the threshold that measures the difference between the historic variability (expressed in the rms of the light curve) and the significance of the latest data point is set such that many nuclear transients are missed as variability may be induced artificially due to the different angles with which {\it Gaia} scans over a galaxy and the observational windows of a rectangular shape (\citealt{2014sf2a.conf..421D}). \begin{figure} \includegraphics[width=\columnwidth]{deltammaxAlertpipe.eps} \caption{Alerting magnitude vs. amplitude. Red dots indicate all transients (not only nuclear events) alerted by the {\it OldSource} detector in AlertPipe between July 2016 and June 2017. The squares in the background show the distribution from the sample found using the skewness -- von Neumann parameter space. The dashed magenta lines indicate the AlertPipe thresholds in the {\it OldSource} detector of delta magnitude of 0.3 and 1.0 magnitude.} \label{fig:deltamAP} \end{figure} \subsection{Future improvements} In this paper we have demonstrated that the skewness -- von Neumann parameter space provides a new window into the discovery of transients with {\it Gaia}, which could be implemented in an improved version of AlertPipe. This naturally bypasses the existing requirement on having two fields-of-view, however, AlertPipe would need to be able to handle the significant number of {\it Gaia} sources which end up with split source IDs, which is non trivial for the current database design. The source astrometry is obtained from a first pass of the On-Ground Attitude determination (OGA1) during the Initial Data Treatment. Using the more accurate astrometry from the second iteration (OGA2) from subsequent data processing will likely provide a boost to the study of transients in galaxy nuclei. The accuracy of the {\it Gaia} coordinates will have improved by 1 to 2 orders of magnitude and this will allow us to determine the offset between the transient and its host nucleus. However, the position of both (transient and host galaxy) must be delivered by {\it Gaia} that makes this relevant to the events detected by the {\it OldSource} detector. This is especially true if the host is present in {\it Gaia} DR2. We notice that several candidates for nuclear transients were rejected due to {\it Gaia} internal cross-match issues, but this should be solved after publishing {\it Gaia} Data Release 2 where the majority of close binary systems should be resolved. | 18 | 8 | 1808.03984 |
1808 | 1808.05793_arXiv.txt | We study the potential of the Square Kilometre Array in the first phase (SKA1) in detecting dark matter annihilation signals from dwarf spheroidals in the form of diffuse radio synchrotron. Taking the minimal supersymmetric standard model as illustration, we show that it is possible to detect such signals for dark matter masses about an order of magnitude beyond the reach of the Large Hadron Collider, with about 100 hours of observation with the SKA1. | 18 | 8 | 1808.05793 |
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1808 | 1808.07043_arXiv.txt | Disintegrating planets allow for the unique opportunity to study the composition of the interiors of \added{small, hot, rocky} exoplanets because the interior is evaporating and that material is condensing into dust, which is being blown away and then transiting the star. Their transit signal is dominated by dusty effluents forming a comet-like tail trailing the host planet (or leading it, in the case of K2-22\added{b}), making these good candidates for transmission spectroscopy. To assess the ability of such observations to diagnose the dust composition, we simulate the transmission spectra \added{from 5-14 $\mu$m } for the planet tail assuming an optically-thin dust cloud comprising a single dust species\added{ with a constant column density scaled to yield a chosen visible transit depth}. We find that silicate resonant features near 10 $\mu$m can produce transit depths that are at least as large as those in the visible. For the average transit depth of 0.55\% in the \textit{Kepler} band for K2-22\added{b}, the features in the transmission spectra can be as large as 1\%, which is detectable with the \textit{JWST} MIRI low-resolution spectrograph \added{in a single transit}. The detectability of compositional features is easier with an average grain size of 1 $\mu$m despite features being more prominent with smaller grain sizes. We find most features are still detectable for transit depths of $\sim 0.3$\% in the visible range. If more disintegrating planets are found with future missions such as the space telescope \textit{TESS}, follow-up observations with \textit{JWST} can explore the range of planetary compositions. | 18 | 8 | 1808.07043 |
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1808 | 1808.07872_arXiv.txt | \csznote{ We analyze the metal accumulation in dwarf and spiral galaxies by following the history of metal enrichment and outflows in a suite of twenty high-resolution simulated galaxies. These simulations agree with the observed stellar and gas-phase mass-metallicity relation, an agreement that relies on large fractions of the produced metals escaping into the CGM. For instance, in galaxies with Mvir ~ 1e9.5 -- 1e10 solar masses, we find that about ~ 85% } We analyze the metal accumulation in dwarf and spiral galaxies by following the history of metal enrichment and outflows in a suite of twenty high-resolution simulated galaxies. These simulations agree with the observed stellar and gas-phase mass-metallicity relation, an agreement that relies on large fractions of the produced metals escaping into the CGM. For instance, in galaxies with $M_{vir} \sim 10^{9.5} - 10^{10} \Msun$, we find that about $\sim$ 85\% of the available metals are outside of the galactic disk at $z = 0$, although the fraction decreases to a little less than half in Milky Way-mass galaxies. In many cases, these metals are spread far beyond the virial radius. We analyze the metal deficit within the ISM and stars in the context of previous work tracking the inflow and outflow of baryons. Outflows are prevalent across the entire mass range, as is reaccretion. We find that between 40 and 80\% of all metals removed from the galactic disk are later reaccreted. The outflows themselves are metal enriched relative to the ISM by a factor of 0.2 dex because of the correspondence between sites of metal enrichment and outflows. As a result, the metal mass loading factor scales as $\eta_{metals} \propto v_{circ}^{-0.91}$, a somewhat shallower scaling than the total mass loading factor. We analyze the simulated galaxies within the context of analytic chemical evolution models by determining their net metal expulsion efficiencies, which encapsulate the rates of metal loss and reaccretion. We discuss these results in light of the inflow and outflow properties necessary for reproducing the mass-metallicity relation. | \label{sec:intro} Galaxies evolve through a balance between gas accretion and outflows. Cosmological accretion of gas from in the inter-galactic media enables the continued growth of halos \citep{Nelson2013}, and metal-poor cold-gas accretion has been detected through absorption on the outskirts of galaxies \citep{Kacprzak2012, Bouche2013,Crighton2013}. Additionally, reaccretion of previously ejected material provides continued fuel and can easily dominate over cosmological accretion in galaxies with halo masses \textgreater$10^{11} - 10^{12} \Msun$ \citep{Oppenheimer10}. This material exists as metal-enriched gas in the circumgalactic media (CGM) prior to its reaccretion \citep{Cheung2016}. Meanwhile, gas loss from galaxies is accomplished through feedback-driven outflows. Such outflows are ubiquitous in high-redshift star forming galaxies and local starburst galaxies \citep[e.g.][]{Heckman1990,Pettini2001,Shapley2003,Martin05, Weiner09,Steidel10,rubin13}, and both semi-analytic models and simulations have found them necessary to reproduce key observations such as the stellar mass-halo mass relation \citep{Henriques2013,scannapieco11,Stinson12,Hopkins2013, White15}. Together with accretion, outflows set the baryonic content within the disk and regulate star formation \citep{Dave11b, Lilly2013, Dekel2014, Christensen2016}. In addition to regulating the baryonic content in galaxies, outflows are key to establishing their metal content. For example, comparisons of the total metals within the interstellar media and stellar disk to the total mass of metals produced predict that 20-25\% of metals remain in the stars and ISM of Milky Way-mass galaxies \citep{Peeples2014} and 6\% remained within the stars and ISM of a dwarf galaxy \citep{McQuinn2015}. As a result, outflows are a leading candidate to regulate the metallicity within the disks of galaxies and establish the mass-metallicity relation (MZR)~\citep[e.g.][]{tremonti04,Finlator08,Ma2016} and its second-parameter dependences on star formation rate and gas content~\citep{Dave11b}. The amplitude and slope of the MZR can be explained by the tendency of outflows to reduce the effective yield and by the greater efficiency of outflows in removing material from low-mass halos in combination with their reduced star formation efficiency. While analytic models can explain the MZR by parameterizing metal inflow and outflow efficiencies \citep{ erb08, Spitoni2010, Peeples2011, Dave2012, Lilly2013} these models generally do not account for the reaccretion of metal-enriched material. Additionally, many of these models assume that outflows share the same metallicity as the ISM, while observations show evidence for metal enrichment compared to the ISM \citep{Chisholm2016}. Understanding the rates of reaccretion and the relative enrichment of outflows are key to understanding the source of the MZR. A corollary to the outflow-driven metal depletion of disk material is the redistribution of metals to the CGM and beyond. Since metals originate primarily in the stellar disks of galaxies, their presence throughout the CGM provides a tracer of the history of inflows and outflows. In particular, strong transport of metals by galactic outflows is indicated by the large, oxygen-rich halos surrounding present-day \citep{Tumlinson2011,Prochaska2011a} and high redshift \citep{Lehner2014} star forming galaxies. Observations of metal line absorption around dwarf galaxies \citep{Bordoloi2014}, around the Andromeda galaxy \citep{Lehner2015}, and throughout the intergalactic medium (IGM) \citep[e.g.][]{Cooksey2013, DOdorico2013, Shull2014} provide additional evidence for outflow-driven enrichment. On the theoretical side, simulations generally require strong outflows from stellar feedback in order to reproduce the rapidly-advancing observations of metal lines around galaxies \citep[e.g][]{Stinson12, Shen13, Ford13a,Hummels13,Suresh2015}. Metal-line absorption observations also provide a range of constraints to the thermal and dynamical state of the CGM and indicate a primarily-bound, multiphase CGM with photoionized and/or collisionally ionized gas embedded in a hotter low-density medium \citep[for a review, see][]{Tumlinson2017}. Galaxy formation simulations can both provide the provenance of metals in the disk and halo and establish the history of metal accretion and outflow. Simulations thus far have primarily focused on examining total baryonic mass loss and reaccretion \citep[e.g.][]{Oppenheimer10,Woods2014,Christensen2016,Muratov2015,AnglesAlcazar2016}. They have tended to converge on mass loading factors with mass scalings between those expected for momentum and energy-conserving winds. Simulations also have tended to agree that recycling of material is common, fuels late-time star formation \citep{Oppenheimer10,Woods2014}, and modifies the angular momentum profile \citep{Brook11b,Ubler2014, Christensen2016}. However, the fate of outflowing gas, including the rates and timescales of outflow reaccretion, are highly model dependent, illustrating the importance of largely unexamined processes happening within CGM. The examination of metals within and surrounding galaxy halos can help delineate between models by, for instance, tracing the eventual distribution of stellar-enriched material. As an example of this type of theoretical investigation, \citet{Shen2011} found satellite progenitors and nearby dwarf galaxies to be the source of 40\% of metals within 3 $R_{vir}$ of a $z = 3$ progenitor of a Milky Way-mass halo. In a different investigation, \citet{Muratov2017} found high recycling rates of metals at early times and in low-mass galaxies, leading to similar metallicities of inflowing and outflowing material within the central halos. In contrast, outflows from their L$^*$ galaxies at low redshift were very weak, leading the to the accumulation of metals within stars. We expand upon these types of studies by following the accumulation of metals within galaxies by tracing gas flows. Following on the work of \citet{Christensen2016}, we use a suite of galaxy-formation simulations to quantify the cycle of metal production, loss, and accretion over two and a half-orders of magnitude in virial mass. By tracking the history of smoothed particle hydrodynamic gas particles, we identify instances of accretion and ejection, determine the eventual location of the metals produced by the galaxy, and measure the metallicity of the outflows. In \S\ref{sec:methods}, we present the suite of simulations and describe the analysis. Results are presented for the redshift zero metal census (\S\ref{sec:census}) and metal distribution (\S\ref{sec:metaldist}), the history of metal cycling (\S\ref{sec:metalhist}), the metallicity of outflows (\S\ref{sec:outflowmetals}), and the metal mass loading factor and (\S\ref{sec:massloading}). These results are discussed in light of the MZR and other work (\S\ref{sec:discuss}), and our conclusions are presented in \S\ref{sec:conclude}. | \label{sec:conclude} In this paper, we analyze the drivers of the metal distribution in a set of twenty simulated galaxies spanning two and a half orders of magnitude in halo mass that match observed characteristics, including the stellar and gaseous MZR. We follow the accumulation of metals in galaxies by tracking their metal production and by identifying instances of metal accretion and loss. This analysis enables us to determine the role of galactic winds and reaccretion in determining the mass of metals within different components of a galaxy and its CGM. \begin{enumerate} \item Gas outflows are highly effective at removing metals from galactic disks. In the lowest mass galaxies, as little as 10\% of the metals produced may remain within the ISM or stars at $z = 0$, while that fraction rises to as much $\sim 50$\% for Milky Way-mass galaxies. This mass trend is dominated by the stars. While the ISM retains a similar range of fractions (between 5 and 25\%) of the produced metals across galaxy mass, the fraction locked in stars rises steeply with galaxy mass. Those metals that do exit the disk of the galaxy are widely dispersed with the majority lying beyond the virial radius. Because of their deeper potential wells, more massive galaxies are generally better able to retain their metals within their virial radius and show lower dispersal of metals. However, this mass trend becomes complicated for dwarf galaxies, as the very lowest mass galaxies are able to retain a moderate fraction of their metals, despite their shallow potentials. In these galaxies, extremely low star formation rates are likely responsible for the reduced metal loss. \item The history of metal enrichment of the ISM and stars largely follows the history of metal production by stars with less than 10\% of metals at $z=0$ coming from externally accreted gas or stars. Large amounts of metals cycle rapidly in and out of the disk with the cumulative history of gas loss sometimes exceeding the mass of metals produced by a factor of three. The majority of these metals (generally between 50 and 80\%), however, are quickly returned to the disk. The fraction of these metals that become dynamically unbound from the disk (``ejected") is much higher in low-mass than high-mass galaxies. These ejected metals are more likely to permanently remain outside of the disk. \item % Ejected material tends to be somewhat more metal rich than the ambient ISM, because gas that receives energy from SNe most likely also received metals from the same stellar population. This effect is largely independent of galaxy mass and is generally stronger at high redshift when the ISM metallicity at a given virial mass would have been lower. Gas that is removed from the disk, but not necessarily considered part of an outflow, also shows some amount of metal enrichment, especially at early times, but the effect is reduced, presumably because this material includes gas not as strongly heated by SNe. Because of dilution by low-metallicity external accretion, the average metallicity of accreted material is lower than outflowing material and, in most cases, the ISM. However, even when excluding external accretion, the metallicity of material reaccreted after being removed from the disk is lower than the metallicity of the removed material. This difference in metallicity indicates a tendency for some highly enriched material to remain outside the disk. \item The metal mass loading factor, $\eta_{metals} = Z_{W} \frac{\dot M_W}{SFR}$, shows a power-law dependency on virial velocity, similar to the standard mass loading factor. However, we observe a flattening of the power from $a = -2.2$ to $a = -0.91$ when comparing the metal mass loading to the standard mass loading. This flattening can be entirely explained by the reduced metallicity of the ISM in lower mass galaxies. Therefore, while low mass galaxies have low metallicity ISM and are no more likely to preferentially eject metals than higher mass galaxies, their exceptionally high mass loading factors still produce high rates of metal loss per stellar mass formed. \item The MZR can be explained using an alternative depiction of the metal expulsion efficiency, $\zeta_{net}= \frac{Z_W \dot M_W - Z_{reaccr} \dot M_{reaccr}}{Z_{ISM} \dot M_{SFR}}$, which scales the net metal loss rate from winds by the star formation rate and ISM metallicity. We find that $\zeta_{net}$ scales as $v_{circ}^{-1.7}$, which is consistent with what would be expected from the observed MZR. We find the outflowing material to be enriched to the same degree relative to the ISM independent of stellar mass, while reaccretion rates of metals removed from the disk are only slightly higher in more massive galaxies. As a result, the metal expulsion efficiency, $\zeta_{net}$ shares a similar, though slightly shallower, scaling with $v_{circ}$ as the mass loading factor. \end{enumerate} These results illustrate how outflows and gas recycling, in combination with accretion and varying star formation efficiency, can together produce the MZR. The simulations naturally produce the mass-loading factors of energy-driven winds, with slight metal enhancement and moderate rates of reaccretion across all halo masses. However, this is not a unique solution to producing the MZR, and further comparisons to the metal content, distribution, and thermodynamic structure of the CGM are necessary to constrain how exactly gas transfer between the ISM and CGM determines the baryonic and metal content of galaxies. | 18 | 8 | 1808.07872 |
1808 | 1808.09929_arXiv.txt | Thirty-three fast radio bursts (FRBs) had been detected by March 2018. Although the sample size is still limited, meaningful statistical studies can already be carried out. The normalised luminosity function places important constraints on the intrinsic power output, sheds light on the origin(s) of FRBs, and can guide future observations. In this paper, we measure the normalised luminosity function of FRBs. Using Bayesian statistics, we can naturally account for a variety of factors such as receiver noise temperature, bandwidth, and source selection criteria. We can also include astronomical systematics, such as host galaxy dispersion measure, FRB local dispersion measure, galaxy evolution, geometric projection effects, and Galactic halo contribution. Assuming a Schechter luminosity function, we show that the isotropic luminosities of FRBs have a power-law distribution that covers approximately three orders of magnitude, with a power-law index ranging from $-1.8$ to $-1.2$ and a cut off $\sim 2\times 10^{44}\,\ergs$. By using different galaxy models and well-established Bayesian marginalisation techniques, we show that our conclusions are robust against unknowns, such as the electron densities in the Milky Way halo and the FRB environment, host galaxy morphology, and telescope beam response. | Fast Radio Bursts (FRBs) are a new type of radio transients, and remain unexplained. The bursts last for a few milliseconds, and show dispersive signatures with peak flux densities ranging from 0.3~Jy to about 100 Jy. The first one (FRB~010724, \citealt{Lorimer07Sci}) was detected serendipitously in the archival data of pulsar survey for Small Magellanic Cloud using the Parkes telescope \citep{MFL06}. Shortly after that, a growing number of FRBs were discovered with Parkes at 1.4 GHz, both in the archival data \citep{Keane12MN,Thornton13Sci, Burke-Spolaor14ApJ} and from the real-time searches \citep{Ravi15ApJ,Petroff15MN, Keane16Nat, Ravi16Sci, Petroff17MN, Bhandari18MN}. FRBs were also detected by Arecibo \citep{Spitler14ApJ}, Green Bank Telescope \citep{Masui15Nat}, UTMOST \citep{Caleb17MN, Farah18MN} and ASKAP \citep{Bannister17ApJ}. At the time of writing this paper, the total number of the reported detections adds up to 33. FRBs are natural celestial probes with a broad range of astrophysical applications. For example, it has been proposed that FRBs could be used to test the Einstein's equivalence principle \citep{Wei15PRL, Tingay16ApJ, Zhang16arXiv}, to constrain the rest mass of photons \citep{Wu16ApJ, Bonetti16PLB, Bonetti17PLB, Shao17PRD}, to detect the baryon contents in the Universe \citep{McQuinn14ApJ}, to probe the cosmological matter distribution \citep{Masui15PRL}, to study the evolution of intergalactic medium (IGM) \citep{Zheng14ApJ} and constrain the dark-energy equation of states \citep{Zhou14PRD, Gao14ApJ}. The origins of FRBs, however, remain mysterious and subject to an intensive debate. Here, we list several proposals in the literature in chronological order: (1) radio pulses from black hole evaporative explosions \citep{Rees77Nat}; (2) superconducting cosmic strings \citep{CSV12, CSS12, Yu14JCAP}; (3) flaring magnetars \citep{PP10, PP13arXiv} or stars \citep{Loeb14MN}; (4) mergers of white dwarfs \citep{Kashiyama13ApJ}; (5) mergers of double neutron stars \citep{Totani2013PASJ, Wang16ApJ}; (6) collapses of neutron stars into black holes \citep{FR14A&A, Zhang14ApJ}; (7) synchrotron masers \citep{Lyubarsky14MN, Ghisellini2017MN, Lu18MN}; (8) binary model of white dwarf and black hole \citep{GDL16}; (9) super-giant pulses from pulsars \citep{Cordes16MN, Connor16MN} ; (10) radio emission from soft gamma-ray repeaters \citep{Pen15ApJ, Katz16ApJ}; (11) axion stars \citep{Iwazaki15PRD}; (12) quark nova \citep{Shand16RAA}; (13) mergers of charged black holes \citep{Zhang16ApJL, Liu16ApJ}; (14) collisions between pulsar and asteroids \citep{Geng15ApJ, Dai16ApJ}; (15) relativistic jet -- cloud interactions \citep{Romero16PRD, Vieyro17A&A}; (16) births of millisecond magnetars \citep{Metzger17ApJ}; (17) `cosmic comb', i.e. magnetosphere -- environment interactions\citep{Zhang17ApJ, Zhang18ApJ}; (18) accretion of black holes \citep{Katz17MN}; (19) star-quakes of compact stars \citep{Wang18ApJ}. To understand the mechanisms of FRBs, the host galaxy information is crucial. At this stage, only the repeating FRB, FRB~121102, had the reliable identification of host galaxy \citep{Spitler16Nat, Scholz16ApJ}. \citet{Chatterjee17Nat} measured its precise position using \emph{Karl G. Jansky Very Large Array}. The optical counterpart was identified as a dwarf galaxy at the redshift of $z=0.193$ \citep{Tendulkar17ApJ}. However, we should be cautious in drawing general remarks on the FRB environment, due to unknown links between repeating and non-repeating FRBs. Statistical analyses are needed to quantify the properties of FRBs as an integrated population. The normalised luminosity function, i.e. the probability density function (PDF) of FRB luminosities, is one particularly important statistics for the FRB intrinsic power output. The computation of the luminosity functions requires not only FRB flux and distances, but also a detailed account of any biases in the sample. For example, without the counterpart identifications, the FRB distances are usually estimated via the dispersion measure (DM). The estimated FRB distance and luminosity are affected by the uncertainties in the DM modelling. It is absolutely necessary to account for these effects in inferring the luminosity function. There are several algorithms to measure the luminosity function (see \citet{Will97} for a review). The non-parametric methods (e.g. \citealt{LB71}) usually require certain uniformity of data coverage to be applicable. The likelihood-based methods \citep{MTA83} or Bayesian methods \citep{KFV08, CLM13} are preferable for the FRB problems, because these algorithms are more flexible in modelling the systematics and less constrained by the conditions of a given sample. In this paper, we aim to measure the normalised FRB luminosity function. To include systematics and unknowns in the statistical inference, we develop a Bayesian framework suitable for the current problem. For most of the known FRBs, there are four main observables relevant to the luminosity function determination: flux density, bandwidth, duration, and dispersion measure. Compared to the other astronomical sources whose luminosity functions are measured, the FRB distance is not directly available. As a result, we have to rely on the dispersion measure to indirectly infer the FRB distance. Our method to measure the FRB luminosity function includes three major steps: (1) mitigate the Galactic foreground contribution of the dispersion measure; (2) model the FRB host galaxy and the cosmological dispersion measure contribution; (3) include dispersion measure models in the Bayesian luminosity function inference, and marginalise the unknowns. The first step is straightforward, as good knowledge on the Galactic electron distribution is available. The second step is to model the effects of some unknown properties on determining the luminosity function. The third step is to use a Bayesian method developed in this paper to `enumerate' all possibilities and include the unknowns in the statistical inference. We can then determine the contribution of the unknowns to statistical errors, e.g. we can make sure that the confidence bounds of inferred parameters contain the uncertainties in the modelling. The paper is organised as follows. In \SEC{sec:frbobs}, we explain how we remove the dispersion measure contribution from the Galactic foreground. In \SEC{sec:meth}, we describe our Bayesian inference method. The likelihood function is built in \SEC{sec:likf}, with detailed modelling of its components in the rest of the subsections of \SEC{sec:meth}. The computational method for posterior evaluation is shown in \SEC{sec:post}. Our results are given in \SEC{sec:res}, with discussion made in \SEC{sec:disc}. For the readers' convenience, we summarise the symbols used throughout this paper in \TAB{tab:not}. | \label{sec:disc} In this paper, we measured the FRB luminosity function using the Bayesian method under different assumptions for the host galaxy type. The Bayesian method helped dealing with the missing information, i.e. the distances of FRBs and beam response, which are difficult to handle otherwise. Assuming the Schechter form for the luminosity function, we measured the power law index and high cut-off luminosity, where $\alpha\simeq-1.5$, and $L^*\simeq10^{44}\,\ergs$. As byproducts, we also used the Bayesian method (see \APP{app:probl}) to infer the most-probable redshift, isotropic luminosity and energy of each source with the values listed in \APP{app:dattab}. The FRB luminosity power-law indices, we measured, range from $-1.8$ to $-1.2$. Such values also agree with (i) the power-law indices of pulsars' giant pulse flux distribution at lower frequency (--1.7, \citealt{KSL12}); (ii) the mean power-law indices of radio emission of pulsars ($-1.65$ to $-2.2$, \citealt{HWX16, JVK18}); (iii) the power-law index of luminosity function of long gamma-ray bursts ($-1.3$ to $-2.3$; \citealt{SZL15,PGS16}) (iv) short gamma-ray bursts ($-1.5$ to $-1.7$, \citealt{SZL15}); (v) compact binary mergers ($-1.2$ to $-1.7$, \citealt{CYZ18}). We can not pin down the radiation mechanisms based on the FRB luminosity function. However, the similarity between it and those of of other astrophysical sources may suggest a common underlying rule of defining burst populations of different kinds. The distance information of FRBs is determined from the DM values. We modelled the DM from three major contributions, i.e. from the Milky way, the IGM, and the FRB host galaxy. We also compared the results to evaluate the effects of Galaxy halo contribution. We showed that the parameters for the luminosity function are rather insensitive to the modeling details. We modelled the electron density distribution functions for two major cases in the paper, i.e. spiral galaxies and elliptical galaxies. The most likely values of $\DMh$ for these two cases are approximately 10 and 15 $\cmpc$, respectively. Such host galaxy DM values are at least one order of magnitude smaller than that of the IGM contribution. Although the most uncertain part in our modelling is the $\DMh$ distribution, the parameters inference for the luminosity function is rather robust as $\DMe\gg\DMh$. We can tolerate the missing information such as the the $\Ha$ filling factor, the stellar distribution in galaxies, halo DM or FRB source DM. The characteristic host galaxy DM values we estimated are less than 100 $\cmpc$. Despite this, considering the scattering of the distribution, our results are still compatible with the values estimated before \citep{Thornton13Sci, XH15RAA, Yang17ApJ}. The average DM value of ETGs we calculated here is higher than that for LTGs. This is mainly due to the stellar distribution and galaxy morphology. The concentration of FRBs in the central region of ETGs produce higher value of DM for the ETGs than for the LTG, where a lower scale height of LTGs leads to a lower DM. For the case of LHEGs, i.e. elliptical galaxies with $\Ha$ luminosity function in \citet{Nakamura04AJ}, the most likely DM host is $\DMh\simeq3000\,\cmpc$. Considering that the observed $\DMh$ is smaller by a factor $(1+z)$ and the roughly linear increase of $\DMi$ with redshift, one obtains that an FRB with $z>2$ would have a $\DMe$ exceeding $2750\,\cmpc$ \citep{Zhang18arXiv} which is larger than the maximum observed $\rm DM_E$ (e.g. $2583.1\,\cmpc$ from FRB~160102, \citealt{Bhandari18MN}). If FRBs all have LHEG hosts, the probability of detecting one FRB with $\DMe\le2750$ is only $\simeq35\%$, as computed by integrating $\fd(\DMh)$ from $0$ to $2750$. Thus there is only a miniscule chance ($8\times10^{-16}$) to observe all 33 FRBs with $\DMe\le 2750$~cm$^{-3}$~pc. We conclude that it is unlikely that the LHEGs are the host galaxies for FRBs, unless all FRBs lie around the galaxy outskirts if they originate in LHEGs. The $\DMh$ distribution function of all the galaxies enables us to infer the corresponding isotropic luminosity and energy of FRB emission as listed in \TAB{tab:frbs}. Using only DM as the distance indicator, our inferred most probable redshift of FRB 121102 ranges from 0.198 to 0.271 at a $2\sigma$ confidence level. This is roughly consistent with the true redshift 0.193 \citep{Tendulkar17ApJ}. The slightly higher value of the inferred redshift may be resulted from the long tail of PDF for $\DMh$ as computed in \SEC{sec:dms}. The excess may also come from an underestimate of the electron density in the Milky Way halo or in the FRB environment. Alternatively, it could be due to the deviation of the mean cosmological $\DMi$ due to the existence of large scale structures, so that the line of sight towards FRB 121102 may have pieced through an over-dense region in the IGM. We expect that more FRBs with optically-measured redshift will help us to clarify these issues. We used M87 and Milky Way as the template galaxies in this study. The choice is made because they are the only two galaxies of each type that previously have accurate measurements on both electron density distributions and $\Ha$ luminosities. As a caveat, both galaxies may not be the typical example of ETGs or LTGs. The M87 is a giant elliptical galaxy, and the Milky way has a relatively low gas fraction \citep{KE2012AR}. We can still use Milky Way and M87 as the reference values, thanks to our scaling method, which accounts for the galaxy size, electron density, and star formation history evolution. We assumed that the FRB distribution in the galaxies follows the stellar distribution. In Milky Way, the steller distribution has low scale height than that of the neutron star distribution. If the FRBs are of a magnetar or pulsar origin, the host galaxy DM may be slightly overestimated here. However, since the host galaxy DM is already smaller than the observed DM, such a second-order small perturbation can be well neglected without affecting the luminosity function inference. We modelled the luminosity distribution using the Schechter function. The measured cut-off luminosity $\log L^*\simeq44.2\,\ergs$ with an error of 0.3 dex indicates that the simple power-law distribution is not precise enough at the high luminosity end. This also helps planning future FRB surveys. For FRBs with distances of $\sim$ 1 Gpc, the size of a radio telescope for FRB survey should be at least 10 meters to get ${\rm S/N}\ge 7$. The possible off-centre position of an FRB in the main beam, without modelling, leads to an underestimate for the FRB luminosity \citep{Niino18ApJ}. We include such an uncertainty through the Bayesian marginalisation. It turns out that the difference in the parameters of inferred luminosity function is not significant between the two cases, regardless of whether or not the beam response is taken into account. Without modelling the beam response, the power-law index of Schechter function is slightly flatter, and the cut-off luminosity is relatively lower. Taking the case of ALGs-YMW16 as an example, we get $\alpha=-1.56^{+0.21}_{-0.20}$ and $\log L^*=44.19^{+0.22}_{-0.24}$ with no beam response modelling, whereas $\alpha=-1.57^{+0.19}_{-0.21}$ and $\log L^*=44.31^{+0.22}_{-0.27}$ with beam response marginalisation. As the difference is still within 1$\sigma$ confidence level, we conclude that the beam response plays a limited role in parameter inference for the current limited sample of FRBs. We could only obtain the upper limits of the lower cut-off luminosity, i.e. $L_0$, due to the limited sample of FRBs (\TAB{tab:frblf}). The current upper limit of $\log L_0<42.0$ is not capable of testing the FRB model yet. In order to measure the true value of $L_0$, a large number of nearby FRBs are required. Due to the unknown spectral shape and width, our luminosity function is based on the reference bandwidth of 1 GHz. This is motivated by the observations of the repeating FRB 121102, which indicates a $\sim$ 1 GHz bandwidth \citep{GSP18}. In general, the parameter $L^*$ in the luminosity function scales with the reference bandwidth. The assumption of a $\sim$ 1 GHz bandwidth can be revised later. Little information is available for the spectrum of FRBs at present, and scintillation may introduce a strong bias in determining the true bandwidth. The measurement in this paper can be further improved, if future observations will provide more information. We expect that the upcoming large field-of-view facilities, e.g. CHIME\citep{Ng17CHIME}, ASKAP\citep{Macquart10PASA}, MeerKAT\citep{Booth12AfrSk} and instruments with higher sensitivity, e.g. ALFABURST survey \citep{Foster17ALFABURST}, FAST \citep{Nan11}, SKA \citep{Macquart15SKA}, and QTT \citep{Wang17}, will provide more opportunities to detect more nearby FRBs and reveal the details of the FRB spectra. | 18 | 8 | 1808.09929 |
1808 | 1808.02877_arXiv.txt | In this Letter, we study the implications of string Swampland criteria for dark energy in view of ongoing and future cosmological observations. If string theory should be the ultimate quantum gravity theory, there is evidence that exact de Sitter solutions with a positive cosmological constant cannot describe the fate of the late-time universe. Even though cosmological models with dark energy given by a scalar field $\pi$ evolving in time are not in direct tension with string theory, they have to satisfy the Swampland criteria $|\Delta\pi|<d\sim\mathcal{O}(1)$ and $|V'|/V>c\sim\mathcal{O}(1)$, where $V$ is the scalar field potential. In view of the restrictive implications that the Swampland criteria have on dark energy, we investigate the accuracy needed for future observations to tightly constrain standard dark-energy models. We find that current 3-$\sigma$ constraints with $c \lesssim 1.35$ are still well in agreement with the string Swampland criteria. However, Stage-4 surveys such as Euclid, LSST and DESI, tightly constraining the equation of state $w(z)$, will start putting surviving quintessence models into tensions with the string Swampland criteria by demanding $c<0.4$. We further investigate whether any idealised futuristic survey will ever be able to give a decisive answer to the question whether the cosmological constant would be preferred over a time-evolving dark-energy model within the Swampland criteria. Hypothetical surveys with a reduction in the uncertainties by a factor of $\sim20$ compared to Euclid would be necessary to reveal strong tension between quintessence models obeying the string Swampland criteria and observations by pushing the allowed values down to $c<0.1$. In view of such perspectives, there will be fundamental observational limitations with future surveys. | Einstein's theory of General Relativity (GR) is still the standard effective field theory for the gravitational interaction below the Planck scale. Having survived a multitude of empirical tests in a wide range of scales, it continues to stand out as the most compelling candidate theory. It thus constitutes the bedrock upon which all of our effective field theories of gravity are constructed (see e.g. \cite{Heisenberg:2018vsk} for a recent review). However, some tenacious challenges remain, concerning in particular its UV completion into a quantum gravity theory and the necessity of enigmatic ingredients in form of dark energy and dark matter. A prevailing view is that GR could be just the low energy limit of the more fundamental string theory, and that some of the IR/UV problems could be cured by particularities of the UV completion. String theory refers to physical models that, instead of describing elementary particles as point-like objects in space-time in the familiar models of quantum field theory, introduces strings as fundamental objects. Originally, string theories were used to describe the strong interaction, where the gluons were interpreted as spatially extended strings between the quarks. String theory has received much interest, however this time as a candidate for a unified theory combining the standard model of particle physics with gravity. Should string theory be the ultimate theory of quantum gravity, the question immediately arises whether the effective field theories of gravity known to us can naturally be embedded into string theory. In this context, the Swampland \cite{Vafa2} emerges as the inhabitable landscape of field theories that are incompatible with quantum gravity. The implications of these Swampland criteria are tremendous (see e.g. \cite{recent} for some discussions). There is evidence that stable de-Sitter vacua in critical string theory do not exist (see e.g. \cite{Obied:2018sgi,noLambda} for some recent references) \footnote{On the other hand, specific constructions of at least metastable de Sitter vacua arising from string theory have been proposed using effective field theory techniques.}. If the late-time universe is dominated by a dark energy scalar field evolving with time, string-theory criteria can be used to constrain these dark-energy models. In this Letter, we discuss the implications of the string Swampland criteria for scalar field dark energy. | In this Letter, we have investigated the implications of the string Swampland criteria for dark-energy models based on a scalar field in view of ongoing and future cosmological observations. Should string theory be the ultimate quantum gravity theory, then, according to the second of the swampland conjectures, de Sitter solutions with a positive cosmological constant cannot account for the fate of the late-time universe. Dark-energy models based on scalar fields evolving in time still have to satisfy certain criteria in order to remain outside the Swampland \footnote{Note that string theory does not naturally lead to scalar fields with the energy scale required to be a candidate for quintessence. Rather, string theory may require radically new ideas to explain dark energy.} . We have studied the observational implications of future surveys on such quintessence models and the associated restrictions on the fate of dark energy. Current limits on $w_0$ and $w_a$ impose already a quite tight upper bound on $w(z)$. According to our analysis, current 2-$\sigma$ limits demand that $w(z) \lesssim -0.91$ at $z \approx 0.3$. This constrains the parameter $\lambda$ to $\lambda \lesssim 0.9$, which is still well in agreement with the string Swampland criteria. The projected limits to be obtained from Stage-4 surveys can be expected to change the situation substantially. Applying 3-$\sigma$ limits taken from the Euclid Red Book, we find that an upper bound on $w(z)$ will be $w(z) \lesssim -0.97$ in the redshift interval $0\le z\le 0.6$. This will restrict the allowed range for $\lambda$ to $\lambda\lesssim 0.4$. Surviving quintessence models would thus be driven into substantial tension with the string Swampland criteria. Should a future survey be able to lower the uncertainties in $w_0$ and $w_a$ to half the uncertainties expected from Euclid, $w(z)$ would have to fall below $-0.985$ in the redshift range $0\le z\le 0.6$ and thus only leave room for small deviations of quintessence from a cosmological constant. The parameter $\lambda$ would then have to fall below $\lambda\lesssim0.3$. These possibly idealised constraints raise the question of how tightly any future survey will ever be able to differentiate between whether the dark energy is a cosmological constant or not. We estimated that constraining $\lambda$ to $\lambda\lesssim0.1$ would require a reduction in the uncertainties of $w_0$ and $w_a$ compared to those expected from Euclid by a factor of $\sim20$. This would correspond to increasing the survey volume by a factor of $\sim400$ compared to that covered by Euclid. In view of such perspectives, there will be fundamental observational limitations on lowering $\lambda$ to $\lambda\lesssim0.1$ with future surveys. | 18 | 8 | 1808.02877 |
1808 | 1808.05720_arXiv.txt | { The Plastic Scintillator Detector (PSD) of the DArk Matter Particle Explorer (DAMPE) is designed to measure cosmic ray charge (Z) and to act as a veto detector for gamma-ray identification. In order to fully exploit the charge identification potential of the PSD and to enhance its capability to identify the gamma ray events, we develop a PSD detector alignment method. The path length of a given track in the volume of a PSD bar is derived taking into account the shift and rotation alignment corrections. By examining energy spectra of corner-passing events and fully contained events, position shifts and rotations of all PSD bars are obtained, and found to be on average about 1mm and 0.0015 radian respectively. To validate the alignment method, we introduce the artificial shifts and rotations of PSD bars in the detector simulation. These shift and rotation parameters can be recovered successfully by the alignment procedure. As a result of the PSD alignment procedure, the charge resolution of the PSD is improved from $4\%$ to $8\%$ depending on the nuclei. | % \label{sect:intro} The DAMPE is a space-borne satellite of China that operates in solar synchronous orbit at an altitude of 500 km for more than two years. The payload of DAMPE is a high-energy cosmic ray detection system equipped with four sub-detectors (Chang. et al. \cite{chang17}): a Plastic Scintillator Detector (PSD) (Yu. et al. \cite{yu}), a Silicon Tungsten tracKer-convertor (STK) (P. Azzarello et al. \cite{azz16}), a BGO calorimeter (BGO) (Zhang. et al \cite{zhang16}), and a NeUtron Detectors (NUD)(He. et al. \cite{he2016}). With this design, DAMPE can measure the charge, the energy and the incoming direction of the cosmic rays. The PSD, as a thin material detector, is designed to detect the charge of the cosmic ray via measuring its energy deposition in plastic scintillator and also serves as a veto detector to identify gamma-rays from charged particles. The STK mounted below the PSD is a silicon-strip tracker with 3 layers of thin tungsten plates inserted below the first, second and third detection layer. By this design, the high-energy gamma-rays can be converted into $e^+/e^-$ pairs and then their trajectories can be reconstructed. STK is also designed to reconstruct the trajectories and measure the absolute charge (Z) of cosmic-ray ions. The BGO is a 3D imaging absolute-absorption calorimeter, which is designed to measure the energy of electrons and gamma-rays from few GeV to 10 TeV and the energy of cosmic ray nuclei from 10 GeV/n to about 200 TeV/n (DAMPE Collaboration \cite{Ele}). The bottom sub-detector of DAMPE is NUD, designed to enhance e/p separation power by detecting neutrons generated by hadronic shower in the BGO. The PSD is composed by 82 plastic scintillator bars arranged into two layers, both layers have 39 bars with a size of $ 824 \times 28 \times 10$ mm$^3 $ and two edge bars with a size of $ 824 \times 25 \times 10$ mm$^3$. Two layers are orthogonal to each other. In order to avoid dead regions, neighbor bars in each layer are staggered by 10 mm as shown in Fig. \ref{fig1}. The other details of the PSD detector structure can be found in (Yu. et al. \cite{yu}). The mean energy deposition (or most probable value of the energy deposition) of a high-energy charged particle in a PSD bar is proportional to its path length (hereafter PL) in the volume of a PSD bar. Therefore, in order to obtain an accurate energy deposition of a charged particle in the PSD, it is important to carry out a detector alignment of all PSD bars. If a PSD bar is not located in its designed position, the measured energy spectrum of minimum ionization particles (MIPs) features a distorted structure due to an incorrect calculation of the PL, especially for the particles that pass only through a corner (corner-passing events). Based on this fact we develop a method to align all PSD bars using the correlation between measured energy spectra and PL. In the paper, we will firstly introduce a method of PSD alignment in Section \ref{sec2}. The validation of this method and the charge resolution improvement are presented in Section \ref{sec3}. The results and the possible application of the alignment method are presented in Section \ref{sec4}. | 18 | 8 | 1808.05720 |
|
1808 | 1808.10838_arXiv.txt | {Type II Cepheids (T2Cs) are radially pulsating variables that trace old stellar populations and provide distance estimates through their period-luminosity (PL) relation.} {We trace the structure of old stellar population in the Galactic bulge using new distance estimates and kinematic properties of T2Cs.} {We present new near-infrared photometry of T2Cs in the bulge from the VISTA Variables in the V{\'i}a L{\'a}ctea survey (VVV). We provide the largest sample (894 stars) of T2Cs with $JHK_s$ observations that have % accurate periods from the Optical Gravitational Lensing Experiment (OGLE) catalog. Our analysis makes use of the $K_s$-band time-series observations to estimate mean magnitudes and individual distances by means of the PL relation. To constrain the kinematic properties of our targets, we complement our analysis with proper motions based on both the VVV and {\it Gaia} Data Release 2.} {We derive an empirical $K_s$-band PL relation that depends on Galactic longitude and latitude: $K_{s0} = (10.66\pm0.02) - (2.21\pm0.03)\cdot(\log{P}-1.2) - (0.020\pm0.003)\cdot l + (0.050\pm0.008)\cdot |b|$ mag; individual extinction corrections are based on a 3D reddening map. Our targets display a centrally concentrated distribution, with solid evidence of ellipsoidal symmetry---similar to the RR Lyr{\ae} ellipsoid---and a few halo outliers up to $\gtrsim$100 kpc. We obtain a distance from the Galactic center of $R_0$=8.46$\pm$0.03(stat.)$\pm$0.11(syst.) kpc. We also find evidence that the bulge T2Cs belong to a kinematically hot population, as the tangential velocity components ($\sigma v_{l*}$=104.2$\pm$3.0 km/s and $\sigma v_{b}$=96.8$\pm$5.5 km/s) agree within 1.2$\sigma$. Moreover, the difference between absolute and relative proper motion is in good agreement with the proper motion of Sgr A* from VLBA measures.} {We conclude that bulge T2Cs display an ellipsoidal spatial distribution and have kinematics similar to RR Lyr{\ae} stars, which are other tracers of the old, low-mass stellar population. T2Cs also provide an estimate of $R_0$ that agrees excellently well with the literature, taking account of the reddening law.} | \label{section:intro} In the Galactic bulge, the Red Clump (RC) stars, which are core-helium burning low-mass stars with ages from intermediate (1$\leq$age<10 Gyr) to old (age$\geq$10 Gyr) and average to high metallicities \citep[{[Fe/H]}$\gtsim$--1.5 dex,][]{cole1998,hill2011}, are used extensively to study their kinematics, chemical abundances, and spatial distributions \citep{mcwilliamzoccali2010,saito2011,gonzalez2015, zoccali2017}. High-resolution spectroscopic studies of RCs suggest that only stars that are more metal-rich than [Fe/H]$\sim$--0.5 dex trace an X-shaped structure that appears to be in the form of a peanut-boxy bulge \citep{ness2013,zoccali2016}. Recently, mid-infrared data obtained with the Wide-field Infrared Survey Explorer (WISE) also showed a clear large-scale X structure \citep{nesslang2016}. In contrast, stars that are more metal-poor than [Fe/H]$\sim$--0.5 dex not only display a centrally concentrated axisymmetric spatial distribution, but also reveal different kinematics \citep{ness2013,valenti2016,zoccali2017}. However, the spatial distribution of the old and more metal-poor population of stars in the bulge, which is traced by RR Lyr{\ae} (RRLs), is still under discussion. \citet{pietrukowicz2015} found an ellipsoidal distribution (axis ratios of 1 : 0.49$\pm$0.02 : 0.39$\pm$0.02), elongated along the same direction as the bar traced by the metal-rich red giants, while \citet{dekany2013} and \citet{kunder2016} found a spheroidal distribution. Type II Cepheids (T2Cs) are old (>10 Gyr), low-mass post-horizontal branch, asymptotic giant branch (AGB) and post-AGB stars. Like the RRLs, T2Cs trace old stellar populations, but they have longer periods (1-80 days), are brighter by 1-3 mag, and their amplitudes can be up to twice as large as those of the RRLs. In his pioneering work that led to the separation of Population I and Population II stars, \citet{baade44} showed that T2Cs are distance indicators that obey a period-luminosity (PL) relation different from that of Classical Cepheids. T2Cs have been widely used in the literature as distance indicators, both in the optical \citep{harris85,nemec1994} and in the near-infrared (NIR) bands \citep{matsunaga06,matsunaga11a,bhardwaj17b,bhardwaj17c}, although not as frequently as RRLs and Classical Cepheids. A key feature of the PL relations is that their intrinsic dispersion becomes smaller from the optical to the NIR \citep{dicriscienzo07}. This means that NIR PL relations are not only more accurate because the reddening is less severe, but they are also intrinsically more precise. Near-infrared photometry of T2Cs in the bulge obtained with the Son of ISAAC (SOFI) telescope was used by \citet{groenewegen08} to estimate the distance of the Galactic center ($R_0$=7.99$\pm$0.09 kpc) using a sample of 38 T2Cs. \citet{bhardwaj17c} matched photometry of the VISTA Variables in the V\'ia L\'actea (VVV) survey \citep{minniti2010,saito2012} with the OGLE III version of the catalog of T2Cs \citep[][335 T2Cs]{soszynski2011} and obtained individual distances. They estimated $R_0$=8.34$\pm$0.03 kpc and ruled out a barred structure. All recent estimates of $R_0$ based on other diagnostics and recent reviews agree that the official IAU value of $R_0$=8.5 kpc is overestimated ($R_0$=8.2$\pm$0.1 kpc, \citealt{blandhawthorn2016}; $R_0$=8.3$\pm$0.2$\pm$0.4 kpc, \citealt{degrijs2016}). Recently, the Optical Gravitational Lensing Experiment (OGLE) IV survey \citep{udalski2015} has generated the largest homogeneous sample of T2Cs known to date, amounting to 924 objects projected toward the Galactic bulge \citep{soszynski2017}. This is almost three times the size of the previous sample. The VVV survey has collected NIR $K_s$-band time series toward the Galactic bulge in a sky area that covers almost the entire OGLE survey area and provides an optimal framework to characterize the structure of the old population of the Galactic bulge with stellar tracers such as RRLs and T2Cs, for which the optical photometry and accurate periods are available from the OGLE survey. The increase of the sample size with respect to previous works is a unique opportunity to achieve new insight into the old stellar population in the bulge, especially for a detailed comparison with RRLs, which has always been hampered by the small sample size of T2Cs. Furthermore, we have the unprecedented opportunity to combine the T2C NIR catalog with the proper motion measurements from VVV itself \citep{contreras2017,smith2018} and {\it Gaia} DR2 \citep{gaia_alldr,gaia_dr2} to constrain the kinematic properties of the old stellar population. The paper is organized as follows: in Section~\ref{section:data} we present our photometric and astrometric databases. We analyze the light curves and derive their properties in Section~\ref{section:lcv}. Section~\ref{section:pls} is dedicated to estimating individual distances of T2Cs and their overall distribution, while in Section~\ref{section:kinematics} we discuss the kinematic properties of our targets. We discuss and summarize our results in Section~\ref{section:conclusions}. | \label{section:conclusions} We have retrieved $K_s$-band light curves from VVV aperture photometry for 894 of 924 T2Cs in the OGLE IV catalog \citep{soszynski2017}. We calculated mean magnitudes and amplitudes based on PLOESS fits \citep{persson2004,braga2018} to the light curves. For BLHs and WVs, we simultaneously estimated individual extinctions and distance moduli, based on a 3D reddening map \citep{schultheis2014} and on a PL relation. The calibration of the PL relation was based on the slope and zero-point of T2Cs in the LMC \citep{bhardwaj17b}, anchored with a late-type eclipsing binaries distance to the LMC \citep{pietrzynski2013}. We found distances ranging from 2.0 to 111.7 kpc kpc, which means that our objects are located in the bulge, in the inner and outer halo, and possibly in the thick disk. The mean individual relative uncertainty is 8.6\%, independent of distance and with a small standard deviation of 1.2\%. The distribution of the individual distances, taking various geometric and selection biases into account, provides an estimate of the distance of the Galactic center $R_0$ of 8.46$\pm$0.03(stat.)$\pm$0.11(syst.), which agrees with the recommended value of 8.3$\pm$0.2(stat.)$\pm$0.4(syst.) kpc \citep{degrijs2016}. Our estimate of $R_0$ does not agree with other estimates with similar methods \citep[$R_0 \approx 8.30$ kpc][]{dekany2013,pietrukowicz2015,bhardwaj17c}, but the difference is consistent with the different reddening law that was adopted (\citealt{alonsogarcia2017} instead of \citealt{nishiyama2009}). We provided solid evidence that the old stellar population in the bulge is ellipsoidal. First, we found a non-negligible dependence of the PL relation on the $l$ coordinate. This has been described before by \citet{groenewegen08}, but their limited sample hampered the precision of the coefficient (--0.028$\pm$0.031 mag/\deg), while ours is more precise (--0.019$\pm$0.003 mag/\deg). Second, we found that at $l\lesssim$--5\deg, the average distance is larger, while at $l\gtrsim$5\deg the average distance of T2Cs is smaller, on a map projected onto the Galactic plane. Third, which is a similar but more quantitative approach as the second point, we found that the distribution of T2Cs at positive $l$ is centered at 8.29$\pm$0.09 kpc, while that of T2Cs at negative $l$ is centered at 8.68$\pm$0.05 kpc. We also adopted proper motions from both {\it Gaia} and VVV itself to constrain the kinematic properties of T2Cs in the bulge. The analysis was restricted to only the sources with a combined statistical error smaller than 2 mas/yr. The power of the synergy between {\it Gaia} and VVV astrometric data is clear when comparing the absolute proper motions from {\it Gaia} with relative proper motions from the VVV. The mean difference (--6.41$\pm$0.02 and 0.12$\pm$0.03 mas/yr in the longitude and latitude direction, respectively) for T2Cs within 2 kpc from the Galactic center is similar within the uncertainties (0.82 mas/yr for PSF and 0.47 mas/yr for {\it Gaia}) to the VLBA estimate of the relative proper motion of Sgr A* (--6.379$\pm$0.026 and --0.202$\pm$0.019 mas/yr). This is reasonable if we assume that the T2Cs of the bulge belong to the kinematically hot, old stellar population \citep{minniti1996,kunder2016}. Another piece of evidence supporting this assumption is that the velocity dispersion in both the longitude and latitude directions agree within almost 1$\sigma$ ($\sigma v_{l*}$=104.2$\pm$3.0 km/s, $\sigma v_b$=96.8$\pm$5.5 km/s.) The difference may be due to contamination by thick-disk stars in the 2 kpc sphere around the Galactic center. It is important to note that while the distribution and kinematics of metal-rich populations in the bulge, tracing the X-shaped structure, have been studied widely, the distributions of more metal-poor populations based on different tracers remain to be investigated in detail. This work on T2Cs provides results that are consistent with RRLs. The spectroscopic follow-up of these objects in the near future will allow us to confirm the differences in their spatial distributions and kinematics to those of metal-rich populations in the Galactic bulge. | 18 | 8 | 1808.10838 |
1808 | 1808.00105_arXiv.txt | Solar flares are one of the most energetic events in the solar system, their impact on Earth at ground level and its atmosphere remains under study. The repercussions of this phenomenon in our technological infrastructure includes radio blackouts and errors in geopositional and navigation systems that are considered natural hazards in ever more countries. Occurrence frequency and intensity of the most energetic solar flares are been taken into account in national programs for civil protection in order to reduce the risk and increase the resilience from Space Weather events. In this work we use the statistical theory of extreme values as well as other statistical methods in order to asses the magnitudes of the most extreme solar flare events expected to occur in a given period of time. We found that the data set under study presents a dual tail behaviour. Our results show that on average we can expect one solar flare greater than X23 each 25 years, that is to say, one such event each two solar cycles. | \label{S-Introduction} One of the most energetic stellar activities is produced by flares \citep{Katsova2018}. In the case of the Sun the flares occur over the solar surface, mainly in active regions \citep{2017ApJ...834...56T}. In these areas, the magnetic structures can produce very large amounts of energy that are released in the magnetic reconnection process \citep{2010ARA&A..48..241B}. This magnetic reconnection has the ability to perform a transformation of energy between magnetic and kinetic energy. The full process that involves solar reconnection and its changes in energy distribution of the system's surroundings is known as a solar flare \citep{1960MNRAS.120...89G}. When the charged particles are accelerated by a flare, they can produce an amount of electromagnetic waves at all wavelengths of the spectrum (in the most energetic, even gamma rays \citep{2014ApJ...787...15A}). Their intensity and dynamic evolution depends on their interactions with their environments \citep{1993ApJ...411..362C}. The emission at soft X-ray wavelengths produced during a solar flare has been used as a measure of the intensity of the whole complex solar flare process \citep{2017SoPh..292..144C}. The X-ray flare classification shows indirectly the amount of energy released during a solar flare event \citep{goes}. The statistical study of occurrences of solar flares using X-ray classification allows us to estimate the magnitudes of the most extreme intensities a solar flare can reach in a given period of time \citep{Koons,Riley}. Knowing the magnitude of the most extreme events is of great interest in the context of Space Weather studies due to the fact that the protocols of civil protection related with this natural phenomenon need information regarding the worst expected scenarios \citep{SWE:SWE20303}. The WSPC/NOAA has regular records of solar flares since 1975 until the present day. From 1975 up to 1978 the records used H$\alpha$ observations, but after 1978 a set of GOES X-ray detectors were used in order to record the solar flare events. SWPC has made efforts in order to maintain the X-ray solar flare classification constant in between changes of detectors installed in the space crafts and they released an homogeneous public list of solar flare records from 1 September 1975 up to 28 June 2017\footnote{The data can be found in \url{https://www.ngdc.noaa.gov/stp/space-weather/solar-data/solar-features/solar-flares/x-rays/goes/xrs/}}. In this work we use the methods of the statistical theory of extreme values (see \cite{Castillo,Coles,deHaan}) as well as \emph{ad hock} statistical methods in order to asses the magnitudes of the most extreme solar flare events that can be expected to occur in the following years. For this purpose we use the NOAA X-ray classification data set from 5 November 1975 to 9 October 2017. In Section \ref{statistic} we set up the statistical problem. Section \ref{rawdata} describes the raw data. In Section \ref{tail} we study the tail behavior of the distribution of the solar flare intensities. We do this by fitting five different models for the distribution under study. | The aim of this work at the beginning was to determine the tail behavior of the probability distribution of the intensity $X$ of solar flares, as depicted by the parameter $\alpha$ in Equation~\ref{ec1}. However, we latter found that the data set under study presents a dual tail behavior. Indeed, we found that the very most extreme values in our data set are less intense than what one would expect on account of the behavior of the rest of the data points. In terms of Figure~\ref{fig6}, the points on the right of the dashed vertical line behave differently from what one would expect from the behavior of the points on the left of the dashed vertical line. This dual tail behavior was confirmed by a hypothesis test in Section~\ref{s3.4}. Now, it is natural to pose the question of whether the attenuation of the intensities is a natural phenomenon pertaining to the physics of the solar activity or whether it is due to a threshold of the measuring instruments \citep{goes}. It should be noted that the dual behavior, or deviation from a pure power law (as in Equation~\ref{ec1}), of the tail of the distribution of the solar flare intensities has already been noticed in the literature (see \cite{Wheatland} and \cite{Aschwanden}). In particular \cite{Wheatland} applied Bayesian methods in order to study flare intensities in a single active region, and found evidence for departure from a pure power law behavior. Our findings are in agreement with Wheatland's. The deviation from a pure power law behavior in empirical data has been observed, not only in studies pertaining solar activity, but in other fields of scientific inquiry as well (see \cite{Chinnery}, for example). As evidence mounts in favour to the claim that pure power laws are not completely suitable for describing extreme events in natural phenomena, alternative models appear in the literature (see \cite{Chakrabarty,Grabchak}) as modelling resources whose usefulness ought to be explored in the future. Our results show that on average we can expect one solar flare greater than X23 each 25 years, that is to say, one such event each two solar cycles. The threshold of \~X20 is important because the energy level of saturation of the GOES spacecraft is X17.1 \citep{goes}. After the flux of solar flares reaches X20 the instruments do not provide any more data. The peak of energy of the solar flare in progress remains unknown until forensic techniques allow us to determine the intensity of the flare (\cite{2005A&A...433.1133K}). However, it is mandatory to know the energy released in order to start the protocols of civil protection for Space Weather. Fortunately, events like Carrington's \citep{2013JSWSC...3A..31C} seem not to be very frequent (X40). Our results show that this kind of extreme events have a return period between 131 years (Log-Logistic/Pareto) and 238 years (Log-logistic/Tempered Pareto). \begin{figure}[b] \begin{center} \includegraphics[width=1.0\textwidth]{f1.eps} \caption{The raw data. The \textit{dashed gray line} represent the density of event occurrence.} \label{fig1} \end{center} \end{figure} \begin{figure} \begin{center} \includegraphics[width=12cm]{f5.eps} \caption{The log-log plot for the 1\,000 most extreme points in the data set. This figure shows that a change in the slope of the fitted \textit{straight line segments} is taking place as we move from \textit{left to right}.} \label{fig5} \end{center} \end{figure} \begin{figure} \begin{center} \includegraphics[width=10cm]{f8.eps} \caption{The distribution of the slope of the fitted regression line to the set of the 28 top extremes in $10^4$ synthetic samples from a log-logistic distribution with $\alpha=1.27$. From this histogram, it follows that the event $\{ \hat{\alpha}>3 \}$ has a small probability of occurrence. Here $\hat{\alpha}$ is computed from the 28 most extreme data points in a sample.} \label{fig8} \end{center} \end{figure} \begin{figure} \begin{center} \includegraphics[width=12cm]{f6.eps} \caption{The log-log plot for the full data set and of $\overline{G}_1(x)$. To the \textit{left} of the \textit{dashed vertical line}, the log-logistic distribution holds. To the \textit{right} of the \textit{dashed line} the Pareto distribution holds.} \label{fig6} \end{center} \end{figure} \begin{figure} \begin{center} \includegraphics[width=12cm]{f7.eps} \caption{The log-log plot for the full data set and of $\overline{G}_2(x)$. To the \textit{left} of the \textit{dashed vertical line}, the log-logistic distribution holds. To the \textit{right} of the \textit{dashed line} the tempered Pareto distribution holds.} \label{fig7} \end{center} \end{figure} \begin{figure} % \centerline{\hspace*{0.015\textwidth} \includegraphics[width=0.515\textwidth,clip=]{f11a.eps} \hspace*{-0.03\textwidth} \includegraphics[width=0.515\textwidth,clip=]{f11b.eps} } \vspace{-0.35\textwidth} % \centerline{\Large \bf % \hspace{0.0 \textwidth} \color{white}{(a)} \hspace{0.415\textwidth} \color{white}{(b)} \hfill} \vspace{0.31\textwidth} % \caption{Large curvature as misbehavior in the threshold approach. } \label{fig11} \end{figure} \begin{figure} \begin{center} \includegraphics[width=11.9cm]{f14.eps} \caption{This is the log-log plot of the 150 most extreme points in the data set. The \textit{blue dashed line} stands for the $\overline{G}_1$ complementary distribution function considered in Subsection~\ref{s3.4}. The right hand part of the tails of the two \textit{dashed green lines} have slopes $-4.33$ and $-2.29$. These slopes correspond to the 95\% confidence interval bounds reported for the parameter $\alpha$ in the second part of the log-logistic/Pareto model in Table~\ref{t3}.} \label{fig12} \end{center} \end{figure} \begin{table} \caption{The distribution of intensities in $1\times10^{-4}$ W m$^{-2}$ of the solar flares.} \label{t1} \begin{tabular}{|c | c | r | c | c | r |} \hline No. & Range & Data & No. & Range & Data\\ \hline\hline 1 & $< 0.05$ (C5) & 62\,364 & 11 & 8 (X8) to 10 (X10) & 11\\ \hline 2 & 0.05 (C5) to 0.1 (M1) & 8\,162 & 12 & 10 (X10) to 12 (X12)& 5\\ \hline 3 & 0.1 (M1) to 0.2 (M2) & 3\,635 & 13 & 12 (X12) to 14 (X14)& 9\\ \hline 4 & 0.2 (M2) to 0.3 (M3)& 1\,221 & 14 & 14 (X14) to 16 (X16)& 3\\ \hline 5 & 0.3 (M3) to 0.5 (M5) & 881 & 15 & 16 (X16) to 18 (X18)& 2\\ \hline 6 & 0.5 (M5) to 1 (X1) & 615 & 16 & 18 (X18) to 20 (X20)& 0\\ \hline 7 & 1 (X1) to 2 (X2) & 263 & 17 & 20 (X20) to 22 (X22)& 2\\ \hline 8 & 2 (X2) to 4 (X4) & 151 & 18 & 22 (X22) to 24 (X24)& 0\\ \hline 9 & 4 (X4) to 6 (X6) & 35 & 19 & 24 (X24) to 26 (X26)& 0\\ \hline 10 & 6 (X6) to 8 (X8) & 10 & 20 & 26 (X26) to 28 (X28)& 1\\ \hline \end{tabular} \end{table} \begin{table} \caption{The list of the most intense events in the data set (magnitude in $1\times10^{-4}$ W m$^{-2}$).} \label{t2} \begin{tabular}{|c|c|c|c|c|c|} \hline No. & Magnitude & Date & No. & Magnitude & Date\\ \hline\hline 1 & 28 (X28) & 4 Nov 2003 & 16 & 12 (X12) & 6 Jun 1982\\ \hline 2 & 20 (X20) & 2 Apr 2001 & 17 & 12 (X12) & 1 Jun 1982\\ \hline 3 & 20 (X20) & 16 Aug 1989 & 18 & 10 (X10) & 20 May 1984\\ \hline 4 & 17 (X17) & 28 Oct 2003 & 19 & 10 (X10) & 17 Dec 1982\\ \hline 5 & 17 (X17) & 7 Sept 2005 & 20 & 10 (X10) & 20 Oct 2003\\ \hline 6 & 15 (X15) & 6 Mar 1989 & 21 & 10 (X10) & 9 Jun 1991\\ \hline 7 & 15 (X15) & 11 Jul 1978 & 22 & 10 (X10) & 25 Jan 1991\\ \hline 8 & 14 (X14) & 15 Apr 2001 & 23 & 10 (X10) & 29 Sept 1989\\ \hline 9 & 13 (X13) & 19 Oct 1989 & 24 & 10 (X10) & 9 Jul 1982\\ \hline 10 & 13 (X13) & 25 Apr 1984& 25 & 9 (X9) & 6 Nov 1997\\ \hline 11 & 13 (X13) & 15 Dec 1982& 26 & 9 (X9) & 22 Mar 1991\\ \hline 12 & 12 (X12) & 15 Jun 1991& 27 & 9 (X9) & 6 Sept 2017\\ \hline 13 & 12 (X12) & 11 Jun 1991& 28 & 9 (X9) & 24 May 1990\\ \hline 14 & 12 (X12) & 6 Jun 1991 & 29 & 9 (X9) & 5 Dec 2006\\ \hline 15 & 12 (X12) & 1 Jun 1991 & 30 & 9 (X9) & 6 Nov 1980\\ \hline \end{tabular} \end{table} \begin{table} \caption{The estimated parameter $\alpha$ in various models. In the column ``Fit" we report the Kolmogorov-Smirnov distance from the empirical distribution to the distribution of the model in turn. For the Pearson chi squared statistic, the null hypothesis that the data fits the model is not rejected for those entries in column ``$\alpha$" marked with an asterisk.} \label{t3} \begin{tabular}{|r|c|c|r|c|} \hline Model & $\alpha$ & Conf. int. & Sample & Fit\\ \hline\hline Log-logistic & $1.35^*$ & (1.25, 1.45) & 500 & 0.032\\ \hline Block & $2.41^*$ & (1.54, 5.57)& 59 & 0.079\\ \hline Threshold & $3.34^*$ & (1.83, 8.77) & 153 & 0.048\\ \hline Log-logistic/Pareto 1 & 1.24 & (1.19, 1.30) & 77\,342 & 0.055 \\ \hline Log-logistic/Pareto 2 & $3.04^*$ & (2.29, 4.33) & 28 & $6\times10^{-5}$\\ \hline Log-logistic/tempered Pareto 1 & 1.36 & (1.34, 1.38) & 68\,724 & 0.029\\ \hline Log-logistic/tempered Pareto 2 & 1.11 & (1.07, 1.14) & 8\,646 & 0.01\\ \hline \end{tabular} \end{table} \begin{table} \caption{Return levels $\times 10^{-4}$ (W m$^{-2}$) for the block approach.} \label{t4} \begin{tabular}{|c|c|c|c|} \hline Period & Return level & Lower bound & Upper bound\\\hline \hline 10 & 16.80 & 10.26 & 23.34\\\hline 25 & 26.49 & 12.39 & 40.58\\\hline 50 & 36.61 & 12.69 & 60.53\\\hline 100 & 50.07 & 12.69 & 89.24\\\hline 150 & 59.92 & 8.32 & 111.5\\\hline \end{tabular} \end{table} \begin{table} \caption{Return levels $\times 10^{-4}$ (W m$^{-2}$) for the threshold approach.} \label{t5} \begin{tabular}{|c|c|c|c|} \hline Period & Return level & Lower bound & Upper bound\\\hline \hline 10 & 19.11 & 12.10 & 26.12\\\hline 25 & 27.55 & 13.75 & 41.36\\\hline 50 & 35.67 & 14.00 & 57.34\\\hline 100 & 45.66 & 12.89 & 78.44\\\hline 150 & 52.55 & 11.35 & 93.75\\\hline \end{tabular} \end{table} \begin{table} \caption{Return levels $\times 10^{-4}$ (W m$^{-2}$) for the log-logistic/Pareto model.} \label{t6} \begin{tabular}{|c|c|c|c|} \hline Period & Return level & Lower bound & Upper bound \\\hline \hline 10 & 17.16 & 13.98 & 21.55 \\\hline 25 & 23.19 & 17.35 & 32.02 \\\hline 50 & 29.12 & 20.38 & 43.40 \\\hline 100 & 36.57 & 23.95 & 58.65 \\\hline 150 & 41.79 & 26.33 & 70.01 \\\hline \end{tabular} \end{table} \begin{table} \caption{Return levels $\times 10^{-4}$ (W m$^{-2}$) for the log-logistic/tempered Pareto model.} \label{t7} \begin{tabular}{|c|c|c|c|} \hline Period & Return level & Lower bound & Upper bound\\\hline \hline 10 & 17.19 & 14.04 & 20.37 \\\hline 25 & 23.17 & 18.51 & 27.94 \\\hline 50 & 28.07 & 22.15 & 34.27 \\\hline 100 & 33.24 & 25.95 & 41.00 \\\hline 150 & 36.36 & 28.21 & 45.04 \\\hline \end{tabular} \end{table} \renewcommand{\arraystretch}{1} \begin{table} \caption{Expected frequency of event occurrence computed by the log-logistic/Pareto model (LL/P) and by the log-logistic/tempered Pareto model (LL/TP).} \label{t8} \begin{tabular}{|c|c|c|c|c|c|} \hline Intensity (W/$\hbox{m}^2$) & NOAA & LL/P& Conf. Int. & LL/TP & Conf. Int.\\\hline \hline $10^{-5}$ (M1) & 2000 & 1894 & (1849, 1940) & 1745 & (1677, 1815)\\\hline $5\times10^{-5}$ (M5) & 350 & 279 & (254, 305) & 283 & (264, 303) \\\hline $10^{-4}$ (X1) & 175 & 119 & (104, 135) & 125 & (115, 136) \\\hline $10^{-3}$ (X10) & 8 & 7.87 & (6.04, 10.04) & 4.00 & (2.77, 5.24) \\\hline $2\times10^{-3}$ (X20) & $<1$ & 4.95 & (1.56, 9.27) & 0.70 & (0.33, 1.15)\\\hline \end{tabular} \end{table} \begin{acks} The authors thank projects for Catedras Conacyt (Conacyt Fellow), Repositorios Institucionales (268273) and Ciencia Basica (254497). \end{acks} \clearpage \newpage | 18 | 8 | 1808.00105 |
1808 | 1808.06746_arXiv.txt | Cubesats and similarly scaled nano-satellites present significant opportunities for hosting both scientific and commercial payloads for Earth sensing and astronomical observations, in particular in the area of rapid-response observations to external triggers. However, one limiting factor to full exploitation of the CubeSat potential in this area lies in the traditional approach of ground-spacecraft communications, which is based on infrequent contact via a limited network of ground stations. An alternative is to leverage existing commercial machine-to-machine orbital networks to transmit the triggers in near-real-time. Here, we present an analysis framework for calculating the likelihood of a time to first contact and the length of contact under minimum guaranteed conditions for these networks. The analysis is then extended to likely operational conditions in orbit, and the results of a comparative trade study of a number of orbital networks are presented, with an emphasis on the applicability to the SkyHopper Space Telescope CubeSat, a nanosatellite astronomical observatory currently undergoing preliminary design. It was found that near-real-time telecommands could be transmitted to SkyHopper within 10min with a likelihood of 62\% using the Globalstar network, or a likelihood of 74\% using the Iridium network, under predicted nominal operational conditions in orbit. Future networks currently under development could improve these figures to reach greater than 98\% coverage with a one second latency. | \vspace{-0.4cm} Traditional Telemetry and TeleCommand (TTC) schemes use either single ground stations, a network of ground stations, or government satellite relay networks such as NASA's Tracking and Data Relay Satellites (TDRS). However, for small missions built around NanoSatellites, the cost of utilising either a large network of ground stations, or TDRS is prohibitively expensive, while a single ground station can only provide regular but infrequent communication windows. Depending on orbit, these windows may occur only twice a day. For missions which require rapid communication for dissemination of observing targets these ground network contacts are insufficient. The opportunity then, presents itself to leverage existing commercial satellite relay networks to achieve the timeliness requirements of a mission. While both the GlobalStar and Iridium networks are currently being investigated on-board small-sats \cite{GEARRSReentry} \cite{TechEdSat}, no comprehensive study of all available networks has been published. To this end, five satellite phone networks are investigated in this paper to assess the potential for near-real-time trigger transmission: GlobalStar, Inmarsat, Iridium, Orbcomm and Thuraya. In addition a ground network of two polar stations, modelled on NASAs Near Earth Network (NEN), is included. The coverage at varying orbital altitudes, time to first contact and contact time, as well as the operational cost of using the networks are compared for each network. As a first application of the method, we analyse the network performance for communication with the SkyHopper Space Telescope Cubesat\footnote{http://skyhopper.space}. \vspace{-0.4cm} | \vspace{-0.4cm} While nano-satellites offer significant potential for Earth and astronomical observation missions, certain mission types require near-real-time telecommand. In this paper, we presented a framework for analysing the probability distribution of contact time lag between a network and a satellite, depending on the network architecture and satellite orbit. A first application to the SkyHopper mission shows that among commercially available solutions, the Iridium and Globalstar networks provide the highest likelihood of delays in commanding, while the geostationary networks, Inmarsat and Thuraya offer extremely long contact windows. For typical nanosatellite orbits, commands can be guaranteed to be delivered to the spacecraft within 1 hour 70\% of the time, while under predicted nominal conditions, commands can be delivered within just a few minutes 70\% of the time, using the Iridium network. For the SkyHopper mission, either of the Iridium or Globalstar networks is estimated to provide delivery of commands with less than 10 min lag at greater than 62\% confidence. This performance can be improved by a combination of the two networks, or by the use of future solutions for near-real time communications such as the Audacy network which should be available by early 2020 \cite{Audacy}. \vspace{-0.6cm} | 18 | 8 | 1808.06746 |
1808 | 1808.01399_arXiv.txt | We analyse interferometric data obtained for Regulus with AMBER (Astronomical Multi- BEam combineR) at high spectral resolution ($\lambda/\delta\lambda \approx 12000$) across the Br$\gamma$ spectral line. The study of the photocentre displacement allows us to constrain a large number of stellar parameters -- equatorial radius $R_{\rm eq}$, equatorial velocity $V_{\rm eq}$, inclination $i$, rotation-axis position angle $PA_{\rm rot}$, and flattening -- with an estimation of gravity-darkening coefficient $\beta$ using previously published theoretical results. We use the Simulation Code of Interferometric-observations for ROtators and CirCumstellar Objects (SCIROCCO), a semi-analytical algorithm dedicated to fast rotators. We chose Regulus because it is a very well-known edge-on star, for which an alternative approach is needed to check the previously published results. Our analysis showed that a significant degeneracy of solution is present.\\ By confronting the results obtained by differential interferometry with those obtained by conventional long-base interferometry, we obtain similar results (within the uncertainties), thereby validating our approach, where $V_{eq}$ and $i$ are found separately. From the photocentre displacement, we can independently deduce $PA_{rot}$. We use two minimization methods to restrict observed stellar parameters via a fast rotator model: a non-stochastic method ($\chi^2$ fit) and a stochastic one (Markov Chain Monte Carlo method), in order to check whether the correct global minimum is achieved particularly with respect to the degeneracies of the gravity darkening parameter $\beta$, where we demonstrate, using a quantitative analysis of parameters, that the estimate of $\beta$ is easier for stars with an inclination angle of around $45^\circ$.\\ | \label{introduction} \subsection{Optical interferometry of rapid rotators} \label{Opt-Interf-rapid_rot} Stellar rotation was measured for the first time by interferometry, from the photocentre displacements by \cite{sl94}, on the slow rotator Aldebaran, which was observed in 1988 at OHP (Observatoire de Haute-Provence) through the 152cm telescope by the Speckle Differential Interferometry method. Results obtained by interferometry on the fast rotators were summarized by \cite{2011SerAJ.183....1J} \& \cite{2012A&ARv..20...51V}. The extreme stellar flattening induced by the rotation was measured by interferometry by \cite{2003A&A...407L..47D} on Achernar ($R_{eq}/R_{pol}=1.56\pm0.05$), using VLTI/VINCI (Very Large Telescope Interferometer/VLT INterferometer Commissioning Instrument) with a dense ($u,v$) coverage. The first image reconstruction of the surface of a fast rotator, showing the gravity darkening effect \citep{1924MNRAS..84..665V, 1924MNRAS..84..684V} was on Altair \citep{2007Sci...317..342M} from CHARA (Center for High Angular Resolution Astronomy) observations, inferring several fundamental parameters: inclination; position angle; effective temperature; and polar and equatorial radii. Inspired by the early works of \cite{1975ApJ...196L..71L}, \cite{1982AcOpt..29..361B} proposed the Differentiel Speckle Interferometry technique, using the chromatic displacement of the speckle photocentre given by the first-order term of the phase of the spatial Fourier transform of the sky brightness according to the MacLaurin series \citep{2001A&A...377..721J}. This technique has been extended to a wider range of wavelengths and applied to long-baseline interferometry by \cite{1988ESOC...29..235P, 1989dli..conf..249P} who established the fundamentals of the Differential Interferometry (DI) technique. This allowed, for the first time, to measure simultaneously the angular separation and the radial velocity difference of the two stellar components of the binary Capella \citep{1992ASPC...32..477P} using the photocentre as a function of the wavelength. The combination of high spatial and high spectral resolution allows us to measure physical properties of fast rotators beyond the diffraction limit, as shown by \cite{2012A&A...545A.130D} and \cite{2014A&A...569A..45H} who used the differential phases from AMBER/VLTI (Astronomical Multi-BEam combineR). Indeed, AMBER \citep{2007A&A...464....1P} is a spectro-interferometric instrument specifically designed to go well beyond the resolution limit \citetext{e.g., \citealt{2007A&A...464...59M,2009A&A...498L..41L}}. Optical interferometry provides several types of measures, as the absolute visibility, differential visibility \& closure phase \citep{2007A&A...464....1P}, but in this paper we focus only on differential phase and vectorial photocentre displacement. \subsection{The fast rotator Regulus } \label{Reg} $\alpha$ Leo A (HR 3982, HD 87901) one of the brightest stars of the sky, is a binary system which brighter primary component is referred as Regulus throughout this paper. Regulus is an edge-on and flattened nearby star which is in rapid rotation. In the following we summarize the spectrophotometric and the interferometric information of our target separately. \subsubsection{Information from spectroscopy and photometry} \label{spec-phot_info} With magnitude $V=1.40$ \citep{2009ApJ...694.1085V}, Regulus has been identified as a fast rotator by \cite{1954ApJ...119..146S}, who determined by spectroscopy its high rotationnal velocity $\vsini=352\pm7.5\kms$ i.e. $96\%$ of its critical velocity. $\alpha$ Leo is a multiple stellar system composed of at least two binaries. The A component of the system ($\alpha$ Leo A, HD 87901) has been recently discovered to be a spectroscopic binary. The brighter companion (Regulus) was classified as a main sequence B7V star by \cite{1953ApJ...117..313J}, and more recently as a sub-geant B8IV star by \cite{2003AJ....126.2048G} of mass $\sim$ 4 $\Msun$ \cite[][and references therein]{2011ApJ...732...68C}. \cite{2008ApJ...682L.117G} argue that the fainter companion of $\alpha$ Leo A is probably a white dwarf or a M4 V star of mass $\sim$ 0.3 $\Msun$ and an orbital period of 40.11 days, and that the magnitude difference of the fainter component with respect to Regulus in the K band is close to $\Delta m_K\approx10$ ($\sim 10^{-4}$ of flux ratio) and 6 ($\sim 4\times 10^{-3}$ of flux ratio), for the cases of a white dwarf and an M4 V star companion, respectively. Thus, the flux of the companion has no influence on our analysis (Br$_\gamma$ and adjacent continuum), as well as on the interferometry presented by \cite{2005ApJ...628..439M} and \cite{2011ApJ...732...68C}. For this reason only an extraordinary activity of the fainter star, reflected as a several order of magnitude enhancement of the Br$_\gamma$ emission, could eventually affect our results. $\alpha$ Leo A has a companion which is in fact a system of two other components (B and C) which together form a binary system \citep{2005ApJ...628..439M}. The B component ($\alpha$ Leo B; HD 87884) is an $\sim$ 0.8 $\Msun$ star of spectral type K2V while the C component is a very faint M4V star with a mass of $\sim$ 0.2 $\Msun$. The Washington Double Star Catalog \citep{2001AJ....122.3466M} lists a D component, also having a common proper motion with the system and a separation of $\approx$3.6$\arcmin$ from the A component while the B-C subsystem is located $\approx$ 3$\arcmin$ from the A component. \cite{2008Ap&SS.318...51I} studied the possible correction of the Keplerian period due to the quadrupole mass moment induced by the oblateness of Regulus. Although this correction could be measured in principle, the total uncertainty in the Keplerian period \citep[0.02 days from][]{2008ApJ...682L.117G} due to the errors in the systems parameters (mostly in the velocity semiamplitude and in the mass of Regulus) is larger than the correction by about two orders of magnitude. Its distance is $d=24.3\pm0.2pc$ according to \cite{2007ASSL..350.....V} and $d=23.759\pm0.045pc$ according to \cite{2009ApJ...694.1085V}. Its mass is $M=4.15\pm0.06\Msun$ from the $Y^2$ stellar evolution model \citep{2001ApJS..136..417Y, 2003ApJS..144..259Y, 2004ApJS..155..667D}, $3.66^{+0.79}_{-0.28}\Msun$ from \cite{2011ApJ...732...68C} according to the oblateness mass method of \cite{2009ApJ...701..209Z} and $3.80\pm0.6\Msun$ from \cite{1990A&AS...85.1015M} according to the evolutionary tracks of \cite{1989A&A...210..155M}. Its age is estimated between 150 Myr \citep{2001A&A...379..162G} and 1 Gyr \citep{2009ApJ...698..666R}. Its effective temperature $T_{\rm eff}$ is $12460\pm200$K according to \cite{1990A&AS...85.1015M} and $11960\pm80$K according to \cite{2003AJ....126.2048G}. Its metallicity $[M/H]$ is $0.0$ according to \cite{2003AJ....126.2048G}. \subsubsection{Interferometric observations of Regulus} \label{Interfero-info} The first interferometric observations of this star were done with the Narrabri Intensity Interferometer by \cite{1974MNRAS.167..121H}. Because of the poor ($u,v$)-plane coverage, only information about its size could be obtained, with an equatorial angular diameter $\diameq=1.32\pm0.06 mas$. Using the CHARA array observations in the K-band, \cite{2005ApJ...628..439M} measured for the first time the inclination of its rotation axis $i=90^\circ$ $^{+0}_{-15}$, and characterized other physical parameters such as: rotation-axis position angle $PA_{rot}=265.5\pm2.8^\circ$; rotational equatorial velocity $V_{eq}=317^{+3}_{-85}\kms$; fractional rotational velocity $\frac{V_{eq}}{V_{crit}}=0.86\pm0.03$; equatorial and polar radii $R_{eq}=4.16\pm0.08\Rsun$ \& $R_{pol}=3.15\pm0.06\Rsun$; equatorial and polar effective temperatures $T_{eq}=10314\pm1000K$ \& $T_{pol}=15400\pm1400K$; mass $M$; luminosity $L$; gravity darkening coefficient $\beta$ \citep[defined as $T_{\rm eff} \propto g_{\rm eff}^{\beta}$ by][where $T_{\rm eff}$ and $g_{\rm eff}$ are local effective temperature and gravity respectively]{1924MNRAS..84..665V}; distance $d$ and interstellar extinction $A_v$. More recently, \cite{2011ApJ...732...68C} used the CHARA/MIRC (Michigan Infra-Red Combiner) instrument to produce maps of $\alpha$ Leo, in the H-band, and deduced inclination angle $i=86.3^\circ$$^{+1.0^\circ}_{-1.6^\circ}$ and gravity darkening coefficient $\beta=0.188^{+0.012}_{-0.029}$ which is consistent (within the uncertainties) with the results of \cite{2005ApJ...628..439M}. So, interferometry revealed that Regulus is an edge-on star with an inclination angle of $i\sim90^\circ$ and rotationnally flattened with an oblateness ratio (equatorial-to-polar radii minus 1; $R_{eq}/R_{pol}-1$) reported between $0.325\pm0.036$ (angular diameter $\diameq=1.65\pm0.02mas$) and $0.307\pm0.030$ ($\diameq=1.61_{-0.02}^{0.03}mas$) \citep{2005ApJ...628..439M, 2011ApJ...732...68C}. Table.~\ref{table2} summarizes the fundamental parameters of Regulus. \begin{table} \begin{minipage}{87mm} \caption{Fundamental stellar parameters of $\alpha$ Leo found in the literature.} \label{table2} \centering \begin{threeparttable} \centering \begin{tabular}{l|c} \hline \hline Parameter & Value\\ \hline \hline Angular diameter ($\diameq$) & $1.65\pm0.02mas$ (1)\\ & $1.61_{-0.02}^{0.03}mas$ (2)\\ \hline Oblateness ratio ($R_{eq}/R_{pol}-1$) & $0.325\pm0.036$ (1)\\ & $0.307\pm0.030$ (2)\\ \hline Distance (d) & $24.3\pm0.2pc$ (3)\\ & $23.759\pm0.045pc$ (4)\\ \hline & $4.15\pm0.06\Msun$ (5)\\ Mass (M) & $3.66^{+0.79}_{-0.28}\Msun$ (2)\\ & $3.80\pm0.6\Msun$ (6)\\ \hline Age & 50 - 200 Myr (7)\\ & $\geq$ 1 Gyr (8)\\ \hline Eff. temperature ($T_{eff}$) & $12460\pm200$K (6)\\ & $11960\pm80$K (9)\\ \hline Metallicity ($[M/H]$) & $0.0$ (9)\\ \hline Rotation-axis position angle $PA_{rot}$ & $265.5\pm2.8 ^\circ$ (10) \\ & $258^{+2}_{-1}$$^\circ$ (11) \\ \hline Rotation-axis inclination angle $i$ & $90^{+0}_{-15}$$^\circ$ (10) \\ & $85.3^{+1}_{-1.6}$$^\circ$ (11) \\ \hline \hline \end{tabular} \begin{tablenotes} \footnotesize $(1)$ \citet{2005ApJ...628..439M} $(2)$ \citet{2011ApJ...732...68C}\\ $(3)$ \citet{2007ASSL..350.....V} $(4)$ \citet{2009ApJ...694.1085V}\\ $(5)$ \citet{2004ApJS..155..667D} $(6)$ \citet{1990A&AS...85.1015M}\\ $(7)$ \citet{2001A&A...379..162G} $(8)$ \citet{2009ApJ...698..666R}\\ $(9)$ \citet{2003AJ....126.2048G} $(10)$ \citet{2005ApJ...628..439M}\\ $(11)$ \citet{2011ApJ...732...68C} \end{tablenotes} \end{threeparttable} \end{minipage} \end{table} \subsection{Structure of the article} \label{Struct_paper} In this paper, we describe the differential interferometry with high spectral resolution observations ($R\simeq12000$) in the K band of the rapid rotator Regulus and we focuss on the parameters that can extracted from the photocentre displacement, eventually in combination with the broadened spectral line profile. We compare this in detail with the parameters obtained from a broad band interferometric images and we discuss the values that are inferred or improved by the comparison and then the combination of both techniques. The present paper is organized as follows: \begin{itemize} \item In Sec.~\ref{obsdatared}, we present the observations and the data reduction of Regulus. \item In Sec.~\ref{geom_params}, we study the photocentre displacement of our target, where we deduce $PA_{\rm rot}$ directly from the observed photocentres displacements. \item In Sec.~\ref{Modeling}, we present the model that was used in order to interpret our measurements and discuss the constraints that they give on the gravity darkening parameter $\beta$ of Regulus. \item In Sec.~\ref{Fitting}, we fit the fundamental parameters of Regulus, using a non-stochastic method ($\chi^2$); and a stochastic one (Markov Chain Monte Carlo method). \item In Sec.~\ref{discus}, we summarize the computed accuracy limits that we could achieve with the quality of our data and we discuss the probability spaces of the couple $(\beta,i)$ of Regulus. \item In Sec.~\ref{conclusions}, we analyse the results and open the discussion to a broader study of fast-rotating stars observed with VLTI-AMBER by DI. \end{itemize} | \label{conclusions} We have presented the differential interferometry data obtained on the rapid rotator Regulus with the high spectral resolution mode of the VLTI instrument AMBER. We have seen that, for K band observations with the VLTI baselines, this target much smaller than the angular resolution $(\lambda/B)$ is not enough resolved for the Rayleigh criteria and for imaging, as all closure phases are equal to zero. It is also not enough resolved for oblateness estimations from absolute visibility measurements, because the visibility is equal to 1 within the error bars for the baselines close to the polar direction and we can only give an upper limit for the polar diameter that is smaller than the equatorial diameter. We have therefore concentrated on the interpretation and model fitting of the differential phases that, on this source much smaller than the standard resolution limit $\lambda/B$, can all be reduced to the vectorial displacement of the photocentre $\vec{\epsilon}(\lambda)$ in the spectral channel $\lambda$ with respect to the photocentre of the target in the continuum. Our data, corresponding to $\sim$30 mn ($\sim$25 mn for the night 2014/03/10 \& $\sim$40 mn for the night 2014/03/12) of open shutter observations on Regulus, yields a typical error per spectral measure (half spectral channel) of 30 $\mu as$ both for the $\alpha$ and $\delta$ components of the vector $\vec{\epsilon}(\lambda)$ and for its polar and equatorial components obtained after a direct computation of the rotation axis position angle. We have basically confirmed the previous interferometric and spectroscopic determinations of the fundamental parameters of Regulus, with our quite different data set and different constraints on the physical parameters. Our $350\pm22 \kms$ velocity measurement is compatible within errors with this used by \cite{2011ApJ...732...68C}, but significantly higher than earlier estimates. This leads to a rotation 88\% of the critical velocity. Our $PA_{\rm rot}$ measurement of $251\pm2^\circ$, which is only deduced from the vectorial photocentre ($\phidiff$), is little bit different of this deduced from squared visibilities, closure phase and triple amplitudes of CHARA/MIRC by \cite{2011ApJ...732...68C}. The vectorial photocentre is very sensitive to $PA_{\rm rot}$ parameter. We don't claim at all that our values of $\beta$ are conclusive, because we are not able to constrain the gravity darkening coefficient from our current data, which are relatively noisy with marginal angular resolution, and without forgetting the fact that Regulus is a edge-on star. Despite the fact that the star was marginally resolved with our observations, we were able, for the first time, to constrain independently (from the $\epsilon_{\alpha}$ \& $\epsilon_{\delta}$) several fundamental stellar parameters, as the $PA_{\rm rot}$ with low uncertainties. This method can be applied to stars which can be only marginally resolved or not angularly resolved at all, because of available baseline lenghts, and especially for rotators with inclinaison angles around of $45^\circ$ and a good SNR. | 18 | 8 | 1808.01399 |
1808 | 1808.04605_arXiv.txt | {The statistical properties of binary stars are clues for understanding their formation process. A radial velocity survey was carried on amongst nearby G-type stars and the results were published in 1991.} {The survey of radial velocity measurements was extended towards K-type stars.} {A sample of 261 K-type stars was observed with the spectrovelocimeter CORAVEL (COrrelation RAdial VELocities). Those stars with a variable radial velocity were detected on the basis of the $P(\chi^2)$ test. The orbital elements of the spectroscopic binaries were then derived.} {The statistical properties of binary stars were derived from these observations and published in 2003. We present the catalogue of the radial velocity measurements obtained with CORAVEL for all the K stars of the survey and the orbital elements derived for 34 spectroscopic systems. In addition, the catalogue contains eight G-type spectroscopic binaries that have received additional measurements since 1991 and for which the orbital elements are revised or derived for the first time. } {} | \label{introduction} The spectrovelocimeter CORAVEL \citep[COrrelation RAdial VELocities,][]{Baranne79} was installed on the Swiss 1-m telescope at the Observatory of Haute-Provence (OHP) from the late 1970s until its decommissioning in 2000. Amongst other programmes, it provided the radial-velocity (RV) measurements exploited in two statistical studies of binarity among the stars in the solar neighbourhood: the study of solar--type stars until G8, and its extension towards the K-type stars. A series of articles has been devoted to these programmes. The first \citep[][Paper~I hereafter]{DMH91} presented the radial-velocity measurements of the sample of F-G type stars; these data led to the orbital elements of several spectroscopic binaries (SBs), and to the statistical properties of solar-type binaries \citep[][DM91 hereafter]{DM91}. Later, \citet[][Paper~III hereafter]{Halbwachs03} extended the statistical investigations to the K-type binaries with periods shorter than ten years, again on the basis of CORAVEL observations. This paper presented the parameters relevant for statistics, namely the periods, the semi-amplitudes, the mass ratios, and the orbital eccentricities of the spectroscopic binaries, excluding the other orbital elements. The long period K-type binaries were eventually studied by \citet{Eggenberger04}. \defcitealias{DMH91}{Paper~I} \defcitealias{DM91}{DM91} \defcitealias{Halbwachs03}{Paper~III} The present paper completes the series by presenting the radial velocity measurements and the full set of orbital elements that gave rise to \citetalias{Halbwachs03}. It will give the orbits we have discovered all the visibility they deserve, so that they are henceforth taken into account in statistical studies, such as that of \citet{Raghavan}. Moreover, they will be available for the validation of the spectroscopic orbits derived from the Radial Velocity Spectrometer of the Gaia satellite \citep{GaiaRVS}. The CORAVEL programme is presented in Sect.~\ref{CORAVEL}, the RV catalogue is in Sect.~\ref{RV}, and the spectroscopic orbital elements are in Sect.~\ref{orbits}. Section~\ref{conclusion} is the conclusion. | \label{conclusion} We have drawn up a catalogue of 5413 RV measurements obtained with CORAVEL for 269 stars, 261 K-type dwarfs, and eight G-type dwarfs of the solar neighbourhood. These measurements were used to detect the SBs on which were based the statistical investigations of \citetalias{Halbwachs03}. We calculated the elements of 44 SB orbits, corresponding to 42 stellar systems. Twenty-one orbits, corresponding to 20 stellar systems, are the first orbits ever published for these stars. All these data will be available through the VizieR service of the Centre de Donn\'ee astronomique de Strasbourg (CDS). The SB orbits and the corresponding RV measurements will also be included in the on-line SB9 catalogue \citep[http://sb9.astro.ulb.ac.be/, ][]{SB9}. | 18 | 8 | 1808.04605 |
1808 | 1808.09331_arXiv.txt | A validation of the cosmic distance duality (CDD) relation, $\eta(z)\equiv (1+z)^2d_A(z)/d_L(z)=1$, coupling the luminosity ($d_L$) and angular-diameter ($d_A$) distances, is crucial because its violation would require exotic new physics. We present a model-independent test of the CDD, based on strong lensing and a reconstruction of the HII galaxy Hubble diagram using Gaussian Processes, to confirm the validity of the CDD at a very high level of confidence. Using parameterizations $\eta(z) = 1 + \eta_0 z$ and $\eta(z) = 1 + \eta_1 z + \eta_2 z^2$, our best-fit results are $\eta_0 = 0.0147^{+0.056}_{-0.066}$, and $\eta_1 = 0.1091^{+0.1680}_{-0.1568}$ and $\eta_2 = -0.0603^{+0.0999}_{-0.0988}$, respectively. In spite of these strong constraints, however, we also point out that the analysis of strong lensing using a simplified single isothermal sphere (SIS) model for the lens produces some irreducible scatter in the inferred CDD data. The use of an extended SIS approximation, with a power-law density structure, yields very similar results, but does not lessen the scatter due to its larger number of free parameters, which weakens the best-fit constraints. Future work with these strong lenses should therefore be based on more detailed ray-tracing calculations to determine the mass distribution more precisely. | \label{sec:intro} The cosmic distance duality (CDD) relation, based on Etherington's theorem (1933), depends on three essential assumptions: (i) that the spacetime is described by a metric theory of gravity; (ii) that photons travel along null geodesics; and (iii) that their number is conserved along the null geodesics. The CDD is commonly written in the form $\eta(z) = 1$, with the definition \begin{align} \eta(z) \equiv (1+z)^2 \frac{d_A(z)}{d_L(z)} \ , \label{equ:cddeta} \end{align} where $d_A(z)$ and $d_L(z)$ are the angular-diameter and luminosity distances, respectively. Many attempts have been made to test the validity of the CDD, using several different kinds of data, and/or assumptions. Typically, the angular-diameter distance $d_A(z)$ is measured using the angular size of galaxy clusters \citep{wei2015clusters,melia2016clusters}, while the luminosity distance $d_L(z)$ is often inferred from Type Ia SNe. For a non-exhaustive set of references, see \citet{2004PhRvD..69j1305B}; \citet{uzan2004distance}; \citet{holanda2010testing}, \citet{holanda2012probing}; \citet{khedekar2011new}; \citet{li2011cosmological}; \citet{nair2011observational}; \citet{lima2011deformed}; \citet{meng2012morphology}; \citet{ellis2013blackness}; \citet{liao2016distance}; \cite{2017arXiv171010929Y}; \citet{hu2018testing}; \citet{melia2018model}. But a principal difficulty with using SNe is that the measurement of $d_L(z)$ is model-dependent. One can easily see this from the definition of the distance modulus $\mu$, which is given as \begin{align} \mu = 5\log d_L - 5 = m_{\mathrm{max}} - M_{\mathrm{max}} \ , \end{align} in terms of the peak magnitude $m_{\mathrm{max}}$ and peak absolute magnitude $M_{\mathrm{max}}$. Every Type Ia SN has almost the same $M_{\mathrm{max}}$, so if $m_{\mathrm{max}}$ is measured, one can obtain the distance modulus $\mu$. The difficulty arises from the scatter in peak magnitudes, which depend rather strongly on the shapes and colors of the SN lightcurves \citep{guy2005salt}. To get $m_{\mathrm{max}}$, one of several fitters must be used to parameterize the light curves. For example, one of the most popular parameterizations is \citep{guy2007salt2} \begin{align} \mu_B (\alpha, \beta, M_B; z) = m_B^{\mathrm{max}} (z) - M_B + \alpha x - \beta c \ , \end{align} where $m_B^{\mathrm{max}}$ is the rest-frame peak magnitude of the $B$ band, $x$ is the stretch factor that describes the effect of the lightcurve shape on $\mu$, and $c$ is the color parameter representing the influence of intrinsic color and reddening due to dust on $\mu$. The so-called `nuisance' parameters $\alpha$, $\beta$, and $M_B$ must be optimized along with all the other parameters in the chosen cosmological model. By now, it is well known that different models are associated with different values of these nuisance variables, so there is no unique way to determine the SN distance moduli in a truly model independent way. It is therefore quite likely that some (or all) of the previously claimed CDD violations may simply be due to unaccounted for influences of the assumed cosmology on $\eta(z)$ \citep{uzan2004distance,holanda2010testing,li2011cosmological}. For a more detailed discussion, see \citet{yang2013improved,melia2009,melia2012SNe,melia2013,wei2015comparative,melia2018model}. In this paper, we steer clear of measurements that require the pre-assumption of particular cosmological models, and instead use strong lenses to measure the ratio of angular-diameter distances, and a reconstruction of the HII galaxy Hubble diagram with Gaussian Processes to obtain the luminosity distance. In the next section, we shall discuss the rationale behind these two kinds of observation, and why one may safely assume model independence in the associated data. Since no model is assumed in any of our analysis, our approach yields a clean measure of the CDD relation. We shall first briefly summarize the methodology of measuring $\eta(z)$ as a function of redshift using strong lensing and HII galaxies in \S~\ref{sec:methd}. We then describe the relevant datasets in \S~\ref{sec:datasets}, and present the results of our analysis in \S~\ref{sec:res}. We shall demonstrate that this combination of observations confirms the CDD at a very high level of confidence. Finally, we present our conclusions in \S~\ref{sec:conc}. | \label{sec:conc} A commonly used method of testing the CDD has been to compare the luminosity distance derived from Type Ia SNe with the angular-diameter distance measured using galaxy clusters. But these are no longer the only standard candles and rulers available today. Seeking alternatives is desirable because the older approach requires the pre-assumption of specific cosmological models in order to optimize the `nuisance' parameters in the distance-redshift relation. This limiting factor can lead to additional uncertainty and bias, which may explain why some earlier work with the CDD has produced conflicting results. A summary of previous inconsistencies may be found in \citet{melia2018model}. In this paper, we have chosen a new combination of standard candles (HII galaxies) and rulers (strong lenses) to test the CDD without the need to pre-assume any particular cosmology. The fact that our analysis shows consistency with the CDD at a very high level of confidence is therefore quite compelling because we have avoided introducing unknown systematics associated with particular models. We do note, however, that we have assumed spatial flatness throughout our analysis, which appears to be consistent with a broad range of cosmological measurements. Nevertheless, this is a caveat to keep in mind, should any new evidence emerge that the Universe is not spatially flat. Another important benefit of our model-independent study is that the additional flexibility in fitting the data otherwise present when cosmology-dependent parameterizations are introduced is absent from our approach. Our test is therefore straightforward and clean because any possible variations in $\eta(z)$ away from $1$ cannot be attributed to the cosmology itself. Were we to find that $\eta(z)\neq 1$, the evidence in favor of new physics would therefore have been stronger with our method than what would be found using model-dependent data. The results in this paper fully confirm another recent model-independent test of the CDD carried out by \citet{melia2018model}. In that work, the standard ruler was provided by compact quasar cores. We are therefore starting to see a consistent pattern of results in which model-independent tests all agree that the CDD is realized in nature. The principal caveat of our work is the irreducible scatter in our CDD data stemming from the use of a simplified single isothermal sphere (SIS) model for the lens. We have attempted to mitigate the impact of an imprecisely known mass distribution in the lens by also considering an extended SIE model, in which the internal structure of the lens is characterized by power laws for the mass and luminosity densities, with two adjustable indices. While this allows greater freedom in modeling the lens itself, however, the downside with such an approach is the additional degeneracy offered by the greater flexibility with the overall optimization of the parameters. Our results for both the simple SIS and the extended SIE lens models confirm the CDD all the way out to $z\sim 2.3$, with a violation no bigger than $\eta_0\sim 0.01-0.015$, in a parameterization $\eta(z)=1+\eta_0z$. But the CDD constraint is actually weaker with the extended SIE lens (i.e., $\eta_0=0.0093^{+0.1520}_{-0.0939}$ versus $\eta_0=0.0147^{+0.056}_{-0.066}$) due to the less precise confidence range of the best fit values resulting from the adjustable power-law distributions. In future work with strong lensing, it would be desirable to determine the lens mass distribution more accurately, e.g., using ray tracing, rather than simply relying on a SIS model, which appears to produce irreducible scatter in the results due to its over-simplification of the lens structure. Eventually, this improvement should yield a better measurement of the angular-diameter distance, allowing us to test the CDD with even higher precision than is available today. | 18 | 8 | 1808.09331 |
1808 | 1808.07606_arXiv.txt | Dust grains of crystalline silicate, which are rarely presented in interstellar space, were found in cometary nuclei. These crystalline silicates are thought to have formed by annealing of amorphous silicate grains or direct condensation of gaseous materials near the Sun in the solar nebula, and incorporated into cometary nuclei in the cold comet-forming region after radial transportation of grains in the solar nebula. Abundances of the crystalline silicate dust grains were therefore expected to be smaller farther from the Sun. We aim to better understand the formation mechanism of minerals incorporated into comet 17P/Holmes based on its mineral abundances. To derive the mineral composition of comet 17P/Holmes, we applied a thermal emission model for cometary dust grains to mid-infrared spectra of comet 17P/Holmes taken with the Cooled Mid-Infrared Camera and Spectrograph (COMICS) mounted on the Subaru Telescope a few days later the great outburst in October 2007. The resulting mass fraction of crystalline silicate, $f_{\rm cry}$, and an olivine-to-pyroxene abundance ratio, $f_{\rm OP}$, are $f_{\rm cry}$ = 0.31 $\pm$ 0.03 and $f_{\rm OP}$ = 1.20$^{+0.16}$/$_{-0.12}$, respectively. Based on a simple consideration of the mixing of dust grains originating in both the interstellar medium and solar nebula, the minerals of 17P/Holmes formed by non-equilibrium condensation. This result is consistent with theoretical and experimental predictions for vaporization and condensation of olivine in the solar nebula. | \label{sec:intro} The 10-$\mu$m emission feature is usually recognized in mid-infrared spectra of cometary dust coma when the comets are at around 0.1-3 au from the Sun \citep{HaBr04}, and the feature could be attributed to sub-$\mu$m-sized silicate grains \citep{Ma70}. Usually this spectral feature is composed of the broad features originating in amorphous silicates and the sub-peaks originating in crystalline silicates (\citealt{Ha94, Cr97, Wo99, Wo04, Ho04, Oo07a}, and references therein). Because there is no clear spectral evidence for crystalline silicates in interstellar medium \citep{Ke04}, the formation location for the crystalline silicate is still under debate. A possibility is that the crystalline silicates in comets were probably produced in the solar nebula from the interstellar amorphous silicate grains, by the process of thermal annealing or condensation in the inner region of the solar nebula (e.g., \citealt{Ha00, Fa00, HaDe02, Ga04}). Other possible mechanisms to make crystalline silicate in the space were proposed: an in situ annealing at a larger distance from the central star by a shock wave \citep{DeCo02, HaDe02}, an annealing inside the clump fragmented from a massive disk \citep{Vo11}, an episodic heating by intense radiation from a central star in outburst phase in surface layer of protoplanetary disks \citep{Ab09, Ju12}. These crystalline silicate grains were incorporated into cometary nuclei that formed far from the proto-Sun ($\sim$10-30 au) after the grains were transported to the outside of the solar nebula due to some mechanisms; e.g., the radial and vertical mixing of materials \citep{Boc02} or the X-wind \citep{Br12}. Therefore, the mass fraction of crystalline silicates with respect to the total (amorphous + crystalline) silicates is expected to be smaller for further distances from the Sun in the solar nebula \citep{Ga01, Boc02}. If the crystalline silicates formed by the direct condensation from gas-phase in the solar nebula \citep{Ga04, De17}, the forming region of the crystalline silicate grains was closer to the Sun than the case of thermal annealing process since the condensation of crystalline silicate grain requires higher temperatures ($\sim$1200-1400 K; \citealt{Ga04}) than the case of thermal annealing process for crystallization ($\sim$800 K; \citealt{Ga01}). In any case, the mass fraction of crystalline silicates in cometary grains is a clue to the place of comet-forming region in the solar nebula and the physical parameters of the solar nebula such as mass accretion rates and viscosity of the disk (e.g., \citealt{Ga01, Ga04, Boc02}). Namely, in addition to the mass fraction of crystalline silicate, abundance ratios of minerals composing cometary grains also give clues to the comet-forming region and the disk parameters. \citet{Ga04} theoretically demonstrated that abundance ratios of various minerals depend on distances from the proto-Sun. For the comet-forming region, olivine was estimated to be more abundant than pyroxene but the abundance ratios between olivine and pyroxene depend on distances from the proto-Sun according to their calculations. \citet{Li08} claimed that the abundance ratio between olivine to pyroxene silicate might be an indicator of system evolutionary age, based on mid-infrared spectroscopic observations of proto-planetary/circumstellar disks. Comet 17P/Holmes is classified as a Jupiter-family comet with an orbital period of 6.9 years. The comet was discovered by E. Holmes on UT 1892 November 6 \citep{Ho92}. The Tisserand criterion with respect to Jupiter is 2.86, computed from its latest orbital parameters listed in Nakano note\footnote{url{http://www.oaa.gr.jp/$\sim$oaacs/nk.htm}} (NK2484). The effective radius of the comet as spherical sphere was estimated as 1.71 km based on imaging photometric observations by the Hubble space telescope \citep{La00}. This value is similar to the effective radius of 1.62 $\pm$ 0.01 km (for the equivalent spherical body) with a large elongation of the nucleus with a large axial ratio of a/b $\geq$ 1.3 from the oscillations of brightness at large distances by using the 3.6-m New Technology Telescope \citep{Sn06}. The comet underwent a great outburst starting at UT 2007 October 23.3 \citep{Wa09, Hs10}, five months after perihelion at 2.05 au on UT 2007 May 5. This outburst reached maximum brightness of $\sim$2-3 mag in $V$-band as total magnitude within two days from the start of the outburst reported by many amateur astronomers. Such a huge outburst, which became brighter by 15 magnitudes, was unlike any other. Other than this outburst in 2007, comet 17P/Holmes had exhibited outbursts in November 1892 \citep{Ho92} and January 1893 \citep{Ba96}. According to \citet{Sc09}, the total amounts of water ice and dust grains during the 2007 apparition of 17P/Holmes were estimated to be $\sim$10$^{10}$ kg and 10$^{11}$ kg, respectively, corresponding to $\sim$0.2\% and at least 1-2\% of the total nucleus volume. \citet{Is10} reported the total mass injected into the coma by the outburst was estimated to be $>$4$\times$10$^{10}$ kg. \citet{Re10} found three different size components of dust grains and estimated total ejecta mass of $\sim$10$^{10}$ kg in the ejecta of the outburst in 2007, based on their mid-infrared spectroscopic and imaging observations by the $Spitzer$ Space Telescope. Note that their mid-infrared observations may slip past very small (sub-micron size) grains. \citep{Ya09} detected not only dust grains but also cold, icy, micron-size grains during the outburst. The enrichment of high-volatile gaseous species such as CH$_{3}$OH, C$_{2}$H$_{6}$, C$_{2}$H$_{2}$ in the coma was also demonstrated immediately after the outburst from the near-infrared spectroscopic observations \citep{DR08}. Some studies have proposed mechanisms for those outbursts but are not yet conclusive; e.g., the phase-transition of amorphous water ice to crystalline water ice with a release of a huge amount of latent heat and gaseous volatiles \citep{Re10}, the vaporization of hypervolatile ice in the cometary nucleus \citep{Sc09}, the avalanche or landslide of comet surface \citep{Boi02, Br04}, and the POP model in which a cometary outburst was triggered by plugging pores and blocking the release of surface gas flow by the recrystallization of water in the surface regolith \citep{de16}. If a large amount of ice of amorphous water and/or hypervolatiles is required for the large-scale outbursts seen in comet 17P/Holmes, its nucleus was considered to form under lower temperature conditions (i.e., further distances from the Sun) than other normal comets in which no large-scale outburst was seen repeatedly. An enrichment of highly volatile species (such as CH$_{3}$OH, C$_{2}$H$_{6}$, C$_{2}$H$_{2}$) compared to other typical comet indicates that comet 17P/Holmes formed under lower-temperature conditions (further from the Sun in the solar nebula) and therefore it contains less crystalline silicate grains compared to other normal comets that did not show large-scale outbursts. To confirm this hypothesis, we derived the mass fraction of crystalline silicates in comet 17P/Holmes based on the mid-infrared low-resolution spectra of the comet taken immediately after its outburst in 2007. | \label{sec:discussion} \subsection{Formation region of dust of comet 17P/Holmes} \label{subsec:formation_region} As listed in Table \ref{tab:table3}, the dust grains of comet 17P/Holmes are compact ($D$ = 3.0), abundant in sub-$\mu$m-sized dust particles ($a_{\rm p} \sim$0.1 $\mu$m and $N \sim$3; a smaller peak size and a smaller power-law index of the size distribution for sub-$\mu$m sized dust grains) compared to other comets. Temporal evolution of size distributions for sub-micron-sized dust grains in comet 17P/Holmes (see Table \ref{tab:table2}) is not significant. It is likely that these properties of dust grains are almost intrinsic and pristine (not affected by the mega-burst events, such as fragmentation). Because a typical grain size increased with time by collisional growth in the solar nebula, such small dust grains in comet 17P/Holmes might form in an earlier phase or in the outer part of the solar nebula. In contrast with small grains (accessed by the 10-$\mu$m silicate feature), it is demonstrated that larger dust grains (mm-size and larger) have a steeper power-law index of size distribution \citep{Is10, Bo12}. In comparison with other comets, comet 17P/Holmes shows smaller $f_{\rm cry}$ among comets found in previous articles (Table \ref{tab:table3}). In general, it is thought that the smaller $f_{\rm cry}$ of a comet indicates further distances from the Sun for the comet-formation (the larger $f_{\rm cry}$ for closer distances from the Sun). We conclude that comet 17P/Holmes formed in the outer region of the solar nebula. An enrichment of hyper-volatile species (such as CH$_{3}$OH, C$_{2}$H$_{6}$, C$_{2}$H$_{2}$) compared to other comets was demonstrated by near-infrared observations immediately after the outburst of the comet \citep{DR08, DR16} may support our hypothesis. Furthermore, some mechanisms of outburst of 17P/Homes were suggested \citep{Re10, Sc09, Boi02, Br04, de16}, and highly volatile abundant species (such as CO, CO$_{2}$) are more likely to contribute to the outburst of 17P/Holmes significantly. \begin{deluxetable*}{ccccccccc}[ht!] \tablenum{3} \tablecaption{Fitting results of the thermal spectra of 17P/Holmes by using the thermal model\label{tab:table3}} \tablewidth{0pt} \tablehead{ \colhead{Comet} & \colhead{UT Date} & \colhead{$r_{\rm H}$} & \colhead{$D$} & \colhead{$a_{\rm p}$ ($\mu$m)} & \colhead{$N$} & \colhead{$f_{\rm cry}$} & \colhead{$f_{\rm OP}$} & \colhead{References} } \startdata 17P/Holmes & Weighted mean & 2.45 & --- & --- & --- & 0.31$\pm$0.03 & 1.20$^{+0.16}/_{-0.12}$ & This work \\ & 2007 Oct 25 & 2.44 & 3.0 & 0.10$^{+0.01}/_{-0.00}$ & 3.34$^{+0.09}/_{-0.03}$ & 0.28$\pm$0.04 & 1.21$^{+0.20}/_{-0.17}$ & This work \\ & 2007 Oct 26 & 2.45 & 3.0 & 0.10$^{+0.03}/_{-0.00}$ & 3.07$^{+0.14}/_{-0.03}$ & 0.30$^{+0.12}/_{-0.10}$ & 1.73$^{+0.75}/_{-0.53}$ & This work \\ & 2007 Oct 27 & 2.45 & 3.0 & 0.10$^{+0.12}/_{-0.00}$ & 2.93$^{+0.44}/_{-0.03}$ & 0.44$^{+0.13}/_{-0.09}$ & 0.93$^{+0.29}/_{-0.21}$ & This work \\ & 2007 Oct 28 & 2.45 & 3.0 & 0.10$^{+0.14}/_{-0.00}$ & 3.00$^{+0.18}/_{-0.02}$ & 0.37$^{+0.08}/_{-0.07}$ & 1.95$^{+0.63}/_{-0.47}$ & This work \\ 9P/Tempel 1 & 2005 Jul 4 (1.0 h\tablenotemark{b}) & 1.51 & 2.857 & 0.3 & 3.7 & 0.13\tablenotemark{c} & 0.92\tablenotemark{c} & 1 \\ & 2005 Jul 4 (1.8 h\tablenotemark{b}) & 1.51 & 2.857 & 0.5 & 3.7 & 0.36\tablenotemark{c} & 7.22\tablenotemark{c} & 1 \\ & 2005 Jul 4 (3.5 h\tablenotemark{b}) & 1.51 & 2.857 & 0.4 & 3.6 & 0.83$\pm$0.10 & 6.5$\pm$1.9 & 2 \\ 73P-B/SW3\tablenotemark{a} & 2006 Apr 29 & 1.11 & 2.727 & 0.5 & 3.4 & 0.45$\pm$0.21 & 0.25$\pm$0.16 & 3 \\ 73P-C/SW3\tablenotemark{a} & 2006 Apr 30 & 1.09 & 2.727 & 0.3 & 3.4 & 0.52$\pm$0.13 & $>$17 & 3 \\ C/1995 O1 & 1996 Oct 11-14 & 2.8 & 2.8 & 0.2 & 3.4 & 0.53$\pm$0.04 & 2.65$\pm$0.51 & 4 \\ & 1997 Feb 14-15 & 1.21 & 2.5 & 0.2 & 3.7 & 0.47$\pm$0.01 & 1.55$\pm$0.07 & 4 \\ & 1997 Apr 11 & 0.97 & 2.5 & 0.2 & 3.7 & 0.62$\pm$0.02 & 2.26$\pm$0.17 & 4 \\ & 1997 Jun 24-25 & 1.7 & 2.7 & 0.2 & 3.7 & 0.56$\pm$0.04 & 1.57$\pm$0.08 & 4 \\ C/2001 Q4 & 2004 May 11 & 0.97 & 3.0 & 0.3 & 3.7 & 0.70\tablenotemark{c} & 3.57\tablenotemark{c} & 5 \\ & 2004 Jun 4 & 1.02 & 3.0 & 0.2 & 3.6 & 0.71\tablenotemark{c} & 6.88\tablenotemark{c} & 6 \\ C/2002 V1 & 2003 Jun 10 & 1.18 & 2.857 & 0.5 & 3.5 & 0.66\tablenotemark{c} & 2.63\tablenotemark{c} & 6 \\ C/2007 N3 & 2009 Mar 3 & 1.45 & 2.73 & 0.9 & 4.2 & 0.43$\pm$0.15 & 0.35$\pm$0.11 & 7 \\ \enddata \tablenotetext{\tiny a}{we refer to the results of the extraction boxes of opt-center (offset of 0 arcsec) for comet 73P/Schwassmann-Wachmann 3, corresponding to apertures B and M in Figure 2 of \citet{Ha11} for fragments B and C of the comet 73P/SW3, respectively.} \tablenotetext{\tiny b}{time after the impact of the Deep Impact Mission in hours.} \tablenotetext{\tiny c}{there is no description about error estimation in the applicable paper.} \tablecomments{$f_{\rm cry}$ and $f_{\rm OP}$ of all comets are calculated by expressions (1) and (2), respectively.} \tablenotetext{\tiny}{{\bf References.} [1] \citet{Ha05}, [2] \citet{Oo07b}, [3] \citet{Ha11} [4] \citet{Ha04} (Erratum of \citealt{Ha02}), [5] \citet{Wo04}, [6] \citet{Oo07a}, [7] \citet{Wo11}. } \end{deluxetable*} \begin{deluxetable}{lcc} \tablenum{4} \tablecaption{mass fraction of crystalline silicate, $f_{\rm cry}$, for each material of 17P/Holmes \label{tab:table4}} \tablewidth{0pt} \tablehead{ \colhead{UT Time} & \multicolumn2c{$f_{\rm cry}$} \\ \colhead{} & \colhead{Olivine} & \colhead{Pyroxene} } \startdata 2007 Oct 25 & 0.36 $\pm$ 0.02 & 0.19 $\pm$ 0.06 \\ 2007 Oct 26 & 0.33 $\pm$ 0.03 & 0.24 $\pm$ 0.17 \\ 2007 Oct 27 & 0.38 $\pm$ 0.02 & 0.56 $\pm$ 0.09 \\ 2007 Oct 28 & 0.36 $\pm$ 0.01 & 0.46 $\pm$ 0.09 \\ Mean & 0.36 $\pm$ 0.01 & 0.34 $\pm$ 0.04 \\ \enddata \end{deluxetable} Here we compared our fitting results with theoretical predictions by \citet{Ga01, Ga04} and \citet{Boc02}, considering equilibrium crystallization near the Sun as dust formation and mass transportation from the inner to outer region in the solar nebula. However, the fraction of crystalline silicate grains predicted by these models may not be reliable as absolute values although their trends (smaller $f_{\rm cry}$ for further distance from the Sun) are reasonable. The formation scenario of crystalline silicate grains in the early solar system and how to derive them into cometary nuclei are still in debate (e.g., \citealt{TaNo15}), and we strongly encourage comparison of further theoretical studies with cometary observations. \citet{Pi16} reported that the proto-planetary disk of HL Tau has small turbulent viscosity coefficient of a few 10$^{-4}$ by radio observations, implying that the turbulence might not be so large as to be able to transport a large amount of crystalline silicate grains to comet-forming region (further from the snowline of water) even if turbulence occurred in the solar nebula. Thus, it is in need of other transport mechanisms to the comet-forming region; e.g., X-wind, vertical mixing in the solar nebula \citep{Sh97}. \subsection{Formation mechanisms of cometary dust} \label{subsec:formation_mechanisms} In the thermal emission model, we assume that mass ratio of olivine originating in interstellar medium (ISM) and solar nebula (SN) is equal to that of pyroxene. If this assumption is incorrect, there is no reason for $f_{\rm cry}$ of olivine to become the same as that of pyroxene in general. We may have to consider different transportation mechanisms for different materials. For instance, grain properties such as typical size and porosity might be different between olivine and pyroxene. In this case, the efficiency of transportation by radiation pressure for olivine is expected to be different from that for pyroxene \citep{TaNo15}. Table \ref{tab:table4} shows that values of $f_{\rm cry}$ for both olivine ($f_{\rm cry,olivine} = m_{\rm CryOl} / (m_{\rm AmoOl} + m_{\rm CryOl})$) and pyroxene ($f_{\rm cry,pyroxene} = m_{\rm CryPy} / (m_{\rm AmoPy} + m_{\rm CryPy})$) in all date are in agreement within 3$\sigma$ error bars. Based on this result, we consider that mass ratio of origin of olivine (ISM/SN mass ratio) is equal to that of pyroxene as working hypothesis to discuss the difference in $f_{\rm cry}$ between olivine and pyroxene. \begin{figure} \plotone{./fig04.eps} \caption{Expanded spectrum of 17P/Holmes taken on 2007 October 25. The pairs of six vertical lines (red, orange, green, blue, purple, and black in the order left to right) indicate peak wavelength of 10.0 $\mu$m, 10.4 $\mu$m, 11.2 $\mu$m, and 11.9 $\mu$m with the Mg/(Mg+Fe) ratio of 100\%, 80\%, 60\%, 40\%, 20\%, and 0\%, respectively. Peak wavelengths of each Mg/(Mg+Fe) ratio are calculated from the expressions listed in Table 3 of \citet{Ko03}. Gray vertical hatch is the absorption band of telluric ozone. \label{fig:fig4}} \end{figure} We assume that dust grains originating in ISM and in the inner solar nebula ($<$0.1 au) have $f_{\rm cry}$ = 0.2\% $\pm$ 0.2\% and $f_{\rm OP}$ = 5.6 considering the observation toward Galactic center sources \citep{Ke04} and $f_{\rm cry}$ = 100\% and $f_{\rm OP}$ = 0.1 derived by theoretical calculation based on the equilibrium condensation condition \citep{Ga04}, respectively. We define two parameters, $x$ and $\gamma$, where the mass ratios of dust from the ISM or from the SN are $x$ (0 $\leq x \leq$ 1) and 1-$x$, respectively, and the mass ratios of dust from the ISM captured directly into the cometary nucleus or changing mineral form (e,g., crystallization and conversion from forsterite to enstatite) in the SN are $\gamma$ (0 $\leq \gamma \leq$ 1) and 1-$\gamma$, respectively. Under these conditions, $f_{\rm cry}$ and $f_{\rm OP}$ of comets are given by \begin{equation} f_{\rm cry,comet} = \gamma xf_{\rm cry,ISM} + (1 - \gamma)xf_{\rm cry,SN} + (1-x)f_{\rm cry,SN} \end{equation} and \begin{equation} f_{\rm OP,comet} = \gamma xf_{\rm OP,ISM} + (1 - \gamma)xf_{\rm OP,SN} + (1-x)f_{\rm OP,SN} , \end{equation} where subscript X of $f_{\rm cry,X}$ and $f_{\rm OP,X}$ means comet as observed value in comet, ISM as interstellar medium, SN as solar nebula, respectively. The first term of the right-hand side of both formula is dust of ISM origin and was incorporated into the comet directly, the second term means dust of ISM origin that was denatured in SN, and the third term is dust originating in SN. We found that there is no solution that satisfies both expressions simultaneously, when we substitute the weighted means of $f_{\rm cry}$ and $f_{\rm OP}$ in 17P/Holmes (e.g., $f_{\rm cry}$ = 0.31 $\pm$ 0.03 and $f_{\rm OP}$ = 1.20$^{+0.16}$/$_{-0.12}$). For other comets for which both $f_{\rm cry}$ and $f_{\rm OP}$ were reported (summarized in Table \ref{tab:table3}), there is no unique solution. These results may support the interpretations that most dust grains of comets including 17P/Holmes did not form by the equilibrium condensation and annealing directly from gas-phase in the solar nebula and/or mass ratio of olivine originating in both ISM and SN are different from those of pyroxene. | 18 | 8 | 1808.07606 |
1808 | 1808.05603_arXiv.txt | {We use stellar kinematics from the latest {\it Gaia} data release (DR2) to measure the local dark matter (DM) density $\rho_{\rm DM}$ in a heliocentric cylinder of radius $R= 150 \ {\rm pc}$ and half-height $z= 200 \ {\rm pc}$. We also explore the prospect of using our analysis to estimate the DM density in local substructure by setting constraints on the surface density and scale height of a thin dark disk aligned with the baryonic disk and formed due to dissipative dark matter self-interactions. Performing the statistical analysis within a Bayesian framework for three types of tracers, we obtain ${\rho_{\rm DM}= 0.016 \pm 0.010}$ M$_\odot$/pc$^3$ for A stars; early G stars give a similar result, while F stars yield a significantly higher value. For a thin dark disk, A stars set the strongest constraint: excluding surface densities (5-12) M$_\odot$/pc$^2$ for scale heights below 100 pc with 95\% confidence. The upper bound of this constraint implies ${\lsim} \, 1\%$ of the Milky Way DM mass is present in a dissipative dark sector. Comparing our results with those derived using {\it Tycho-Gaia} Astrometric Solution (TGAS) data, we find that the uncertainty in our measurements of the local DM content is dominated by systematic errors that arise from assumptions of our dynamical analysis in the low $z$ region. Furthermore, there will only be a marginal reduction in these uncertainties with more data in the {\it Gaia} era. We comment on the robustness of our method and discuss potential improvements for future work. } \begin{document} | The second release of data collected by the European Space Agency's {\it Gaia} telescope provides the positions and proper motions, with unprecedented precision, of more than one billion sources in the Milky Way (MW)~\cite{2016A&A...595A...1G,2018arXiv180409365G,2018arXiv180409366L,2018arXiv180409368E,2018arXiv180409367R,2018arXiv180409371S, 2018arXiv180409369C, 2018arXiv180409372K, 2018arXiv180409376L}. With the release of line-of-sight velocities for about seven million stars, DR2 also allows, for the first time, a dynamical analysis with a self-consistent measurement of the 6D phase space for a stellar population. DR2 presents an exciting opportunity to use the vertical velocity and number density distributions of different populations of stars that trace the gravitational potential for precisely determining the total matter density, including baryons and dark matter (DM), in the local solar neighborhood. Significant progress has been made in modeling the local baryon budget (interstellar gas, stars, stellar remnants) and its uncertainties~\cite{Flynn:2006tm, Bovy:2013raa, Mckee:2015, Kramer:2016dew} since Oort's early estimate~\cite{1932BAN.....6..249O} of the baryon density. Meanwhile, kinematic methods for estimating the local DM density rely on constraining the total matter content using motions of tracers after assuming a model for the baryons and attributing any additional density, within uncertainty, to DM. These methods are based on: {\it a)} the Jeans analysis that reduces the collisionless Boltzmann equation for the phase space distribution function into a set of moment equations by integrating over all velocities, and {\it b)} the Poisson equation which uses the total matter density in all components to calculate the gravitational potential. In this work, we primarily focus on the 1D distribution function method developed by Refs.~\cite{1989MNRAS.239..571K,1989MNRAS.239..605K,1989MNRAS.239..651K, 1993AIPC..278..580F, 1994MNRAS.270..471F} and used by Refs.~\cite{Holmberg:1998xu, 2004MNRAS.352..440H} to constrain the local DM density with data from the {\it Hipparcos} satellite~\cite{vanLeeuwen:2005yx}. However, the approximations of isothermality and decoupling of radial and vertical motions in this method are only valid up to scale height $z \, {\sim} 1 \ {\rm kpc}$. Therefore, for using tracer data at high $z$, Refs.~\cite{garbari:2012, Bovy:2012tw} adopt the more general moment-based method to estimate the DM density. A non-parametric formulation of the moment-based method, described by Ref.~\cite{Silverwood2016} and implemented in Ref.~\cite{Silverwood2017}, uses SDSS/SEGUE G stars in a heliocentric cylinder with $ R \, {\sim} 1 \ {\rm kpc}$ and $0.5 \ {\rm kpc} \, {\lsim} \, |z| \, {\lsim} \, 2.5 \ {\rm kpc}$, grouped by age, namely $\alpha$-young and $\alpha$-old stars, as tracers. Ref.~\cite{Bovy:2013raa} also uses SDSS/SEGUE G star data between $4 \ {\rm kpc} \, {\lsim} \, R \, {\lsim} \, 9 \ {\rm kpc}$ and $0.3 \ {\rm kpc} \, {\lsim} \, |z| \, {\lsim} \, 3 \ {\rm kpc}$ to constrain the stellar and DM density through action-based distribution function modeling \cite{Bovy2012a}. Their analysis incorporated the age information of tracers in a more sophisticated manner by constructing mono-abundance populations~\cite{Bovy2012d} that consist of stars with similar elemental abundances. The above discussion is by no means an exhaustive review of the different attempts at measuring the local DM density (most notably, it doesn't address dynamical measurements made by, for example, Refs.~\cite{Salucci:2010qr, Pato:2015dua}); instead, we refer interested readers to Refs.~\cite{2013PhR...531....1S, Read:2014qva}. Besides determining the local DM density, one could apply the data of stellar kinematics to constrain more exotic DM distributions. A recent example is dissipative DM scenarios, in which self-interactions among a sub-dominant component of DM dissipates energy to form a (possibly thin) dark disk (DD), co-rotating with the baryonic disk~\cite{Fan:2013tia, Fan:2013yva}. Possible effects of such a DD and variants of dissipative DM scenarios have been studied further in Refs.~\cite{CyrRacine:2012fz, McCullough:2013jma, Fan:2013bea, Randall:2014lxa, Fischler:2014jda, Foot:2014uba, Randall:2014kta, Reece:2015lch, Foot:2016wvj, Shaviv:2016umn, Kramer:2016dqu, Kramer:2016dew, Agrawal:2017rvu, Agrawal:2017pnb, Buckley:2017ttd, DAmico:2017lqj, Caputo:2017zqh, Vattis:2018aen, Outmazgine:2018orx, Alexander:2018lno}. The formation of the disk from DM self-interactions is highly debatable~\cite{Ghalsasi:2017jna} and numerical simulations using a cooling prescription (as in Ref.~\cite{Rosenberg:2017qia}) are still absent. Inspite of these uncertainties regarding its formation, it is still worthwhile to use the stellar data to test the simplest possibility of a thin DD aligned with the baryonic disk and parametrized by only two parameters: the surface density, $\Sigma_{DD}$ and a scale height, $h_{DD}$. This has been carried out using {\it Hipparcos} data in Ref.~\cite{Kramer:2016dqu} and {\it Tycho--Gaia} Astrometric Solution (TGAS), a joint solution combining {\it Tycho}-2 catalog with early {\it Gaia} data, in Ref.~\cite{Schutz:2017tfp}. \begin{figure*} \begin{center} \includegraphics[width=0.9\linewidth]{flowchart.pdf} \end{center} \caption{Flowchart of our analysis.} \label{fig:flowchart} \end{figure*} In this article, we work with the second {\it Gaia} data release (DR2)~\cite{2018arXiv180409365G} to estimate the local DM density as well as constrain thin DD models assuming that the DD is aligned with the baryonic disk. We follow the method in Refs.~\cite{Holmberg:1998xu, Kramer:2016dqu, Schutz:2017tfp} and use A, F and early G dwarf stars in the {\it Gaia} catalog as the tracers. In Section~\ref{sec:data}, we discuss the details of {\it Gaia} DR2 and the empirically determined survey selection function (Section~\ref{sec:Selection_Function}), which we use to construct the vertical number density profile and midplane velocity distribution for each tracer population in Sections~\ref{sec:numberdensity} and \ref{sec:midplane_velocity} respectively. Our fiducial analysis is described in Section~\ref{sec:analysis}. We use the 1D distribution function method summarized in Section~\ref{sec:PJ_Theory} to construct the equilibrium number density for the parameters of our mass model described in Section~\ref{sec:mass_model}. In Sections~\ref{sec:basic_setup} and \ref{sec:BM_discuss}, we introduce a Bayesian framework for comparing our predicted density with data for each tracer population while taking into account the uncertainties due to potential non-equilibrium effects. The important steps of our analysis are outlined as a flowchart in Fig.~\ref{fig:flowchart}. While our method is not new, we obtain interesting results, some of which are quite different from those based on TGAS. We present our results for the local DM content using {\it Gaia} DR2 in Section~\ref{sec:localdm_no_dd} and~\ref{sec:localdm_dd}, and list various sources of systematic uncertainties in the context of our method in Sections~\ref{sec:density_val}-\ref{sec:bary_dm}. Differences between constraints derived using DR2 and TGAS are discussed in Section~\ref{sec:localdm_tgas}. We conclude and comment on future directions in Section~\ref{sec:outlook}. | \label{sec:outlook} We apply the 1D distribution function method to {\it Gaia} DR2 and use stellar kinematics in the solar neighborhood to constrain the local DM density and properties of a thin DD aligned with the baryonic disk by performing our analysis within a Bayesian framework. We adopt young A, F, and early G stars as tracers as they have shorter equilibration timescales and consequently are expected not to be strongly affected by disequilibria. Using A stars gives an estimate of $\rho_{\rm DM}= 0.016 \pm 0.010$ M$_\odot$/pc$^3$ and sets the strongest constraint on the thin DD, excluding $\Sigma_{DD} \, {\gsim}$ (5-12) M$_\odot$/pc$^2$ depending on the scale height with 95\% confidence. This upper bound is used to constrain the amount of dissipative DM in the galaxy: a thin DD with $\Sigma_{DD} \, {\lsim} \, 12$ M$_\odot$/pc$^2$ and a scale radius ${\sim} 3 \, \rm{kpc}$ contains ${\lsim} \, 1\%$ of the total DM mass in the Milky Way \cite{Fan:2013tia}. While we obtain similar results from early G stars, F stars seem to prefer a higher value of the local DM content. Even though the distributions derived from DR2 are consistent with those from TGAS data within uncertainties, the allowed DM density and parameters of DD model are quite different for all tracers. In light of these results, we address the origins of the differences and discuss the robustness of our kinematic analysis. Our results also suggest that we need a better understanding of the physical origin of the systematic uncertainties, which we include in our analysis to account for the asymmetry in the midplane velocity distributions of tracers. One possibility is that with complete data for radial velocities, we could define the midplane region using the $z$-cut instead of the $b$-cut and obtain a more precise determination of the velocity distribution. Another possibility is to take a closer look at local disequilibria and their effects on traditional kinematic methods. Although we do not find any statistically significant evidence for non-equilibrium in the vertical density and velocity distributions in our samples, several analyses based on DR2 seem to suggest various sources of disequilibria at distances larger than the heliocentric cylinder we consider. In terms of baryon modeling, it could be useful to find a self-consistent, data-driven approach to determine the baryon distributions instead of assuming the isothermal Bahcall model. One way to achieve this would be to construct the mass density for stars directly from the data rather than treating it as an isothermal disk. For a more precise determination of the local DM density, a dynamical analysis could be performed using tracers at heights greater than the scale height of the stellar disk to minimize the latent degeneracy between baryons and DM. However, besides modeling effects of disequilibria, an analysis at larger scale height has to go beyond the 1D method and must include terms that couple the motions of tracers in different directions. We also see a degeneracy between parameters of ordinary DM and thin DD in the marginalized posteriors obtained through MCMC sampling. To break the degeneracy, we would need to distinguish between their effects on tracers by developing new observables and modeling priors that reflect these differences. | 18 | 8 | 1808.05603 |
1808 | 1808.07430_arXiv.txt | {Scattering and decay processes of thermal bath particles involving heavy leptons can dump hot axions in the primordial plasma around the QCD phase transition. We compute their relic density, parameterized by an effective number $\Delta N_{\rm eff}$ of additional neutrinos. For couplings allowed by current bounds, production via scattering yields $\dN \lesssim 0.6$ and $\dN \lesssim 0.2$ for the cases of muon and tau, respectively. Flavor violating tau decays to a lighter lepton plus an axion give $\Delta N_{\rm eff} \lesssim 0.3$. Such values of $\dN$ can alleviate the tension between the direct local measurement of the Hubble constant $H_0$ and the inferred value from observations of the Cosmic Microwave Background, assuming $\Lambda$CDM. We analyze present cosmological data from the Planck collaboration and baryon acoustic oscillations with priors given in terms of the axion-lepton couplings. For axions coupled to muons, the tension can be alleviated below the 3$\sigma$ level. Future experiments will measure $\Delta N_{\rm eff}$ with higher precision, providing an axion discovery channel and probing the role of hot axions in the $H_0$ tension.} \begin{document} | \label{sec:intro} Axion-like particles (ALPs) are motivated candidates for extremely light and weakly-coupled degrees of freedom beyond the Standard Model (SM). The motivation is notably robust for the QCD axion, as the Peccei-Quinn (PQ) mechanism is an elegant solution to the strong CP problem~\cite{Peccei:1977hh,Peccei:1977ur} and it could account for the observed dark matter (DM) abundance~\cite{Wilczek:1977pj,Weinberg:1977ma,Kim:1979if,Shifman:1979if,Zhitnitsky:1980tq,Dine:1981rt,Preskill:1982cy,Abbott:1982af,Dine:1982ah}. Axion phenomenology is quite broad, detection strategies are multiple and complementary. Direct interactions with SM particles, unavoidably present for the QCD axion, are the subject of present and future searches~\cite{Arik:2008mq,Irastorza:2011gs,Graham:2013gfa,Budker:2013hfa,Irastorza:2013dav,Armengaud:2014gea,Kahn:2016aff,Barbieri:2016vwg,Melcon:2018dba,Du:2018uak}. This paper focuses on a peculiar cosmological imprint: scatterings and decays of thermal bath particles can produce relativistic axions throughout the expansion history of our universe~\cite{Turner:1986tb,Berezhiani:1992rk,Masso:2002np,Graf:2010tv,Brust:2013xpv,Salvio:2013iaa,Baumann:2016wac,Ferreira:2018vjj}. The resulting growth in the radiation energy density, traditionally parameterized by $\dN$ additional neutrino species, can be probed by observations of the Cosmic Microwave Background (CMB) and Baryon Acoustic Oscillations (BAO)~\cite{doi:10.1080/00107514.2013.818063, Lesgourgues:2014zoa, Archidiacono:2013fha}. The current best-fit value from the latest Planck 2018+BAO data \cite{Planck2018_cosmopars}, $N_\text{eff}=2.99 \pm 0.17$, is in agreement with the SM prediction, $N_\text{eff(SM)} = 3.046$~\cite{Dodelson:1992km,Hannestad:1995rs,Dolgov:1997mb,Mangano:2005cc,deSalas:2016ztq}. Ongoing and forthcoming experimental efforts make the present time remarkably exciting. Current experiments~\cite{SPT-3G} will reach a sensitivity $\dN\sim 0.06$ in the near future, whereas future CMB-S4 surveys forecast an improvement down to $\dN\sim 0.024$~\cite{Abazajian:2016yjj}. As a useful benchmark, a $2 \sigma$ detection with the latter sensitivity would require hot axions that were in thermal equilibrium at early times and then decoupled at temperatures $T \lesssim 10 \, {\rm GeV}$~\cite{Ferreira:2018vjj}. Combined with the complementary information from direct searches, this makes the next few years very promising for the quest for axions. Recently, a mismatch between low and high redshift determinations of the Hubble constant $H_0$ triggered further interest in $\dN$. The value inferred from CMB observations (assuming $\Lambda$CDM), $H_0=67.27\pm 0.60$ km $\text{s}^{-1}$ Mpc$^{-1}$~\cite{Planck2018_cosmopars}, is in $3.6\sigma$ tension with direct measurements from supernovae, $H_0= 73.52 \pm 1.62$ km $\text{s}^{-1}$ Mpc$^{-1}$ \cite{Riess:2018byc}. The latter also include data from the latest GAIA release \cite{2018arXiv180409365G}. Such a tension, although weaker, also appears in other independent measurements of $H_0$ that probe the local universe and observations at high redshift (see, {\it e.g.}, Ref.~\cite{Bernal_baccus,Dhawan:2017ywl}). Notably, BAO are a complementary probe that depends on an assumed expansion history of the Universe and on the product $r_{\rm s} h$~\cite{Aubourg_2015,StandardQuantities}, where $r_{\rm s}$ and $h$ are the sound horizon at radiation drag and the reduced Hubble constant $h=H_0/100$ km $\text{s}^{-1}$ Mpc$^{-1}$, respectively. An independent determination of $r_{\rm s}$ would make it possible to infer the value of $H_0$ using BAO. Measurements of $H_0$ from BAO using $r_{\rm s}$ from the CMB \cite{Alam:2016hwk,Zarrouk:2018vwy} or from primordial deuterium measurements~\cite{Cooke_deu2017}, as in \cite{Addison_bao13,Addison_H0bao}, are also in tension with the direct local measurement. The combination of CMB and BAO leads to a 3.46$\sigma$ disagreement. The tension persists even if we parameterize the cosmic expansion at low redshifts in a model independent way~\cite{Bernal:2016gxb}, since supernova type~Ia measurements (up to $z\sim 1$) force $H(z)/H_0$ to be $\Lambda$CDM-like with deviations $ \lesssim 5\%$ at low redshift. Exotic dark energy models that change the expansion history at low redshift to reconcile the two $H_0$ measurements seem disfavored (see also~\cite{Mortsell:2018mfj}). Remarkably, having $\dN>0$ is a way to decrease the tension~\cite{RiessH0_2016,Bernal:2016gxb}; a combined fit~\cite{Planck2018_cosmopars} with the new 2018 Planck data, including lensing, BAO and the direct $H_0$ measurement~\cite{2018ApJ...855..136R} leads to $N_{\rm eff}=3.27\pm0.15$ and $H_0=(69.32\pm0.97) {\rm \, km \, s^{-1} \, Mpc^{-1}}$. Motivated by the above considerations, we study thermal production of axions coupled to heavy leptons. This mechanism is mostly active at temperatures around the lepton mass~\cite{Turner:1986tb}, hence close to the QCD phase transition (QCDPT).~\footnote{Bounds on axion-electron couplings~\cite{Viaux:2013lha} lead to unobservably small $\dN$. Couplings to heavy quarks ($c$, $b$ and $t$) give $\dN$ observable in the near future~\cite{Ferreira:2018vjj}, whereas couplings to light quarks ($u$, $d$ and $s$) require a treatment of hadronic bound states.} The resulting comoving relic abundance is quite large, since it is inversely proportional to the number of relativistic degrees of freedom that sharply decreases around this epoch. The consequence is of a twofold nature: future CMB experiments can observe axions; hot axions produced via this mechanism can alleviate the current $H_0$ tension. An estimate of $\dN$ via these couplings was provided in Refs.~\cite{Turner:1986tb,Brust:2013xpv,Baumann:2016wac}. Here, we perform a careful analysis by computing the full cross-sections and solving the Boltzmann equations to compute the precise value of $\dN$. Our results are valid also when a complete thermalization is never reached. With these results in hand, we re-analyze CMB and BAO data to reassess the $H_0$ tension in light of this theoretical framework. We present axion interactions with heavy leptons in Sec.~\ref{Lag}. After establishing the experimental bounds on such interactions, we compute in Sec.~\ref{Sec:production} the axion relic density for allowed couplings. The effect of such a hot axion population is parameterized by a number $\dN$ of additional neutrinos, as we quantify in Sec.~\ref{sec:dN}. Up to this point, we only assume axion derivative coupling to leptons. In Sec.~\ref{sec:QCD} we discuss the results in the context of the QCD axion. Then, we study how these hot axions and their theory based priors affect the tension in the measured value of $H_0$ in Sec.~\ref{sec:cosmo}, and we give our conclusions in Sec.~\ref{sec:Conclusion}. We also provide detailed appendices with our calculations. | \label{sec:Conclusion} We have studied hot axions production through scatterings and decays of heavy leptons (muon and tau). Our results hold for a generic ALP as well as for the QCD axion. Axion production via fermion scatterings was originally proposed in Ref.~\cite{Turner:1986tb}, previous estimates of $\dN$ due to leptons were given in Refs.~\cite{Brust:2013xpv,Baumann:2016wac}. Here, we have computed cross sections and decay rates, we have evaluated the thermal averages of these quantities to find the temperature dependent production rates and we have numerically solved the Boltzmann equation to compute the axion abundance. As a final result, we have found the axion contribution to the effective number of additional neutrinos $\dN$ as a function of the axion couplings. As summarized in Fig.~\ref{figura}, couplings within the allowed parameter space can lead to large signals: $\dN^\text{max} \simeq 0.6$ for muon scatterings, $\dN^\text{max} \simeq 0.3$ for tau decays and $\dN^\text{max} \simeq 0.2 $ for tau scatterings. These scenarios provide a well motivated particle physics framework for $\dN >0$ that can be further explored by current and future experiments~\cite{SPT-3G,Abazajian:2016yjj}. Given such results, we have also investigated the consequences for the current $3.6\sigma$ ($3.46\sigma$, combining with BAO) tension between high and low redshift measurements of the Hubble constant $H_0$. As already explored, {\it e.g.}, in \cite{Bernal:2016gxb}, large values of $\dN$ can alleviate the tension between the two datasets. Here, we have performed a similar analysis but with theory-based priors by choosing a flat prior in $\log(f/c_\ell)$, with the axion-lepton effective coupling defined as $c_\ell/f$. Our findings feature some dependence on the maximum value considered in the prior distribution, $(f/c_\ell)_{\rm max}$, since the probability of having very small $\dN$ is enhanced at large $(f/c_\ell)_{\rm max}$. Values $\dN \gtrsim 0.2$ are disfavored by the latest Planck 2018 temperature and polarization data combined with BAO data, so the tension cannot be completely erased by a non-vanishing $\dN$. In the scenarios studied here and summarized in Tab.~\ref{tab:tension}, we have also generically found moderate improvement in mitigating the $H_0$ tension. A notable exception is for hot axions produced via muon scattering and $(f/c_\ell)_{\rm max} \simeq 3\times 10^7$, where the tension is reduced to $2.75 \sigma$. Our findings provide a theoretically motivated origin for $\dN >0$ and motivate further studies of UV complete models with axions coupled to heavy leptons. Hot axions produced around or below the QCDPT open an exciting window to observe extremely weakly-coupled pseudo-scalars. Forthcoming results from CMB surveys and direct searches make the future of these hypothetical dark components of our universe very bright. | 18 | 8 | 1808.07430 |
1808 | 1808.10546_arXiv.txt | Faraday tomography is thought to be a powerful tool to explore cosmic magnetic field. Broadband radio polarimetric data is essential to ensure the quality of Faraday tomography, but such data is not easy to obtain because of radio frequency interferences (RFIs). In this paper, we investigate optimum frequency coverage of Faraday tomography so as to explore Faraday rotation measure (RM) due to the intergalactic magnetic field (IGMF) in filaments of galaxies. We adopt a simple model of the IGMF and estimate confidence intervals of the model parameters using the Fisher information matrix. We find that meaningful constraints for RM due to the IGMF are available with data at multiple narrowbands which are scattered over the ultra-high frequency (UHF, 300~MHz -- 3000~MHz). The optimum frequency depends on the Faraday thickness of the Milky Way foreground. These results are obtained for a wide brightness range of the background source including fast radio bursts (FRBs). We discuss the relation between the polarized-intensity spectrum and the optimum frequency. | \label{s1} \begin{figure*}[t] \begin{center} \includegraphics[width=160mm]{f01.eps} \end{center} \caption{ Radio frequency environment around the Kashima 34-m antenna of the National Institute of Information and Communications Technology (NICT) in Japan. The blue and red lines show the instantaneous (sweep time 2.05 second) and 5 minutes max-hold (i.e. maximum during 5 minutes) spectra, respectively. The band characteristic of the receiver is removed from the spectra, so that the vertical axis is the relative radio power with respect to the detection limit. } \label{f01} \end{figure*} Magnetic field is a fundamental element of the Universe and it affects formation and evolution of astronomical objects. Centimeter radio polarimetry is one of the promising tools to study cosmic magnetism (see \cite{han17, aka18} for reviews); synchrotron intensity, its linear-polarization vector, and Faraday rotation measure (RM) provide us with properties of magnetic field in galaxies and AGN jets, and they reveal detailed structures of magnetized plasma such as the interstellar medium (ISM) and intergalactic medium (IGM). Cosmic magnetism is one of the key sciences for the Square Kilometre Array (SKA) \citep{joh15}. Faraday RM synthesis or Faraday tomography (\cite{bur66, bre05}) grows up progressively in radio polarimetry. There are a lot of successful applications to the ISM \citep{sak18}, galaxies \citep{mao17}, radio lobes \citep{osu18}, quasars \citep{and15, and16}, and galaxy clusters \citep{oza15}. Furthermore, discovery of Fast radio bursts (FRBs) fosters momentum of the study of cosmic magnetism. As at April 2018, seven linearly-polarized FRBs are published in the literature (see \cite{cal18}). For example, \citet{mic18} observed FRB121102 and found strongly-magnetized medium with ${\rm RM} \sim O(10^5)$~${\rm rad~m^{-2}}$, implying an environment similar to that around a super massive black hole (SMBH). It has been predicted that the cosmic web is permeated with a large amount of magnetized IGM. \citet{aka14a} studied possible situations to estimate RM due to the intergalactic magnetic field (IGMF) by means of Faraday tomography, and demonstrated that the ultra-high frequency (UHF) band is promising to maximize the capability of Faraday tomography for the study. \citet{rav16} applied Faraday tomography to FRB150807 and derived an upper limit of the IGMF strength $<21$~nG, which does not contradict theoretical predictions (e.g., \cite{aka16, vaz18}). The above studies demonstrate the capability of Faraday tomography for a wide RM range of diffuse, compact, and even time-domain radio sources. Because wider frequency coverage gives better quality of Faraday tomography (e.g., \cite{aka14a}), a modern wideband observation makes Faraday tomography feasible. However, obtaining a seamless dataset over broad bandwidth is difficult. One of the essential reasons is radio frequency interferences (RFIs). Centimeter wavelength is commonly used in industry, such as broadcasting, mobile phone, wireless communication, and radar. Figure~\ref{f01} shows an example of RFIs (see Appendix for observational details). Appreciable signals are all RFIs against radio astronomy. These RFIs easily saturate amplifiers, produce artificial higher-harmonic signals, and alter the band characteristics fatally; they make signal processing unreliable. Persistent RFIs can be cut by frequency filters at an early stage of a receiver system, but this means that we never obtain astronomical signal at the frequencies. Although many large radio telescopes are located at countryside with low human population, radio frequency environment rapidly changes as human lifestyle improves\footnote{Indeed, ``Sky Muster" RFIs at ATCA 15 mm band and ``BSAT-4a" RFIs at VERA K band are very recently appeared.}. SKA-MID antennas will be constructed in radio-quiet districts in South Africa, but economic growth in South Africa would impact on radio frequency environment at the site in future. In this paper, we investigate the optimum frequency of Faraday tomography to explore RM due to the IGMF. Although \citet{aka14a} briefly considered RFIs on the SKA sites, a more comprehensive study about frequency dependence on Faraday tomography could maximize the chance to discover the IGM and IGMF through Faraday tomography. This paper is organized as follows. We explain our model and calculation in Section 2. The results are shown in Section 3 followed by discussion and summary in Section 4. | \label{s4} We found that the optimum frequency depends on the thickness of the Faraday spectrum for the foreground Milky Way emission. The optimum frequencies are $\sim 500$~MHz, $\sim 650$~MHz, $\sim 800$~MHz, $\sim 950$~MHz for $\delta\phi_{\rm MW}=$ 2, 4, 6, 8 ${\rm rad~m^{-2}}$, respectively, and it reaches $\sim 1400$~MHz if $\delta\phi_{\rm MW} = 15$~${\rm rad~m^{-2}}$. We find that the optimum frequency is close to the frequency at which foreground Milky Way emission is significantly depolarized at the observer frame (figure~\ref{f02}). Such depolarization is seen in the polarized-intensity spectrum, for example, at $\lambda^2 \sim 0.2$~${\rm m}^2$ in the bottom panel of figure~\ref{f02}. This depolarization is classified into differential Faraday rotation depolarization \citep{sok98, ars11}. \citet{ars11} investigated the optimum wavelength ($\lambda_{\rm opt}$) of the maximum polarized emission according to differential Faraday rotation, and proposed the equation of the optimum wavelength as \begin{equation}\label{eq06} |\sin k | - k |\cos k| = 0, \end{equation} where we consider a flat spectral index (the case of $\alpha=0$ in \cite{ars11}) and $k=2 |RM| \lambda_{\rm opt}^2$. The optimum frequencies that we found is in broadly agreement with the solution of $k= 2.0288~({\rm radian})$ for an effective RM value of $|RM|\sim 1.3 \delta \phi_{\rm MW}$. Therefore, the optimum frequencies can be explained by the depolarization theory. The above depolarization frequency, i.e. the optimum frequency, depends on the model Faraday spectrum of the Milky Way; we have considered a Gaussian shape and the intrinsic polarization angle is constant. We can consider more complicated, realistic FDFs of polarized sources \citep{ide14b}. However, this work focuses on a typical, global solution of the optimum frequency. A specific model is beyond the scope of this work and it will be considered in a separate paper. Nevertheless, if an actual Faraday spectrum of the Milky Way deviates from the Gaussian, the depolarization frequency can change. Note that the constraint on $RM_{\rm IGMF}$, i.e. the gap between MW and BG, primarily depends on the edge of the Faraday spectrum of the Milky Way rather than a detailed profile of the Faraday spectrum of the Milky Way. Throughout this paper, the total intensity is independent on the frequency so that a flat spectral index is considered. If we consider a steep spectrum, the intensity of the P$_*$ band becomes brighter and the signal to noise ratio becomes better by several times. This may result in better constraint on $RM_{\rm IGMF}$, because we obtain better quality of data at the P$_*$ band. We will address this effect more quantitatively in future, since Faraday tomography considering a non-zero spectral index is under development. The intensity ratio between the background and foreground sources does not significantly change the results, and exceptionally bright (Jy-level) background sources are available to this work of exploring the IGMF. Therefore, our method can be applicable for background, linearly-polarized FRBs. Meanwhile, detection of the Milky Way foreground would be more challenging. We have considered the signal-to-noise ratio of 10 for MW in each 1 MHz channel. If the noise level is higher, it seriously impacts on the detection of the Milky Way. Moreover, an interferometric observation may suffer from the missing flux of diffuse foreground emission. Although we introduced RFIs in Kashima as an example, our results do not depend on where and how the polarized intensity spectrum is obtained. Therefore, our results can be applicable to other current radio facilities and even the future telescopes such as the SKA. A possible recipe to confront this Milky Way foreground issue would be that we combine another single-dish observation of diffuse Milky Way foreground. Comparison between on-source and off-source observations is also useful, where the off-source observation measures a nearby sky sharing almost the same foreground. These follow-up observations confirm the diffuse foreground and decide the edge of the Faraday spectrum of the Milky Way. If we have two background sources located closely each other, we can apply another methodology, case (ii), discussed in \citet{aka14a}. In summary, we studied optimum frequencies to constrain Faraday rotation measure (RM) due to the IGMF by means of Faraday tomography. The frequency resolution of 1 MHz has been considered throughout this work. Using a simple model and Fisher information matrix, we find that multiple narrowband data in the UHF provides a reasonable constraint on the RM due to the IGMF. With data at 1400 MHz and 1600 MHz, RM$_{\rm IGMF} \sim 10$~${\rm rad~m^{-2}}$ toward a high Galactic latitude is detectable with less than 10~\% error, if we choose the center frequency of the P$_*$ band around 400 -- 700 MHz with a 40 MHz bandwidth. \vskip 12pt This work was supported in part by JSPS KAKENHI Grant Numbers JP15H05896 (KT), JP16H05999 (KT), JP16K13788 (T. Aoki), JP17K01110 (T. Akahori, KT), and Bilateral Joint Research Projects of JSPS (KT). Numerical computations were carried out on PC cluster at Center for Computational Astrophysics, National Astronomical Observatory of Japan. \appendix | 18 | 8 | 1808.10546 |
1808 | 1808.02165_arXiv.txt | { We develop a full four-dimensional numerical code to study scalar gravitational radiation emitted from binary systems and probe the Vainshtein mechanism in situations that break the static and spherical symmetry, relevant for binary pulsars as well as black holes and neutron stars binaries. The present study focuses on the cubic Galileon which arises as the decoupling limit of massive theories of gravity. Limitations associated with the numerical methods prevent us from reaching a physically realistic hierarchy of scales; nevertheless, within this context we observe the same power law scaling of the radiated power as previous analytic estimates, and confirm a strong suppression of the power emitted in the monopole and dipole as compared with quadrupole radiation. Following the trend to more physically realistic parameters, we confirm the suppression of the power emitted in scalar gravitational radiation and the recovery of General Relativity with good accuracy. This paves the way for future numerical work, probing more generic, physically relevant situations and sets of interactions that may exhibit the Vainshtein mechanism. } | \label{sec:intro} Understanding the physical origin of the observed accelerated expansion of the Universe has led to an explosion of theoretical dark energy and modified gravity models, which incorporate different types of screening mechanisms \cite{Khoury:2010xi,deRham:2012az,Jain:2013wgs,Brax:2013ida}. These screening mechanisms provide a means by which fields that are active at cosmological scales, potentially significantly modifying the behavior of the gravitational force, are hidden from solar system/astrophysical/lab tests of gravity. A large class of theoretical models rely on the Vainshtein screening mechanism \cite{Vainshtein:1972sx} (see \cite{Babichev:2013usa} for a review) which was originally proposed in the context of massive theories of gravity. The essential features of this screening mechanism are captured in the simpler context of Galileon theories \cite{Nicolis:2008in} and pioneering works on how the Vainshtein mechanism manifests itself in models of massive gravity were presented in \cite{Deffayet:2001uk,Dvali:2002vf,Lue:2003ky,Babichev:2009us,Babichev:2009jt,Babichev:2010jd}. Galileons, a class of scalar field theories that exhibit the nonlinearly realized Galileon symmetry $\pi \to \pi +v_\mu x^\mu +c$, arise naturally from many theories of dark energy like DGP \cite{Luty:2003vm,Nicolis:2004qq}, Ghost-free massive gravity \cite{deRham:2009rm,deRham:2010gu,deRham:2010ik,deRham:2010kj}, Bigravity models \cite{Hassan:2011zd,Fasiello:2013woa} and other theories of massive gravity \cite{Bergshoeff:2009hq,deRham:2011ca} or higher dimensional gravity \cite{deRham:2010eu}. The Galileon arises as the helicity zero mode of the hard or soft massive graviton, whose nonlinear interactions dominate over the usual helicity two interactions of General Relativity (GR). The Galileon effective theory thus describes the region in which the helicity zero mode may be nonlinear, while the helicity two modes are still in the weak field region. It is thus sufficient to work with this effective theory in order to describe systems that principally test weak field Einstein gravity, most notably the orbital decay of binary pulsars \cite{Hulse:1974eb,Taylor:1982zz,Taylor:1989sw}. Further, these scalar field theories are interesting in their own right as intermediate scale effective field theories \cite{Nicolis:2008in} and may admit many infrared (IR) completions (\eg the covariant Galileon \cite{Deffayet:2009wt} is a distinct IR completion from massive gravity \cite{deRham:2010ik,deRham:2010kj}, having the same decoupling limit). These theories automatically incorporate the Vainshtein mechanism in a way that is well understood for static sources \cite{Nicolis:2004qq,Nicolis:2008in}. However the Vainshtein mechanism is altogether less well understood in time-dependent systems. Binary pulsar systems are of particular importance since their orbital decay rate provides ones of the most precise tests of GR and its extensions \cite{Hulse:1974eb,Taylor:1982zz,Taylor:1989sw}. Binary black hole, neutron star systems, and possibly black hole-neutron star binaries \cite{Seymour:2018bce} are now equally important as sources of directly observed gravitational waves by Advanced Ligo/Virgo \cite{Abbott:2016blz,TheLIGOScientific:2017qsa}. In the case of binary pulsars, the sources may be treated as approximately non-relativistic since the orbital velocity is typically small. In this non-relativistic limit it was shown in \cite{deRham:2012fw}, that for the cubic Galileon, based on analytic approximations, that the screening of scalar gravitational radiation from binary pulsars is less effective than it is for static sources. This is due to the introduction of a new length scale associated with the dynamic time scale of the orbit $\op^{-1}$. In addition to the usual static Vainshtein suppression, there is an enhancement inversely proportional to the velocity of the orbiting system. Although \cite{deRham:2012fw} principally considered binary pulsars, similar expectations would hold if the sources were binary black holes or neutron stars --- at least in the region in which curvatures are small --- as the mechanism by which the Vainshtein mechanism works is largely only sensitive to the mass of the source and not on its precise nature. In the region of parameter space that is relevant for cosmological purposes, this screening is still strong enough to be below current constraints from binary pulsar systems. It has been suggested that adding the quartic or quintic Galileon terms, which more naturally arise in the context of (hard) massive gravity theories, may weaken the screening enough to be constraining \cite{deRham:2012fg}. But the perturbation theory that worked for the cubic Galileon failed for the quartic and quintic Galileon due to the approximations made \cite{deRham:2012fw,deRham:2012fg} (see also \cite{Berezhiani:2013dca} for a discussion of this point). Thus, to explore these systems we require the use of better analytic approximations or numerical methods. In this work we shall use a full three-dimensional, time stepping numerical code that bypasses the need of any analytic approximation or assumption, and is hence usable for any system irrespective of its symmetry (or absence thereof). The code, a heavily modified version of \emph{Grid and Bubble Evolver} (GABE) \url{http://cosmo.kenyon.edu/gabe.html} \cite{Child:2013ria}, uses a second order finite differencing scheme on a fixed Cartesian grid and integrates in time with an explicit second order (or optionally fourth order) Runge-Kutta method. The major modifications include adding a spherical harmonics power computation module, altered boundary and initial conditions, and changes to how the equations of motion are handled in order to deal with the non-linear equations of the Galileons (see \sect\ref{sec:numerics} for a discussion). Importantly, by using a Cartesian grid we are not making any assumptions about the symmetry of the system and thus solve the full non-linear equations exactly. This means that the results are an independent way of computing the power radiated by these binary systems from what was done in \cite{deRham:2012fw,deRham:2012fg}. As a first step we consider the case of the cubic Galileon (decoupled from gravity) coupled to a binary system, whose trace of the stress energy is simulated by a pair of orbiting localized Gaussians on Keplerian orbits. The situation naturally applies to binary pulsars, but can be easily modified to describe black hole binaries or neutron stars, in the regime where they are sufficiently far away that the local metric determined by the helicity two modes between the two black holes/neutron stars remains in the weak field limit\footnote{Although black holes are themselves nonlinear solutions, what is relevant is the contribution to the metric in the vicinity of one black hole generated by the other. For sufficient separations this coupling between the two black holes may be well approximated by a linear analysis, at least for the helicity two modes. Clearly when the black holes and neutron stars are close, \ie close to the merger regime, then these approximations will break down. However in this situation the Galileon decoupling limit approximation would not be valid.}. In the absence of the Vainshtein mechanism, the predicted scalar gravitational radiation would be of a comparable magnitude to the tensor radiation of General Relativity, something which would be immediately ruled out by constraints on binary pulsars. The Galileon interactions capture the nonlinearities which are expected to suppress the scalar radiation relative to the usual tensor contribution. The case of the cubic Galileon coupled to a binary system has previously been studied analytically in \cite{deRham:2012fw,Chu:2012kz}, and comparison with the analytic results allows us to probe the accuracy of the code. Alternatively, the success of the numerical code and its agreement with the analytic results allow us to confirm the validity of the analytic estimations performed in \cite{deRham:2012fw,deRham:2012fg}. Having confirmed the validity of the code, this now opens up the possibility to extend the analysis to other interactions that exhibit the Vainshtein mechanism (\eg~quartic and quintic Galileon interactions and other kinetic types of interactions as in k-mouflage \cite{Babichev:2009ee}) and to other physically relevant situations. An interesting feature of the analytic results performed in \cite{deRham:2012fw} is the realization that in the time-dependent system, the Vainshtein suppression for the total power radiated in the scalar (dominated by the quadrupole) goes as $(\op \bar r)^{-1}(\op\rv)^{-3/2}$, instead of the expectation of $(\bar r/ \rv)^{3/2}$ from static sources. This represents an actual enhancement going as $v^{-5/2}$, since for realistic systems like the Hulse-Taylor pulsar the orbital velocity $v\sim 10^{-3}$ \cite{Taylor:1989sw}. In the context of the Hulse-Taylor binary system, the overall Vainshtein suppression is still manifest $(\op \bar r)^{-1}(\op\rv)^{-3/2}\ll 1$ and the overall power emitted in the scalar cubic Galileon is negligible compared to that in GR. Therefore no observable effects would be detected for physical values of the parameters. This situation could change in the future as binary pulsars with different orbital frequencies/eccentricities are discovered and longer time measurements are made. When considering black hole mergers or other astrophysical binaries, we typically expect these systems to be even more relativistic and therefore the Vainshtein suppression to be more pronounced. As a proof of principle we may consider a very non-relativistic scenario $v=(\op \bar r)\ll 1$, so that even though one may be inside the Vainshtein radius (that is $(\bar r/ \rv)^{3/2} \ll 1$) the hierarchy of scales is such that the power emitted by a `Vainshtein-ly' screened binary system could be higher than what it would have been in the absence of Vainshtein screening. This means the hierarchy of scales involved for that to happen are not representative of any realistic physical system that we know, so at the moment this simply appears as a mathematical possibility which has not been realized in nature. Nevertheless it raises the question of whether the Vainshtein screening could in some context amplify the power emitted (see \cite{Ogawa:2018srw} for an interesting scenario where anti-screening has been identified numerically). A very natural concern is whether this enhancement of the power for sufficiently `slow' binary systems while remaining in the Vainshtein region could be an artifact of the assumptions employed in \cite{deRham:2012fw}, where a hierarchy of scales between the Vainshtein radius and the inverse frequency scale was assumed. Clearly as $\op \to 0$, there is no notion of power emitted and the system reduces to a static one where the Vainshtein suppression goes as $(\bar r/ \rv)^{3/2}$. In this work we therefore explore the validity of the analytic results \cite{deRham:2012fw} in a region where the hierarchy of scales may not be very strong $(\op ^{-1} \not\ll \rv)$. Based on the numerical results (that make no a priori assumptions on the scaling), we recover a scaling of the power emitted which is precisely in agreement with the analytic ones from perturbation theory, and does indeed manifest an enhancement of the power emitted while the source remains within the Vainshtein region. As the hierarchy of scales tends towards a more physically relevant regime, we observe a scaling of the power that remains again in perfect agreement with the analytic results and that suppresses the overall power emitted in the scalar gravitational radiation, providing a good handle on the Vainshtein screening from cubic Galileons in such systems. In \sect\ref{sec:Context} we briefly review how Galileons emerge from infrared models of gravity, and summarize the Vainshtein mechanism and the expected power emitted in the Cubic Galileon. We leave the details of the derivation for \app\ref{app:galrad}, where we first review the method for calculating the power from the effective action and then derive the analytic monopole, dipole, and quadrupole power emitted through scalar gravitational radiation in the cubic Galileon. We see that the power radiated is indeed dominated by the quadrupole, while the dipole is suppressed by at least 7 orders of magnitude (see \tbl\ref{table1}). \Sect\ref{sec:numerics} then focuses on the numerical work we performed with the Galileon theories. We enumerate the difficulties of the numerical approach due to the many conflicting dynamical scales and the non-linearities associated with the Galileons in binary systems. We continue on to discuss how the power calculated in the simulation differs from that computed from perturbation theory due to the lack of hierarchy between the size of the source, the inverse frequency scale, and the Vainshtein radius. We perform tests of the numerics with the well understood Klein-Gordon (free field and thus no Vainshtein mechanism) matching within 1\% of the well-known analytic results when the scales are well resolved by the numerics. We then continue on to work with the cubic Galileons where we expect a Vainshtein mechanism to affect the amount of power emitted (in a multipole dependent manner). Despite the lack of strong hierarchy between the size of the source, the inverse frequency scale, and the Vainshtein radius, the scaling of the power with angular velocity and Vainshtein radius is in good agreement with the analytic results summarized in \sect\ref{sec:Context}. Finally we summarize our results in \sect\ref{sec:conc}. We refer to \app\ref{app:rescale} for the rescaling used in the program. | \label{sec:conc} We have successfully performed full four dimensional simulations of a cubic Galileon coupled to a binary system on Keplerian orbits, and computed the resulting radiated scalar gravitational power. Our numerical results exhibit a power law dependence on the parameters $\op$ and the Vainshtein radius $\rv$, relative to the GR result of the form \begin{equation} \left.\frac{P^\text{cubic}_2}{P^\text{GR}_2}\right|_{\rm numeric} \propto (\op)^{-2.49 } (\rv)^{-1.44}. \end{equation} This is in very good argument with the scaling predicted by the perturbative analytic results derived in \cite{deRham:2012fw} (\ie \eqn\eqref{eqn:cu_power_rat}) assuming a strong hierarchy between $\bar r\ll\op^{-1}\ll\rv$. Even though the numerical simulations are performed with a relatively mild hierarchy, the consistency of the two lends strong support to the validity of the scaling implied by the analytic approximations performed in \cite{deRham:2012fw} over a much wider range of parameter space than is accessible numerically. At present the large hierarchies relevant to realistic binary pulsar systems for the cosmologically motivated choice of the scale $\Lambda$ are beyond the reach of our simulations, purely for reasons of storage and computational time. Nevertheless the agreement between the numerical and analytic results in the regime in which both are expected to be valid allows us to be confident in the analytic scaling. An important future step will be to apply these simulations to the quartic and quintic Galileon theories. These theories are of more direct relevance to hard massive gravity theories. In this case the approximations performed in \cite{deRham:2012fg} simply broke down, and so at present we do not have an analytic result to check or to make predictions. Analogous numerical simulations will fill this gap, and it is plausible that an improved approximate analytic treatment may have a better regime of validity. We will leave these considerations to future work.\\ \noindent{\textbf{Acknowledgments:}} We would like to thank Matthew Johnson, Luis Lehner, Andrew Matas, James Mertens, Leo Stein, and Jun Zhang for useful discussions. CdR and AJT would like to thank the Perimeter Institute for Theoretical Physics for hosting them during the final part of this work. The work of CdR and AJT is supported by an STFC grant ST/P000762/1. CdR thanks the Royal Society for support at ICL through a Wolfson Research Merit Award. CdR is supported by the European Union's Horizon 2020 Research Council grant 724659 MassiveCosmo ERC-2016-COG and by a Simons Foundation award ID 555326 under the Simons Foundation's Origins of the Universe initiative, `\textit{Cosmology Beyond Einstein's Theory}'. AJT thanks the Royal Society for support at ICL through a Wolfson Research Merit Award. JTG is supported by the National Science Foundation, Grant No. PHY-1719652. Numerical simulations were performed on equipment provided by the Kenyon Department of Physics and the National Science Foundation. This work also made use of the High Performance Computing Resource in the Core Facility for Advanced Research Computing at Case Western Reserve University. \begin{appendix} | 18 | 8 | 1808.02165 |
1808 | 1808.06166_arXiv.txt | We study a gas rich merging dwarf system KUG 0200-096. Deep optical image reveals an optically faint tail with a length of 20 kpc, giving a visual impression of tidal antenna similar to NGC 4038/39. The interacting dwarf galaxies have B-band absolute magnitudes -18.06 and -16.63 mag. We identify a young stellar clump of stellar mass of 2$\times$10$^{7}$ M$_{\sun}$ at the tip of the antenna, possibly a Tidal Dwarf Galaxy (TDG). The putative TDG candidate is quite blue with $g-r$ color index of -0.07 mag, whereas the interacting dwarf galaxies have $g-r$ color indices 0.29 and 0.19 mag. The TDG is currently forming star at the rate of 0.02 M$_{\sun}$/yr. We obtained HI 21 cm line data of KUG 0200-096 using the GMRT to get a more detailed view of neutral hydrogen (HI) emission in interacting dwarf galaxies and its TDG. Evidence of merger between the dwarf galaxy pair is also presence in HI kinematics and morphology where we find the HI contents of interacting pair is disturbed, forming a extended tail toward the TDG. The HI velocity field shows strong gradient along the HI tidal tail extension. We present a comparative study between the Antennae galaxy, NGC 4038/39, and KUG 0200-096 in both optical and HI gas properties and discuss possible origin of KUG 0200-096 TDG. | The Antennae (NGC 4038/39) is the text book example of a pair of merging disk galaxies \citep{Arp66,Struck99}. As such, they have been studied in detailed observation \citep{Amram92,Kunze96,Neff00,Gao01,Gordon01,Hibbard01,Whitmore05,Zhang10,Whitmore14} and reproduced in various numerical simulations \citep{Toomre72,Barnes88,Karl10,Teyssier10,Renaud15,Lahen18}. The system displays a prominent pair of tidal tails which extend a projected distance of 20' and the two merging disks are visibly distinct \citep{Schweizer78}. The latter has been assumed to be an indication of an early merger state, putting the system in the first place of the \cite{Toomre77} merger sequence. One important aspect of the Antennae system has been the discovery of ongoing formation of tidal dwarf galaxies at the tip of the antenna \citep{Mirabel92,Hibbard01}. Tidal Dwarf Galaxies (TDGs) are the most massive sub-structures born in gas-rich mergers. Their total mass is typical to that of dwarf galaxies (M$_{*}$ $<$ 10$^{9}$). Made out of tidal material (gas and stars) ejected from galaxies into the intergalactic medium, they are independent gravitationally bound systems, usually supported by rotation. Their dynamical status qualifies them as ``galaxies" \citep[see][for a review on TDGs]{Duc12} although being nascent galaxies they are not necessarily always in dynamical equilibrium \citep{Lelli15}. The importance of TDGs among the dwarf population is rather controversial. The idealized numerical simulations of galaxy-galaxy collisions made by \cite{Bournaud06} suggests that only a fraction of massive TDGs might survive longer enough to evolve as independent galaxies. Objects that are not kicked out from their parent's potential well are subjected to dynamical friction and gradual reduction of orbital energy plunges them into the host system where they could suffer destructive tidal force of the host galaxies \citep{Mayer01,Fleck03}, or can be destabilized by effect of ram pressure \citep{Smith13}. While massive galaxy interactions have been studied in great detail in the past, very little is known about the evolution of dwarf-dwarf interactions and mergers. There has been growing interest in the dwarf-dwarf interaction in recent literatures and in the last few years number of studies have presented observational evidences of the merging dwarf galaxies \citep{Paudel15, Stierwalt15,Pearson16,Delgado12}. In addition, many star-bursting dwarf galaxies show disturbed HI kinematics as a signature of tidal interactions \citep{Lelli14}. Studying formation TDGs in the dwarf-dwarf merging systems is important as the TDGs born out of colliding dwarf galaxies is expected to have different environment. Owing to the shallow potential well of host galaxies, the new born TDGs are subjected to a significantly low level of harsh tidal force from the parent. Low mass galaxies also have lower level of X-ray emission therefore less strong ram pressure stripping effect. This provides a higher survival probability of TDGs born out of dwarf galaxy collisions. \begin{table*} \caption{Photometric properties} \begin{tabular}{ccCc R cCcCccc} \hline Galaxy & RA & Dec & m$_{r}$ & g-r & m$_{FUV}$ & M$_{B}$ & M$_{*}$& SFR & M$_{HI}$ & z & V$_{r}$\\ & h:m:s & d:m:s & mag & mag & mag & mag & log(M$_{\sun}$) & log(M$_{\sun}$/yr)& log(M$_{\sun}$) & &km/s \\ \hline D1 & 02:02:38.77 & -09:22:13.2 & 15.76 & 0.29 & 18.06 & -18.06 & 9.27 & -0.68 & 9.3 & 0.018 & 5376 \\ D2 & 02:02:39.95 & -09:22:43.0 & 17.32 & 0.19 & 19.28 & -16.63 & 8.74 & -1.17 & \nodata &\nodata & 5355 \\ TDG & 02:02:39.14 & -09:23:23.4 & 20.36 & -0.07 & 20.62 & -13.93 & 7.27 & -1.70 & 8.1 & \nodata & 5441 \\ \hline \end{tabular} \label{phtab} \\ The magnitudes are corrected for galactic extinction and the B-band magnitude is obtained from the SDSS $g-$band magnitude using conversion formula provided by the SDSS. Typical error on the SDSS magnitudes are 0.01 mag. The stellar mass, given in column 8, is derived from the $r-$band luminosity and the mass-to-light ratio obtained from \cite{Bell03} for the color $g-r$ where we expect an conservative error on our estimate is 0.2 dex. We list star-formation rate in column 9, which is estimated from the FUV magnitudes. In column 10, we give HI mass, were the D1 value represents HI mass of entire the system. Radial velocity measured from HI kinematics is listed in the last column. \end{table*} In this paper we present an unique example of a low mass (M$_{*}$ $\approx$4$\times$10$^{9}$) Antennae system, KUG 0200-096, where we identify a TDG at the tip of tidal tail. Throughout the paper we adopt a luminosity distance of 68 Mpc (m - M = 34.1) and a scale of 0.32 kpc/arcsec valid for H0 = 71 km/s Mpc$^{-1}$. | We have presented the case of a merging dwarf galaxies pair in an isolated environment where we found a TDG is forming at the tip of tidal stream similar to the well known interacting system NGC 4038/39 (The Antennae galaxy). The TDG has a stellar mass M$_{*}$ = 1.9$\times$10$^{7}$ M$_{\sun}$ which is 0.5 percent of entire merging system. It is located at 22 kpc sky projected distance from main merging galaxy. It is blue with $g-r$ color index of -0.07 mag and gas rich with HI gas-to-light ratio of log(M$_{HI}$/L$_{B}$) = 0.4. \subsection{Tidal interaction and Star-formation} KUG 0200-096, no doubt, provides a great example of gas-rich merging dwarf galaxies. The interacting galaxies, D1 and D2 have B-band absolute magnitudes -18.06 and -16.63 mag which is similar to that of LMC/SMC pair, a well known interacting dwarf galaxies in our near vicinity. Several physical properties of KUG 0200-096, i.e., color, metal content, and SFR, are fairly similar to the typical BCDs and there is little doubt that its star formation activity is affected, if not triggered, by the interaction. LMC/SMC do not host star-forming region out-side of galactic main body, although it shows substantial extension of gas structure \citep{Onghia16} . However, in our previous publication \citep{Paudel17}, we have identified a TDG located in a stellar bridge between two interacting dwarf galaxies, which have similar star-formation and physical properties to LMC/SMC. We find that on average KUG 0200-096 has typical star-forming properties of BCDs. In Figure \ref{sfr}, we show a relation between star-formation rate and B-band magnitude for star-forming galaxies. The comparison sample is taken from \cite{Lee09}, who study star forming activity of local volume ($<$11 Mpc) star-forming galaxies where star-formation rates are also derived from FUV flux. We find no enhanced star-formation for overall B-band magnitude of the system compared to a sample of star-forming galaxies from the local volume. However, the TDG is located at the upper edge of scattered data. The gas mass fraction of the TDG is high compared to the overall gas mass fraction of the system but due to large beam size of GMRT observation, we can not rule out that the possibility of some contamination from the host galaxies. In any case, the value of gas to stellar mass ratio of the TDG, log(M$_{HI}$/M$_{*}$) = 0.83, is perfectly scaled with the relation of log(M$_{HI}$/M$_{*}$) and stellar mass, see \cite{Popping15} Figure 3. The elongated tails are a clear sign of tidal interaction of nearly equal-mass gas rich-disk galaxies. According to Toomre sequence \citep{Toomre72}, KUG 0200-096 is probably in early-stage of interaction. In comparison to the Antennae system, both interacting galaxies are clearly separated which may hint that the interaction in KUG 0200-096 is young compare to the interaction between NGC 4038 and NGC 4039. Numerical simulations of this interacting system could help to assess time scale of interaction and further morphological evolution of the TDG and merging remanent but they are beyond the scope of this paper. \begin{figure} \includegraphics[width=8cm]{sfr_ant.pdf} \caption{Relation between SFR and B-band absolute magnitude. The black dot represents overall KUG 0200-096 and blue dot for its TDG. NGC 4038/39 TDG is shown in red dot. The comparison sample, gray dots, is taken from \cite{Lee09}. } \label{sfr} \end{figure} \subsection{Formation of TDGs during dwarf-dwarf merger} Merging probability of low mass galaxies decreases in low redshift universe and as a result the chance of formation of tidal dwarf galaxies by merger of dwarf galaxies is also low \citep{Lucia06}. This makes low redshift dwarf-dwarf mergers an interesting phenomena to study. It has been found that dwarf-dwarf interactions are more likely to happen in isolated environments than in the groups or clusters \citep{Stierwalt15,Paudel18}. Given that dwarf galaxies located in isolated environment are gas rich, mostly Blue Compact Dwarf galaxy (BCD) type, it is not surprising that TDGs may form frequently in these gas rich dwarf-dwarf mergers. In fact, we also identified a new born TDG in our previous study of merging dwarf galaxies although they are in group environment \citep{Paudel15,Paudel17c}. KUG 0200-096 is noteworthy for the presence of a stellar stream hosting clump of star formation at the tip, and in that respect it resembles the system involving massive colliding galaxies, such as the Antennae (NGC 4038/39). In the study of numerous massive interacting galaxies, evidence of in situ star formation occurring in gas-rich collisional debris has been reported \citep{Mundell04,Mello08,Peterson09}. Such regions are also believed to be a nursery of super star clusters or TDGs \citep{Duc94,Duc07,Duc14,Paudel15,Paudel17}. These evidences of observation are also supported by idealized and cosmological numerical simulations of galaxies, where massive and compact super star clusters are seen forming in tidal tails \citep{Bournaud08,Renaud15,Ploeckinger18} and some of them may have evolved independently and survive against internal feedback and external tidal shear. The most massive and extended of them may become independent TDGs \citep{Wetzstein07}. The same phenomena seem to also occur in dwarf-dwarf major mergers. We find a blue stellar clump at the tip of stellar stream which hosts star forming region. It has a stellar mass of 1.9$\times$10$^{7}$ M$_{\sun}$. As its parent galaxies are low-mass systems with shallow potential wells, one may speculate that it will survive longer than in an environment of systems involving massive merging galaxies, e.g. NGC 4038/39. In comparison, NGC 4038/39 TDG is brighter than KUG 0200-096 TDG with a V-band absolute magnitude of -15.3 mag and has star-formation rate 0.03 M$_{\sun}$/yr \citep{Mirabel92}. Currently, KUG 0200-096 TDG is forming stars at the rate of 0.02 M$_{\sun}$/yr and both follow a the scaling relation of SFR and blue band absolute magnitude defined by normal galaxies, see Figure \ref{sfr}. \cite{Hibbard01} presented a high resolution HI mapping of NGC 4038/39 and its TDG Candidates. They found the HI morphology possess plenty of tidal features and substructures. The HI tail nearly follow the extension of both antennae observed in the optical. In KUG 0200-096, we find that gaseous tail does not overlap with the optical counterpart. However note that our spatial resolution is not sufficient enough to resolve the HI tail, to confirm whether it has a substructure. The TDG HI velocity is closer to that of D1 and a careful examination of Figure \ref{chmap} reveals that the extended HI emission towards the TDG actually originates in D1 (see higher velocity channel maps). It is possible that the HI tail actually emerges from D1, and that the stellar stream and the HI extension do not have the same point of origin, but are just projected on each other on sky. That can also explain the observed offset between the HI tail and the stellar stream and the lower relative line of sight velocity between D1 and the TDG. In this scenario, the TDG may be located at the end of the HI tail but the location of the TDG near the tip of stellar stream maybe a chance projection. Like the Antennae, KUG 0200-096, may be hosting two antennae -one is the gas poor stellar stream originated from D2 and another is HI tail originated from D1. In any case, to confirm this we certainly need to study a higher resolution and better signal to noise ratio HI map. Probably, most interesting difference between NGC 4038/39 TDG and KUG 0200-096 TDG is that the latter is significantly more compact compared to the former. NGC 4038/39 system is more massive and the TDG has a 15 kpc diameter whereas KUG 0200-096 TDG diameter is 2.5 kpc. \cite{Weilbacher18} identify multiple sub-clumps in the NGC 4038/39 TDG and detected multiple HII regions. In that sense, KUG 0200-096 TDG is morphologically more similar to BCDs, typically xBCDs \citep{Drinkwater91} and in contrast morphological properties of NGC 4038/39 TDG is comparable to a typical dwarf irregular galaxy (dIr). | 18 | 8 | 1808.06166 |
1808 | 1808.01746_arXiv.txt | Turbulence dissipation is an important process affecting the energy balance in molecular clouds, the birth place of stars. Previously, the rate of turbulence dissipation is often estimated with semi-analytic formulae from simulation. Recently we developed a data analysis technique called core-velocity-dispersion (CVD), which, for the first time, provides direct measurements of the turbulence dissipation rate in Taurus, a star forming cloud. {The thus measured dissipation rate of $(0.45\pm 0.05)\times 10^{33}~{\rm erg~s^{-1}}$ is similar to those from dimensional analysis and also consistent with the previous energy injection rate based on molecular outflows and bubbles.} | \label{sec:intro} In molecular clouds, turbulence is a ubiquitous process playing a crucial role in the star formation \citep{2004ARA&A..42..211E,2007ARA&A..45..565M}. {Although turbulence can generate high-density structures and thus enhance the effect of gravity in local and relatively small scales, it is generally treated as a pressure term, which counteracts gravity, retarding cloud cores from collapsing to form stars. In regions with strong apparent turbulence, such as those near the Galactic center, the star formation efficiency is clearly damped~\citep{2017A&A...603A..89K}.} Gas cores with comparable gravitational energy and turbulence energy can also form stars after the latter is dissipated if there is no continuous turbulence energy injection \citep{Gao2015}. Therefore, turbulence energy dissipation rate is a key parameter to determine the time scale of star formation. Turbulence energy can be injected by differential rotation of galactic disk \citep{1981ApJ...246L.151F}, galactic disk tidal force \citep{2015MNRAS.446..973F}, large-scale gravitational instabilities in galactic disks \citep{2003ApJ...590..271E,2010MNRAS.409.1088B}, stellar feedback \citep{2012ApJ...752..146L}, supernova explosions \citep{2005A&A...436..585D,2009ApJ...704..137J,2016ApJ...822...11P}, and fluctuations in Galactic synchrotron radiation \citep{2017MNRAS.466.2272H}. The injected energy will cascade down to small scales and dissipate through viscous processes (in this case at Kolmogorov scale) or low velocity shocks \citep{2012ApJ...748...25P}. The dissipation of turbulence energy evolves from viscosity dominated ($\mathcal{M}_{\rm s}\lesssim 1$) to shock dominated ($\mathcal{M}_{\rm s}\gtrsim 10$) regimes with increasing rms sonic Mach number $\mathcal{M}_{\rm s}$ \citep{2004PhRvL..92s1102P} \footnote{{In the weakly ionized interstellar media like molecular clouds, the dominant dissipation mechanism of MHD turbulence is ion-neutral collisional damping, i.e. ambipolar diffusion \citep{Lithwick2001,2016ApJ...833..215X}.}}. In a typical molecular cloud, such as Taurus, $1\lesssim \mathcal{M}_{\rm s}\lesssim 10$, so both dissipation mechanisms could take effect. The typical Kolmogorov scale in molecular clouds is estimated to be $10^{-5}\sim 10^{-4}$ pc, considering thermal viscosity \citep{2007ARA&A..45..565M,2011ApJ...737...13K}. There have been some unsuccessful attempts to probe these scales with the slope change of the turbulence energy spectrum \citep{2008ApJ...677.1151L}. It is conceivable, but rarely attempted, to measure the turbulence dissipation rate by observing the excess emission from gas, presumably excited by turbulence. Goldsmith et al.\ (2010) detected NIR H$_2$ emission across the boundary of the Taurus molecular cloud. Since H$_2$ transitions are hundreds of Kelvins above the ground state, such emission remains a mystery, given the much lower temperature and the lack of UV source in Taurus. One possibility is excitation by shocks. Recent studies suggest that the shock induced turbulence energy dissipation can be traced by mid-$J$ (e.g. $J$=6-5) CO lines \citep{2014MNRAS.445.1508P}. However, neither NIR H$_2$ emission nor mid-$J$ CO lines can be readily observed to trace viscous dissipation. The turbulence dissipation rate $\dot{E}_{\rm diss}$ in molecular clouds was also estimated with semi-analytical formulae based on numerical simulations \cite{Mac_Low1999}. Such estimate is in essence equivalent to dimensional analysis \citep{2007ARA&A..45..565M}. The key parameter in these analysis is the turbulence dissipation time, which is on the order of the turbulence crossing time. For Taurus clouds, these methods gave an estimate $\dot{E}_{\rm diss} \sim$ $0.7\times 10^{33}-3.8\times 10^{33}~{\rm erg~s^{-1}}$ \citep{Li2015,2012MNRAS.425.2641N}. There are other observation-based ways to estimate the turbulence energy dissipation rate, e.g., with structure functions \citep{1995tlan.book.....F}. The calculation of structure functions needs three-dimensional position and velocity, rarely available in astronomical observations. We have developed a new method, namely the core-velocity-dispersion (CVD) method, to study the cloud and turbulence structures \citep{Qian2012}. In the present work, we developed a CVD-based method to estimate the turbulence dissipation rate $\dot{E}_{\rm diss}$ in a thin and face-on cloud. As an example, the turbulence energy dissipation rate in Taurus molecular cloud was estimated. | \label{sec:discussion} We used molecular cores as an approximate and practical tracer of turbulent flow in molecular clouds. {Although the internal motion of the cores (smaller scales) could be dominated by compressive motions, we do not expect CVD to be of much bias in this regard, as only the collective properties (peak velocity of the whole core) are being used. } {There exit other statistical methods for obtaining structure functions (of Faraday rotation measure) in terms of the projected separation for both thin and thick clouds~\cite[e.g.][]{Lazarian2016}. These structure functions are calculated with the integrated value (along the line of sight) at each point on the projected plane. On the contrary, CVD is calculated based on collective characteristics of each core, thus having a much better localization property than existing methods. For example, when cores overlap with each other in projected 2D space, CVD can resolve them in the spectral dimension, which was demonstrated in~\cite{Qian2012}.} \QL{CVD, in the current incarnation, is limited by the lack of knowledge of separation along the line of sight. Previously, \cite{Qian2015} looked into the effect of using projected distance. The main obvious conclusion is the fact that as long as there is any correlation existing between CVD and the projected distance, the cloud cannot be thick. In a thick cloud, the motion of cores at different locations should have no dependence whatsoever on the projected distance between them. Indeed, we explored this projection effect and quantified the thickness of Taurus to be smaller than the 1/8 of the cloud transverse scale, i.e., Taurus is thin! The recipe presented in this paper utilizing CVD thus only works for a thin cloud. } \QL{As shown in Figure 2, if the thickness of the cloud is still in the inertial range, then the procedure described in the manuscript can still yield accurate results even if the projected distance is comparable to the thickness. However, when the thickness is large, i.e., larger than the large-eddy size, the proportionality between CVD$^2$ and $S_{\rm tt}$ given by Eq.~\ref{CVD} will be destroyed and the method cannot be used, unless detailed information along the line of sight is known.} The scaling laws in Eq.~\ref{sll} and \ref{stt} are rigorously correct only for incompressible turbulence. For compressible turbulence, no simple analytical relation exists. The fact that the relative motion between dense cores seem to follow the general cloud turbulent flow, i.e.\ CVD mimics the original Larson's law \cite{Larson1981} allows us to trust the CVD measurement to the degree that it does not deviate from incompressible turbulence by orders of magnitude. At small scales, the estimates of the structure function $S^2_{\rm tt}$ with CVD will be affected by the finite thickness $h$ of the cloud. The structure function at relatively larger scales $L\sim 5-10 \ \rm pc$ was used to estimate the turbulence energy dissipation rate (Fig. \ref{cvd}). In this sense the estimates of the turbulence energy dissipation rate in this paper does not depend on either the energy injection or the dissipation mechanism, but only relies on the scaling laws of turbulence energy cascade. This energy dissipation rate derived from the CVD method can be further used to estimate the turbulence decay rate/time in cloud cores \cite{Gao2015}. In this case how the decay of turbulence in cloud cores facilitates the star formation activity can be qualitatively studied. It is also interesting to compare the energy dissipation rate with the cloud cooling rate. The cooling rate per H$_2$ molecule is about $10^{-27}\ \rm erg/s$ for a volume density of $10^3\ \rm cm^{-3}$~\cite{1995ApJS..100..132N}. The total cooling rate for Taurus molecular cloud is then $\sim 9.0\times 10^{33}\ \rm erg/s$ for an average volume density of $10^3\ \rm cm^{-3}$. This cooling rate is higher than the turbulence dissipation. The transverse structure function $S^2_{\rm tt}$ can be estimated with core velocity dispersion (CVD) in a thin and face-on molecular cloud. The ratio $S^2_{\rm tt}/{\rm CVD}^2$ is found to be $2.0\pm 0.2$, based on fractional Brownian motion model. The measured turbulence energy dissipation rate of $(0.45\pm 0.05)\times 10^{33}~{\rm erg~s^{-1}}$ for scales between 5 and 10 pc matches previous observational estimates. Such a dissipation rate is also consistent with the energy injection rate from star formation feedback at relatively smaller scales between 0.05-0.5 pc \citep{Li2015}. An empirical picture of the turbulence in Taurus molecular cloud is that, the majority of energy injection happens at cloud complex scales ($>10$ pc), the energy then cascades through the intermediate scales down to clump scales while counterbalance the gravity in rough viral equilibrium. It finally reaches a dynamic balance with star formation feedback at small scales of clumps and cores. | 18 | 8 | 1808.01746 |
1808 | 1808.03433_arXiv.txt | We directly compare predictions of dwarf galaxy properties in a semi-analytic model (SAM) with those extracted from a high-resolution hydrodynamic simulation. We focus on galaxies with halo masses of $10^9<M_\mathrm{vir}/\mathrm{M}_\odot\lesssim10^{11}$ at high redshift ($z\ge5$). We find that, with the modifications previously proposed in \citet{qin2018}, including to suppress the halo mass and baryon fraction, as well as to modulate gas cooling and star formation efficiencies, the SAM can reproduce the cosmic evolution of galaxy properties predicted by the hydrodynamic simulation. These include the galaxy stellar mass function, total baryonic mass, star-forming gas mass and star formation rate at $z\sim5-11$. However, this agreement is only possible by reducing the star formation threshold relative to that suggested by local observations. Otherwise, too much star-forming gas is trapped in quenched dwarf galaxies. We further find that dwarf galaxies rapidly build up their star-forming reservoirs in the early universe ($z>10$), with the relevant time-scale becoming significantly longer towards lower redshifts. This indicates efficient accretion in cold mode in these low-mass objects at high redshift. Note that the improved SAM, which has been calibrated against hydrodynamic simulations, can provide more accurate predictions of high-redshift dwarf galaxy properties that are essential for reionization study. | Reionization refers to an important process after the Big Bang, during which the intergalactic medium (IGM) was transiting from neutral hydrogen to its ionized state \citep{Wyithe:2004kb}. According to the observed galaxy sample at high redshift \citep{Bouwens2014,Bouwens2015,Stefanon2016arXiv161109354S,Oesch2016ApJ...819..129O}, this process can only happen with ionizing photons from much fainter galaxies taken into account \citep{Robertson2013ApJ...768...71R,duffy2014low,Bouwens:2015hk,Liu2016}. Although there are still some debates on other possible sources that can dominate the high-redshift photon budget such as active galactic nuclei (AGN; \citealt{Giallongo2015A&A...578A..83G,Madau2015ApJ...813L...8M,Qin2017c,Hassan2018MNRAS.473..227H}), dwarf galaxies that are beyond our observational capabilities are generally thought to have driven the Epoch of Reionization (EoR). In this context, understanding the formation of these unobserved objects is crucial to studying the EoR and can only be probed at this stage with theoretical simulations. Hydrodynamic simulations evolve dark matter and baryonic particles simultaneously and provide direct insights into the relevant astrophysical process \citep{Vogelsberger2014,Schaye2014,Hopkins2014MNRAS.445..581H,feng2015bluetides}. However, resolving dwarf galaxies within a cosmological volume for reionization studies usually involves more than a few billions of particles, which remains computationally challenging at this stage. A more efficient method is to apply semi-analytic models (SAMs; \citealt{croton2006many,Somerville2008,guo2011dwarf,Henriques2015}) to \textit{N}-body simulations \citep{springel2005simulating,Boylan_Kolchin2009,Klypin2010,Garrison2017} that only consider collisionless particles. Using the halo properties inherited from the parent simulation, SAMs approximate baryonic physics such as gas accretion, star formation and feedback using simplified scaling relations. These relations are motivated directly from physical processes, or empirically from observational results and more complicated numerical techniques such as hydrodynamic simulations and radiative transfer calculations. The semi-analytic prescriptions that indirectly model galaxy formation introduce free parameters to describe efficiencies which are inevitably accompanied by parameter degeneracies \citep{Mutch2013,Henriques2013}. This can make their predictions sometimes controversial, and potentially disconnected from the \textit{true} behaviour in the universe. An alternative to validate SAMs in the absence of observations at high redshift is to compare their results against hydrodynamic calculations that start from identical cosmological initial conditions. The goal of this work is to capture emergent behaviours from the hydrodynamic simulations (e.g. large-scale mass removal by winds from supernova events in individual star-forming sites) and to improve the parametrised modelling in SAMs as so to replicate these processes. Under the assumption that hydrodynamic simulations, which model the details of galaxy formation in a more physically realistic manner, are a more natural description of the astrophysical phenomenon and hence more representative of real galaxies, we can explore the semi-analytic prescriptions for quantities that are, in practice, unobservables and potentially reveal improper assumptions or missing physics in SAMs. \citet{Guo2016} compared the \textsc{l-galaxies} \citep{Cole2000,Bower2006} and \textsc{galform} \citep{Springel2005,Henriques2015} SAMs with the \textit{EAGLE} hydrodynamic simulations \citep{Schaye2014}, and concluded that the models can reproduce the stellar mass function predicted by \textit{EAGLE}. However, discrepancies were also found in the efficiencies of stellar and AGN feedback as well as the prediction of stellar mass-metallicity and size relations. \citet{Mitchell:2017je} also used \textit{EAGLE} to assess \textsc{galform} and found the angular momentum as well as the baryon cycling might not be properly traced in the SAM, leading to inaccurate predictions of galaxy sizes. \citet{Stevens:2017fi}, on the other hand, investigated cooling of Milk Way-like galaxies in \textit{EAGLE} and addressed the necessity of updating the cooling prescription employed in most SAMs (see recent updates of the cooling model in \citealt{Hou2018mnras.475..543h,Hou2018arxiv180301923h}). We, in the previous paper \citep{qin2018}, also found that the cooling prescription needs revision for more accurate modelling of \textit{low-mass} galaxies at \textit{high redshift} and proposed an alternative modification to the current prescription, avoiding the introduction of a new model. \citet{Cote:2017uh} recently extended the comparison from cosmological simulations of smoothed particles to zoom-in simulations \citep{Bryan2014ApJS..211...19B} of a system with a total mass of ${\sim}10^9\mathrm{M}_\odot$ (one main halo and two satellites; \citealt{Wise2012MNRAS.427..311W,Wise2012ApJ...745...50W,Wise2014MNRAS.442.2560W}), and investigated the difference in dwarf galaxy formation between a SAM and hydrodynamic simulation. They found their SAM, which employs a different prescription of gas accretion, was successful in reproducing the hydrodynamic calculation of star formation history but with a prediction of a much narrower distribution of metallicity compared to the hydrodynamic result. This is the second paper following the work of \citet{qin2018}, where we investigated the performance of SAMs when applied to high-redshift dwarf galaxies. We used the {\meraxes} SAM \citep{Mutch2016a} as an example and focused on gas accretion, cooling and star formation with reionization and supernova feedback isolated. We compared the stellar and gas masses with a high-resolution hydrodynamic simulation from the {\smaug} suite \citep{duffy2010impact,duffy2014low}, and found that, in the SAM, \begin{enumerate} \item due to the lack of hydrostatic pressure in parent \textit{N}-body simulations, inheriting halo properties directly from the dark matter halo merger trees overestimates the total mass of haloes hosting dwarf galaxies; \item the assumption that, in the absence of feedback, haloes consists of a baryonic reservoir with a mass of $\Omega_\mathrm{b}/\Omega_\mathrm{m}$ of their total mass is not accurate for dwarf galaxy formation modelling and can lead to a significant overestimation of the total baryonic mass; \item star formation modelled by consuming the total gas disc in a few dynamical times of that disc cannot capture the evolutionary path of star formation implemented in hydrodynamic simulations; and \item gas accreted by dwarf galaxies is cold, the median temperature of which is significantly lower than the halo virial temperature, and the current cooling prescription is not representative of this process. \end{enumerate} Accordingly, we proposed modifications to SAMs, seeking for consistency with hydrodynamic simulations in calculations of the evolution of stellar and gas components of dwarf galaxies. In this work, we include these modifications as well as feedback from reionization and supernovae, and investigate whether the updated SAM can broadly agree with the hydrodynamic calculation of dwarf galaxies in the presence of feedback. We start with a brief review of the {\meraxes} SAM as well as the modifications proposed in \citet[hereafter \citetalias{qin2018}]{qin2018} and the {\smaug} hydrodynamic simulation suite in Section \ref{sec:models}. We then present and discuss our comparison results in Section \ref{sec:results}. Conclusions are given in Section \ref{sec:conclusion}. In this work, we adopt the Chabrier initial mass function (IMF; \citealt{chabrier2003galactic}) in the mass range of $0.1-120\mathrm{M}_\odot$ and cosmological parameters from \textit{WMAP7} ($\Omega_{\mathrm{m}}, \Omega_{\mathrm{b}}, \Omega_{\mathrm{\Lambda}}, h, \sigma_8, n_s $ = 0.275, 0.0458, 0.725, 0.702, 0.816, 0.968; \citealt{Komatsu2011}) in all simulations. | \label{sec:conclusion} Following \citet[][\citetalias{qin2018}]{qin2018}, we further investigate the semi-analytic modelling prescriptions of galaxy formation that are commonly adopted in the literature. In this work, we include supernova feedback and homogeneous reionization background in both the {\meraxes} SAM \citep{Mutch2016a} and \textit{Smaug} high-resolution hydrodynamic simulation \citep{duffy2014low}, and make comparisons between the stellar and gas reservoirs predicted by these two methods. We focus on galaxies with $10^9\mathrm{M}_\odot<M_\mathrm{vir}\lesssim 10^{11}\mathrm{M}_\odot$. With the modifications previously proposed in \citetalias{qin2018} including adjustments to halo masses from the merger trees, suppression of baryon fractions accounting for hydrostatic pressures, and the modulation of time-scales for the transition of gas from hot to star-forming ($t_\mathrm{transition}$) and from star-forming to stars (depletion time-scale, $t_\mathrm{sf({\htwo})}$), we find that the current SAM is able to reproduce the hydrodynamic calculation of the cosmic evolution of galaxies with $M_\mathrm{vir}>10^{10}\mathrm{M}_\odot$ at high redshift. This includes the stellar mass function, total baryonic mass, star-forming gas mass and SFR between $z=5-11$. However, in less massive galaxies ($10^9\mathrm{M}_\odot<M_\mathrm{vir}< 10^{10}\mathrm{M}_\odot$) with SFR calculated using the total star-forming gas, we identify a significant amount of star-forming gas stored in quenched galaxies due to the imposed mass threshold of star formation. After reducing the threshold, the SAM successfully mimics the evolution of dwarf galaxies in the hydrodynamic simulation. We also investigate a second star formation prescription, which splits the star-forming gas disc into molecular and atomic hydrogen and forms stars from molecules \citep{Lagos2011}. Fixing the depletion time-scale of $\htwo$ inferred from a previous study of the {\smaug} hydrodynamic simulation \citep{Duffy2017}, we find that, with only calibrations of the gas transition rate and supernova efficiencies, the SAM can also reproduce the dwarf galaxy properties calculated by the hydrodynamic simulation. In addition, we find that when reionization and supernova feedback are included, dwarf galaxies tend to accrete a significant amount of star-forming gas at early times ($z>10$), which quickly becomes suppressed towards lower redshifts. Future work needs to take this into account and incorporate modelling of cold-mode accretion to study dwarf galaxies in the early universe. | 18 | 8 | 1808.03433 |
1808 | 1808.06020_arXiv.txt | Pulsar-timing arrays (such as NANOGrav) will detect the ensemble gravitational-wave signal from many inspiraling supermassive black-hole binaries throughout the Universe within the next $3-7$ years \citep{2016ApJ...819L...6T}. The statistical properties of this signal will reflect the dynamical history of these supermassive black-holes as they evolve to form a bound system and reach milliparsec orbital separations. NANOGrav will constrain the environments of supermassive black-hole binaries \citep{tss17} through detailed studies of the gravitational-wave strain spectrum in the nanohertz band. The ngVLA has a goal resolution of $30$ mas, equivalent to $\sim 300$ pc at $z\sim 1$. This resolution is sufficient to distinguish dual AGN, allowing estimates of SMBH merger rates to be formed that can act as prior constraints for PTA analysis. Furthermore, if intercontinental VLBI capabilities are added to the ngVLA it will achieve $\sim$pc-scale resolution out to $z\sim0.1$, allowing a large measured sample of SMBH binary separations to inform whether binaries stall at $\sim$pc separations or are driven efficiently to the sub-pc regime by dynamical interactions. Intercontinental VLBA will also allow the influence radius of supermassive black-holes to be resolved, which will drastically improve mass estimation and thus the fidelity of derived $M_\mathrm{BH}-M_\mathrm{bulge}$ relationships. Finally, ngVLA observations of galactic gas reservoirs, large-scale inflows, and gas dynamics in central galactic regions will provide crucial insight into the physics of SMBH binary accretion processes. There are several current and future facilities that will complement the ngVLA in the aforementioned areas. These include the \textit{Atacama Large Millimeter/submillimeter Array} (ALMA)\footnote{http://www.eso.org/public/usa/teles-instr/alma}, whose angular resolution at $230$ GHz with baseline length of $15$ km is $15$ mas, which is comparable to the ngVLA's goal resolution of $30$ mas at $40$ GHz with a baseline length of $180$ km \citep{ngVLA8}. Likewise, the ngVLA should have similar angular resolution to adaptive-optics corrected images from the \textit{Giant Magellan Telescope} (GMT)\footnote{\href{https://www.gmto.org}{https://www.gmto.org}}, \textit{Thirty Meter Telescope} (TMT)\footnote{\href{https://www.tmt.org}{https://www.tmt.org}}, and the \textit{Extremely Large Telescope} (ELT)\footnote{https://www.eso.org/public/usa/teles-instr/elt}. This range of complementarity in angular resolution will allow for validation and comparison studies of sub-kiloparsec galaxy scales out to $z\sim 1$, which is sufficient to distinguish dual AGN for the purpose of estimating SMBH binary merger rates. On the other hand, the SKA\footnote{\href{https://www.skatelescope.org}{https://www.skatelescope.org}} is expected to operate concurrently with the ngVLA, but its point-source sensitivity will be only be half that of the ngVLA in their frequency-overlap band \citep{ngVLA8}. Similarly, JWST\footnote{\href{https://www.jwst.nasa.gov}{https://www.jwst.nasa.gov}} will operate at similar wavelengths to GMT/TMT/ELT, but have lower angular resolution than those instruments and the ngVLA. NANOGrav gravitational-wave analysis, in concert with observations by the ngVLA and complementary facilities, will paint a multi-messenger portrait of galaxy evolution and the dynamics of the most massive black-holes in the Universe. | 18 | 8 | 1808.06020 |
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1808 | 1808.10220_arXiv.txt | Structure formation in young protoplanetary disks is investigated using a one-dimensional model including the formation and the evolution of disks. Recent observations with ALMA found that a ring-hole structure may be formed in young protoplanetary disks, even when the disk is embedded in the envelope. We present a one-dimensional model for the formation of a protoplanetary disk from a molecular cloud core and its subsequent long-term evolution within a single framework. Such long-term evolution has not been explored by numerical simulations due to the limitation of computational power. In our model, we calculate the time evolution of the surface density of the gas and the dust with the wind mass loss and the radial drift of the dust in the disk. We find that the MHD disk wind is a viable mechanism for the formation of ring-hole structure in young disks. We perform a parameter study of our model and derive condition of the formation of ring-hole structures within $6\times 10^5$ years after the start of the collapse of the molecular cloud core. The final outcome of the disk shows five types of morphology and this can be understood by comparing the timescale of the viscous diffusion, the mass loss by MHD disk wind and the radial drift of the dust. We discuss the implication of the model for the WL~17 system, which is suspected to be an embedded, yet transitional, disk. | Protoplanetary disks are thought to be the birth place of planets. The density and temperature structures strongly affect the formation processes of planets. Recent observations reveal the detailed structures of protoplanetary disk. They have found that spiral structure \cite[e.g.][]{2012ApJ...748L..22M, 2013ApJ...762...48G, 2015A&A...578L...6B, 2016Sci...353.1519P}, non-axisymmetric structure \cite[e.g.][]{2013Sci...340.1199V,2013Natur.493..191C,2013PASJ...65L..14F, 2015PASJ...67..122M}, and ring-like structure \cite[e.g.][]{2007A&A...469L..35G,2010ApJ...725.1735I,2012ApJ...747..136I,2013ApJ...775...30I, 2016PhRvL.117y1101I, 2011ApJ...732...42A,2011ApJ...729L..17H,2012ApJ...758L..19H,2012ApJ...753...59M,2012ApJ...760L..26M, 2015ApJ...808L...3A,2016ApJ...820L..40A, 2016ApJ...829L..35T,2017arXiv171105185F,2017A&A...600A..72F,2017ApJ...840...23L} are formed in protoplanetary disks. The observed disk structures will give us some clues to reveal the planet formation scenario. Especially, the disk structure formed in the early evolutionary phase is important to understand the disk evolution process and the initial condition of the planet formation. Recently, \cite{2017ApJ...840L..12S} have found that the young protoplanetary disk around WL~17 in $\rho$ Ophiuchus star forming region exhibits a ring structure. The dust continuum emission from the disk shows a ring-like structure with a central hole. This is similar to the structure of transitional disks. However, one remarkable difference between standard transitional disks and the WL~17 system is the disk age. Observations of \cite{2017ApJ...840L..12S} suggest that WL~17 is still covered by the envelope. It is regarded as class I YSO \cite[]{2009A&A...498..167V,2009ApJ...692..973E}, whose age is suggested to be $\lesssim 0.5$~Myr \cite[]{2009ApJS..181..321E}, despite large uncertainty. WL~17 appears like transitional disks, which are one or two orders of magnitude older. The ring-hole structure formation mechanism has not been well studied for such young disks. One possibility is the gap formation by (an) unseen planet(s), but it may be very difficult to form planets at such an early stage. The observations of WL~17 suggest that there may be another mechanism that results in the ring structure formation at the early evolutionary phase. In this work, we investigate the early phase of disk formation and evolution to explore the possibility of forming small scale structures in young disks. The important processes for the early evolution of the disk are the gravitational collapse of the cloud core, the angular momentum transfer due to the gravitational instability within the disk, the growth and the radial drift of dust particles, and the effect of the magnetic fields. The disk formation through the gravitational collapse of the cloud core and the angular momentum transfer due to the gravitational instability are investigated well by using three-dimensional numerical simulations \cite[e.g.][]{1998ApJ...508L..95B,2007ApJ...670.1198M,2010ApJ...714L..58T,2010ApJ...718L..58I,2011MNRAS.416..591T}. The subsequent long-term evolution of the star-disk systems have been investigated separately from the star formation phase by assuming some initial disk model. This is partly due to the limitation of computational power of solving all the way from the initial star and disk formation from the molecular cloud core to the final dispersal. However, the evolution of young disks that we focus on in this work may strongly depend on the final outcome of the star-disk formation, which is the initial condition of the disk evolution. The pioneering theoretical work on the formation and evolution of protoplanetary disks is done by \cite{1981Icar...48..353C} \cite[see also][]{1983Icar...53...26C}. They have developed the model for the evolution of a one-dimensional viscous accretion disk including the effect of the infall from the cloud core onto the disk. Following their work, \cite{1994ApJ...431..341S} have done the detailed investigation on the trajectory of accreting gas onto the disk in the early stage of the star-disk formation. \cite{2005A&A...442..703H} have performed the model calculations with large parameter space of the viscosity, temperature and rotation rate of the cloud cores and compared the results with the observed disk structures. The model calculations show that the disk becomes gravitationally unstable in their formation phase \cite[]{1994ApJ...421..640N}. \cite{2010ApJ...713.1143Z} have investigated the disk evolution due to the angular momentum transfer caused by the gravitational instability and MRI using two-layered disk model. These previous studies used the mass accretion rate from the cloud core onto the disk obtained from the self-similar solution of singular isothermal sphere \cite[]{1977ApJ...214..488S}, which gives time-independent mass accretion rate. In \cite{2013ApJ...770...71T}, we have constructed the analytical model providing the time-dependent mass accretion rate from the cloud core with arbitrary radial density profile. The importance of dust growth and radial drift at the initial disk formation stage has been pointed out recently. Two-dimensional numerical simulations show that dust particles grow and drift inward resulting in the small dust-to-gas mass ratio during the disk formation \cite[]{2018arXiv180106898V}. The dust growth and the decrease of dust-to-gas ratio in the disk formation stage are also indicated by the steady accretion disk model \cite[]{2017ApJ...838..151T}. Such behavior of dust particles in the disk formation phase will affect the following disk evolution. To investigate the early evolution of disks, we need to deal with the gravitational collapse of the cloud core, the disk evolution caused by the angular momentum transfer, and the dust evolution in the disk comprehensively. In addition, we investigate the effects of the disk wind that can be driven by the magneto-rotational instability \cite[]{2009AIPC.1158..161S, 2010ApJ...718.1289S} on the formation of small scale structures in young disks. \cite{2010ApJ...718.1289S} investigated the mass loss rate due to the disk wind. They parameterized the mass loss by $\Cw$, where ${\dot \Sigma}_{\rm wind}=\Cw \Sigma \Omega$. According to the three-dimensional local MHD simulations, $\Cw\sim 10^{-5}-10^{-3}$. The timescale of wind mass loss scales with the local Kepler time of the disk, which is faster at inner radii. Therefore, if $\Cw$ is constant in the disk, the wind mass loss is efficient in the inner region and the inner hole is naturally formed by the wind. In such disks, it is expected that the dust concentrates around the inner edge of the disk and the ring-hole structure in dust distribution will be formed. The long-term evolution of the disks including the disk wind is investigated by using one-dimensional models treating isolated disks that have already been formed \cite[e.g.][]{2016A&A...596A..74S,2016A&A...596A..81P}. These work, however, did not calculate the disk formation phase. In this paper, we calculate the formation and the evolution of protoplanetary disks within a single framework. We extend the model provided in \cite{2013ApJ...770...71T}, which takes into account the time-dependent mass accretion, to include the effect of the dust and the wind mass loss on the disk evolution. We show that various gas and dust distributions can be formed by the MRI disk wind in young protopranetary disks. This paper is organized as follows. The model for the formation and the evolution of the disk is explained in Section \ref{method}. In Section \ref{result}, we show the result obtained from the model. Section \ref{discussion} and \ref{conclusion} are discussion and conclusion. | \label{conclusion} In this work, we investigate the ring structure formation in young disks by the wind mass loss caused by the MRI turbulence. To calculate the ring formation in young disks, we use a one-dimensional disk model that treats the formation and evolution of disks in a single framework. In this model, the strength of the turbulence, the mass loss rate by the disk wind, and dust size (Stokes number) are treated as parameters, and the dependence of the disk evolution on these parameters are investigated. Main results obtained in this work are summarized as follows: \begin{itemize} \item We find five types of disk structures as a result of disk formation and evolution model including the effects of wind mass loss: the ring structure, the dust gap disk, the filled disk, the dust poor disk, and the normal filled disk are obtained from our model calculations. \item The disk evolution is characterized by the timescales of the viscous diffusion $t_{\rm vis}$, the wind mass loss $t_{\rm wind}$, and the radial drift of the dust $t_{\rm drift}$. The formation of various disk structures can be understood by comparing these timescales. \item The relation between the timescales and the disk structure is summarized in Figure \ref{fig:type_table}. When $t_{\rm wind}<t_{\rm vis}$, the inner hole structure is formed in the gas disk. In this case, the dust can concentrate at the pressure maximum of the gas disk and the ring structure is formed when $t_{\rm drift} < t_{\rm wind}< t_{\rm vis}$. In the case $t_{\rm wind}< t_{\rm vis} < t_{\rm drift}$ or $t_{\rm wind}< t_{\rm drift} < t_{\rm vis}$ in the entire disk, the gas inner hole expands faster than the radial drift of the dust. As a result, the dust disk remains in the inner hole of the gas disk and the filled dust disk is formed. If the disk satisfies $t_{\rm wind}< t_{\rm vis} < t_{\rm drift}$ or $t_{\rm wind}< t_{\rm drift} < t_{\rm vis} $ only in the inner region and the outer region satisfies $t_{\rm drift} < t_{\rm wind}< t_{\rm vis}$, the dust gap disk is formed. When $t_{\rm vis}<t_{\rm wind}$, the inner hole structure of the gas disk is not formed. In this case, the dust ring structure is not formed either. If $t_{\rm drift}<t_{\rm vis}<t_{\rm wind}$, the dust in the disk drifts inward faster than the gas disk evolution. As the result, the dust poor disk is formed. On the other hand, the dust-to-gas mass ratio is almost same as the initial value (0.01) if $t_{\rm vis}<t_{\rm wind} < t_{\rm drift}$ or $t_{\rm vis}< t_{\rm drift}< t_{\rm wind}$. This case is categorized as the normal filled disk. \item When the ring structure is formed, the ring radius and the dust-to-gas mass ratio at the radius of the dust surface density maximum increase with time. When the dust-to-gas mass ratio becomes larger than about unity, the back reaction from the dust to the gas becomes efficient and the ring radius does not increase after that. \item To explain the ring structure observed around WL~17, $\alpha_{\rm MRI} \lesssim 10^{-3}$, $\Cw \sim 10^{-4}$ and St$\sim 0.1$ are suitable based on our model calculations. The resultant structures obtained by using these parameters are consistent with the ring radius, the dust surface density of the ring, formation of the inner hole in the dust disk, and mass loss rate estimated from observations. Our model calculations suggest that the dust-to-gas mass ratio at the ring is larger than 0.1 and the outer radius of the gas disk is larger than that of the dust ring. \end{itemize} In reality, the dust radius, strength of the turbulence, and the mass loss rate due to the wind are given by the result of the dust growth and the evolution of the magnetic field in the disk. The calculation of these values with the disk evolution is required to obtain more realistic structure of the disk, which will be the subject of our future work. We thank Takeru K. Suzuki for fruitful discussions and his valuable comments. This work was supported by NAOJ ALMA Scientific Research Grant Numbers 2016-02A. TM is supported by JSPS KAKENHI Grant Nos. 17H01103, 15H02074, and 26800106. \appendix | 18 | 8 | 1808.10220 |
1808 | 1808.07817_arXiv.txt | We present a new numerical model called CAFE-PakalMPI with the capability to solve the equations of classical magnetohydrodynamic (MHD) and to obtain the multispecies whose ionization states are calculated through statistical equilibrium, using the approximation of non-local thermodynamic equilibrium (NLTE) in three dimensions with the multiprocessor environment. For this, we couple the Newtonian CAFE MHD code introduced in \cite{CAFE_code} with PakalMPI presented in \cite{De_la_Luz_et_al_2010}. In this model, Newtonian CAFE solves the equations of ideal MHD under the effects of magnetic resistivity and heat transfer considering a fully ionized plasma. On the other hand, PakalMPI calculates the density of ionization states using the NLTE approximation for Hydrogen, electronic densities and H-, the other species are computed by the classical local thermodynamic equilibrium (LTE) approximation. The main purpose of the model focuses on the analysis of solar phenomena within the chromospheric region. As an application of the model, we study the stability of the C7 equilibrium atmospheric model with a constant magnetic field in a 3D environment. According to the results of the test, the C7 model remains stable in the low chromosphere, while in the range $z\in$[1.5,2.1] Mm we can observe the propagation of a wave produced by the changes in the ionization rate of H. | The solar atmosphere is the most dynamic layer of the Sun, here the fully or partially ionized plasma interacts with magnetic fields, this interaction can produce different kind of waves, jets, and other phenomena. In particular, the chromosphere-transition region is a zone of the solar atmosphere, where the plasma is not fully ionized. Here the plasma goes from nearly neutral to almost completely ionized, the temperature increases from a few thousand Kelvin to a million Kelvin, and the dynamics goes from being dominated by the gas pressure to being dominated by the magnetic pressure \citep{Golding_et_al_2016}. The most abundant element in the Sun is the Hydrogen, its ionization does not obey local thermodynamic equilibrium (LTE) or instantaneous statistical equilibrium because the timescale is long compared with important hydrodynamical timescales, especially of magneto-acoustic shocks \citep{Leenaarts_et_al_2007}. Usually, Hydrogen ionization is calculated from the condition of statistical equilibrium, that is, the equality of the ionization and recombination rates for the local temperature, density, and radiation. Statistical equilibrium assumes infinitely fast rates and an adjustment to the local thermodynamic and radiation state \citep{Carlsson&Stein_2002}. An account of the non-equilibrium ionization states involves solving a set of rate equations for the atomic population densities. The transition rate coefficients involve frequency integrals over the intensity, therefore a general description must necessarily take into account the radiative transfer, which in 1D is computationally doable, but in 2D or 3D a detailed treatment of radiative transfer is challenging. The reasons mentioned above have motivated the development of numerical tools capable to solve the full MHD equations coupled with the calculation of non-equilibrium ionization states including radiative processes mimicking the most realistic conditions in the solar atmosphere \citep{Hansteen_et_al_2007,Leenaarts_et_al_2007}. For this purpose a number of codes have been developed, for instance, the {\it Bifrost} code \citep{Gudiksen_et_al_2011}, designed to simulate stellar atmospheres from the convection zone to the corona, solves the MHD equations considering the effects of thermal conduction and full radiative transfer in the LTE and NLTE approximations. {\it MURaM} code \citep{Vogler_2004,Rempel_et_al_2009}, developed for applications in the solar convection zone and photosphere, solves the non-ideal MHD equations considering the non-local and non-grey radiative transfer and takes into accounts the effects of partial ionization. {\it CO$^{5}$BOLD} code \citep{Freytag_et_al_2012}, designed for realistic simulations that take into account the detailed microphysics of the solar or stellar surfaces layers, solves the hydrodynamic and MHD equations including gravity, radiative energy exchange and Hydrogen ionization. {\it ANTARES} code \citep{Muthsam_et_al_2010}, developed for simulations in stellar hydrodynamics with radiative transfer and realistic microphysics in 1D, 2D and 3D. Unlike the complexity of the codes mentioned above, in this paper, we present a model that do not solve the equations of full radiative MHD, neither it solves the rate equations for the atomic populations densities coupled with the MHD equations, instead, our model first solves the equations of classical MHD under the effects of magnetic resistivity and heat transfer of a fully ionized plasma in order to obtain post-processing the dynamics of multispecies whose ionization states are calculated through statistical equilibrium, using the LTE and NLTE approximations. The numerical model consists of two parts: i) Newtonian CAFE solves the MHD equations for a fully ionized gas to obtain the total mass density and temperature for each time-step, ii) PakalMPI uses the total mass density and temperature calculated with CAFE to obtain the ionization states using the NLTE approximation for Hydrogen, electronic densities and H-, the other species are computed by classical LTE approximation each time-step. The overview of Newtonian CAFE code, Pakal code and the coupling of the CAFE-PakalMPI model is described in Section \ref{CAFE_pakal_model}. The study of the stability of the C7 model and the convergence are presented in Section \ref{results}. In Section \ref{conclusions}, we present our conclusions and final comments. | \label{conclusions} We have developed a new numerical model capable to solve the MHD equations and the multispecies, whose ionization state densities are calculated by the NLTE approximation for HI, HII, H- and $n_e$ and the other species are computed by classical LTE approximation. The test of our code show the stability of the very complex process to obtain the densities for the species that compose the solar chromosphere. We prove that the numerical dispersion of our approximation is very low ($\approx 1\%$), the integration time of the MHD equations for the magnetohydrostatic system lasted approximately 5 days using 32 processors and the calculation of the ionization densities lasted a day and half using 32 processors. Our results show a hydrostatic chromospheric model, which is similar to previous results published in \cite{2016A&A...585A...4C}. The low region of the chromosphere remains unalterable but the transition region presents fast changes in temperature and density. The high difference in temperature between the corona and the chromosphere can not be sustained with the initial conditions of our model. We observe a fast thermalization in the high region of the chromosphere. At altitudes between 1.5 and 2.5 Mm over the photosphere the gradient of temperature increase with the time, producing a high rate of ionization at these altitudes. However, between 0 and 1.5 Mm the chromosphere is very stable for at least 300 s. The wave observed in the electronic density profiles are similar to the observed in the time-sequence of the intensities of submillimeter and millimeter continua in the solar chromosphere \citep{Loukitcheva_et_al_2004}. In such paper, the authors considered a collection of submillimeter and millimeter wave observed brightness temperatures $T_{b}$ of the quiet Sun and compared it with brightness temperatures computed from the standard static models of Fontella, Avrett and Loeser (FAL) \citep{Fontella_et_al_1993} and the dynamic simulations of Carlsson \& Stein \citep{Carlsson&Stein_1992,Carlsson&Stein_1995,Carlsson&Stein_1997,Carlsson&Stein_2002}. In particular, the dynamic simulation is based on a initial atmosphere in radiative equilibrium disturbed by waves driven trough the atmosphere by a subphotospheric piston. These waves propagate and increase in amplitude and form shocks above a height of 1 Mm. In contrast to the previous model, in our simulation we do not perturb the equilibrium state, instead, we let to evolve the system and find the appearance of disturbances in the electronic density that resemble oscillatory behavior. In the case of our model, this wave can be explained as the result of local changes in the rate of ionization of HII in NLTE. The thermalization process in the top of the chromosphere could play an important rol in local changes of electronic density at these altitudes. Finally, the model presented in this paper represents the initial step to develop a more general model capable to solve the full MHD equations coupled with the ionization state equations in NLTE and the calculation of the emission by solving the radiative transfer equation. | 18 | 8 | 1808.07817 |
1808 | 1808.07483_arXiv.txt | We present moderate (${\sim}5^{\prime\prime}$) and high angular resolution (${\sim}1^{\prime\prime}$) observations of $^{12}\rm{CO\,}(J=2-1)$ emission toward nearby, interacting galaxy NGC~3627 taken with the Submillimeter Array (SMA). These SMA mosaic maps of NGC~3627 reveal a prominent nuclear peak, inter-arm regions, and diffuse, extended emission in the spiral arms. A velocity gradient of ${\sim}400$--$450$~km~s$^{-1}$ is seen across the entire galaxy with velocity dispersions ranging from $\lesssim 80$~km~s$^{-1}$ toward the nuclear region to $\lesssim 15$~km~s$^{-1}$ in the spiral arms. We also detect unresolved $^{13}\rm{CO\,}(J=2-1)$ line emission toward the nuclear region, southern bar end, and in a relatively isolated clump in the southern portion of the galaxy, while no $\rm{C}^{18}O(J=2-1)$ line emission is detected at a $3\sigma$~rms noise level of 42~mJy~beam$^{-1}$ per 20~km~s$^{-1}$ channel. Using RADEX modeling with a large velocity gradient approximation, we derive kinetic temperatures ranging from ${\sim}5$--$10$~K (in the spiral arms) to ${\sim}25$~K (at the center) and H$_2$ number densities from ${\sim}$400--1000~cm$^{-3}$ (in the spiral arms) to ${\sim}$12500~cm$^{-3}$ (at the center). From this density modeling, we find a total H$_2$ mass of $9.6\times10^9~M_{\odot}$, which is ${\sim}50\%$ higher than previous estimates made using a constant H$_2$--CO conversion factor but is largely dependent on the assumed vertical distribution of the CO gas. With the exception of the nuclear region, we also identify a tentative correlation between star formation efficiency and kinetic temperature. We derive a galactic rotation curve, finding a peak velocity of ${\sim}207$~km~s$^{-1}$ and estimate a total dynamical mass of $4.94 \pm 0.70 \times 10^{10} M_{\odot}$ at a galactocentric radius of ${\sim}6.2$~kpc ($121^{\prime\prime}$). | \label{sec:intro} In addition to being fascinating in their own right, nearby galaxies are crucial to the understanding of galaxy evolution and interactions. They allow us to directly resolve the region around active nuclei and individual star-forming clouds over the full galactic disk. This comprehensive view allows us to connect the small-scale physics of the interstellar medium (ISM) and star formation to the disk-wide processes that drive galaxy evolution. It is also well known that interactions and mergers are an important process in galaxy evolution as shown by the increased merger rate in the early Universe \citep[e.g.,][]{Bridge10}. Gravitational interactions between galaxies significantly alter the morphology, luminosity, color, size, star formation rate, and mass distribution in a relatively short period of time. The NGC~3627 system is the natural choice for a representative nearby galaxy and interaction study because of its environmental richness and large suite of complementary data sets ranging from X-ray through radio wavelengths. Most of the previous extragalactic molecular gas studies used the $^{12}$CO($J=1-0$) line to trace molecular gas mass. With the recent abundance of $^{12}$CO($J=2-1$) observations, a quantitative understanding of the $^{12}$CO($J=2-1$) / $^{12}$CO($J=1-0$) line ratio (R$_{21/10}$) is required to compare results across the literature. Adding to the urgency of this, $^{12}$CO($J=2-1$) and $J=3-2$ lines are now regularly observed from z $\sim$ 1--3 galaxies \citep[e.g.,][]{Tacconi13, Tacconi17}, where they are used (with assumptions) to trace the total gas supply. A quantitative understanding of the variation in the R$_{21/10}$ line ratio is required to discern what is driving the changes in R$_{21/10}$ observed on large (galaxy) scales as well as to make rigorous statements about the behavior of molecular gas across galaxy populations. Additionally, observations of the full $^{12}$CO($J=2-1$) distribution in the very outer arms and inter-arm regions of a spiral galaxy such as NGC~3627 are particularly informative. Typical spiral galaxies do not contain conspicuous $^{12}$CO($J=2-1$) emission in the arm and inter-arm regions as observed in NGC~3627. Comparing this gas tracer with the large existing set of complementary data for NGC~3627 will provide a unique opportunity to study gas conditions in a very wide range of environments. NGC 3627\footnote{A summary of basic astronomical information is presented in Table \ref{tab:basic_info}.} (M66) is a spiral galaxy [RC3 type SAB(s)b; \citealt{Vaucouleurs91}] in the Leo Triplet galaxy group and displays signatures of a LINER/Seyfert 2-type nuclear activity in its spectrum \citep{Ho97, Peng98}. Optical broadband images of NGC~3627 reveal a weak optical bar, two prominent, asymmetric spiral arms, and large-scale dust lanes \citep{Arp66, Ptak06}. The perturbed morphology of the western arm, which appears to be displaced from the plane of the galaxy, provides evidence for recent interaction with neighboring galaxy NGC~3628 \citep[e.g.,][]{Rots78, Haynes79, Soida01}. The close proximity (${\sim}11\rm{\,Mpc}$) and high inclination (${\sim}61^{\circ}$) of NGC~3627 allow for an excellent view onto its spiral structure and pronounced dust patterns, making the galaxy an attractive candidate for investigating post-interaction galactic evolution. As a result, NGC 3627 has been studied in a wide range of continuum and spectroscopic observations -- in \ion{H}{1} \citep{Zhang93, Haan08}, CO (e.g., \citealt{Reuter96, Regan01, Helfer03, Kuno07, Leroy09, Warren10, Morokuma15, Beuther17, Donaire17, Cormier18}), H$\alpha$ \citep{Chemin03}, HCN / HCO$^+$ (\citealt{Krips08, Murphy15, 2017MNRAS.466...49J, Gallagher18}), UV \citep{Calzetti15}, and X-ray emission \citep{Georgantopoulos02, Soida12}. Multiple CO and radio continuum (327 MHz, 1.4 GHz, and 2.64 GHz) mapping observations \citep[e.g.,][]{Paladino08, Paladino09, Haan09, Nikiel13} have revealed that the majority of molecular gas in the galaxy is localized in a narrow bar structure ${\sim}300\rm{\,pc}$ in width with emission peaking at the nuclear position and extending along the leading edges of the bar, forming two broad peaks at the bulge ends and trailing off into the spiral arms. However, in atomic \ion{H}{1} emission, NGC 3627 appears to have a spiral structure free of any bar-like signatures \citep{Haan08, Walter08}. An inner ring (${\sim}30^{\prime \prime}$--$60^{\prime \prime}$) is also reported in $^{12}$CO($J=1-0$) \citep{Regan02} and H\textsc{$\alpha$} \citep{Chemin03} observations. An elongated inner ring along the north-south direction is also seen in GALEX far-ultraviolet (FUV) emission \citep[$\lambda_{\rm{eff}}=1516$ \AA;][]{Gil07} and surrounds a net depression in nuclear FUV emission with the $^{12}$CO bar-like structure being contained inside this FUV hole \citep{Casasola11}. Molecular transitions from HCN and HCO$^+$ have also been detected, indicating the presence of high density gas \citep{Gao04, Krips08, Murphy15}. Based on X-ray observations, \citet{Soida12} proposed a recent collision of NGC 3627 and a dwarf companion galaxy to explain some of the spiral arm distortions. NGC 3627 has a relatively high molecular gas fraction relative to other local star-forming galaxies \citep[e.g.,][]{Casasola04, Saintonge11} with molecular gas mass being comparable to atomic gas \citep{Helfer03, Walter08}. \citet{Zhang93} have suggested that the high H$_2$/\ion{H}{1} mass ratio of NGC 3627 is likely the result of tidal interaction and \ion{H}{1} stripping by companion galaxy NGC 3628. NGC 3627 exhibits X-ray characteristics reflective of a galaxy that has recently undergone a starburst \citep{Dahlem96}. Intense star formation activity has been observed in the nucleus and at both ends of the galactic bar \citep{Warren10}. Low levels of star formation have also been seen in H\textsc{$\alpha$} in the western arm, while the eastern arm contains a more vigorous star-forming region in its inner section \citep{Smith94}. Higher velocity dispersions have been measured in the southern bar end relative to the northern end \citep[e.g,][]{Zhang93, Chemin03, Dumke11}. The southern bar end also has been found to exhibit double line profiles in CO \citep{Beuther17} as well as harbor an unexplained magnetic field orientation, which does not follow the underlying optical spiral arm structure \citep{Soida01}. However, in a recent CO line analysis, \citet{Watanabe11} found that the star formation rates of the southern and northern bar ends are elevated relative to all other regions of NGC 3627 but are not substantially different than one another. \citet{Beuther17} conclude that the active star formation in the bar-arm interaction regions of NGC 3627 is the result of crossing gas orbits and colliding gas clouds, piling up dense gas that can then collapse and undergo intense star formation. \indent We present moderate (${\sim}5^{\prime \prime}$) and high angular resolution (${\sim}1^{\prime \prime}$) observations of $^{12}\rm{CO\,} (J = 2 - 1)$ and $^{13}\rm{CO\,} (J = 2 - 1)$ line emission toward NGC~3627 taken with the Submillimeter Array (SMA). These observations represent the most complete, in terms of overall spatial coverage and resolution, of $^{12}$CO ($J = 2 - 1$) emission toward NGC~3627. In Section \ref{sec:observations}, we present the SMA observations and describe the imaging process. We discuss CO morphology and kinematics and present detailed emission maps in Section \ref{sec:CO_Emission_section}. In Section \ref{sec:13CO_section}, we discuss the detection of the $^{13}\rm{CO\,} (J = 2 - 1)$ isotopologue and then we use RADEX modeling to derive physical properties of the molecular gas in Section \ref{sec:modeling_section}. We derive a galactic rotation curve and dynamical mass estimates in Section \ref{sec:rotation_curve_section} and summarize our results in Section \ref{sec:conclusions}. \begin{deluxetable}{lcccc} \tablecaption{Basic Astronomical Properties of NGC\,3627\label{tab:basic_info}} \tablehead{\colhead{Property} & \colhead{Value} & \colhead{Ref.}}\vspace{-0.3cm} \startdata R.A. (J2000) & $11^h20^m15.02^s$ & 1 \\[-0.1cm] Dec. (J2000) & $+12^{\circ}59^{\prime}29^{\prime \prime}.50$ & 1\\[-0.1cm] Classification & SAB(s)b & 2 \\[-0.1cm] Arm Class & Grand Design Spiral (7) & 3\\[-0.1cm] Nucleus & LINER/Sy2 & 4\\[-0.1cm] Distance (Mpc) & $10.57 \pm 0.73$ & 5 \\[-0.1cm] Linear Scale (pc arcsec$^{-1}$) & 51 & 5\\[-0.1cm] Inclination ($^{\circ}$) & $61.3$ & 1 \\[-0.1cm] Position Angle ($^{\circ}$) & $178.0 \pm 1$ & 1 \\[-0.1cm] V$_{\rm{hel}}$ ($\rm{km}\,\rm{s}^{-1}$) & $744$ & 1 \\[-0.1cm] Stellar Mass ($M_{\odot}$) & $10.23 \times 10^{10}$ & 6\\[-0.1cm] Orbital Mass ($M_{\odot}$) & $(1.45 \pm 0.39) \times 10^{12}$ & 6 \\[-0.1cm] H$_2$ Mass ($M_{\odot}$) & $9.6 \times 10^{9}$ & 7 \\ \enddata \tablecomments{(1) \citet{Casasola11}; (2) \citet{Vaucouleurs91}; (3) \citet{Elmegreen87}; (4) \citet{Ho97}, (5) \citet{Lee13}; (6) \citet{Karachentsev14}; (7) This work.} \end{deluxetable} | \label{sec:conclusions} We presented a $^{12}$CO($J=2-1$) emission map for NGC 3627. Based on CO observations of the interacting spiral galaxy NGC 3627 with the SMA, we conclude the following: \begin{itemize} \item We find enhanced emission and velocity dispersions in the nuclear region and bar ends with more diffuse emission and smaller dispersion in the spiral arms. We find a velocity gradient of ${\sim}400$--$450$ km s$^{-1}$ across the entire galaxy. \item We detected unresolved $^{13}$CO($J=2-1$) emission in the galactic center, southern bar end, and in an isolated clump of emission in the south of NGC 3627. Typical integrated-line intensity ratios of $^{12}$CO / $^{13}$CO are ${\sim}2.5$--$4$ with elevated ratios corresponding to regions with higher $\Sigma_{\rm{SFR}}$. No C$^{18}$O($J=2-1$) emission was detected in NGC~3627 down to a $3\sigma$ rms noise level of 42 mJy beam$^{-1}$ per 20~km~s$^{-1}$ channel. \item Using archival BIMA $^{12}$CO($J=1-0$), we produced a $R_{21/10}$ line ratio map for NGC~3627. Moderate ratio gas (${\sim}$0.4--0.6) was found in the nuclear region and the bar ends, while the spiral and inter-arm regions often exhibited substantially lower ratios ${\sim}$0.2. High ratio gas was also observed in the southernmost end of the extended western arm (${\gtrsim}$\,1.0) and in the clump (${\sim}$0.6--0.7), indicative of warm and dense molecular material likely due to previous tidal interaction with NGC~3628. \item Using the $J=2-1$ and $J=1-0$ transitions of $^{12}$CO, we produced a map of beam-averaged kinetic temperature and $n_{\rm{H}_2}$ at physical scales of ${\sim}$50 pc under non-LTE conditions. Kinetic temperatures ranged from ${\sim}5$--$10$ K in the spiral arms to ${\sim}25$ K in the nuclear region. A similar trend was found for $n_{\rm{H}_2}$ with values spanning ${\sim}$400--1000 cm$^{-3}$ to ${\sim}$12500 cm$^{-3}$. \item For all regions except the center, we find a tentative SFE-$\rm{T}_{\rm{K}}$ correlation and no correlation between SFE and $n_{\rm{H}_2}$. Since molecular gas density is believed to control spatial variations observed in SFE, the lack of this latter correlation is potentially surprising but with the important caveat that our RADEX analysis based on two $^{12}$CO lines, i.e. $J=2-1$ and $J=1-0$, may allow us to only probe a narrow range of gas volume densities. \item We derived a rotation curve for NGC 3627 out to a galactocentric radius of ${\sim}6.2$ kpc. Assuming an intermediate geometry between a flat disk and spherical distribution, we estimated dynamical mass as a function of galactocentric radius. By using the median velocity of ${\sim}192$ km s$^{-1}$ of the flat portion of the rotation curve at large galactocentric radius ($>3.5$ kpc), we report an upper limit of $M_{\rm{dyn}} = 5.75 \times 10^{10} M_{\odot}$ for the entire galaxy. \end{itemize} | 18 | 8 | 1808.07483 |
1808 | 1808.07951_arXiv.txt | {We present estimates of the signal to be expected in quiescent solar conditions, as would be obtained with the COronal Spectrographic Imager in the EUV in its coronagraphic mode (COSIE-C). COSIE-C has been proposed to routinely observe the relatively unexplored outer corona, where we know that many fundamental processes affecting both the lower corona and the solar wind are taking place. The COSIE-C spectral band, 186--205~\AA, is well-known as it has been observed with Hinode EIS. We present Hinode EIS observations that we obtained in 2007 out to 1.5~\rsun, to show that this spectral band in quiescent streamers is dominated by \ion{Fe}{xii} and \ion{Fe}{xi} and that the ionization temperature is nearly constant. To estimate the COSIE-C signal in the 1.5--3.1~\rsun\ region we use a model based on CHIANTI atomic data and SoHO UVCS observations in the \ion{Si}{xii} and \ion{Mg}{x} coronal lines of two quiescent 1996 streamers. We reproduce the observed EUV radiances with a simple density model, photospheric abundances, and a constant temperature of 1.4 MK. We show that other theoretical or semi-empirical models fail to reproduce the observations. We find that the coronal COSIE-C signal at 3~\rsun\ should be about 5 counts/s per 3.1\arcsec\ pixel in quiescent streamers. This is unprecedented and opens up a significant discovery space. We also briefly discuss stray light and the visibility of other solar features. In particular, we present UVCS observations of an active region streamer, indicating increased signal compared to the quiet Sun cases. } | It is now quite well established that the outer corona, i.e. the region between 1.5 and 3~\rsun\ (as measured from Sun center) is the place where many fundamental processes are taking place. For example, this is the region where the solar wind and the Coronal Mass Ejections (CME) become accelerated, see e.g. \cite{abbo_etal:2016, zhang_dere:2006, temmer_etal:2017}. This is also the region where CME driven waves steepen into shocks, and as the Alfven Mach number increases from 1, the post shock plasma in a CME can transition from low beta to high beta. The outer corona is also the region where the small-scale complex topology of the magnetic field, which shapes the observed features of the lower corona, becomes more simple and radial. It is also a transition region from collisional fluid to collisionless plasma. The outer corona could also contain, depending on which model one considers, the transition region from the low-$\beta$ collisional inner corona to the high-$\beta$ collisionless heliosphere and the nascent solar wind. For example, Fig.~6 of the MHD model of \cite{vasquez_etal:2003} shows that $\beta=1$ around 2\rsun\ in a streamer, while $\beta$ in a coronal hole is very small out to the edge of the plot at 4~\rsun. Other models put the $\beta=1$ further out, closer to the Alfven point. The outer corona is also the region where processes such interchange reconnection between closed and open structures is likely taking place, even in the quiet Sun \citep[see, e.g.][]{fisk:2003}. The complex mix of quiet Sun areas and coronal holes creates a complex topological system. MHD simulations of the quiet corona around the time of the 2008 eclipse run by the Predictive Science group with the MAS code \cite{mikic_etal:2007} showed the presence of a multitude of separators and quasi-separatrix layers in the outer corona, which have a profound effect on the magnetic connectivity between a point in the heliosphere and its source region, the so-called S-web \citep{antiochos_etal:2011}. When even a single active region is present, the topology of the global magnetic field of the outer corona becomes more complex. For example, magnetic field modeling of a few isolated active regions observed in 2007 showed the presence of null points at about 2~\rsun\ \citep{delzanna_etal:2011_outflows}. A model of interchange reconnection taking place at these null points was able to reproduce the location and strength \citep[see also][]{bradshaw_etal:11} of the so-called `coronal outflows', regions mostly connected to sunspots that show upflowing plasma at temperatures above 2~MK \citep[see, e.g.][]{delzanna:08_flows,harra_etal:08,doschek_etal:08}. This interchange reconnection process in active regions is in principle capable of injecting coronal plasma that was originally present in the closed hot (3 MK) loops of active region cores into the heliosphere, and forming as a by-product the large-scale cooler (1 MK) loops which fan out of sunspots and connect to the entire surroundings of an active region. It is clear that the density at 2~\rsun\ is so low that whatever occurs at 2~\rsun\ is virtually invisible in on-disk observations, against the bright background of the inner corona. Therefore, off-limb high-resolution observations around 2~\rsun\ are needed. Yet, this region is relatively unexplored, with the exception of a few ground-based observations during total eclipses, and SoHO/UVCS \citep{uvcs} spectra. Such observations returned plenty of scientific results, but were not capable of monitoring the dynamical evolution of the outer corona. Narrow-band high-resolution images of the solar corona in the visible taken during total eclipses have shown us the outer corona in its full glory, with many complex open and closed structures, and features that are difficult to comprehend. For example, intriguing features have appeared in recent images of the \ion{Fe}{xi} visible forbidden line \citep{habbal_etal:2011}. Such images of the coronal lines show more clearly the outer structures not only because of the natural occulter (the Moon), but also because the intensities of the visible forbidden lines decay more slowly with height, compared to the allowed lines. This occurs partly because forbidden line intensities decay linearly with the electron density, and partly because the photospheric radiation increases their intensities via photo-pumping \citep[see, e.g.][ for a recent discussion of visible and infrared forbidden lines]{delzanna_deluca:2017}. Unfortunately, these eclipse images are very short snapshots of a corona that we know is highly dynamic. Above 2~\rsun, coronagraphic observations in the visible with e.g. the SoHO Large Angle Spectroscopic COronagraph (LASCO) C2 \citep{brueckner_etal:1995}, and further out from STEREO HI, have provided valuable information about the dynamic outer corona. However, the spatial resolution and sensitivity have not been comparable with ground-based eclipse observations. Future coronagraphs in the visible (e.g. Proba-3, Aditya) will provide significant improvements, but such instruments will always be limited by the fact that the coronal signal is more than 6 orders of magnitude weaker than the photospheric one. The Solar Orbiter Metis coronagraph will be an improvement on the current situation as it will provide images both in the visible and in the UV, in the hydrogen Ly $\alpha$ above 1.6~\rsun. The Solar Orbiter EUI full-disk images are in principle also capable of observing the outer corona, but outside the currently planned remote-sensing windows (during close encounter). Around 2~\rsun, we had only a few images of the outer corona. LASCO/C1 observed the outer corona up to 3\rsun\ in the green and red visible forbidden lines, but returned only limited results \cite[see, e.g.][]{mierla_etal:2008}. We have plenty of observations from the X-rays to the UV of the lower corona in a range of spectral lines and broad-bands, typically up to 1.3~\rsun, but the nearly exponential decay of the electron density with radial distance, typical of stationary isothermal plasma, reduces dramatically the XUV signal, as the intensities of dipole-allowed transitions are proportional to the square of the electron density. There are a few off-points from SDO AIA (mainly for comet observations) where we see structures in the EUV out beyond 2\rsun. Recently, also SUVI obtained some off-limb images. We also have PROBA2/SWAP EUV images (in a band around 174~\AA\ dominated by \ion{Fe}{x} and \ion{Fe}{ix}) which clearly show a very complex and dynamical behavior of open and closed structures, especially above active regions. The dynamical behavior is mainly seen in the SWAP Carrington movies. The effects of an emerging AR are also clearly evident in AIA, see e.g. \cite{schrijver_higgins:2015}. To overcome the lack of detailed information about the outer corona, an instrument of novel design has recently been proposed to fly on the International Space Station: the COronal Spectrographic Imager in the EUV (COSIE). The instrument has two modes of operation. The spectrograph mode will mainly be used on-disk: it is a full Sun slitless imaging spectrograph (COSIE-S). The other one is a broad-band imager which will mainly be used in its coronagraphic mode (COSIE-C), although it can also observe on-disk, with a filter to reduce the bright on-disk signal. The key issue here is that building a coronagraph in the visible has always been a challenge, given the huge dynamic range between the disk intensity and the signal of the outer corona. This is not the case in the EUV, so an instrument such as COSIE can naturally observe both on-disk and off-limb. In a nutshell, COSIE-C is designed to observe the corona in an EUV spectral band between 186 and 205~\AA\ with a high-cadence, a large field of view (FOV) of 6.6\rsun\ $\times$ 6.6\rsun, and a spatial resolution of 3.1\arcsec. Such a resolution is really high, when compared to previous coronagraph images, and comparable to that of many current EUV instruments observing on-disk, such as Hinode EIS. This spectral range is well-known, as it has been routinely observed with many previous and current space-based instruments such as SoHO/CDS, SoHO/EIT, TRACE, STEREO/SECCHI, SDO/AIA, SDO/EVE, GOES SUVI. In particular, detailed spectroscopic observations with Hinode/EIS have been studied in detail, and over the past years all the main lines have been identified, and the atomic data has been recalculated and benchmarked, as reviewed in \cite{delzanna:12_atlas}. The quiet solar corona (at about 1 MK) is emitting in this spectral region strong coronal lines from iron ions: \ion{Fe}{x}, \ion{Fe}{xi}, and \ion{Fe}{xii}. The main aim of the coronagraphic mode of COSIE is to provide continuous monitoring of the outer corona with high-cadence and sensitivity, so dynamic events such as CME or waves in the outer corona can be studied in detail. In order to achieve this, a key technical issue regards the ability to observe simultaneously the inner corona from the solar limb to the outer corona. In principle the instrument is capable of observing up to 4.6~\rsun\ from Sun center, along the diagonal of its square FOV of 6.6$\times$6.6~\rsun, however as we have discussed the most interesting (and relatively unexplored) region of the solar corona is between 1.5 and 3.3~\rsun\, which is what will be observed within the square FOV of COSIE. The main aim of the present paper is therefore to provide observational evidence of how the brightness of the EUV corona changes with radial distance from the Sun in the 1.5--3~\rsun\ range, so we can provide relatively accurate estimates of the COSIE signal up to these distances. We mainly consider quiet Sun streamers in this paper as they are much less variable than e.g. those above active regions, or signals in other features. However, we provide one example of the signal above an active region streamer, and provide some comments of what signals could be expected for other regions. Our main aim might seem relatively simple, from an observational or modeling perspective, but it is not, as we discuss here. We know, from past observations, that the intensity of the disk in the EUV is comparable with the intensity of the inner corona as observed off-limb, but the behavior of the coronal lines up to about 3~\rsun\ was not known until the present study. Within the literature, we have not found direct measurements of coronal lines up to 3~\rsun\ in the quiet Sun. The only study of the outer corona is that by \cite{goryaev_etal:2014}, where however a streamer above an active region was observed. PROBA2/SWAP images in the 174~\AA\ band were analyzed up to 2~\rsun. The broad-band signal decreased by about 3 orders of magnitude in this range. Interestingly, a nearly isothermal corona of 1.4~MK was needed to explain the observations, as we have also found here. In principle, we could supplement the lack of observations with modeling. However, being an unexplored region, we really do not know how the fundamental plasma parameters vary with radial distance. Even considering the most simple case of a streamer in the quiet Sun, very different estimates of densities and temperatures have been published. A few are shown in Figure~\ref{fig:ne_te}, together with those we have adopted for our modeling, as described below. \begin{figure}[htbp] \centerline{\includegraphics[width=7.0cm, angle=0]{ne_te.ps}} \caption{The radial profiles of the electron densities and temperatures we assumed for modeling quiet Sun streamers (black lines), compared to some literature values.} \label{fig:ne_te} \end{figure} Such variable densities and temperatures produce very different estimates of coronal line radiances, as discussed in \cite{andretta_etal:2012} and as also shown here. Furthermore, an issue which greatly complicates any modeling effort is the variability of the elemental abundances, as we discuss below. An ideal instrument which explored the outer corona is SoHO/UVCS because it observed many coronal lines formed at similar temperatures as those that contribute to the COSIE-C band. There are, of course, many published results from SoHO/UVCS (and the previous similar SPARTAN instrument) which observed the outer corona from 1996 even as far as 10~\rsun. However, the instrument sensitivity was such that only the brightest lines, the H I Lyman $\alpha$ and the \ion{O}{vi} doublet, had enough signal in the outer corona. These lines are largely affected by photoexcitation as the density of the outer corona decreases, so they are visible to greater distances, but are not ideal to estimate the signal in the collisionally-dominated EUV lines, as we will discuss below. There are many published results from UVCS, where several coronal EUV lines were observed. However, they were mostly at distances of about 1.5~\rsun, see e.g. \cite{raymond_etal:97,parenti_etal:2000}. We have therefore searched the UVCS database to try and find some observations that could be useful for our purpose. We have identified a few observations of the Mg X and Si XII coronal lines in the 1.4--3~\rsun\ range, which we could use to build a model to estimate the COSIE Iron line count rates in the 186--205~\AA\ range. Several new and interesting results are obtained from this analysis. At lower heights, we present the analysis of a unique Hinode EIS off-limb observation where coronal lines (the same that would be observed by COSIE) were visible up to 1.5~\rsun. The paper is organized as follows: Section~2 describes the UVCS observations of quiescent streamers and their analysis. Section~3 describes the way we have modeled the UVCS observed radiances, and the predictions we make for the COSIE-C signal in the outer corona. Section~4 briefly discusses the EIS off-limb observation, while Section~5 briefly discusses stray light issue. Section~6 provides some comments on the expected COSIE-C signal in other regions and features, while Section~7 draws the conclusions. | An EUV coronagraph such as COSIE has a great potential for novel observations of the outer corona, especially between 1.5 and 3~\rsun, where many fundamental processes affecting the generation of the solar wind but also the evolution of structures in the low corona are taking place. This region is still nearly unexplored. We have a long heritage of coronagraph observations at visible wavelengths, which has notorious difficulties in reducing the disk light and in observing the corona close to the Sun. This is not an issue for COSIE in its coronagraph mode, which will allow, with the use of a filter, both on-disk and off-limb observations out to 3.3~\rsun. It might seem surprising, but we have shown here that significant signal in lines collisionally excited is still present at least until 3.1~\rsun, i.e. 2~\rsun\ above the photosphere. The few SOHO UVCS observations we were able to find clearly show this. We are confident about these results. We have also clearly shown that the same EUV lines of the COSIE instrument have been easily observed by Hinode EIS in a relatively quiet period in 2007 out to 1.5~\rsun. The main conclusion of the present study is that, on the basis of current baseline designs, COSIE-C exposures of the order of tens of seconds will be sufficient to observe quiescent streamers at 3~\rsun, if stray light issues will turn to be negligible. This is unprecedented (except for images during eclipses), when considered in combination with a spatial resolution of 3\arcsec. We have briefly discussed stray light issues. As there are discrepancies in assessing the stray light in current instruments, and stray light very much depends on the actual optical layout and components (e.g. filters), we defer an assessment to a future study, noting that our current knowledge on the basis of AIA and EIS performances indicates that COSIE-C, with improved filters and micro-roughness of the main optical surface, should have a negligible stray light. A secondary main conclusion is that variations in the outer corona of the electron densities and temperatures obtained from models fail to explain the observations. It is surprising, but the ionization temperature of the outer corona seems to be relatively constant out to 3.1~\rsun, at least on the basis of the few observations from UVCS and EIS presented here. That the temperature is nearly isothermal and constant with height was known from previous SOHO CDS and Hinode EIS measurements, but only up to 1.2~\rsun. The relative radiances between the coronal lines and the \ion{H}{i} Lyman $\beta$ observed by UVCS indicate photospheric abundances in the quiescent 1996 streamers, using the present model. The Lyman $\beta$ line is mostly collisional at 1.4~\rsun, so this result is largely independent on the photoexcitation and the density distribution of the plasma. It is interesting to note that this result is in agreement with a recent revision by \cite{delzanna_deluca:2017} of SOHO SUMER observations near the solar limb during the same period. There are however several complexities related to the interpretation and modelling of the UVCS Lyman and \ion{O}{vi} lines that are significantly photo-excited, and which could also affect these elemental abundance measurements at larger distances, so these abundance measurements will need to be confirmed. | 18 | 8 | 1808.07951 |
1808 | 1808.02573_arXiv.txt | Various models of structure formation can account for various aspects of the galaxy formation process on different scales, as well as for various observational features of structures. Thus, the investigation of galaxies orientation constitute a standard test of galaxies formation scenarios since observed variations in angular momentum represent fundamental constraints for any model of galaxy formation. We have improved the method of analysis of the alignment of galaxies in clusters. Now, the method allows to analyze both position angles of galaxy major axes and two angles describing the spatial orientation of galaxies. The distributions of analyzed angles were tested for isotropy by applying different statistical tests. For sample of analyzed clusters we have computed the mean values of analyzed statistics, checking whether they are the same as expected ones in the case of random distribution of analyzed angles. The detailed discussion of this method has been performed. We have shown how to proceed in many particular cases in order to improve the statistical reasoning when analyzing the distribution of the angles in the observational data. Separately, we have compared these new results with those obtained from numerical simulations. We show how powerful is our method on the example of galaxy orientation analysis in 247 Abell rich galaxy clusters. We have found that the orientations of galaxies in analyzed clusters are not random. It means that we genuinely confirmed an existence of the alignment of galaxies in rich Abells' galaxy clusters. This result is independent from the clusters of Bautz-Morgan types. | Solving the problem of the structures formation is one of the most significant issue of modern extragalactic astronomy. Many authors investigated the scenarios of structures formation since \citet{Peebles69,Zeldovich70}. New scenarios are mostly modifications and improvements of the older ones \citep{Lee00,Lee01,Lee02,Navarro04,MO05,Bower05,Trujillio06,Brook08,Paz08,Shandarin12,Codis12,Varela12,Giah14,Blaz15}. The final test of veracity in a given scenario is the convergence of its predictions with observations. One of the possibilities of such test is analysing the angular momenta of galaxies. Investigating the orientation of galaxy planes in space is of great importance since various scenarios of cosmic structures formation and evolution predict different distributions of galaxies angular momentum, \citep{Peebles69,Doroshkevich73,Shandarin74,Silk83,Catelan96,Li98,Lee00,Lee01,Lee02,Navarro04,Trujillio06,Zhang13}, i.e. provide distinct predictions concerning the orientation of objects at different levels of structure -- in particular clusters and superclusters of galaxies. Our model assumes that normals to the planes of galaxies are their rotational axes, which seems to be quite reasonable, at least for the spiral galaxies. Various models can account for various aspects of the galaxy formation process on different scales, as well as for various observational features of structures. This provides us with a method for testing scenarios of galaxy formation. In other words, the observed variations in angular momentum give us simple but fundamental test for different models of galaxy formation \citep{Rom12,Joachimi15,kiess15}. From the observational point of view it is not very difficult to investigate the distribution of the angular momenta for the luminous matter i.e. real galaxies and their structures. One should note however, that in real Universe, observed luminous matter of galaxies are surrounded by dark matter halos that are much more extended and massive. Direct observation of dark mater halos and theirs angular momentum is not so easy. Fortunately, there are the relation between luminous and dark mater (sub)structures. As a result we have a dependence between dark matter halos and luminous matter (real galaxies) orientation \citep{Trujillio06,Paz08,Pe08,Bett10,Paz11,Kim11,Varela12}. Recently, the analysis of the Horizon-AGN simulation shows the similar dependence \citep{Okabe18,Codis18}. It means that the analysis of angular momentum of luminous matter gives us also information about angular momentum of the total structure hence the analysis of the angular momentum of real (luminous) galaxies is still useful as a test of galaxy formation. The investigation of the galaxies orientation in clusters are also very important with regard to investigation of weak gravitational lensing For more detailed discussion of the significance of this problem see \citet{heav00,Hey04,kiess15,sg15,Codis16}. Since the angular momenta of galaxies and also the directions of galaxy spin are usually unknown, the orientations of galaxies are investigated instead. In order to acquire this, either the distributions of galaxy position angles only \citep{h4} or the distributions of the angles giving the orientation of galaxy planes \citep{Jaaniste78,f4} are examined. Many authors investigated the orientation of galaxies in different scales. The review of the observational results on the problem of galaxies orientation and structures formation was presented both theoretically \citep{Sch09} and observationally \citep{g2011a,Rom12}. One of the most meaningful aspects of the problem of the origin of galaxies involves the investigation of the orientation of galaxies in clusters. During the analysis of the angular momentum of a galaxy cluster, in principle we should take into account that total angular momentum of the cluster could come from both the angular momentum of each galaxy member and from the rotation of the cluster itself. However, one should note that there is no evidence for rotation of the groups and clusters of galaxies themselves. So, it is commonly agreed that such structures do not rotate (for example \citet{regos89,dia97,dia99,rines03,Hwang07,Tovm15}, see however \citet{kal05} for the opposite opinion). An especially important result is obtained by \citet{Hwang07}. They have examined the dispersions and velocity gradients in 899 Abell clusters and have found a possible evidence for rotation in only six of them. This allowed us to conclude that any non-zero angular momentum in groups and clusters of galaxies should arise only from possible alignment of galaxy spins. Moreover, the stronger alignment means the larger angular momentum of such structures. For many years, astrophysisicts payed a lot of attention to the orientation of galaxies in clusters. It was investigated both theoretically (see for example \citep{Ci94,Ci98}) and observationally. Generally, summarizing the research results provided by various authors, it can be stated that we have no satisfactory evidence for the alignment of galaxies in groups and poor clusters of galaxies, while there is an ample evidence of this kind for rich clusters of galaxies \citep{Godlowski05} (see also \citet{g2011a} for an improved analysis and \citet{sg15} for review). Thus, an interesting problem arises if there are any dependence on the alignment to the mass of the structure. \citet{Godlowski05} suggested that the alignment of galaxies in clusters should increase with the mass of the cluter. Thus, \citet{Godlowski05} hinted that the alignment should increase with the number of objects (richness) in a particular cluster, too. These suggestions were later confirmed by \citet{Aryal07}. These autors analyzed a total of 32 clusters of different richness and BM types. They confirmed that the alignment is changing with the richness and moreover that they change with BM type of the clusters. However, one should note that both \citet{Godlowski05} and \citet{Aryal07} investigations were qualitative only. The next step is to test this hypothesis also quantitatively. This was the reason why \citet{g10a} examined the orientation of galaxies in clusters both qualitatively and quantitatively. In this paper it was found that the alignment of galaxy orientation increased with numerousness of the cluster. However, the problem that we may obtain is whether we found a significant alignment in analyzed sample of 247 rich Abell clusters, or increasing alignment with cluster richness only. For this reason \citet{g2011b} analyzed the distribution of position angles using $\chi^2$ test, Fourier test and autocorrelation test as well as Kolmogorov test, showing that it is not random. In the present paper, following the ideas of \citet{g2011b} and \citet{Panko13}, we improved method allows us to analyze the distributions both of the position angles $p$ and distribution of two angles giving spatial orientation of galaxies. We denote $\delta_D$ angle (the polar angle between the normal to the galaxy plane and the main plane of the coordinate system) and the $\eta$ angle (the azimuth angle between the projection of this normal onto the main plane and the direction towards the zero initial meridian), see Figure \ref{fig:f0} for geometry of the angles. The main idea of our method is to analyze the distributions of these angles using statistical tests. We have analyzed in more details and improved the statistical tests used in \citep{g2011b} as well as introduce new statistical tests into the method. We analyzed how the tests changes if expected values of galaxies in particlular bins varies (as in the case of analysis of the $\delta_D$ angle). It slightly changes for autocorelation test but it is very important for Fourier test and Kolmogorov-Smirnov test. We also introduced to our improved method of investigation of galaxy alignment in clusters, the "control tests" that neglects a possible asymmetry of the distribution according to main coordinate plane. The idea of such tests is to analyze only the difference between more "parallel" or more "perpendicular" orientation according to the coordinate system main plane (or main direction towards the zero initial meridian in the case $\eta$ angle). We have checked how the Kolmogorov-Smirnov test behaves in the investigation of the orientation of galaxies in cluters. We have introduced alternative tests, namely Cr\'amer-von Mises and Watson that showed more explicitly that the allignment truly exists. Usually the effect of aligment of galaxies in structures is not very strong and its analysis requires precise statistical considerations. In such a case it is very important to verify that no other observational systematics can affect. To avoid a problem with the possible impact on the obtained results by data systematics, we think it is necessary to test the method on a well-tested sample of galaxy clusters. We have decided to use a sample of the galaxy clusters selected on the basis of the PF catalog \citep{Panko06}. Hence, on the example of analysis of position angles in 247 rich Abell clusters we show how the method works in case of observational data. For our sample of 247 clusters, we computed the mean values of the analyzed statistics. Our null hypothesis $H_0$ is that the mean value of the analyzed statistics is as expected in the case of random distribution of analyzed angles. At first, we have compared the theoretical prediction with the results obtained from numerical simulations. Later, they are compared with the results obtained from the real sample of the 247 Abell clusters. Separately, we analyzed the sample when only galaxies brighter than $m_3+3^m$ were considered. Moreover, we decided to analyze if there are any differences in alignment of galaxies in the clusters belonging to different Bautz-Morgan (BM) types. In order to exclude the case that the obtained results comes from errors in observational measurements, we have used two separate methods. We have analyzed the sample assuming random errors in position angles and additionaly we have used jackknife method especially to investigate the possible influence of background objects. The novelty of our approach is to gather many methods of analysis of statistics of all angles $p$, $\delta_D$ and $\eta$ not only for some particular galaxy clusters but also for big samples of clusters. Unfortunately, such approach inevitably turns to the analysis on a case by case basis. That is why in each case, we point out possible difficulties and show which method has to be used. At first glance, most cases looks very similar, but one has to be careful not to omit the crucial differences. The advantage of the approach is that by analysing much more data at once, we are able to draw more general conclusions. | The motivating theoretical goal of the project has been to give an improvement in the discrimination among different models of galaxy formation. A general idea has been to anylyze the angular momentum of galaxies in clusters and check if the results agree with scenarios predictions. That is why, in this paper we have focused on how we perform the analysis of the alignment of galaxies in clusters. In the original method presented in \citet{g2011b} the distributions of the position angles for galaxies in each cluster were analyzed using statistical tests: $\chi^2$ test, Fourier tests, Autocorrelation test and Kolmogorov test. The mean value of the analyzed statistics was compared with theoretical predictions as well as with results obtained from numerical simulations. The method allows to check if the mean value of analyzed statistics is the same as expected in the case of random distribution of the position angles of galaxies. In the present paper we have analyzed this method in detail, giving proposal of some significant improvements and introducing new statistical tests into the method. We have considered how the tests changes if we assume various expected values of galaxies in bins. In particular, in the autocorrelation test, the values of statistics slightly changes. However, in the Fourier test, not only the formulas for coefficients changes but also the coefficients need not be independent and we consider this in the analysis. We have also analyzed the properties of the Kolmogorov-Smirnov test applied to the analysis of the alignment of galaxies in clusters and finally we introduced control tests to all considered tests. In all cases the theoretical predictions have been compared with the numerical simulations. The second major advantage of the present paper in comparison to the previous investigation is that our analysis allowed us to expand the investigation of alignment of galaxies in cluster from the analysis of position angles only to the angles giving spatial orientation of galaxy planes, that has never been done before. The main difference is that the analysis of the position angles gives the information about orientation of galaxies only for edge-on galaxies. The analysis of the spatial orientation has allowed us to include all galaxies especially face-on galaxies. The difficulty that arises is that during the process of deprojection of the spatial orientation of galaxies from its optical images we obtain two possible orientations, and because we are able to find which solution is correct only in the small number of galaxies, both solutions must be taken into account during further analysis. Another crucial problem during analysis of the angles giving spatial orientation of galaxies is that if for any reason we exclude from analysis any type of galaxies (for example face-on galaxies), then the theoretical distribution of analyzed angles will be modified, even in the case when the distribution of galaxy planes is random and isotropic. In this case, a random distribution of analyzed angles which is the base of comparison with the real one must be, in practice, obtained from numerical simulations. This problem was analyzed for example by \citet{g2,Ar00} (for modern analysis see for example \citet{Flin11}). However, we have noticed that nobody took care of the fact that both obtained solutions for orientations are not independent of each other (only \citet{Panko13} made a remark that such problem could arise). Consequently, nobody analyzed if statistical test gives, even for "random" distributions, the same values of statistics as it is predicted in the case of the position angles. Our results have clearly showed that the expected mean values of the statistics for $\delta_D$ and $\eta$ angles varied from that obtained during analysis of the position angles. It means that "theoretical random distribution" must be modified this time. This is very easy to be observed in the example of nearly face-on galaxies, when both possible orientations are similar, which leads to the situation that both obtained values for $\delta_D$ and $\eta$ angles are similar. As a result in this case the first solution strongly affects the second. This phenomenon is also responsible for the results that during analysis of the spatial orientation of galaxies, we found the significant difference between the case when we assumed the real coordinates for galaxies in clusters and the case of the analysis of the fictitious clusters with coordinates distributed around the whole celestial sphere. In the paper we have analyzed this problem theoretically as well as show using numerical simulations how it affects the real data. In this paper we have analyzed the sample of 247 rich Abell clusters containing at least 100 members using a significantly improved method of the investigation of the orientation of galaxies in clusters. We found that the mean values of tested statistics, obtained on the base of the analyzed sample, significantly deviated from the expected in the case of the random distributions. As a result, we could conclude that the orientations of galaxies in analyzed clusters are not random. It means that we genuinely confirmed an existence of the alignment of galaxies in rich Abells' galaxy clusters, suggested by \citep{g2011b}, especially by results of Cr\'amer-von Mises and Watson tests. Moreover, we have shown that the above results are not due to errors in measurement of position angles nor influence of background objects (also done by jackknife method). The results that the aligmnent is increasing with richness and is observed in rich clusters supports the scenarios that predict such a thing (Li model, tidal torque scenario in the hierarchical clustering model). The other scenarios like Zeldovich pancakes \citep{Zeldovich70} and primordial turbulence \citep{Silk83} cannot explain such alignment and hence are not supported by our results. It was natural to expect that observed alignment could not be connected with equatorial plane. Indeed, the obtained values of $\Delta_{11}/\sigma(\Delta_{11})$ statistics does not show any deviation from zero, as predicted in such cases. This result was obtained both in the case of the analysis in Equatorial and Supergalactic coordinate systems, what means that observed alignment is also not connected with Local Supercluster plane. This result is generally independent from the clusters Bautz-Morgan types. Only cluster type BM II-III shows possible deviation from results obtained for other morphological types especially if we compare BM II-III with BM II type clusters. Our result clearly confirmed \citet{g10a} opinion that, contrary to the suggestions of \citet{Aryal04,Aryal05b,Aryal05c,Aryal06,Aryal07}, the alignment of the orientation of galaxies is only weakly correlated with their morphological types according to the classification of Bautz-Morgan (BM). One should note that in the paper \citet{Bier15}, during the analysis of the Binggeli effect for sample of 6188 galaxy clusters also selected from \citet{Panko06} catalogue, the differences was found with the Binggeli effect for BM type II clusters. It sugests that both of this observations could be connected with different morphological populations of the clusters i.e. the late type clusters (BM II-III and BM III) are spiral-rich clusters. It is important to note that the present observational results obtained for the sample of rich Abell clusters is based only on the analysis of the positions angles. It is due to the fact that we have no information connected with morphological type of members galaxies, hence the process of deprojection of the spatial orientation of galaxies from its optical images is a source of errors that are difficult to be controlled. Therefore, during the comparison of the real data with theoretical predictions and numerical simulations, we have concentrated on the analysis of the position angles only and postponed the spatial analysis of the real clusters to future studies. We should point out that our method of analysing the spatial angles is now well developed theoretically. In the future studies, we will investigate the real samples also with the analysis of the distribution of the angles $\delta_D$ and $\eta$ giving spatial orientation of galaxies. The future investigation will be possible with the use of information about frequency of galaxy occurrence in clusters with particular morphological types, since galaxy proportions with different spectral types can be estimated on the basis of density profiles in cluster \citet{Dre80,cal12,coe12,Hoy12}. We are also planning to extend our research to fewer galaxy clusters. Finally, we would like to conclude that now we have a well-tested method of studying the orientation of galaxies in clusters that can be used for research on other data sets, such as these from the new Kilo-Degree Survey. | 18 | 8 | 1808.02573 |
1808 | 1808.02090_arXiv.txt | Consider a planar three-body system of gravitating bodies: a central massive binary and a much less massive particle orbiting around the binary. Thus, the particle's orbit is circumbinary. The orbital eccentricities of circumbinary particles are periodically forced on secular and local orbital timescales \citep{MN04,PLT12,DS15,AIR17}. Therefore, the circumbinary orbits cannot be permanently circular. This phenomenon provides a natural universal mechanism of internal tidal friction and heating in circumbinary planets (CBP) \citep{S17AJ}. In this article, we describe and consider the planetary escape process that takes place as a result of this shrinkage. Indeed, a particle in the slowly shrinking circumbinary orbit enters eventually the chaotic zone around the central binary, and therefore escapes. Once in the zone, the particle escapes inevitably, being subject to the chaotic diffusion along the ``staircase'' of overlapping integer mean-motion resonances (between the binary and the particle), up to crossing the separatrix between the bound and unbound dynamical states \citep{S15}. We show that the effect of tidal decay may explain, at least partially, the observed lack of CBP of close-enough (with periods $< 5$~days) stellar binaries. What is more, on longer timescales (greater than the age of the Universe but well within stellar lifetimes), it may provide massive liberation of chemically evolved CBP. | We have shown that circumbinary planetary systems are subject to universal tidal decay (shrinkage of orbits), caused by the forced orbital eccentricity inherent to them. CBP are ejected from parent systems when they enter the circumbinary chaotic zone. On shorter timescales (less than the current age of the Universe), the proposed effect may explain, at least partially, the observed lack of CBP of close-enough (with periods $< 5$~days) stellar binaries. On longer timescales (greater than the age of the Universe but well within stellar lifetimes), it may provide massive liberation of chemically evolved CBP. It should be underlined that the phenomenon of tidal decay of circumbinary systems is universal. Here we have considered planetary systems, but in fact it is present in any circumbinary system of gravitating bodies. While this paper was in the reviewing process, a preprint was posted and an article published \citep{FBG18}, where the lack of close isolated binaries with CBP was also explained by the destabilization of CBP orbits due to their entering the circumbinary chaotic zone. However, the mechanism of the entry is different from that considered above. In our scenario, the planetary orbit slowly shrinks, while the chaotic zone stays constant in size. In the tidal scenario of \cite{FBG18}, in contrast, the chaotic zone swells (as the binary's orbit widens, due to the tidal transfer of angular momentum from the stellar rotation), while the planetary orbit stays constant in size. It should be noted that the tidal scenario of \cite{FBG18} is relevant to an early (pre-main-sequence) stage of evolution of the host star (see Section~\ref{sec_ep}), whereas our scenario acts on much longer timescales, and is thus capable of providing the escape of chemically evolved planets. | 18 | 8 | 1808.02090 |
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1808 | 1808.07497_arXiv.txt | The InfraRed Imaging Spectrograph (IRIS) is a first-light instrument for the Thirty Meter Telescope (TMT) that will be used to sample the corrected adaptive optics field by NFIRAOS with a near-infrared (0.8 - 2.4 $\mu$m) imaging camera and Integral Field Spectrograph (IFS). In order to understand the science case specifications of the IRIS instrument, we use the IRIS data simulator to characterize photometric precision and accuracy of the IRIS imager. We present the results of investigation into the effects of potential ghosting in the IRIS optical design. Each source in the IRIS imager field of view results in ghost images on the detector from IRIS's wedge filters, entrance window, and Atmospheric Dispersion Corrector (ADC) prism. We incorporated each of these ghosts into the IRIS simulator by simulating an appropriate magnitude point source at a specified pixel distance, and for the case of the extended ghosts redistributing flux evenly over the area specified by IRIS's optical design. We simulate the ghosting impact on the photometric capabilities, and found that ghosts generally contribute negligible effects on the flux counts for point sources except for extreme cases where ghosts coalign with a star of $\Delta$m$>$2 fainter than the ghost source. Lastly, we explore the photometric precision and accuracy for single sources and crowded field photometry on the IRIS imager. | \label{sec:intro} % With three giant segmented-mirror telescopes (GSMTs) on the horizon new scientific opportunities will be available to astronomers by utilizing unprecedented spatial resolution and photometric capabilities. The Thirty Meter Telescope (TMT) instrumentation aims to sample the diffraction-limit of a 30-m aperture with high precision relative photometric accuracy. The InfraRed Imaging Spectrograph (IRIS)\cite{Larkin1,Larkin2,Larkin3,Larkin4} is a first-light instrument for TMT which will enable the use of TMT for near-infrared (0.8 - 2.4 $\mu$m) imaging and integral-field spectroscopy (IFS). The imager operates with a plate scale of 0.004\arcsec per pixel with a maximum field of view of 34\arcsec x 34\arcsec \cite{imager1,imager2}. IRIS imager and integral field spectrograph will utilize a real-time data reduction pipeline\cite{Walth2} for data processing. Characterizing and predicting IRIS's photometric precision and accuracy is imperative for aiding the design of the instrument, adaptive optics system, and the data reduction system\cite{Walth1}. It is important to understand potential "ghost images" (i.e., additional point or resolved images generated by a multi-layer optical system) introduced by the optical design of IRIS that can directly impact the astrometric and photometric accurary. The photometric budget is defined by the combination of telescope optics, Narrow Field InfraRed Adaptive Optics System (NFIRAOS)\cite{adc1} \cite{adc2}, and IRIS optical system. Each of these optical systems have the potential capability to redistribute flux on the detector field of view (FoV), causing ghost images from science sources. The observed source magnitude and its optical location determines the brightness and location of the ghosts that are observed by the detector. It is critical to understand the different ghost impacts to properly characterize scenarios that may cause ghost images to adversely affect IRIS science cases. The existence of ghost images also increases the overall noise seen within the detector FoV. Ghost images therefore have a direct impact on the total photometric precision achievable by IRIS, and increases the base noise-floor of the imager and IFS. It is therefore necessary to analyze the overall impact of ghosts on photometric precision and noise-floor contribution in order to accurately define IRIS's photometry budget. We use the IRIS data simulator\cite{Wright1,Wright2,Wright3} to analyze the photometric effects of ghost images on the IRIS imager. Using the current specifications for IRIS's optical design and filter characteristics, we simulate the effects of the optical system, IRIS throughputs, sky background, ghosts, and Poisson noise. We simulated point sources of magnitudes 1-25 (Vega), with 25 iterations per simulation, including and not including the effects of IRIS in order to assess the photometric error for each simulated source using aperture photometry and Point Spread Function (PSF) fitting routine \textit{Starfinder}\cite{starfinder}. \textit{Starfinder} returns photometric results by extracting a PSF from the image, and fitting that PSF to each source in the field to estimate flux values for each source. This paper reports on the results of the simulations done, analyzing the specific ghost image photometry for both isolated ghost images and point source, binary, and crowded field science cases with the IRIS imager. | We have presented the results of ghost image analysis by simulating point sources with the IRIS imager and calculating the associated photometric error using PSF-fitting and aperture photometry. We explored the photometric impact of individual ghost image types (wedge filter, ADC, entrance window) from the IRIS optical design, and find that ghosts from the wedge-shaped filters have the highest associated photometric error. We investigated various multiple-point-source scenarios where ghost images could adversely affect photometry, and find that ghosts generally contriburte negligible effects (less than 1\% photometric error) except in extreme cases where ghosts coalign with a star of $\Delta$m$>$2 fainter than the ghost source. We report the impact on IRIS's photometric accuracy from ghosting to to be 0.292$\pm$0.005\% for a single point source. We report preliminary values for IRIS's photometric precision for a single point source at varying magnitudes and integration time, with a potential for $<$1$\%$ precision at $m_Y$=25 and $\sim$ 3\% at $m_K$=25, and anticipate further development of IRIS's predicted photometric specifications in parallel with IRIS's advancing design. We plan further development of these simulations with the addition of systematic errors and PSF variability, which will better predict the photometric performance of IRIS. | 18 | 8 | 1808.07497 |
1808 | 1808.02798_arXiv.txt | We develop a cosmological model based on action-dependent Lagrangian theories. The main feature here is the nonconservation of the energy momentum tensor due to the nontrivial geometrical construction of the theory. We provide the basic set of equations necessary to study both the cosmological background expansion as well as the linear matter perturbation growth. We show that the simplest realization of the Universe as described by only one component is not viable as expected from the existing correspondence between this model and the case of viscous cosmological fluids. However, modeling the energy content of the Universe as composed by two pressureless fluids, i.e., one a typical cold dark matter fluid and the other a pressureless “dark energy” fluid which is responsible for driving the late-time acceleration expansion, is qualitatively compatible with observational data. | The real world is pervaded by all sorts of dissipative processes; what imposes a suitable description of these phenomena on any consistent physical process is its respective scope. Nonetheless, this issue is usually left outside the standard variational formulations. In the classical mechanics context, one usually approaches it by using the so-called Rayleigh dissipation function, which is a useful tool to deal with dissipative forces with linear dependence in the velocity \cite{rayleigh,goldstein,whittaker}. Although this method provides a quite handy procedure for the description of such friction forces, the Rayleigh’s function method pos- sesses clear limitations. For instance, it fails in addressing a broader class of dissipating cases existing in nature with more general dependencies upon the velocity and the history of the system. Besides, this function arises as a correction in the Euler-Lagrange equation which does not affect at all the underlying variational formalism. In fact, it was demonstrated in \cite{bauer} that the Rayleigh's function is prohibited from emerging from a variational principle, unless the dissipative coefficient is not a constant anymore. These limitations are significant and point towards the search for possible extensions. Many attempts at incorporating dissipation effects into the traditional principle of least action were made over the last century. They basically rely on the use of time-dependent Lagrangians \cite{timeD}, auxiliary coordinates that describe the reverse-time system \cite{Morse}, or a fractional derivatives formalism \cite{fracD,fracD1}. However, these proposals face serious conceptual (or operational) obstacles which can undermine them as feasible alternatives, as they can either plague the theory with nonphysical Lagrangians or give rise to nonlocal differential operators, whose implementation introduces an undesirable complexity to the study of some problems. Another noteworthy alternative dates back to the 1930s, where G. Herglotz presented an elegant variational treatment to this issue by assuming action-dependent Lagrangians in the context of classical mechanics. In his approach, Herglotz was indeed successful in describing the class of dissipative systems whose motion is damped by a friction force, characterized by the aforementioned term proportional to velocity \cite{herglotz}. Most important, the Herglotz variational formulation is free from the conceptual and practical obstacles found in the others approaches. In recent works, one of us, in partnership with some colleagues, extended the original Herglotz formalism to a covariant language \cite{lazo,lazo2}. As it was shown in the work \cite{lazo}, such a covariant generalization laid the cornestone for the construction of a new theory of gravity in which dissipative effects would be a natural consequence, coming from first principles and having a purely geometric origin. In this vein, the authors derived explicitly, from a generalized action, the modified field equations of the theory. Additionally, they showed that the dynamics of the model shall include a nonstandard conservation law for the energy-momentum tensor as a consequence of the breaking of diffeomorphism resulting from the incorporation of dissipative processes into the description of the gravity. Some possible effects of this theory of gravity on the cosmological environment were explored by us in a further work \cite{fabris}. There we verified an analogy of the background dynamics arising in this model with that one of a bulk viscous cosmology in the Eckart formalism. This feature provided us an immediate mapping between the coefficient of bulk viscosity with the coupling parameter encoding the modification of gravity. We also addressed the evolution of matter perturbations at the linear level, which allowed us to glimpse a possible way out to avoid some drawbacks faced by the viscous model, perhaps leading to the reconcilement of the obtained pattern of perturbations growth with the expected background dynamics. In this work, we deepen our previous study on the cosmological aspects of the action-dependent gravity by investigating its viability in light of some important obser- vational data. As a first step, we consider a model endowed with a single matter fluid whose conservation departs from the usual one due to the geometriclike dissipative effects induced by this modified gravity. Due to the inviability of the single component model shown below, we model in Sec. III a cosmological model in which there are two pressureless components. One of them remains obeying the conservation equation while the second couples to the nontrivial geo- metrical construction of the action-dependent Lagrangian theory and therefore yields to an accelerated expansion. We investigate the viability of this model and then present our conclusion in Sec. IV. | In this work we have explored the cosmological consequences of the idea of action-dependent Lagrangians. Although this idea relies in the realm of nonstandard approaches for covariant theories of gravity, it has been deeply analyzed in the recent literature. By constructing the cosmological solutions of a FLRW expansion and the scalar perturbations around it we have demonstrated in this work what kind of cosmologies appear in this scenario. In particular, by sourcing the resulting field equations with a single fluid which due to the intrinsic features of the theory does not obey the usual conservation law, it is worth noting that the effective dynamics resembles that one of a bulk viscous fluid. As widely explored in the literature and also shown here in Sec. II such proposal of a single component driven the FLRW expansion does not provide a viable description of observational data. In Sec. III we introduced the strategy of splitting the total energy momentum tensor into two pressureless components. One of then does not couple to the geometric sector of the theory while the other one does. The former behaves therefore as a typical cold dark matter fluid while the latter plays the role of an effective dark energy fluid yielding to a consistent accelerated expansion at late times. From the astroparticle point of view our viable model can be composed by two distinct dark matter-like particles. Also, by analyzing the growth of scalar perturbations in such double pressureless components we find a reasonable agreement with available data. Our main goal in this work was to set up the a viable cosmological model composed by two pressureless fluids, one of them coupled to the geometrical sector via the features imposed by the action-dependent Lagrangian formalism. Though we have provided here only the qualitative aspects of the model but demonstrated its viability, we hope that a full statistical analysis with an enlarged set of observational data will provide the best-fit parameters of this model. We hope to present these results in a future work. \noindent {\bf Acknowledgements:} Partial financial support by FAPES, CNPq and CAPES (Brazil) is gratefully acknowledged. | 18 | 8 | 1808.02798 |
1808 | 1808.09675_arXiv.txt | Spectroscopic orbits are computed for inner pairs in 9 hierarchical multiple systems (HIP 19639, 60845, 75663, 76816, 78163, 78416, 80448, 84789, and HD~105080) and for one simple binary HIP~61840. All subsystems are double-lined, and their periods range from 2.27 to 30.4 days. Five spectroscopic binaries are twins with equal masses. Each hierarchical system is discussed individually, providing estimates of outer periods, masses, orbital inclination, and axial rotation. For systems with three resolved visual components (HIP 60845 and 80448), the outer and inner visual orbits are determined, complementing the description of their architecture. The radial velocities of HIP~75663A, 76816B, and 78163B are found to be variable with long periods, implying that these hierarchies are 2+2 quadruples. The period-eccentricity relation for spectroscopic subsystems is discussed. | \label{sec:intro} This paper continues the series on spectroscopic orbits of stars belonging to hierarchical systems \citep{paper1,paper2,paper3}. It is motivated by the need to improve statistics of orbital elements in stellar hierarchies. Statistics will inform us on the processes of their formation and dynamical evolution, as outlined in the previous papers of this series. This work augments the collection of observational data on stellar hierarchies assembled in the multiple star catalog \citep[MSC;][]{MSC}. \begin{deluxetable*}{c c rr l cc rr r c } \tabletypesize{\scriptsize} \tablecaption{Basic parameters of observed multiple systems \label{tab:objects} } \tablewidth{0pt} \tablehead{ \colhead{WDS} & \colhead{Comp.} & \colhead{HIP} & \colhead{HD} & \colhead{Spectral} & \colhead{$V$} & \colhead{$V-K$} & \colhead{$\mu^*_\alpha$} & \colhead{$\mu_\delta$} & \colhead{RV} & \colhead{$\overline{\omega}$\tablenotemark{a}} \\ \colhead{(J2000)} & & & & \colhead{Type} & \colhead{(mag)} & \colhead{(mag)} & \multicolumn{2}{c}{ (mas yr$^{-1}$)} & \colhead{(km s$^{-1}$)} & \colhead{(mas)} } \startdata 04125$-$3609 &A & 19639 & 26758 & F3V & 7.12 & 1.12 & 61 & 24 & 35.40 & 7.94 \\ &B & 19646 & 26773 & F2IV & 7.91 & 0.92 & 61 & 24 & 35.98 & 8.03 \\ 12059$-$4951 &A & \ldots & 105080 & G3V & 9.13 & 1.43 & 25 &$-$15 & 50.19 & 9.99 \\ &B & \ldots & 105081 & G0V & 9.18 & 1.42 & 31 &$-$20 & 50.09 & 7.20 \\ 12283$-$6146 &A & 60845 & 108500 & G3V & 6.82 & 1.64 & 71 &$-$160 & 40.02 & 19.93 \\ &D & \ldots & \ldots & \ldots& 13.70 & 4.54& 73 &$-$169 & \ldots & 19.91 \\ 12404$-$4924 &A & 61840 & 110143 & G0V & 7.60 & 2.00 & $-$28 & $-$112 & 6.76 & 18.59 \\ 15275$-$1058 &A & 75663 & 137631 & G0 & 8.14 & 1.35&$-$65 &$-$35 & $-$56.0 v & 9.29 \\ &B & \ldots & \ldots & G0 & 9.21 & 1.50&$-$61 &$-$35 & $-$56.82 & 7.69 \\ 15410$-$1449 &A & 76816 & 139864 & F8V & 9.47 & 1.62&$-$26 &$-$1 & $-$38.94 & 3.23 \\ &B & \ldots & \ldots & \ldots& 9.74 & 2.51&$-$25 &$-$2 & $-$50.9 v & 3.15 \\ 15577$-$3915 &A & 78163 & 142728 & G3V & 9.04 & 1.54& 17 & 6 & 9.41 & 10.42 \\ &B & \ldots & \ldots & \ldots&10.30 & \ldots& 31 & 4 & 6.78 v & 13.57 \\ 16005$-$3605 &A & 78416 & 143215 & G1V &8.65 & 1.32 &$-$26 &$-$41 & 1.60 & 9.33 \\ &B & \ldots & \ldots & \ldots&9.32 & 1.31 &$-$28 &$-$41 & 1.43 & 9.31 \\ 16253$-$4909 &AB& 80448 & 147633 & G2V & 7.5? &2.3? &$-$95 &$-$94 & $-$2.08 & 19.66 \\ 17199$-$1121 &A & 84789 & 156769 & F2 & 9.11 & 1.37& 6 & 13 & 5.62 & 5.33 \\ &B & \ldots & \ldots & \ldots& 9.89 & 1.37& 5 & 12 & 5.97 & 5.34 \enddata \tablenotetext{a}{Proper motions and parallaxes are from the {\it Gaia} DR2 \citep{Gaia,Gaia1}.} \end{deluxetable*} The systems studied here are presented in Table~\ref{tab:objects}. Only one of them (HIP 61840) is a simple binary belonging to the 67 pc sample of solar-type stars; others contain from three to five components and are also relatively close to the Sun. Their principal components are main sequence stars with spectral types from F2V to G3V. The data in Table~\ref{tab:objects} are collected from Simbad and {\it Gaia} DR2 \citep{Gaia}, the radial velocities (RVs) are determined here (variable RVs are marked by 'v'). The structure of this paper is similar to the previous ones. The data and methods are briefly recalled in Section~\ref{sec:obs}, where the new orbital elements are also given. Then in Section~\ref{sec:obj} each system is discussed individually. The paper closes with a short summary in Section~\ref{sec:sum}. | \label{sec:sum} \begin{figure} \plotone{fig14.eps} \caption{Eccentricity vs. period for members of hierarchical systems studied here (large triangles) and for 467 spectroscopic binaries from the MSC with primary masses from 0.5 to 1.5 solar (crosses). The dashed line shows the locus of HIP~78416 Aa,Ab for evolution with constant angular momentum, $P(1 - e^2)^{3/2} = {\rm const}$. \label{fig:pe} } \end{figure} Probably by accident, the periods of 9 spectroscopic systems within hierarchical multiples are equally divided between three distinct groups: (i) circular orbits with $P \approx 2.3$ days, (ii) intermediate periods between 6 and 9 days, circular or nearly circular, and (iii) eccentric orbits with periods from 21 to 30 days. Figure~\ref{fig:pe} places these orbits on the period-eccentricity plot. The plus signs are 467 spectroscopic binaries with primary masses from 0.5 to 1.5 ${\cal M}_\odot$ from the MSC \citep{MSC}. When the eccentric orbits of the group (iii) are tidally circularized, their periods will match those of group (ii), suggesting that these subsystems could be formed by a common mechanism, such as Kozai-Lidov cycles with dynamical tides \citep{Moe2018}. The periods of group (i) are substantially shorter, so their formation history could be different. Six out of the 10 double-lined binaries studied here are twins with mass ratio $q > 0.95$, while the lowest measured mass ratio is 0.67. If the mass ratios were uniformly distributed in the interval (0.7, 1.0), where double lines are detectable, the fraction of twins would be only 0.15, whereas in fact it is 0.6. It is established that twins correspond to a well-defined peak in the mass ratio distribution of solar-type spectroscopic binaries \citep{twins}. They are believed to be formed when a low-mass binary accretes a major part of its mass. The mass influx also creates conditions for formation of additional components, building stellar hierarchies ``from inside out''. Thus, twins are naturally produced as inner components of multiple systems in the process of mass assembly. The goal of this study was to determine unknown periods of spectroscopic subsystems in several multiple stars. Although this goal is reached, I discovered RV variability of other visual components (HIP 75663A, 76816B, and 78163B), converting these triples into 2+2 quadruples. The periods of new subsystems, presumably long, remain unknown so far. | 18 | 8 | 1808.09675 |
1808 | 1808.01939_arXiv.txt | We show that relativistic magnetohydrodynamics (MHD) can be recast as a novel theory of superfluidity. This new theory formulates MHD just in terms of conservation equations, including dissipative effects, by introducing appropriate variables such as a magnetic scalar potential, and providing necessary and sufficient conditions to obtain equilibrium configurations. We show that this scalar potential can be interpreted as a Goldstone mode originating from the spontaneous breaking of a one-form symmetry, and present the most generic constitutive relations at one derivative order for a parity-preserving plasma in this new superfluid formulation. | 18 | 8 | 1808.01939 |
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1808 | 1808.04388_arXiv.txt | \begin{centering} \vspace{5pt} {\bf Abstract}\\[4pt] \end{centering} \noindent Observing gamma rays using ground-based atmospheric Cherenkov telescopes provides one of the only probes of heavy weakly interacting dark matter. A canonical target is the thermal wino, for which the strongest limits come from searches for photon lines from annihilations in the Galactic Center. Irreducible finite energy resolution effects motivate refining the prediction for a wino signal beyond the photon line approximation; recently, modern effective field theory techniques have been utilized to obtain a precise calculation of the full photon energy spectrum from wino annihilation. In this paper, we investigate the implications for a realistic mock H.E.S.S.-like line search. We emphasize the impact of including the non-trivial spectral shape, and we carefully treat the region of interest, presenting results for choices between $1^{\circ}$ and $4^{\circ}$ from the Galactic Center. Projected limits for wino masses from $1$-$70$ TeV are interpreted as a constraint on the wino annihilation rate, or alternatively as the minimum core size required such that the wino is not excluded. If there is a thermal wino, H.E.S.S. will be able to probe cores of several kpc, which would begin to cause tension between this dark matter candidate and astrophysical observations/simulations. | \label{sec:model} In this section, we review several prerequisites for understanding the search for DM annihilations from the GC. We discuss the DM density distribution with an emphasis on the variations that will be utilized to explore the dependence of our results on this uncertain quantity, the relevant aspects of ground based observations with IACTs, and the recent work yielding the precision calculation of the photon spectrum used in the results that follow. \subsection{Dark Matter Density Distribution} The integrated photon flux due to the pair annihilation of DM particles from a region of solid angle $\Delta \Omega$ is computed using \begin{equation} \frac{\text{d}\Phi^{\rm DM}_{\gamma}}{\text{d}E}\big(\Delta\Omega,E\big) = \frac{\langle\sigma\, v\rangle}{8\,\pi \, m_{\DM}^2}\frac{\text{d}N_{\gamma}(E)}{\text{d}E}\,J\big(\Delta\Omega\big) \,, \label{eqn:flux} \end{equation} where $\langle\sigma\, v\rangle$ is the total annihilation cross section to all final states with a photon, $\text{d}N_\gamma/\text{d}E$ is the photon spectrum per such annihilation, $m_\DM$ is the DM mass, and $J(\Delta \Omega)$ is the integrated $J$-factor over a given region of interest (ROI) of solid angle size $\Delta \Omega$, defined by\footnote{We define the $J$-factor such that it carries units of [${\rm GeV}^2\cdot {\rm cm}^{-5}$]. Importantly, we do not associate the $\text{d} \Omega$ appearing in the definition with a unit of ${\rm sr}$, as this integral can be immediately identified as originating from a volume integral. Similarly, the $1/(4\, \pi)$ embedded in Eq.~\eqref{eqn:flux}, which originates in the surface area over which the flux from a given DM annihilation is diffused, is taken to be dimensionless. A common alternative to this convention is to associate these quantities with a ${\rm sr}$ and ${\rm sr}^{-1}$ unit, respectively, as explained in detail in App.~A of~\cite{Lisanti:2017qoz}. Note in either convention the units for the flux $\Phi_{\gamma}^{\DM}$ is identical. } \begin{equation} J\big(\Delta \Omega\big) \equiv \int_{\Delta \Omega} \text{d}\Omega \int_0^\infty \text{d}s\, \rho_\DM\big(r(s,\theta)\big)^2 \,, \label{eqn:jfac} \end{equation} where $\rho_\DM$ is the mass density of the DM. We take the standard observer centered coordinate system so that $r = \big(s^2 +r_{\odot}^2-2\,r_{\odot}\,s\, \cos\theta \big)^{1/2}$, where $s$ is the distance along the line of sight from the observer to the annihilation point, $r_{\odot} = 8.5 \text{ kpc}$ is the distance from the Sun to the GC, and $\theta$ is the angle between the direction of observation and the Galactic center. If no significant excess over the background is found, indirect searches for DM annihilation signals can be interpreted as an upper limit on the annihilation cross section into a specific final state and for a given value of the $J$-factor, which parameterizes the DM density in the ROI. If the DM density in that region is not well-known, constraints should be interpreted as a joint limit on the particle physics and the astrophysics; for a particular particle physics model, constraints on the $J$-factor can be set, and if those constraints are inconsistent with the known constraints on the DM density distribution, we can then say that the model is ruled out. The DM density in the region close to the GC has large uncertainties, because the density of visible baryonic matter is expected to dominate that of DM at small Galactocentric radii. On the observational front, this means that going from gravitational measurements of the total mass density to limits on the DM density requires careful modeling of the baryonic component and has associated large systematic uncertainties, see, for instance,~\cite{Iocco:2015xga,2017MNRAS.465.1621P}. Simulation-based predictions including hydrodynamics and feedback physics in addition to the gravitational effects for the expected DM abundance have large uncertainties due to the effects of baryonic physics, and at sufficiently small Galactocentric distances, the resolution limit of simulations also becomes relevant. These issues currently prevent them from making robust predictions for the DM profile at radii smaller than a few kpc. \begin{figure}[t] \begin{center} \includegraphics[width=0.48\textwidth]{Plots/Fig1} \caption{$J$-factors for different DM profiles as a function of the angle $\theta$ from the GC. The DM profiles chosen in this study are the Einasto profile (solid line) and cored profiles with core radii of 0.3 kpc (dashed line), 1 kpc (dotted line), 3 kpc (dotted-dashed line), and 5 kpc (long-dashed-dotted line).} \label{fig:profiles} \end{center} \end{figure} Simulations including only DM particles and neglecting the baryonic physics give rise to cusped density profiles that rise steeply toward the GC; such profiles are often parameterized by the Navarro-Frenk-White (NFW) \cite{Navarro:1995iw} or Einasto \cite{1965TrAlm...5...87E} profiles \begin{equation} \rho_\DM(r) = \rho_0 \begin{cases} \Big[ \Big(\frac{r}{r_s}\Big) \Big(1 + \left(\frac{r}{r_s}\right )^2 \Big)\Big]^{-1} & \text{NFW} \\ \\ \exp\left[-\frac{2}{\alpha} \left( \left(\frac{r}{r_s}\right)^\alpha - 1 \right)\right] & \text{Einasto}\end{cases} \,, \end{equation} where $r$ is the distance from the GC and $r_{\rm s}$ is a scale radius determined from the simulation. In this work we will use the Einasto profile as our baseline for a cusped DM profile, with the same parameters as in~\cite{Abramowski:2011hc}: explicitly, we choose $\alpha = 0.17$, $r_s=20$ kpc \cite{Pieri:2009je}, and $\rho_0$ chosen so that $\rho_\DM(r_{\odot}) = 0.39$ GeV/cm$^3$. This last choice is based on estimates of the DM density at the position of the Earth \cite{Catena:2009mf}. Once baryonic matter is included in simulations, the short-distance cusp can be flattened out, producing a ``cored'' profile. For Milky-Way-sized galaxies, the scale within which the DM density is flattened can be of order 1 kpc~\cite{Chan:2015tna}. Depending on the modeling of baryonic physics within the simulation, DM cores in Milky Way-like galaxies extending to $\sim5$ kpc can be obtained~\cite{Mollitor:2014ara}. On the observational front, the total DM mass in the Galactic Bulge region can be estimated from measurements of Bulge stellar populations \cite{2015MNRAS.448..713P}, and disfavors a NFW profile with a core size exceeding $\sim 2$ kpc \cite{Hooper:2016ggc}. However, a recent analysis using a dynamical modeling of the Galactic bulge, bar, and disk favors a shallow cusp or core in the Bulge region~\cite{2017MNRAS.465.1621P}. In order to account for possible kpc-sized DM cores in the GC, we will empirically parameterize a core of varying sizes by using the Einasto profile described above for $r > r_\text{\rm c}$, and setting $\rho_\DM(r) = \rho_\DM(r_\text{c}) = \rho_{\rm Einasto}(r_\text{\rm c})$ for $r < r_\text{\rm c}$. The normalization of the profile is such that $\rho_\DM(r_{\odot})$ = $\rho_\odot$. We plot the $J$-factor versus the angular distance, $\theta$, between the GC and the observation direction, for the Einasto profile and several choices of the core size, in Fig.~\ref{fig:profiles}. \subsection{Ground-Based Observations with IACTs} \label{sec:GCTeV} The most promising avenue for experimental tests of wino DM lies in indirect detection; since the relevant mass scales are high, IACTs have sensitivity to the annihilation products of thermal winos. Furthermore, the cross section for high-mass weakly-interacting DM annihilation can be strongly enhanced at low velocities by the nonperturbative Sommerfeld enhancement \cite{Hisano:2003ec,Hisano:2004ds}, which also enhances the gamma-ray line signal relative to the continuum emission. The enhancement effect is large for the thermal wino, and so the pure wino DM presents a particularly attractive target for gamma-ray line searches with IACTs \cite{Fan:2013faa,Cohen:2013ama}. Current arrays of IACTs like H.E.S.S., MAGIC, and VERITAS consist of 2-to-5 telescopes on the ground. The differential flux sensitivity achieved is $10^{-12}\text{ TeV}^{-1}\text{ cm}^{-2}\text { s}^{-1}$ at $\sim1\text{ TeV}$, about 1\% of the Crab flux~\cite{Aharonian:2006pe}. Based in Namibia near the tropic of Capricorn, the H.E.S.S. observatory is particularly well located to observe the central region of the Milky Way. Phase I of H.E.S.S. consists of four 12 m-diameter telescopes and reaches an angular resolution of 0.06$^{\circ}$ (68\% containment radius) and an energy resolution $\Delta E/E$ of 10\% above 300 GeV~\cite{2009APh32231D}. The GC region harbors numerous VHE gamma-ray emissions: they include H.E.S.S. J1745-290~\cite{Aharonian:2004wa,Aharonian:2009zk} a strong emission coincident with supermassive black hole Sagittarius A*, the supernova/pulsar wind nebula G0.9+0.1~\cite{Aharonian:2005br}, the supernova remnant H.E.S.S. J1745-303~\cite{Aharonian:2008gw}, and a diffuse emission extending along the Galactic plane~\cite{Aharonian:2006au,HESS2014sla,Abramowski:2016mir}. The H.E.S.S. observatory has carried out a deep observation program of the GC region from 2004 to 2014. The rich observational dataset obtained from H.E.S.S. phase I has been used to look for continuum~\cite{Abramowski:2011hc,Abdallah:2016ygi} and line~\cite{Abramowski:2013ax,Abdallah:2018qtu} signals from DM annihilations. Standard analyses of H.E.S.S.-I observations of the GC region provided $\sim 250$ hours of live time in the inner 1$^{\circ}$ of the GC with a mean zenith angle of about 20$^{\circ}$ that yields an energy threshold of 160 GeV. The energy-dependent gamma-ray acceptance reaches $\sim3\times10^5\text{ m}^2$ above 1 TeV, with a typical hadronic rejection factor of about $10$. A rejection factor of $10$ corresponds to an efficiency of $10\%$, where the efficiency is defined by the number of events passing the overall event selection cuts. In order to face the challenging standard astrophysical backgrounds, a robust approach consists of masking these regions from the data analysis for DM searches, as successfully applied in~\cite{Abdallah:2016ygi,Abdallah:2018qtu}. Once this has been performed, the dominant background in the GC region consists of misidentified CR hadrons (protons and nuclei), electrons,\footnote{The CR electron spatial distribution is assumed isotropic. No significant anisotropy of the VHE CR electrons is found in Fermi-LAT observations on any angular scale~\cite{Abdollahi:2017kyf}.} and Galactic diffuse emission. The dominant flux of CR hadrons interacting in the Earth's atmosphere generates hadronic showers which include electromagnetic sub-showers from neutral pions decaying into photons. Hadronic showers can be efficiently discriminated from the shower initiated by primary gamma-rays, requiring a stereoscopic view of the event and using morphological and timing parameters of the shower image in the camera. The incoming CR hadron flux is much larger than the CR electrons and gamma-ray fluxes, so that a fraction of the hadron flux cannot be rejected due to the finite hadron rejection power of the instrument. \begin{figure*}[t] \begin{center} \raisebox{0.75cm}{ \includegraphics[width=0.45\textwidth]{Plots/Fig2a}} \hfill \includegraphics[width=0.45\textwidth]{Plots/Fig2b.pdf} \caption{{\it Left:} A schematic of the key physical contributions for the precision calculation as relevant for a H.E.S.S. search for heavy wino annihilation. DM particles, $\chi$, annihilate to the detected photon (in red), which recoils against a jet (in blue), \emph{i.e.} a collimated spray of electroweak radiation. Low energy isotropic radiation (in green) also yields important physical effects. The winos collide with non-relativistic velocity, and the exchange of electroweak gauge bosons gives rise to the Sommerfeld enhancement (in purple). {\it Right:} The spectral shape of the endpoint contribution at NLL (orange solid line), as compared to the line which is a pure delta function (black dashed line). The theoretical spectrum is shown for a thermal wino of mass $m_{\rm DM} = 2.9$ TeV. } \label{fig:theory_setup} \end{center} \end{figure*} The measurement of the residual background in the GC region is complex~\cite{Abdallah:2016ygi,Abdallah:2018qtu}. An accurate background determination can be obtained for each observation where the background events are recorded in a region symmetric to the signal region from the pointing position. This allows the signal region and background regions to have the same sky acceptance and solid angle size, and thus does not require further offline normalization. This technique is very well suited in the case where a strong emission gradient is expected between the signal and background regions. However, the technique weakens for diffuse emission that is extended on the scale of the field of view of the instrument. In particular, for dark matter searches, this technique is proven to be very efficient in the case of cuspy DM density distributions, but fails for cored profiles with flattened density within 100 parsec or more of the GC. In order to avoid this limitation, one can extract the residual background energy distribution from extragalactic observations~\cite{Abdallah:2016ygi}. In this case, the residual background is extracted from blank fields at high Galactic latitudes in the same observation conditions as for the GC dataset. Alternatively, Monte Carlo simulations can be used to predict the residual background rate, since they can be performed in the same observational conditions and telescope configurations as for the GC dataset, allowing for reduced systematic uncertainties~\cite{Holler:2017ynz}. \subsection{Precision Signal Spectrum} Due to the large backgrounds, the most striking signal for DM annihilation is a line signal, which in this energy range should not arise from astrophysical backgrounds. However, because of the finite energy resolution of IACTs, it is impossible to measure only the line spectrum; gamma-ray photons from DM annihilation to $\gamma\, \gamma$ or $\gamma\, Z$ will inevitably be accompanied by gamma-ray photons from other annihilation final states, and these cannot be distinguished on an event-by-event basis if the energy of the resulting photon varies by an amount less than the energy resolution of the telescope. Furthermore, if a smooth background model is included in the fit, as in \cite{Abramowski:2013ax}, the unaccounted-for presence of lower-energy signal photons could potentially bias the background model. Thus to obtain precise and accurate constraints, it is important to have a theoretical prediction for the full photon spectrum to compare with the data, rather than simply comparing constraints on an isolated line to a theoretical prediction for the strength of the line signal. Obtaining a reliable prediction for the photon spectrum from wino annihilation is complicated by the presence of the hierarchical scales $\mW$, and $m_\DM$, and in the endpoint region -- where \begin{equation} z\equiv \frac{E}{m_\DM}\,, \label{eq:z} \end{equation} is close to 1 -- by the presence of non-trivial phase space restrictions. By looking for a line signal within the H.E.S.S. resolution of $m_\DM$, the final state must be close to a two-body decay, for which $z=1$. More precisely, it must consist of collimated energetic radiation recoiling against the detected photon (an electroweak jet), as well as additional low energy radiation. Any additional radiation would make the energy of the photon far from $m_\DM$. This configuration is shown on the left of Fig.~\ref{fig:theory_setup}. This restriction introduces perturbative Sudakov double logarithms $\aW\log^2(m_\DM/\mW)$~\cite{Hryczuk:2011vi,Baumgart:2014vma,Bauer:2014ula,Ovanesyan:2014fwa,Baumgart:2014saa,Baumgart:2015bpa,Ovanesyan:2016vkk}, and $\aW\log^2(1-z)$ \cite{Baumgart:2014vma,Baumgart:2014saa,Baumgart:2015bpa,Baumgart:2017nsr} , as well as Sommerfeld enhancement terms of the form $\left( \aW m_\DM/\mW \right )^k$ \cite{Hisano:2003ec,Hisano:2004ds,Cirelli:2007xd,ArkaniHamed:2008qn,Blum:2016nrz}. Here $\aW$ is the weak fine structure constant. Reliable predictions for the shape of the distribution in the endpoint region require that all these effects are resummed to all orders in perturbation theory. Once the perturbative series is reorganized in this manner, it again converges rapidly due to the smallness of the electroweak coupling, $\aW$, allowing precise theoretical predictions for the photon spectrum. In \cite{Baumgart:2017nsr}, an effective field theory (EFT) framework was developed for the calculation of the photon spectrum in the endpoint region for heavy DM annihilation. It combines non-relativistic EFTs for the description of the annihilating DM, and soft-collinear effective theory (SCET) \cite{Bauer:2000yr, Bauer:2001ct, Bauer:2001yt}, as well as its multi-scale extensions \cite{Bauer:2011uc,Larkoski:2014tva,Procura:2014cba,Larkoski:2015zka,Pietrulewicz:2016nwo}, and extensions to include massive gauge bosons~\cite{Chiu:2007yn,Chiu:2008vv,Chiu:2007dg}, for the treatment of the final state radiation. This EFT allows the photon energy spectrum to be computed precisely, properly incorporating both the Sommerfeld and Sudakov effects (as well as their interplay) to all orders, and allowing for reliable uncertainty estimates. Using this EFT, an analytic form for the photon energy spectrum in the endpoint region for annihilating pure wino DM was derived in~\cite{Baumgart:2017nsr} at leading logarithmic (LL) accuracy, and this calculation was extended in~\cite{us:NLL} to next-to-leading logarithmic (NLL) accuracy, greatly reducing the theoretical uncertainty. For simplicity, we present the final formula for the LL photon spectrum in this region, as this allows us to illustrate the general features of its shape in the endpoint region with a simple expression. We then briefly comment on how this is modified by additional logarithmic corrections at NLL. We refer the reader to~\cite{Baumgart:2017nsr} for the derivation of the LL result, and \cite{us:NLL} for the analytic form of the spectrum at NLL. In the endpoint region, the photon spectrum at LL accuracy can be written as a function of $z$ and $m_\DM$ as \newpage \begin{widetext} \begin{align} \left(\frac{\text{d} \sigma}{\text{d}z}\right)^{\text{LL}} \!&=\, 4 \,|s_{0\pm}|^2\, \hat \sigma^\text{LL}_{\text{line}}\, \delta(1-z) + \frac{2\,\aW}{\pi} \frac{\hat \sigma^\text{LL}_{\text{line}}}{1-z} \,e^{\frac{4\,\aW}{\pi} \, L_J^2(z)} \bigg\{ F_1 \Big( 3\,L_S(z) - 2\,L_J(z) \Big) e^{\frac{-3\,\aW}{\pi}\, L^2_S(z) } - 2\,F_0\, L_J(z) \bigg\} \,. \label{eq:resummed} \end{align} \end{widetext} This provides a simple analytic expression describing both the line contribution, which is given by the first term in Eq.~(\ref{eq:resummed}) proportional to $\delta(1-z)$, as well as the endpoint contribution which is given by the second term, and is a non-trivial function of $z$, describing the steeply falling spectrum. The line contributions were first calculated with resummation in \cite{Bauer:2014ula,Ovanesyan:2014fwa,Ovanesyan:2016vkk}, while it is the shape of the spectrum away from $z=1$ that is the primary contribution of \cite{Baumgart:2017nsr}. On the right of Fig.~\ref{fig:theory_setup} we show the spectrum of photons associated with the endpoint for the thermal wino. We now describe each of the components of Eq.~(\ref{eq:resummed}) in turn. Both terms are multiplied by the exclusive line cross section (without Sommerfeld effects), which at leading logarithmic accuracy is given by \begin{align} \hspace{-2pt}\hat \sigma^\text{LL}_{\text{line}}= {\pi \,\aW^2\, \sin^2 \thetaW \over 2\,m_\DM^2\, v} \exp\left[- \frac{4\,\aW}{\pi}\, \ln^2 \left( \frac{\mW}{2\,m_\DM} \right) \right] \,, \label{eq:partonicSigmaLine} \end{align} where $\thetaW$ is the Standard Model weak mixing angle. This can be computed within an EFT framework by considering charged wino annihilation into both $\gamma\,\gamma$ and $\gamma\,Z$. The exponential appearing in this formula is the massive Sudakov form factor \cite{Collins:1989bt} and is due to the exchange of virtual electroweak bosons. This process is then mapped onto the neutral wino initial state by non-trivial mixing due to the Sommerfeld enhancement involving the exchange of a ladder of gauge bosons with one or more $W^\pm$ bosons, \emph{i.e.}, $s_{0\pm} \neq 0$. The energy dependence of the photon spectrum in the endpoint region, which is crucial to our analysis, is described by a $1/(1-z)$ power law growth towards the endpoint, modified by the logarithms \begin{align} L_J(z)&= \ln \left({\mW/m_\DM \over 2 \, \sqrt{1-z} }\right) \Theta\!\left(1-\frac{\mW^2}{4\, m_\DM^2}-z \right)\,, \nonumber \\[5pt] L_S(z)&= \ln \left({\mW/m_\DM \over 2 \, (1-z) }\right) \Theta\!\left( 1-\frac{\mW}{2\,m_\DM}-z \right) \,, \label{eq:LJLS} \end{align} associated with additional radiation in the final state. The $\Theta$ functions are set by the kinematics, and cut off the divergence in the $1/(1-z)$ growth before reaching the $z=1$ endpoint. Importantly, this power law form is unmodified beyond LL, with higher order corrections simply dressing this result with additional logarithms. Furthermore, we find that these higher order corrections are of the anticipated size, and that their primary utility is to reduce the theoretical uncertainty. Finally, the non-perturbative Sommerfeld effect is captured by a non-relativistic quantum mechanics calculation of the matrix element of the $S$-wave combination for the annihilating neutral winos $(\chi^0\chi^0)_S$, \begin{align} \label{eq:wavefunction} &\Big\langle 0 \Big|\, \chi^{0\,T}\, i\sigma_2 \,\chi^{0\,\,} \,\Big| \big(\chi^0 \chi^0\big)_S \Big\rangle = 4 \sqrt{2} \, m_\DM\, s_{00} \,, \nonumber \\[5pt] &\Big\langle 0 \Big|\, \chi^{+T}\, i\sigma_2 \,\chi^-\, \Big| \big(\chi^0 \chi^0\big)_S \Big\rangle= 4\, m_\DM \,s_{0\pm} \,, \end{align} where $\sigma_2$ is the second Pauli matrix, and we have used the standard notation $\chi^0 = \chi^3$ and $\chi^\pm = (\chi^1 \mp i \chi^2)/\sqrt{2}$ for the neutral and charged wino states respectively. Then $s_{00}$ ($s_{0\pm}$) provides the enhancement for a neutral wino initial state and a perturbative Feynman diagram involving neutral (charged) wino annihilation. For a detailed discussion, see \emph{e.g.}~\cite{Beneke:2012tg, Cohen:2013ama}. For the LL line contribution, only the $s_{0\pm}$ contributes, since we are only matching the EFT to the full theory at tree level, which implies that the only non-zero diagrams are due to charged wino annihilation to $\gamma\,\gamma$ and $\gamma\,Z$. However, in the endpoint region, the non-trivial combinations \begin{align} F_0 &= \frac43 \,\big|s_{00}\big|^2 + 2 \,\big|s_{0\pm}\big|^2 + {4\, \sqrt{2} \over 3}\, \Re\Big (s_{00} \,s^*_{0\pm} \Big) \,, \nonumber \\[3pt] F_1 &= - \frac43 \,\big|s_{00}\big|^2 + 2\,\big |s_{0\pm}\big|^2 - {4 \sqrt{2} \over 3}\, \Re\Big(s_{00}\, s^*_{0\pm}\Big ) \,, \end{align} also appear, where $\Re(\dots)$ gives the real part of the argument. This occurs because three-body processes yielding $W^+\,W^-\,\gamma$ can now appear, and this channel is non-zero for both charged and neutral wino annihilation. Beyond LL the structure of both the line and endpoint spectrum contributions become more sophisticated, but the basic ingredients discussed here, and type of logarithms that are resummed, remain the same. For our analysis here we will make use of the full NLL results from \cite{us:NLL}. These results for the case of wino DM can in principle be straightforwardly extended to other heavy WIMP DM. The calculation presented here is based on an EFT expansion that requires that the resolution is much greater than $\mW/(2\,m_\DM)$. We find that our calculation is not reliable below $\sim 1$ TeV. For masses below this value we would have to match our prediction onto an EFT that is valid in the low mass region. This is beyond the scope of the current paper, and therefore we only consider $m_\DM\gtrsim1$ TeV. However, due to the high quality data in this low mass region, we believe it would be interesting to consider, and we intend to pursue this in future work. For recent EFT work relevant to DM masses below $1$ TeV, see~\cite{Beneke:2018ssm}. \subsection{Interpreting the Signal Prediction} \label{sec:interpretSignal} Given this precise prediction for the gamma-ray spectrum resulting from wino annihilation, it is worth revisiting the procedure for converting this into a flux prediction that can be used to probe wino annihilation. In particular, one of our goals here is to understand the extent to which including the spectral shape -- the \emph{endpoint} spectrum in the results below -- impacts the limits one would set, when compared to the limits derived assuming that only the line contribution is relevant. To this end, we define the \emph{line} annihilation cross section $\sigma_\text{line}$ to be half the coefficient of $\delta\big(E - m_\DM\big)$ in the expression for $\text{d}\sigma/\text{d}E$. For example, in the LL case we take the differential spectrum in Eq.~(\ref{eq:resummed}) and derive that $\sigma_\text{line} = 2 \,|s_{0\pm}|^2\, \hat \sigma_{\text{line}}$, where $\hat \sigma_\text{line}$ is defined in Eq.~(\ref{eq:partonicSigmaLine}), and we have used Eq.~(\ref{eq:z}) to convert $z$ into $E$. We emphasize that the NLL result is used for all numerical analysis in what follows. Our conventions are such that the contribution to $\text{d}\sigma/\text{d}E$ is normalized as $2\, \sigma_\text{line}\, \delta\big(E - m_\DM\big)$, where the factor of 2 accounts for the presence of two photons in exclusive $\chi\, \chi \rightarrow \gamma\, \gamma$ annihilations. For line events, we must also include the branching rate to $\gamma\,Z$ though, giving $\sigma_\text{line} = \sigma(\chi\,\chi \rightarrow \gamma\, \gamma) +(1/2)\,\sigma(\chi\,\chi\rightarrow\gamma\,Z)$. The analysis can then be interpreted as either a constraint on $\langle \sigma\,v \rangle_\text{line}$, or as a constraint on the DM profile using the predicted wino rate. In order to include the endpoint spectrum in the analysis, we take the NLL analog of Eq.~(\ref{eq:resummed}), subtract the contribution proportional to the line $\delta\big(E - m_\DM\big)$, and normalize to $\sigma_\text{line}$. This yields an analytic prediction for the endpoint spectral shape that we will refer to as $\big(\text{d}\mathcal{N}_{\gamma}(E)/\text{d}E\big)^\text{endpoint}$, specifically \begin{align} \bigg(\frac{\text{d}\sigma}{\text{d}E}\bigg)^{\text{NLL}} = \sigma_\text{line} \left[2\, \delta\big(E - m_\DM\big) + \left(\frac{\text{d}\mathcal{N}_\gamma}{\text{d}E}\right)^\text{endpoint}\right ]\,.\notag\\[3pt] \end{align} Note that the use of a new notation for the spectrum, $\text{d}\mathcal{N}_{\gamma}/\text{d}E$, rather than $\text{d} N_{\gamma}/\text{d}E$ as appeared in Eq.~\eqref{eqn:flux} is deliberate, and is designed to emphasize that we are using a spectrum normalized to the line cross section. \begin{table*}[tb] \centering \renewcommand{\arraystretch}{1.8} \setlength{\tabcolsep}{4.5pt} \setlength{\arrayrulewidth}{1.3pt} \begin{tabular}{|r|c|c|c|c|c|c|c|} \hline \multirow{2}{*}{$i$-th ROI}\hspace{15pt} & \multirow{2}{*}{$\begin{array}{c}\text{Solid angle:}\\[-5pt] \Delta\Omega_i\,\,\, \big[10^{-4} \text{ sr}\big]\end{array}$} & \multicolumn{6}{c|}{$J\text{-factor: } J_i \big(\Delta\Omega_i \big) \,\,\, \big[10^{20} \text{ GeV}^2\text{ cm}^{-5}\big]$} \\[2pt] \cline{3-8} & & Einasto & $r_\text{c} = 0.3 \text{ kpc}$ & $r_\text{c} = 0.5 \text{ kpc}$ & $r_\text{c} = 1 \text{ kpc}$ & $r_\text{c} = 3 \text{ kpc}$ & $r_\text{c} = 5 \text{ kpc}$\\ \hline 1: $\big(\bar \theta_1= 0.3^\circ\big)$ & 0.31 & 3.76 & 1.08 & 0.60 & 0.23 & 0.035 & 0.012 \\ 2: $\big(\bar \theta_2=0.4^\circ\big)$ & 0.50 & 5.16 & 1.14 & 0.97 & 0.38 & 0.056 & 0.019 \\ 3: $\big(\bar \theta_3=0.5^\circ\big)$ & 0.69 & 6.15 & 2.40 & 1.34 & 0.52 & 0.078 & 0.026\\ 4: $\big(\bar \theta_4=0.6^\circ\big)$ & 0.88 & 6.89 & 3.04 & 1.71 & 0.66 & 0.099 & 0.033\\ 5: $\big(\bar \theta_5=0.7^\circ\big)$ & 1.08 & 7.45 & 3.67 & 2.07 & 0.81 & 0.12 & 0.040\\ 6: $\big(\bar \theta_6=0.8^\circ\big)$ & 1.27 & 7.88 & 4.29 & 2.43 & 0.95 & 0.14 & 0.047\\ 7: $\big(\bar \theta_7=0.9^\circ\big)$ & 1.46 & 8.20 & 4.90 & 2.79 & 1.09 & 0.16 & 0.055\\ \vdots\hspace{31pt} & \vdots & \vdots &\vdots & \vdots& \vdots & \vdots & \vdots\\ 37: $\big(\bar \theta_{37}=3.9^\circ\big)$ & 7.55 & 8.78 & 8.78 & 8.78 & 5.23 & 0.88 & 0.28 \\ \hline \end{tabular} \caption{ \small \label{tab:tableJ} Definitions of the $i$-th ROI together with the corresponding solid angle size, and value of the $J$-factor for the several DM profiles considered here. For brevity, we only show the first $7$ ROIs, which are used in the $1^\circ$ analysis, and then skip to the 37th since this is the largest ROI considered in this work, and is used in the $4^\circ$ analysis. } \end{table*} Finally, in addition to investigating the impact of the perturbative endpoint spectrum, we will also include the contribution to the gamma-ray flux from processes where the hard annihilation is to a Standard Model final state that generates a spectrum of photons due to its subsequent decay. We refer to this as \emph{continuum} in the results below. In the wino example, both the neutral and charged winos can annihilate to $W^\pm$ bosons, which then decay. Note that both parton level processes must be included, since again they will be accessed by the mixing due to the Sommerfeld effect. The Sommerfeld enhanced annihilation rate is then convolved with the final state photon spectrum provided by the PPPC 4 DM ID~\cite{Cirelli:2010xx}. Following the same logic as with the endpoint spectrum, we take the PPPC results, and normalize them to $ \sigma_\text{line}$ in order to derive the continuum spectral shape, referred to as $\big(\text{d}\mathcal{N}_{\gamma}(E)/\text{d}E\big)^\text{continuum}$. Now we are setup to compare the three levels of approximation -- $(i)$ line, $(ii)$ line + endpoint, and $(iii)$ line + endpoint + continuum. Revisiting Eq.~(\ref{eqn:flux}), \begin{equation} \hspace{-5pt}\frac{\text{d}\Phi^{\rm DM}_{\gamma}}{\text{d}E}\big(\Delta\Omega,E\big) = \frac{\langle \sigma\,v \rangle_\text{line} \,J\big(\Delta\Omega\big)}{8\,\pi\,m_{\rm DM}^2} \left[\frac{\text{d}\mathcal{N}_\gamma(E)}{\text{d}E}\right] \,,\!\! \label{eqn:fluxRevisited} \end{equation} where we have grouped the combination $\big(\langle \sigma\,v \rangle_\text{line} \,J\big(\Delta\Omega\big)\big)$ since this is the quantity that can be constrained using IACT data, and the bracketed term encodes the spectral shape, \begin{align} \notag\\[-5pt] \left[\frac{\text{d}\mathcal{N}_{\gamma}(E)}{\text{d}E}\right] = \begin{cases} 2\,\delta\big(E - m_\DM\big) & (i) \\ \\ (i) + \left(\frac{\text{d}\mathcal{N}_{\gamma}(E)}{\text{d}E}\right)^\text{endpoint} & (ii) \\ \\ (ii) + \left(\frac{\text{d}\mathcal{N}_{\gamma}(E)}{\text{d}E}\right)^\text{continuum}& (iii) \\[5pt] \end{cases}\,. \end{align} Note that having a broad spectrum impacts the analysis, as it can lead to contamination outside of the energy window associated with a signal of a given mass. Note that a benchmark choice $\langle \sigma\,v \rangle_\text{line} = 10^{-27} \text{cm}^3/\text{s}$ is used for many of the plots below, since this is about where the projected limit for a $3 \text{ TeV}$ wino will lie. Now that we have a clear understanding of the various aspects of the signal prediction, we are ready to explain the procedure used for deriving our expected limits. | \label{sec:summary} We show the prospects for wino DM over a mass range from 1 TeV up to 70 TeV using VHE gamma-ray observations of the GC that rely on the most-up-to-date EFT computation of the annihilation spectrum of winos. We build realistic mock data simulations of H.E.S.S.-I-like observations of the GC region and implement spectral and spatial analysis of the VHE emissions. We compute the sensitivity to wino DM using a binned likelihood test statistic ratio using the spectral and spatial information of signal and background. Various DM density distributions in the GC region are considered including DM density cores up to 5 kpc. We show that ($i$) the line contribution to the wino annihilation spectrum drives the overall limits in the TeV mass range, ($ii$) the endpoint contribution significantly improves the sensitivity compared to the line-only signal with increasing importance for higher DM masses, and ($iii$) the continuum contribution is sub-dominant compared to the line and endpoint contributions but becomes more relevant as the DM mass increases. The present sensitivity of H.E.S.S.-I-like observations is able to provide strong constraints on wino DM. We show for the case of winos constituting 100\% of the DM that strong constraints on the DM density in the central region of the Milky Way can be obtained using H.E.S.S.-I-like observations of the GC region. In particular, DM cores up to several kpc radii could be excluded for 2.3 and 9 TeV masses respectively, where the Sommerfeld effect strongly enhances the annihilation cross section through resonances. We additionally provided a sensitivity projection for H.E.S.S.-I-like observations of the GC region using the IGS strategy. For the thermal wino mass of 2.9 TeV, we find a projected core size limit of approximately 4.5 kpc. This makes clear that future searches will provide a decisive test for the wino under reasonable assumptions for how the DM is distributed in the center of the Milky Way at kpc scales. | 18 | 8 | 1808.04388 |
1808 | 1808.02989_arXiv.txt | \noindent We report the first detection of a credible progenitor system for a Type Ic supernova (SN~Ic), SN~2017ein. We present spectra and photometry of the SN, finding it to be similar to carbon-rich, low-luminosity SNe Ic. Using a post-explosion Keck adaptive optics image, we precisely determine the position of SN 2017ein in pre-explosion \hst\ images, finding a single source coincident with the SN position. This source is marginally extended, and is consistent with being a stellar cluster. However, under the assumption that the emission of this source is dominated by a single point source, we perform point-spread function photometry, and correcting for line-of-sight reddening, we find it to have $M_{\rm F555W} = -7.5\pm0.2$~mag and $m_{\rm F555W}-m_{\rm F814W}$=$-0.67\pm0.14$~mag. This source is bluer than the main sequence and brighter than almost all Wolf-Rayet stars, however it is similar to some WC+O- and B-star binary systems. Under the assumption that the source is dominated by a single star, we find that it had an initial mass of $55\substack{+20\\-15}~M_{\odot}$. We also examined binary star models to look for systems that match the overall photometry of the pre-explosion source and found that the best-fitting model is a $80$+$48~M_{\odot}$ close binary system in which the $80~M_{\odot}$ star is stripped and explodes as a lower mass star. Late-time photometry after the SN has faded will be necessary to cleanly separate the progenitor star emission from the additional coincident emission. | \label{sec:introduction} In the past three decades, there have been over $20$ detections of pre-explosion counterparts to core-collapse supernovae \citep[SNe; for a review, see][]{smartt09}. Most of these counterparts are red supergiant (RSG) progenitor stars of Type II-P SNe (SNe with a ``plateau'' in their light curves consistent with recombination emission from an extended hydrogen envelope), which agrees with predictions from star formation and stellar evolution that suggest low-mass RSG progenitors stars should be relatively common. There are, however, mixed results in finding the progenitor systems of other SN sub-types, with identified progenitor systems for some SNe~IIn \citep[SNe with narrow lines of hydrogen in their spectra, e.g., SNe~2005gl and 2009ip;][]{gal-yam+07,smith+10} and SNe~IIb \citep[SNe with transient hydrogen lines in their spectra, with progenitor star detections for SNe~1993J, 2008ax, 2011dh, 2013df, and 2016gkg;][]{aldering+94,woosley+94,crockett+08,maund+11,vandyk+14,dessart+11,dessart+15,kilpatrick+17}. The progenitor stars of SNe~Ib/c (which have no hydrogen in their spectra, or helium in the case of SNe~Ic) have been comparatively elusive and only one credible pre-explosion counterpart has been identified so far in the literature \citep[the SN~Ib iPTF13bvn;][]{cao+13}. In part, the paucity of pre-explosion SN counterparts for SNe~Ib/c is because they only make up $\sim$20\% of transients discovered in volume-limited surveys \citep[e.g., LOSS;][]{li+00,li+11,smith+11,shivvers+17}, and the incidence of SNe with deep, high-resolution pre-explosion imaging is even smaller. However, as more nearby SNe are discovered, especially those with pre-explosion {\it Hubble Space Telescope} (\hst) imaging, the growing sample of upper limits on SN~Ic progenitor systems in particular has placed strong constraints on predictions from stellar evolution and SN explosion models \citep[with deep limits on counterparts for SNe~2002ap, 2004gt, and 2007gr;][]{gal-yam+05,crockett+07,crockett+08,maund+16}. This evidence suggests that some non-RSG SN progenitor stars are either intrinsically less luminous than RSGs or heavily obscured by dust in the \hst\ optical bands typically available for pre-explosion imaging. These stars may highly-stripped by stellar winds, and although they may be comparable in luminosity to RSGs, their SEDs peak predominantly in the ultraviolet \citep[and outside of optical or infrared bands in which pre-explosion imaging is typically available; for a review of SN~Ib/c progenitor studies see][]{eldridge+13}. Dust obscuration is a distinct possibility for high-mass SN progenitor stars, as some high-mass RSGs are observed to have optically thick circumstellar dust \citep[e.g., SN~2012aw;][]{kochanek+12}. In addition, high-mass SN progenitor stars ought to explode promptly, perhaps close to the dusty environments where they form \citep[see, e.g., analysis of SN environments in][]{kuncarayakti+13,galbany+16,galbany+17}. Because SNe~Ic are the explosions of massive stars without significant hydrogen or helium in their outer layers, their progenitor star must be significantly stripped by stellar winds or a companion star. Highly-stripped Wolf-Rayet (WR) stars are therefore good candidates for SN~Ic progenitor stars \citep{yoon+10,yoon+12,yoon+17}. WR stars undergo radiatively-driven mass loss at rates exceeding $10^{5}~M_{\odot}~\text{yr}^{-1}$ \citep[although exact mass-loss rates are highly uncertain;][]{maeder+87,hamann+95,smith+14}, and so observational and theoretical evidence suggest that some pre-SN WR stars ought to be hydrogen- and helium-deficient \citep[][]{podsiadlowski+02,woosley+93,steiner+05}. However, radiatively-driven winds are highly metallicity-dependent and WR stars tend to form in high-metallicity environments; indeed, the Small Magellanic Cloud exhibits a decreased WR-to-O-star ratio relative to Solar neightborhood \citep{hainich+15}. Predicted mass-loss rates for WR stars at Solar metallicities indicate that single WR stars may have high pre-SN masses \citep{meynet+05} and very few of these stars are predicted to be helium-poor \citep[e.g.,][]{yoon+15}. This finding is in tension with predictions that they are SN~Ic progenitor stars given SN~Ic rates and their typical ejecta masses \citep{drout+11,taddia+15}. One alternative is that, if WR stars are a likely channel for producing SNe~Ic, most SN progenitor systems are interacting binaries in which a WR star has been stripped by a companion. This hypothesis is supported by the fact that many late-type WR stars are observed to be in close binaries with O-type stars \citep[e.g., WR~104;][]{tuthill+99} as well as the fact that the overall binary fraction for Milky Way WR stars is 40\% \citep{vanderhucht+01}. If SNe~Ic mostly come from low-mass WR stars in close binaries or in dusty environments, this would explain their non-detection in optical pre-explosion imaging to date, and so examples with deep detection limits can be used to verify or rule out this possibility. In this paper we discuss the SN~Ic 2017ein discovered in NGC~3938 on 2017 May 25 by \citet{arbour+17}. Deep imaging starting 2~days before discovery and continuing for 2~weeks afterward indicated that SN~2017ein rose quickly after discovery \citep{atel10481}, which suggests that it was discovered very soon after explosion. Here we report photometry, spectroscopy, and high-resolution adaptive optics imaging of SN~2017ein. We demonstrate that SN~2017ein is most consistent with carbon-rich SNe~Ic, although the source exhibits strong Na\I\ D lines at the redshift of NGC~3938 and is significantly reddened. We use relative astrometry between our high-resolution and pre-explosion imaging, we find a single, luminous, blue source consistent with being the progenitor system of SN~2017ein, although that source appears extended and may be a blend of multiple point sources. By comparing this source to Galactic supergiants and evolutionary tracks, we investigate channels that could produce the SN~2017ein progenitor system. While we were preparing this manuscript, \citet{vandyk+18} published another analysis of SN~2017ein and its pre-explosion imaging. The authors came to similar conclusions about the nature of SN~2017ein and its photometric and spectroscopic similarity to carbon-rich SNe~Ic. They identified the same source in pre-explosion imaging as the potential progenitor system of SN~2017ein and concluded the SN likely had a very massive ($>45~M_{\odot}$) progenitor star. Throughout this paper, we assume a distance to NGC~3938 of $m-M=31.17\pm0.10$ \citep{tully+09} and Milky Way extinction of $A_{V}=0.058$ \citep{schlafly+11}. | \label{sec:conclusions} We present pre-explosion imaging and high-resolution imaging, photometry, and spectroscopy of the SN~Ic 2017ein. We find: \begin{enumerate} \item Spectra and light curves of SN~2017ein are remarkably similar to carbon-rich, low-luminosity SNe~Ic such as SN~2007gr and unlike SNe~Ic such as 2011bm. At the same time, matching the continuum and peak $V$-band luminosity of SN~2017ein to SN~2007gr requires roughly $A_{V}=1.2$~mag of host extinction. We also detect strong Na\I\ D absorption at the approximate redshift of NGC~3938. These spectral characteristics suggest that the progenitor system contained very little hydrogen or helium, but also that it may have had an intrinsically high carbon abundance in its outer layers, as has been suggested for some WC stars. \item The location of SN~2017ein as determined from high-resolution laser guide star adaptive optics imaging is consistent with a single source in pre-explosion \hst/WFPC2 imaging. The source is marginally extended in the \hst/WFPC2 images and there may be non-uniform background emission at this location. \item Accounting for the extended source and host extinction, photometry from the pre-explosion is consistent with single stars with masses up to $75~M_{\odot}$, but with a preferred mass of $55~M_{\odot}$. However, most of these stars, which include O- and B-type supergiants and WN stars, are hydrogen-rich, and so are unlikely SN~Ic progenitor stars. \item Comparison to highly-stripped WR star binaries indicates that the only systems that match the colors and luminosity of PSF1 are WC+O and B star binaries. We find that a $80$+$48~M_{\odot}$ BPASS model can explain some of the parameters of SN~2017ein and the pre-explosion counterpart and produces a star whose terminal state is roughly consistent with predictions of SN~Ic progenitor stars. Additional modeling is needed to explore the full ramifications of this evolutionary pathway and the precise terminal state of such a system. \item Nebular spectroscopy of SN~2017ein will be critical for measuring the true carbon abundance in the ejecta. Late-time imaging of the site of SN~2017ein will also be important for measuring the extent to which the SN~2017ein progenitor star contributed to emission from the pre-explosion source. \end{enumerate} \bigskip\bigskip\bigskip \noindent {\bf ACKNOWLEDGMENTS} \smallskip \footnotesize We thank Raj Chowdhury and Bella Nguyen for help with Nickel observations. We also thank David Coulter, C\'{e}sar Rojas-Bravo, and Matthew Siebert for help with Shane and Mayall observations. The UCSC group is supported in part by NSF grant AST--1518052, the Gordon \& Betty Moore Foundation, the Heising-Simons Foundation, and by fellowships from the Alfred P.\ Sloan Foundation and the David and Lucile Packard Foundation to R.J.F. Some of the data presented herein were obtained at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and NASA. The observatory was made possible by the generous financial support of the W. M. Keck Foundation. We wish to recognise and acknowledge the cultural significance that the summit of Mauna Kea has within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. Some of the data in this publication were calibrated using object catalogs from the Pan-STARRS1 Surveys. The Pan-STARRS1 Surveys (PS1) and the PS1 public science archive have been made possible through contributions by the Institute for Astronomy, the University of Hawaii, the Pan-STARRS Project Office, the Max-Planck Society and its participating institutes, the Max Planck Institute for Astronomy, Heidelberg and the Max Planck Institute for Extraterrestrial Physics, Garching, The Johns Hopkins University, Durham University, the University of Edinburgh, the Queen's University Belfast, the Harvard-Smithsonian Center for Astrophysics, the Las Cumbres Observatory Global Telescope Network Incorporated, the National Central University of Taiwan, the Space Telescope Science Institute, the National Aeronautics and Space Administration under Grant No. NNX08AR22G issued through the Planetary Science Division of the NASA Science Mission Directorate, the National Science Foundation Grant No. AST-1238877, the University of Maryland, Eotvos Lorand University (ELTE), the Los Alamos National Laboratory, and the Gordon and Betty Moore Foundation. The {\it Hubble Space Telescope} (\hst) is operated by NASA/ESA. The \hst\ data used in this manuscript come from programme GO-10877 (PI Li). Some of our analysis is based on data obtained from the \hst\ archive operated by STScI. We acknowledge the use of public data from the {\it Swift} data archive. Some of the data presented in this manuscript come from the Kitt Peak National Observatory (KPNO) 4-m telescope through programme 2017A-0306 (PI Foley). KPNO is operated by the Association of Universities for Research in Astronomy, Inc. (AURA) under cooperative agreement with the National Science Foundation. The Nickel and Shane telescopes are operated by the University of California and Lick Observatories. Some of the data presented in this manuscript come from UCO/Lick programmes 2017Q2-N007, 2017Q3-N005 (PI Kilpatrick) and 2017A-S011, 2017B-S018 (PI Foley). This work makes use of observations performed by the Las Cumbres Global Telescope Network through programme 2017AB-012 (PI Kilpatrick). \textit{Facilities}: Keck (ESI/LRIS/OSIRIS), LCOGTN (SINISTRO), Mayall (KOSMOS), Nickel (Direct), Shane (Kast), {\it Swift} (UVOT) | 18 | 8 | 1808.02989 |
1808 | 1808.07081_arXiv.txt | We investigate the possibility that multiple populations in globular clusters arise as a natural by-product of massive star-cluster formation. We use 3D radiative hydrodynamics simulations for the formation of young massive clusters to track their chemical self-enrichment during their first 5 Myr. These clusters form embedded within filamentary Giant Molecular Clouds by a combination of gas accretion and rapid merging of protoclusters. Chemical enrichment is a dynamic process happening as the young cluster assembles, so that the original (1P) and enriched (2P) subpopulations of stars form almost simultaneously. Here we test two simple and opposite extremes for the injection of enriched material into the intracluster gas: we assume either \emph{continuous injection} in a way that tracks the star formation rate; or \emph{sudden injection} by a single instantaneous event. Using helium abundance as a proxy for the enrichment, we find that realistic multiple population features can be reproduced by injecting a total helium mass amounting to a few percent of the cluster's total mass. The differences in individual growth histories can lead to widely differing 1P/2P outcomes. These models suggest that dual or multiple populations can emerge rapidly in massive star clusters undergoing the typical mode of star cluster formation. | Globular clusters are massive ($10^{4-7} M_{\odot}$) and low-metallicity ([Fe/H] $\lesssim 0$) star clusters that formed in the early universe \citep{Kruijssen} and are present in all large galaxies. An outstanding puzzle is that many globular clusters host multiple populations of stars with distinct chemical abundance patterns \citep{Maclean2015,BastianLardo}. The existence of multiple populations (MPs) in globular clusters (GCs) is nearly ubiquitous regardless of cluster metallicity and has been identified through abundance anomalies of proton capture elements such as C-N, Na-O, Mg-Al, and Na-F anticorrelations, as well as He abundance spreads \citep{Carretta2009,Piotto,Mucciarelli2014,Milone2018}. In many cases, GCs host two distinct stellar populations roughly in a 1:1 ratio \citep{Caretta2009b} but with wide variety; 3 or more populations have been observed in several clusters, and in some cases the abundance distribution resembles a roughly continuous spread with a few indentifiable clumps \citep{Milone2017}. MPs characterized by these abundance anomalies have not been seen in massive clusters younger than $\sim$2 Gyr \citep{BastianLardo}. MPs show abundance spreads primarily in light elements and rarely, for example, in Fe which is associated with Type II supernovae, suggesting that hot hydrogen burning is needed \citep{DeniHart}. Exotic origin scenarios for MPs have been proposed, but to date all encounter serious problems \citep{BastianLardo, Renzini, Bastian2015}. Several different concepts have attempted to explain the origin of MPs by invoking special conditions in the early universe or special processes \citep{DErcole2016,Renzini}. A number of scenarios assume that the enriched second population of stars (2P) forms out of material released by particular members of a first population (1P) such as AGB stars, massive stars, or binaries. However, these models have difficulty producing enough material to create the second populations (the ``mass budget" problem), often cannot reproduce the observed chemical abundance patterns in detail, and require a time lag of anywhere from $5 - 100$ Myr between populations, which is not supported by observations of either GCs or young massive clusters \citep{Nardiello,BastianLardo,Martocchia}. To the limits of current measurements the 1P and 2P subpopulations have similar ages. Before discussing stellar populations in GCs, we briefly summarize the main features of the current observations and theory of cluster formation. The birth sites of star clusters are now known to lie within overdense regions, known as clumps, within Giant Molecular Clouds (GMCs). Observational studies of star clusters across three decades in cluster mass ($10^3-10^5 M_{\odot}$) indicate that they have a continuum of physical properties that indicates a common formation mechanism \citep[see, e.g., the review of][]{krumholz2018}. The velocity dispersion of gas within GMCs is supersonic \citep{Larson1981}, usually interpreted as evidence for supersonic turbulence. GMCs consist of networks of filamentary structures \citep[e.g.][]{Andre2014} and many simulations over the last two decades have shown that filaments are readily created by gas compression that occurs at the intersection of shock waves in these supersonic turbulent conditions (eg. review \citep{MaclowKlessen}. Stellar clusters are observed to form in the filaments, typically in clumps at the intersections (``hubs") of such systems of filaments \citep{Myers2009}. Moreover, gas accretion into forming star clusters occurs by observed filamentary flows \citep{Kirk2013}. GCs form in the most massive GMCs since such clouds have more massive clumps \citep{HarrisPudritz1994, Reina-CamposKruijssen2017}. As an example, \cite{Johnson2015} observed a clump within a massive GMC in the Antennae galaxies of the order of $5 \times 10^6 M_{\odot}$, in the right range to give birth to a young GC. Cluster formation is terminated by feedback, which for the most massive clusters, involves a variety of processes including radiative feedback (\cite{Dale2005, Murray2010}). Our previous numerical simulations, which include radiative feedback effects of the forming clusters on their host GMCs, showed that there is a universal scaling relation between the maximum cluster mass in a GMC, and the GMC mass: $M_{max} \propto M_{GMC}^{0.78}$, across three decades of GMC mass \citep{NatAst}. Thus, while earlier models of cluster formation viewed them as isolated entities each with their own peculiarities (open clusters, globular clusters, associations, etc), modern observations clearly indicate that they are parts of an extended hierarchical formation process that continues up to GMC scales and beyond. It is within the context of this observationally grounded, physical picture of cluster formation that we can now begin to address the nature of the stellar populations in GCs. The observational fact that MPs are found in many or most GCs suggests that they may be a byproduct of this normal mode of star cluster formation for the most massive clusters. The purpose of this paper is therefore to begin exploring whether or not MPs with realistic ranges of properties can plausibly emerge within this standard picture of cluster formation. To build a quantitative description, we start with our 3D radiative hydrodynamics (RHD) simulations of massive star-cluster formation within 10$^7$ M$_{\odot}$ GMCs \citep{Howard2017-2,NatAst}. In these, massive clusters grow rapidly (within $\simeq 5$ Myr) through an almost equal combination of rapid gas accretion from their natal GMC filaments, and merging with other protoclusters, and are certainly not isolated systems. The details of these radiation hydrodynamics simulations are laid out in \citet{NatAst}. These RHD simulations have not yet included any details of the chemistry or enrichment of the young stars inside the clusters. Into these simulations, we therefore add one extra ingredient that will allow enrichment of the intracluster gas, but in a way that will not require a large temporal spread between 1P and 2P formation. As an initial trial of such a model, we investigate two simple extreme cases. The first mechanism assumes that enriched material is added continuously during cluster growth at a rate that tracks the star formation rate within the forming cluster. This case will be referred to below as the \emph{continuous injection} (CI) alternative. The second mechanism assumes an instantaneous single addition of enriched material, which is referred to below as \emph{sudden injection} (SI). Most importantly, both mechanisms need to be active during the early stages of cluster formation before supernovae have cleared the intracluster gas, implying that the 1P and 2P stars form almost concurrently. Unlike previous scenarios, neither of these cases is viewed as happening within monolithic, isolated, single protoclusters; such objects are too idealized to represent real cluster formation and in any case lead to serious interpretive problems (see above). Instead, we use simple post-processing techniques within our RHD simulations of GMCs to track the He abundance $Y$ of the stars and gas within their young star clusters, where $Y$ is used as a proxy for the overall level of chemical enrichment. At this early stage of our modelling, the two opposite cases of enrichment rate that we calculate (CI and SI) are motivated by two general types of enrichment processes that have been discussed in the recent literature on He and various proton-capture elements that arise from the evolution of the most massive stars in clusters. The two opposite extremes that we calculate, as noted above, are intended to bracket the current uncertainties in the theory of massive star formation and evolution (see Section 4 below for additional discussion of some of the possibilities). In section 2, a short review of the RHD simulations of giant GMCs is provided, followed by the details for our computation of the internal enrichment of a model young massive cluster, under both of our extreme cases. In Section 3 the numerical results of each case are shown, along with brief comparisons with observations. In Section 4, we suggest some possibilities for the types of massive stars responsible for the enrichment. Finally, in Section 5 we give some additional discussion, a summary of the method and its advantages, and prospects for future work. | Overall, our simulations indicate that sudden-injection enrichment generally produces higher $\Delta Y$ and more discrete abundance distributions. On the other hand, gradual injection yields abundance distributions that are somewhat more continuously populated and with smaller $\Delta Y$. As discussed above, the SI route is less easily linked with a convincing or well understood stellar source. It is encouraging, however, that both routes are capable of producing realistic 1P/2P ratios, as well as abundance distributions for the 2P (enriched) population that can be complex and varied. In many real GCs, just two clearly distinct populations appear, as shown in the chromosome maps for the survey of 57 Milky Way GCs \citep{Milone2017,Milone2018}. But it is already clear that this is not a universal outcome. For many other GCs in the observational survey, 3 or more subpopulations are evident, and even cases of rather continuous abundance distributions with no clearly identifiable `gaps', within the limits set by the observational uncertainties. The enrichment model presented here is capable of producing this range of outcomes. In both cases our model allows us to place firm constraints on the total amount of He required to reproduce the range of observed stellar He spreads and population ratios. It can easily be generalized to models of other stellar processes that are characterized by injection of large quantities of He into star-forming gas over finite but short time intervals. The OSB approach would require a large fraction (25 - 50\%) of each O star's mass to be released as He, and it leaves behind $\Delta Y$ distributions that are usually not sharply bimodal. In the SI calculation, the results we have so far indicate that the (postulated) SMS at the centre of the cluster needs to be at least 3-7\% of the cluster's mass, corresponding to $\sim$1000-10,000 M$_{\odot}$ in a 10$^{5-6}$ M$_{\odot}$ cluster. Observations that can better quantify the He abundance distributions will be able to distinguish more strongly the different enrichment paths. In this paper, we have outlined a preliminary investigation of one particular route to producing multiple stellar populations within massive star clusters. At this point it is worth summarizing what we believe to be the major advantages of the approach: \begin{enumerate} \item{} Perhaps most importantly, this interpretation of MPs is built on a rigorous, quantitative RHD model for star cluster formation within GMCs. MPs are seen as emerging as a direct byproduct of normal cluster formation without supposing that GC formation occurred in a fundamentally different way than does lower-mass cluster formation. \item{} Both the 1P (pristine) and 2P (enriched) populations form actively and simultaneously with the first star formation within young massive clusters, explaining in a natural way why no detectable age difference between them should be seen in real GCs. \item{} The classic ``mass budget" problem is essentially avoided entirely because clusters are built within GMCs rather than starting as isolated monolithic gas clouds. Here, the entire GMC provides the much bigger reservoir of gas that a young cluster in formation can draw from through inflow along gaseous filaments. \item{} The complex assembly of clusters within GMCs has an in-built stochasticity, which means that different clusters can experience radically different growth histories \citep{NatAst} as expected in turbulent, filamentary clouds. Large differences in final chemical abundance patterns can therefore emerge between clusters quite naturally, but this variety itself is an important point of agreement with observations \citep{Renzini, Piotto, Milone2017}. The large cluster-to-cluster variety in the 1P/2P ratio and the degree of enrichment ($\Delta Y$) are also natural outcomes of this general model. \item{} More massive clusters are better able to hold on to their gas reservoirs during formation, so the relative numbers of 2P stars should increase with cluster mass, consistent with observations \citep{Milone2017}. \item{} Even in this simple form, the model is quantitative enough to constrain either the numbers and ejected He mass fractions of O-star binaries or (much more speculatively) the mass range of SMSs. \end{enumerate} The modelling route presented in this paper explores two opposite limiting cases (smooth, continuous enrichment versus a large one-time event). We regard this work as a first promising step that can be developed further in a number of ways. For example, in practice young massive star clusters will certainly hold a significant population of O stars in close binaries, but an SMS could also be present, so a combination of the two mechanisms could be explored. A true SMS may shed mass more continuously through its short lifetime \citep{Gieles2018}, so a delta-function pulse is certainly simplistic. The true enrichment rate ($dY/dt$) need not be either artificially smooth nor a delta-function, given the stochastic nature of the gas inflow, merging with other protoclusters, and random sampling of the IMF. Further work will also need to be done to compute more detailed abundance patterns of the light elements that are represented here only by the He enrichment. RHD modelling of the young clusters also needs to be continued past 5 Myr through the supernova stage to track their survival. In addition, the model GMCs have relatively high star formation efficiencies \citep{NatAst} that might be reduced to more realistic levels by including several additional processes (winds, supernovae, and magnetic fields). This step may in turn reduce the net $\Delta Y$ yields, which would bring more of the model clusters into the observationally valid range. Finally, on the computational frontiers, new techniques may allow RHD simulations to be performed at sufficient resolution to resolve individual star formation in the young clusters themselves. This will open up new capabilities in understanding what happens to the gas reservoir held by a protocluster and ultimately how realistic our general model can be. | 18 | 8 | 1808.07081 |
1808 | 1808.00945_arXiv.txt | Over half of all observed hot subdwarf B (sdB) stars are found in binaries, and over half of these are found in close configurations with orbital periods of 10$ \,\rm{d}$ or less. In order to estimate the companion masses in these predominantly single-lined systems, tidal locking has frequently been assumed for sdB binaries with periods less than half a day. Observed non-synchronicity of a number of close sdB binaries challenges that assumption and hence provides an ideal testbed for tidal theory. We solve the second-order differential equations for detailed 1D stellar models of sdB stars to obtain the tidal dissipation strength and hence to estimate the tidal synchronization time-scale owing to Zahn's dynamical tide. The results indicate synchronization time-scales longer than the sdB lifetime in all observed cases. Further, we examine the roles of convective overshooting and convective dissipation in the core of sdB stars and find no theoretical framework in which tidally-induced synchronization should occur. | Hot subdwarf B (sdB) stars are compact sub-luminous stars. They have surface temperatures between $20\,000$ and $40\,000\,\rm{K}$ and surface gravities $5<\log_{10}(g_{\rm{surf}}/\rm{cm\,s^{-2})}<6$. The sdBs were first observed by \cite{discoveryzwicky} and their spectra were quantified by \cite{specdefinition}. The stars are helium core burning with low-mass hydrogen envelopes. Typically the stars spend around $150$ $\rm{M \, yrs}$ in their He burning phase. They are thought to be the cores of red giant branch (RGB) stars exposed by close binary-star interaction \citep{han1}. One of the proposed mechanisms for sdB formation is common envelope ejection. The sdBs produced in this manner are in binary systems with orbital periods less than $10\,\rm{d}$. Observations suggest that about half of the observed sdB systems lie in such configurations \citep{napiow04binaryf,binaryfraction2011}. The close sdB binaries are spectroscopically single-lined with either white dwarf (WD) or low-mass main-sequence dM companions. Eclipsing post-common envelope sdBs with a dM companion are referred to as HW Vir type systems. Unless it is eclipsing it is generally not possible to find the inclination of a system. This means it can be difficult to estimate the component masses. By assuming tidal synchronization, a spectroscopic measurement of the projected rotation velocity and an assumed radius constrains the rotational period, the orbital inclination, and hence the companion mass. \cite{methodtsyncfirst} first applied this method to the subdwarf O star HD49798. \cite{geier2010} further applied the same technique to a sample of 51 close sdB stars. The fact that so many sdBs are in close binaries makes them an ideal test bed for tidal dissipation theories. These theories have always been controversial for stars with convective cores and radiative envelopes such as sdBs. Two competing theoretical prescriptions for dissipation in such stars are given by \cite{zahn75,zahn77} and \cite{tassoul}. \cite{zahnbad} demonstrated that the dynamical tide proposed by \cite{zahn77} is too inefficient to describe the observed level of synchronization of some early main-sequence spectroscopic binaries, particularly when the fractional radius of the convective region is below 0.05. \cite{tassoul} address this efficiency issue by suggesting that pumping across the Ekman boundary provides a mechanism for tidal dissipation. \cite{tassoulbad} dispute the physical validity of Tassoul's mechanism. In Section 2 we review the current observations and previous calculations of tidal synchronization. In Section 3 we address the methods used for this paper, first by reviewing tidal theory, then by discussing the numerical methods and finally the stellar models. In Section 4 we present our results and Section 5 concludes. | The goal of this study was to find synchronization time-scales for short period sdB binary systems. A grid of sdB models was created with the STARS code for a variety of progenitor masses, envelope masses and treatments of convection. Previous studies have predominantly used Zahn's theory of dynamical tides with a scaling from main-sequence models to find the synchronization times. Recalculating the tidal coefficient $E_2$ for the grid of sdBs shows scaling from main-sequence models overpredicts $E_2$ by a factor of at least 3000. The synchronization time-scales should be several orders of magnitude longer. As a result, estimates of Zahn's dynamical tide synchronization time-scales are longer than EHB lifetimes, even for the extreme case of \cdthirty. The sdB stars have convective cores which provide a mechanism for tidal dissipation. By solving Clairaut's equation the tidal synchronization times owing to turbulent convection have been calculated. Initial calculations of the convective tides predicted that the three sdB systems with the most massive WD companions should be synchronized. Closer examination revealed that the orbital period is typically shorter than the convective turnover time. This causes the convective dissipation of the tides to be damped and become substantially less efficient. The damping coefficient depends on the turnover time for viscous elements within the star and is calculated with mixing length theory. The damping factor causes estimates of synchronization time-scales to increase by several orders of magnitude so that no sdB binary systems are conclusively predicted to be synchronized. Traditional mixing length theory predicts a singularity at the stellar centre. The effects on tidal synchronization time-scales when the mixing length was altered to remove this singularity were examined. Reducing the mixing length to avoid the central singularity generally increases the synchronization time because the estimated viscosity decreases. The optimal case for tidal dissipation is to reduce the mixing such that the convective turnover time is slightly shorter than the orbital period so that the tidal dissipation is not damped. Even in this case synchronization is not achieved because the viscosity is substantially reduced and the tidal interactions are less efficient. The rotational periods of sdB stars at the TAEHB were calculated to investigate the impact of the tides. The models with the optimally chosen mixing length and with envelope masses less than $0.01\,\rm{M_\odot}$ are most substantially affected by the tides. The convective region accounts for a larger fractional volume in the sdBs with the lowest mass envelopes so tides are more effectively dissipated. With the theoretical framework presented, tidal synchronization times for EHB stars are long, but not excessively so, compared with nuclear lifetimes. With evidence from asteroseismology that convective core sizes may be larger than those predicted by classical convection theory, and with the possibility that the tides could induce differential rotation with the EHB star, these avenues of exploration still open. | 18 | 8 | 1808.00945 |
1808 | 1808.02347_arXiv.txt | With the recent discoveries of terrestrial planets around active M-dwarfs, destruction processes masking the possible presence of life are receiving increased attention in the exoplanet community. \\ We investigate potential biosignatures of planets having Earth-like (N$_2$-O$_2$) atmospheres orbiting in the habitable zone of the M-dwarf star \adl. These are bombarded by high energetic particles which can create showers of secondary particles at the surface. % We apply our cloud-free 1D climate-chemistry model to study the influence of key particle shower parameters and chemical efficiencies of \nox\ and \hox\ production from cosmic rays. We determine the effect of stellar radiation and cosmic rays upon atmospheric composition, temperature, and spectral appearance. % Despite strong stratospheric \oz\ destruction by cosmic rays, smog \oz\ can significantly build up in the lower atmosphere of our modeled planet around \adl\ related to low stellar UVB. \nto\ abundances decrease with increasing flaring energies but a sink reaction for \nto\ with excited oxygen becomes weaker, stabilizing its abundance. \meth\ is removed mainly by Cl in the upper atmosphere for strong flaring cases and not via hydroxyl as is otherwise usually the case. Cosmic rays weaken the role of \meth\ in heating the middle atmosphere so that H$_2$O absorption becomes more important. We additionally underline the importance of HNO$_3$ as a possible marker for strong stellar particle showers.\\ In a nutshell, uncertainty in \nox\ and \hox\ production from cosmic rays significantly influences biosignature abundances and spectral appearance. | \label{sec:introduction} Cool M-dwarf stars are favored targets in exoplanetary sciences due to their high abundance in the Solar neighborhood, a close-in Habitable Zone (HZ), hence short orbital periods, and a high planet/star contrast. For an overview see e.g. \citet{kasting1993,scalo2007, shields2016}. There are however drawbacks. Planets lying in the close-in HZ could be tidally-locked (e.g. \citet{selsis2000,kasting1993}) and could be bombarded by high levels of energetic particles \citep{griessmeier2005}. An additional drawback for M-star planet habitability is the long, bright, pre-main-sequence phase of the parent star, which may devolatilize planets that would later reside in their habitable zones (e.g. \citet{luger2015, ramirez2014, tian2015}). Nevertheless, planets in the HZ of M-dwarf stars could represent the first opportunity to detect atmospheric properties and even biosignatures of rocky extrasolar planets. There are numerous relevant model studies e.g. in 1D \citep{segura2003,segura2005,segura2010,rugheimer2015,kopparapu2013,grenfell2012,tabataba2016} and in 3D \citep{shields2013,shields2016,leconte2013,yang2014,godolt2015,kopparapu2016}. Interpretation of such potential future observation heavily relies on our detailed understanding of atmospheric physical, chemical, and biological processes and their interaction with different electromagnetic radiation and High Energetic Particles (\hep), such as Galactic Cosmic Rays (\gcr s) and Stellar Energetic Particles (\sep s). The latter has received only limited attention in the exoplanets community so far. Our general understanding of the redistribution of incoming \hep s into secondary particles in so-called air showers through the atmosphere and their influence upon atmospheric chemistry dates back to theoretical work done in the early 1980s \citep{rusch1981, solomon1981}. A more recent study by \citet{airapetian2016} investigated the production of \nto\ via \sep s for the early Earth. While our own star is comparably quiescent, many M-dwarfs show high activities, in which flares with energies comparable to the devastating Carrington event on Earth in the 19th century regularly occur up to a few times a day and orders of magnitude higher energetic events have been observed, e.g. for the here studied M-dwarf \adl\ \citep{atri2017, hawley1991}.\\ Recently, potentially Earth-like planets have been found in the HZ around M-dwarfs (Proxima Cen b, LHS1140 b, and TRAPPIST-1 d-f) which may be studied in further detail with upcoming instrumentation on e.g. the James Webb Space Telescope (\jwst ) \citep{gardner2006} and E-ELT \citep{kasper2010}. The impact of \hep s upon atmospheric chemistry needs further investigation. \hep -induced ion-pairs react with molecular oxygen, molecular nitrogen, and water to cascade into nitrogen oxides (\nox =N+NO+NO$_2$+NO$_3$) and hydrogen oxides (\hox =H+HO+HO$_2$) \citep{porter1976,rusch1981,solomon1981}. \nox\ and \hox\ catalytically destroy ozone (\oz ) in the lower and upper stratosphere respectively \citep{crutzen1970} but can form \oz\ in the troposphere due to the so-called smog mechanism \citep{haagen1952}. They are stored and released from reservoir molecules such as HNO$_3$, depending on e.g. UV radiation and temperature. Recent model studies have shown that \hep\ induced \nox\ and \hox\ from particle showers can indeed significantly reduce \oz\ in an Earth-like atmosphere \citep{grenfell2012, griessmeier2016, tabataba2016}. Production rates of around 1.27 \nox\ \citep{porter1976, rusch1981} and 2.0 \hox\ \citep{solomon1981} per ion pair have been assumed in numerous atmospheric studies, but recent ion-chemistry studies by e.g. \citet{sinnhuber2012} and \citet{verronen2013} have pointed out that the uncertainties in these complex chemical coupling coefficients might be under-estimated, especially when additionally taking into account negative ion chemistry. When conducting numerical studies of rocky planets around active M-dwarfs, such uncertainties can have a major impact on atmospheric abundances - including species influenced by biogenic processes. \oz, for example is removed catalytically by \hox\ and \nox. Also, \meth\ is usually removed by OH, a member of the \hox\ family.\\ Based on the above, the main motivation of this work is to compare the influence of different M-dwarf stellar flaring energies to that of the uncertainties in atmospheric \nox-\hox\ production efficiencies from incoming \sep s and show their impact on overall climate and spectral features in transit observations. In Section \ref{sec:methods} we describe the models used for this work and motivate the modeled scenarios, in Section \ref{sec:results} we briefly describe our results, before discussing and comparing them to other relevant works in Section \ref{sec:discussion}. Finally, in Section \ref{sec:conclusion} we draw our conclusions. | \label{sec:conclusion} We have performed atmospheric simulations of virtual Earth-like planets around the flaring M-star \adl\ with our cloud-free 1D climate-chemistry model and have compared the influence of flaring strength with the uncertainty ranges of chemical \nox-\hox\ production efficiencies. \\ New chemical insights found in this work are: \begin{itemize} \item \textbf{\nox-\hox:} The chemical production efficiencies f$_{\nox}$ and f$_{\hox}$ can significantly influence biosignature chemistry and abundances in our model, as well as stratospheric temperatures, and are therefore potentially important for Earth-like planets around M-dwarf stars like \adl. In the Earth's atmosphere, on the other, the influx of \sep s has a much stronger effect than f$_{\nox}$ and f$_{\hox}$, which makes the empirical determination of the latter challenging. \item \textbf{HNO$_3$:} Spectroscopic transit measurements of exoplanets may be able to help constrain their stellar environments by looking at e.g. HNO$_3$ features above 10 microns together with infrared \oz, H$_2$O and \meth\ absorption bands. Especially the measurement of the HNO$_{3}$ features at 17 and 21 microns would hint towards high f$_{\hox}$ production. \item \textbf{Cl:} We introduce and discuss a change of the major \meth\ sink in the stratosphere from OH (lower stratosphere) to Cl (upper stratosphere). This may also become important for worlds with e.g. high volcanic chlorine emissions. \item \textbf{\oz:} We show that on Earth the UVB radiation from the Sun (G-star) is sufficient to limit global tropospheric smog \oz\ abundances even in hypothetical high flaring Sun scenarios, while we confirm lower atmospheric build-up of \oz\ for Earth-like planets around active M-stars like \adl, as has been modeled in multiple studies e.g. \citet{segura2005,grenfell2012,tabataba2016}. \item \textbf{\nto:} Atmospheric \nto\ abundance runs into 'saturation' for flaring cases regardless of stellar spectrum, flaring strengths, or stratospheric \oz\ levels. \nto\ reactions e.g. with O($^1$D) in addition to diffusion processes within the atmosphere counteract the \oz\ - UV - \nto\ coupling (See Fig. \ref{Biosignatures-Diagram.pdf}). Hence, destruction of \nto\ by cosmic rays is ineffective in our model. \end{itemize} Additionally, in our model OH is the major sink for \meth\ in the lower to mid atmosphere and is directly produced by \sep s, but we find that around high flaring solar-like stars atmospheric \oz\ abundances can significantly drop, which itself is a major source of tropospheric OH production (see Fig. \ref{Biosignatures-Diagram.pdf}). This lack of OH from \oz\ can outweigh OH production from \sep s, subsequently causing unexpectedly high \meth\ abundances (see figures \ref{AllStars.pdf} and \ref{Biosignatures-Diagram.pdf}). Furthermore, in absence of other sources HNO$_2$ can become the main OH source throughout our whole model atmosphere for high flaring host star cases. Further work on this is needed to see for which range of planetary atmospheres HNO$_{2}$ may become important.\\ We would like to emphasize once more that \nox\ and \hox\ produced by cosmic rays can become important when studying Earth-like atmospheres around active M-stars. | 18 | 8 | 1808.02347 |
1808 | 1808.05421_arXiv.txt | The \OII\ 3726$+$3728\AA\ emission line doublet is often used to estimate star formation rates within the host galaxies of active galactic nuclei (AGN), as it is known to be strongly excited by star formation, but is only weakly excited in the broad and narrow line regions of AGN. However, within AGN host galaxies, \OII\ can also be excited in low-density gas located at appreciable distances from the nucleus, but still ionized by the AGN. These AGN extended emission line regions (EELRs) can contribute significant flux to integrated spectra, even in the presence of luminous AGN. Here, we identify EELRs by the presence of the \NeV\ 3426\AA\ emission line, which, like \OII, is not strongly excited in the inner regions of AGN, but is a prominent emission line in the lower density EELRs. Critically, unlike \OII, \NeV\ is not excited by star formation. Therefore, when strong \NeV\ is present in an AGN spectrum, the flux from the EELR is not negligible, implying the \OII\ flux is contaminated by emission from the EELR, and is not a good measure of star formation. After removing objects with EELRs identified by \NeV, the \OII\ flux in the host galaxies of radio-loud AGN is found to be higher than that within radio-quiet AGN, which could either indicate higher star-formation rates, or the presence of moderate-velocity shocks. Being mindful of EELRs for upcoming large-area spectroscopic surveys, particularly those tied to radio continuum surveys, will be important for determining star formation rates in AGN host galaxies. | \label{sec:Introduction} The current star formation rate (SFR) is an important observational metric of a galaxy. It tells us the evolutionary state by indicating whether a galaxy is actively building its stellar mass, or if it is passively ageing. Star formation (SF) can also be triggered by, for example, mergers and/or accretion of gas, thus providing information about a galaxy's environment. The SF history is a crucial underlying factor in establishing the diversity of galaxies we observe throughout the universe. There are many observables at a variety of wavelengths that serve as calibrated proxies for the current SF in galaxies. These include optical emission lines such as \Halpha, and continuum flux at ultra-violet, far-infrared and radio wavelengths. A comprehensive comparison of various SF tracers can be found in \citet{Hopkins2003}. The relationship between the flux in the \OII\ 3726+3728\AA\ emission line doublet (hereinafter referred to as \OII) and the SFR has been calibrated by a number of authors (e.g. \citealt{Kennicutt1998}, \citealt{Hopkins2003}, and \citealt{Kewley2004}). Although it suffers from dust extinction due to the relatively short emission wavelength, \OII\ is very useful for estimating SFRs, as it is a strong, easily identified feature and can be seen in moderate resolution optical spectra out to high ($z<1.5$) redshifts. Complicating the measurement of SF in some galaxies is the presence of active galactic nuclei (AGN), as the flux from the AGN outshines the stellar component, at some wavelengths by orders of magnitude. In order to measure SF in AGN host galaxies, we require a quantity whose production is dominated by SF rather than AGN. As shown in the quasar composite of \citet{VandenBerk2001}, \OII\ is only weakly excited in the narrow line region (NLR) of AGN, and not at all in the broad line region (BLR). However, it is seen in emission from even moderate levels of star formation. It has therefore been used to estimate the SFR in quasar host galaxies (e.g. \citealt{Ho2005}, \citealt{Kim2006}, \citealt{Kalfountzou2012}, \citealt{Matsuoka2015}, and \citealt{Vergani2017}). \subsection{Extended Emission Line Regions} AGN are sometimes seen to ionize gas in the host galaxy beyond the BLR and NLR. While the NLR is typically confined to within $<$1 kiloparsec (kpc) of the nucleus, extended emission line regions (EELRs) \footnote{We use the term EELR to describe extended emission beyond the NLR which, based on emission line diagnostics, is shown to be gas ionized by the AGN. We do not use the sometimes quoted alternative term `extended narrow line region' (ENLR), to reinforce the fact that the gas is distinct from the NLR.} can extend throughout the entire host galaxy, spanning tens of kpcs in some cases (e.g. \citealt{Fosbury1982}, \citealt{Spinrad1984}, \citealt{Stockton1987}, \citealt{Fu2009a}, \citealt{Villar2011}, \citealt{Husemann2013}, \citealt{Liu2013a}, \citealt{Liu2013b}, \citealt{Liu2014}, \citealt{Harrison2014}, \citealt{Husemann2014}). EELRs are spatially and kinematically distinct from the classic AGN NLRs. The gas within an EELR can show velocity differences of several hundreds, to $>$1000 \kms\ with respect to the systemic velocity of the galaxy, but also very low velocity dispersion (\citealt{Fu2009a}, \citealt{Husemann2013}). Early studies employing long-slit spectroscopy focused on the EELRs within the host galaxies of bright radio sources (\citealt{Spinrad1984}, \citealt{Unger1987}), investigating possible triggering mechanisms for the nuclear activity, including mergers with gas-rich galaxies. However, EELRs are also now known to be observed in AGN hosts with no appreciable nuclear radio emission and no indications of recent merger activity (e.g. \citealt{Villar2011}, \citealt{Husemann2013}). Integral field unit (IFU) spectrographs with wide spectral range and substantial (tens of arcseconds) fields-of-view have been very useful for studying the kinematics and spatial extent of EELRs. \citet{Fu2009a} observed eight galaxies known to have extended emission, covering a wide wavelength range, to investigate EELR clouds on small spatial scales. The ratios of their detected emission lines all indicate AGN as the source of the ionizing flux, rather than star formation. Similarly, \citet{Husemann2013}, \citet{Liu2013a}, \citet{Harrison2014} and \citet{Husemann2014} also use IFU observations to isolate regions throughout AGN host galaxies where the gas has been ionized by the AGN. Many EELR studies focus on optical emission lines such as \OIII\ 5008\AA\ and \Halpha. However, for studies with wider wavelength coverage, or observations of objects at higher redshifts, shorter wavelength lines are also visible throughout the EELR, including strong \OII. The lower density of the interstellar gas, of a few hundred cm$^{-3}$ (\citealt{Husemann2014}), with respect to the order of magnitude higher density of the AGN NLR, enables \OII\ emission, which has a low critical density, to arise within the EELR (\citealt{Villar2011}). \citet{Fu2009b} observed a subset of their EELR galaxies in the infrared (IR), but found no signatures of SF in the hosts, thus confirming the \OII\ flux detected in the IFU spectra should be attributed to the EELR, rather than SF. As both SF and AGN EELRs independently result in strongly excited \OII\ throughout the full extent of galaxies via unrelated mechanisms, we require some diagnostic to distinguish between the different ionizing sources. While emission line ratios can be broadly used to determine the source of the ionizing radiation (i.e. AGN or SF), the most commonly used spectral features, including \Halpha, are shifted out of the optical spectroscopic range by moderate, $z\sim 0.4$, redshifts. Fortunately, there are specific emission lines which are weak in both the high-density medium surrounding AGN as well as star-forming regions. One such emission line is \NeV\ 3426\AA, which, due to its high ionization potential (97 eV, compared to 13.6 eV for \OII, \citealt{Lide2005}) is not excited by star formation, but as with \OII, is very weak in AGN spectra \citep{VandenBerk2001}. Indeed, for observations of EELRs with spectral coverage of these short wavelengths, \NeV\ is often seen as well (e.g. \citealt{Spinrad1984}, \citealt{Fu2009a}). Here we investigate the subset of AGN which, in addition to emission lines common in either SF or AGN spectra, also show strong \NeV\ in their spectra, and argue that it indicates the presence of an EELR, or in some cases, extreme shocks. Since EELRs also strongly excite \OII, the \OII\ in these objects should not be used as a proxy for SF, since the flux in this emission line is contaminated. In Section~\ref{sec:ratios} we investigate the diagnostic power of emission line ratios in determining the properties of the ionizing radiation, and where objects with detected \NeV\ emission lie in this parameter space. Section~\ref{sec:case} is a case study of star formation in SDSS quasar host galaxies. A discussion and conclusions are presented in Section~\ref{sec:discussion}. Concordance cosmology with $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$ (thus $h\equiv H_{0}$/[100 km s$^{-1}$ Mpc$^{-1}$]$=0.7$), $\Omega_{m} = 0.3$, $\Omega_{\Lambda} = 0.7$ is assumed throughout. We use laboratory wavelengths for emission lines, taken from Table 2 in \citet{VandenBerk2001}. | \label{sec:discussion} We have shown that strong \NeV\ emission detected in individual spectra is an indicator of the presence of extended emission excited by an AGN, either via photoionization, or in some extreme cases, via high-velocity AGN-driven shocks. This emission line is not strongly excited within the usual AGN broad or narrow line regions, or by star formation. Table~\ref{tab:sources} summarizes the relevant sources of ionizing radiation, and whether they excite \OII\ and \NeV. As \OII\ is also excited within EELRs, we conclude that for objects with \NeV\ emission, the \OII\ flux is contaminated and cannot be used as a measure of SF. \begin{table} \centering \caption{Sources of ionizing radiation within AGN host galaxies, and whether \OII\ and \NeV\ are strongly excited by each. $^{\star}$\NeV\ is only excited by shocks at high ($>$600\kms) velocities.} \label{tab:sources} \begin{tabular}{lcc} \hline Source & \OII & \NeV \\ \hline SF & Strongly & Weakly \\ AGN BLR & $\times$& $\times$ \\ AGN NLR & Weakly & Weakly \\ AGN EELR & Strongly & Strongly \\ Shocks & Strongly & Strongly$^{\star}$ \\ \hline \end{tabular} \end{table} The fraction of SDSS DR7 quasars removed due to \NeV\ contamination ranges from 70~per~cent at low redshift, to 20~per~cent and 10~per~cent for RL and RQ quasars, respectively, at $z=1.3$. After removing objects with \NeV\ detected in individual spectra, stacking the remaining spectra reveals a \OII/\NeV\ flux ratio elevated above that expected from pure AGN emission for the RL quasars, indicating an excess of \OII\ emission within these objects. The excess of \OII\ flux in RL quasars can be attributed to star formation, but may also be partly due to moderate-velocity shocks from the radio jets acting within the host galaxy, which do not excite \NeV. Therefore, the SFRs shown in Fig.~\ref{fig:stackedSF} are upper limits. The fact that the EELR diagnostic line \NeV\ is so close to \OII\ in wavelength is particularly fortuitous, as they are both visible over the same redshift range, and when taken as a ratio, it does not suffer severely from dust reddening, even though they are in the rest-frame $u$-band. If there is appreciable amounts of dust in the quasar host galaxies, it will serve to decrease the strengths of both the \OII\ and \NeV. Therefore, there may be some objects whose \NeV\ has been attenuated below the threshold for detection in individual spectra, in which case these objects were not removed from the stacked samples used for Fig.~\ref{fig:stackedEW} and Fig.~\ref{fig:stackedSF} and the \OII\ measurements are still contaminated. However, as the SDSS quasars are optically selected, the dust reddening of individual objects is low (\citealt{Maddox2012}), thus we expect this effect to be small. A number of large-area radio continuum surveys are either already underway, or are about to come online. In the Northern hemisphere the Low-Frequency Array (LOFAR, \citealt{vanHaarlem2013}) is surveying 2$\pi$ steradians of the northern sky (\citealt{Shimwell2017}). Spectroscopy for this survey will be supplied by the new spectrograph to be installed on the William Herschel Telescope (WHT), the WHT Enhanced Area Velocity Explorer (WEAVE, \citealt{Dalton2014}), providing redshifts for radio continuum-selected sources (\citealt{Smith2016}). Covering the Southern hemisphere, the Evolutionary Map of the Universe (EMU, \citealt{Norris2011}) will be undertaken with the Australian SKA Pathfinder (ASKAP, \citealt{Johnston2008}). This will be coupled with the spectroscopic redshift survey Taipan (\citealt{daCunha2017}). While the SFR for radio-detected galaxies can be determined directly from the radio flux, computing SFR from \Halpha, and to higher redshifts, \OII, will serve as a useful check. In addition, the radio flux can not be used as a SFR indicator for radio-loud quasars, leaving \Halpha\ and \OII\ as the only SF diagnostics available from planned large-area spectroscopic surveys. Other calibrated SF indicators, such as flux at far-infrared (FIR) wavelengths, would require additional observations not currently planned. The importance of understanding and identifying the contribution from EELRs to the \OII\ emission line, as measured by the presence of \NeV, is essential for these projects. | 18 | 8 | 1808.05421 |
1808 | 1808.05617_arXiv.txt | Nearby galaxy surveys have long classified X-ray binaries (XRBs) by the mass category of their donor stars (high-mass and low-mass). The \nustar\ observatory, which provides imaging data at E $>10$ keV, has enabled the classification of extragalactic XRBs by their compact object type: neutron star (NS) or black hole (BH). We analyzed \nustar/\chandra/\xmmn\ observations from a \nustar-selected sample of 12 galaxies within 5 Mpc having stellar masses ($M_{\star}$) $10^{7-11}$ \msun\ and star formation rates (SFR) $\approx0.01-15$ \sfr. We detect \numsrc\ \nustar\ sources to a sensitivity of $\approx10^{38}$ \es. Using \nustar\ color-intensity and color-color diagrams we classify \numns\ of these sources as candidate NS and \numbh\ as candidate BH. We further subdivide BH by accretion states (soft, intermediate, and hard) and NS by weak (Z/Atoll) and strong (accreting pulsar) magnetic field. Using 8 normal (Milky Way-type) galaxies in the sample, we confirm the relation between SFR and galaxy X-ray point source luminosity in the \full and \hard keV energy bands. We also constrain galaxy X-ray point source luminosity using the relation $L_{\rm{X}}=\alpha M_{\star}+\beta\text{SFR}$, finding agreement with previous work. The XLF of all sources in the \full and \hard keV energy bands matches with the $\alpha=1.6$ slope for high-mass XRBs. We find that NS XLFs suggest a decline beginning at the Eddington limit for a 1.4 \msun\ NS, whereas the BH fraction shows an approximate monotonic increase in the \full and \hard keV energy bands. We calculate the overall ratio of BH to NS to be $\approx1$ for \full keV and $\approx2$ for \hard keV. | \label{sec:intro} Until the launch of the first focusing telescope to operate at E $>10$ keV, the \nustarf\ \citep[\nustar;][]{harrison06-13}, we knew very little about the behaviour and nature of extragalactic black hole (BH) and neutron star (NS) populations at harder energies. In the absence of an X-ray bright supermassive BH, the total X-ray emission of a galaxy above 2 keV is dominated by X-ray binaries (XRBs), classified as low-mass (LMXB) or high-mass (HMXB) based on their donor star. Previous studies of nearby galaxies in the soft X-ray band ($0.5-10$ keV) by, e.g.\ \chandra\ and \xmmn\ \citep[e.g.][]{stiele10-11,mineo01-12,mineo01-14,long06-14,haberl02-16,peacock02-16} have revealed important new information on compact object populations, such as strong correlations between properties of XRBs and galaxy star formation rate (SFR), stellar mass, and metallicity \citep[e.g.][]{basu-zych02-16}. Extrapolation of these local-Universe measurements as well as supporting measurements at high-redshift \citep{lehmer07-16} have indicated a possible significant role of XRBs in heating the Intergalactic Medium (IGM) of the early Universe \citep[e.g.][]{fragos10-13, mesinger04-14, pacucci09-14, madau05-17, sazonov06-17, das07-17}. However, there are questions about the extragalactic XRB population that are difficult to answer at E $<10$ keV, including whether compact objects are BH or NS. The rich suite of thousands of \rxte\ (\rxtet) PCA spectra of BH/NS XRBs in the Milky Way galaxy provide critical diagnostics in the \full keV band of both compact object type (BH vs. NS) and accretion state \citep[e.g.][]{maccarone01-03, mcclintock04-06, done12-07}. With \nustar, for the first time, we are able to leverage the knowledge gained from compact objects in our own galaxy by applying these harder X-ray diagnostics to extragalactic populations. The hard X-ray coverage with \nustar\ is crucial for distinguishing different types of accreting binaries, such as BH/NS XRBs and accreting pulsars. Compact object diagnostics have already been successfully applied to characterize XRBs in several nearby galaxies observed by \nustar. These studies include simultaneous {\em NuSTAR/Chandra/XMM-Newton/Swift} studies of the nearby star-forming galaxies NGC 253 \citep{lehmer07-13, wik12-14} and M83 \citep{yukita06-16}, as well as Local Group galaxy M31 \citep{maccarone06-16, yukita03-17, lazzarini06-18}; for a description of the \nustar\ galaxy program please see \citet{hornschemeier-16}. Using \full keV color-color and color-intensity diagnostics, these studies have shown that the starburst galaxies are dominated by luminous BH-XRB systems, mostly in intermediate accretion states. Specifically, ultraluminous X-ray sources (ULXs) with $3-30$ keV spectra indicative of super-Eddington accretion (e.g.\ \citealt{gladstone08-09}) appear to dominate the hard X-ray emission of starburst galaxies \citep{walton12-13, bachetti10-14, rana02-15, lehmer06-15}. Meanwhile, M31 has a significant contribution from NS accretors (pulsars and low-magnetic field Z-type sources; \citealt{maccarone06-16, yukita03-17}). As expected, the pulsars trace the young stellar population in the spiral arms and the Z-type sources are concentrated in globular clusters and the bulge/field of the galaxy. \nustar\ data were crucial to the reclassification of previously identified BH candidates in M31 globular clusters as NS, based on their hard X-ray spectra \citep{maccarone06-16}. \nustar\ has previously resolved the XRB population in 3 galaxies. Thus, it is now time for a broader investigation of the relationship between the properties of a galaxy and the X-ray source types and accretion states as determined from hard X-ray observations. Specifically, what is the relationship between galaxy properties such as the stellar mass and recent star formation rate/history and compact object type/accretion state as determined from hard X-ray diagnostics? To estimate the number of BH and NS that will be formed in a galaxy requires binary population synthesis, and a detailed understanding of concepts such as supernova explosions, which is not well understood \citep[e.g.][]{pejcha03-15}. Alternatively, we can use observational data and methods to determine the BH fraction and its dependence on X-ray luminosity and specific star formation rate (sSFR). With \nustar\ we can measure local-galaxy SEDs over $0.5-30$ keV that are applicable to high-$z$ galaxies detected by \chandra. One of our goals is to determine what sources are contributing to the $0.5-30$ keV emission. Furthermore, we would like to be able to predict, based on galaxy properties such as star formation rate/history and stellar mass, what the distribution of binaries and their emitting properties are. Achieving this goal is rather complicated, as there are parameters such as the duty cycle that result in a broad range of population properties for different stellar ages, etc. One approach to this complicated problem is to make direct measurements over a variety of galaxy properties. Each snapshot view of an individual galaxy measures the state of the overall population, giving us a constraint on duty cycles \citep{binder01-17}. Hard X-ray diagnostics allow us to determine the distribution of BH spectral states, similar to Galactic BH studies \citep[e.g.][]{tetarenko02-16}. Using this approach, we can obtain baseline estimates of XRB formation rate, duty cycles, spectral states, and galaxy SEDs. Understanding these properties at E $>10$ keV is critical to compare to the results of XRB evolution in the $0.5-10$ keV bandpass. \nustar\ is well-matched to the rest-frame energies of high-$z$ galaxies at $z=3-4$ probed by \chandra\ and is thus a new window into XRB evolution. The X-ray luminosity function (XLF) represents the distribution of sources in a galaxy based on their luminosity. Seminal studies of LMXBs in elliptical galaxies \citep[e.g.][]{gilfanov03-04, zhang10-12} and HMXBs in spiral galaxies \citep[e.g.][]{grimm03-03, mineo01-12} found that their XLFs were (approximately) universal when normalizing by the stellar mass and SFR of a galaxy, respectively (see \citealt{gilfanov-04} for a summary). Small variations in the power law slope and cutoff are dependent on factors such as metallicity \citep{basu-zych02-16} and star formation history \citep{lehmer12-17}. We will investigate how scaling \nustar\ XLFs by SFR compares with results from \chandra/\xmmn\ studies. To date, studies of the XLFs of nearby galaxies have mostly focused on LMXB or HMXB populations. However, certain XLF characteristics can be attributed to compact object types \citep{lutovinov05-13}, such as the break at $\sim$few$\times10^{38}$ \es\ corresponding to the Eddington limit for NS. This break is often argued to reflect the transition from a population of NS to BH XRBs \citep{sarazin12-00, kim08-04, wang09-16}. \nustar\ is well-suited to distinguish between BH and NS accretors, therefore allowing a first-look at BH-only and NS-only XLFs. In addition, this can elucidate how the $0.5-30$ keV SED of galaxies depends on the compact object type and accretion states of BH and NS. Our goals are to study the hard X-ray properties of the XRB population of 12 nearby galaxies ($<5$ Mpc) using joint \nustar\ and \chandra/\xmmn\ data. We will use knowledge of galaxy parameters such as SFR and stellar mass to investigate the connection between XRB populations and host galaxy properties. In Section \ref{sec:sample} we describe the sample selection and calculation of SFR and stellar mass for galaxies in the sample. In Section \ref{sec:obs} we summarize the \nustar, \chandra, and \xmmn\ observations. In Section \ref{sec:data} we outline our analysis methods, which focus on the PSF fitting procedure for \nustar\ data. In Section \ref{sec:resdis} we present \nustar\ diagnostic diagrams, XLFs, and scaling relations, and discuss their implications. We summarize our conclusions in Section \ref{sec:con}. | \label{sec:con} Using a \nustar-selected sample of 12 late-type and dwarf galaxies, we investigated the \full keV properties of the XRB population. With novel diagnostic methods that leverage the E $>10$ keV energy band, we were able to distinguish between compact object types and accretion states via hardness-intensity and color-color diagrams. Specifically, we were able to classify \numclass/\numsrc\ sources in the \full keV energy band: \numbh\ BH and \numns\ NS. This is a marked improvement from previous studies where identifying extragalactic XRB compact object types and spectral states has only been possible for a select few of the brightest systems. We studied the relationship between BH and NS and the sSFR of a galaxy. A Spearman's Rank test on the BH fraction, $N_{\rm{BH}}$ / ($N_{\rm{BH}}+N_{\rm{NS}}$), versus sSFR gave a $p$-value of 0.072 and coefficient $r_{s}=0.56$, indicating weak monotonicity but no correlation. Including M31, which is dominated by NS and has low-sSFR, we obtained a $p$-value of 0.028 and coefficient $r_{s}=0.63$. The data suggests that BH dominate star-forming galaxies and NS dominate low-sSFR galaxies. However, due to the varying sensitivity and completeness of our sample, we require improved statistics to investigate this further. Similarly, while there were indications from the data, in agreement with theoretical expectations, that accreting pulsars dominate at high-sSFR and Z/Atoll sources were prevalent at low-sSFR, no statistically significant correlation was found. We found that most BH were identified with the hard accretion state, regardless of host galaxy sSFR, similar to the Galactic BH-HMXB Cygnus X-1. Subsequent analysis of the rich multiwavelength data sets using UV/optical/IR catalogs (in combination with the \nustar-\chandra/\xmmn\ data we analyzed) can help confirm the nature of these sources. We classified galaxies as BH, NS, and ULX-dominated if $>70$\% of their total \full or \hard keV X-ray point source emission came from one of these groups. We found that galaxies with sSFR $\gtrsim2\times10^{-9}$ yr$^{-1}$ were all ULX-dominated, which included all four dwarf galaxies in our sample as well as the starburst galaxies IC 342 and M82. Most galaxies were BH-dominated, whereas in the \full keV energy band, only M31 and NGC 4945 were NS-dominated. We confirmed the \lx-SFR correlation from previous studies by investigating the 8 normal (Milky Way-type) galaxies in the \nustar\ sample with SFR $0.3-12.5$ \sfr. The best-fitting parameters for the relation $\log L_{\text{X}}=\log A+B\log \text{SFR}$ can be found in Table \ref{tab:paramlxsfr}. The \full keV result was nearly identical to previous work in the $2-10$ keV range despite the use of different galaxy samples. We constrained the correlation of X-ray luminosity with SFR and stellar mass using the relationship $L_{\rm{X}}=\alpha M_{\star}+\beta\text{SFR}$. We determined the best-fitting values (see Table \ref{tab:params}) for $\alpha$ and $\beta$ based on the 8 normal (Milky Way-type) galaxies in the \nustar\ sample. In particular, the four dwarf galaxies had increased \lx/SFR compared to normal galaxies, based on past scaling relations. This is not surprising as these dwarf galaxies were selected as ULX hosts, and as such are a biased sample. Studying an unbiased sample of dwarf galaxies would help determine a statistically significant \lx/SFR relation. With the introduction of new programs/observatories to identify faint dwarf galaxies in the optical (e.g.\ Dragonfly, \citealt{abraham01-14}; Dark Energy Survey, \citealt{the-des-collaboration08-16}), multiwavelength studies incorporating X-ray emission would improve our constraints on X-ray source populations in the low-mass regime. We measured galaxy XLFs in the \full and \hard keV energy bands, including the first E $>10$ keV extragalactic XLF for an ensemble of galaxies. We determined that the combined XLF of all sample galaxies in each energy band followed that of the canonical HMXB XLF found by previous studies at E $<10$ keV. Using the classifications for BH and NS, we produced cumulative BH and NS XLFs in the \full and \hard keV energy bands. The \full and \hard keV NS XLFs each indicated a decline beginning at $\approx10^{38}$ and $\approx6\times10^{37}$ \es, respectively, attributable to the 1.4 \msun\ NS Eddington limit. Increased sensitivity and completeness in the \hard keV energy band is required to confirm the decline seen in the \hard keV NS XLF. Using our classifications we investigated the characteristics of BH and NS at different \lx, with a focus on behavior near the Eddington limits. We calculated the overall BH to NS ratio, finding $N_{\rm{BH}}/N_{\rm{NS}}\approx1$ (\full keV) and $N_{\rm{BH}}/N_{\rm{NS}}\approx2$ (\hard keV), over a \full and \hard keV luminosity range for all detected sources of $\sim10^{37}-10^{40.5}$ and $\sim10^{37}-10^{40}$ \es, respectively. We found that the \full keV ratio of BH to NS increased from a value of 1 past the \full keV Eddington limit for a 1.4 \msun\ NS and reached a maximum value of 15 near the \full keV Eddington limit for a 10 \msun\ BH. However, while the total number of accreting stellar-mass BH may decrease beyond the 10 \msun\ BH $L_{\rm{Edd}}$, an improved statistical sample is required to determine its validity. To extend to larger \lx\ we investigated the BH fraction, $N_{\rm{BH}}$ / ($N_{\rm{BH}}$+$N_{\rm{NS}}$), finding approximately monotonic increase beyond the Eddington limits for a 1.4 \msun\ NS in both energy bands. We found evidence for a decrease in the BH fraction below 40\% beyond the \full keV Eddington limit for a 1.4 \msun\ NS (the data point was coincident with the bolometric Eddington limit for a 1.4 \msun\ NS). A larger sample with uniform completeness for \lx$<10^{38}$ \es\ is required to determine if NS cluster near their Eddington luminosities. This galaxy sample was biased towards late-type/spiral galaxies and contained no early-type galaxies, meaning that we did not offer a comprehensive view of older stellar populations. Future observations that focus on building a sample of elliptical galaxies would be of great interest. Such galaxies have inherently faint LMXB populations and thus require extended observing campaigns. However, the nearest candidate, Cen A, is problematic due to its AGN, and most giant elliptical galaxies are located at $d\gtrsim10$ Mpc, prohibiting resolved XRB studies with \nustar. Therefore, to improve our understanding of the XRB population in elliptical galaxies at E $>10$ keV requires a next generation hard X-ray telescope. Using our XRB classifications and XLFs enables comparison with binary population synthesis modeling \citep{fragos09-09, sorensen01-17} that predicts the NS and BH XRB populations in these galaxies. Expanding the range of sSFR coupled with increased sensitivity has the potential to profoundly impact the study of accreting compact objects. | 18 | 8 | 1808.05617 |
1808 | 1808.06614_arXiv.txt | Following the current debate on the fate of SN-condensed dust grains, here we explore by means of three-dimensional hydrodynamical simulations the interaction of dusty supernova remnants (SNRs) with the shocked winds of neighboring massive stars within young massive stellar clusters (SSCs). As a comparison, we have also explored the evolution of supernova remnants in the diffuse ISM with constant density. Since the hydrodynamics of SNRs is intimately related to the properties of their immediate environment, the lifecycle of dust grains in SNRs within SSCs is radically different from that in the diffuse ISM. Moreover, off-centered SNRs evolving in the steep density gradient established due to a star cluster wind experience a blowout phase: shell fragmentation due to protruding Rayleigh-Taylor instabilities and the venting of SN ejecta. Our main finding is that clustered SN explosions will cause a net increase in the amount of dust in the surroundings of young massive stellar clusters. Our analysis considers the multiple dust processing resulting from the passage of the SN reverse shock, including its reflection at the SNR's center, the injection of shocked stellar winds within the respective remnant's volume and the effect of secondary forward shocks produced in sequential SN explosions. In the simulations, we have on-the-fly calculated the rates of thermal sputtering and dust-induced radiative cooling provided an initial distribution of grain sizes and dust content. Fast-moving elongated dusty SN ejecta resemble mushroom clouds violently ascending in a stratified atmosphere after volcanic super-eruptions, where the pyroclasts carried by the clouds are wind-driven and eventually accumulate into the vast surroundings. | Dust nucleation and subsequent growth in the ejecta of core-collapse supernovae (SNe) are nowadays recognized as highly efficient and rapid processes \citep{Cernuschietal1967,TodiniandFerrara2001}. To put it into perspective, it has been estimated that up to $0.4$ to $1.1$ M$_{\odot}$ of dust can be formed per SN, as in the case of several nearby supernova remnants (SNRs) like SN 1987A \citep{Indebetouwetal2014,Matsuuraetal2014}, Cassiopeia A \citep{DeLoozeetal2017,Bevanetal2016} and G54.1+0.3 \citep{Temimetal2017,Rhoetal2017}. Modeling dust injection by SNe is key to understand the process of dust enrichment of galaxies at all cosmic epochs \citep[needed to explain the otherwise puzzling amounts of dust observed in galaxies during cosmic reionization,][]{Laporteetal2017}. However, the fate of the dust grains formed in SN ejecta remains a matter of intense debate as they must be processed by the shocks that are formed from the interaction of freely-expanding SNRs with their ambient medium: a forward shock which sweeps the surrounding ISM, and a reverse shock which moves through the SN ejecta. The former might lead to the destruction of the pre-existent dust in the surroundings, while the latter might be capable of inducing the destruction of a large fraction of the ejecta dust primarily via thermal sputtering. Notwithstanding, while there are thoughtful arguments that put into question the efficiency of dust destruction by SN shocks \citep[e.g.][]{Jonesetal2011,Ferraraetal2016}, recent chemical and dynamical evolution models have favored the notion that only a minor fraction of this dust ($\lesssim 10\%$ by mass) will be able to survive the passage of the respective SN reverse shock \citep[see e.g.][]{Nozawaetal2007,Micelottaetal2016,Bocchioetal2016}. For instance, \citet{Bocchioetal2016} (hereafter B16) followed the dust mass, chemical composition and grain size distribution of several dust species in SN ejecta taking into account the grain dynamics during the free-expansion, Sedov-Taylor and snowplough phases of the evolution of SNRs. The influence of turbulent ISM magnetic fields was studied in detail by \citet{Fryetal2018}, whom, remarking the importance of Rayleigh-Taylor instabilities, showed that charged grains which are kinematically-decoupled from the ejecta gas might be impeded to traverse the contact discontinuity which separates the shocked ejecta from the shocked ISM. To offer further insights, it is crucial to tackle the evolution of dusty SNRs within young massive stellar clusters, \textit{where prolific dust enrichment is indeed localized} \citep{Consiglioetal2016,Leroyetal2018} and \textit{most massive stars are formed} \citep{LadaandLada2003}, with only a small percentage of them being ejected as runaways \citep{Khorramietal2016,Portegies-Zwartetal2010}. Nonetheless, additional complications arise. Firstly, the gas reinserted amidst these clusters is already shocked before the occurrence of first supernova explosion \citep[e.g.][]{Wunschetal2017} and thus the corresponding sputtering rate is high. Secondly, their high SN rate implies that the cluster and its surroundings will be frequently crossed by successive SN forward shocks (FSs), fostering the destruction of some fraction of the dust (from all sources) in the vicinity. Thirdly, the magnetohydrodynamic turbulence which is driven by the feedback of massive stars may trigger grain shattering \citep{Hirashitaetal2010}, producing an excess of small grains that can be rapidly sputtered in a thermalized medium. In this respect, \citet[][hereafter MST16 \& MWP17, respectively]{MartinezGonzalezetal2016,MartinezGonzalezetal2017} pointed out that the presence of near- to mid-infrared (NIR-MIR) excesses in SSCs, such as those observed in several clusters in M 33 and SBS 0335-052 \citep{Relanoetal2016,Reinesetal2008}, are likely to be an indication of dust efficiently produced by SNe, stochastically heated to high temperatures by electronic collisions and photon absorptions, and readily destroyed in frequent ionic collisions. \citet[][]{TenorioTagleetal2015b} \citep[see also][]{Silichetal2017} have proposed a mechanism by which the SN yields (heavy metals) can be discharged out of a dense proto-globular cluster without polluting the gas within the cluster. Stressing that SN explosions are not synchronized events, they followed the evolution of SN blast waves in a strong density gradient. In the case of off-centered remnants, the interaction with the gas left-over from star formation will lead to the SNR's elongated growth and subsequently to a blowout phase: the development of Rayleigh-Taylor (RT) instabilities and shell fragmentation; thus allowing the venting of the SN ejecta to the ISM. The rapid post-blowout hydrodynamical evolution which leads to milder conditions (i.e. rapid decline in density and temperature) in SN ejecta, make the supernova blowout scenario a viable mechanism by which SN-condensed dust grains could be injected into the circumcluster medium before significant thermal sputtering takes place. As we are interested more on SSCs, here we will explore a similar scenario by means of three-dimensional hydrodynamical simulations, which will enable us to follow the development and growth of RT instabilities and the breaking of the SNR's spherical symmetry. This scheme will also permit us to include the further dust processing which occurs when the SN reverse shock reaches the center of an SNR and generates a secondary forward shock \citep{TenorioTagleetal1990,Micelottaetal2016}; while we will not, however, take into account the full dynamics of the grains relative to the gas. Furthermore, we will include the effects of the additional radiative cooling mediated by frequent collisions between gas particles and dust grains in such extreme conditions. The main difference with the hydrodynamical model presented by \citet[][]{TenorioTagleetal2015b} is that the SNRs in the present work will interact with mass-loaded shocked stellar winds in an SSC rather than with the dense gas ($\sim 10^7$ cm$^{-3}$) in a still-embedded proto-globular cluster. The paper is organized as follows: in Section \ref{sec:scheme} we formulate the star cluster model, the implementation of the star cluster wind and the insertion of SNe with the help of three-dimensional hydrodynamical simulations; then we briefly describe the injection and destruction of dust grains in our simulations and the additional cooling induced by their presence. In Section \ref{sec:models}, we introduce several models to test the relevance of the rapid evolution of isolated SNRs and SNRs within SSCs and the various implications regarding the efficiency of dust injection into the unshocked ISM. Finally, in Section \ref{sec:conclusions} we discuss our results and outline our main conclusions. | \label{sec:conclusions} We have presented three-dimensional hydrodynamical simulations of dusty SNRs evolving in the steep density gradient established by shocked stellar winds in young massive clusters. In the model, dust grains are injected and advected with the SN ejecta. Thus, we have derived the fraction of SN dust mass which is injected into the circumcluster medium after shock-processing. We also explored the case of SNRs evolving in the diffuse ISM finding an excellent agreement with the results obtained by \citet{Bocchioetal2016} within the first $10^4$ years. In these cases, the log-normal grain size distribution results in $\sim 10\%$ more surviving dust mass in comparison to the standard \citetalias{MRN1977} distribution. However, the hydrodynamical evolution of SNRs within massive clusters, where most massive stars reside, is radically different to that of isolated SNRs. This evolution strongly depends on the interplay between the rapid adiabatic expansion and cooling after blowout, and the replenishment of mass and energy in the same region occupied by the SNR via shocked stellar winds \citep{Silichetal2017}. Thus, in many cases, the SN reverse shock crosses already thermalized ejecta and the fate of SN dust grains is sealed at early post-explosion times. The fraction of surviving SN dust mass is a strong function of the cluster's mass and the position of the SNe with respect to the cluster's center. These dependences are manifested as follows: \begin{enumerate} \item In the event of well-centered SNe, as much as $30\%$ of the condensed dust mass survives the shock processing they encounter. Destruction is primarily induced by the interaction with shocked stellar winds rather than due to the crossing of the reverse shock. Given that a large fraction of the SN-condensed grains are efficiently heated to hundreds of K and subsequently destroyed, they should be observed transiently at NIR-MIR wavelengths \citep{MartinezGonzalezetal2016,MartinezGonzalezetal2017}. There are two sub-cases for well-centered explosions according to the cluster's mass: \begin{enumerate} \item For young stellar clusters with masses $\sim 10^6$ M$_{\odot}$, gas cooling induced by gas-grain collisions is very efficient as the grains are injected in multiple events. As originally shown by \citet{TenorioTagleetal2013,TenorioTagleetal2015}, this radically alters the thermodynamics of the star cluster wind. \item For young stellar clusters with masses $\lesssim 10^5$ M$_{\odot}$, the average wind number density within the cluster is low and a larger fraction of the dust mass is able to stream out of the cluster (as much as $\sim 30\%$ in the present models). \end{enumerate} \item Off-centered SNe experience a blowout phase where the remnant elongates in the direction of decreasing wind density and protruding RT instabilities develop. As the SN ejecta cools down and lowers its density faster than in the case of an SN occurring in the diffuse ISM, thermal sputtering becomes less efficient and a non-negligible amount, up to $50\%$ in the present models, of the dust mass could be injected into the circumcluster medium. \end{enumerate} Given the assumed stellar density profile, which approximates the (deprojected) King stellar surface density profile \citep{King1962,Ninkovic1998}, a large fraction of the massive stars reside outside the core of the cluster, and thus, also a large fraction of SNe should experience a blowout phase. \textit{We have found that a larger dust mass fraction is capable of surviving shock-processing in the case of SNe occurring in SSCs than in the case of isolated spherical SNRs. Not only that, we predict that clustered SN explosions will cause a net increase in the amount of dust in the surroundings of young massive stellar clusters} after having survived shock-processing due to the cluster wind, the passage of the reverse and secondary forward shocks and the crossing of subsequent SN forward shocks. As the large quantities of dust observed in galaxies at extreme redshifts require both, efficient dust formation and efficient dust survival rates, our model might serve to alleviate the tension between observational and theoretical expectations, the so-called ``dust budget problem'' \citep{Ferraraetal2016}. The grains which stream out of the cluster will face milder conditions: far from the bulk of the starlight and the hot environment within the cluster. These grains should accumulate in the circumcluster medium and manifest persistently at MIR-FIR wavelengths. We plan to quantify this emission and compare it with observations of nearby massive clusters and blue compact dwarf galaxies. We have restricted our analysis to the evolution of multiple SNRs in clusters as massive as $10^6$ M$_{\odot}$. Clusters which are even more massive have a higher SN rate and a larger gravitational potential; under certain conditions (e.g. enhanced radiative cooling), they might be able to retain a non-negligible fraction of SN ejecta, leading to enhanced Fe abundances in secondary stellar generations. Addressing this scenario is also in our plans. We finally note that the fast-moving dusty ejecta in the blowout phase resemble various aspects of volcanic mushroom clouds upwardly expanding into a stratified atmosphere after violent super-eruptions \citep[see the hydrodynamical simulations of such events by][]{Costaetal2018}. In that case, the pyroclasts\footnote{The term pyroclastic derives from the greek roots {\it pyros}, meaning ``fire'' and {\it klastos}, meaning ``broken in pieces''.}, i.e. ashes and cinders, carried by the unstable cloud are wind-driven and eventually accumulate into continent-size regions. | 18 | 8 | 1808.06614 |
1808 | 1808.08924_arXiv.txt | A recent laboratory experiment of ideal magnetohydrodynamic (MHD) instabilities reveals four distinct eruption regimes readily distinguished by the torus instability (TI) and helical kink instability (KI) parameters \citep{Myers2015}. To establish its observational counterpart, we collect 38 solar flares (stronger than GOES class M5 in general) that took place within 45$^{\circ}$ of disk center during 2011$-$2017, 26 of which are associated with a halo or partial halo coronal mass ejection (CME) (i.e., ejective events), while the others are CMEless (i.e., confined events). This is a complete sample of solar events satisfying our selection criteria detailed in the paper. For each event, we calculate decay index $n$ of the potential strapping field above the magnetic flux rope (MFR) in and around the flaring magnetic polarity inversion line (a TI parameter), and the unsigned twist number $T_w$ of the non-linear force-free (NLFF) field lines forming the same MFR (a KI parameter). We then construct a $n-T_w$ diagram to investigate how the eruptiveness depends on these parameters. We find: (1) $T_w$ appears to play little role in discriminating between confined and ejective events; (2) the events with $n\gtrsim0.8$ are all ejective and all confined events have $n\lesssim0.8$. However, $n\gtrsim0.8$ is not a necessary condition for eruption, because some events with $n\lesssim0.8$ also erupted. In addition, we investigate the MFR's geometrical parameters, apex height and distance between footpoints, as a possible factor for the eruptiveness. We briefly discuss the difference of the present result for solar eruptions with that of the laboratory result in terms of the role played by magnetic reconnection. | A solar flare is primarily considered to be a low-atmosphere tracer of magnetic explosions/eruptions. In terms of outcome, there are two types of flares: ejective and confined \citep{Moore2001}. Ejective flares are accompanied by coronal mass ejections (CMEs), while confined flares do not have associated CMEs. A magnetic flux rope (MFR), characterized by a twisted and writhed topological structure, is thought to be a fundamental structure underlying the phenomenon of CMEs. Although the initiation mechanisms are still under debate, it is now common to explain the onset conditions of a MFR eruption in the context of two ideal magnetohydrodynamic (MHD) instabilities, the torus instability (hereafter TI; \citealt{Kliem2006}) and helical kink instability (hereafter KI; \citealt{Torok2004}). TI and KI are mainly controlled by the structural properties of the strapping magnetic field (i.e., the ambient field that runs perpendicular to the MFR) and the guide magnetic field (the ambient field that runs toroidally along the MFR), respectively. Simply put, TI occurs when the strapping field above the MFR declines with height at a sufficiently steep rate, as quantified by decay index $n$ \citep{Torok2005, Kliem2006, Torok2007}. The TI onset criterion of $n\geq n_\mathrm{crit}=1.5$ was first derived analytically by \citet{Bateman1978} and some MHD simulations have found similar values \citep{Kliem2006, Aulanier2010}. A number of other analytical/numerical studies suggest that this critical index $n_\mathrm{crit}$ may lie in a wider range of $0.5<n_\mathrm{crit}<2$ \citep{Fan2007, Demoulin2010, Fan2010, Olmedo2010, Zuccarello2015}. KI, on the other hand, occurs when a MFR is twisted by more than a critical value. The minimum critical twist $\Phi_\mathrm{crit}$ found among analytical/numerical studies is 2.5$\pi$ (corresponding to 1.25 field line windings about the rope axis) \citep[e.g.,][]{Hood1981, Baty2001, Torok2003, Fan2003}. The slow decay of strapping field with height may help confine MFRs, and, in some simulation cases, allows MFRs to build up twist for developing KI \citep{Fan2007}. Observationally, investigations on what magnetic factors determine the likelihood of ejective/confined eruptions have largely focused on one or two aspects: the decay index $n$ of the potential strapping field \citep[e.g.,][]{LiuY2008, Guo2010, Liu2010a, Nindos2012, Baumgartner2018}, and/or the non-potentiality of active regions (ARs) such as free magnetic energy, relative magnetic helicity, magnetic twist, etc. \citep[e.g.,][]{Nindos2004, Falconer2006, Falconer2009, Tziotziou2012, Lee2016, Toriumi2017}. It has been found in some well-studied cases that confined flares are often hosted by ARs with stronger strapping field and weaker non-potentiality in comparison to ejective ones \citep{Sun2015, Jing2015}. For ejective events, the CME speed shows a positive correlation with the decay index of hosting ARs \citep{Xu2012b, Cui2018}. It is worth noting that an unprecedented laboratory experiment designed to study the Sun-like line-tied MFRs reveals four distinct eruption regimes which are readily distinguished by the TI and KI parameters (\citealt{Myers2015}; see their Figure 2). In the four regimes MFRs are either eruptive, stable, failed kink (i.e., torus-stable MFRs that exceed the kink threshold fail to erupt), or failed torus (i.e., kink-stable MFRs that exceed the torus threshold fail to erupt). Such an experimental result on the TI and KI has direct implications for eruptions in the solar corona, and its observational counterpart remains to be established, which is the motivation of this study. In this paper, we present the TI vs. KI parameter diagram, established from a statistical study using solar observations together with the coronal field extrapolation techniques. The goal of this study is to improve our understanding of the requirements for a solar eruption: what the trigger/driver mechanisms might be, and what, if any, onset criteria must be reached. | The previous laboratory experiment reveals that the eruptiveness of MFRs is dependant on the interplay between the TI and KI, as represented by the $n-T_w$ diagram. In this paper we intended to establish a solar counterpart to the diagram, by which we may be able to tell the likelihood of a CME based on the observed $n$ and $T_w$ parameters. The key results are summarized and discussed as follows: First, the TI quantified by $n$ appears to play an important role in differentiating between ejective and confined flares. However, the TI onset criteria ($n\geq n_\mathrm{crit}=\sim0.75$) found here is not a necessary condition for CMEs. Some MFRs in the TI-stable regime still manage to break through the strong strapping field and evolve into CMEs. It therefore implies that an additional trigger and driving mechanism may be involved in solar eruptions. A very likely candidate for the alternative process is magnetic reconnection during solar flares. Actually there are a number of analytical/numerical models invoking magnetic reconnection in the mechanism of CMEs. For instance, in the magnetic breakout model \citep{Antiochos1999}, magnetic reconnection leads to the progressive transformation of the magnetic configuration, allowing a MFR to burst open. In the tether-cutting reconnection model \citep{Moore2001}, magnetic reconnection below a MFR ``cut''s the ``tether''s of the strapping field to unleash CMEs. Such non-ideal MHD processes are absent in the laboratory experiment which was designed to simulate eruptions solely in terms of an ideal MHD process. Second, it is unclear in this study whether the KI represented by $T_w$ is a major factor for solar eruption. Two MFRs with the highest value of $T_w>1.2$ erupted, but many other MFRs with smaller values of $T_w$ were also able to erupt, and we tend to believe that KI is less influential. We consider two possible caveats. The first concerns the ongoing debate whether a helical magnetic structure pre-exists before an eruption \citep{Low1994, Chen1989, Fan2004} or is formed in the course of an eruption via magnetic reconnection \citep{vanBallegooijen1989, Amari2000, MacNeice2004}. There are observational evidence in favor of each scenario \citep[e.g.,][]{Dere1999, Qiu2007, Liu2010b, Song2014, Wang2015, Yan2015, Gopalswamy2017}. In the latter case, it is not surprising that $T_w$ derived from the pre-eruption magnetic field may be underestimated and can not correctly predict the eruptiveness. The second possibility is that helical KI could result in the deformation of a MFR, but may not allow a huge expansion of the MFR to produce a CME \citep{Green2018}. In this sense we may consider that KI might be capable of initiating a filament eruption and a flare, but may not be the key factor in driving a CME into the heliosphere. Third, the laboratory experiment by \citet{Myers2015} shows that there can be both failed TI and failed KI events. Namely, MFRs have more difficulty in eruption than solar community believed. This is contrary to our results that even the TI-stable ($n<0.75$) ones can erupt and CMEs can occur regardless of the KI parameter $T_w$. As mentioned earlier, we speculate that magnetic reconnection, which was absent in the laboratory experiment, may be the factor causing the differences between the laboratory and the present solar observations, if it alleviates the difficulties in eruption. The differences between the laboratory results and our results may also arise from multiple sources of assumptions and approximations of this study in contrast to the lab experiment. In the present study, the TI and KI parameters $n$ and $T_w$ are not directly measured in observations, but rather estimated from MFRs in NLFF field models. The identification of MFRs relies on the quality of NLFF field extrapolation. Although the up-to-date NLFF field extrapolation technique employed here was evaluated thoroughly in comparison with a 3D radiative MHD model and was found to offer a reasonably high accuracy of the coronal field reconstruction \citep{Wiegelmann2010a, Wiegelmann2010, Fleishman2017}, the direct validation of NLFF fields still cannot be performed due to the lack of the coronal magnetic field diagnostics. We'd like to add a caution that NLFF field extrapolation has intrinsic limitations associated with the force-free assumption and is subject to numerous uncertainties in the data reduction and modeling process which are not reflected in our results. It may be that $n$ and/or $T_w$ could not be accurately calculated under the observational limits. In addition, the KI parameter $T_w$ is derived from and averaged over individual field lines, assuming that it's related to the winding of field lines around the axis, but actually the twist of a MFR could be underestimated by its built-in assumption. Finally, we'd like to mention that the present statistical study is a step forward to access the role of the TI and KI in solar eruptions. Detailed studies of the pre-to-post flare magnetic configuration are also needed to better understand the underlying physics, which will be conducted in the future. | 18 | 8 | 1808.08924 |
1808 | 1808.10353_arXiv.txt | {% Baikal-GVD is a next generation, kilometer-scale neutrino telescope under construction in Lake Baikal. It is designed to detect astrophysical neutrino fluxes at energies from a few TeV up to 100 PeV. GVD is formed by multi-megaton subarrays (clusters). The array construction was started in 2015 by deployment of a reduced-size demonstration cluster named "Dubna" . The first cluster in it's baseline configuration was deployed in 2016, the second in 2017 and the third in 2018. The full-scale GVD will be an array of $\sim$10.000 light sensors with an instrumented volume about of 2 cubic km. The first phase (GVD-1) is planned to be completed by 2020-2021. It will comprise 8 clusters with 2304 light sensors in total. We describe the design of Baikal-GVD and present selected results obtained in 2015 – 2017. } | \label{intro} The deep underwater neutrino telescope Baikal Gigaton Volume Detector (Baikal-GVD) is currently under construction in Lake Baikal [1]. Baikal-GVD is formed by a three-dimensional lattice of optical modules (OMs) arranged at vertical load-carrying cables to form strings. The telescope has a modular structure and consists of functionally independent clusters - sub-arrays comprising a total of 288 OMs each and connected to shore by individual electro-optical cables. The first, reduced size cluster named “Dubna” has been deployed in Lake Baikal and was operated during 2015. In April 2016, this array has been upgraded to the baseline configuration of a GVD-cluster, which comprises 288 optical modules attached at 8 strings at depths from 750 m to 1275 m. In 2017 and 2018 the second and the third GVD-clusters were deployed, increasing the total number of operating optical modules to 864 OMs. During Phase-1 of Baikal-GVD implementation an array consisting of eight clusters will be deployed by 2020-2021. Since each GVD-cluster represents a multi-megaton scale Cherenkov detector, studies of neutrinos of different origin are allowed with early stages of construction. | The ultimate goal of the Baikal-GVD project is the construction of a km3-scale neutrino telescope with implementation of about ten thousand light sensors. The array construction was started by deployment of reduced-size demonstration cluster named "Dubna" in 2015, which comprises 192 optical modules. The first cluster in it's baseline configuration was deployed in 2016 and the second one in 2017. After deployment of the third GVD-cluster in April 2018 Baikal-GVD comprises the total of 864 OMs arranged at 24 strings and becomes, at present, the largest underwater neutrino telescope. The modular structure of Baikal-GVD design allows studies of neutrinos of different origin with early stages of construction. Analysis of data collected in 2015-2017 allows for extraction of a sample of upward through-going muons as clear neutrino candidates and the identification of the first two promising high-energy cascade events - candidates for events from astrophysical neutrinos. The search for neutrinos assosiated with GW170817 with Baikal-GVD allows to derive upper limits on the neutrino spectral fluence from this source. The commissioning of the first stage of the Baikal neutrino telescope GVD-1 with an effective volume 0.4 km$^3$ is envisaged for 2020-2021. \begin{acknowledgement} {\it This work was supported by the Russian Foundation for Basic Research (Grants 16-29-13032, 17-02-01237).} \end{acknowledgement} | 18 | 8 | 1808.10353 |
1808 | 1808.07013_arXiv.txt | It has been suggested by \cite{Weinberg:2013pbi} that an instability due to the nonlinear coupling of a neutron star's tide to its $p$- and $g$-modes could affect the gravitational-wave phase evolution of a neutron-star binary. \cite{Weinberg:2015pxa} suggests that this instability can turn on as the gravitational waves pass through the sensitive band of ground-based detectors, although the size of the effect is not known. The discovery of the binary neutron star merger GW170817 provides an opportunity to look for evidence of nonlinear tides from $p$-$g$ mode coupling. We compute Bayesian evidences that compare waveform models that include the $p$-$g$ mode coupling to models that do not. We find that the consistency between GW170817 signal and the $p$-$g$ mode model reported by \cite{abbott2019constraining} is due to a degeneracy between the phenomenological waveform used to model the effect of nonlinear tides and the standard post-Newtonian waveform. We investigate the consistency of the GW170817 signal with regions of the parameter space where the effect of nonlinear tides is not degenerate with the standard model. Regions of the nonlinear tide parameter space that have a fitting factor of less than 99\% (98.5\%) are disfavored by a Bayes factor of 15 (25). We conclude that regions of the parameter space where nonlinear tides produce a measurable effect are strongly disfavored and improved theoretical modeling will be needed if future observations are to constrain nonlinear tides from $p$-$g$ mode coupling in neutron stars. | \label{sec:intro} The discovery of the binary neutron star merger GW170817 \citep{TheLIGOScientific:2017qsa} has given us a new way to explore the physics of neutron stars. Recent studies have measured the star's tidal deformability and placed constraints on the equation of state of the neutron stars~\citep{TheLIGOScientific:2017qsa,Tews:2018iwm,Most:2018eaw,Raithel:2018ncd,de2018tidal,Abbott:2018exr,Abbott:2018wiz,Radice:2018ozg,LIGOScientific:2019eut,Capano:2019eae}. \cite{Weinberg:2013pbi} have suggested that the star's tidal deformation can induce nonresonant and nonlinear daughter wave excitations in $p$- and $g$-modes of the neutron stars via a quasi-static instability. This instability would remove energy from a binary system and possibly affect the phase evolution of the gravitational waves radiated during the inspiral. Although \cite{Venumadhav:2013nla} concluded that there is no quasi-static instability and hence no effect on the inspiral, \cite{Weinberg:2015pxa} claims that the instability can rapidly drive modes to significant energies well before the binary merges. However, the details of the instability saturation are unknown and so the size of the effect of the $p$-$g$ mode coupling on the gravitational waveform is not known~\citep{Weinberg:2015pxa}. The discovery of the binary neutron star merger GW170817 by Advanced LIGO and Virgo provides an opportunity to determine if there is evidence for nonlinear tides from $p$-$g$ mode coupling during the binary inspiral. Since the physics of the $p$-$g$ mode instability is uncertain, \cite{Essick:2016tkn} developed a parameterized model of the energy loss due to nonlinear tides. This model is parameterized by the amplitude and frequency dependence of the energy loss, and the gravitational-wave frequency at which the instability saturates and the energy loss turns on. For plausible assumptions about the saturation, \cite{Essick:2016tkn} concluded that $> 70\%$ of binary merger signals could be missed if only point-particle waveforms are used, and that neglecting nonlinear tidal dynamics may significantly bias the measured parameters of the binary. Bayesian inference can be used to place constraints on nonlinear tides during the inspiral of GW170817. An analysis by \cite{abbott2019constraining} computed Bayes factors that investigate whether the GW170817 signal is more likely to have been generated by a model which includes nonlinear tides or one which does not. \cite{abbott2019constraining} find a Bayes factor of order unity, and conclude that the GW1701817 signal is consistent with both a model that neglects nonlinear tides and with a model that includes energy loss from a broad range of $p$-$g$ mode parameters. However, the prior space used in this analysis includes a large region of parameter space where the amplitude of the effect produces a gravitational-wave phase shift that is extremely small. In this case, a waveform that includes $p$-$g$ mode parameters will have a likelihood that is identical to the likelihood of the waveform without the $p$-$g$ mode instability. The $p$-$g$ mode model extends the standard waveform model by adding additional parameters that describe the nonlinear tidal effects. However, when including new parameters in a hypothesis if the likelihood does not vary across large portions of the prior volume for these new parameters relative to the likelihood of the original model, then the Bayes factor will not penalize this additional prior volume, nor will it penalize any extraneous parameters in the model (see e.g. \cite{kass1995bayes,hobson2010bayesian}). We examine the prior space of the $p$-$g$ mode model used by \cite{abbott2019constraining} and find that although the $p$-$g$ model model contains regions that are not consistent with the standard model, there are large regions of the prior space where the likelihood is high because the $p$-$g$ mode model is degenerate with the standard model. These regions of the prior space dominate the evidence and hence the Bayes factor neither favors nor disfavors the inclusion of $p$-$g$ mode parameters. We investigate a variety of different prior distributions on the $p$-$g$ mode parameters beginning with a prior distribution that is similar to that tested in~\cite{abbott2019constraining} and includes large regions of the parameter space that produce a negligible gravitational-wave phase shift. When comparing the evidence for this model with the standard waveform model used by \cite{de2018tidal} we find a Bayes factor of order unity, as expected. We then investigate a prior distribution in which the $p$-$g$ mode instability parameters are constrained to induce a phase shift to the waveform that is greater than $0.1$ radians. This phase shift is calculated from the time the waveform enters the sensitive band of the detector to the time when the waveform reaches the innermost stable circular orbit. We choose this threshold to exclude trivial regions of the parameter space that produce a non-measurable effect. However, we again find a Bayes factor of order unity when compared to the model hypothesis that does not model the $p$-$g$ mode instability. Investigation of these results showed that this is due to parameter degeneracies between the $p$-$g$ mode model and the intrinsic parameters of the standard waveform model. Finally, we reduce the prior space to contain only the regions where the $p$-$g$ mode waveform is not degenerate with the standard model by computing the fitting factor~\citep{Apostolatos:1995pj} of $p$-$g$ signals against a set of standard waveforms. We do this to restrict the region of parameter space to that where the $p$-$g$ effect is \emph{measurably} distinct from a model that neglects nonlinear tides. We calculate the Bayes factor as a function of the fitting factor. We find that as the $p$-$g$ mode parameter space is restricted to exclude regions that have a high fitting factor with standard waveforms, the Bayes factor decreases significantly. Regions of the nonlinear tide parameter space that have a fitting factor of less than 99\% (98.5\%) are strongly disfavored by a Bayes factor of 15 (25). While certain prior distributions of $p$-$g$ mode parameters are consistent with the data, we find that these distributions are ones that contain large regions of non-measurable parameter space either because the effect produced is too small to measure, or the effect is degenerate with other parameters of the standard model. We conclude that the consistency of the GW170817 signal with the model of \cite{Essick:2016tkn} is due to degeneracies and that regions where nonlinear tides produce a measurable effect are strongly disfavored. | In this paper, we have used the observation of GW170817 and the model of~\cite{Essick:2016tkn} to look for evidence of nonlinear tides from $p$-$g$ mode coupling during the inspiral~\citep{Weinberg:2013pbi,Weinberg:2015pxa,Zhou:2018tvc}. Over the broad prior space, we find a Bayes factor of unity which gives an inconclusive result on whether nonlinear tides are favored or disfavored in GW170817, consistent with \cite{abbott2019constraining}. This Bayes factor can be interpreted as stating that there is insufficient evidence to change our prior beliefs about the credibility of the $p$-$g$ mode hypothesis after the observation of GW170817. A closer examination of the posterior distribution lead us to conclude that nonlinear tides are consistent with the signal GW170817 because they either cause very small phase shifts to the waveform, or the nonlinear tides must enter the waveform in a way that is degenerate with the other intrinsic parameters of GW170817. Regions of the nonlinear tide parameter space that have a fitting factor of less than 99\% (98.5\%) are disfavored by a Bayes factor of 15 (25). We find that waveforms from a $p$-$g$ mode instability with overlap $>98.5$ \%, tend to either induce a very small phase shifts to the waveform or are degenerate with other intrinsic parameters of GW170817. This leads us to conclude that modeling GW170817 with nonlinear tidal parameters may not offer advantages over using a simpler model. We conclude that the consistency of the GW170817 signal with the model of \cite{Essick:2016tkn} is due to parameter degeneracy and that regions where nonlinear tides produce a measurable effect are strongly disfavored. In principle, one could improve our analysis by separately parameterizing the amplitude, turn-on frequency, and frequency evolution for each star as in~\cite{abbott2019constraining}. However, we find our results to be broadly consistent with~\cite{abbott2019constraining}, and so we do not expect these to affect the main conclusion of our paper. Further improvements to the parametric model of $p$-$g$ mode instability could include a higher order post-Newtonian expansion of the instability, or further understanding of the instability's interaction with neutron star magnetic fields~\citep{Weinberg:2015pxa}. Nonlinear tides are poorly understood and the contribution from other stellar oscillation modes may yet contribute to a more accurate picture of the interior dynamics of neutron stars~\citep{Andersson:2017iav}. Current models of the gravitational-wave phase shift caused by nonlinear tides from the $p$-$g$ mode instability suffer from parameter degeneracies with the other intrinsic parameters of a neutron star binary. A measurement of the binary's chirp mass that is independent of gravitational-wave observations would break this degeneracy. However, for a system like GW170817, this would require measurement of the binary's chirp mass to a precision greater than $\sim 0.02 \%$ using an electromagnetic counterpart, which is implausible. Absent improved theoretical understanding of nonlinear tides from $p$-$g$ mode coupling, it is unlikely that future observational constraints will be able to significantly improve our knowledge of these physical processes. | 18 | 8 | 1808.07013 |
1808 | 1808.01129_arXiv.txt | \citehuang have analysed the population of 15 known galactic \dns[s] regarding the total masses of these systems. They suggest the existence of two sub-populations, and report likelihood-based preference for a two-component Gaussian mixture model over a single Gaussian distribution. This note offers a cautionary perspective on model selection for this data set: Especially for such a small sample size, a pure likelihood ratio test can encourage overfitting. This can be avoided by penalising models with a higher number of free parameters. Re-examining the \dns total mass data set within the class of Gaussian mixture models, this can be achieved through several simple and well-established statistical tests, including information criteria (AICc, BIC), cross-validation, Bayesian evidence ratios and a penalised EM-test. While this re-analysis confirms the basic finding that a two-component mixture is consistent with the data, the model selection criteria consistently indicate that there is no robust preference for it over a single-component fit. Additional \dns discoveries will be needed to settle the question of sub-populations. | The population of galactic \dns[s] -- or \bns[s], as the \gw community prefers to call them -- is of high interest as a locally accessible predictor for the population of merging binaries in the wider Universe, which has recently become accessible to \gw observations with LIGO and Virgo~\citep{TheLIGOScientific:2017qsa}. Traditionally, a lot of work has focused on using the observed galactic sample to predict coalescence rates \citep[see][ and references therein]{Aasi:2013wya}, though the distribution of component masses has also been studied~\citep{Schwab:2010jm,Zhang:2010qr,Ozel:2012ax,Kiziltan:2013oja}. In a recent paper, \huang (in the following: \huangshort) have considered the total gravitational masses $\Mtot$ of \numdns known \dns[s]. $\Mtot$ is of special interest in predicting the fate of binary merger remnants and for studies of the nuclear \eos~\citep{Baiotti:2016qnr,Margalit:2017dij,Ma:2017yva,Abbott:2017dke,Abbott:2018hgk,Abbott:2018exr}. \huang point out an apparent bimodality in the distribution of $\Mtot$, and with the help of \gmm[s] and a likelihood ratio test, they arrived at a $2\sigma$ preference for two components over one. In this note, I suggest additional statistical tests not originally considered by \huang, and caution against relying on likelihood-ratio tests alone, especially when applied to small data sets. Hence, let us re-evaluate the suggested preference for a two-component \gmm fit to the observed \dns $\Mtot$ distribution with a series of simple tests. First, for completeness, (i) visual inspection of the data set (Sec.~\ref{sec:data}) and (ii) \gmm fitting and likelihood-ratio tests (Sec.~\ref{sec:gmms}) are briefly summarised. The additional hypothesis test methods include (iii) information criteria (AICc and BIC) that penalise underconstrained parameters (Sec.~\ref{sec:aicbic}), (iv) a cross-validation test to understand the impact of individual \dns systems on model selection (Sec.~\ref{sec:crossval}), (v) Bayesian evidence computation through nested sampling (Sec.~\ref{sec:nestsamp}), and (vi) a penalised EM-test (Sec.~\ref{sec:em-test}). To provide more context for the model selection results, the same criteria are also applied on additional examples: simulated larger $\Mtot$ data sets (appendix~\ref{sec:appendix-sims}) and a physically different, but statistically not dissimilar data set of \ns spins from~\citet{Patruno:2017oum} (appendix~\ref{sec:appendix-lmxbs}). \vspace{-0.5\baselineskip} | The distribution of total masses $\Mtot$ of Galactic \dns systems shows some apparent bimodality, which can be fit with a two-component \gmm as shown by \huang. A pure likelihood ratio test prefers those two components over one, with \huang estimating the significance of this preference as $2\sigma$. As a first step towards testing if this indeed points to two distinct underlying populations of astrophysical objects, while it is my understanding that \huang are also working on a more sophisticated analysis, in this note I have kept their initial \gmm assumption, but considered more robust model selection criteria: Neither the frequentist information criteria (AICc and BIC) considered in Sec.~\ref{sec:aicbic}, which amend the likelihood ratio test with a penalty for the higher number of free parameters in multi-component \gmm[s], nor a Bayesian evidence ratio test (Sec.~\ref{sec:nestsamp}) find any robust preference for more than one component. The various \gmm fitting methods employed here still all agree with \huang that a two-component \gmm certainly provides `a good fit' to the data; the scenario is not ruled out either and, as pointed out by \huang, could have interesting consequences for stellar evolution models and \gw astronomy. But it appears that the present set of known \dns[s] is simply too small, and some systems' masses are not constrained well enough, to robustly decide between one or two components. It will be interesting to revisit this model selection problem once additional \dns systems are observed, as expected in great numbers from upcoming surveys e.g. with MeerKAT~\citep{Bailes:2018azh} and the SKA~\citep{Smits:2008cf}; or to combine the Galactic sample with \gw observations of extragalactic mergers, as suggested by \huang in the second half of their paper. (Though~\citet{Pankow:2018iab} suggests that GW170817~\citep{TheLIGOScientific:2017qsa} might not be consistent with the same population as the galactic \dns[s].) In the meantime, the simple reanalysis in this note has certainly not exhausted the full potential of the present data set. One could also consider distribution functions beyond the \gmm family. (\huang already suggested a \gmm plus uniform distribution.) And since the cross-validation analysis suggests that, for the current small data set size, a few systems can have a large effect on any inference of the underlying distribution, revisiting individual systems' mass measurements -- or even their identity as \dns[s] -- could also improve the situation. For example, \mbox{J1811--1736} has the widest uncertainty in the \huang data set (\mbox{$\Mtot=2.57\pm0.10$}); referring back to the original studies of \citet{Lyne:1999hu} and \citet{Corongiu:2006rd}, its rather low companion mass means that while it is generally accepted as a \dns, this identification might not be completely iron-clad. A combined reanalysis of total and component masses could also be promising in constraining the model selection problem, and more sophisticated statistical techniques could be applied to deal with possible selection effects. Data sets used in this note (reproduced from \citehuang and \citet{Patruno:2017oum}) and \cpnest posterior samples are provided as \href{https://arxiv.org/src/1808.01129/anc}{ancillary files} of the arXiv preprint. \vspace{-\baselineskip} | 18 | 8 | 1808.01129 |
1808 | 1808.04821_arXiv.txt | For a sample of $\sim$80 local ($0.02 \leq z \leq 0.1$) Seyfert-1 galaxies with high-quality long-slit Keck spectra and spatially-resolved stellar-velocity dispersion (\s) measurements, we study the profile of the [OIII]$\lambda$5007\AA~emission line to test the validity of using its width as a surrogate for \s. Such an approach has often been used in the literature, since it is difficult to measure \s~for type-1 active galactic nuclei (AGNs) due to the AGN continuum outshining the stellar-absorption lines. Fitting the [OIII] line with a single Gaussian or Gauss-Hermite polynomials overestimates \s~by 50-100\%. When line asymmetries from non-gravitational gas motion are excluded in a double Gaussian fit, the average ratio between the core [OIII] width ($\sigma_{\rm {[OIII],D}}$) and~\s~is $\sim$1, but with individual data points off by up to a factor of two. The resulting black-hole-mass-$\sigma_{\rm {[OIII],D}}$ relation scatters around that of quiescent galaxies and reverberation-mapped AGNs. However, a direct comparison between \s~and $\sigma_{\rm {[OIII],D}}$ shows no close correlation, only that both quantities have the same range, average and standard deviation, probably because they feel the same gravitational potential. The large scatter is likely due to the fact that line profiles are a luminosity-weighted average, dependent on the light distribution and underlying kinematic field. Within the range probed by our sample (80-260\,km\,s$^{-1}$), our results strongly caution against the use of [OIII] width as a surrogate for \s~on an individual basis. Even though our sample consists of radio-quiet AGNs, FIRST radio-detected objects have, on average, a $\sim$10\% larger [OIII] core width. | \label{intro} The relationship between the masses of supermassive black holes (BHs) and the properties of their host galaxies has been amongst the most active research areas in contemporary astrophysics, hinting at a co-evolution between BHs and galaxies \citep[for a recent review see, e.g.,][]{kor13}. Such a co-evolution can be explained either by mutual growth via mergers or by feedback from the active galactic nucleus (AGN) in an evolutionary stage when the BH is growing through accretion. AGNs are thus promising probes towards understanding the origin of these BH mass (\mbh) scaling relations. Unfortunately, the AGN emission (featureless non-stellar continuum plus emission lines) often outshines the host galaxy, making it difficult to measure the host-galaxy properties. In particular, measuring stellar-velocity dispersion ($\sigma_{\star}$), which, of all host galaxy properties, seems to scale the tightest with the BH mass \citep{bei12, sha16}, is hampered by the contaminating AGN continuum and emission lines. To mitigate this problem, several studies have suggested to use the width of the [OIII]$\lambda$5007\AA~emission line (hereafter [OIII]) originating in the narrow-line region (NLR) as a surrogate for $\sigma_{\star}$, assuming that the NLR is gravitationally bound to the bulge and thus, that the gas kinematics follows the bulge potential \citep[e.g.,][]{ter90,whi92,nel96,nel00,shi03,bor03,gre05,net07,sal07,sal13}. However, while the [OIII] emission line is a prominent line that can be easily measured in AGNs out to large distances, it is also known to often have asymmetric line profiles due to non-gravitational gas kinematics such as outflows, infalls, or interaction with radio jets. In particular, it is known to often display a blue wing \citep[e.g.,][]{hec81,der84,whi85,wil85,mul13,woo16}, generally interpreted as a signature of outflows with dust preferentially hiding one cone behind the stellar disk. For that reason, some studies have excluded the [OIII] blue wing, as well as any radio sources and galaxies undergoing tidal interactions. The \mbh~was found to scale with the width of the [OIII] line ($\sigma_{\rm {[OIII]}}$), albeit with a large scatter \citep[e.g.,][]{nel96,gre05}. Other studies have suggested the use of different emission lines, such as [SII]$\lambda\lambda$6716, 6731 \citep[e.g.,][]{kom07,ho09} that have a lower ionization potential and do not suffer from substantial asymmetries, or mid-infrared lines \citep[e.g.,][]{das08,das11}, but the scatter is comparable to that of the core of the [OIII] line. While all studies confirm the original findings by \citet{nel96}, i.e.~a moderately strong correlation between $\sigma_{\star}$ and $\sigma_{\rm {[OIII]}}$ but with real scatter, the origin of the scatter remains unclear. No dependencies have been found with AGN luminosity, host galaxy morphology, star formation rate, or local environment \citep{gre05,ric06}. However, unlike the original study by \citet{nel96}, very few previous studies have measured both properties, $\sigma_{\star}$ and $\sigma_{\rm {[OIII]}}$, directly and simultaneously for a given sample, mainly due to the difficulties of measuring stellar-velocity dispersion in type-1 active galaxies. Often, conclusions are instead drawn by comparing the \mbh-$\sigma_{\rm {[OIII]}}$ relation for type-1 galaxies to the \mbh-$\sigma_{\star}$ relation for quiescent galaxies \citep{nel00,kom07}, or by comparing \mbh~derived from $\sigma_{\rm {[OIII]}}$ to \mbh~derived from reverberation mapping \citep{nel00} or the virial method using H$\beta$ \citep{bor03}. \citet{bon05} predict $\sigma_{\star}$ indirectly from the Faber-Jackson relation and conclude, from studying the $M_{\rm host}$-$\sigma_{\rm {[OIII]}}$ relationship for a sample of 21 radio-quiet quasars, that $\sigma_{\rm {[OIII]}}$ is on average consistent with $\sigma_{\star}$. Similarly, \citet{sal15} find agreement with the Faber-Jackson relation when using the width of the [OIII] emission line as a proxy for stellar-velocity dispersion, supporting the general utility of the [OIII] line width as a surrogate for $\sigma_{\star}$ in statistical studies. \citet{gru04} and \citet{wan01} use $\sigma_{\rm {[OIII]}}$ to investigate their \mbh~distributions of narrow-line Seyfert-1 galaxies (NLSy1). \citet{gre05} compare $\sigma_{\rm {[OIII]}}$ to $\sigma_{\star}$ directly, but for a sample of type-2 Seyfert galaxies. Similarly, \citet{woo16} use a sample of 39,000 type-2 AGNs at $z<0.3$ from SDSS and find a broad relation between [OIII] and $\sigma_{\star}$, but with [OIII] being wider by 30-40\% since wings are not excluded from the fit. However, for a sub-sample of AGNs for which the [OIII] profile is well fitted by a single Gaussian model, \citet{woo16} find that the velocity dispersion is comparable to the stellar-velocity dispersion. \citet{ric06} use spatially-resolved HST/STIS spectra for a sample of mostly type-2 Seyfert galaxies and find that NLR line widths underestimate $\sigma_{\star}$. Other studies have assumed that $\sigma_{\rm {[OIII]}}$ traces $\sigma_{\star}$ and used it to probe cosmic evolution \citep{shi03, sal13}. Also, most studies cited above use the width of the entire [OIII] emission line, possibly including non-gravitational motion, even though already \citet{nel96} showed that the [OIII] line profile base and wings do not correlate as tightly with stellar-velocity dispersion as the [OIII] core \citep[similar conclusions were also reached by][]{gre05}. Thus, despite the widespread use of $\sigma_{\rm {[OIII]}}$ as a substitute for $\sigma_{\star}$, caution is in order. We have recently presented a baseline of the \mbh-$\sigma_{\star}$ relation for active galaxies for a sample of 65 Seyfert-1 galaxies in the local Universe selected from the Sloan Digital Sky Survey (SDSS) \citep[][]{ben15}. SDSS images are used to determine host-galaxy morphology and AGN luminosity free of host-galaxy contamination. High signal-to-noise ratio Keck spectra yield H$\beta$ line width to estimate \mbh~and spatially-resolved stellar-velocity dispersion \citep[][]{ben11a,har12}. Thus, our sample is uniquely suited to study the direct relationship between $\sigma_{\star}$ and $\sigma_{\rm {[OIII]}}$ for a homogeneous sample of local Seyfert-1 galaxies. Moreover, we make use of the spatially-resolved Keck spectra to isolate the nuclear line profile and to probe spatial dependencies. We compare the resulting \mbh-$\sigma_{\rm {[OIII]}}$ relation to the \mbh-$\sigma_{\star}$ relation \citep{ben15} and look for trends with host galaxy and nuclear properties. The paper is organized in the following manner. Section~\ref{sample} summarizes the sample selection, observations, and data reduction. Section~\ref{analysis} describes the analysis of the data. Section~\ref{results} discusses the derived quantities and results. Section~\ref{summary} concludes with a summary. Note that the paper presents, first, a traditional approach focused on velocity dispersion ratios and their correlations to \mbh, and then discusses the correlation between kinematic estimators directly and the shortcomings of conclusions based solely on ratios. Throughout the paper, a Hubble constant of H$_{\rm{0}}$ = 70 km s$^{-1}$, $\Omega_{\rm{\lambda}}$ = 0.7, and $\Omega_{\rm{M}}$ = 0.3 are assumed. | \label{results} We here compare the resulting widths for [OIII] and [OII] with $\sigma_{\star}$. All 81 objects have at least one $\sigma_{\rm {[OIII]}}$ measurement. Quantities necessary for comparison of $\sigma_{\rm {[OIII]}}$ and $\sigma_{\star}$ for aperture spectra within the effective bulge radius are available for 62 of the 81 objects and, thus, the \mbh-$\sigma_{\rm {[OIII]}}$ relation is compared directly to the M$_{\rm{BH}}$-$\sigma_{\star}$ relation for that sub-sample of 62 objects \citep{ben15}. Likewise, when including $\sigma_{\rm [OII]}$ within the effective radius in the comparison, a total of 62 objects are compared. \subsection{[OIII] profile} The double Gaussian fit reveals information on the general [OIII] line profile. For 66\% of objects/spectral rows, the double Gaussian fitting resulted in the fitting of a blue wing (-500\,km\,s$^{-1}$ $\le$ $v$ $\le$ -25\,km\,s$^{-1}$). For 22\% of objects/spectral rows, a Gaussian redshifted compared to the central core was fitted, implying a red wing (25\,km\,s$^{-1}$ $\le$ $v$ $\le$ 500\,km\,s$^{-1}$). For 12\% of objects/rows, the second Gaussian fitted a broader central component (-25\,km\,s$^{-1}$ $\le$ $v$ $\le$ 25\,km\,s$^{-1}$). The histogram of the velocity offset of the second Gaussian (the wing component) compared to the central core Gaussian is shown in Figure~\ref{figure:oiiihisto}, including all objects and spectral rows. The average velocity offset for the blue wing is -155$\pm$7\,km\,s$^{-1}$, and for the red wing 124$\pm$13\,s$^{-1}$, respectively. While these results are overall comparable with those of \citet{woo16} for a sample of $\sim$39,000 type-2 AGNs in SDSS, we find an even higher fraction of kinematic signatures for outflows, likely because of the type-1 nature of our objects for which the viewing angle is favorable to see outflows. Indeed, the average [OIII] profile for type-1 AGNs, as determined from a sample of $\sim$10,000 AGNs from SDSS, shows a strong blue wing that can be well fitted by a broad second Gaussian component (average velocity offset of -148\,km\,s$^{-1}$) \citep{mul13}. Figure~\ref{figure:oiiibroadnarrow} shows examples of the broadest and the narrowest [OIII] emission line profile. The [OIII] line shows rotation in at least 17\% of objects with rotational velocities up to $\sim$$\pm$250\,km\,s$^{-1}$, matching those of the stellar rotation curve \citep{har12}. In 15\% of objects do we see evidence for HII regions in the outer spectra, as traced by a sudden peak in [OIII] along with an increase in the H$\beta$/[OIII] ratio. However, since H$\alpha$ is not covered by our spectra, we cannot verify the origin of the ionization of these regions and thus do not further discuss them here. There is a small fraction of objects ($\sim$7\%) that shows evidence for a change in the [OIII] profile as a function of distance from the center, with the majority showing a red wing on one side of the galaxy center and a blue wing on the other, and some galaxies with the blue wing only present on one side of the galaxy center (Figure~\ref{figure:oiiichange}). Other than that, we do not find any trends with distance from the center. For example, the ratio of broad (wing) [OIII] to narrow (core) [OIII] does not change significantly as a function of radius (when fitted by a double Gaussian); nor does the width of the broad [OIII] component change with radius. Part of this is likely due to the fact that (i) the spectra are restricted to the central few kpc, given the S/N ratio, and (ii) that the central 1-2 kpc are unresolved due to the ground-based seeing. (The 1" width of the long-slit was chosen to match the seeing. 1" corresponds to 0.43kpc for the smallest redshift of z=0.021 of our sample, to 1.8kpc for the largest redshift of z=0.097, and to 1.1kpc for the average redshift of z=0.058.) \begin{figure} \includegraphics[width=\linewidth]{oiii_histo.eps} \caption{ Histogram of the velocity offset of the second Gaussian (the wing component) compared to the central core Gaussian for all objects and spectral rows.} \label{figure:oiiihisto} \end{figure} \begin{figure} \includegraphics[width=\columnwidth]{oiii_BN.eps} \caption{ Examples of central [OIII] profiles for broadest [OIII] line (2222-0819, $\sigma_{\rm [OIII], GH}$ = 514 km\,s$^{-1}$) and narrowest [OIII] line (1605+3305, $\sigma_{\rm [OIII], GH}$ = 127 km\,s$^{-1}$). For comparison, the local continuum was subtracted and the peak flux scaled to 1. \label{figure:oiiibroadnarrow}} \end{figure} \begin{figure} \includegraphics[width=\linewidth]{oiii_change.eps} \caption{ Examples of two objects that show a spatially changing [OIII] emission line profile (center = black; ``negative'' offset from center = solid lines; ``positive'' offset from center = dash-dotted lines; red = 0.68'' distance; blue = 1.62'' distance; green = 2.84'' distance). Left: 1535+5754, observed at a position angle (p.a.) of 100deg, with broader lines further out from the center (``negative''=south-east). This galaxy does not show a strong rotation curve \citep{ben11a}. Right: 1554+3238, observed at a p.a. of 80deg. In addition to the rotation curve ($\pm$200km/s) also visible from the stellar-absorption lines \citep{ben11a}, the object shows a blue wing on the ``positive'' side of the center (south-east) and a red wing on the ``negative'' side (north-west).} \label{figure:oiiichange} \end{figure} \subsection{Comparison between [OIII] line width and $\sigma_{\star}$} We compare the [OIII] line width ($\sigma_{\rm [OIII]}$) derived from the three different fitting methods (single Gaussian, double Gaussian using the central core component only, and Gauss-Hermite polynomials) with the stellar-velocity dispersion ($\sigma_{\star}$). In Figure~\ref{figure:ratiodistance}, the resulting $\sigma_{\rm [OIII]}$/$\sigma_{\star}$ ratio is shown as a function of distance from the center for both the spatially resolved spectra (left panels) as well as the aperture spectra within the bulge effective radius (right panels). In summary, the results show that both the single Gaussian fit as well as the Gauss-Hermite polynomial fit result in an overestimation of $\sigma_{\star}$ by on average 50-100\% (see Table~\ref{table:sigmacompare}). In other words, the entire [OIII] line is broader by $\sim$75\% compared to $\sigma_{\star}$. However, when line asymmetries are fitted by a second Gaussian and excluded, then the central core [OIII] emission-line width is a good tracer of $\sigma_{\star}$ (mean ratio 1.06$\pm$0.02 for spatially-resolved spectra; mean ratio 1.02$\pm$0.04 for spectra within aperture of effective radius \footnote{Note that we list the standard deviation of the mean.}), but with individual data points off by up to a factor of two. Another approach to exclude line asymmetries would be to consider only the width (i.e., $\sigma$) of the first pure Gaussian term in the Gauss-Hermite polynomial fit. Note that the first term (an original symmetric Gaussian) can represent most of the core of the line profile, while the rest of the series (Gaussian multiplied by Hermite polynomials) represents deviations to better describe the observed data profile. The resulting mean ratio with $\sigma_{\star}$ is then reduced to 1.25$\pm$0.04. While this is significantly lower than using the width (i.e., line dispersion) of the full profile of the fit, it still overestimates $\sigma_{\star}$ by $\sim$25\%. This is likely due to the fact that the higher order series terms can have negative values which might then be compensated for by the Gaussian, resulting in an overestimation of the width by the Gaussian component \citep[see also,][]{woo18}. Within the uncertainties, our data do not show a strong dependency of the $\sigma_{\rm [OIII]}$/$\sigma_{\star}$ on distance from the galactic center for any of the three fitting methods. At first sight, this might indicate that the influence of outflows is not necessarily more dominant in the central regions. However, given the S/N ratio, we do not probe regions outside the central few kpc. Moreover, given the ground-based seeing of $\sim$1-1.5$\arcsec$ of these Keck long-slit spectra and given the redshift range of our sample, the central 1-2 kpc are essentially unresolved (as mentioned above). We probe the dependency of the $\sigma_{\rm [OIII]}$/$\sigma_{\star}$ ratio on the velocity offset of the second Gaussian, the wing component, with respect to the central core Gaussian component (using the spatially resolved data). For the majority of the objects and rows, the [OIII] profile has a blue wing (see previous section). Fitting this wing with a separate Gaussian results in $\sigma_{\rm [OIII],D}$/$\sigma_{\star}$ = 1.06$\pm$0.02. For objects/rows with a red wing, the core component $\sigma$ ratio is $\sigma_{\rm [OIII],D}$/$\sigma_{\star}$ = 1.01$\pm$0.03. For objects/rows for which the second Gaussian component fitted a broader underlying central component, $\sigma_{\rm [OIII],D}$/$\sigma_{\star}$ = 1.05$\pm$0.06. However, in all three cases, if these non-gravitational kinematic (blueshifted/redshifted/broad central) components are not excluded from the fit by a second Gaussian, they result in an overestimation of $\sigma_{\star}$. For a single Gaussian fit, $\sigma_{\star}$ is overestimated by 50$\pm$4\% for blueshifted wings, by 45$\pm$7\% for redshifted wings, and by 49$\pm$8\% for central broadening. A Gauss-Hermite Polynomial leads to an overestimation of 92$\pm$5\% for blue wings, 94$\pm$10\% for red wings and 82$\pm$12\% for broader central components. This shows the necessity of fitting a double Gaussian for all types of [OIII] profiles (blue wing, red wing or broader center) and considering only the narrow core component as a surrogate for $\sigma_{\star}$. We also checked for dependencies of the $\sigma_{\rm [OIII]}$/$\sigma_{\star}$ ratio on the velocity shift of the entire [OIII] profile compared to the H$\beta$ absorption line from stars. The only noticeable trend is that a handful of objects/rows with large $\sigma_{\rm [OIII]}$/$\sigma_{\star}$ ratio in the core [OIII] (as fitted by the double Gaussian) are among those with large blueshifted [OIII] lines with an offset of at least -150\,km\,s$^{-1}$. However, while [OIII] can be offset by -300\,km\,s$^{-1}$ to 200\,km\,s$^{-1}$, there is no strong trend between the velocity shift and the $\sigma_{\rm [OIII]}$/$\sigma_{\star}$ ratio, regardless of fitting method. The width of the [OIII] wing (when fitted by a double Gaussian) is larger by an average factor of 2.95$\pm$0.06 compared to the [OIII] core, without showing a trend with distance from the center or overall velocity shift of the [OIII] line with respect to the H$\beta$ absorption line. This result is consistent with \citet{woo16} for a sample of $\sim$39,000 type-2 AGNs from SDSS. To look for a possible physical origin of the scatter, we test dependencies of the $\sigma_{\rm [OIII]}$/$\sigma_{\star}$ ratio on other AGN and host-galaxy parameters, taken from our previous publications \citep{ben15,run16}. In particular, we probe the relationship between the $\sigma_{\rm [OIII]}$/$\sigma_{\star}$ ratio and BH mass, as well as $L_{5100}$~luminosity, but do not find a relationship. Likewise, there is no correlation between the $\sigma_{\rm [OIII]}$/$\sigma_{\star}$ ratio and the [OIII]/H$\beta_{\rm narrow}$ flux ratio, host-galaxy morphology, or host-galaxy inclination. This is in line with results by \citet{ric06} who also did not find any trends in residuals when compared to host galaxy and nuclear properties. While our sample consists of radio-quiet objects, we discuss the effect of radio jets further below. Note that while integral-field spectroscopic studies have found increasing evidence of galaxies with kinematically de-coupled stellar and gaseous components with fractions as large as $\sim$30-40\% in elliptical and lenticular galaxies \citep[see e.g.,][and references therein]{sar06,dav11,bar15}, the larger survey of MaNGA finds only 5\% of kinematically misaligned galaxies \citep{jin16}. Moreover, out of these, 90\% reside in early-type galaxies. Given our sample of pre-dominantly late-type galaxies ($\sim$77\% with host galaxies classified as Sa or later; \citealt{ben15}), we expect a negligible fraction of kinematically de-coupled galaxies in our sample. Indeed, the overall gas rotation curve (as traced by [OIII]) matches that of the stellar rotation curve \citep{har12}, with rotational velocities up to $\sim$$\pm$250\,km\,s$^{-1}$. \begin{table*} \caption{Ratios of [OIII] width to stellar-velocity dispersion depending on fitting method and distance from center. Col. (1): Extraction of spectra. Col. (2): Fitting method of [OIII] emission line. Col. (3): Mean and uncertainty (of the mean) of the resulting ratio of [OIII] width ($\sigma_{\rm [OIII]}$) to stellar-velocity dispersion ($\sigma_{\star}$) for all measurements. Col. (4): Same as Col. (3), but for distance from center of 0-2 kpc. Col. (5): Same as Col. (3), but for distance from center of 2-4 kpc. Col. (6): Same as Col. (3), but for distance from center of 4-6 kpc (4-10 kpc in case of reff). } \label{table:sigmacompare} \begin{tabular}{ccccccc} \hline Spectrum & [OIII] Fit & Mean Ratio & Mean Ratio & Mean Ratio & Mean Ratio\\ & & Total & Bin 1 & Bin 2 & Bin 3 \\ (1) & (2) & (3) & (4) & (5) & (6)\\ \hline Spatially resolved & Double Gaussian & 1.06$\pm$0.02 & 1.06$\pm$0.02 & 1.03$\pm$0.04 & 1.4$\pm$0.3\\ & Single Gaussian & 1.49$\pm$0.03 & 1.45$\pm$0.03 & 1.6$\pm$0.1 & 2.2$\pm$0.5\\ & Gauss-Hermite Polynomials & 1.95$\pm$0.05 & 1.85$\pm$0.04 & 2.2$\pm$0.1 & 2.8$\pm$1\\ Within effective radius & Double Gaussian & 1.02$\pm$0.04 & 1.06$\pm$0.05 & 1.08$\pm$0.07 & 0.83$\pm$0.05\\ & Single Gaussian & 1.42$\pm$0.07 & 1.5$\pm$0.1 & 1.5$\pm$0.1 & 1.1$\pm$0.1\\ & Gauss-Hermite Polynomials & 1.74$\pm$0.08 & 1.8$\pm$0.1 & 1.8$\pm$0.1 & 1.4$\pm$0.1\\ \hline \end{tabular} \end{table*} \begin{figure*} \includegraphics[scale=0.42]{distancespatall.eps} \includegraphics[scale=0.42]{distancereffall.eps} \caption{ Ratio of [OIII] width ($\sigma_{\rm [OIII]}$) to stellar-velocity dispersion ($\sigma_{\star}$) as function of distance from galaxy center. Left: for spatially-resolved spectra. Right: for aperture spectra integrated over effective bulge radius. Red data points show average ratio within distance bins 0-2 kpc, 2-4 kpc and 4-6 kpc (4-10 kpc for right panel), respectively.} \label{figure:ratiodistance} \end{figure*} \subsection{Including [OII] in the comparison} We compare the [OII] line width ($\sigma_{\rm [OII]}$) with the [OIII] line width ($\sigma_{\rm [OIII]}$) derived from the three different fitting methods (single Gaussian, double Gaussian using the central component only, and Gauss-Hermite polynomials) and with the stellar-velocity dispersion ($\sigma_{\star}$), in all cases as derived from spectra of the central row or within the bulge effective radius (since these are the only spectra with [OII] width measurements, given the lower S/N of [OII]). Figure~\ref{figure:oiii_oii} shows examples of a direct comparison the [OIII] and [OII] profiles. In Figure~\ref{figure:oii_comparison}, the resulting ratios are shown as a function of [OII] width ($\sigma_{\rm [OII]}$) for the aperture spectra within the bulge effective radius. Table~\ref{table:oiicompare} summarizes the average ratios, both overall as well as a function of [OII] width. To summarize, the [OII] width is smaller than the entire [OIII] line (as represented by fits using a single Gaussian or Gauss-Hermite polynomials), since the [OIII] line has prominent blue and red wings. When these wings are excluded in a double Gaussian fit and when comparing the narrow core component of [OIII] with [OII], the widths are more comparable, but the [OII] line is broader (on average by 17\%). This can be attributed to wings that also appear in the [OII] emission line, especially for larger widths: while for 90\,km\,s$^{-1}$ $<$ $\sigma_{\rm [OII]}$ $<$ 140\,km\,s$^{-1}$, the average ratio is 1.02$\pm$0.03, the [OII] is wider by 12\% for velocities 140\,km\,s$^{-1}$ $<$ $\sigma_{\rm [OII]}$ $<$ 190\,km\,s$^{-1}$ and even up to 28\% wider for velocities 190\,km\,s$^{-1}$ $<$ $\sigma_{\rm [OII]}$ $<$ 240\,km\,s$^{-1}$. This shows that while the lower ionization line has generally less prominent wings from outflows (or infalls), they are nevertheless present, especially for wider lines. The same trend is observed when comparing $\sigma_{\rm [OII]}$ and $\sigma_{\star}$. It is thus recommended to also fit the [OII] emission line with a double Gaussian to exclude inflows and outflows as well, i.e., using the same strategy as for the [OIII] fitting. However, given that [OII] is already a blended doublet line, the fitting of a double Gaussian to each individual line is difficult, especially with low spectral resolution and S/N which can often lead to the fitting of noise in the spectrum instead, as our data showed. Thus, using [OIII] is the better choice between both lines. Our comparison cautions the use of low S/N emission lines (or spectra) such as [OII] for which the fitting of wings is more challenging. Note that the results for [OII] determined from the central spectra are within the uncertainties of those within the bulge effective radius and thus not further discussed here. While the [SII] emission lines have also been found to be a good substitute for $\sigma_{\star}$ \citep{gre05, kom07}, our spectral range does not cover these lines and we cannot make a direct comparison. However, we suspect that [SII], also a line with a lower ionization potential (23 eV), will behave similarly to [OII]. \begin{figure} \includegraphics[width=\columnwidth]{oiii_oiiT.eps} \caption{ Examples of central [OIII] emission line (solid line) compared to [OII] (dash-dotted line). For comparison, the local continuum was subtracted and the peak flux scaled to 1. Since [OII] is a blended doublet line, it is broader than [OIII] in all cases. Blue wings seen in [OIII] are also present in [OII] (e.g., 1355+3834), but sometimes noisy, given the fainter [OII] line (e.g., 2327+1524). \label{figure:oiii_oii}} \end{figure} \begin{figure} \includegraphics[width=\columnwidth]{oii_comparison.eps} \caption{ Ratio of [OII] width ($\sigma_{\rm [OII]}$) to [OIII] width as fitted with different methods ($\sigma_{\rm [OIII, D]}$, $\sigma_{\rm [OIII, S]}$, and $\sigma_{\rm [OIII, GH]}$) as well as to stellar-velocity dispersion ($\sigma_{\star}$; lower panel) as function of [OII] width ($\sigma_{\rm [OII]}$), for aperture spectra within the effective bulge radius. Red data points show average ratio within velocity bins 90-140 km\,s$^{-1}$, 140-190 km\,s$^{-1}$, and 240-290 km\,s$^{-1}$, respectively. } \label{figure:oii_comparison} \end{figure} \begin{table*} \caption{Ratios of [OII] width to [OIII] width and stellar-velocity dispersion depending on fitting method and [OII] width. Col. (1): Extraction of spectra. Col. (2): Fitting method of [OIII] emission line. Col. (3): Mean and uncertainty (of the mean) of the resulting ratio of [OII] ($\sigma_{\rm [OII]}$) width to [OIII] width ($\sigma_{\rm [OIII]}$) and stellar-velocity dispersion ($\sigma_{\star}$) for all measurements. Col. (4): Same as Col. (3), but for [OII] width between 90-140\,km\,s$^{-1}$. Col. (5): Same as Col. (3), but for [OII] width between 140-190\,km\,s$^{-1}$. Col. (6): Same as Col. (3), but for [OII] width between 190-240\,km\,s$^{-1}$. } \label{table:oiicompare} \begin{tabular}{ccccccc} \hline Spectrum & [OIII] Fit & Mean Ratio & Mean Ratio & Mean Ratio & Mean Ratio\\ & & Total & Bin 1 & Bin 2 & Bin 3 \\ (1) & (2) & (3) & (4) & (5) & (6)\\ \hline Effective radius (aperture) & Double Gaussian & 1.17$\pm$0.04 & 1.02$\pm$0.03 & 1.12$\pm$0.03 & 1.28$\pm$0.07\\ & Single Gaussian & 0.86$\pm$0.02 & 0.81$\pm$0.04 & 0.84$\pm$0.02 & 0.87$\pm$0.07\\ & Gauss-Hermite Polynomials & 0.70$\pm$0.02 & 0.64$\pm$0.04 & 0.69$\pm$0.02 & 0.73$\pm$0.06\\ & Stellar-Velocity-Dispersion & 1.15$\pm$0.04 & 0.95$\pm$0.05 & 1.13$\pm$0.04 & 1.4$\pm$0.1\\ \hline \end{tabular} \end{table*} \subsection{Black Hole Mass - $\sigma_{\rm [OIII]}$ relation} We here compare the resulting \mbh-$\sigma_{\rm [OIII]}$ relations with the ``true'' \mbh-$\sigma_{\star}$ relation taken from \citet{ben15}. For comparison samples, we include quiescent galaxies \citep[][72 objects]{mcc13} and reverberation-mapped AGNs \citep[][29 objects; adopting the same virial factor as for our sample; $\log f = 0.71$]{woo15}. The results show that the total [OIII] emission line (as fitted by either a single Gaussian and even more extreme for Gauss-Hermite Polynomials) overestimates $\sigma_{\star}$ and the points scatter to the right of the relation (Fig.~\ref{figure:mbhsigma}, bottom panels). However, when based on $\sigma_{\rm [OIII],D}$, our sample follows the same \mbh-$\sigma_{\star}$ scaling relationship. The systematic offset for the full [OIII] line width is significant, especially since it is of the same order as the expected evolutionary trend out to $z=1-2$ \citep[e.g.,][]{ben10,ben11b} and in the opposite direction. In other words, using the width of the full [OIII] line as surrogate for $\sigma_{\star}$ (e.g., by simply fitting a single Gaussian) in an attempt to study the evolution of the \mbh-\s~relation as done by e.g., \citet[][]{sal13} will suggest a null result, even though there actually is significant evolution. Since our sample spans a small dynamical range in BH mass (6.7$<$$\log$\mbh$<$8.2), and given the uncertainties of \mbh~of 0.4 dex, we cannot determine the slope of the relationship independently. Instead, we fit the data by the linear relation \begin{eqnarray} \log (M_{\rm BH}/M_{\odot}) = \alpha + \beta \log(\sigma/200\,\rm km\,s^{-1}) \end{eqnarray} taking into account uncertainties and keeping the value of $\beta$ fixed to the corresponding relationships of quiescent galaxies (5.64 for \cite{mcc13} and 4.38 for \citet{kor13}) or reverberation mapped AGNs \citep[][3.97]{woo15}. The resulting zero point and scatter of the distribution are comparable to that of the quiescent galaxies. Table~\ref{fits_relations} summarizes the results for $\sigma_{\rm [OIII],D}$, including a comparison to a quiescent galaxies sample taken from \citet[][51 objects; pseudo bulges and mergers excluded]{kor13}. Note that the intrinsic scatter depends on the uncertainties of the measurements. For the quiescent galaxy sample, \mbh~was derived from the kinematics of gas and/or stars within the gravitational sphere of influence of the BH; for the comparison AGN sample, \mbh~was derived more directly through reverberation mapping. Thus, for those samples, the uncertainty on \mbh~is significantly lower, on average 0.2dex and 0.15dex, respectively, compared to 0.4dex for the single-epoch method used for our sample. If, for example, for our \mbh-$\sigma_{\star}$, we artificially assumed an uncertainty of \mbh~of 0.17dex, the scatter would increase from 0.19 to 0.39 (for a fixed slope of 3.97), in other words, comparable to the 0.41 scatter of the reverberation-mapped AGN sample of \citet{woo15}. Thus, the most direct comparison of scatter is between the scatter of the \mbh-\s~relation from \citet{ben15} and that of the \mbh-$\sigma_{\rm [OIII],D}$ relation here, since these are the identical samples with the same uncertainties in the \mbh ~measurements. Independent of assumed fixed slope, we find a smaller scatter in the \mbh-$\sigma_{\rm [OIII],D}$ This is likely due to the fact that $\sigma_{\rm [OIII],D}$ covers a smaller dynamic range than $\sigma_{\star}$ (both within the effective bulge radius); however, since the scatter is within the range of uncertainties, we do not discuss this here further. \begin{table*} \caption{Fits to the local \mbh-$\sigma$ relation, $\log (M_{\rm BH}/M_{\odot}) = \alpha + \beta \log (\sigma / 200 {\rm km\,s}^{-1})$. Col. (1): Sample and sample size in parenthesis. Col. (2): Mean and uncertainty on the best-fit intercept. Col. (3): Mean and uncertainty on the best-fit slope. Col. (4): Mean and uncertainty on the best-fit intrinsic scatter. Col. (4): References for fit. Note that the quoted literature uses FITEXY with a uniform prior on the intrinsic scatter, so our fits assume the same. The results from ``this paper'' are based on using $\sigma_{\rm [OIII],D}$ as surrogate for $\sigma_{\star}$. $^a$ Relation plotted as dashed lines in Fig.~\ref{figure:mbhsigma} and used as fiducial relation when calculating residuals.} \label{fits_relations} \begin{tabular}{llcccc} \hline Sample & $\alpha$ & $\beta$ & Scatter & Reference\\ (1) & (2) & (3) & (4) & (5)\\ \hline Quiescent Galaxies (72) & 8.32$\pm$0.05 & 5.64$\pm$0.32 & 0.38 & \citealt{mcc13}$^a$\\ Quiescent Galaxies (51) & 8.49$\pm$0.05 & 4.38$\pm$0.29 & 0.29 & \citealt{kor13}\\ Reverberation-mapped AGNs (29) & 8.16$\pm$0.18 & 3.97$\pm$0.56 & 0.41$\pm$0.05 & \citealt{woo15}\\ AGNs (66) & 8.38$\pm$0.08 & 5.64 (fixed) & 0.43$\pm$0.09 & \citealt{ben15}\\ AGNs (66) & 8.20$\pm$0.06 & 4.38 (fixed) & 0.25$\pm$0.10 & \citealt{ben15}\\ AGNs (66) & 8.14$\pm$0.06 & 3.97 (fixed) & 0.19$\pm$0.10 & \citealt{ben15}\\ \hline AGNs (62) & 8.41$\pm$0.07 & 5.64 (fixed) & 0.25$\pm$0.11 & this paper (based on $\sigma_{\rm [OIII],D}$)\\ AGNs (62) & 8.23$\pm$0.06 & 4.38 (fixed) & 0.14$\pm$0.09 & this paper (based on $\sigma_{\rm [OIII],D}$)\\ AGNs (62) & 8.16$\pm$0.06 & 3.97 (fixed) & 0.12$\pm$0.08 & this paper (based on $\sigma_{\rm [OIII],D}$)\\ \hline \end{tabular} \end{table*} \begin{figure*} \includegraphics[scale=0.42]{localmsreff.eps} \includegraphics[scale=0.42]{localmsreffoiii.eps} \includegraphics[scale=0.42]{MBHS.eps} \includegraphics[scale=0.42]{MBHGH.eps} \caption{ \mbh-\s~relation. Upper left panel: ``True'' \mbh-\s~relation for 65 objects presented in \citet{ben15} (red open pentagons), reverberation-mapped AGNs \citep[blue;][]{woo15}, and a sample of quiescent local galaxies \citep[black;][with the black dashed line being their best fit]{mcc13}. The error on the BH mass for our sample is 0.4 dex and shown as a separate point with error bar in the legend, to reduce confusion of data points. We assume a nominal uncertainty of the stellar-velocity dispersion of 0.04 dex. Upper right panel: The same as in the left panel, but for $\sigma_{\rm [OIII],D}$ (from aperture spectra within effective bulge radius; our sample only) instead of $\sigma_{\star}$ (as shown in the left panel, also derived within effective bulge radius). Lower panels: The same as in the upper right panel, but using the [OIII] width as fitted by a single Gaussian (left panel) and Gauss-Hermite polynomials (right panel), in both cases clearly overestimating the ``true'' $\sigma_{\star}$. } \label{figure:mbhsigma} \end{figure*} \subsection{Comparison with FIRST} We searched the Very Large Array (VLA) Faint Images of the Radio Sky at Twenty-Centimeters (FIRST) catalog for radio detection. While our sample is radio quiet, out of the 62 objects in the M$_{\rm{BH}}$-$\sigma_{\star}$ relation, 21 have been detected in FIRST, 37 objects have not been detected (FIRST detection limit $\sim$ 1mJy), and 4 objects are outside of the survey area. While for objects not detected in FIRST the ratio $\sigma_{\rm [OIII],D}$/$\sigma_{\star}$ is comparable to the overall average of our sample, i.e., close to 1, (1.05$\pm$0.02 for spatially-resolved data and 0.99$\pm$0.04 for aperture spectra within the bulge-effective radius, respectively), radio-detected objects have a larger width of [OIII], overestimating $\sigma_{\star}$ by 13\% (the ratio is 1.13$\pm$0.03 for spatially-resolved data and 1.13$\pm$0.06 for effective-radius integrated spectra). When probing the broadening as a function of distance from the center, we see a trend that it is more pronounced towards the nucleus. We color-code objects accordingly in the \mbh-$\sigma$ relations (Figure~\ref{figure:radiooiiiwidth}). In the M$_{\rm{BH}}$-$\sigma_{\star}$ relation, objects detected in radio vs. those undetected by FIRST do not form distinct populations. However, when using the width of the core [OIII] emission line (as traced by a double Gaussian, excluding the wing component), there is a trend of objects detected in FIRST having larger widths, especially those with lower \mbh. Our results show that the radio emission, even in these radio-quiet objects, has an effect on the [OIII] emission, broadening its dispersion, even for the core component. This effect has also been observed in radio-loud emission-line galaxies, where the [OIII] central component shows a strong trend of increasing line width with increasing central [OIII] peak shift (i.e., outflow velocity), likely due to strong jet-cloud interactions across the NLR \citep{kom18}. \begin{figure*} \includegraphics[scale=0.42]{radioMBHoiii.eps} \includegraphics[scale=0.42]{radioM.eps} \caption{ Same as in Figure~\ref{figure:mbhsigma}, upper panels, but now distinguishing between objects detected in FIRST (magenta) and those with only upper limits (darkgreen); (literature samples shown in black). (No error bars shown to reduce confusion.) While there is no trend with $\sigma_{\star}$ (left panel), the radio does have a broadening effect on the [OIII] emission line (right panel), even when only considering the core of the line. } \label{figure:radiooiiiwidth} \end{figure*} \subsection{The potential and limitations of [OIII] width as a surrogate for $\sigma_{\star}$} Overall, the results presented above are in agreement with those of previous studies, concluding that the width of the narrow core of the [OIII] emission line can be used as a replacement for $\sigma_{\star}$, albeit with a large scatter \citep{nel00,gre05}, when considering only the central [OIII] component \citep{kom07,woo16}, when excluding sources with a blueshifted central [OIII] component since these objects show strong additional line broadening \citep{kom08}, and when excluding objects with strong radio emission \citep{kom18}. The resulting \mbh-$\sigma_{\rm [OIII],D}$ correlation scatters around the known relation of quiescent galaxies. However, when a direct comparison is made by plotting $\sigma_{\star}$ against $\sigma_{\rm [OIII],D}$, either from spatially-resolved data or integrated within an aperture of the effective bulge radius, there is no strong correlation between the two (Figure~\ref{figure:sigma_sigmaoiii}; Pearson linear correlation coefficients of 0.25 for spatially-resolved data and 0.41 for aperture data; same results for Spearman rank correlation coefficient; see also \citet{liu09}). This holds for both the radio-detected objects in the sample as well as the ones not detected in FIRST. Instead of a direct correlation between $\sigma_{\star}$ and $\sigma_{\rm [OIII],D}$, our data show that they cover the same range, and that their average and standard deviation are similar. Since we did not select on either quantity, but purely on H$\beta$ width\footnote{Note that we are limited by our spectral resolution of 88\,km\,s$^{-1}$.}, this indicates a physical connection and that they feel the same overall gravitational potential. As a consequence, the ratio of $\sigma_{\rm [OIII],D}$ to $\sigma_{\star}$ is close to one with a small deviation of the mean. And since we start out with a \mbh-\s~relation that follows that of quiescent galaxies and reverberation-mapped AGNs, this naturally results in \mbh-$\sigma_{\rm [OIII],D}$ that scatter around the same relation. Given the large uncertainty in \mbh~based on single-epoch masses (0.4 dex), a factor of 2 in $\sigma_{\rm [OIII],D}$/$\sigma_{\star}$ is still not that large. At first sight, the absence of a strong correlation could be due to the fact that we cover a relatively small dynamic range in \mbh, especially given the large uncertainty in \mbh: the range covered is roughly twice the uncertainty. However, this is not true for measurements of $\sigma$: For $\sigma_{\star}$, our sample has a factor of $\sim$3 in dynamic range with a relatively small uncertainty (the range covered is roughly seven times the uncertainty). Thus, the fact that we do not find a close correlation is significant. While we cannot exclude that adding galaxies with larger $\sigma_{\star}$ would result in a trend, especially when considering mainly elliptical galaxies for which the underlying kinematic field is simpler, our sample consisting of AGNs hosted in mostly spiral galaxies \citep[77\% classified as Sa or later;][]{ben15} does not exhibit a significant correlation between $\sigma_{\star}$ and $\sigma_{\rm [OIII],D}$. 8 objects have been conservatively classified as having a pseudo-bulge \citep{ben15}. These objects are not amongst any particular outliers in the $\sigma_{\star}$ and $\sigma_{\rm [OIII],D}$ plots. However, the sample size is small and the classification based on SDSS images for which a morphological classification is difficult, given the presence of the bright AGN point source. We will re-visit the question of pseudo-bulges with higher-resolution images (HST-GO-15215; PI: Bennert). We consider the careful fitting of a double Gaussian, excluding wings and the use of the narrow core component for estimation of the [OIII] width, a robust approach; the $\sigma_{\star}$ measurements were taken with an equally great care \citep{har12}. Our sample further has the advantage of high S/N spatially-resolved spectra, allowing a direct comparison of $\sigma_{\rm [OIII],D}$ and $\sigma_{\star}$ for the same object, using the same spectra and the same aperture. Thus, the reason for the scatter is likely a physical one. Generally speaking, both absorption and emission line profiles are a luminosity-weighted line-of-sight average, depending on the light distribution and the underlying kinematic field which can be different between gas and stars. Also, there may still be effects of outflows, inflows, and anisotropies not accounted for in the double Gaussian fitting of [OIII]. Finally, radiation pressure would only act on gas, not on stars. Our data do not allow to single out any of these as the main cause of the scatter. Given the high quality of our kinematic data, both in terms of S/N of the spectra as well as the detailed fitting, the fact that we do not find a close correlation between $\sigma_{\rm [OIII],D}$ and $\sigma_{\star}$ strongly cautions against the use of $\sigma_{\rm [OIII],D}$ as a surrogate for $\sigma_{\star}$ on an individual basis, even though as an ensemble they trace the same gravitational potential. \begin{figure*} \includegraphics[scale=0.42]{SS_spat.eps} \includegraphics[scale=0.42]{SS.eps} \caption{ Direct comparison between $\sigma_{\star}$ and $\sigma_{\rm [OIII],D}$. The dashed line indicates the one-to-one relation. The left panel shows the result from spatially-resolved data, the right panel from aperture data, integrated within the effective bulge radius. } \label{figure:sigma_sigmaoiii} \end{figure*} | 18 | 8 | 1808.04821 |
1808 | 1808.08187_arXiv.txt | This paper reports on the validation and mass measurement of \thisplanet, a sub-Neptune orbiting a quiet G9V star. Using {\it K2} data from campaigns C5 and C16, we find this planet to have a period of $50.818947\pm 0.000094$\,days and a radius of $2.41\pm0.12$\,\rearth. We followed this system with HARPS-N to obtain 67 precise radial velocities. A combined fit of the transit and radial velocity data reveals that \thisplanet\ has a mass of $14.8\pm3.1$\,\mearth. Its bulk density ($5.7_{-1.4}^{+1.6}$\,g\;cm$^{-3}$) implies that this planet has a significant envelope of water or other volatiles around a rocky core. \thisplanet\ likely formed in a similar way as the cores of the four giant planets in our own Solar System, but for some reason, did not accrete much gas. The planetary mass was confirmed by an independent Gaussian process-based fit to both the radial velocities and the spectroscopic activity indicators. \thisplanet\ belongs to only a handful of confirmed {\it K2} exoplanets with periods longer than 40 days. It is among the longest periods for a small planet with a precisely determined mass using radial velocities. | Both the {\it Kepler} mission and its revived version, the {\it K2} mission, have discovered thousands of exoplanets, uncovering an exciting diversity in the exoplanet population \citep[e.g. ][]{Morton16,Mayo18}. The modified {\it K2} mission differs from the original {\it Kepler} mission in that it does not stare at the same field, but instead visits multiple fields in the Ecliptic Plane, each for about 80 days. This limited timespan makes the mission sensitive to short-period planets only. Only a handful of K2 exoplanets with periods longer than 40 days (half the timespan of a \emph{K2} campaign) have been reported and validated\footnote{according to \url{http://archive.stsci.edu/k2/published_planets/}}. The planet with the longest period within the K2 campaign timespan is K2-118\,b. It has a period of $50.921$\,days and a radius of $2.49$\,\rearth \citep{Dre17}. The faintness of the star ($V\sim14$) impedes obtaining precise radial velocities (RVs). The other validated long-period exoplanets from K2 are the three outer planets (each showing a monotransit) in the five-planet system orbiting HIP41378 \citep{Vand16b} with estimated periods of $156, 131, 324$\,days for planets d, e, and f, respectively. No mass measurements have been reported on this system yet. Precise and accurate masses for planets similar to Earth in size with a variety of orbital periods are essential to understand the transition between rocky and non-rocky planets for small planets. Recently, a gap was found around 2\,\rearth in the distribution of planetary radii of Kepler planets \citep[e.g. ][]{Ful17,Zeng17b,VanE18,Ful18}. Planets with radii below that gap are most likely rocky or Earth-like in composition. However, without a value for the planetary mass, the composition of the planets above the gap remains uncertain. Having a well-characterised sample of small planets spanning a broad variety of parameters, such as orbital period, planetary mass, planetary radius, and various stellar parameters (mass, radius, chemical abundances, ...), can shed light on the formation and evolution history of these planets. This can include their formation location (in terms of the snow line), the amount of planetary migration, and the effects of photo-evaporation amongst other scenarios. In this paper we report on a four-sigma mass measurement of \thisplanet. This planet was labeled as a small planetary candidate in \citet{Mayo18} with an orbital period of $50.82$\,days and a preliminary planetary radius above the radius gap. This paper is structured as follows. Section \ref{observations} describes the obtained data, both from photometry and spectroscopy. We validate the transit in Section \ref{valid}. Stellar properties, including stellar activity indicators, are discussed in Section \ref{star}. Sections \ref{combo} and \ref{rv} describe the two analyses we performed on the light curve and RVs. Finally, we discuss and conclude in Section \ref{concl}. | 18 | 8 | 1808.08187 |
|
1808 | 1808.01545_arXiv.txt | We study the hydrodynamics and nucleosynthesis in the double-detonation model of Type Ia supernovae (SNe~Ia) and the interaction between the ejecta and a surviving white dwarf (WD) companion in the double-degenerate scenario. We set up a binary star system with $1.0M_\odot$ and $0.6M_\odot$ carbon-oxygen (CO) WDs, where the primary WD consists of a CO core and helium (He) shell with $0.95$ and $0.05M_\odot$, respectively. We follow the evolution of the binary star system from the initiation of a He detonation, ignition and propagation of a CO detonation, and the interaction of SN ejecta with the companion WD. The companion (or surviving) WD gets a flung-away velocity of $\sim 1700$~km~s$^{-1}$, and captures $^{56}$Ni of $\sim 0.03M_\odot$, and He of $3 \times 10^{-4}M_\odot$. Such He can be detected on the surface of surviving WDs. The SN ejecta contains a ``companion-origin stream'', and unburned materials stripped from the companion WD ($\sim 3 \cdot 10^{-3}M_\odot$), although the stream compositions would depend on the He shell mass of the companion WD. The ejecta has also a velocity shift of $\sim 1000$~km~s$^{-1}$ due to the binary motion of the exploding primary WD. These features would be prominent in nebular-phase spectra of oxygen emission lines from the unburned materials like SN~2010lp and iPTF14atg, and of blue- or red-shifted Fe-group emission lines from the velocity shift like a part of sub-luminous SNe~Ia. We expect SN~Ia counterparts to the D$^6$ model would leave these fingerprints for SN~Ia observations. | \label{sec:introduction} The progenitor system of Type Ia supernovae (SNe~Ia) is one of the biggest mysteries in astronomy and astrophysics. It is generally thought that an SN~Ia is powered by thermonuclear explosion of a carbon-oxygen (CO) white dwarf (WD). However, the progenitor system is yet to be confirmed. Since a single CO~WD never starts exploding spontaneously, an exploding CO~WD must have a companion star. The stellar type of the companion star has been under debate. There is a famous and long-standing dichotomy between single degenerate \citep[SD; e.g.][]{2018SSRv..214...67N} and double degenerate \citep[DD][]{1984ApJS...54..335I,1984ApJ...277..355W} scenarios, where the companion star is a main-sequence or red-giant star in the SD scenario, or is an another WD in the DD scenario. Other scenarios are also suggested, such as the core degenerate scenario \citep{2011MNRAS.417.1466K}. Recent observations have revealed some significant constraints on the SD scenario. Red-giant stars are absent in the pre-explosion images of SN~2011fe and SN~2014J \citep[][respectively]{2011Natur.480..348L,2014ApJ...790....3K}, which are the closest SNe~Ia in these decades. No main-sequence star has been detected in a supernova remnant LMC SNR 0509-67.5 \citep{2012Natur.481..164S,2017ApJ...837..111L}, although spin-up/spin-down models can explain the non-detection \citep{2011ApJ...730L..34J,2011ApJ...738L...1D,2012ApJ...756L...4H,2015ApJ...809L...6B}. However, we should note some SNe~Ia indicate signals supporting the SD scenario. For example, PTF11kx has given a signature of the interaction of supernova (SN) ejecta and circumstellar matter \citep{2012Sci...337..942D}, and iPTF14atg and SN~2012cg exhibit the interaction of SN ejecta and non-degenerate companion stars \citep[][respectively]{2015Natur.521..328C,2016ApJ...820...92M}, although for iPTF14atg and SN~2012cg these detections have been contested \citep[e.g.][respectively]{2016MNRAS.459.4428K,2018ApJ...855....6S}. SNe~Ia may have several types of progenitor systems, although they may be dominated by a single type of a progenitor system. The DD scenario suffers from the following problem, if one assumes the DD systems are dominant progenitor systems for SNe~Ia. Super-Chandrasekhar DD systems, whose total mass is more than Chandrasekhar mass, merge at a fewer rate than the SN~Ia event rate \citep[e.g.][]{2014ARA&A..52..107M}. In the violent merger model \citep{2010Natur.463...61P}, the primary CO~WD in a DD system is ignited by hydrodynamical effects, and hence super-Chandrasekhar DD systems is not necessarily needed. However, \cite{2015ApJ...807..105S,2016ApJ...821...67S} have shown that the violent merger model works well only when DD systems have the companion mass of $\gtrsim 0.8 M_\odot$; the DD systems are super-Chandrasekhar DD systems. Although \cite{2015ApJ...800L...7K,2017ApJ...840...16K} have suggested spiral instability after DD mergers drives thermonuclear explosions, Sato et al's results have indicated the spiral instability can apply only to super-Chandrasekhar DD systems. \cite{2016MNRAS.462.2486F} have numerically demonstrated thermonuclear explosion of the primary WDs in detached DD systems, and found that the successful DD systems are super-Chandrasekhar DD systems. Another solution could be collisional DD models \citep{2009MNRAS.399L.156R,2009ApJ...705L.128R,2010MNRAS.406.2749L,2015MNRAS.454L..61D}. \cite{2012arXiv1211.4584K} have argued DD collisions in triple systems can account most of SNe~Ia, but it has been controversial \citep{2014ARA&A..52..107M}. If we take into account sub-Chandrasekhar DD systems, whose total mass is less than the Chandrasekhar mass, the total merger rate of super- and sub-Chandrasekhar DD systems would be comparable to the SN~Ia event rate. Such DD systems may explode as SNe~Ia with the aid of helium (He) ignition -- the double detonation model. Originally, the double detonation model has been suggested as a derivative of the SD scenario, since the companion star is a non-degenerate star, such as a He star \citep{1982ApJ...257..780N,1986ApJ...301..601W,1990ApJ...354L..53L,1990ApJ...361..244L}. \cite{2010ApJ...709L..64G} and \cite{2013ApJ...770L...8P} have shown that the primary CO~WD in a DD system can possibly experience CO detonation driven by He detonation. In particular, a DD system in \cite{2010ApJ...709L..64G} is a sub-Chandrasekhar DD system. The double detonation model in DD systems requires only a small amount of He, $\lesssim 0.01M_\odot$, since the He detonation is triggered by hydrodynamical effects of shock compression. This model is called ``Dynamically Driven Double-Degenerate Double-detonation (D$^6$) model'' by \cite{2018arXiv180411163S}, or ``helium-ignited violent merger model'' by \cite{2013ApJ...770L...8P}. Hereafter, we refer to this model as D$^6$ model for simplicity. The D$^6$ model is more advantageous than the double detonation model in SD systems in the following reason. Since the double detonation model in SD systems requires such a large amount of He as $\gtrsim 0.1M_\odot$ \citep{1982ApJ...257..780N}, this model is predicted to leave behind signature of the He detonation \citep{1994ApJ...423..371W,2011ApJ...734...38W}; actually the signature has been found \citep{2017Natur.550...80J,2018ApJ...861...78M}, although the observations of such signature has been rare. The distinct point of the D$^6$ model from other DD models is that the companion WD can survive thermonuclear explosion of the primary WD \citep{2013ApJ...770L...8P}. The DD system is so close that the surviving WD gets hypervelocity (HV) $\gtrsim 10^3$~km after the primary WD explodes. Recently, \cite{2018arXiv180411163S} have found out three HV~WDs from Gaia DR2. If the D$^6$ model can explain all the SNe~Ia in the Milky Way Galaxy, one should find $\sim 30$ HV~WDs within $1$~kpc of the Sun. The number of HV~WDs are fewer than expected. However, if more HV~WDs would be found in near future, it would support the statement that the D$^6$ model is a major origin of SNe~Ia. If the D$^6$ model would be the case for a significant fraction of SNe~Ia, it is important to study SN ejecta and surviving WD of the D$^6$ model. \cite{2010ApJ...709L..64G} and \cite{2013ApJ...770L...8P} have not followed WD explosion although they have investigated the merging process of DD systems, and He detonation. \cite{2015MNRAS.449..942P} have focused only on ejecta-companion interaction, manually setting up the blast wave of SN~Ia explosion. There are several studies for the interaction of SN ejecta with non-degenerate companions \citep[e.g.][]{2008A&A...489..943P,2012A&A...548A...2L,2013ApJ...774...37L}. Therefore, we numerically follow the following sequence of events: the He detonation, CO detonation, WD explosion, and ejecta-companion interaction by means of Smoothed Particle Hydrodynamics (SPH) simulation coupled with nuclear reactions. Although we treat super-Chandrasekhar DD system, we believe super- and sub-Chandrasekhar DD systems have common features in the D$^6$ explosion. Our paper is structured as follows. In section~\ref{sec:method}, we present our SPH simulation method and initial conditions. In section~\ref{sec:result}, we show simulation results. In section~\ref{sec:discussion}, we compare our results with observations of SNe~Ia and HV~WDs. In section~\ref{sec:summary}, we summarize this paper. | \label{sec:discussion} \subsection{Exploding primary white dwarf} First, we discuss SN~Ia counterparts to the D$^6$ model, based on the results shown in section~\ref{sec:result}. Here, we ignore products yielded by the He detonation, which are the heavy Si-group elements (Ar, Ca, and Ti) with such high velocities, $\gtrsim 2 \cdot 10^4$~km~s$^{-1}$. This reason is as follows. We set the He shell of the primary WD to be so thick ($\sim 0.05M_\odot$) that the He and CO detonation easily occurs in our simulation. However, the D$^6$ model would succeed when the He shell is $\lesssim 0.01M_\odot$ \citep{2010ApJ...709L..64G,2013ApJ...770L...8P}, and would indicate smaller signal of these products than our simulation results. The most prominent feature in our SN ejecta is the companion-origin stream (see Figure~\ref{fig:viewSnrPosition}). Owing to the presence of the stream, the abundance of unburned materials has two peaks in velocity space from specific viewing angles, as seen in Figure~\ref{fig:AngleDependent}. The lower velocity component of the unburned materials (a few $10^3$~km~s$^{-1}$) can be observed as oxygen emission lines in nebular-phase spectra. Such oxygen emission lines have been observed in SN~2002cx-likes \citep{2006AJ....132..189J,2007PASP..119..360P}, and a part of SN~2002es-like SN~2010lp \citep{2013ApJ...775L..43T,2013ApJ...778L..18K} and iPTF14atg \citep{2016MNRAS.459.4428K}. In SN~2002cx-likes, possibly explained by pure-deflagration explosion \citep{2005A&A...437..983K} \citep[but see][]{2013A&A...559A..94W}, $^{56}$Ni prevails from the inner to outer ejecta, while our SN ejecta confines $^{56}$Ni to the inner parts with $\lesssim 10^4$~km~s$^{-1}$. Hence, we rule out SN~2002cx-likes as D$^6$ explosion candidates. SN~2010lp and iPTF14atg could be promising counterparts to D$^6$ model, since their light curves and spectral evolutions are consistent with the explosion of sub-Chandrasekhar mass WDs \citep[][respectively]{2013ApJ...778L..18K,2016MNRAS.459.4428K}. Although iPTF14atg have ultraviolet (UV) pulse due to collision of the SN ejecta with the non-degenerate companion \citep{2015Natur.521..328C}, the UV pulse could be explained by surface radio activity of $^{56}$Ni produced by the He detonation \citep{2016MNRAS.459.4428K}. However, our SN ejecta may be inconsistent with SN~2010lp. SN~2010lp has both blue- and red-shifted oxygen emissions in its nebular spectra. On the other hand, our SN ejecta would have either of blue- or red-shifted oxygen emissions, since the companion-origin ejecta stream propagates in one direction from the explosion center. Note that it may be difficult to identify these oxygen emissions, since the companion-origin stream has small mass ($\sim 3 \cdot 10^{-3}M_\odot$). We need to study nebular-phase spectra of the D$^6$ model by performing radiative transfer calculations \cite[e.g.][]{2010ApJ...708.1703M,2017ApJ...845..176B}. We should note that these oxygen feature may not be observed in the following reason. The companion WD should have a He shell in reality, although it does not have in our setup. If the He shell has more than $\sim 3 \cdot 10^{-3}M_\odot$, the SN ejecta may strip only the He shell, not the CO core. Therefore, the companion-origin stream can consist of He materials. Nevertheless, the companion-origin stream can contain CO because of mixing of CO into the overlying He shell during common envelope phase and merging process as mentioned for the He shell of a primary WD. Even though the He shell has more than $\sim 3 \cdot 10^{-3}M_\odot$, C+O is likely mixed in the companion-origin stream. Moreover, a companion WD with larger mass has smaller He shell mass. Eventually, the presence and absence of oxygen futures depends on detail binary parameters of progenitor systems. In future, we will investigate compositions of companion-origin streams in the cases of companion WDs with different total and He shell masses. Another feature is the velocity shift of SN ejecta due to the binary motion of the primary WD, $\sim 10^3$~km~s$^{-1}$. \cite{2011MNRAS.413.3075M} have shown iron and nickel emission lines can be tracers of such a velocity shift. \cite{2018MNRAS.tmpL.101D} have compiled cobalt emissions in nebular spectra of various SNe~Ia, and have found the cobalt emissions are both blue- and red-shifted in SNe~Ia with $-19<M_{\rm B}<-18$ (SN~2007on, SN~2003hv, and SN~2003gs), and either blue- or red-shifted in those with $M_{\rm B}>-18$ (SN~2016brx, SN~2005ke, SN~1999by, and SN~1991bg). Although they have attributed these blue- and red-shifted features to the collisional DD model \citep{1989ApJ...342..986B,2010MNRAS.406.2749L,2009MNRAS.399L.156R,2010ApJ...724..111R,2009ApJ...705L.128R,2012ApJ...759...39H,2015MNRAS.454L..61D}, SNe~Ia with either of blue- or red-shifted Fe-group emissions can be also explained by the D$^6$ model. \subsection{Surviving white dwarf companion} Hereafter, we describe issues related to the surviving WD. \cite{2017ApJ...834..180S} have discussed post-supernova winds blown by radioactive $^{56}$Ni on the surfaces of surviving WDs. We can compare our results with their surviving CO~WD model with $0.6M_\odot$. They have modeled the surface of the surviving CO~WD, such that the mass of radioactive $^{56}$Ni is $0.0003$ -- $0.03M_\odot$, and the entropy of its surface is $1$ -- $3 \cdot 10^8$~erg~g$^{-1}$~K$^{-1}$. As we obtain the $^{56}$Ni mass and entropy on the surface of the surviving WD to be $\sim 0.03M_\odot$, and $4.5 \cdot 10^8$~erg~g$^{-1}$~K$^{-1}$, our simulation results are consistent with their modeling, although materials on the surface in our results are slightly less bound than those in their models. Thus, SN~2011fe could not be explained by the D$^6$ model, which is the same conclusion as theirs. This is because SN~2011fe would be more luminous than observed if it contained a surviving WD. We discuss the surface abundance of the surviving WD. First, we consider the surface pollution by interstellar medium (ISM) and interstellar objects (ISOs). The surviving WD could accrete ISM through the Bondi-Hoyle-Lyttleton accretion. The accreting mass is estimated as \begin{align} M_{\rm acc} &\sim \dot{M}_{\rm acc,BHL} T_{\rm disk} \nonumber \\ &\sim \frac{\pi \rho_{\rm ism} G^2 M_{\rm wd}^2}{c_{\rm s,ism}^4} \frac{h_{\rm disk}}{v_{\rm wd}} \nonumber \\ &\sim 1.0 \cdot 10^{20} \left( \frac{n_{\rm ism}}{1 \mbox{cm}^{3}} \right) \left( \frac{c_{\rm s,ism}}{20\mbox{km~s}^{-1}} \right)^{-4} \left( \frac{h_{\rm disk}}{200\mbox{pc}} \right) \nonumber \\ &\times \left( \frac{M_{\rm wd}}{0.6M_\odot} \right) \left( \frac{v_{\rm wd}}{2000 \mbox{km~s}^{-1}} \right)^{-1} \; \mbox{[g]}, \end{align} where $\dot{M}_{\rm acc,BHL}$ is the mass accretion rate through Bondi-Hoyle-Lyttleton accretion, $T_{\rm disk}$ is time the surviving WD spending in the Galactic disk, $\rho_{\rm ism}$, $n_{\rm ism}$, and $c_{\rm s,ims}$ are, respectively, ISM mass density, number density, and sound speed, $h_{\rm disk}$ is the scale height of the Galactic disk, and $M_{\rm wd}$ and $v_{\rm wd}$ are the mass and velocity of the surviving WD. Note that this estimate constrains on the upper limit of the accreting mass \citep{2006ApJ...638..369K}. Moreover, we estimate a collision rate of the surviving WD with ISOs like 1l/`Oumuamua \citep{2017Natur.552..378M}. The estimate method is the same as in \cite{2018PASJ...70...80T}. Then, the surviving WD collides with ISOs at most once, and accrete the ISO mass of $\sim 10^{13}$~g at most. Eventually, the surviving WD accretes ISM and ISO mass much less than materials captured from the SN ejecta by several orders of magnitude. Hence, ISM and ISOs cannot pollute the surface of the surviving WD. As shown in section~\ref{sec:swd}, the surviving WD captures $^{56}$Ni of $\sim 0.03M_\odot$, and He of $\sim 3 \cdot 10^{-4}M_\odot$. The $^{56}$Ni of $\sim 0.03M_\odot$ actually includes the mushroom-shaped, unburned materials of $\sim 0.01M_\odot$. Even if the unburned materials are not numerical artifacts, they cannot be regarded as anomaly, since the surviving WD also has similar unburned materials. The $^{56}$Ni will undergo radioactive decay. The $^{56}$Ni decay products could be identified as anomalous abundances. However, the decay products do not necessarily stay on the surface of the surviving WD, since they will receive sedimentation \citep{1986ApJS...61..197P,1992ApJS...82..505D}. Note that they can keep their position due to radiative levitation \citep{1995ApJS...99..189C,1995ApJ...454..429C}. It must be necessary to perform sophisticated numerical calculation to follow the time evolution of the surviving WD if we know whether the decay products stay on the surface of the surviving WD. Here, we do not perform such calculations. The surviving WD certainly has He on its surface, since He does not experience sedimentation. However, it is difficult to assess whether a HV~WD is a surviving WD against the D$^6$ explosion on the basis of the detection of He in the following two reasons. First, the detection of He on a HV~WD cannot be the smoking-gun evidence that the HV~WD is a surviving WD against the D$^6$ explosion. Since WDs generally have He on its surface, a HV~WD gets its HV through mechanism other than the D$^6$ explosion. Second, the non-detection of He on a HV~WD's surface does not always deny the HV~WD is a D$^6$ candidate. He on a HV~WD's surface can be seen only when the HV~WD has high temperature on its surface. In fact, \cite{2018arXiv180411163S} did not found He on the surface of their HV~WDs (WD1) for this reason. In summary, we can say a HV~WD is not a surviving WD against the D$^6$ explosion only if He is not detected despite of high temperature on the HV~WD's surface. We reemphasize WD1 in \cite{2018arXiv180411163S} can be a surviving WD against the D$^6$ explosion despite of the non-detection of He, since WD1's surface has low temperature. Since LP~40--365 (or GD~492) has high abundance of Mn and other iron group elements, it is thought to be a WD candidate surviving against the Type Iax explosion \citep{2017Sci...357..680V,2018ApJ...858....3R,2018MNRAS.479L..96R}. Such abundance pattern could not be reconciled with the D$^6$ explosion, since the D$^6$ explosion involves sub-Chandrasekhar mass WD. In order to study features of SN ejecta and surviving WD in the D$^6$ model, we perform SPH simulation of a binary star system with $1.0M_\odot$ and $0.6M_\odot$ CO~WDs, where the primary WD has a He shell with $0.05M_\odot$ mixed with C+O. The primary WD undergoes thermonuclear explosion following the He detonation on the shell and the CO detonation in the core. The SN ejecta collides with the companion WD, and the interaction of the SN ejecta with the companion WD form the ejecta shadow and companion-origin stream. \cite{2015MNRAS.449..942P} have also found out such ejecta shadows in their simulations for their D$^6$ models. The companion WD survives the explosion of the primary WD, and flies away at velocity of $\sim 1700$~km~s$^{-1}$ as the surviving WD. The SN ejecta has typical features of the double detonation explosion on average. However, there are two different features from the double detonation explosion. (1) First, the SN ejecta strips materials of the companion WD, and contains the companion-origin ejecta consisting of C+O. The companion-origin ejecta can make oxygen emission lines in nebular-phase spectra. Therefore, SN~Ia counterparts to the D$^6$ model can be a part of SN~2002es-likes, such as SN~2010lp and iPTF14atg which have oxygen emission lines in their nebular-phase spectra. Note that the compositions of the companion-origin ejecta may depend on the He shell mass of the companion WD. In future, we will investigate this dependence. (2) Second, the SN ejecta has velocity shift of $\sim 1000$~km~s$^{-1}$ due to the binary motion of the exploding primary WD. This velocity shift can result in blue- or red-shifted Fe-group emission lines in nebular-phase spectra seen in sub-luminous SNe~Ia, such as SN~2016brx, SN~2005ke, SN~1999by, and SN~1991bg. The surviving WD certainly has He on its surface. The He originates from residuals of He produced by photo-dissociation at the center of the primary WD. However, since WDs generally have He on their surfaces, the presence of He could not be the smoking-gun evidence of surviving WDs against the D$^6$ explosion. The surviving WD also has $^{56}$Ni decay products on its surface just after it survives the explosion of the primary WD. However, the decay products would experience sedimentation and radiative levitation. In order to determine the surface abundance of the surviving WD, we should follow the long-term evolution of the surviving WD. Finally, we summarize observational features of SNe~Ia under the D$^6$ explosion. At an early time, its light curve may show a UV pulse due to radioactive nuclei yielded by the He detonation. At the maximum-light time, its spectra indicate Si absorption lines similarly to ordinary SNe~Ia. At late times, in the nebular-phase, oxygen emission lines can be observed, where the oxygen originates from the companion-origin stream stripped by the SN ejecta. From specific viewing angles, blue- or red-shifted Fe-group emission lines can be also seen due to the binary motion of the exploding primary WD. | 18 | 8 | 1808.01545 |
1808 | 1808.01773_arXiv.txt | The X-ray luminosity function of distant ($3<z<5.1$) unabsorbed quasars has been measured. A~sample of distant high-luminosity quasars ($10^{45} \leq \LX210 < 7.5 \times 10^{45}$~erg/s in the 2--10~keV energy band) from the catalog given in Khorunzhev~et~al.~(2016) compiled from the data of the {3XMM-DR4} catalog of the {XMM}-Newton serendipitous survey and the Sloan Digital Sky Survey (\SDSS) has been used. This sample consists of \Nmysamplelum \ sources. Most of them (\Nspeccnt) have spectroscopic redshifts $\zspec\geqslant 3$. The remaining ones are quasar candidates with photometric redshift estimates $\zphot\geqslant 3$. The spectroscopic redshifts of eight sources have been measured with \AZT33IK and \BTA \ telescopes. Owing to the record sky coverage area ($\simeq 250$~sq.~deg at X-ray fluxes $\sim 10^{-14}$~erg/s/cm$^{2}$ in the 0.5-2~keV), from which the sample was drawn, we have managed to obtain reliable estimates of the space density of distant X-ray quasars with luminosities $\LX210 > 2 \times 10^{45}$~erg/s for the first time. Their comoving space density remains constant as the redshift increases from $z=3$ to $z=5$ to within a factor of 2. The power-law slope of the X-ray luminosity function of high-redshift quasars in its bright end (above the break luminosity) has been reliably constrained for the first time. The range of possible slopes for the quasar luminosity dependent density evolution model is $\gamma_2=2.78^{+0.00}_{-0.04}\pm0.20$, where initially the lower and upper boundaries of $\gamma_2$ with the remaining uncertainty in the detection completeness of X-ray sources in \SDSS, and subsequently the statistical error of the slope are specified. | A reliable measurement of the X-ray luminosity function of high-luminosity active galactic nuclei (AGNs, hereafter quasars) and its evolution at $z\gtrsim 3$ is one of the most important components of the research on the growth history of supermassive black holes and the evolution of massive galaxies in the Universe. The samples of \XMM\ and {Chandra} extragalacitc X-ray surveys (representative fluxes $\FX0.52\lesssim~10^{-15}$~erg/s/cm$^2$ and areas about one~sq.~deg) turn out to be insufficiently large for the evolution of distant quasars to be studied \citep{civano12, vito14}. The addition of sources from shallower extragalactic surveys ($\FX0.52\sim 10^{-14}$--$10^{-13}$~erg/s/cm$^2$) covering much larger areas (tens of square degrees \citealt{ueda14,aird15,georgakakis15}) improves the situation. \cite{vito14} constructed and extensively studied the luminosity function of quasars at \mbox{$z>3$} with luminosities $\LX210<10^{45}$~erg/s in the 2--10~keV band based on the combined data of several deep X-ray surveys with a total area \mbox{$\simeq 3.3$~sq.~deg}. Using data from the {XMM-XXL} survey with an area of 18~sq.~deg (typical fluxes of sources $\FX0.52\simeq5\times10^{-15}$~erg/s/cm$^2$, \citealt{menzel16}), \cite{georgakakis15}, obtained statistically significant estimates of the quasar luminosity funcion at $z>3$ for even higher luminosities ($\LX210\gtrsim 10^{45}$~erg/s). \cite{ueda14} studied the evolution of the X-ray luminosity function of AGNs based on the collection of data from a large set of X-ray surveys, including the {ROSAT} all-sky survey. The {ROSAT} sample of sources includes several quasars with a very high luminosity ($\LX210>10^{46}$~erg/s) at $z>3$, which allowed the space density of such very luminous and distant quasars to be constrained. This estimate turned out to be in agreement with the predictions of the empirical luminosity function model obtained from samples of sources with a much lower luminosity ($L_{X,2-10}<10^{45}$~erg/s). \cite{kalfonzou14} compiled a catalog of quasars at $z>3$ on an area of $\simeq 33$~sq.~deg based on the archival data of individual nonoverlapping Chandra pointings over the entire time of its operation. Using this catalog, they were able to estimate the space density of distant quasars with luminosities $L_{X,2-10}>5\times 10^{44}$~erg/s and to exclude some of the empirical luminosity function models. However, the size of this sample is still insufficient for a detailed study of the population of most luminous \mbox{($\LX210>10^{45}$~erg/s)} and distant ($z>3.5$) quasars. The data from the XMM-Newton X-ray telescope accumulated over 15 years constitute a serendipitous sky survey \citep{watson09} with a total coverage of $\sim$800~sq.~deg and a sensitivity $\FX0.52\sim 5\times 10^{-15}$~erg\,s$^{-1}$\,cm$^{-2}$ \citep[the {3XMM-DR4} fourth data release of serendipitous source catalog\footnote{\url{http://heasarc.gsfc.nasa.gov/W3Browse/xmm-newton/xmmssc.html}},][]{watson09}. Based on these data, one can produce an X-ray sample of quasars at $z>3$ that exceeds the existing samples by several times \citep{kalfonzou14,georgakakis15} and obtain more rigorous constraints on the luminosity function model parameters. This is the goal of our paper. We made an attempt to find new candidates for distant quasars among the X-ray sources of the {3XMM-DR4} catalog as described in \citep{khorunzhev16,khorunzhev17}. Our goal was to obtain a sample of X-ray quasars at $z>3$ as complete as possible in {XMM}-Newton serendipitous survey fields at Galactic latitudes $|b|>20^\circ$ using photometric data from the Sloan Digital Sky Survey \citep[\SDSS,][]{alam15} as well as the infrared {\2MASS} \citep{cutri03} and {\WISE} \citep{wright10}. The total area of the overlap between these surveys is 300~sq.~deg. The photometric redshift estimates ($\zphot$) had been done by \cite{khorunzhev16AA} and a catalog (\K16) of 903 candidates for distant quasars (presumably of type 1) selected by photometric redshift had been compiled. The catalog includes both previously known quasars (with measured spectroscopic redshifts $\zspec>3$) and new quasar candidates (with photometric redshift estimates $\zphot>2.75$). The additional table of the \K16 \ catalog presents 63 known X-ray quasars with $\zspec>3$ that did not pass the photometric selection of quasar candidates. The first results of our spectroscopic identification of new quasar candidates from the \K16 catalog, based on which we made a quantitative estimate of the purity of this catalog, are presented in \cite{khorunzhev17}, \cite{khorunzhev17b}. The additional selection was shown to provide an increase in the number of new sources at $z>3$ relative to the existing spectroscopic sample of quasars: by $\sim 20$\% for optically bright ($\zmag<20$) and X-ray ($\LX210\gtrsim 10^{45}$~erg/s) luminous sources and by $\sim 50$\% for fainter sources. In this paper we use data from the K16 catalog to measure the space density of luminous ($\LX210>10^{45}$~erg/s) quasars at $z>3$ and to obtain rigorous constraints on the slope of the luminosity function $\gamma_2$ in its bright end. In our calculations we used the following cosmological constants, the same as those in Vito~et~al.~(2014), whose results are actively used below: $H_0=70$~km/s/Mpc, $\Omega_{\rm m}=0.27$, $\Omega_\lambda=0.73$. | In this paper we obtained estimates of the \mbox{X-ray} luminosity function for type~1 quasars for a sample of 101 sources with luminosities $L_{X,2-10} \geqslant 10^{45}$~erg/s from our catalog \citep{khorunzhev16}. The LDDE, LADE, ILDE, and PDE luminosity function models describe equally well the density distribution of unabsorbed quasars. The constraints on the bright end slope of the X-ray luminosity luminosity function ($\gamma_2 =\gamaod\pm \gamaoderr$ for the LDDE model) were improved. The values of $\gamma_2$ and other model parameters depend on the choice of a quasar incompleteness correction for the \K16\ catalog. As the correction increases, the slope $\gamma_2$ becomes steeper and the break luminosity grows. The necessity of taking into account this correction stems from the fact that only for sources with $\zmag<20.5$ we can make photometric redshift estimates using the entire set of SDSS filters, thus improving the reliability and accuracy of $\zphot$. In this case, some of the X-ray luminous quasars at $z>3$ turn out to be fainter than the chosen optical threshold and will be missed in the selection. Most of the \K16 \ sources selected by $\zphot$ are spectroscopically confirmed SDSS quasars. The sample of distant X-ray quasars at luminosities $L_{X,2-10}>10^{45}$~erg/s can be expanded by 20\% by the method of searching for new candidates for distant quasars described in \cite{khorunzhev16}. These candidates are confirmed by the spectroscopic observations performed at the following telescopes: \AZT33IK\ \citep{kamus02} with the ADAM low-resolution spectrograph \citep{afanasev16,burenin16} and \BTA\ with the SCORPIO-I \citep{afanasev05} and SCORPIO-II \citep{afanasev11,afanasev12} spectrographs (see~\cite{khorunzhev17,khorunzhev17b,khorunzhev18b}). The produced X-ray sample of luminous quasars at $z>3$ is one of the most extensive in sky coverage area and number of luminous sources. It can be used as a reference one to estimate the completeness and purity of the methods for the selection of distant quasars and to test the algorithms for optical identifications of X-ray sources from the planned SRG all-sky survey \citep{pavlinsky11,merloni14}. An X-ray quasar at $z=3$ with a \mbox{0.5--2~keV} flux $\simeq10^{-14}$~erg/s/cm$^2$ has a 2--10 keV luminosity $\LX210\simeq 10^{45}$~erg/s. This means that $\zphot$ in SDSS fields can be obtained for $\gtrsim 50$\% of the X-ray quasars at $z\sim3$ found in the SRG/eROSITA survey \citep{merloni14} with fluxes $\gtrsim 10^{-14}$~erg/s/cm$^2$, which corresponds to the average sensitivity of a four-year survey over the sky. It will be possible to refine the break luminosity \mbox{($L_*\simeq 4\times10^{44}$~erg/s)} using the data of deep SRG survey fields near the poles of the ecliptic, where a sensitivity $\FX0.52\simeq 2\times10^{-15}$~erg/s/cm$^2$ will be achieved. We are planning to expand the existing sample of distant X-ray quasars through new X-ray (\XMM) and optical (\SDSS, Pan-STARRS) data, to improve the selection methods (see, e.g., \citealt{mescherakov15}, and to continue the program of their spectroscopic identification with the AZT-33IK and BTA telescopes. | 18 | 8 | 1808.01773 |
1808 | 1808.07891_arXiv.txt | Phosphorus is a crucial element in prebiotic chemistry, especially the P$-$O bond, which is key for the formation of the backbone of the deoxyribonucleic acid. So far, PO had only been detected towards the envelope of evolved stars, and never towards star-forming regions. We report the first detection of PO towards two massive star-forming regions, W51 e1/e2 and W3(OH), using data from the IRAM 30m telescope. PN has also been detected towards the two regions. The abundance ratio PO/PN is 1.8 and 3 for W51 and W3(OH), respectively. Our chemical model indicates that the two molecules are chemically related and are formed via gas-phase ion-molecule and neutral-neutral reactions during the cold collapse. The molecules freeze out onto grains at the end of the collapse and desorb during the warm-up phase once the temperature reaches $\sim$35 K. The observed molecular abundances of 10$^{-10}$ are predicted by the model if a relatively high initial abundance of phosphorus, 5$\times$10$^{-9}$, is assumed. | The detection of new interstellar molecules related with prebiotic chemistry in star-forming regions will allow us to make progress on understanding how the building blocks of life could originate in the interstellar medium (ISM). However, there is a key ingredient that still evades detection: phosphorus, P. This element is essential for life (\cite[Maci\'a et al. 1997]{macia97}), because it plays a central role in the formation of P-bearing macromolecules such as phospholipids, which are the structural components of all cellular membranes, or adenotryphosphate (ATP), responsible for the transfer of energy in cells (\cite[Pasek et al. 2005]{pasek05}). Especially important to basic biochemistry is the P$-$O bond, fundamental for many relevant biological molecules such as phosphate esters, which are essential for the formation of the backbone of the genetic macromolecule deoxyribonucleic acid (DNA). Phosphorus is thought to be synthesized in massive stars and injected to the ISM through supernova explosions (\cite[Koo et al. 2013]{koo13}). It has a cosmic abundance of P/H$\sim$ 3$\times$10$^{-7}$ (\cite[Grevesse \& Sauval 1998]{grevesse98}). It has been detected towards atmospheres of stars (\cite[Caffau et al. 2011]{caffau11}, \cite[Roederer et al. 2014]{roederer14}, \cite[Caffau et al. 2016]{caffau16}), but it is barely detected in the ISM. The ion P$^+$ has been detected in several diffuse clouds (\cite[Jura \& York 1978]{jura78}), and only a few simple P-bearing species (PN, PO, CP, HCP, C$_2$P, PH$_3$) have been identified towards the circumstellar envelopes of very evolved objects (\cite[Tenenbaum et al. 2007]{tenenbaum07}, \cite[De Beck et al. 2013]{debeck13}, \cite[Ag\'undez et al. 2014]{agundez14}). In star-forming regions, only PN has been detected so far (\cite[Turner \& Bally 1987]{turner87}; \cite[Ziurys \& Friberg 1987]{ziurys87}; \cite[Fontani et al. 2016]{fontani16}). Previous searches of PO towards star-forming regions (\cite[Sutton et al. 1985]{sutton85}; \cite[Matthews et al. 1987]{matthews87}) were unsuccessful. \begin{figure*} \centering \includegraphics[scale=0.45]{PO-3mm-W51.eps} \caption{Spectrum observed at 3 mm towards W51. The PO transitions are indicated with blue vertical lines. The lower panels show zoom-in views of the PO transitions. The red line is the LTE fit.} \label{figure-PO-W51-3mm} \end{figure*} Despite the prebiotic importance of P-bearing molecules, their chemistry is still poorly understood in the ISM. The few theoretical models devoted to P-chemistry disagree in both the chemical formation pathways and the predictions of the abundances of PO and PN. While some works (\cite[Millar et al. 1987]{millar87}, \cite[Adams et al. 1990]{adams90}, \cite[Charnley \& Millar 1994]{charnley94}) suggest that PN should be more abundant than PO, other studies involving theoretical modeling and laboratory experiments (e.g. \cite[Thorne et al. 1984]{thorne84}) predict that P should be found mainly in the form of PO. To constrain the chemical models and understand the chemistry of phosphorus in the ISM, astronomical detections of PO and PN in star-forming regions are required. We present observations searching for PN and PO towards two massive star-forming regions: the W51 e1/e2 and the W3(OH) complexes (with luminosities of 1.5$\times$10$^{6}$ and 1.0$\times$10$^{5}$ L$_{\odot}$, respectively). | We report the first detection of the key prebiotic molecule PO towards two star-forming regions: W51 e1/e2 and W3(OH). The derived molecular abundances of PO are $\sim$10$^{-10}$ in both sources. We have found an abundance ratio PO/PN of 1.8 and 3 for W51 e1/e2 and W3(OH), respectively, in agreement with the values estimated for evolved stars. Our chemical modeling indicates that the two molecules are chemically related and are formed via gas-phase ion-molecular and neutral-neutral reactions during the cold collapse. The molecules freeze out onto grains at the end of the collapse, and evaporate during the warm-up phase once the temperature reach $\sim$35 K. Similar abundances of PO and PN are expected during a period of $\sim$5$\times$10$^4$ yr at the early stages of the warm-up phase, when the temperature is in the range 35$-$90 K. The observed molecular abundances require a relatively high initial abundance of atomic phosphorus of 5$\times$10$^{-9}$, 25 times higher than the $"$low-metal$"$ P-abundance typically used in dark cloud chemical models. | 18 | 8 | 1808.07891 |
1808 | 1808.07543_arXiv.txt | We present an analysis of the two-point peculiar velocity correlation function using data from the CosmicFlows catalogues. The Millennium and MultiDark Planck 2 N-body simulations are used to estimate cosmic variance and uncertainties due to measurement errors. We compare the velocity correlation function to expectations from linear theory to constrain cosmological parameters. Using the maximum likelihood method, we find values of $\Omega_m= 0.315^{+0.205}_{-0.135}$ and $\sigma_8=0.92^{+0.440}_{-0.295}$, consistent with the Planck and Wilkinson Microwave Anisotropy Probe CMB derived estimates. However, we find that the cosmic variance of the correlation function is large and non-Gaussian distributed, making the peculiar velocity correlation function less than ideal as a probe of large-scale structure. | \label{sec:introduction} The peculiar velocity field is a sensitive probe of mass fluctuations on large scales and a powerful tool for constraining cosmological parameters. However, the precision measurement of the velocity field is limited by the error in the measurement of radial distance. Many methods have been introduced to measure the distance with the smallest possible error, such as Tully-Fisher (TF) \citep{TullyFisher1977}, Faber-Jackson \citep{FaberJackson1976}, and the Fundamental Plane (FP) \citep{DjoDav1987,DreLynBurDav1987}. These methods do not directly measure radial distance. Rather they estimate the distance modulus, which is proportional to the logarithm of the distance. While the errors in the distance modulus are Gaussian, the distances themselves have a non-Gaussian error distribution which may bias the results. To address this issue, \citet{WatFel2015a} introduced a new unbiased estimator of the peculiar velocity that gives Gaussian distributed errors. In this paper we will use this unbiased estimator together with the measured redshift \citep[see also][]{DavScr2014,TulCouSor2016} to derive the velocity correlation function. The fractional observational errors of the radial distances are typically of the order of $\approx 20$\% \citep[e.g.][]{MasSprHay2006, SprMasHay2007, TulCouDol2013}, and peculiar velocities tend to have errors proportional to the distances, which may be large. Because of this large error, a single peculiar velocity measurement is not a good approximation of the velocity of a galaxy. However, statistical ensembles, especially the low-order moment statistics, may be a good estimator of the cosmic velocity field and thus a good tracer of the underlying mass distribution in the Universe \citep[e.g.][]{FelWat2008,WatFelHud2009,FelWatHud2010, DavNusMas2011,NusBraDav2011,MacFelFer2011,Turnbull2012,MacFelFer2012,Nusser2014,SprMagColMou2014,JohBlaKod2014,ScrDavBlaSta2015}. Many recent studies have focused on the bulk flow, which is the lowest order statistic of the velocity field and is generally thought of as the average of peculiar velocities in a volume \citep[e.g.][]{AbaFel2012,Nusser2014,KumWanFelWat2015,SeiPar2016,ScrDavBla2015,Nusser2016}. Bulk flows are typically calculated using one of two popular methods. The first is the maximum likelihood estimate (MLE) method \citep[e.g.][]{Kaiser1988,WatFel2007}. The MLE formalism estimates the bulk flow as a weighted average of the sample velocities, with the weights calculated to minimize its overall uncertainty given the positions, velocities and errors distributions in the catalogue. The formalism reduces the entire data set to three numbers, namely the components of the bulk flow vector. Since the particular data and error distribution in the surveys analysed are unique to each catalogue, it is difficult to compare the bulk flow calculated using this method between independent surveys. The other popular formulation is the minimum variance (MV) method \citep{WatFelHud2009, FelWatHud2010, ScrDavBlaSta2015}. The MV formalism minimizes the differences between the actual observational data and an 'ideal' survey that may be designed to probe a volume in a particular way. It can be used with Gaussian-weighted \citep{AgaFelWat2012} or tophat-weighted ideal survey distributions \citep{DavNusMas2011,HofNusCorTul2016}. Because it uses a standard ideal survey bulk flow as a reference, it easily lends itself to direct comparisons between independent surveys. Another approach to studying the large-scale velocity field is the pairwise velocity statistic ($v_{12}$) \citep{FerJusFel1999, JusFerFelJaf2000, FelJusFer2003, Hellwing2014, HelBarFre2014}, which is the mean value of the peculiar velocity difference of a galaxy pair at separation $\vect{r}$. Recent studies show that it can also be used to detect the kinetic Sunyaev-Zeldovich effect \citep{ZhaFelJus2008, HanAddAub2012, PlanckXXXVII2015}. In this paper we will use a different approach to probe the cosmic velocity field, namely the peculiar velocity correlation function. It was first introduced by \citet{Gorski1988} and further elucidated in \citet{GorDavStr1989}. In subsequent studies, the velocity correlation function has shown potential for providing interesting constraints on cosmological parameters \citep[e.g.][]{JafKai1995,ZarZehDekHof1997,JusFerFelJaf2000,BorCosZeh2000, AbaErd2009, NusDav2011, OkuSelVla2014, HowStaBla2017}. At the time that the original studies of velocity correlation were done, the small sizes of peculiar velocity catalogues limited its usefulness. The recent availability of large, calibrated catalogues of peculiar velocities and large-scale cosmological simulations suggests that it is worthwhile to revisit the velocity correlation function as a cosmological probe. Here, we will present a feasibility study of this statistic for the study of the large-scale-structure given state-of-the-art peculiar velocity catalogues CosmicFlow-2 (CF2) \citep{TulCouDol2013} and CosmicFlows-3 (CF3) \citep{TulCouSor2016}. In particular, we assess the magnitude of the cosmic variance expected in the correlation function using mock catalogues extracted from the Millennium Simulation\footnote{\label{note1}http://gavo.mpa-garching.mpg.de/Millennium/}. Using this result, we show that the velocity correlation functions calculated from CF2 and CF3 are consistent with the standard cosmological model. However, our results suggest that the particular cosmological volume we live in is on the higher end of the cosmic variance, suggesting that the $\sim$150$h^{-1}$Mpc radius region around us has greater large-scale motions than one would expect on average. The organization of this paper is as follows: in section~\ref{sec:catalogue}, we detail our use of N-body simulations to generate mock surveys. In section~\ref{sec:correlation_function}, we introduce the velocity correlation statistic and the methods used to calculate it. In Section~\ref{sec:error}, we discuss the use of the Monte Carlo method for error analysis. In section~\ref{sec:CR}, we show the results for the velocity correlation function using the CosmicFlows catalogues. In Section~\ref{sec:linear}, we explore using the correlation function to constrain cosmological models. Section~\ref{sec:conclusion} concludes this paper. | \label{sec:conclusion} In principle, the velocity Correlation Function is a powerful statistical tool in exploring the Peculiar Velocity Field and through that, the mass distribution on cosmological scales. We have shown that on average the correlation function calculated from simulated catalogues recovers the expected signal from linear theory, thus demonstrating that it is an unbiased statistic. Since the statistical error in the correlation function is significantly smaller than the cosmic variance, the velocity correlation function does a reasonable job dealing with the large uncertainty inherent in the determination of peculiar velocities of galaxies and groups. However, the non-Gaussian nature of the cosmic variance and redshift distortion put limits on how well we can use this statistic to constrain cosmological parameters. We have calculated the velocity correlation functions for the CosmicFlows-2 and CosmicFlows-3 catalogues and shown that they are consistent with expectations from the standard cosmological model. In addition, we have used our results together with linear theory to constrain the cosmological parameter $\Omega_m$ and $\sigma_8$. In constraining the cosmological parameters, we have assumed Gaussian distributed errors, while the simulations have clearly shown that the error distribution of the cosmic variance has distinct non-Gaussian tails. Furthermore, since the cosmic variance is smaller at larger separations, the covariance matrix gives more weight at larger separations, where skewness is most pronounced and thus, may introduce systematic biased parameter estimations. In addition, redshift distortions give rise to the mismatch between CosmicFlows correlations and linear predictions and thus may contribute further bias to parameter constrains. To mitigate this effect we have used a weighting scheme that combines the effects of cosmic variance and redshift distortion, which appears to be both more stable and less biased. Future studies that account for the non-Gaussian distribution of cosmic variance may result in more robust constraints, particularly with regard to uncertainties in parameter estimation. The systematically larger velocity correlations observed in this study, especially in closer bins, using both the CosmicFlows-2 and and CosmicFlows-3 compilations is consistent with the observed bulk flows from these and other catalogues that is on the larger end of the expected range given the predictions from the $\Lambda$CDM model with CMB derived parameters. However, this excess may also arise from local inhomogeneities in our local volume. | 18 | 8 | 1808.07543 |
1808 | 1808.02220_arXiv.txt | {The total gas mass is one of the most fundamental properties of disks around young stars, because it controls their evolution and their potential to form planets. To measure disk gas masses, CO has long been thought to be the best tracer as it is readily detected at (sub)mm wavelengths in many disks. However, inferred gas masses from CO in recent ALMA observations of large samples of disks in the 1--5 Myr age range seem inconsistent with their inferred dust masses. The derived gas-to-dust mass ratios from CO are between one and two orders of magnitude lower than the ISM value of $\sim$ 100 even if photodissociation and freeze-out are included. In contrast, \textit{Herschel} measurements of hydrogen deuteride line emission of a few disks imply gas masses in line with gas-to-dust mass ratios of 100. This suggests that at least one additional mechanism is removing \ce{CO} from the gas-phase. } {Here we test the suggestion that the bulk of the CO is chemically processed and that the carbon is sequestered into less volatile species such as \ce{CO2}, \ce{CH3OH} and \ce{CH4} in the dense, shielded midplane regions of the disk. This study therefore also addresses the carbon reservoir of the material which ultimately becomes incorporated into planetesimals.} {Using our gas-grain chemical code we performed a parameter exploration and follow the CO abundance evolution over a range of conditions representative of shielded disk midplanes. } {Consistent with previous studies, we find that no chemical processing of CO takes place on 1--3 Myr timescales for low cosmic-ray ionisation rates, $< 5\times 10^{-18}$ s$^{-1}$. Assuming an ionisation rate of $10^{-17}$ s$^{-1}$, more than 90\% of the CO is converted into other species, but only in the cold parts of the disk below 30 K. This order of magnitude destruction of CO is robust against the choice of grain-surface reaction rate parameters, such as the tunnelling efficiency and diffusion barrier height, for temperatures between 20 and 30 K. Below 20 K there is a strong dependence on the assumed efficiency of H tunnelling.} {The low temperatures needed for CO chemical processing indicate that the exact disk temperature structure is important, with warm disks around luminous Herbig stars expected to have little to no \ce{CO} conversion. In contrast, for cold disks around sun-like T Tauri stars, a large fraction of the emitting CO layer is affected unless the disks are young ($< 1$ Myr). This can lead to inferred gas masses that are up to two orders of magnitude too low. Moreover, unless CO is locked up early in large grains, the volatile carbon composition of the icy pebbles and planetesimals forming in the midplane and drifting to the inner disk will be dominated by CH$_3$OH, CO$_2$ and/or hydrocarbons.} | The total gas mass is one of the most fundamental parameters that influences protoplanetary disk evolution and planet formation. Interactions of the gas and dust set the efficiency of grain-growth and planetesimal formation \citep[e.g. ][]{Weidenschilling1977, Brauer2008, Birnstiel2010,Johansen2014}, while interactions of planets with the gaseous disk leads to migration of the planet and gap formation \citep[see, e.g.][for reviews]{Kley2012,Baruteau2014}. Significant amounts of gas are needed to make giant Jovian-type planets. All of these processes depend sensitively on either the total amount of gas or the ratio of the gas and dust mass. Dust masses can be estimated from the continuum flux of the disk, which is readily detectable at sub-millimeter (mm) wavelengths. However, the main gaseous component \ce{H2} does not have any strong emission lines that can trace the bulk of the disk mass, so that other tracers need to be used. Emission from the \ce{CO} molecule and its isotopologues is commonly used as a mass tracer of molecular gas across astronomical environments \citep[for reviews see, e.g. ][]{Dishoeck1987,Bolatto2013,Bergin2017}. \ce{CO} is resistant to photodissociation because it can self-shield against UV-photons and is thus a molecule that can trace \ce{H2} in regions with low dust shielding \citep{Dishoeck1988,Viala1988,Lee1996,Visser2009}. \ce{CO} also has, in contrast with \ce{H2}, strong rotational lines, coming from states that can be populated at 20 K, the freeze-out temperature of \ce{CO}. In most astronomical environments, \ce{CO} is also chemically stable due to the large binding energy of the C--O bond. This chemical stability means that \ce{CO} is usually the second most abundant gas-phase molecule and the main volatile carbon reservoir in molecular astronomical environments. Thus, the recent finding that CO emission from protoplanetary disks is very weak came as a big surprise and implies that CO may be highly underabundant \citep{Favre2013,Bruderer2012, Du2015, Kama2016,Ansdell2016}. Is CO transformed to other species or are the majority of disks poor in gas overall? By extrapolating the chemical behaviour of \ce{CO} from large scale astronomical environments to protoplanetary disks it was expected that only two processes need to be accounted for in detail to determine the gaseous \ce{CO} abundance throughout most of the disk: photodissociation and freeze-out of \ce{CO} \citep{Dutrey1997,vanZadelhoff2001,vanZadelhoff2003}. This was the outset for the results reported by \cite{Williams2014} who computed a suite of disk models with parametrised chemical and temperature structures, to be used for the determination of disk gas masses from the computed line emission of CO isotopologues. This method was expanded by \cite{Miotello2014, Miotello2016} who calculated the temperature, \ce{CO} abundance and excitation self-consistently using the thermo-chemical code DALI\footnote{\url{http://www.mpe.mpg.de/~facchini/DALI/}} \citep{Bruderer2012, Bruderer2013}. \cite{Miotello2016} used a simple gas-grain network that includes CO photodissociation, freeze-out and grain-surface hydrogenation of simple species, but no full grain surface chemistry. DALI also computes the full 2D dust and gas temperature structure, important for determining the regions affected by freeze-out and emergent line emission. Because emission from the main \ce{CO} isotopologue \ce{^{12}C^{16}O} is often optically thick, most observations target the rarer \ce{CO} isotopologues. These do not necessarily follow the highly abundant \ce{^{12}C^{16}O} as \ce{^{12}C^{16}O} can efficiently shield itself from photodissociating UV radiation at lower \ce{H2} column densities compared with the less abundant isotopologues. As such, the rarer isotopologues are dissociated over a larger region of the disk, an effect known as isotope-selective photodissociation \citep[see, for example][]{Visser2009}. The combined effects of the different temperature structure and isotope-selective photodissociation change the emission strengths of the CO isotopologues by up to an order of magnitude compared with the predictions of \cite{Williams2014}. When either of these model predictions including photodissociation and freeze-out are applied to ALMA observations of large samples of disks, still low gas masses are determined: inferred gas masses are close to, or lower than, the calculated dust mass from the same observations instead of the expected 100:1 ratio \citep{Ansdell2016, Miotello2017, Pascucci2016, Long2017}. While it is possible that these disks are indeed very gas depleted, independent determinations of the gas masses such as from far-infrared \ce{HD} data \citep[see, e.g.][]{Bergin2013,McClure2016,Trapman2017} and mass accretion rates \citep{Manara2016b} imply that the \ce{CO}/\ce{H2} abundance ratio is likely much lower than expected, at least in the CO emitting part of the disk. Multiple mechanisms have been proposed to explain this low \ce{CO} abundance, both chemical and physical. A physical argument for the low \ce{CO} abundances comes from the vertical mixing of the gas together with settling of dust. \cite{Kama2016} argued that the low CO abundance in the upper emitting layers of the outer disk can be explained by the constant vertical cycling of gaseous \ce{CO}. Every vertical cycle some \ce{CO} will freeze-out onto grains that have grown and settled below the \ce{CO} snow surface. These larger grains do not cycle back up again to the warmer regions where CO can be returned to the gas. They show that the \ce{CO} abundance can be significantly lowered over the disk lifetime. This mechanism also predicts a strong anti-correlation between age and measured \ce{CO} abundance. The mechanism can explain the destruction of \ce{CO} in the warm layers, such as reported by \cite{Schwarz2016} and at the same time explain the lower than expected \ce{H2O} abundances found in the outer disk of TW Hya and other disks by \cite{Hogerheijde2011} and \cite{Du2017}. However this mechanism cannot explain the low abundance of \ce{CO} inside of the \ce{CO} iceline, the radial location of the snow surface at the midplane, as inferred by \cite{Zhang2017} for TW Hya. Alternatively, there are various chemical mechanisms that destroy \ce{CO}, sometimes referred to as `chemical depletion'. Some of the proposed chemical pathways start with the destruction of gaseous \ce{CO} by \ce{He+}, leading to the formation and subsequent freeze-out of \ce{CH4} \citep{Aikawaflow1999, Eistrup2016} or, when computed at slightly higher temperatures, the gas-phase formation of \ce{C2H2} and subsequent freeze-out and further chemical alteration on the grain-surface \citep{Yu2016}. Another pathway to destroy \ce{CO} is the reaction with \ce{OH} to form \ce{CO2}, either in the gas-phase \citep{Aikawaflow1999}, or on the grain-surface \citep{Furuya2014, Reboussin2015, Drozdovskaya2016, Eistrup2016, Schwarz2018}. The formation of \ce{CO2} through the grain-surface route seems to be most efficient at temperatures around 25 K, just above the freeze-out temperature of \ce{CO}. A third pathway to destroy \ce{CO} is the hydrogenation of \ce{CO} on the dust grain-surface forming \ce{CH3OH} \citep{Cuppen2009, Yu2016, Eistrup2018}. All of these models start with a high abundance of CO and modify the abundance through chemical processes. Alternatively there models that do not have CO initially as they assume that, due to some reset process, the gas is fully ionised or atomic at the start of the calculation \citep{Eistrup2016,Molyarova2017}. Due to the high abundance of \ce{OH} during the transition of atomic to molecular gas, \ce{CO2} can be efficiently formed. At low temperatures (< 50 K) \ce{CO2} becomes the most abundant carbon bearing species. All of these \ce{CO} destruction processes are driven by dissociating or ionising radiation, either UV photons, X-rays or cosmic-rays. In regions where UV photons and X-rays are not able to penetrate, cosmic-rays drive the chemistry, so that the chemical timescales of \ce{CO} processing are strongly dependant on the cosmic-ray ionisation rate \citep{Reboussin2015, Eistrup2018}. Indeed, \cite{Eistrup2016} show that chemical evolution during the disk lifetime in the dense midplane is negligible if the only source of ionisation is provided by the decay of radioactive nuclides. In line with these results, \cite{Schwarz2018} find that even in the warm molecular layers either a cosmic-ray ionisation rate of $10^{-17}$ s$^{-1}$ or a strong X-ray field is needed to significantly destroy \ce{CO}. High cosmic-ray ionisation rates are not expected if the proto-stellar magnetic field is sufficiently strong to deflect galactic cosmic-rays \citep{Cleeves2015}. The goal of this paper is to study the chemical pathways that can destroy \ce{CO} in those regions of the disk that are sufficiently shielded from UV photons such as that near the disk midplane. The effectiveness and timescale of \ce{CO} destruction pathways as functions of temperature, density and cosmic-ray ionisation rate are investigated for comparison with the increasing number of ALMA surveys of CO in disks in the 1-10 Myr age range. We also study the effect of the assumed grain-surface chemistry parameters, in particular the tunnelling barrier width and the diffusion-to-binding energy ratio. To be able to do this study in an as general sense as possible we do not restrict ourselves to any specific disk structure but instead perform a parametric study of temperature, density and cosmic-ray ionisation rate over a range representative of a significant portion of the disk mass. | We performed a kinetic chemical modelling study of the destruction of \ce{CO} under UV shielded, cold (< 40 K) and dense ($10^{6}$--$10^{12}$ cm$^{-3}$) conditions. Both grain-surface and gas-phase routes to destroy \ce{CO} are considered and their efficiencies and timescales evaluated using a gas-grain chemical network. Furthermore we studied the effects of the assumed ice diffusion speed (through the diffusion-to-binding energy ratio, $f_\mathrm{diff}$) and the assumed \ce{H} and \ce{H2} tunnelling efficiency (through the tunnelling barrier width, $a_\mathrm{tunnel}$) on the evolution of the \ce{CO} abundance both in the gas-phase and on the grain-surface. Our findings can be summarised as follows: \begin{itemize} \item \ce{CO} destruction is linearly dependent on the assumed \ce{H2} ionisation rate by energetic particles (cosmic-rays, X-rays) over a large region of the considered physical parameter space. Only high enough cosmic-ray ionisation rates, $> 5\times 10^{-18}$ s$^{-1}$ can destroy CO on a $<$3 Myr timescale (Sec.~\ref{sssc:COdestr}, Fig.~\ref{fig:CR_CO}). \item The chemical processing of \ce{CO} is most efficient at low temperatures. A relation between disk temperature and measured \ce{CO} abundance is expected. The coldest disks would have the lowest \ce{CO} abundances; in contrast, flaring disks around luminous Herbig stars should close to canonical CO abundances. \item At low temperatures, hydrogenation of \ce{CO} is efficient when \ce{CO} is fully frozen-out, leading to a reduction of the total \ce{CO} abundance by $\sim$ 2 orders assuming $\zeta_{\ce{H2}} = 10^{-17}$ s$^{-1}$. This route is only weakly dependent on the temperature and density, as long as \ce{CO} is fully frozen-out (Sec.~\ref{ssc:phys_param}, Fig.~\ref{fig:dens_temp_chem}). \item At temperatures of 20--30 K, just above the desorption temperature of \ce{CO}, formation of \ce{CO2} from the reaction of \ce{CO} with \ce{OH} on the ice is efficient. The \ce{CO} abundance can be reduced by two orders of magnitude in 2--3 Myr for $\zeta_{\ce{H2}} = 10^{-17}$ s$^{-1}$. The formation of \ce{CO2} is more efficient at higher densities (Sec.~\ref{ssc:phys_param}, Fig.~\ref{fig:dens_temp_chem}). \item Gas-phase destruction of \ce{CO} by \ce{He+}, eventually leading to the formation \ce{CH4} and \ce{H2O}, only operates on timescales $> 5$ Myr for $\zeta_{\ce{H2}} = 10^{-17}$ s$^{-1}$. Furthermore, this pathway is only effective at low densities ($< 10^9$ cm$^{-3}$) and in a small range of temperatures (15 -- 25 K). As such, this pathway is not important in the context of protoplanetary disk midplanes (Sec.~\ref{ssc:phys_param}, Fig.~\ref{fig:dens_temp_chem}). \item The assumed tunnelling barrier width ($a_\mathrm{tunnel}$) strongly influences the speed of \ce{CO} hydrogenation with efficient tunnelling leading to fast hydrogenation of \ce{CO}. The \ce{CO} destruction due to \ce{CO2} formation is only weakly dependent on the assumed chemical parameters. Only when $f_\mathrm{diff}$ is increased above 0.35 can this reaction be slowed down (Sec.~\ref{ssc:Chemical_param}, Fig.~\ref{fig:All_chemstudy}). \item \ce{CO2}, \ce{CH3OH} and, on a longer timescale, \ce{CH4} are all abundantly formed in the regions where \ce{CO} is destroyed. Observations of anomalously high abundances of \ce{CH4}, \ce{CH3OH} or \ce{CO2} either in infrared absorption spectroscopy towards edge-on systems, or in infrared emission from the inner disk, can help in distinguishing the chemical pathway responsible for \ce{CO} destruction. \item Vertical mixing can bring gas from the warmer, \ce{CO} richer layers to the lower colder layers, where \ce{CO} can be converted into \ce{CH3OH} or \ce{CO2}. This would allow for the chemical processes to also lower the \ce{CO} abundance in the higher, warmer layers of the disk. \end{itemize} In conclusion, chemical reprocessing of \ce{CO} can have a significant impact on the measured disk masses if sufficient ionising radiation is present in the cold ($< 30$ K) regions of the disk. Further modelling of individual disks will have to show to which degree chemical processes are important and where other physical processes will need to be invoked. | 18 | 8 | 1808.02220 |
1808 | 1808.00013_arXiv.txt | Cosmic acceleration may be due to modified gravity, with effective field theory or property functions describing the theory. Connection to cosmological observations through practical parametrization of these functions is difficult and also faces the issue that not all assumed time dependence or parts of parameter space give a stable theory. We investigate the relation between parametrization and stability in Horndeski gravity, showing that the results are highly dependent on the function parametrization. This can cause misinterpretations of cosmological observations, hiding and even ruling out key theoretical signatures. We discuss approaches and constraints that can be placed on the property functions and scalar sound speed to preserve some observational properties, but find that parametrizations closest to the observations, e.g.\ in terms of the gravitational strengths, offer more robust physical interpretations. In addition we present an example of how future observations of the B-mode polarization of the cosmic microwave background from primordial gravitational waves can probe different aspects of gravity. | Acceleration of the cosmic expansion is a signal of new physics: a cosmological constant vacuum energy, a new scalar field, or new laws of gravity. As we extend the standard model into new theories, we must ensure that the foundation is sound and internally consistent. In particular, the theory should be free of pathologies such as ghosts and instabilities. For modified gravity, there is a wide class within effective field theory, Horndeski gravity the most general scalar-tensor theory with second order equations of motions, that has four free functions of time in addition to the cosmic background expansion. These can also be viewed as four property functions, describing properties of the scalar and tensor sectors and their mixing \cite{Bellini:2014fua}. Parametrization of these functions in a physically meaningful way -- with a clear connection to observables and a sound theoretical foundation -- has been a challenging task fraught with pitfalls \cite{1512.06180,1607.03113}. (Also see, e.g., \cite{1705.01960,1505.00174,koyama} for some theory characteristics dealing with the field definitions rather than the property functions.) Here we examine this in terms of sensitivity and characteristics, concentrating on stability from the theoretical side, while also investigating the impact of very general observational considerations such as agreement with general relativity at early times and possessing characteristics consistent with the late time expansion history (e.g.\ a de Sitter limit). Recently, \cite{Kennedy:2018gtx} has proposed the interesting idea of using stability, in terms of the sound speed of scalar perturbations, as the quantity to parametrize and deriving the property function behavior from this. In our analysis of the function space, and its relation to stability, we can assess the utility and generality of that approach, in addition to elucidating the characteristics of the property function space. Furthermore, we explore the sensitivity to the parametrization used on the physical results and constraints. For example, \cite{Mueller:2016kpu} demonstrated that the strength of modified gravity constraints could vary by almost two orders of magnitude depending on time dependence and priors assumed. This is a key question for the utility and robustness of comparing theory quantities such as property functions or sound speed to observables such as growth and clustering of matter structure and light deflection (gravitational lensing). In Section~\ref{sec:scan} we scan through property function space and elucidate the relation between stability and functional parametrization, and also give an example of an observational effect by calculating the B-mode CMB polarization signature of the property functions. We discuss specific theories in Section~\ref{sec:cases} and compare to analytic stability results. Section~\ref{sec:difq} examines the approach of using an explicitly stable parametrization of sound speed to map out the stable regions of property function space. In Section~\ref{sec:obs} we discuss observationally related issues such as the implications for the modified Poisson equation gravitational strengths $\gm$ and $\gl$, and the impact of a general relativity past and de Sitter asymptotic future on acceptable parametrizations and stability. We conclude in Section~\ref{sec:concl}. | \label{sec:concl} Modified gravity as an explanation for cosmic acceleration is a highly attractive concept, and has been connected to the observations in an increasingly sophisticated manner in recent years. If one wants to extract general physical characteristics of the theory, rather than working within one specific theory (with a particular functional form assumed, and particular values for the parameters assumed), then approaches such as effective field theory or property functions or modified Poisson equations are quite useful. However, these all contain functions that themselves need to be parametrized. Even before engaging in detailed calculations of such parametrized theories one must check that the theory is sound: lacking ghosts and instability. We examined in some detail the relation between the functional parametrization in the property function approach and the stability of the theory: the relation is not trivial. In particular, we showed how the stability evolves with redshift, picking out different regions of parameter space that can have complex structure (see Fig.~\ref{fig:amaba}). The final allowable stable part of parameter space is the intersection of stability for all redshifts. This can exhibit disconnected islands and also shows significant sensitivity to the time dependent form assumed for the property functions, even for the case where only two property functions contribute. Such sensitivity raises questions about the utility of the property function (or EFT) approach to give robust, general conclusions about modified gravity. Exploring this further, we considered a power law time dependence and studied the change in stability region as a function of power law index $s$. We derived various analytic expressions for the stability conditions and related them to two modified gravity theories: $f(R)$ gravity and No Slip Gravity. No Slip Gravity has the interesting property that it is a bounding theory: no theory that lies beyond No Slip Gravity in the relation $\al_B=-r\al_M$, i.e.\ with $r>2$, is stable for $s>3/2$ for all $\almnow<0$. The property function $\al_M$ is particularly interesting since it affects gravitational wave propagation, as well as density perturbations. We exhibited its effect on CMB B-mode polarization from primordial gravitational waves (and late time lensing), illustrating how it scales the power (as it does for late universe gravitational waves as well). A derived property from the property functions is the sound speed of scalar perturbations. We examined the implications of various parametrizations of the property functions on the sound speed, finding a great diversity in its behaviors -- power law dependence giving large $c_s$ at early times, bounded but nonmonotonic variation, both concave and convex variation -- all within the stability criterion and coming from simple power law time dependence of the property functions. This is directly relevant to the attractive idea by \cite{Kennedy:2018gtx} that one could start with enforcing stability by choosing a positive sound speed and then deriving the form of the property function $\al_B(a)$ preserving stability. That is, since $c_s$ is a function of $\al_M$ and $\al_B$, one can choose any two and determine the third function. However, our finding that simple $\al_i$'s give complicated $c_s$ casts some doubt on the approach of parametrizing $c_s(a)$. To explore this, we chose several forms of $c_s(a)$ (and $\al_M(a)$) and calculated the resulting $\al_B(a)$. We found that even if we chose a form $c_s(a)$ close to that predicted from a full theory such as $f(R)$ or No Slip Gravity, the reconstructed $\al_B$ and overall modified gravity was not faithful to the original. It broke essential physical characteristics such as injecting slip into No Slip Gravity or breaking the relation $\al_B=-\al_M$ in $f(R)$ gravity. Moreover, this stability approach was pure but not complete -- it did indeed guarantee stability but it did not (with reasonable guesses for the parametrized function $c_s(a)$) generate standard theories such as $f(R)$ gravity. Another relevant question is whether this stability approach is efficient. Removing the need for a stability check in the Boltzmann code saves computational time, but adding an extra differential equation to solve (and possibly increasing the overall number of parameters because one may have to account for $\al$, or $\al_K$, while it can mostly be ignored in the standard approach) compared to the standard approach of parametrizing $\al_B$ and $\al_M$ can cost time. We checked this and found there was no significant time savings from the stability approach, even when it did not involve an increased number of parameters. Finally, we investigated the impact of observational constraints on allowable parametrizations. One would like to impose that general relativity is restored in the early universe, so all the $\al_i$ go to zero. We explored the resulting implications on the sound speed. Similarly, one might look for a de Sitter state in the asymptotic future, and we discussed its implications on the property functions and sound speed. A useful parametrization that encompasses both these conditions is the ``hill'' form, and we compute $c_s$ and $\al_B$ in this case. We motivated use of the combination $q=\al c_s^2$ which enters the equation for $\al_B$, and showed this can be reasonably fit by the hill form, and in turn the reconstructed $\al_B$ looks qualitatively, if not quantitatively, similar to the input truth. However, we demonstrated that even small inaccuracies in the reconstructed $\al_B$, from residuals of the parametrization of the sound speed, can give rise to significant physical flaws. The denominators of the gravitational strengths $\gm$ and $\gl$ can spuriously pass through zero, giving pathologies. Combined with the lack of fidelity in preserving physical characteristics of known theories such as $f(R)$ and No Slip Gravity, and indeed the difficulty including them using straightforward parametrizations of the sound speed, this means that parametrization in terms of property functions or EFT is highly nontrivial, notwithstanding stability considerations. Parametrizations from the theory side, while undeniably attractive, unfortunately are found to be subject to issues of functional sensitivity and lack of robustness. However, there is a reasonable solution by moving closer to the observables. The gravitational strengths $\gm$ and $\gl$ entering the modified Poisson equations, directly related to growth of matter structure and light deflection, have been demonstrated to give robust and highly accurate descriptions of the observables, as well as key indicators to theory characteristics \cite{Denissenya:2017thl,Denissenya:2017uuc}. Such simple, model independent parametrizations as binning in redshift of these functions can be a highly useful first step in uncovering signatures of modified gravity. | 18 | 8 | 1808.00013 |
1808 | 1808.01824_arXiv.txt | In the past decades, photometric survey of open clusters have produced extensive rotation period measurements on stars with different ages. The results indicate a spin-up phase of stellar rotation during pre-main sequence contraction followed by a spin-down near the ZAMS, and during further evolution on the main sequence. These measurements also show (e.g. Barnes 2003; Meibom et al. 2009; 2011) a bimodal distribution of stellar rotation in young open clusters. Young Sun-like stars tend to group into two distinct populations of slow and fast rotators. These populations lie on narrow sequences in diagrams where the measured rotation periods of the members of a stellar cluster are plotted against their B $-$ V colors. One sequence consists of stars that form a diagonal band of increasing rotation period with increasing B - V color. In young clusters, another sequence of fast rotators is also observed. Few stars lie in the intervening gap between these two sequences. Beyond the age of about 500 Myrs, the two groups of fast and slow rotators converge towards a single distribution of angular velocities. Modelling the evolution of stellar rotation has been the subject of many studies. However, very few papers (Brown 2014) have intended to account for the bimodal distribution of stellar rotation observed in young open cluster. The present work addresses the origin and evolution of this bimodal distribution. Section 2 describe a phenomenological model of stellar rotation. Section 3 present simulation results of the rotation evolution a Sun-like star as a function of its initial rotation rate. It compares the evolution of a normal distribution of stellar rotation periods with measurements of stellar rotation in open clusters of various ages. The results are discussed in Sect. 4. | The appearance of a bimodal distribution of rotation periods in young open clusters and the correlation of the rotation sequences with X-ray emission and Li abundance point towards a scenario where Sun-like stars with a rapid enough rotation after circumstellar disk dispersion experience a short episode of large rotational braking in their early evolution. This catastrophic event is driven by a sudden increase of the mass loss rate due to stellar winds at Rossby number included between 0.13 and 0.3. The resulting increase of the braking torque induces a large rotational shear at the bottom of the convective zone. It occurs on stars with ages included between 20-30 Myrs and $\sim$ 600 Myrs depending on their initial rotation rate after dispersion of their circumstellar disk, thus accounting for the bimodal distribution of stellar rotation observed in clusters with those ages. | 18 | 8 | 1808.01824 |
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1808 | 1808.03010_arXiv.txt | The California-Kepler Survey (CKS) catalog contains precise stellar and planetary properties for the \Kepler\ planet candidates, including systems with multiple detected transiting planets (``multis'') and systems with just one detected transiting planet (``singles,'' although additional planets could exist). We compared the stellar and planetary properties of the multis and singles in a homogenous subset of the full CKS-Gaia catalog. We found that sub-Neptune sized singles and multis do not differ in their stellar properties or planet radii. In particular: (1.) The distributions of stellar properties \mstar, [Fe/H], and \vsini\ for the \Kepler\ sub Neptune-sized singles and multis are statistically indistinguishable. (2.) The radius distributions of the sub-Neptune sized singles and multis with $P > 3$ days are indistinguishable, and both have a valley at $\sim1.8~\rearth$. However, there are significantly more detected short-period ($P < 3$ days), sub-Neptune sized singles than multis. The similarity of the host star properties, planet radii, and radius valley for singles and multis suggests a common origin. The similar radius valley, which is likely sculpted by photo-evaporation from the host star within the first 100 Myr, suggests that planets in both singles and multis spend much of the first 100 Myr near their present, close-in locations. One explanation that is consistent with the similar fundamental properties of singles and multis is that many of the singles are members of multi-planet systems that underwent planet-planet scattering. | Comparisons between planetary systems with multiple planets and those with just one known planet have long been used to probe planet formation. A decade after the discovery of the first multi-planet system around a main sequence star \citep{Butler1999}, \citet{Wright2009} conducted a statistical study of 28 multi-planet systems, all of which were discovered and characterized with radial velocities. They compared the multi-planet systems to systems with only one known planet and found that multi-planet systems were spaced uniformly in log-period (unlike the single-planet systems) and typically had lower eccentricities and \msini\ values than the single-planet systems. More recently, the \Kepler\ Mission \citep{Borucki2010} has detected hundreds of multi-planet systems \citep{Latham2011,Lissauer2011_multis,Fabrycky2014,Lissauer2014,Rowe2014}. In the \Kepler\ multi-planet systems, multiple planet candidates transit the star, resulting in measured orbital periods, planet-to-star radius ratios, and transit durations for each planet. The vast majority of the \Kepler\ planet candidates in multis are bona-fide planets, based on statistical arguments \citep{Lissauer2012,Lissauer2014}. The \Kepler\ multi-planet systems differ from the previously studied RV multi-planet systems in that \Kepler\ was sensitive to smaller (lower-mass) planets. The majority of the \Kepler\ single-planet and multi-planet systems have sub-Neptune sized planets rather than giant planets \citep{Latham2011}. Also, \Kepler\ only detected transiting planets. In systems with multiple transiting planets, the planets are very likely nearly coplanar by virtue of the fact that they all transit \citep{Lissauer2011_multis}. However, not all multi-planet systems must be nearly coplanar. A sufficiently non-coplanar system might result in only one transiting planet detected by \Kepler, although multiple planets might exist. The systems with just one detected transiting planet (``singles'') might belong to the tail of a single underlying distribution that describes systems with multiple detected transiting planets (``multis''). On the other hand, a high fraction of the singles might belong to a population with different formation conditions or a different dynamical history. We would like to understand whether the \Kepler\ singles and multis differ in their orbital and physical parameters. Some orbital parameters of interest include multiplicity, orbital periods, eccentricities, and inclinations. Physical parameters of interest include host star mass, metallicity, and rotation velocity, as well as planet radius and mass. If the singles differ from the multis in their underlying distributions of orbital and/or physical parameters, such a distinction likely points to a divergence in the planet formation and/or evolution of the \Kepler\ singles versus multis. Past research has considered the hypothesis that a large fraction of the \Kepler\ singles belong to a distinct population from the multi-planet systems. Some examples of a distinct population are a dynamically hot population (high mutual inclinations and eccentricities for the singles) or a population with wider spacing in the orbital period ratios for the singles than is typical for the multis. \citet{Lissauer2011_multis} found that the typical mutual inclinations in the \Kepler\ multis were $< 10^\circ$ and noted that these small mutual inclinations seemed inconsistent with the large number of observed singles. \citet{Hansen2013} explored the multiplicity vectors and period distributions of the \Kepler\ singles and multis through a model of \textit{in situ} planet formation. They found that the number of \Kepler\ singles is too high to result from an in situ formation scenario (although the authors required each system to have at least three initially coplanar planets). In another study that required a minimum number of planets per system, \citet{Ballard2016} found that there is an excess of singles among the \Kepler\ M-dwarfs. \citet{Xie2016} used stellar spectra from LAMOST and the transit durations from \Kepler\ lightcurves to estimate of the mean eccentricities and inclinations for singles and multis. They found that the mean eccentricity of the singles was $\sim0.3$, whereas the multis were on nearly circular orbits ($e = 0.04\pm0.04$). In a sample of stars with asteroseismically determined properties, \citet{vanEylen2018_ecc} also found higher eccentricities for the singles than the multis. However, other studies have found no need for a large fraction of the singles to have distinct underlying architectures. \citet{Ford2011} found that the prevalence of TTVs in singles was consistent with the multis, suggesting that many singles belong to compact, multi-planet systems. \citet{Tremaine2012} explored a variety of possible orbital geometries and found that no separate population was needed to explain the apparent excess of \Kepler\ singles, if high mutual inclinations were allowed in a small fraction of the multis. \citet{Fang2012} modeled the transit duration ratios as well as the transiting planet multiplicity. They found that an underlying distribution in which most multi-planet systems have mutual inclination distributions of $<3^\circ$, and 75\% of systems have 1-2 planets with $P < 200$ days (like the solar system) describes the observed planet multiplicities and transit duration ratios. \citet{Gaidos2016} found that with improved stellar parameters and an exponentially-distributed number of planets per star, the large number of M dwarf singles compared to multis announced in \citet{Ballard2016} could be reconciled. \citet{Zhu2018} used spectra from LAMOST to measure the properties of \Kepler\ planet candidate host stars, giving special attention to the differences between multis and singles that did and did not exhibit transit timing variations. They found that the stellar properties of the singles and multis did not differ substantially. \citet{MunozRomero2018} compared the metallicities determined by the California-\textit{Kepler} Survey (described below) for singles and multis and found no significant differences in the stellar metallicities. We push the comparison of the fundamental properties of the \Kepler\ singles versus multis into new regions of parameter space by leveraging the precise stellar and planetary parameters of The California-\Kepler\ Survey (CKS) combined with Gaia DR2. CKS obtained high-resolution (R=60,000) spectra for 1305 \Kepler\ systems with transiting planets \citep{Petigura2017}. The improved stellar and planetary parameters \citep[][CKS II]{Johnson2017} enable a more accurate and precise characterization of the \Kepler\ systems than was previously available, yielding 2025 transiting planet candidates with precise radii and host star properties. Fulton \& Petigura 2018 (CKS VII) revised the stellar properties and planet radii based on parallaxes from the Gaia DR2 catalog \citep{GaiaDR2}. CKS and Gaia have dramatically improved the characterization of the \Kepler\ stellar radii, metallicities, masses, and rotations, as well as the planet radii and equilibrium temperatures, compared to what was available before the CKS project \citep[e.g.,][]{Brown2011}. In this paper (CKS VI), we use the refined stellar and planetary properties presented in CKS VII to compare a large, homogeneous, high-purity sample of \Kepler\ singles and multis. Where applicable, we also examine how the stellar and planetary properties of the multis differ for system with 2, 3, and 4 or more transiting planets. In section \ref{sec:sample}, we discuss the cuts to the CKS catalog needed to generate homogenous samples for comparison. In section \ref{sec:stars}, we compare the distributions of the stellar properties for the singles vs. the multis. In section \ref{sec:planets}, we compare the distributions of the planet radii and orbital periods for the singles vs. the multis. We conclude in section \ref{sec:conclusion}. | \label{sec:conclusion} In this paper, we explored how the physical properties of the CKS systems containing multiple detected transiting planets (multis) compare to systems with just one detected transiting planet (singles). Although other studies have examined the relationships between stellar and/or planetary properties and planet multiplicity, our study presents three advantages: (1) The CKS-Gaia dataset enables the largest, most accurate, and most precise comparison of the fundamental host star properties of the \Kepler\ singles and multis so far. (2) As a result of stringent magnitude, detection threshold, and false positive cuts, our comparison of singles vs. multis suffers from fewer observational biases than other studies. (3) In addition to comparing the properties of singles vs. multis, we compare the host star and planet properties as a function of the number of transiting planets. Our conclusions are as follows: \begin{enumerate} \item \textit{The distributions of stellar mass, metallicity, and projected rotation velocity do not differ significantly for the singles and multis.} The lack of a relationship between stellar physical properties and apparent planet multiplicity suggests that any physical process that preferentially creates ``singles'' occurs late in planet formation and in a manner that is not related to the properties of the host star. Also, stellar properties are not particularly useful in predicting the number of transiting planets around a star. \item \textit{Transiting planets of various multiplicities exhibit a valley in the radius distribution at $\sim1.8~\rearth$}. The statistically indistinguishable size distributions of small planets ($\rpl < 4\rearth$) in singles and multis suggests that the acquisition of and subsequent evaporation of a volatile envelope around the planetary core is the same for the singles and multis. Because photo-evaporation happens within the first 100 Myr and is only effective at short orbital periods, the singles and multis likely arrive near their present orbital distances within the first 100 Myr. \item \textit{For the sub-Neptune sized planets, there is a significant ($p = 0.001$) excess of short-period singles ($P < 3$ days) compared to multis.} However, among the multis, the $\Ntp = 2$, 3, and 4+ orbital period distributions are the same, suggesting that geometrical bias alone is unlikely to explain the excess of short-period singles. False positives are unlikely to mimic such a large number of planets considering the high purity of our sample. The excess of short-period planets could also be the hallmark of a late mechanism, such as planet-planet scattering or tidal migration, that produces more singles than multis at very short orbital periods. \item Hot Jupiters are almost always single transiting planets that orbit high-mass and high-metallicity host stars, but these systems are intrinsically rare. \end{enumerate} Our main finding is that host star properties and planet radii for the majority of the sub-Neptune sized planets have no strong relationship with whether there are multiple transiting planets in the system. The similarity of the singles and multis suggests that they have a common origin. The majority of the singles with $P > 3$ days and $\rpl < 4 \rearth$ likely belong to multi-planet systems with higher mutual inclinations than the CKS multis. Perhaps the singles are multi-planet systems that have undergone planet-planet scattering, resulting in systems of multiple planets in which only one planet transits. By contrast, the multis likely had dynamically quieter histories, as evidenced by their current low-entropy states. How might planet composition relate to multiplicity? At face value, the observed radii and orbital periods of the singles and multis are inconsistent with the prediction in \citet{Dawson2016}. That study predicted that planets in multi-planet systems should have preferentially volatile-rich compositions, whereas the singles should preferentially have rocky compositions. The idea behind the prediction was that the multi-planet systems form a little bit earlier than the singles, while the gas disk is a little bit denser, which contributes to both (1) eccentricity damping, resulting in more circular (hence, stable) orbits for the planets, and (2) more gas-rich planets. In contrast, we observe that both the singles and multis include significant populations of planets larger and smaller than the transition from rocky to volatile-rich planets at $\sim1.8~\rearth$. In other words, there is no evidence that the singles are preferentially rocky, or that the multis are preferentially gas-rich. Furthermore, if there were an especially large population of rocky singles with $10 < P < 30$ days and $1~\rearth < \rpl < 1.8~\rearth$, many such planets would have been detected in the \Kepler\ Mission and included in our sample. Perhaps the singles and multis generally form in compact multi-planet systems. These planets can either form early while gas is abundant (forming volatile-rich planets) or later when there is less gas (forming rocky planets). In either case, whatever small eccentricities the planets have acquired will grow after the gas disk dissipates, sometimes leading to dynamical instability. For instance, \citet{Obertas2017} found that the Lyapunov time can vary by a couple orders of magnitude for compact multi-planet systems based on slight differences in the initial orbital conditions\footnote{Although it might not be obvious how the initial conditions of some compact multi-planet systems (but not others) lead to eventual instability, \citet{Tamayo2018} found that machine learning techniques are able to predict the outcomes of N-body simulations with high fidelity.}. Thus, whether a multi-planet system becomes unstable on a timescale of gigayears might be primarily assigned at birth. On the other hand, external influences such as passing stars might also play a role in dynamically disrupting initially coplanar multis \citep{Spalding2016}. The singles have higher eccentricities on average than the planets in multis \citep{Xie2016, vanEylen2018_ecc}, suggesting that dynamical heating and perhaps instability play a role in the formation of the singles. Single sub-Neptunes at very short orbital periods likely have a different dynamical history than the multis. The singles with $P < 3$ days likely belong to multi-planet systems in which some combination of planet-planet scattering, secular chaos, and/or tidal inspiral has moved the innermost planet close to its star. Additional measurements, especially of the impact parameters, orbital obliquities, eccentricities, and masses of the planets in both singles and multis with $P < 3$ days, will clarify which astrophysical process best explains the apparent excess of sub-Neptune sized singles at short orbital periods. | 18 | 8 | 1808.03010 |
1808 | 1808.01709_arXiv.txt | { We have selected a sample of nearby galaxies from Sloan Digital Sky Survey Data Release 7 (SDSS DR7) to investigate the physical properties variation from blue cloud to green valley to red sequence. The sample is limited in a narrow range in color-stellar mass diagram. After splitting green valley galaxies into two parts---a bluer green valley (green 1) and a redder one (green 2) and three stellar mass bins, we investigate the physical properties variation across the green valley region. Our main results are as following: (i) The percentages of pure bulge and bulge-dominated/elliptical galaxies increase gradually from blue cloud to red sequence while the percentages of pure disk and disk-dominated/spiral galaxies decrease gradually in all stellar mass bins and different environments; (ii) With the analysis of morphological and structural parameters (e.g., concentration (C) and the stellar mass surface density within the central 1Kpc ($\Sigma_{1}$)), red galaxies show the most luminous and compact cores than both green valley and blue galaxies while blue galaxies show the opposite behavior in all stellar mass bins. (iii) A strong negative (positive) relationship between bulge-to-total light ratio (B/T) and specific star formation rate (sSFR) ($D_{4000}$) is found from blue to red galaxies. Our results indicate that the growth of bulge plays an important role when the galaxies change from the blue cloud, to green valley, and to the red sequence. | % \label{sect:intro} A bimodal distribution of optical color \citep{Strateva2001, Blanton2003}, ultraviolet-optical color \citep{Salim2007}, morphologies \citep{Driver2006} or star formation rates (SFR) \citep{Kauffmann2003a, Kauffmann2003b} of galaxies has been unambiguously observed. In color-magnitude diagrams (CMD), the galaxies are divided into ``red sequence'' and ``blue cloud''. Generally speaking, red sequence contains old and quiescent galaxies \citep{Kauffmann2003a}, while blue cloud mainly contains blue star-forming disky galaxies \citep{Kaviraj2014a, Kaviraj2014b}. The galaxies between red sequence and blue cloud are called ``green valley'', which are considered as a transition population. \cite{Wyder2007} and \cite{Jin2014} found that the two-Gaussian fitting to the galaxies in CMD is not sufficient which suggests that the green valley is not a simple mixture of red sequence and blue cloud. Therefore, green valley galaxies can provide us with crucial clues to connect red sequence and blue cloud in terms of star formation quenching and evolution of galaxies. Previous studies have shown that since $z\sim1$ the number and stellar mass of blue galaxies are almost constant while the stellar mass of red galaxies have increased by a factor of 2 $\sim$ 4 \citep{Bell2004, Faber2007}. This scenario supports that the existence of red galaxies requires certain quenching mechanisms to stop or weaken the star formation in blue galaxies \citep{Bell2004}. Different quenching mechanisms have been proposed to explain the transition from blue to red galaxies, such as major mergers \citep{Springel2005,Di Matteo2005}, AGN and supernovae feedback \citep{Di Matteo2005, Nandra2007, Marasco2012}, morphological quenching \citep{Martig2009, Martig2013} and environment quenching \citep{Peng2010}. The quenching mechanisms mentioned above propose restrictions on the galaxy structure. These restrictions provide us with an additional approach to understand the connection between the changes of galaxy structure and the locations that galaxy resides in. For example, elliptical galaxies are more concentrated due to internal or external processes than spiral galaxies. This leaves a motivation for us to explore the connection between the morphology/structure and the star formation activity. \cite{Pan2013} have shown that green valley galaxies have the lower (higher) Gini coefficient (G) (second order moment (M20)) than red galaxies but higher (lower) G (M20) than blue galaxies. The average value of asymmetry parameter (A) for green valley galaxies is also between both red and blue galaxies. Moreover, the strong connection between morphological/structural parameters and star formation activity has been demonstrated in the studies of local galaxy surveys \citep{Kauffmann2003b, Mendez2011, Fang2013}. Some structural thresholds, such as critical stellar mass, stellar surface mass density and S\'{e}rsic index reflect the transformation from blue galaxies to red galaxies. Galaxies above these thresholds tend to be old or quiescent while galaxies tend to be young or active below these thresholds \citep{Kauffmann2003b, Driver2006, Schiminovich2007, Bell2008}. \cite{Cheung2012} and \cite{Fang2013} used a structural parameter, $\Sigma_{1}$ (the stellar mass surface density within the central 1kpc), to investigate whether there is a difference between different galaxy colors, and found $\Sigma_{1}$ is a better indicator for the sequence of galaxies than other parameters. \cite{Bait2017} used multiwavelength data to study the dependence of star formation on the morphology and environment in local universe. Their results suggested that the morphology of galaxies correlates with star formation although environmental effects on morphology are weak. This work will focus on the physical properties variation from blue cloud to green valley to red sequence in a narrow range of optical color and different stellar mass bins. In particular, we try to investigate whether there is a monotonic variation in morphology/structure and star formation, so we split green valley galaxies into two populations (i.e., green 1 and green 2). We also expect to have a further understanding to the environmental effects on star formation activity in the local universe. The outline of the paper is as follows. In section 2, we introduce our sample selection and data. The results and discussions are presented in section 3 and 4. Finally, we present our conclusions in section 5. The cosmological constants adopted in this work are $\Omega_m=0.3, \Omega_\Lambda=0.7$ and $H_0=70$ $km^{-1}$ $s^{-1}$ $Mpc^{-1}$. | \label{sect:conclusion} Using a sample drawn from SDSS DR7, together with the data from other public catalogs, we have investigated the physical parameters variation for blue, green 1, green 2 and red galaxies. In addition, we have also discussed the importance of bulge in quenching star formation. Our main conclusions are as follows: (1) The blue galaxies have the lowest fraction of bulge/elliptical galaxies and the highest fraction of disk/spiral galaxies. However, the fraction is opposite in red galaxies. Green valley galaxies have intermediate fraction. The morphologies show continuous change from blue cloud to green valley to red sequence. It is found that the monotonic variation is independent of the stellar mass and environments. (2) The fraction of green valley galaxies is almost constant while the ratio between blue and green valley galaxies decrease as we move from low to high environmental richness, which indicates that the effects of environment is not on the timescale of crossing green valley but on the time when the blue galaxies start to transform their morphologies. (3) The medians of morphological parameters (C, M20 and G) and structural parameters (n, B/T and $\Sigma_{1}$) show the monotonic decreasing or increasing tendency from blue to green 1 to green 2 to red galaxies, suggesting that the buildup of bulge component play an important role during the morphological transformation. (4) For the galaxies with the same color, massive galaxies have more prominent bulge property. The quenching rate is higher in high environmental richness than that of low environmental richness if we consider the percentage of red galaxies as a indicator for quenching rate in different environments. (5) Combining additional blue and red galaxies, we find there is a strong negative (positive) relationship between B/T and sSFR ($D_{4000}$). In the same range of $u-r$ color, massive galaxies have higher B/T when given the sSFR or $D_{4000}$. The monotonic variation suggests that the physical progresses of strengthening bulge are also the ones of weakening the star formation. | 18 | 8 | 1808.01709 |
1808 | 1808.08302_arXiv.txt | The probabilistic approach to turbulence is applied to investigate density fluctuations in supersonic turbulence. We derive kinetic equations for the probability distribution function (PDF) of the logarithm of the density field, $s$, in compressible turbulence in two forms: a first-order partial differential equation involving the average divergence conditioned on the flow density, $\langle \nabla \cdot {\bs u} | s\rangle$, and a Fokker-Planck equation with the drift and diffusion coefficients equal to $-\langle {\bs u} \cdot \nabla s | s\rangle$ and $\langle {\bs u} \cdot \nabla s | s\rangle$, respectively. Assuming statistical homogeneity only, the detailed balance at steady state leads to two exact results, $\langle \nabla \cdot {\bs u} | s \rangle =0$, and $\langle {\bs u} \cdot \nabla s | s\rangle=0$. The former indicates a balance of the flow divergence over all expanding and contracting regions at each given density. The exact results provide an objective criterion to judge the accuracy of numerical codes with respect to the density statistics in supersonic turbulence. We also present a method to estimate the effective numerical diffusion as a function of the flow density and discuss its effects on the shape of the density PDF. | Supersonic turbulence in molecular clouds plays a crucial role in the process of star formation. The probability distribution function (PDF) of density fluctuations in supersonic turbulence has been extensively investigated \citep[e.g.][]{Vazquez-Semadeni94,Padoan+97,Nordlund+Padoan99,Molina+12} and widely used in theoretical models of star formation \citep{Krumholz+McKee05,Padoan+Nordlund11sfr,Hennebelle+Chabrier11,Federrath+Klessen12}. In star formation models based on turbulent fragmentation, the shape of the density PDF, particularly its high-density tail, is of particular importance, due to its impact on the star formation rate and the predicted stellar initial mass function \citep[e.g.][]{Padoan+97,Padoan+Nordlund02,Hennebelle+Chabrier08,Padoan+Nordlund11sfr}. Numerical simulations of isothermal supersonic turbulence with solenoidal forcing have shown that the density PDF is generally consistent with a lognormal distribution, whereas changes in the equation of state \citep{Passot+Vazquez-Semadeni98,Scalo+98}, the forcing pattern \citep{Federrath+08,Federrath+10}, and the inclusion of gravity \citep{Collins+11,Kritsuk+11} all induce variations in the PDF shape. The density PDFs used in star-formation models are usually based on results from numerical simulations. The theoretical understanding of the origin of such PDFs is still incomplete, with most interpretations of numerical results being heuristic or qualitative. For example, the usual argument that the log-normal distribution is the consequence of a multiplicative process of successive, independent compressions and expansions is purely phenomenological. Also, it is not clear how artificial numerical diffusion that exists in all simulations affects the PDF shape. In this Letter, we study the density statistics from first principles, by deriving kinetic equations of the density PDF. Exact results corresponding to the detailed balance of probability fluxes at steady state are derived using the assumption of statistical homogeneity only (\S 2). We stress that, due to strong nonlinearity, exact results in turbulence are very rare, with the known examples being Kolmogorov's celebrated 4/5 law and similar ones in different flow cases \citep[e.g.][]{Yaglom49, Politano+Pouquet98, Galtier+Banerjee11}. The exact results are used to test the accuracy of numerical simulations in \S~3, and our conclusions are summarized in \S 4. | 18 | 8 | 1808.08302 |
|
1808 | 1808.06259_arXiv.txt | We revisit calculations of the X-ray emission from the warm-hot intergalactic medium (WHIM) with particular focus on contribution from the resonantly scattered cosmic X-ray background (CXB). If the significant part of the CXB emission is resolved into point sources, the properties of the WHIM along the line of sight are recorded in the absorption lines in the stacked spectrum of resolved sources \textit{and} in the emission lines in the remaining diffuse signal. For the strongest resonant lines, this implies a factor of $\sim30$ boost in emissivity compared to the intrinsic emissivity over the major part of the density-temperature parameter space region relevant for WHIM. The overall boost for the 0.5-1 keV band is $\sim4$, declining steeply at temperatures above $10^{6}$ K and over-densities $\delta\gtrsim100$. In addition to the emissivity boost, contribution of the resonant scattering changes relative intensities of the lines, so it should be taken into account when line-ratio-diagnostics from high resolution spectra or redshift determination from low resolution spectra are considered. Comparison between WHIM signatures in X-ray absorption and emission should allow differentiating truly diffuse gas of small overdensity from denser clumps having small filling factor by future X-ray missions. | \label{s:introduction} ~~~~~~ Gravitational collapse of over-dense regions, proceeding via unilateral compression into sheets followed by formation of filaments and compact virialized objects (\citealt{1970A&A.....5...84Z}; see \citealt{2012ARA&A..50..353K} for a recent review), provides steady heating to the baryonic content of the Universe \citep{1972A&A....20..189S}, so that approximately half of its mass attains temperature in between $10^5$ K and $10^7$ K, and some $\sim 15$ \% gets virialized and heated to $10^7$ K, at $z\lesssim 1$ \citep{1999ApJ...514....1C,2001ApJ...552..473D,2006ApJ...650..560C}. The latter portion constitutes the baryonic content of massive galaxy clusters and groups, i.e. intracluster medium (ICM), and it is readily observed via intense thermal X-ray emission and the Sunyaev-Zel'dovich distortion of the Cosmic Microwave Background (CMB, e.g. \citealt{1981ASPRv...1....1S}), thanks to a combination of high temperature and relatively high density ($\gtrsim 200$ the critical density). Detection of the former portion, the warm-hot intergalactic medium (WHIM), however, turns out to be extremely challenging, since its thermal emission is not only weak, due to relatively low density ($\sim 10$ the mean density), but also falls into extreme ultraviolet range, that is both heavily absorbed and contaminated by the interstellar medium (ISM) of our own Galaxy \citep[e.g.][]{2007ARA&A..45..221B}. Besides that, being progressively heated to temperatures comparable to or higher than the hydrogen ionization potential, and also photoionized by the cosmic UV and X-ray background radiation, the WHIM matter at $z\lesssim 1$ lacks sufficient amount of neutral hydrogen to be seen via Ly$\alpha$ absorption in spectra of distant quasars, as the matter at higher redshifts ($z\gtrsim2$) does (\citealt{1998ARA&A..36..267R,2013AJ....145...69L}; see \citealt{2009RvMP...81.1405M} and \citealt{2016ARA&A..54..313M} for recent reviews). As a result, the dominant role of WHIM in the baryonic budget of the Universe at $z\sim0$ is most clearly demonstrated by inability of the all observationally accounted baryons in the Local Universe to amount more than a half of $\Omega_{b}$, measured by the combination of CMB observations with the primordial nucleosynthesis models \citep{1998ApJ...503..518F,2016A&A...594A..13P}. Fortunately, by $z\sim0$, WHIM is already substantially enriched by metals expelled from the star-forming galaxies via intensive galactic-scale outflows, so that it is characterized by typical metallicity $Z\sim0.1Z_{\odot}$ \citep[e.g.][]{2001ApJ...560..599A,2009MNRAS.399..574W}. For gas temperature $\lesssim 10^7$ K, some of the most abundant metals like oxygen and iron stay not fully-ionized, even when photo-ionization by cosmic X-ray background radiation (CXB) is taken into account \citep{2016MNRAS.459..310R}. As a result, X-ray emission of WHIM at such temperatures is dominated by line emission from these atoms \citep{2003PASJ...55..879Y,2010MNRAS.407..544B}. Additionally, thanks to resonant scattering in these lines, WHIM should also reveal itself in absorption against bright background UV/X-ray sources {\citep[e.g.][and references therein]{2007ARA&A..45..221B,2008SSRv..134...25R,2018Natur.558..406N}}. The latter possibility is of particular importance for detection of the gas with the lowest density, since the absorption amplitude is proportional to the column density of the intervining matter (as far as it doesn't become saturated), i.e. it scales linearly with the gas number density, while the thermal emission is proportional to the emission measure, so it scales as the gas density squared \citep[e.g.][]{2008SSRv..134..405P}. In fact, resonant scattering is not a true absorption process by itself, so that intensity decrement in the direction of the bright background sources is compensated by the intensity increment in all other directions (see Fig. \ref{fig:sketch}). Thus, the net effect would cancel out after integration over all directions in the case of isotropic incident radiation field, such as CXB, for instance. That is, a cloud that scatters isotropic radiation field is not visible neither as a bright nor a dark patch on the sky. However, a large fraction of CXB is composed of bright points sources, which can be resolved individually, and hence excluded at some level from any given aperture (as illustrated on the right panel in Fig. \ref{fig:sketch}). The remaining signal will then contain both the unresolved part of the background radiation (with similar absorption features as in resolved part, of course) plus resonantly scattered CXB, which is spatially extended, so that the excluded regions contribute only a small portion of overall signal from the filament. As a result, intrinsic thermal emission from the WHIM will be supplemented by resonantly scattered CXB, which might turn out to be of comparable amplitude, especially for regions of relatively low temperature and density, as has been pointed out by \citet{2001MNRAS.323...93C}. This effect not only significantly boosts possibility for WHIM detection in X-ray emission, but also changes the ratios of resonant lines to the continuum (i.e. their equivalent widths) and to forbidden lines, which is potentially an important diagnostic tool for physical conditions in this medium \citep{2001MNRAS.323...93C,2008SSRv..134..405P}. {In this paper, we revisit the calculations of \citet{2001MNRAS.323...93C} and supplement them with calculations performed using the publicly-available photo-ionization code \texttt{Cloudy} \citep{2017RMxAA..53..385F} in order to produce a spectral model for the X-ray emission from the WHIM illuminated by CXB. We analyse basic characteristics of the predicted emission and compare it with the case when contribution of resonant scattering is ignored. The main results of \citet{2001MNRAS.323...93C} are confirmed and quantified in more detail over an extensive region of the density-temperature phase space relevant for the WHIM. We discuss the emissivity boost and corrections in diagnostics of WHIM based on emission-absorption comparison (e.g. \cite{2008SSRv..134..405P}) and emission line ratios induced by the contribution of the resonant scattering. We take advantage of the density-temperature distribution of the metal mass in the Local Universe extracted from a cosmological hydrodynamical simulation coupled with the ionization state calculations to highlight the parameter regions, for which these predictions are of the highest importance from the observational point of view. We illustrate the predicted X-ray emission from fiducial sheet-like (mean over-density $\sim5$) and filament-like (mean over-density $\sim30$) structures in the Local Universe in more detail, and briefly consider their detectability with the future X-ray missions. } The structure of the paper reads as follows: in Section 2, we describe the calculations of the X-ray emission and present the resulting spectral model. In Section 3, we explore properties of this model and compare it to the purely thermal emission and emission from photoionized medium with no account for the resonant scattering. In Section 4, we discuss observational requirements and detectability of the predicted signal with current and future generations of X-ray observatories. We end up summarizing the main conclusions in Section 5. \centerline{} Throughout the paper, we adopt the base \textit{Planck} flat $\Lambda$CDM cosmology with the Hubble parameter $h=H_{0}/(100~\km~\s^{-1}~\Mpc^{-1})=0.678$ and the baryonic density parameter $\Omega_{b}h^2=0.0223$ \citep{2016A&A...594A..13P}. Consequently, the mean baryonic density of the Universe equals $ \left<\rho_b\right>=\Omega_{b}\,\frac{3H_0^2}{8\pi G}=4.2 \times 10^{-31} \g~\cm^{-3}$ in what follows ($G$ is the gravitational constant). This corresponds to the mean number density of hydrogen $ \left<n_H\right>=1.8\times10^{-7}$ cm$^{-3}~\left(\frac{X}{0.7}\right)$ at $z\sim 0$, the hydrogen mass fraction $X=0.7$ for solar metallicity and $X=0.75$ for primordial abundances. In what follows, we will adopt $n_0=2\times10^{-7}$ cm$^{-3}$ as a fiducial value for simplicity. \begin{figure} \centering \label{fig:sketch} \includegraphics[bb=50 220 600 680,width=1.0\columnwidth]{./figs/sketch_igm_img_re.pdf} \caption{{A schematic illustration of three main signatures of a WHIM layer in X-ray spectral band containing a triplet of lines from some helium-like ion (e.g. O VII). \textit{Left panel} depicts a WHIM layer illuminated by distant X-ray sources (shown as stars). The layer is seen (by a distant observer to the right) in intrinsic X-ray emission produced inside it (top), resonant absorption in the spectra bright background sources (middle), and resonant scattering of the isotropic CXB emission (bottom).} Amplitude of the first effect is proportional to the emission measure integrated over line-of-sight, for the latter two it is proportional to the integrated number density (i.e. column density) of the layer. \textit{Middle panel} demonstrates X-ray spectra characteristic for the the aforementioned signatures, with the prominence of the features strongly exaggerated for clarity in each case. Namely, the emission component (E) possesses both signification continuum and line emission, with comparable amplitude of the resonant(r), intercombination(i) and forbidden(f) triplet components. {The absorption spectrum (A) is the spectrum of the background source (black dashed line) with imposed resonant absorption lines at the redshift of the WHIM layer and also slight decrement in the continuum.} The scattered component (S) is heavily dominated by resonantly scattered emission lines of the most abundant ions, with no contribution from forbidden lines and very weak (Thomson scattered) continuum component. \textit{Right panel} illustrates the observing strategy for detection of the scattered component: emission of bright background sources (marked with crossed red circles), containing the absorption features induced by a WHIM filament (sky projection of which is shown as the hatched region), should be removed from the aperture, so that residual signal will contain both unresolved fraction of the CXB and emitted (E) plus scattered (S) radiation from the WHIM.} \end{figure} | \label{s:conclusions} We revisited the calculation of the X-ray emission from WHIM including contribution of the resonantly scattered CXB emission. The latter should be taken into account given that significant fraction of CXB itself is resolved into point sources and excluded from the X-ray emission integrated over some aperture. We confirm the general conclusions of the previous study by \citet{2001MNRAS.323...93C}, and quantify some of the predictions in more detail. The overall boost of emission in the most prominent resonant lines (O VII, O VIII and Ne IX) equals $\sim 30$, {and it is pretty much uniform across almost the whole region of the density-temperature diagram relevant for WHIM.} After averaging over broader spectral bands, it diminishes to $\sim 6$ for 30-eV wide smoothing and to $\sim 4$ when the full 0.5-1 keV band is considered. The ratio of the scattered to intrinsic emission, however, declines steeply at temperatures above $T\sim 10^6$ K for the whole range of considered densities and at over-densities $\delta \gtrsim 100$ for temperatures between $10^4$ K and $10^5$ K. The comparison between absorption in spectra of background point sources with the total emission from WHIM might be used as a diagnostic for gas properties in this regions of the parameter space. The resonantly scattered O VII line at 0.574 keV should dominate the total (scattered plus intrinsic) emission from the full triplet of He-like oxygen at 0.56-0.58 keV, contributing $\sim80$\% of the flux in this band for the bulk portion of the parameter space relevant for WHIM. Relative contribution of this single line to the total emission in 0.5-1 keV turns out to be $\sim 40$\%, with steep decline, however, at lower densities and higher temperatures, where O VIII takes over as oxygen's most abundant ionization specie. The two brightest lines of O VII and O VIII have comparable intensity for the major portion of WHIM, including the parameters of the reference sheet-like and filament-like structures, which we illustrated in more detail. A significant detection of WHIM in emission might be achieved by an X-ray instrument with effective area $\sim 1000$ deg$^2$ at 0.5-1 keV after 10$^6$s-long exposure over an 1-deg$^2$ region of the sky, which is fully covered by a filament-like structure (Thomson optical depth $\sim 10^{-4}$) at redshift $\sim0.1$. We demonstrate this by simulations of such an observation using the actual response functions of the \textit{SRG}/eROSITA telescope and the full X-ray background (cosmic X-ray background plus Galactic diffuse soft X-ray emission). Future X-ray missions provide great opportunities to study WHIM. This refers both to large area X-ray surveys and deep small area observations with X-ray calorimeters. For the former, the signal can be detected by a cross-correlation of stacked (absorption and emission) X-ray signal with certain tracers of overdensities in the large-scale structure, while for the latter detection (and potentially diagnostics) of the prominent individual filaments at $z\sim0.1$ is the primary goal. The calculated table model, suitable for use in numerical simulations and data analysis, is publicly available at \href{https://www.mpa-garching.mpg.de/~ildar/igm/}{https://www.mpa-garching.mpg.de/\~\,ildar/igm/}. {The extracted scattered-to-intrinsic emissivity ratios are also provided there in the form of FITS images. } | 18 | 8 | 1808.06259 |
1808 | 1808.04762_arXiv.txt | {} {With the aim of performing an analysis of the orientations of galaxy pair systems with respect to the underlying large-scale structure, we study the alignment between the axis connecting the pair galaxies and the host cosmic filament where the pair resides. In addition, we analyze the dependence of the amplitude of the alignment on the morphology of pair members as well as filament properties. } {We build a galaxy pair catalog requiring $r_p < 100\kpc$ and $\Delta V < 500 \kms $ within redshift $z<0.1$ from the Sloan Digital Sky Survey (SDSS). We divided the galaxy pair catalog taking into account the morphological classification by defining three pair categories composed by elliptical-elliptical (E-E), elliptical-spiral (E-S) and spiral-spiral (S-S) galaxies. We use a previously defined catalog of filaments obtained from SDSS and we select pairs located closer than $1\mpc$ from the filament spine, which are considered as members of filaments. For these pairs, we calculate the relative angle between the axis connecting each galaxy, and the direction defined by the spine of the parent filament. } {We find a statistically significant alignment signal between the pair axes and the spine of the host filaments consistent with a relative excess of $ \sim$ 15\% aligned pairs. We obtain that pairs composed by elliptical galaxies exhibit a stronger alignment, showing a higher alignment signal for pairs closer than 200 $\kpc$ to the filament spine. In addition, we find that the aligned pairs are associated with luminous host filaments populated with a high fraction of elliptical galaxies. The findings of this work show that large scale structures play a fundamental role in driving galactic anisotropic accretion as induced by galaxy pairs exhibiting a preferred alignment along the filament direction.} {} | Clusters, filaments, sheets and voids are the building blocks of the cosmic web. It is believed that galaxies in filaments represent about the half of the baryon mass in the universe \citep[][and references therein]{fil3}. In addition, \cite{Libeskind18} performed a comparative analysis of twelve different methods devised to classify the cosmic web finding a good agree between the most of the thechinque with a mass fraction in filaments between $\approx 10\%$ and $\approx 60\%.$ According to the models of hierarchical structure formation, galaxy clusters grow through repeated mergers with other groups and clusters of galaxies \citep{zel, kat, bond, jen, col}. These processes occur anisotropically along preferential directions, indicating that galaxy clusters are fed through filaments containing individual galaxies and galaxy systems \citep{kod,ebe,pim}. The distribution and abundance of filaments may affect the properties of galaxies inhabiting these structures \citep[e.g.][]{f2,f3}. Also, the large scale environments have influence on the formation and evolution of dark matter, it is therefore important to understand the correlations between the properties of halos and the topology of the cosmic web, because it may gives us valuable information about the physics of galaxy formation. On this topic, using simulations \citet{zha} conducted an investigation on the spins of the dark matter halos and the direction of the cosmic filaments. The authors found that both, the spins and the main axes of halos in filaments with masses $M\leq 10^{13} M_{\odot}$, are preferably aligned with the direction of the filaments, while spins and the main axes of halos in sheets tend to remain parallel. Also the authors found that with increasing halo mass the major axis tends to be more strongly aligned with the direction of the filament, but the alignment between the halo spin and filament becomes weaker for higher halo mass. In the similar direction, \cite{chen} from the study of the high-resolution hydrodynamical cosmological simulations, found that the galaxy alignment signal along filaments increases significantly with the subhalo mass. Moreover, \cite{Libeskind12} using dark matter cosmological simulation examined the large-scale orientation of substructures and haloes with respect to the cosmic web, finding that the orbital angular momentum of subhaloes tends to align with the intermediate eigenvector of the velocity shear tensor for all haloes in knots, filaments and sheets. In addition, \cite{tyl13} show that the spin axis of spiral galaxies is found to align with the host filament, and also the minor axes of ellipticals are found to be preferentially perpendicular to hosting filaments. Furthermore, there exists a relation between satellite galaxies and their large scale environment \citep[][and references therein]{shao}. For example, based on numerical simulations \citet{bar15} predict a statistical excess of satellite galaxies with main axis aligned in the direction of the central galaxy. Evidence of this relation can be seen in the satellite population of M31, suggesting that tidal effects may have played an important role on its evolution. On this line, by using simulations and observational data \citet{fil1} study the dependence on the alignment amplitude of satellite galaxies with respect to filaments, finding a statistically significant alignment signal between satellite position and filament axis. Also, they show that this alignment depends on the color/luminosity of the system, then it is stronger when the primary and satellite galaxies are brighter. In addition, the authors suggest that the alignment signal may be a consequence of the satellite accretion via streams along the direction of the filaments. In addition, \citet{guo} showed that the satellite luminosity function of galaxies in filaments is significantly higher than those of galaxies not in filaments. The authors also found that the filamentary structures can increase the abundance of the brightest satellites by a factor of $\approx$2, and this is independent of the primary galaxy magnitude. One the observational side, the pioneer study of \citet{lam88} accounts for a preferential distribution of bright galaxies according to their environment. The authors studied a sample of bright galaxies in rich clusters finding that, at scales up to 15 $\mpc$, galaxy counts are consistently higher in the direction of the major axes of bright clusters. In this line, \citet{bin} found that the orientation of the galaxy distribution in two neighboring clusters tend to be similar and, in addition, the brightest cluster galaxies has a tendency to be aligned with the distribution of galaxies in the system. On the other hand, \citet{don06} developed a study of a sample of luminous red galaxies (LRGS) extracted from the fourth release of Sloan Digital Sky Survey (SDSS) within the redshift range 0.4 $<z <$0.5. They found a clear sign of alignment between the orientations of the LRGS and the distribution of galaxies within 1.5 $\mpc$. This alignment effect is present only for red tracers while the orientation of the LRGS is anti-correlated with the population of blue neighboring galaxies. These results could indicate the existence of a preferential direction of accretion in clusters, which also promotes the orientation of the brightest galaxies in the system. \citet{zha1} show that the major axes of galaxies in filaments are preferentially aligned with directions of the filaments, while galaxies in sheets have its major axes parallel aligned to the plane of the sheets. The strength of this alignment signal is stronger for red central galaxies, in agreement with results found for dark matter halos in N-body simulations \citep{lib,fil1}, suggesting that central red galaxies are well aligned with their host halos. These results are consistent with the works of \cite{hirv} and \cite{f3} who find a preferential alignment of red galaxies with the axis of SDSS filaments in the catalog of \cite{fil}. Galaxy interactions play an important role on the formation of the galaxies, since affect almost every aspect of the evolution of these objects \citep{alo06, woods07, elli10, lam03,lam12, mesa}. The presence of a close galaxy companion drives a clear enhancement in galaxy morphological asymmetries, and this effect is statistically significant up to projected separations of at least 50 $\kpc$ \citep{patton16}. Galaxy mergers have a relevant impact on the star formation activity since can trigger starbursts, and affect the galaxy stellar mass function \citep{gama}. The large scale environment can also affect the properties of interacting galaxies in pairs. On this line, \cite{fil2} calculated the angle between the line connecting galaxies of a sample of pairs with separations up to 1$\mpc$ and the direction of its host filament. The authors found that loose pairs, (i.e. pairs with projected separations greater than 300 $\kpc$) have a clear signal of alignment, with at least 25\% excess of aligned pairs, when compared with a random distribution. Motivated by these results, in this work we study the properties of close pair galaxies and the relation of a preferred orientation of pairs with respect to the underlying larger structures in which they are immersed. This paper is structured as follows: Section 2 describes the data used in this work, a detailed description of the catalog of filaments, and the procedure used to construct the pair catalog, explaining the classification process of the sample. In Section 3 we performed an study of the relative orientation of pairs of galaxies and filaments. Finally in Section 4, we summarize our main conclusions. Throughout this paper we adopt a cosmological model characterized by the parameters $\Omega_m=0.3$, $\Omega_{\Lambda}=0.7$ and $H_0=100 \kms \rm Mpc ^{-1}$. | We use a spectroscopic sample derived from SDSS and select galaxy pairs considering projected distances $r_p < 100 \kpc$ and radial velocity differences $\Delta V < 500 \kms$, within $z<0.1$. With the aim of understanding the impact of morphology in our studies, we have used the Galaxy Zoo catalog to divide the samples into pairs composed of two elliptical galaxies (E-E), one elliptical and one spiral (E-S) and two spiral galaxies (S-S). In order to study the presence of alignment effects with larger structures we use the filament catalog of \citet{fil} and select galaxy pairs located closer than 1$\mpc$ from the filament spine. To avoid the effect of the peculiar velocities on the results, we work in projection on the plane of the sky. We compute the angle $\alpha$ between the axis connecting pair members and the direction axis of filaments and we consider the distribution function $F(\alpha)$ to measure the excess of pairs. We measure the ratio $\beta$ of aligned to antialigned pairs as a suitable measure of the preferred orientation effect with typical values $\beta \sim 1.15$. Finally we study the dependence of the alignment on filament and galaxy properties Our main results can be summarized as follow: \begin{itemize} \item We determine that pairs composed by two elliptical galaxies tend to be strongly aligned with the parent filament spine. This tendency increases for pairs closer to the axis of the filament. \item Pairs composed by two spiral galaxies show a weaker alignment signal that increases for pairs at larger separations from the filament. \item The global properties of the filaments affect significantly the pair alignment signal: \item The number of galaxies in host filaments is higher for associated aligned E-E pairs compared to host filaments in antialigned E-E systems. \item On the other hand, aligned S-S pairs reside in less luminous filaments than non-aligned S-S systems. \\ \end{itemize} We obtain a significant dependence of the relative alignment of the pair orientations and nearby filaments on both galaxy morphology and distance to the filament. Close galaxy pairs show a preference to be aligned with the filamentary structures, particularly those formed by elliptical galaxies. In general, reported alignment signal in previous works have concerned samples of ellipticals \citep{lam88,don06}. We stress the fact that here, samples of close pairs composed by spirals also exhibit significant alignment signals. The alignment of galaxy pairs and filaments was studied in \cite{fil2}, where the effects are detected for pairs of galaxies with relative separations of the order of 1$\mpc$. In the present paper, this reported alignment signal extends to close galaxy pairs at significantly smaller relative separations $\sim 30 - 60 kpc$ which are near the process of undergoing a merger event. In this context, the findings of this study show evidence that the global environment has a key role in driving accretion of galaxy pairs with a preferred orientation along the filament direction. Moreover, we find that the morphology of pair member galaxies and global features of the cosmic filaments are important issues to take into account in this large-scale -- galaxy interplay. | 18 | 8 | 1808.04762 |
1808 | 1808.06883_arXiv.txt | Many aspects of the evolution of stars, and in particular the evolution of binary stars, remain beyond our ability to model them in detail. Instead, we rely on observations to guide our often phenomenological models and pin down uncertain model parameters. To do this statistically requires population synthesis. Populations of stars modelled on computers are compared to populations of stars observed with our best telescopes. The closest match between observations and models provides insight into unknown model parameters and hence the underlying astrophysics. In this brief review, we describe the impact that modern big-data\index{big data} surveys will have on population synthesis\index{population synthesis}, the large parameter space problem that is rife for the application of modern data science algorithms, and some examples of how population synthesis is relevant to modern astrophysics. | \label{IzzardSec1}The evolution of binary stars\index{binary stars} is often quite different to their solitary cousins \citep{2017PASA...34....1D}. Many types of stars \emph{only} form in binaries, and their properties allow investigation into astrophysics that is impossible to probe in single stars. Good examples include type Ia supernovae, merging neutron stars (kilonovae) and black holes, the most massive stars e.g.~blue stragglers, many massive Wolf-Rayet stars, X-ray binaries, long and short gamma-ray bursts, barium/CH/CEMP-\emph{s} stars, Algols with peculiar surface chemistry for their evolutionary state, W~Uma contact binaries, low-mass helium stars and related sdB and sdO stars, sequence D variable stars, and thermonuclear novae. Other processes in which binaries are implicated include the formation of asymmetric planetary nebulae, many post-(asyptotic-)giant-branch stars with circumbinary discs and overmassive stars in the Galactic thick disc. Triple and higher-multiple systems are often hierarchical, e.g.~triples are effectively a long-period outer binary system containing a short-period inner binary. Many aspects of binary-star physics remain highly uncertain. Mass transfer, accretion and loss, and associated angular momentum redistribution, e.g.~common envelope evolution, dominates our ignorance of binary-star evolution. One way to improve our knowledge is to apply statistical techniques to compare our stellar models to the stars we observe. This process is called stellar population synthesis. In this brief review, we consider the challenge and opportunity that big data from modern astronomy brings, the basic technique of binary star population synthesis, and the associated parameter space problem. As an example of binary population synthesis we consider how the chemical ejecta from stars depend on model input parameters. We also show an example of how modelling metal-poor stars led to better astrophysical understanding. Finally, we report on recent population synthesis work, in particular with respect to predicting and understanding the discovery of gravitational waves from stellar and black-hole mergers. | 18 | 8 | 1808.06883 |
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1808 | 1808.08857_arXiv.txt | The understanding of black holes in loop quantum gravity is becoming increasingly accurate. This review focuses on the possible experimental or observational consequences of the underlying spinfoam structure of space-time. It adresses both the aspects associated with the Hawking evaporation and the ones due to the possible existence of a bounce. Finally, consequences for dark matter and gravitational waves are considered. | The Planck length is $10^{15}$ times smaller than scales probed at colliders. Linking quantum gravity with observations is therefore extremely hard (see, {\it e.g.}, \cite{Barrau:2017tcd} for a recent review and \cite{Hossenfelder:2009nu,Liberati:2011bp,AmelinoCamelia:2008qg} for complementary viewpoints). Most works devoted to the connection of quantum gravity with experiments are focused on cosmology or astroparticles physics. In the cosmological sector, the main goal consists in calculating scalar and tensor power spectra (see, {\it e.g.}, \cite{eucl3,Agullo2}), together with the background dynamics (see, {\it e.g.}, \cite{Martineau:2017sti,Ashtekar:2011rm}). In the astroparticle physics sector, the main idea is to investigate the possible consequences of the granular structure of space (see, {\it e.g.}, \cite{Vasileiou:2015wja} for a recent investigation). \\ Although black holes (BH) have been intensively studied in quantum gravity, those investigations were mostly disconnected from observations and focused on consistency issues. Recovering, at the leading order, the Bekenstein-Hawking entropy is, for example, obviously a major requirement for all tentative theories (see, {\it e.g.}, \cite{G.:2015sda} and references therein). Curing the central singularity -- understood as a classical pathology -- is another one (see, {\it e.g.}, \cite{Ashtekar:2005qt,Bojowald:2005ah}). Solving the information paradox (see, {\it e.g.}, \cite{Mathur:2009hf} and references therein) would also be highly desirable (this is clearly connected to the previous issues).\\ In this article, we focus on black holes as possible probes for loop quantum gravity (LQG). We begin by a very short summary of the basics of black hole physics in this framework. We then switch to consequences for the Hawking evaporation, considering different possible perspectives. The quite recent (within the LQG setting) hypothesis of black holes bouncing into white holes is presented with the possible associated signals. Finally, we critically review the possible links with dark matter and conclude with prospective for gravitational waves. | The description of black holes in loop quantum gravity has much improved in the last years. A globally consistent picture is now emerging. In this article we have reviewed its possible experimental consequences.\\ The main results are the following :\\ \begin{itemize} \item First, the Hawking evaporation spectrum should be modified in its last stages. We have shown that it could not only allow for the observation of a clear signature of LQG effects but also, in principle, to the discrimination between different LQG models. In particular, holographic models lead to specific features. The value of the Barbero-Immirzi parameter could even by measured. \item Second, attempts to calculate the greybody factors were presented. They should keep a subtle footprint of the polymerization of space and of the existence of a non-vanishing minimum area gap. \item Third, it was emphasized that a local quantum gravity perspective would lead to an observable modification to the Hawking spectrum (line structure), even {\it arbitrarily} far away from the Planck mass. This prediction is not washed out by the secondary emission from the BH. \item Fourth, a model with BHs bouncing into white holes with a characteristic time proportional to $M^2$ was presented and shown to have astrophysical consequences. It can be fine-tuned to explain whether fast radio bursts or the Fermi gamma-ray excess, depending on the values of the parameters. The possible associated background was also studied. A specific redshift dependence allows to discriminate the model from other possible explanations. \item Fifth, the possibility of having a large amount of dark matter in the form of white holes appearing after a quantum gravitational tunneling is presented together with possible weaknesses and future improvements of the model. \item Sixth, observable effects on gravitational wave detections associated with the BHs spin distribution expected are presented. \item Seventh, promising prospects for quasinormal modes are outlined. \end{itemize} It could be that black holes will play a major role in making quantum gravity become an experimental science. | 18 | 8 | 1808.08857 |
1808 | 1808.04624_arXiv.txt | {\textls[-10]{The nucleus of \hen is a short orbital period (4.2\,h) spectroscopic binary, whose status as potential supernovae type Ia progenitor has raised some controversy in the literature. We present preliminary results of a thorough analysis of this interesting system, which combines quantitative non-local thermodynamic (non-LTE) equilibrium spectral modelling, radial velocity analysis, multi-band light curve fitting, and state-of-the art stellar evolutionary calculations. Importantly, we~find that the dynamical system mass that is derived by using all available {He} $\Roman{Rco}{\scriptsize \text{II}}$ lines does not exceed the Chandrasekhar mass limit. Furthermore, the individual masses of the two central stars are too small to lead to an SN\,Ia in case of a dynamical explosion during the merger process.}} \keyword{binaries: spectroscopic; stars: atmospheres; stars: abundances; supernovae} \begin{document} | The detection and study of progenitor systems of type Ia supernovae (SN Ia) are crucial to understand the exact explosion mechanism of these important cosmic distance indicators. Although there is a general consensus that only the thermonuclear explosion of a white dwarf can explain the observed features of those events, the nature of their progenitor systems still remains elusive. In the single-degenerate model, a white dwarf accretes material from a non-degenerate companion and explodes when it reaches the Chandrasekhar mass limit \cite{Whelan73}. An alternative scenario is the double-degenerate model, in which the explosion is triggered during the merger process of two white dwarfs \cite{Iben84, Webbink84, Shen2015}. Identifying progenitors for the latter scenario is particularly interesting in view of the applicability of SN Ia as standardisable candles, as the merging system could exceed or fall below the Chandrasekhar limit~significantly. \citet{Santander15} have claimed to have discovered the first definite double-degenerate, super-Chandrasekhar system that will merge within a Hubble time, namely the central stars of the planetary nebula (CSPN) \hen. They found a photometric period of 4.2\,h and that \Ionw{He}{2}{5411} is double lined and time variable. By fitting the radial velocities (RVs) and light curves, they~concluded that the system consists of two pre-white dwarfs with equal masses of 0.88\,\Msol.~In this case, the~system would merge within 700 million years making, \hen one of the best SN Ia progenitor candidates~known. This scenario has since been challenged by \citet{Garcia-Berro16}, who criticized the strong mismatch between the luminosities and radii of both pre-white dwarf components as derived by \cite{Santander15} with the predictions from stellar evolution models \cite{Bloecker91}. In addition, Reference \cite{Garcia-Berro16} suggested that the variable \Ionw{He}{2}{5411} line might instead be a superposition of an absorption line plus an emission line, possibly arising from the nebula, the irradiated photosphere of a close companion, or a stellar wind. Since this would question the dynamical masses derived by \cite{Santander15}, Reference \cite{Garcia-Berro16} that repeated the light curve fitting and showed that the light curves of \hen may also be fitted well by assuming an over-contact binary system that consists of two lower mass (i.e., masses of 0.47\,\Msol\ and 0.48\,\Msol) stars. Thus, Reference \cite{Garcia-Berro16} concludes that the claim that \hen provides observational evidence for the double degenerate scenario for SN Ia is~premature. Given the potential importance of \hen as a unique laboratory to study the double degenerate merger scenario, it is highly desirable to resolve this debate. Therefore, we use an improved approach for the analysis of this unique object by combining quantitative non-LTE spectral modelling, RV~analysis, multi-band light curve fitting, and state-of-the art stellar evolutionary calculations. | Our preliminary results suggest that the dynamical system mass of Hen\,2-428 that is derived by using all available \Ion{He}{2} lines does not exceed the Chandrasekhar mass limit. Furthermore, the individual masses of the two central stars are also too small to lead to an SN\,Ia in case of a dynamical explosion during the merger process. Further investigations on the red-excess and atmospheric models that consider opacities of heavy metals are mandatory for a reliable analysis this intriguing system. \vspace{6pt} \authorcontributions{Writing---Original Draft Preparation, N.R.; Writing---Review and Editing, N.R., N.L.F, V.S., M.A.B., S.L.C., S.G., M.M.M.B., and S.T.; Visualization, N.R. and V.S.; Supervision, N.R., S.L.C., and M.A.B.; Data reduction: S.T.; Spectral Analysis: N.L.F. and N.R.; Light curve modeling: V.S.; RV analysis: N.L.F. and S.G.; mass determination: N.R.; Interpretation of the results: N.R., S.G., V.S., M.M.M.B, S.T.} \funding{N.R. is supported by a RC1851 research fellowship. N.L.F. is supported by an Science and Technology Facilities Council Studentship.} | 18 | 8 | 1808.04624 |
1808 | 1808.06621_arXiv.txt | Despite considerations of mass loss from stellar evolution suggesting otherwise, the content of gas in globular clusters seems poor and hence its measurement very elusive. One way of constraining the presence of ionized gas in a globular cluster is through its dispersive effects on the radiation of the millisecond pulsars included in the cluster. This effect led Freire et al. in 2001 to the first detection of any kind of gas in a globular cluster in the case of 47 Tucanae. By exploiting the results of 12 additional years of timing, as well as the observation of new millisecond pulsars in 47 Tucanae, we revisited this measurement: we first used the entire set of available timing parameters in order to measure the dynamical properties of the cluster and the three-dimensional position of the pulsars. Then we applied and tested various gas distribution models: assuming a constant gas density, we confirmed the detection of ionized gas with a number density of $n= 0.23\pm 0.05$ cm$^{-3}$, larger than the previous determination (at $2\sigma$ uncertainty). Models predicting a decreasing density or following the stellar distribution density are highly disfavoured. We are also able to investigate the presence of an intermediate mass black hole in the centre of the cluster, showing that is not required by the available data, with an upper limit for the mass at $\sim 4000$ M$_{\odot}$. | Globular clusters (GCs) are known to harbour a very large population of pulsars. Currently 150 pulsars have been found in 28 globular clusters\footnote{For an up-to-date number visit: \href{http://www.naic.edu/\textasciitilde pfreire/GCpsr.html}{http://www.naic.edu/\textasciitilde pfreire/GCpsr.html}}. Almost all of them are millisecond pulsars (MSPs) which have been recycled during accretion from a binary companion. In globular clusters, the number of low-mass X-ray binaries (and their products, MSPs) per unit stellar mass is much greater than in the Galactic field \citep{1975ApJ...199L.143C}. This results from the very high central densities of the clusters which increases the chance of close stellar encounters. These encounters can lead to the formation of new binaries containing a neutron star where accretion can occur and the neutron star can be recycled. The same processes that form these binaries can also destroy \citep{VerbuntFreire2014} them or the neutron star can ablate the companion with its strong wind. For these reasons many MSPs in globular clusters are isolated. Thanks to their abundance and to their rotational stability, MSPs can be used as unparalleled probes of the gravitational potential and environment of globular clusters. MSPs have been used to constrain the properties of the parent clusters \citep{Phinney1993, Anderson1993, Freire2003, Prager2017} and to study the presence of intermediate mass black holes (IMBHs) \citep{Kiziltan2017a, Kiziltan2017b, Freire2017, Perera2017, Prager2017}. 47 Tucanae (also known as NGC 104, hereafter 47 Tuc) is one of the most prominent globular clusters. The main properties of the cluster are listed in Table \ref{tab47tuc}. The central density has been measured from the proper motion central velocity dispersion $\sigma_{\mu,0}= 0.574 \pm 0.005$ mas yr$^{-1}$ \citep{Watkins2015} and the angular core radius $\theta_c=26.5\pm 1.16$ arcsec \citep{Bellini2017} using equation 1-34 in \cite{Spitzer1987}: \begin{equation} \label{vel_disp_eq} \rho_0=\frac{9\sigma_{\mu,0}^2}{4 \pi \rm{G} \theta_c^2} = (7.5 \pm 0.2) \times 10^{4}\, \rm{M}_{\odot}\, \rm{pc}^{-3}; \end{equation} which is accurate to $\sim 0.5$ per cent for clusters in which the tidal radius is much larger than the core radius \citep{Freire2017}. This value is significantly lower than that derived by \cite{Freire2017} because of the significant difference in $\theta_c$. We chose to use the value of $\theta_c=26.5\pm 1.16$ arcsec \citep{Bellini2017} as it is the most recent value obtained combining the surface brightness profile with the kinematic information both along the line of sight and in proper motion taken from state-of-the-art Hubble Space Telescope observations. Radio observations of 47 Tuc led to the discovery of 25 pulsars \citep{Manchester1990, Robinson1995, Camilo2000, Pan2016} and to the phase-coherent timing solution of 23 of them \citep{Freire2001a, Freire2003, Ridolfi2016, Freire2017,Freire2018}. All of these pulsars are MSPs and have spin periods shorter than 8 ms. Fifteen of them are in binary systems. Except for 47 Tuc X, they all reside within 1 arcminute of the centre. \cite{Ridolfi2016} and \cite{Freire2017} also provided the values of the period derivative, the second period derivate and proper motion for 22 pulsars. For 10 of the binary pulsars it was also possible to measure the orbital period derivative. \begin{table} \caption{Main properties of the globular cluster 47 Tuc. $^1$\protect\cite{McLaughlin2006}, $^2$\protect\cite{Bogdanov2016}, $^3$\protect\cite{Harris2010}, $^4$\protect\cite{Bellini2017}, $^5$\protect\cite{Gratton1997}, $^6$\protect\cite{Gnedin2002}, $^7$\protect\cite{Watkins2015}, $^{8}$\protect\cite{Freire2001b} } \begin{center} \begin{tabular}[t]{c c c} \hline {Parameter} & {Value} & {References} \\ \hline \hline Centre RA (J2000)& $00^{\rm h} 24^{\rm m} 05^{\rm s} .67 \pm 0^{\rm s} .07$ & $^1$ \\ Centre Dec (J2000) & $-72^{\circ} 04'52".62 \pm 0".26 $ & $^1$ \\ Distance from Sun & $4.53 \pm 0.04$ kpc & $^2$ \\ Metallicity & $-0.72$ dex & $^3$ \\ Mass & $(8.4 \pm 0.4) \times 10^5 \rm{M}_{\odot}$ & $^4$ \\ Tidal radius & 42'.3 (55.4 pc) & $^3$\\ Core radius & 26".5 (0.58 pc) & $^4$ \\ Age & $10.0 \pm 0.4 $ Gyr& $^5$ \\ Escape velocity at core & 68.8 km/s & $^6$ \\ Central velocity dispersion& $0.574\pm 0.005 $ mas yr$^{-1}$ & $^7$ \\ Central density& $75000 \pm 2000\, \rm{M}_{\odot}\, \rm{pc}^{-3}$ & $^{4,7}$ \\ Central ICM density & $0.067\pm 0.015 \, \rm{cm}^{-3}$ & $^{8}$ \\ \hline \end{tabular}\\ \label{tab47tuc} \end{center} \end{table} 47 Tuc was the first GC where evidence of the presence of ionized gas was found \citep{Freire2001b}. This discovery was made possible by the study of the dispersion measure (DM) differences between each pulsar. The DM causes a frequency dependent delay of the time of arrival of the pulses and is caused by the presence of free electrons along the line of sight. The DM was seen to be higher for pulsars farther along the line of sight compared to ones closer to the observer. This was interpreted as being due to the presence of an ionized component of constant density of the intracluster medium (ICM). Because of the large errors it was not possible to discriminate between various distribution models of the gas. Despite this detection, there is very little additional evidence of any kind of interstellar medium inside globular clusters. This is a long standing problem in the astrophysics of GCs \citep{Smith1990, VanLoon2006, Barmby2009}. The only certain detection of neutral gas in a globular cluster was made in M15: an HI cloud of $\sim 0.3\, \rm{M}_{\odot}$ and $9 \pm 2 \times 10^{-4}\, \rm{M}_{\odot}$ of dust \citep{Evans2003, Boyer2006, VanLoon2006}. This amount of gas and dust is very small if compared to what is expected to be emitted by the evolved stars of the cluster, i.e. $\sim 10^{-6}\, \rm{M}_{\odot} \rm{yr}^{-1}$ \citep{McDonald2011}. A fast clearing mechanism for the dust is necessary to explain the discrepancy between the observations and the predictions. This clearing mechanism could be caused by pulsar winds \citep{Spergel1991}, fast winds from red giants \citep{Smith2004}, classical novae \citep{Moore2011} or by white dwarfs \citep{McDonald2015a}. A more detailed modelling of the gas density could in principle be used as a tracer of the origin and evolution of the gas itself. Furthermore, it has been suggested that the distribution of gas could be influenced by the presence of an intermediate mass black hole \citep{Pepe2016}, thus allowing us to put additional constraints on its presence. Stringent upper limits have been put in the past on the mass of the central IMBH in 47 Tuc between 1000 - 5000 M$_{\odot}$ from both kinematic methods and radio continuum observations \citep{McLaughlin2006, Maccarone2008, Maccarone2010, Lu2011}. Recently a claim of an IMBH of 2200 M$_{\odot}$ was put forward \citep{Kiziltan2017a, Kiziltan2017b} using pulsar observations. However, using updated results, \cite{Freire2017} deemed the claim unnecessary; the same result was obtained by \cite{2018arXiv180703307M} using detailed measurements of the normal stars in the cluster; and a similar conclusion was derived for a larger number of globular clusters from radio continuum surveys \citep{2018ApJ...862...16T}. An independent method for testing for the presence of the IMBH might help to solve this question. The aim of this paper is to test various distribution models for the ionized gas inside the globular cluster 47 Tuc using the new timing results presented in \cite{Ridolfi2016}, \cite{Freire2017} and \cite{Freire2018}: they were obtained from a much longer data-span ($16$ yrs as compared to $4$ yrs) than that available at the time of the original detection \citep{Freire2001b}. The analysis is made using a Markov Chain Monte Carlo (MCMC) algorithm first used to determine the dynamical parameters of Terzan 5 \citep{Prager2017}. Since the core radius and the velocity dispersion of 47 Tuc are well constrained thanks to optical observations, this algorithm can be used to accurately measure the line-of-sight position of the pulsars and to test the presence of an IMBH using the equations described in Section \ref{section_theory}. The algorithm itself is described in Section \ref{section_mcmc}. With the three-dimensional positions of the pulsars and their measured values of DM we test the presence of ionized gas with different distributions in Section \ref{section_gas}. In Section \ref{section_discussion} and Section \ref{section_conclusion} we discuss the results and derive the conclusions. | In this paper we used the new timing results of the millisecond pulsars associated with the globular cluster 47 Tuc to perform a detailed modelling of the dynamics of the cluster. We measured the properties of the cluster, found an upper limit on the mass of a possible IMBH at the centre and the position along the line of sight of the pulsars. By using this information and the observed DMs of the pulsars, we tested the presence of ionized gas following different distributions. The model with the highest statistical likelihood has a constant density distribution in the region populated by the pulsars, with a density of $n_g= 0.23 \pm 0.05$ cm$^{-3}$. Other models invoking a gas density distribution that follows the stellar distribution or a radially decreasing distribution are disfavoured. The proposed explanation for how a region of constant gas density can be maintained in the centre of the cluster is that the thermal and the radiative pressure provide the necessary support against the gravitational collapse. However, more detailed modelling of the gas injection and of the energy input must be developed to test this model. Finally, we used the derived information about the density and temperature profiles for the gas in order to put upper limits on the mass of a putative IMBH at the centre of 47 Tuc. The presence of a massive central black hole in 47 Tuc will also be better constrained in the future when the effects of an IMBH on the jerks of the pulsars close enough to the latter are included. | 18 | 8 | 1808.06621 |
1808 | 1808.08911_arXiv.txt | % {The reionization of the Universe ends the dark ages that started after the recombination era. In the case of H, reionization finishes around $z\sim 6$. Faint star-forming galaxies are the best candidate sources of the H-ionizing radiation, although active galactic nuclei may have also contributed. We have explored whether the termination regions of the jets from active galactic nuclei may have contributed significantly to the ionization of H in the late reionization epoch, around $z\sim 6-7$. We assumed that, as it has been proposed, active galactic nuclei at $z\sim 6$ may have presented a high jet fraction, accretion rate, and duty cycle, and that non-thermal electrons contribute significantly to the pressure of jet termination regions. Empirical black-hole mass functions were adopted to characterize the population of active galactic nuclei. From all this, estimates were derived for the isotropic H-ionizing radiation produced in the jet termination regions, at $z\sim 6$, through inverse Compton scattering off CMB photons. We find that the termination regions of the jets of active galactic nuclei may have radiated most of their energy in the form of H-ionizing radiation at $z\sim 6$. For typical black-hole mass functions at that redshift, under the considered conditions (long-lasting, common, and very active galactic nuclei with jets), the contribution of these jets to maintain (and possibly enhance) the ionization of H may have been non-negligible. We conclude that the termination regions of jets from active galactic nuclei could have had a significant role in the reionization of the Universe at $z\gtrsim 6$.} | The reionization of the Universe put an end to the dark ages that followed the recombination era, after the big bang. The reionization epoch took place mostly around $z\sim 6-8$ \citep[e.g.][]{madau17}, although the nature of the dominant H-ionizing source(s) is still uncertain. Nowadays the favoured candidates are faint star-forming galaxies, in particular the fainter and less massive ones, so long as the escape fraction of the ionizing radiation is not too low \citep[e.g.][]{stark16}\footnote{Nevertheless, it has been proposed that brighter galaxies may have dominated if the reionization duration was relatively short, and their escape fraction high enough \citep[e.g.][]{sharma18}.}. Another possible candidate ionizing source at high redshift is accretion in active galactic nuclei (AGN) \citep[e.g.][]{arons70,donahue87,shapiro87,meiksin93}. Such sources are presently considered to be a minor player in the reionization of H \citep[e.g.][]{hopkins07,onoue17,parsa18,mcgreer18}, although they are still sometimes discussed as potentially important reionization sources \citep[e.g.][]{giallongo15,grazian18}, or as indirect factors in the reionization process \citep[e.g.][]{seiler18,trebitsch18,kakiichi18}. In addition, a so far undetected, faint AGN population, as for instance accreting intermediate mass black holes in the centre of gas-rich dwarf galaxies, cannot be discarded as significant reionization sources \citep{silk17}. Accretion radiation is not the only form of energy output in AGN. In particular, jets may have actually dominated the energy output of AGN at $z\sim 6$: (i) the jetted AGN fraction may have been close to one (at least for the most massive black-hole AGN - e.g. \citealt{sbarrato15}-, in principle valid as well for less massive black-hole sources); and (ii) the AGN central black holes may have been accreting close to the Eddington limit with a high duty cycle \citep[e.g.][]{willott10,shankar13}. Since jets may be as powerful as accretion radiation, or even more \citep[e.g.][]{ghisellini14,sbarrato16}, their role in the reionization of the Universe should be considered: if a substantial fraction of the jet energy went to H-ionizing photons, they may have contributed as much as AGN accretion, and perhaps even more. It is worth mentioning that very energetic cosmic rays may also be efficiently accelerated in AGN jets, and the impact of these cosmic rays on the ionization and heating of the intergalactic medium (IGM) at very high redshift deserves detailed studies \citep[see, in the context of POPIII microquasars, e.g.][]{tueros14,douna18,romero18}. However, cosmic ray production in AGN jets is uncertain. First, only the most energetic cosmic rays can diffuse out of the radio lobes; otherwise, they stay in the lobes and cool through adiabatic losses. Secondly, it is not known which fraction of the jet energy is in the form of very energetic cosmic rays. On the contrary, the termination region of jets are known to be filled by non-thermal electrons that could dominate pressure \citep[e.g.][]{croston18}, and these electrons could have efficiently emitted H-ionizing photons via inverse Compton (IC) scattering off CMB photons at high redshift. As bulk velocities in the lobes are relatively low, this emission would have been rather isotropic. Therefore, in this work we have focussed on the AGN jet termination regions, in particular the lobes inflated by shocked jet material, as sources of H-ionizing photons in the late phase of the reionization epoch. The beamed emission from the (relativistic) jet smaller scales is not considered here, although it could be included if it were very efficient turning jet energy into hard photons. This work is a first step in exploring the role of AGN jets and their termination regions in the late reionization epoch, say at $z\sim 6-7$; more accurate predictions are left for future studies. The paper is structured as follows: In Sect.~2, we introduce the prescriptions adopted for the black-hole population at high redshift, together with the assumptions adopted to derive the black-hole jet power. In Sect.~\ref{jt}, we show that the termination regions of the jets of AGN at $z\gtrsim 6$ could have been powerful sources of (isotropic) H-ionizing radiation, and estimate how much this radiation may have contributed to (maintain) the reionization of the Universe. Finally, we discuss our results and conclude with a summary in Sects.~\ref{dis} and \ref{sum}, respectively. | \label{dis} As mentioned in the previous section, to compare with $\dot\epsilon_{\rm obs}$, $\dot\epsilon_{\rm ion}$ was assumed to be constant for the whole duration of the reionization epoch, but this was likely not the case. For $z>6-7$, in addition to the expected decline of the empirical black-hole mass function, the masses of the growing black holes, and thus their $L_{\rm j}$, are expected to be lower. On the other hand, the actual number of (jetted) black holes may have been larger than those empirically inferred. For instance, it has been proposed that active IMBH may have been present in all gas-rich dwarf galaxies at $z\gtrsim 6$, producing outflows and possibly jets (e.g. \citealt{silk17,barai18}; see however \citealt{latif18}). In fact, \cite{salvador17} predicts $\sim 100$ times more IMBH ($M\sim 10^5$~M$_\odot$) at redshift $z\sim 6$ than the empirical mass functions \citep{willott10}. If for instance 10\% of those IMBH untraced by empirical studies had produced jets with $L_{\rm j}\sim 0.1\,L_{\rm Edd}$, their contribution to $\dot\epsilon_{\rm ion}$ may have been relevant, also at $z>6-7$. In addition to the AGN luminosity function, or the black-hole mass function, the accretion rate (comoving) density ($\dot\rho_{\rm acc}$) at $z\sim 6$ derived from X-ray background studies can set a limit on the ionizing role of AGN jet lobes assuming $\dot\epsilon_{\rm ion}<\dot\rho_{\rm acc}\,c^2$. \cite{vito18} estimates the accretion rate (comoving) density at $z\sim 6$ as $\dot\rho_{\rm acc}\sim 10^{-6}\,$M$_\odot$~yr$^{-1}$~Mpc$^{-3}$ \citep{vito18}, which yields $\dot\epsilon_{\rm ion}<6\times 10^{40}$~erg~s$^{-1}$~Mpc$^{-3}$, leaving room for $\dot\epsilon_{\rm ion}$ to significantly contribute to $\dot\epsilon_{\rm obs}$. The estimate for $\dot\rho_{\rm acc}$ is model dependent and possibly too conservative, but is not inconsistent with AGN jets being relevant H-ionizing sources at $z\gtrsim 6$. The black-hole mass (comoving) densities of the adopted mass functions are $\sim 10^2-10^3$~M$_\odot$, again consistent with the (model-dependent) limits found in the literature \citep[$\sim 10^3-10^4$~M$_\odot$; e.g.][]{hopkins07,salvaterra12,comastri15,cappelluti17}. It is worth checking, in the jetted AGN scenario studied here, whether $\dot\epsilon_{\rm ion}$ could represent a non-negligible contribution to the reionization without overcoming the cosmic radiation background (CRB). For the black-hole mass functions adopted here, one can provide an estimate of the broadband flux as: \begin{equation} F_\Omega\sim 0.1\dot\epsilon_{\rm ion}V_{\rm c}/16\pi^2 r_{\rm c}^2(1+z)^2= \end{equation} $$ =0.1\dot\epsilon_{\rm ion}r_{\rm c}/12\pi(1+z)^2\approx (1-4)\times 10^{-4}(\chi/0.1)\, {\rm nW}\,{\rm m}^{-2}\,{\rm srad}^{-1}, $$ where $V_{\rm c}=4\pi r_{\rm c}^3/3$ is the comoving volume of the Universe, $z=6$, and a constant source behaviour with $z$ has been assumed. The value of $F_\Omega$ is an order of magnitude estimate, but can be directly compared with the spectral energy distribution ($\lambda F^{\rm obs}_{\Omega\lambda}$) shown for example in Fig.~3 in \cite{cooray16}, which presents values ranging from $\sim 10^{-3}$ (radio, gamma rays) to $10^{-1}$~nW~m$^{-2}$~srad$^{-1}$ (X-rays). For instance, for an optimistic $\chi\sim 0.1$, the jetted AGN at $z\gtrsim 6$ may have had an important role ionizing the Universe without violating (but somewhat contributing to) the CRB. Future, more accurate studies should be devoted to further explore the possibility that AGN jets may have contributed to the reionization of the IGM at $z\sim 6$ and earlier. In particular, more detailed population prescriptions, including yet untraced IMBH populations, an empirically derived accretion-jet relation for high-$z$ AGN, and a more precise model for the jet lobe dynamics and emission (including the shocked IGM thermal component), would give better constraints on the importance of AGN jets ionizing (and heating) the IGM at $z\gtrsim 6$. It would also be valuable to explore the impact of jet lobe relativistic electrons on the CMB spectrum through IC, as arcminute-scale small distortions may be expected \citep[e.g.][]{yamada99,malu17}. | 18 | 8 | 1808.08911 |
1808 | 1808.08559.txt | An effective practical model with two characteristic parameters is presented to describe both of the tidally induced shape and spin alignments of the galactic halos with the large-scale tidal fields. We test this model against the numerical results obtained from the Small MultiDark Planck simulation on the galactic mass scale of $0.5\le M/(10^{11}\,h^{-1}\,M_{\odot})\le 50$ at redshift $z=0$. Determining empirically the parameters from the numerical data, we demonstrate how successfully our model describes simultaneously and consistently the amplitudes and behaviors of the probability density functions of three coordinates of the shape and spin vectors in the principal frame of the large scale tidal field. Dividing the samples of the galactic halos into multiple subsamples in four different mass ranges and four different types of the cosmic web, and also varying the smoothing scale of the tidal field from $5\,h^{-1}$Mpc to $10,\ 20,\ 30\,h^{-1}$Mpc, we perform repeatedly the numerical tests with each subsample at each scale. Our model is found to match well the numerical results for all of the cases of the mass range, smoothing scale and web type and to properly capture the scale and web dependence of the spin flip phenomenon. | \label{sec:intro} The physical properties of the observed galaxies in the universe is a reservoir of information on the conditions under which they formed, the evolutionary processes which they went through, and the interactions in which they are involved. Although the local conditions and processes at the galactic scales must have had the most dominant impact on the galaxies, the non-local effects beyond the galactic scales are also believed to have contributed partly to their physical properties \citep[e.g.,][]{PS17}. Subdominant as its contribution is, the non-local effects on the galaxies are worth investigating, since it may contain valuable independent information on the galaxy formation and the background cosmology as well. The non-local effects on the galaxies are manifested by the correlations between the galaxy properties and the large-scale environments. Among various properties of the galaxies that have been found correlated with the large-scale environments, the shape and spin alignments of the galaxies with the large-scale structures (collectively called the galaxy intrinsic alignments) have lately drawn considerable attentions, inspiring vigorous extensive studies \citep[see][for recent reviews]{review1,review2,review3}. It is partially because the galaxy intrinsic alignments, if present and significant, could become another systematics in the measurements of the extrinsic counterparts caused by the weak gravitational lensing \citep[see][and references therein]{TI15}. The other important motivation for the recent flurry of research on this topic is that the origin of the galaxy intrinsic alignments is amenable to the first order perturbation theory and thus a rather fundamental approach to this topic is feasible \citep[e.g.,][]{hea-etal00,LP00,cat-etal01,cri-etal01,LP01,por-etal02,HZ08,bla-etal11,bla-etal15,tug-etal18}. In the first order Lagrangian perturbation theory \citep{zel70,buc92}, the minor (major) eigenvectors of the inertia momentum tensors of the proto-galactic regions are perfectly aligned with the major (minor) eigenvectors of the local tidal tensors around the regions. Several $N$-body simulations have indeed detected the existence of strong correlations between the inertia momentum and local tidal tensors at the proto-galactic sites \citep{LP00,por-etal02,lee-etal09}. Since the tidal fields smoothed on different scales are cross correlated, the eigenvectors of the inertia momentum tensors of the proto-galactic regions are expected to be aligned with those of the large-scale tidal fields. The major eigenvectors of the inertia momentum tensors of the proto-galaxies correspond to the most elongated axes of their shapes, while the minor eigenvectors of the large-scale tidal tensors correspond to the directions along which the surrounding matter become minimally compressed. Henceforth, this expectation based on the first order Lagrangian perturbation theory basically translates into the possible alignments between the galaxy shapes and the most elongated axes of the large-scale structures such as the axes of the filaments, the signals of which have been detected by several numerical and observational studies \citep[e.g.,][and references therein]{alt-etal06,hah-etal07b,zha-etal09,zha-etal13,che-etal16}. In the linear tidal torque (LTT) theory that \citet{dor70} formulated by combining the first order Lagrangian perturbation theory with the Zel'dovich approximation \citep{zel70}, the anisotropic tidal field of the surrounding matter distribution originates the spin angular momentum of a proto-galaxy provided that its shape departs from a spherical symmetry. The generic and unique prediction of this LTT theory is the inclinations of the spin vectors of the proto-galaxies toward the intermediate eigenvectors of the large scale tidal field \citep{LP00}, which has also garnered several numerical and observational supports \citep[e.g.,][]{nav-etal04,tru-etal06,hah-etal07a,LE07,wan-etal11,zha-etal15,che-etal16}. The recently available large high-resolution $N$-body simulations that covered a broad mass range, however, limited the validity of the LTT prediction to the mass scale of $M\ge M_{\rm t}\sim 10^{12}\,h^{-1}\,M_{\odot}$, showing that on the mass scale below $M_{t}$ the spin vectors of dark matter halos at $z=0$ are aligned not with the intermediate but rather with the minor eigenvectors of the large scale tidal field, similar to the axes of the halo shapes \citep{ara-etal07,hah-etal07b,paz-etal08,zha-etal09,cod-etal12,lib-etal13,tro-etal13,dub-etal14,vee-etal18}. This difference in the spin alignment tendency between the low and high mass scales were found most conspicuous in the filament environments: the spin axes of the galactic halos with masses lower (higher) than $M_{\rm t}$ measured at $z=0$ tend to be parallel (perpendicular) to the elongated axes of their host filaments, in contradiction with the LTT prediction. The transition of the spin alignment tendency at $M_{\rm t}$ is often called "spin flip" phenomenon \citep{cod-etal12} and the break-down of the LTT prediction below $M_{\rm t}$ has also been witnessed in recent observations \citep{tem-etal13,TL13,hir-etal17,che-etal18}. The detection of this spin-flip phenomenon puzzled the community and urged it to find a proper answer to the critical question of what the origin of this phenomenon is. What has so far been suggested as a possible origin includes the major merging events, mass dependence of the merging and accretion processes, assembly bias, vorticity generation inside filaments, web-dependence of the galaxy formation epochs, nonlinear tidal interactions, geometrical properties of the host filaments and etc \citep{BF12,LP12,cod-etal12,lib-etal13,wel-etal14,cod-etal15,lai-etal15,BF16,WK17,vee-etal18}. Although these previously suggested factors were believed to play some roles for the occurrence of the spin-flip phenomenon, none of them are fully satisfactory in explaining all aspects of the spin-flip phenomenon including the dependence of the transition mass scale $M_{\rm t}$ on the types of the cosmic web, redshifts, and scales of the filaments. The occurrence of the spin-flip phenomenon basically implies that for the case of the galaxies with masses $M\le M_{t}$, the tendency of the spin alignments with the large scale tidal field becomes similar to that of the shape alignments. Thus, it is suspected that whatever caused the spin-flip phenomenon, it should be linked to the shape alignments with the large scale tidal field. To address these remaining issues, what is highly desired is an effective model that can describe consistently and simultaneously both of the galaxy shape and spin alignments. Here, we attempt to construct such a model by modifying the original LTT theory and to explore if the shapes of the galaxies also show any transition of the alignment tendency like the spin counterparts The organization of this Paper is as follows. A refined analytic model for the galaxy shape alignments is presented in Section \ref{sec:shape_model} and tested against the numerical results in Section \ref{sec:shape_test}. An effective model for the tidally induced spin alignments is presented in Section \ref{sec:spin_model} and tested against the numerical results in Section \ref{sec:spin_test}. A discussion over the possible application of this model as well as a summary of the results is presented in Section \ref{sec:con}. Throughout this Paper, we will assume a Planck universe whose total energy density is dominantly contributed by the cosmological constant ($\Lambda$) and the cold dark matter (CDM) \citep{planck13}. | \label{sec:con} To study the large-scale tidal effect on the spin and shape orientations of the galaxies and the spin-flip phenomenon, we have considered three different analytic models, the \texttt{model I}, \texttt{model II} and \texttt{model III}. The \texttt{model I}, Equation (\ref{eqn:model1}), which was originally developed by \citet{LP00} based on the LTT theory, describes the alignment tendency between the galaxy spin vectors, $\hat{\bf s}$, and the intermediate eigenvectors, $\hat{\bf u}_{2}$, of the large-scale tidal field, ${\bf T}$. The \texttt{model II}, Equation (\ref{eqn:hpro_axis_dt}), has been constructed here to describe the alignments (anti-alignments) of the galaxy shapes, $\hat{\bf e}$, with the minor (major) eigenvectors, $\hat{\bf u}_{3}$ ($\hat{\bf u}_{1}$) of ${\bf T}$. This model is based on the first order Lagrangian perturbation theory according to which the major principal axes of the inertia momentum tensors of the galactic halos are perfectly aligned with the minor principal axes of the local tidal tensors in the Lagrangian regime. The \texttt{model III}, Equation (\ref{eqn:model3}), is a practical formula constructed by combining the \texttt{model I} and \texttt{model II} to describe simultaneously the tidally induced shape and spin alignments. The \texttt{model I} (\texttt{model II}) carries a single parameter, $c_{t}$ ($d_{t}$), which measures the strength of the alignment with $\hat{\bf u}_{2}$ ($\hat{\bf u}_{3}$). The \texttt{ model III} carries two parameters, $c_{t}$ and $d_{t}$, whose relative ratio determines the transition mass scale for the occurrence of the spin-flip. The first parameter, $c_{t}$, would reach the maximum value of unity, if the inertia momentum tensors of the galaxies are uncorrelated with the surrounding tidal tensors, while the second parameter, $d_{t}$, will attain the value of unity if the two tensors are perfectly correlated. These parameters can be empirically determined by Equation (\ref{eqn:cdt_sol}) directly from the measured values of $\hat{\bf e}$ and $\hat{\bf s}$ in the principal frame of $\hat{\bf T}$ without resorting to any fitting procedure. To numerically test the three analytic models, we have utilized the density fields and the Rockstar halo catalogs extracted from the SMDPL simulations \citep{smdpl}. Constructing the unit traceless tidal tensor, $\hat{\bf T}$, smoothed on the scale of $R_{f}=5\,h^{-1}$Mpc from the density fields given on the $512^{3}$ grids that constitute the simulation box of volume $400^{3}\,h^{-3}\,{\rm Mpc}^{3}$ and selecting the galactic halos in the mass range of $0.5\le M/(10^{11}\,h^{-1}\,M_{\odot})\le 50$ from the Rockstar catalog, we have first numerically obtained the probability density functions of the tidally induced shape alignments, $\{p(\vert\hat{\bf u}_{i}\cdot\hat{\bf e}\vert)\}_{i=1}^{3}$ (see Figures \ref{fig:pro_axis}-\ref{fig:pro_axis_filter}). The numerical results have clearly shown that $\hat{\bf e}$ has a tendency to be strongly aligned (anti-aligned) with $\hat{\bf u}_{3}$ ($\hat{\bf u}_{1}$) but no correlation with $\hat{\bf u}_{2}$. Investigating the dependence of the strength of the tidally induced shape alignments on $M$, $R_{f}$, and the type of the cosmic web, it has been found that the more massive galactic halos yield stronger $\hat{\bf u}_{3}$-$\hat{\bf e}$ alignments ($\hat{\bf u}_{1}$-$\hat{\bf e}$ anti-alignments) and that the increment of $R_{f}$ weakens the alignment tendency (see Figures \ref{fig:dt_f}). These numerical results are consistent with what the previous works already found \citep{joa-etal13,zha-etal13,che-etal16,hil-etal17,xia-etal17,pir-etal18}. The strongest (weakest) $\hat{\bf u}_{3}$-$\hat{\bf e}$ alignments are found from the void (knot) galactic halos (see Figures \ref{fig:pro_axis_knot}-\ref{fig:pro_axis_void}), which seem inconsistent with the previous numerical result that the DM halos showed the strongest shape alignments in the knot environments \citep{xia-etal17}. This inconsistency has been ascribed to the different classification schemes used in the two analyses. The sheet galactic halos yield much stronger shape alignment tendency than the knot and filament galactic halos in the whole mass range, which result is consistent with what \citet{hah-etal07a} found. In the lowest and low mass range ($0.5\le M/[10^{11}\,h^{-1}\,M_{\odot}]<5)$, the knot and filament galactic halos show similar strengths of the shape alignments. In the medium-mass ($5\le M/[10^{11}\,h^{-1}\,M_{\odot}]<10$) and high-mass ($10\le M/[10^{11}\,h^{-1}\,M_{\odot}]<50$) ranges, the shape alignments of the filament galactic halos become stronger than the knot counterparts (Figure \ref{fig:dt_web}). These numerical results imply that the void and sheet galactic halos retain best the tidally induced shape alignments, while the evolution of the galactic halos in the dense environments like the knots and filaments has an effect of deviating the directions of their shapes from the tidally induced inclinations. The comparison with the numerical results revealed the success of the \texttt{model II} in describing the amplitudes and behaviors of $\{p(\vert\hat{\bf u}_{i}\cdot\hat{\bf e}\vert)\}_{i=1}^{3}$, for all of the cases of $M$, $R_{f}$ and the type of the cosmic web. For the shape alignments, the \texttt{model III} turns out to be identical to the \texttt{model II}. For all of the four cases of the web type, the increment of $R_{f}$ has been found to decrease the strength of the tidally induced shape alignments but improve the agreements between the \texttt{model III} and the numerical results (Figure \ref{fig:pro_axis_filter_knot}-\ref{fig:dt_webf}). We interpret this result as an evidence supporting the scenario that the nonlinear evolution has an effect of diminishing the strength of the tidally induced shape alignments. In a similar manner, we have numerically determined the probability density functions of the tidally induced spin alignments, $\{p(\vert\hat{\bf u}_{i}\cdot\hat{\bf s}\vert)\}_{i=1}^{3}$, explored their dependences on $M$, $R_{f}$ and the web type, and compared the results with the three analytic models. The tidally induced spin alignments have been found significant but quite weak compared with the shape alignments (Figures \ref{fig:pro_spin}-\ref{fig:cdt_spin_m}), consistent with the results from the previous works \citep[e.g.,][]{hah-etal07b,for-etal14,zha-etal15}. The occurrence of the spin-flip phenomenon has been witnessed. For the case of $R_{f}=5\,h^{-1}$Mpc, the lowest-mass, low-mass and medium-mass galactic halos show strong $\hat{\bf u}_{3}$-$\hat{\bf s}$ alignments and negligible $\hat{\bf u}_{2}$-$\hat{\bf s}$ alignments, while the high-mass galactic halos exhibit strong $\hat{\bf u}_{2}$-$\hat{\bf s}$ alignments, which results have confirmed the claims of the previous works \citep{ara-etal07,paz-etal08,zha-etal09,cod-etal12,lib-etal13,tro-etal13,dub-etal14,che-etal16,vee-etal18}. However, we have noted that the spin-flip does not occur abruptly at a certain fixed transition mass scale. Rather it is a gradual transition of the spin alignment tendency that proceeds over a broader mass range, depending on $R_{f}$ (Figures \ref{fig:pro_spin_filter}-\ref{fig:cdt_spin_f}). For the case of $R_{f}=5\,h^{-1}$Mpc, the high-mass galactic halos have been found to yield stronger $\hat{\bf u}_{2}$-$\hat{\bf s}$ and weaker but significant $\hat{\bf u}_{3}$-$\hat{\bf s}$ alignments, while the medium-mass galactic halos exhibit strong $\hat{\bf u}_{3}$-$\hat{\bf s}$ and much weaker $\hat{\bf u}_{2}$-$\hat{\bf s}$ alignments. For the case of $R_{f}\ge 10\,h^{-1}$Mpc, however, the high-mass galactic halos exhibit stronger $\hat{\bf u}_{3}$-$\hat{\bf s}$ and weaker but significant $\hat{\bf u}_{2}$-$\hat{\bf s}$ alignments. The strengths of the tidally induced spin alignments have been also found to sensitively vary with the types of the cosmic web (see Figures \ref{fig:pro_spin_knot}-\ref{fig:pro_spin_filter_void}), which supports the claim of \citet{lib-etal13}. The strongest (weakest) signals of the tidally induced spin alignments have been found from the sheet (void) galactic halos, while the filament galactic halos have been found to have stronger spin alignments than the knot counterparts in the whole mass range ( Figures \ref{fig:ct_spin_web}-\ref{fig:dt_spin_web}). These results are inconsistent with the observational finding of \citet{zha-etal15} that the knot galaxies exhibited the strongest signals of the spin alignments. We have suspected that this inconsistency might be related to the construction of the tidal field from the galaxy groups and the determination of the spin axes of the galaxies from their stellar components in the observational analysis. Determining empirically $\langle c_{t}\rangle$ and $\langle d_{t}\rangle$ from the numerical data (Figures \ref{fig:ct_spin_web}-\ref{fig:dt_spin_web}) and defining the condition for the occurrence of the spin flip as $\langle c_{t}\rangle > \langle d_{t}\rangle$, we have quantitatively investigated how the occurrence and the transition mass scale, $M_{t}$, of the spin-flip phenomenon depend on the size and type of the cosmic web and found the following: \begin{enumerate} \item Regardless of the web type, the transition mass scale, $M_{t}$, of the spin-flip increases with the increment of $R_{f}$. \item The knot galactic halos do no show any spin-flip phenomenon. That is, the unit spin vectors, $\hat{\bf s}$, of the knot galactic halos are always preferentially aligned with $\hat{\bf u}_{3}$ rather than with $\hat{\bf u}_{2}$ in the whole mass range, regardless of the value of $R_{f}$ (Figure \ref{fig:pro_spin_filter_knot}). \item For the case of the filament galactic halos, the spin flip occurs around $M_{\rm t}\sim 5\times 10^{12}\,h^{-1}\,M_{\odot}$ when $R_{f}=5\,h^{-1}$Mpc. At the larger scale of $R_{f}>5\,h^{-1}$Mpc, the value of $M_{\rm t}$ exceeds the galactic mass scales, i.e, $M_{\rm t}> 5\times 10^{12}\,h^{-1}\,M_{\odot}$ (Figure \ref{fig:pro_spin_filter_fil}). \item In the sheet environment, the transition mass scale has a lower value than in the filaments: $M_{\rm t}\sim 10^{12}\,h^{-1}\,M_{\odot}$ when $R_{f}=5\,h^{-1}$Mpc. Only when $R_{f}$ reaches $30\,h^{-1}$Mpc, the value of $M_{\rm t}$ becomes larger than the galactic mass scale (Figure \ref{fig:pro_spin_filter_sheet}). \item The void galactic halos yield the lowest transition mass scale, $M_{\rm t}\sim 5\times 10^{11}\,h^{-1}\,M_{\odot}$ when $R_{f}=5\,h^{-1}$Mpc. At the larger scales, the number of the void galactic halos is too low to produce any significant signals (Figure \ref{fig:pro_spin_filter_void}). \end{enumerate} It is interesting to note that our results on the web and mass dependence of the spin-flip phenomenon are consistent with the theoretical explanation of \citet{cod-etal15}, according to which the misalignments between the inertia momentum and tidal tensors in the anisotropic environments like the filaments and sheets are largely responsible for the occurrence of the spin flip. In line with their theoretical explanation, we interpret no occurrence of the spin flip in the knot environments as an evidence for the stronger alignments between the two tensors in the dense environments. In other words, in the knot regions where the tidal tensors are more isotropic, the inertia momentum and tidal tensors may be more strongly aligned with each other, which plays a role in suppressing the occurrence of the spin-flip of the knot galaxies. It has also been clearly demonstrated in the current work that the \texttt{model III} succeeds in describing consistently and simultaneously the numerical results of the tidally induced shape and spin alignments for all of the cases of $M$, $R_{f}$ and type of the cosmic web, while the \texttt{model I} and \texttt{model II} fail. Showing that the \texttt{model III} works better as $R_{f}$ increases, we have ascribed the slight mismatches between the numerical results and the \texttt{model III} to the inaccuracies caused by the approximations of $p({\bf s}\vert{\bf T})$ as a multivariate Gaussian distribution and $\hat{\bf T}$ as a Gaussian random field made in the construction of the \texttt{model III}. We also suspect that the uncertainties in the measurements of $\hat{\bf s}$ and $\hat{\bf e}$ caused by the simple assumptions of each galactic halo having a perfect ellipsoidal shape and no substructure in a completely relaxed dynamical state must contribute to the mismatches. We conclude that the \texttt{model III} is an effective practical model for the spin and the shape alignments of the galactic halos with the large-scale tidal fields, providing an analytic tool with which the condition of the spin flip occurrence as well as its dependence on the properties of the large-scale structures can be quantitatively described. Its good accord with the numerical results supports the scenario that the occurrence of the spin flip phenomenon is associated more with the geometrical properties of the large-scale tidal field as well as the interactions of the galactic halos with the cosmic web rather than with the physical processes during the nonlinear evolution \citep[see][]{lib-etal13,cod-etal15,WK17,vee-etal18}. Given that the \texttt{model III} is expressed in terms of the linear quantities, it may provide another independent probe of the background cosmology. For this purpose, however, a couple of back-up works will have to be done. First, as suspected in our analysis, differences in the schemes used to to construct the tidal fields, to measure the shape and spin axes of the galaxies, and to classify the cosmic web would yield different patterns in the dependence of the tidally induced shape and spin alignments on the sizes and types of the cosmic web. Thus, it will be necessary to test the robustness of the \texttt{model III} against the variations of the schemes. Second, it will be also essential to examine its validity using the numerical results for alternative cosmologies such as models with modified gravity, coupled dark energy, massive neutrinos, primordial non-Gaussianity, anisotropic inflation and so forth. Our future work is in this direction. | 18 | 8 | 1808.08559 |
1808 | 1808.02002_arXiv.txt | { Conventional approaches to cosmology inference from galaxy redshift surveys are based on $n$-point functions, which are under rigorous perturbative control on sufficiently large scales. Here, we present an alternative approach, which employs a likelihood at the level of the galaxy density field. By integrating out small-scale modes based on effective-field theory arguments, we prove that this likelihood is under perturbative control if certain specific conditions are met. We further show that the information captured by this likelihood is equivalent to the combination of the next-to-leading order galaxy power spectrum, leading-order bispectrum, and BAO reconstruction. Combined with MCMC sampling and MAP optimization techniques, our results allow for fully Bayesian cosmology inference from large-scale structure that is under perturbative control. We illustrate this via a first demonstration of unbiased cosmology inference from nonlinear large-scale structure using this likelihood. In particular, we show unbiased estimates of the power spectrum normalization $\sigma_8$ from a catalog of simulated dark matter halos, where nonlinear information is crucial in breaking the $b_1-\sigma_8$ degeneracy. } \begin{document} | \label{sec:intro} One of the prime difficulties in inferring cosmology from the observed large-scale structure is the nonlinear and nonlocal connection between the matter density and tidal field, whose evolution is well understood, and the density of observed tracers such as galaxies. One possible approach, which offers the advantages of theoretical robustness as well as controlled systematic uncertainties is the \emph{effective field theory (EFT) approach}, or equivalently, \emph{general bias expansion} (see \cite{baumann/etal:2012,carrasco/etal:2012} for details on the EFT approach in large scale structure, and \cite{biasreview} for a review of bias). In this approach, the fully nonlinear small-scale modes are integrated out, leading to a relation between the observed galaxy density and operators constructed out of the matter density field multiplied by free coefficients, the bias parameters. Predictions are obtained by truncating the expansion at fixed order in perturbations and spatial derivatives. This approach is inherently limited to scales larger than the scale where the matter density field becomes nonlinear, at a wavenumber $\knl \sim 0.25 \iMpch$ for a standard $\Lambda$CDM cosmology at today's epoch, and larger than the characteristic spatial scale $R_*$ involved in the formation of the tracer considered. Standard techniques for applying this approach to observational or simulated data sets make use of summary statistics, in particular $n$-point correlation functions in real or Fourier space on large scales (large $r$, small $k$), where $n \geq 2$ (see Sec.~4.1 of \cite{biasreview} for an overview). However, in order to extract information from galaxy statistics beyond the linear regime in this approach, higher-order $n$-point functions, such as the three- and four-point functions, are crucial. By combining 2- and 3-point functions, for example, the degeneracy between linear bias and the amplitude of matter fluctuations can be broken, allowing for cosmological constraints. Unfortunately, both estimators and the associated covariances for higher-order $n$-point functions become increasingly difficult to handle with growing $n$, both technically and computationally, due to the rapidly increasing size of the data vectors. In particular, survey-specific systematic effects such as the mask, varying survey depth, or fiber-collision effects, need to be incorporated into the model for each $n$-point function. A possible alternative is to attempt to forward-model the galaxy density field itself, without resorting to summary statistics. This approach offers the advantage of a much more straightforward incorporation of systematic effects. Starting from early attempts based on galaxy peculiar velocities \cite{ 1999ApJ...522....1D, 1992ApJ...391..443N}, this approach is being pursued increasingly actively \cite{ 2008MNRAS.389..497K, 2010MNRAS.407...29J, 2010MNRAS.409..355J, 2013MNRAS.432..894J, 2013ApJ...772...63W, 2013ApJ...779...15J, 2014ApJ...794...94W, 2015MNRAS.446.4250A, 2017JCAP...12..009S, modi/etal, 2018arXiv180611117J}. The forward model for matter, together with the perturbative bias expansion, provides us with a ``mean tracer field'' in a certain sense. But which likelihood should be used to compare this mean field with the observed galaxy density field? Answering this question, on which little theoretical progress has been made so far, in the context of the EFT is the goal of this paper. The studies cited above assumed localized likelihoods in real space either of a specific form, or modeled using neural networks. As we will see, the EFT approach arrives at a somewhat different result. After reviewing the program of Bayesian LSS inference, we derive the likelihood in the EFT approach by integrating out small-scale modes. As we will see, this yields an approximately Gaussian likelihood in Fourier space. We restrict to the rest-frame density of a tracer, i.e. neglect redshift-space distortions. Hence, even though our results in principle apply to any LSS tracer, we will refer to the tracer as ``halos'' throughout the paper, as the most straightforward application consists of halos identified in N-body simulations. We further derive and investigate the maximum-likelihood point when phases are fixed, and show that it corresponds to matching the halo-matter cross-power spectrum at next-to-leading (1-loop) order, and the halo-matter-matter bispectrum at leading order (tree level). We derive the precise conditions that need to be met to ensure an unbiased result. Finally, we show that this approach naturally includes fully nonlinear reconstruction of the baryon acoustic oscillation (BAO) feature. BAO reconstruction refers to the fact that the displacements of galaxies from their initial (Lagrangian) positions lead to a damping of the oscillatory BAO feature. These displacements are dominated by large-scale modes which are inferred jointly with the cosmological parameters when following a forward-modeling approach. This offers another clear advantage over approaches based on $n$-point functions. The outline of the paper is as follows. \refsec{forward} reviews the Bayesian posterior for large-scale structure, focusing on large scales, which involves the prior on the initial conditions, forward model for matter, and deterministic bias expansion. \refsec{GaussF} then derives the remaining missing ingredient, the conditional probability of the observed halo field given the final matter density field and bias parameters. The following sections examine the ramifications of this likelihood: in \refsec{PT}, we study the maximum-likelihood point of this likelihood, and derive its relation to $n$-point correlation functions in the EFT. \refsec{nongauss} shows how non-Gaussian corrections to the likelihood are suppressed. Finally, \refsec{reco} describes how the likelihood presented here incorporates BAO reconstruction. We then turn to a preview of numerical results based on this likelihood in \refsec{results}, which will be described in an upcoming publication \cite{paperII}, and conclude in \refsec{concl}. The appendices contain colletions of important relations as well as details on the calculations presented in the main text. \subsection*{Notation} Our notation largely follows that of \cite{biasreview}. In particular, our Fourier convention and short-hand notation is \ba f(\vk) \equiv\:& \int d^3 \vx\, f(\vx) e^{-i\vk\cdot\vx} \equiv \int_{\vx} f(\vx) e^{-i\vk\cdot\vx} \vs f(\vx) \equiv\:& \int \frac{d^3 \vk}{(2\pi)^3}\, f(\vk) e^{i\vk\cdot\vx} \equiv \int_{\vk} f(\vk) e^{i\vk\cdot\vx}\,. \ea Wavenumbers that are integrated over (loop momenta) will further be denoted as $\vp, \vp',\vp_i\cdots$. Primes on Fourier-space correlators indicate that the momentum conserving Dirac delta $(2\pi)^3 \d_D(\vk_1+\vk_2+\cdots)$ is to be dropped. We will also frequently use the nonlocal derivative operator \be \Del_{ij} \equiv \left(\partial_i\partial_j \nabla^{-2} - \frac13 \d_{ij} \right) \label{eq:Deldef} \ee which is defined via its action on fields in the Fourier representation. We will typically deal with the filtered density field, employing an unspecified smoothing kernel $W(\vx)$ with Fourier-space counterpart $W(\vk)$. We will further use $W_{\L}$ for a sharp-$k$ filter: \be W_{\L}(\vk) = \thH(\Lambda-|\vk|)\,. \ee The matter and rest-frame halo (or galaxy) density perturbations are given by \be \d(\vx,\tau) \equiv \frac{\rho(\vx,\tau) - \rhob(\tau)}{\rhob(\tau)} \quad\mbox{and}\quad \d_h(\vx,\tau) \equiv \frac{n_h(\vx,\tau) - \avng(\tau)}{\avng(\tau)}\,. \ee Correspondingly, we denote the filtered matter and halo fields as $\d_W$ and $\d_{h,W}$, respectively. We reserve the notation $\d_\L,\ \d_{h,\L}$ for fields filtered with a sharp-$k$ filter on the scale $\L$. Since the matter density is related to the potential $\Phi$ through the Poisson equation \be \lapl\Phi = \frac32 \Om \cH^2 \d\,, \ee this allows us to combine the matter density perturbation and tidal field $K_{ij}$ into a tensor $\Pi^{[1]}$: \be \Pi^{[1]}_{ij}(\vx,\tau) \equiv \frac{2}{3\Om\cH^2} \partial_{x,i}\partial_{x,j}\Phi(\vx,\tau) = K_{ij}(\vx,\tau) + \frac13 \delta_{ij}\d(\vx,\tau)\,, \label{eq:Pi1} \ee which contains $\d = \tr \Pi^{[1]}$ and $K_{ij}$ as the trace-free part of $\Pi^{[1]}_{ij}$. All of these quantities are the evolved, nonlinear quantities. We further use the following notation for perturbative order: \begin{itemize} \item[$O^{(n)}\,:$] Operator evaluated at $n$-th order in perturbation theory. \item[$O^{[n]}\,:$] Operator whose \emph{lowest-order} contribution is at $n$-th order in perturbation theory. \end{itemize} All fields are implicitly assumed to be filtered on the grid scale used in the forward model. We will denote the filtering explicitly only when it is contrasted with unfiltered fields (specifically in \refsec{PT}). Finally, we let $\vec{\d}$ stand for a field defined on a grid: $\vec{\d}=\{ \d(\vx_i) \}_{i=1}^{N_g^3}$, where $N_g$ is the number of grid cells on one side, and throughout latin indices $i,j,k,...$ label grid cells, while greek indices $\alpha, \beta,...$ label parameters. The set of cosmological parameters will usually be denoted as $\theta$, while we reserve $\lambda$ for ``nuisance'' parameters (e.g. moments of the likelihood). We will also frequently use the corresponding discrete fields defined in the Fourier domain. For this, we adopt the standard box normalization: \ba \d(\vk) &= \sum_i^{N_g^3} \d(\vx_i) e^{-i \vk\cdot \vx_i} \label{eq:boxnorm}\\ \d(\vx) &= \frac1{N_g^3}\sum_{\vk}^{k_{\rm Ny}} \d(\vk_i) e^{i\vk_i\cdot \vx} \quad\mbox{where}\quad \vk \in (n_x, n_y, n_z) k_F\,,\ k_F = \frac{2\pi}{\Lbox}\,, \nonumber \ea and $n_i \in \{-N_g/2, \cdots N_g/2\}$ while $k_{\rm Ny} \equiv N_g k_F/2$. We will often use \be \fsum{\vk} \equiv \sum_{\{n_x,n_y,n_z\} \neq \{0,0,0\}}^{n_x^2+n_y^2+n_z^2 \leq (k_{\rm max}/k_F)^2}\,. \ee Correlators of box-normalized fields obey \ba \< X(\v{n} k_F) Y(\v{n}' k_F) \> &= \frac1{\Lbox^3} \d_{\v{n},-\v{n}'} P_{XY}(\v{n} k_F) \vs \< X(\v{n} k_F) Y(\v{n}' k_F) Z(\v{n}'' k_F) \> &= \frac1{\Lbox^6} \d_{\v{n}+\v{n}',-\v{n}''} B_{XYZ}(\v{n} k_F, \v{n}' k_F) \,, \label{eq:corrbox} \ea where \be \d_{\v{n},\v{n}'} \equiv \d_{n_x n_x'}\d_{n_y n_y'}\d_{n_z n_z'}\,, \ee and $P_{XY}$ ($B_{XYZ}$) are the cross-power spectrum (bispectrum) respectively. We have neglected the averaging over the finite $k$-space bin of width $k_{\rm Ny}$ on the r.h.s. of \refeq{corrbox}. | \label{sec:concl} We have derived an EFT-based likelihood for the galaxy density field that allows for cosmological inference from galaxy clustering with rigorously controlled theoretical systematics, without the need for measuring arbitrarily higher order $n$-point functions. In our concrete application, we have restricted to a second-order bias expansion, including the leading higher-derivative bias contribution as well as scale-dependent stochasticity. While we have not included projection effects such as redshift-space distortions, and thus referred to the tracers as halos, the bias expansion is fully general and also holds for galaxies. At this order, and when combined with a 2LPT forward model for matter, our posterior self-consistently combines the following sources of cosmological information in large-scale structure: \begin{enumerate} \item The leading- and next-to-leading order power spectrum, and leading-order bispectrum. In particular, this breaks the bias-amplitude degeneracy, which is perfect at linear order, using the second-order displacement term. Further, it allows for improved constraints on the slope (spectral index) of the linear power spectrum by extending the range in wavenumber $k$ useable for robust constraints. \item Fully resummed BAO reconstruction using 2LPT displacements, both at the power spectrum and bispectrum levels. \item Correct description of curvature and tidal effects on the local BAO scale. This effectively includes information from the 4-point function as well as higher-order statistics, through the highly nontrivial posterior in \refeq{post}. \end{enumerate} These probes translate into constraints on cosmological parameters. The first point allows for direct constraints on $\sigma_8$ and parameters which control the shape of the power spectrum, such as $\Omega_m$, $n_s$, and the sum of neutrino masses. The second and third allow for constraints on the expansion history and thus dark energy equation of state. Quantifying the precise information content in terms of parameter constraints will be the subject of upcoming work. We have also presented concrete numerical results validating the theoretical derivation. Using halo catalogs obtained from N-body simulations as physical tracers of the underlying matter density field, we computed maximum-likelihood estimates of bias parameters and $\sigma_8$. Evaluating the performance of our EFT likelihood approach at different scales and halo mass ranges, we demonstrated that the input parameter values can be consistently recovered. This is the first demonstration of unbiased cosmology inference for forward-modeling approaches to date. The likelihood can be straightforwardly extended to include higher-order bias contributions. While this might mean that non-Gaussian corrections to the likelihood and more noise fields need to be considered, the fundamental approach remains the same, and the dimensionality of the inference problem hardly changes. Thus, this approach appears much more feasible than explicitly measuring ever higher $n$-point functions. Further, many additional types of cosmological physics can be included straightforwardly, such as the scale-dependent bias induced by primordial non-Gaussianity, as well as multiple tracers within the same volume. The main missing ingredient for the application to actual data are nontrivial survey geometries (mask) and redshift-space distortions. The former leads to a nondiagonal noise covariance as discussed in \refsec{GaussF}, which however only needs to be determined once. Redshift-space distortions can be treated in the EFT approach as well \cite{perko/etal:2016,fonseca/etal:2018,pkgspaper}, and can be included in forward modeling correspondingly.\footnote{Redshift-space distortions allow for a measurement of the growth rate in the linear regime. However, if line-of-sight-dependent selection effects are present, then this information is removed due to a degeneracy with a new bias parameter. Then, nonlinear information is necessary to infer the growth rate \cite{pkgspaper}, in close analogy with the discussion in \refsec{sigma8} here.} We leave all of these developments to future work. | 18 | 8 | 1808.02002 |
1808 | 1808.00007_arXiv.txt | We present an update to the chemical enrichment component of the smoothed-particle hydrodynamics model for galaxy formation presented in \cite{S05} in order to address the needs of modelling galactic chemical evolution in realistic cosmological environments. Attribution of the galaxy-scale abundance patterns to individual enrichment mechanisms such as the winds from asymptotic giant branch (AGB) stars or the presence of a prompt fraction of Type Ia supernovae is complicated by the interaction between them and gas cooling, subsequent star formation and gas ejection. In this work we address the resulting degeneracies by extending our implementation to a suite of mechanisms that encompasses different IMFs, models for yields from the aforementioned stars, models for the prompt component of the delay-time-distribution (DTDs) for Type Ia SNe and metallicity-dependent gas cooling rates, and then applying these to both isolated initial conditions and cosmological hydrodynamical zoom simulations. We find DTDs with a large prompt fraction (such as the bimodal and power-law models) have, at $z=0$, similar abundance patterns compared to the low-prompt component time distributions (uniform or wide Gaussian models). However, some differences appear, such as the former having systematically higher [X/Fe] ratios and narrower [O/Fe] distributions compared to the latter, and a distinct evolution of the [Fe/H] abundance. | Over the last decades, the accumulation of precise and detailed observations on the chemical properties of the Milky Way and external galaxies have enabled substantial improvements in our understanding of cosmic structure formation. Chemical information of the gas and stars in galaxies, as well as in the intergalactic medium, and for different redshifts, are important tools to complement observations of stellar masses, current star formation rates, magnitudes/colors, morphologies. Important physical processes occurring in galaxies leave imprints on their chemical properties, which carry information on their past star formation rates, the occurrence of mergers and accretion, and the amount of dust that obscures their light. In particular, the chemical properties of the stellar component provide information of the galaxies at different cosmic times, enabling the reconstruction of the galaxy's formation history. In previous decades the importance of chemical evolution in cosmological simulations has often been seen to be subordinate to the greater uncertainties of the feedback and star formation history of the universe which can modify the stellar populations by an order of magnitude, rather than, for example, the changes in Initial Mass Function (IMF), SNIa rates, stellar life-times and chemical yields, IMF changes being the most significant with changes up to a factor 2 in the fraction of massive stars \citep[e.g.][]{Vincenzo_2016}. Nevertheless, our knowledge of the galactic chemical profile has made significant progress, with large increases in both the number of species for which we can measure abundances \citep[e.g. the neutron-capture elements of ][]{Sneden_2008} and the number of stars for which $\alpha$-element abundances can be measured \citep[e.g. ][]{Gilmore_2012}. Combining this with the kinematical data from the European Space Agency's Gaia satellite will lead to tight constraints on the galactic chemical evolution of our own galaxy. On the other hand, observations of the chemical properties of external galaxies, both locally and at higher redshifts, now enable study of the evolution of chemical species in galaxies and to look for links between their chemical and dynamical properties, allowing a better understanding on their formation histories. Accompanying the observational evidence has been increasingly detailed models of the stellar populations, their nucleosynthesis, the evolution and subsequent enrichment of the ISM and the mechanisms by which this is mixed on galactic scales. Consequently, the constraints have the potential to distinguish not just between e.g. the choice of a \citet{Salpeter_1955, Kroupa01} or \citet{Chabrier03} IMF by their effect on the fraction of massive stars (e.g. \citealp{Francois_2004}), but on a whole interrelated network of processes that incorporate different models for the SNII yields \citep[e.g.][]{WW95, P98, Limongi_2003}, the AGB ejecta \citep{M01, Marigo_2008}, the type Ia nucleosynthesis \citep{Nomoto_1997,Iwamoto_1999}, the existence of a prompt component for the aforementioned \citep{Mannucci_2006}, the metal dependent life-times of all their progenitors \citep[][P98 hereafter]{P98}, and how rapidly the metals are mixed \citep[e.g.][]{Aumer_2013} or ejected \citep{Creasey_2015} and their effects on gas cooling \citep{W09}. As observations of chemical properties increase in amount and detail, a theoretical understanding of the building up of the chemical history of galaxies is needed, and in fact over the last decades a great deal of effort has been made in the field of galactic chemical evolution. The evolution of chemical species in galaxies is intimately linked to the star formation process, the amount of feedback that injects energy into the medium and the mixing and redistribution of chemical elements. Thanks to the advances in numerical techniques and computer processing capabilities we are gradually improving beyond one zone/ordinary differential equation models of \citet{Matteucci_1994, Matteucci_2001, Vincenzo_2017} or hydrodynamical models based on N-body trees \citep{Minchev_2013} or the semi-analytical models of \citet{Nagashima_2005, Arrigoni_2010, Yates2013,DeLucia_2014} to simulations in a full cosmological context of the chemical properties of galaxies. These have evolved from the chemical enrichment modules of \cite{Mosconi01, Lia02, Kawata03, Tornatore04, Okamoto05, S05}, and most of the current state-of-the-art simulation codes consider the production and distribution of metals, their mixing and their effects on the gas cooling process \citep[][]{Crain2009, Schaye2015, Illustris}, which have enabled substantial improvements in our understanding of cosmic structure formation and the chemical evolution of the Universe. Models of chemical enrichment are, however, subject to many uncertainties, as the shape and universality of the IMF, the chemical yields of SNII, the nature of the SNIa progenitors and the SNIa delay distribution are still under debate. In three-dimensional hydrodynamical simulations there is still no consensus on the precise mechanism by which one should implement feedback to regulate star formation by reheating or preventing the gas from cooling, which results into uncertainties in the level, distribution and timing of the chemical production and mixing. Constraints have, however, been more forthcoming with simulations which restrict the geometry such as \citet{Francois_2004, Matteucci_2009} on the effects of the Type Ia distribution, and \citet{Cote_2017} w.r.t. inflows and outflows. The lack of fundamental theories for several complex physical processes has lead to a number of models which for a same set of initial conditions predict galaxies with different properties \citep{S12}. Despite these problems, such simulations have proven to be extremely useful for an understanding of the effects of given physical processes (mergers, interactions, mixing, instabilities, accretion) on the properties of galaxies, allowing a better interpretation of observational results both at the present time and at higher redshifts. In this work we attempt to restrict our focus away from the emphasis on feedback and onto the application of the many advances that have been made in the chemical enrichment mechanisms and related processes since the pioneering works of \citet{Katz_1992,Steinmetz94,Raiteri96} and \citet{Mosconi01}. In particular we are interested in the effects of updates to type II SN yields, cooling, AGB enrichment and the time distribution of type Ia SNe. We present an update of the chemical model of \cite[][S05 hereafter]{S05} and evaluate the effects of the different assumptions on the chemical properties of the baryons, using both idealized and cosmological initial conditions. The paper is organized as follows. In Section~\ref{sec:implem} we present the chemical model; in Section~\ref{sec:isolated} we use idealized initial conditions to isolate and test the effects of the different assumptions for the IMF, chemical yields, delay time distributions of SNIa and cooling. We present results for cosmological simulations in Section~\ref{sec:cosmo}; and in Section~\ref{sec:conclu} we summarize our conclusions. | \label{sec:conclu} In this work we present several updates to the \cite{S05} model of galaxy formation focusing on the chemical enrichment mechanisms, and perform a series of controlled experiments to disentangle their interacting effects. We implement additional models for the IMF, the type II SNe yields, included ejecta from asymptotic-giant branch stars, updated the cooling tables and added several models of the delay-time distribution of type Ia SNe. In order to validate these models without performing an overwhelming series of calculations we utilised two different sets of ICs. The first was an isolated galaxy which had the advantages of simplicity both in the computational sense and in the star formation history which allowed us to disentangle the effects of the various processes we introduced. For the second we used the cosmological zoom of the very well studied Aquarius-C halo, which has a much more realistic formation history and subsequent star formation episodes, with the increased complexity in computation and analysis that involves, and as such we simulated only the latest combination of IMF, AGB and SNII yields, and varied the delay time distribution. With this cosmological initial condition we have also tested the effects of varying the cooling function which also considers the coupling to the UV background field. Our primary results are as follows. \begin{enumerate} \item{The ascending fraction of high mass stars in the Salpeter, Kroupa and Chabrier IMF respectively drives corresponding trends in stellar mass and metallicity, with the Kroupa IMF almost universally being the intermediate case. The Chabrier IMF has the most feedback and forms the least mass of stars, a trend that is seen in both the isolated and cosmological simulations. This higher feedback is, however, not sufficient to prevent it also producing the highest metallicities in both gas and stars.} \item{The effect of WW95 vs. P98 SNII yields varies greatly between elements. Most significant is for N and Mg, with the P98 model producing excesses of the order of 1 dex in the stellar abundances, and at a lower level for Ne and O, with differences of about 0.5 dex in the isolated simulations. These correspond to large changes in the [O/Fe] ratios of about 0.5 dex with P98.} \item{AGB stars return chemicals to the ISM in a time distribution with both a significant prompt fraction and relatively heavy tails. We settled on a default implementation of 3 enrichment periods of 100 Myr, 1 Gyr and 8 Gyr per star particle, the latter two capturing the extended phases and the earliest for the material that is quickly recycled. The AGB stars are particularly effective at polluting with C and N, and our isolated simulations exhibit increases in [C/Fe] and [N/Fe] stellar ratios by $0.6$~dex and $0.4$~dex respectively. } \item{The effects of switching from \citet{SD93} to \citet{W09} cooling functions can be observed in the phases of star formation in isolated and cosmological halos, however the net effect is rather modest, likely due to the tight self-regulation of star formation via stellar feedback. } \end{enumerate} The delay time distribution of type Ia SNe is of particular interest as it affects the typical time-scales of the iron release, and this effect might still be imprinted in the properties of the chemical properties of the stars and gas in galaxies. In our various implementations of SNIa models, the distribution of SNIa takes the form of a very extended process along with the presence or absence of a prompt component. Since this can interplay with hierarchical formation and gas accretion, we tested these both with isolated and cosmological simulations. In the following we summarize our main results. \begin{enumerate} \item{The prompt component is maximized in our power-law and bimodal models, and at the other extreme the narrow Gaussian (centered on 1 Gyr) has the least prompt SNIa. The models with the prompt component produce the highest stellar element ratios at late times, in both the isolated and cosmological simulations.} \item{The models with the prompt component also exhibit the narrower [O/Fe] distributions at the present time, suppressing the long tail to low stellar [O/Fe] ratios. Conversely at the opposite extremes the narrow Gaussian and uniform delay-time distributions create the lowest overall [O/Fe] ratios, also having a low [O/Fe] tail.} \end{enumerate} This work marks an essential step in linking the observed abundance patterns of stars in our own galaxy to its star formation, gas evolution and enrichment history over cosmic time. The most immediate application of this work is for more detailed studies of the abundances of individual elements and their distribution, both spatially and in terms of age in our galaxy. Having a validated cosmological model also allows the examination of the statistics of galaxies, for example the analysis of dispersions and gradients of $\alpha$ and Fe, and the correlations with star formation. This will require some additional simulation effort to provide a significant sample of galaxies evolved in a $\Lambda$CDM cosmology to capture the effects of diverse formation histories and environments, and we leave this to a future paper. | 18 | 8 | 1808.00007 |
1808 | 1808.05622_arXiv.txt | We report a synchronized kinematic shift of \civ and \siv broad absorption lines (BAL) in a high-ionization, radio-loud, and X-ray bright quasar SDSS-J092345$+$512710 (at \zem\ $\sim$ 2.1627). This quasar shows two broad absorption components (blue component at $v \sim 14,000$~\kms, and red component at $v \sim 4,000$~\kms\ with respect to the quasars systemic redshift). The absorption profiles of \civ and \siv BAL of the blue component show decrease in outflow velocity with an average deceleration rate of $ -1.62_{-0.05}^{+0.04}$~\cmss\ and $-1.14^{+0.21}_{-0.22}$~\cmss\ over a rest-frame time-span of 4.15 years. We do not see any acceleration-like signature in the red component. This is consistent with dramatic variabilities usually seen at high velocities. During our monitoring period the quasar has shown no strong continuum variability. We suggest the observed variability could be related to the time dependent changes in disk wind parameters like launching radius, initial flow velocity or mass outflow rate. | \label{lab:xbal_intro} Outflows are ubiquitous and appear to be the main source of active galactic nucleus (AGN) feedback which regulates black hole growth and host galaxy evolution as well as enrich the intergalactic and circumgalactic medium around galaxies \citep[see][]{Ostriker2010ApJ...722..642O,Kormendy2013ARA&A..51..511K}. These feedback processes most likely drive the well-known observed correlation between the supermassive black hole (SMBH) mass and physical properties of the host galaxy, along with the steep decline in the number density of galaxies at high masses \citep{Ferrarese2000ApJ...539L...9F,Hopkins2005ApJ...630..705H, Ostriker2010ApJ...722..642O}. The signatures of strong outflows are directly observed in roughly 20 percent of quasar population via broad ultraviolet-resonance absorption lines (BAL) \citep[][]{Weymann1991ApJ...373...23W, Trump2006ApJS..165....1T,Gibson2009ApJ...692..758G}, spanning a large range of outflow velocities from $1000$~\kms\ up to several $10,000$ \kms\ \citep[e.g.,][]{Weymann1991ApJ...373...23W,Hamann1997ApJ...478...87H,Hidalgo2011MNRAS.411..247R,Rogerson2016MNRAS.457..405R}. Nonetheless, many aspects of quasar outflows remain poorly understood, including the gas geometry, acceleration mechanism(s) and their influence on the host galaxy and its environments. \par BAL variability study is a promising technique for constraining the structure and location of the associated wind. In recent systematic studies of BAL variability, BAL troughs are commonly observed to show variability in absorption strength \citep{Barlow1994PASP..106..548B,Lundgren2007ApJ...656...73L,Gibson2008ApJ...685..773G,Capellupo2011MNRAS.413..908C,Capellupo2012MNRAS.422.3249C,Capellupo2013MNRAS.429.1872C,Ak2013ApJ...777..168F,Vivek2014MNRAS.440..799V,Grier2015ApJ...806..111G,McGraw2018MNRAS.475..585M} and/or profile, also known as ``transient BALs'' showing emergence or disappearance of BAL features \citep[e.g.,][]{Hamann2008MNRAS.391L..39H,Hall2011MNRAS.411.2653H, Hidalgo2011MNRAS.411..247R,Ak2012ApJ...757..114F,Vivek2016MNRAS.455..136V, McGraw2017MNRAS.469.3163M}, over a broad range of rest-frame timescales, ranging from months to years. However, the signature of acceleration (e.g., kinematic shift of absorption profile) are more scarce, reported only a few times \citep[e.g.,][]{Vilkoviskij2001MNRAS.321....4V, Rupke2002ApJ...570..588R, Gabel2003ApJ...595..120G, Hall2007ApJ...665..174H,Joshi2014MNRAS.442..862J, Grier2016ApJ...824..130G}. Mechanisms proposed to understand the observed BAL variability invoke changes in ionization state and/or the movement of individual clouds or substructures in the outflow, across our line of sight \citep{Lundgren2007ApJ...656...73L,Hall2007ApJ...665..174H,Hamann2008MNRAS.391L..39H}. In most cases the BAL variations are not found to be correlated with the optical continuum variations. Also not all velocity components seen in absorption show correlated variations. This hints towards mechanisms other than photoionization induced variations. However, a unified understanding of BAL variations is an ongoing endeavor. \par The BAL features are widely believed to be formed in ``disk winds'', launched from the surface of the accretion disk at 10-100 light days from the central SMBH (of $\sim 10^9 \rm \ M\odot$) mainly driven by radiative forces \citep{Arav1994ApJ...427..700A,Murray1995ApJ...451..498M, Proga2000ApJ...543..686P,Higginbottom2014ApJ...789...19H}. In radiation-driven scenarios, the wind is efficiently accelerated to high velocities by invoking the shielding gas close to the base of outflow which prevents the UV-absorbing gas from becoming overionized by nuclear X-ray and extreme-ultraviolet (UV) photons \citep{Murray1995ApJ...451..498M,Proga2000ApJ...543..686P}. The above paradigm is challenged by the observed flows having high velocities and lower degree of ionization in X-ray bright mini-BAL (typical full width half maximum of $500 - 2000$~\kms) quasars \citep{Hamann2008MNRAS.391L..39H,Hamann2013MNRAS.435..133H}. These observations favor the substructured flow, involving tiny dense clouds with a low volume filling factor, driven out by radiative forces while being confined by magnetic pressure ~\citep{Rees1987QJRAS..28..197R,Baskin2014MNRAS.445.3025B,Matthews2016MNRAS.458..293M}. It suggests that the strong radiative shielding gas may not be universal component of quasar outflows for accelerating the gas to high speeds. This idea is also supported by the recent high-energy X-ray observations showing that a large fraction, $\sim 6-23\%$, of BAL quasars among the general BAL quasar population are perhaps intrinsically X-ray weak in nature \citep{Luo2013ApJ...772..153L,Luo2014ApJ...794...70L,Teng2014ApJ...785...19T,Liu2018ApJ...859..113L}. The emerging picture of BAL outflows suggests that the mini-BALs and BALs arise from the same quasar wind, where BALs form in the main part of the outflow near the accretion disc plane while mini-BALs form along sightlines at higher latitudes \citep{Ganguly2001ApJ...549..133G,Hamann2008MNRAS.391L..39H,Hamann2013MNRAS.435..133H}. This also explains the observed X-ray bright nature of mini-BALs. Interestingly, a new population of X-ray bright BAL quasars is recently discovered in X-ray surveys \citep[e.g.,][]{Giustini2008A&A...491..425G,Gibson2009ApJ...692..758G, Streblyanska2010AIPC.1248..513S,Liu2018ApJ...859..113L} which further possess major challenges to the models of BAL outflows. Note that, if the X-ray bright BAL quasars are preferentially originated in a structured flow viewed along the sight lines of higher latitudes than one would naively expect to see the combination of line shift and line strength variability in this subclass. Indeed, in our recent efforts to probe the variability nature of X-ray bright BAL quasars we have find two such rare cases of kinematic shift and strength variability of the \civ BAL trough \citep[e.g.,][]{Joshi2014MNRAS.442..862J}. Given the X-ray weak nature of general population of BAL quasars, a systematic study of this rare population of X-ray bright BAL quasars will provide important observational constrains on the BAL geometry and the physical mechanisms for launching and accelerating the quasar outflows. Here, we report the detection of a deceleration-like signature in \civ and \siv BAL outflows towards X-ray bright quasar J092345$+$512710. This source is part of our on going monitoring program of BAL spectral variability in rare X-ray bright BAL quasars \citep[see,][]{Joshi2014MNRAS.442..862J}. This paper is organized as follows. Section~\ref{lab:xbal_obsand dataredu} describes the observations and data reduction. In Section~\ref{lab:xbal_RES}, we present the results of our analysis followed with the discussion and conclusion in Section~\ref{lab:xbal_DnC}. | \label{lab:xbal_DnC} We report on the kinematic shift of \civ and \siv BAL profiles in a radio-loud quasar SDSS-J092345$+$512710. This quasar belongs to a rare sub-class of X-ray bright BAL quasars with a neutral hydrogen column density of $N_{\rm H} < 3 \times 10^{22}\ \rm cm^{-2}$ and an optical to X-ray spectral index, $\alpha_{ox}$, of $\sim- 1.59$ which is greater than the typical $\alpha_{ox}$ measured for soft X-ray weak quasars, i.e., $\alpha_{ox} < -2$ \citep{Giustini2008A&A...491..425G}. In addition, the difference between observed $\alpha_{ox}$ and expected $\alpha_{ox}$ from the UV luminosity, i.e., $\Delta \alpha_{ox}$, is found to be 0.04 \citep{Giustini2008A&A...491..425G}. We detect an average acceleration-like signature of $-1.62 \pm ^{0.04}_{0.05}$~\cmss\ and $-1.14 \pm ^{0.21}_{0.22}$~\cmss\ in \civ and \siv BALs trough, respectively, over a rest-frame time span of 4.15 years (Table~~\ref{lab:bal_kinematics}). We do not find any acceleration signature for the ``red'' component. Interestingly, we find that the measured deceleration for ``blue'' \civ BAL may not be constant over time, the rate of deceleration between Epoch 2 and 4 is more rapid, about factor 1.4 (significant at only $1.3\sigma$ level) higher than Epoch 1 and 2. \par In the handful of previous studies of the kinematic shift in individual objects \citet{Vilkoviskij2001MNRAS.321....4V}, \citet{Rupke2002ApJ...570..588R} and \citet{Hall2007ApJ...665..174H} have measured a positive kinematic shift (i.e., acceleration) of BAL with a typical acceleration rate of $0.035 \pm 0.016$ \cmss, $0.08 \pm 0.03$ \cmss, and $\sim 0.0154 \pm 0.025$ \cmss\ respectively. The first detection of negative kinematic shift (i.e., deceleration) was reported by \citet{Gabel2003ApJ...595..120G} in the Seyfert galaxy NGC 3783. Using the multi-epoch observations they have found a synchronous kinematic shift of C~{\sc iv}, Si~{\sc iv}, and \nv\ absorption features, while preserving the absorption profile, with a varying deceleration rate which raises from $-0.1 \pm 0.03$~\cmss\ to $-0.25 \pm 0.05$~\cmss\ in the later interval. In addition, \citet{Joshi2014MNRAS.442..862J} have detected relatively larger deceleration rate of $-0.7 \pm 0.1$~\cmss\ and $-2.0 \pm 0.1$~\cmss\ of \civ BAL in two X-ray bright BAL quasars over rest-frame time-spans of 3.11 and 2.34 years. Recently, \citet{Grier2016ApJ...824..130G} have performed the first systematic search for BAL acceleration using three epoch SDSS spectra of 140 BAL quasars over timescales of 2.5$-$5.5 years and found only 3 cases, two acceleration and one deceleration, of monolithic velocity shift showing an overall lack of widespread BAL acceleration. They have measured an average acceleration/deceleration rate of $0.63^{+0.14}_{-0.13}$~\cmss, $0.54 \pm 0.04$~\cmss\, and $-0.83^{+0.19}_{-0.24}$~\cmss, which is comparable to the present study. Using the upper limits for \civ BAL acceleration and deceleration in 76 BAL troughs, over a rest-frame timescales of 2.5 to 5.5 years, they show that the majority of BALs exhibit stable mean velocities to within about 3 per cent. Interestingly, for all the three cases they have found that the wind acceleration rate is not constant over time. The observed kinematic shift in BAL can be produced due to several reasons e.g., actual line-of-sight acceleration of a shell of material from an intermittent outflow, directional shift in the outflow, and changes in velocity dependent quantities such as ionization state, or covering factor \citep{Hall2002ApJS..141..267H, Gabel2003ApJ...595..120G,Hall2007ApJ...665..174H}. At first, we consider the simplest case of changing radiation energy which will lead to decrease in the injected momentum and thus will slow down the wind. As mentioned before using the CRTS light curve we find that our source shows a negligible variation, less than order of half a magnitude, over the time spanned by our observations. In addition, we do not see any significant variation in the emission line flux which responds to the continuum. So we conclude that change in radiative energy/momentum may not be the primary source of deceleration. Secondly, we consider the possibility of gravitational force for the bulk radial deceleration of the flow. Most of the accretion disk wind models predicts the BAL outflow at a typical distance of 0.01 pc from the central source \citep{Murray1995ApJ...451..498M, Proga2000ApJ...543..686P} whereas the observations suggest that the outflows are located at much larger distances in the range of parsecs to several kilo-parsecs. Recently, using \siv BAL trough \citet{Xu2018arXiv180501544X} have shown that more than 75\% of BAL outflows are at $> 100$ pc \citep[see also,][and references therein]{Arav2018ApJ...857...60A}. Given the central black hole mass of J092345$+$512710 to be $3 \times 10^{9}\ \rm M\odot$ \citep{shen2011ApJS..194...45S} and assuming the absorbing cloud is at a typical distance of 1 pc from the central ionizing source where gravity in mainly dominated by the black hole, at the SDSS Epoch 1 we find that the escape velocity is much lower ($\sim 5081$~\kms) than the average outflow speed of 15,000 \kms. It indicates that the observed deceleration-like signature of the outflow is very unlikely caused by the deceleration of continual flow due to gravitational force. Simultaneous coverage of different ions at high signal to noise and spectral resolution is needed to constrain the distance of the absorbing gas from the quasar. Alternatively, if we consider the absorbing clouds to be at larger distances, it is quite plausible that the outflowing gas may interact with the ambient material in the host galaxy which in turn may cause the deceleration of BAL winds \citep[see,][]{Leighly2014ApJ...788..123L}. In such scenario the interaction will also change the ionization and thermal state of the gas thereby introducing profile variation. This is not evident from the observed BAL profiles. In view of the fact that majority of BALs are stable within 3 percent of their mean velocity \citet{Grier2016ApJ...824..130G} argue that BAL cloud may not have traveled sufficiently far to interact with the ambient medium. However, more such examples of BAL deceleration will be crucial to test this scenario. In disk wind models the BAL profiles are not only produced in a steady smooth wind \citep{Murray1995ApJ...451..498M}, but also in a unsteady clumpy flows \citep{Proga2000ApJ...543..686P,Proga2004ApJ...616..688P} and magnetically confined disk wind \citep{Arav1994ApJ...427..700A,deKool1995ApJ...455..448D}. In case of steady wind whose density and velocity as a function of radius is governed by force equation (balance between radiative acceleration and gravity) and mass and momentum conservation at all radial distances `$r$', it is known that parameters such as the initial injection position and velocity of the wind and mass outflow rate will alter velocity and density at a radial distance from quasar (for example, see the basic set of equations given in Section 2.4.2 of \citealt{Borguet2010A&A...515A..22B}). Therefore, even if the quasar luminosity does not change, any time variation in the initial condition of the wind (launch radius, initial velocity, and mass outflow rate) can lead to acceleration or deceleration signatures (see Section 4.1 of \citealt{Grier2016ApJ...824..130G}). It may be noted that, absorption profile change introduced by the radial velocity profile change and associated density profile change due to continuity equation will also produce ionization change effects (even when radiation field remains constant). Therefore, while producing velocity shift one will have independent constraints from the equivalent width and equivalent width ratio variations. \citet{Proga2000ApJ...543..686P} have shown that the outflowing disk wind is self shielded and unsteady which is radially accelerated to relatively high velocities by the UV radiation. In addition, such a wind can be spatially inhomogeneous having velocity and time dependent partial coverage \citep[see also,][for 3D axisymmetric disc wind simulations]{Dyda2018MNRAS.475.3786D,Dyda2018MNRAS.478.5006D}. This variable covering factor may also lead to the apparent profile shape variations \citep{Proga2012ASPC..460..171P}. Interestingly, \citet{Waters2016MNRAS.460L..79W} have shown that the acceleration by the radiation pressure is very efficient when flux is time-dependent, it can lead to a large change of about factor two in the net acceleration for a small flux variation of $\sim$ 20\%. Given the fact that the optical flux is not a perfect tracer of the line-driving flux and quasars may show a higher amplitude variability in UV \citep{Welsh2011A&A...527A..15W} it is possible that observed deceleration-like signature can also be produced in the time dependent disk winds. Alternatively, the disk wind may involve many small self-shielded clouds with low volume filling factor, driven out by radiative force while being confined by magnetic pressure \citep{deKool1995ApJ...455..448D}. Not only the magnetically confined clouds can maintain a roughly constant density and ionization across the acceleration region \citep{dekool1997ASPC..128..233D}, but also have superthermal velocity dispersion and therefore only a few of them can explain the observed broad and smooth BAL profiles \citep{Bottorff2000MNRAS.316..103B}. In view of the X-ray bright nature of J092345+512710 the BAL troughs may well favor the small clouds scenario, instead of the homogeneous radial outflows, producing the observed velocity shift due to acceleration or change in covering factor and/or optical depth. Therefore, disentangling these effects would further require variability follow-ups. Also, high signal to noise spectra at higher spectral resolution will be important to further constrain the outflow models. In our BAL variability studies of X-ray bright BALQSOs till now we have found significant profile shift in three cases. \citet{Joshi2014MNRAS.442..862J} reported two cases of deceleration-like signatures. Unlike in the present case the deceleration is also accompanied by profile shape variation in other two cases. However, there are few common trends we notice in all three cases: (i) The decelerating components typically have large ejection velocities (i.e., $> 10,000$~\kms); (ii) There is associate absorption at low velocities without showing any signatures of acceleration; and (iii) the optical quasar continuum has not varied appreciably. Additionally, it is intriguing that we have not yet seen acceleration signatures in X-ray BALs combined to the fact of high frequency of occurrence of deceleration-like signatures in them compared to the X-ray weak typical BALs. Confirmation of this trend in large sample will be interesting for understanding the physical origin of X-ray loudness (or weakness) of BAL quasars. | 18 | 8 | 1808.05622 |
1808 | 1808.00970_arXiv.txt | We present an X-ray stacking analysis of $\sim$75,000 star-forming galaxies between $0.1<z<5.0$ using the \textit{Chandra} COSMOS Legacy survey to study the X-ray emission of low-luminosity active galactic nuclei (AGN) and its connection to host galaxy properties. The stacks at $z<0.9$ have luminosity limits as low as $10^{40}-10^{41}$ erg s$^{-1}$, a regime in which X-ray binaries (XRBs) can dominate the X-ray emission. Comparing the measured luminosities to established XRB scaling relations, we find that the redshift evolution of the luminosity per star formation rate (SFR) of XRBs depends sensitively on the assumed obscuration and may be weaker than previously found. The XRB scaling relation based on stacks from the \textit{Chandra} Deep Field South overestimates the XRB contribution to the COSMOS high specific SFR (sSFR) stacks, possibly due to a bias affecting the CDF-S stacks because of their small galaxy samples. After subtracting the estimated XRB contribution from the stacks, we find that most stacks at $z>1.3$ exhibit a significant X-ray excess indicating nuclear emission. The AGN emission is strongly correlated with stellar mass but does not exhibit an additional correlation with SFR. The hardness ratios of the high-redshift stacks indicate that the AGN are substantially obscured ($N_{\mathrm{H}}\sim10^{23}$ cm$^{-2}$). These obscured AGN are not identified by IRAC color selection and have $L_X\sim10^{41}-10^{43}$ erg s$^{-1}$, consistent with accretion at an Eddington rate of $\sim10^{-3}$ onto $10^7-10^8 M_{\odot}$ black holes. Combining our results with other X-ray studies suggests that AGN obscuration depends on stellar mass and an additional variable, possibly the Eddington rate. | \label{sec:intro} A key ingredient of galaxy evolution that is still not fully understood is the relationship between the formation of stars and the growth of the supermassive black hole (BH). The processes regulating galaxy and BH growth are thought to be linked across cosmic time because of the observed correlations at $z=0$ between BH mass and the large scale properties of galaxies such as stellar mass ($M_*$; e.g., \citealt{magorrian98}; \citealt{ferrarese00}; \citealt{gebhardt00}; \citealt{haring04}; \citealt{mcconnell13}; \citealt{kormendy13}), and the striking resemblance between the cosmic histories of star formation and BH accretion (e.g., \citealt{hopkinsbeacom06}; \citealt{silverman08}; \citealt{aird10}). Different physical mechanisms that could trigger BH growth have been proposed, including major galaxy mergers (\citealt{sanders88}; \citealt{dimatteo05}; \citealt{hopkinsp06}), and secular processes driving gas inflow (\citealt{englmaier04}; \citealt{hopkinshernquist06}; \citealt{hopkins10}). However, uncertainties remain with regards to what extent these different processes contribute to BH growth, and how their relative importance varies with redshift and different levels of BH accretion (see reviews by \citealt{alexander12} and \citealt{heckman14}). \par Numerous studies have investigated the relationship between the accretion of active galactic nuclei (AGN) and the star formation rates (SFRs) of their host galaxies. In high-luminosity AGN ($L_{\mathrm{bol}}\gtrsim10^{45}$ erg s$^{-1}$), a strong correlation between the AGN luminosity, a proxy for the BH accretion rate (BHAR), and SFR is observed (e.g. \citealt{lutz08}; \citealt{bonfield11}; \citealt{mor12}; \citealt{rosario12}); major mergers may drive the high SFRs and BHARs in these galaxies. \par However, the BHARs of lower luminosity AGN and the SFRs of their host galaxies exhibit at most a weak correlation when compared on a source by source basis (e.g. \citealt{shao10}, \citealt{rosario12}). Several studies do observe a strong correlation between the SFR and mean BHAR of moderate luminosity AGN binned by SFR (e.g. \citealt{rafferty11}; \citealt{mullaney12}; \citealt{chen13}; \citealt{azadi15}; \citealt{rodighiero15}; \citealt{lanzuisi17}), but when the galaxies are binned by BHAR, there is no correlation between BHAR and mean SFR (\citealt{rosario12}; \citealt{lanzuisi17}), and the mean SFRs of moderate luminosity AGN hosts are consistent with those of inactive galaxies (e.g. \citealt{santini12}; \citealt{bongiorno12}; \citealt{mullaney15}; \citealt{suh17}). It has been suggested that the apparent contradictions in these observed trends result from the shorter variability timescale of the BHAR when driven by secular processes compared to the galaxy-averaged SFR (\citealt{hickox14}; \citealt{volonteri15}). Whether BH accretion on average is linked to star formation in moderate luminosity AGN remains a matter of debate, as some studies argue that the BHAR in these sources is more strongly connected to the stellar mass of the host galaxy than its SFR \citep{yang17}. \par In addition to AGN variability, a factor which can complicate investigations of the relationship between BH and galaxy growth is obscuration. Obscured AGN may be missed by surveys which probe the rest-frame UV, optical, or near-IR wavelengths, and in X-ray surveys, where they are more easily detected, they can be mistaken for intrinsically lower-luminosity AGN. If any systematic trends exist between the obscured AGN fraction, host galaxy properties, or BHAR, the measured relationships between BHAR and SFR may be biased. While some studies find no correlation between AGN obscuration and SFR (\citealt{rosario12}; \citealt{delmoro16}), others observe such a correlation both in the low-luminosity \citep{castro14} and high-luminosity regimes \citep{chen15}. Conflicting results also exist on the correlation between AGN obscuration and the specific SFR (sSFR$=$SFR/$M_*$) of the host galaxy. \citet{juneau13} find that the obscured AGN fraction increases with sSFR, while \citet{lanzuisi17} find the opposite relation and argue that the obscured AGN sample used by \citeauthor{juneau13} is contaminated. \par Improving our understanding of the connection between BH and galaxy growth requires large galaxy and AGN samples so as to be able to account for AGN variability and to elucidate any trends that may exist between AGN obscuration and host galaxy properties. The 2.2 deg$^2$ \textit{Chandra} COSMOS-Legacy survey \citep{civano16} and associated multi-wavelength coverage of the COSMOS field offer an excellent opportunity to investigate BH-galaxy evolution. Some of the aforementioned studies were based on X-ray selected AGN samples from the COSMOS-Legacy survey (\citealt{suh17}, \citealt{lanzuisi17}), which contains 4016 X-ray detected sources \citep{civano16}. \par Other studies have pushed below the sensitivity threshold of this survey using X-ray stacking techniques in order to probe low-luminosity AGN. Through X-ray stacking of early-type galaxies (ETGs), \citet{paggi16} find enhanced AGN emission in ETGs with lower stellar masses, and evidence for highly absorbed AGN emission at $z\sim1.2$. Performing a similar analysis with dwarf galaxies with $M_*<10^{9.5} M_{\odot}$, \citet{mezcua16} discover an X-ray excess above the expected contribution of X-ray binaries (XRBs), which is consistent with emission from intermediate-mass BHs ($M\sim10^5 M_{\odot}$) that are likely obscured at $z>0.8$. \par In this paper, we present a complementary stacking study of the X-ray emission of star-forming galaxies in the COSMOS field, focused on the low and moderate luminosity ($L_X\sim10^{40}-10^{43}$ erg s$^{-1}$) AGN population. Due to the low average X-ray luminosities reached by our stacks, the XRB contribution can be comparable to or even dominant over the AGN emission. Some studies have shown that the XRB luminosity per SFR and per stellar mass increases with redshift (\citealt{basu13}; \citealt{lehmer16}, hereafter \citetalias{lehmer16}; \citealt{aird17a}, hereafter \citetalias{aird17a}), a trend which is attributed to the formation of more luminous XRBs in lower-metallicity environments (\citealt{dray06}; \citealt{linden10}; \citealt{fragos13}; \citealt{brorby16}). Thus, in this paper, we also discuss the constraints we can place on the redshift evolution of XRBs. \par We describe our star-forming galaxy sample in \S\ref{sec:sample}. Our X-ray stacking analysis and the spectral models we use to calculate rest-frame X-ray luminosities are described in \S\ref{sec:stacking} and \S\ref{sec:specmodel}, respectively. We estimate the XRB contribution to the stacked X-ray emission in \S\ref{sec:xrb}, and compare our results to previous studies of XRB scaling relations in \S\ref{sec:xrbdiscussion}. We discuss the relationship between BH activity and host galaxy properties in \S\ref{sec:agnactivity}, and evidence for an obscured AGN population at $z>1.3$ in \S\ref{sec:obscured}. In \S\ref{sec:conclusion}, we summarize our conclusions and consider how the next generation of X-ray telescopes could improve our understanding of low-luminosity AGN and XRB populations. Throughout this work, we assume a $\Lambda$CDM cosmology with $\Omega_{\mathrm{m}}$ = 0.3, $\Omega_{\Lambda}$ = 0.7, and $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$. | \label{sec:conclusion} We have performed an X-ray stacking analysis of star-forming galaxies that fall below the \textit{Chandra} COSMOS Legacy survey sensitivity limit in order to study low-luminosity AGN populations, their obscuration, and their connection to host galaxy properties. Splitting our sample of 75,000 galaxies by redshift, sSFR, and $M_*$ results in 92 bins, 68 of which are detected with $\geq3\sigma$ confidence in the $0.5-2$~keV band. This X-ray stacking analysis allows us to probe X-ray luminosities as low as $10^{40}-10^{41}$ erg s$^{-1}$ at $z<1.3$ and $10^{41}-10^{42.7}$ erg s$^{-1}$ at $z>1.3$, which are up to two orders of magnitude fainter than the COSMOS sensitivity limit. This study provides insights into both low-luminosity AGN and XRB populations in star-forming galaxies, which are summarized below: \par 1. The hardness ratios and a comparison of the rest-frame $2-10$~keV luminosities derived from the observed $0.5-2$ and $2-8$~keV bands were used to determine an observationally motivated spectral model to convert the stacked count rates into X-ray luminosities. This spectral model includes substantial obscuration ($N_{\mathrm{H}}\sim10^{22}-10^{23}$ cm$^{-2}$), which increases at $z>1.3$ in low and mid-sSFR galaxies. The typical spectral models assumed in similar studies of AGN and XRB populations are relatively unobscured ($N_{\mathrm{H}}<10^{21}$ cm$^{-2}$) and result in calculated X-ray luminosities that are up to 0.6 dex lower than those determined using our observationally motivated spectral model. This difference in the derived luminosities demonstrates the importance of constraining the spectral model in these studies. \par 2. The $L_X$/SFR of the high-sSFR stacks does not exhibit significant redshift dependence. Since such galaxies are expected to be dominated by HMXBs, our high-sSFR stacks suggest that the redshift evolution of $L_{\mathrm{HMXB}}$ per SFR is weaker than found by \citetalias{lehmer16}, but consistent with the local \citetalias{lehmer10} relation and the weaker redshift dependence found by \citetalias{aird17a}. This result depends on the obscuration assumed in the spectral model used to derive the X-ray luminosities. If we adopt an unobscured spectral model, a significant $z$-dependence is measured, but such a spectral model is inconsistent with the observed hardness ratios. \par 3. We find that, regardless of the spectral model assumed, the X-ray luminosities of our high-sSFR stacks are overestimated by the $z$-dependent XRB scaling relation from \citetalias{lehmer16}. The overestimation of the XRB luminosities of most of the COSMOS high-sSFR stacks and some of the mid-sSFR stacks can be partly explained by the fact that the CDF-S stacks contain too few galaxies to adequately sample the bright end of the XRB luminosity functions, causing them to be biased to higher values. The COSMOS stacks, which consist of hundreds to thousands of galaxies, are large enough to adequately sample the HMXB luminosity function but are still insufficient to sample the bright end of the LMXB luminosity function, which is very steep. \par 4. The spectral constraints and X-ray luminosities of the low and mid sSFR stacks at $z>1.3$ provide evidence for the presence of an obscured AGN population. Stacking studies of elliptical and dwarf galaxies in the COSMOS field similarly revealed an obscured AGN population at $z\gtrsim1$. \par 5. Most of the stacks exhibiting residual AGN emission are at $z=1.3-2.3$. For the stacks in this redshift range, the average black hole accretion rate increases with $M_*$, but does not show a significant correlation with SFR once the mass dependence has been taken into account. This result for the low-luminosity AGN population is consistent with results from \citet{yang17}, which includes moderate and high luminosity AGN. \par 6. Assuming the average bolometric correction factor and Eddington ratio for AGN with $L_X<10^{43}$ erg s$^{-1}$ in local star-forming galaxies ($k_{\mathrm{bol}}=16$ and $\lambda_{\mathrm{Edd}}=10^{-3}$), we find that the AGN in our $z>1.3$ stacks have $M_{BH}\sim10^7-10^9 M_{\odot}$. The $M_{\mathrm{BH}}-M_*$ relation for these AGN falls between the \citet{reines15} relation for nearby AGN and the $z$-dependent evolution measured for higher-luminosity AGN by \citet{merloni10}. \par 7. Less than 4\% of the galaxies in each of the COSMOS stacks are identified as AGN via their IRAC colors. It is not plausible that these IRAC-selected AGN dominate the X-ray emission of the stacks; removing these AGN from our stacks does not significantly impact the measured hardness ratios or X-ray luminosities. Thus, IRAC color selection is not sufficient for identifying obscured AGN with $L_X<10^{43}$ erg s$^{-1}$. \par Future wide, deeper X-ray surveys will be crucial for studying obscured AGN and the evolution of XRB populations. In order to measure the obscured fraction as a function of AGN luminosity and to reduce systematic uncertainties associated with the spectral model used to calculate X-ray luminosities, a key improvement that could be made by the next generation of X-ray telescopes is to measure spectra, or at least hardness ratios, for individual low-luminosity sources. \par The Advanced Telescope for High ENergy Astrophysics (\textit{ATHENA}; \citealt{nandra13}) is an ESA large mission expected to be launched in 2028. A multi-tier survey strategy being planned for the \textit{ATHENA} Wide-Field Imager is expected to detect $\sim$400,000 X-ray sources down to $f_X\approx3-7\times10^{-17}$ erg cm$^{-2}$ s$^{-1}$. Based on the number-count distribution measured in the 7 Ms CDF-S field \citep{luo17}, the vast majority of these sources would be AGN. These surveys will allow studies of the BH-host galaxy connection and the obscured AGN fraction for $L_X\gtrsim10^{43}$ erg s$^{-1}$ out to $z\sim6$ with AGN samples a factor of $\sim50$ larger than are currently available. \par A few thousand normal galaxies are expected to be detected in the \textit{ATHENA} surveys, constituting a factor of $>10$ increase over current samples. However, since the fluxes of most of these galaxies will be close to the confusion limit of \textit{ATHENA}, their X-ray properties will be difficult to disentangle due to source blending, and it will not be possible to associate many of them to unique multi-wavelength counterparts due to \textit{ATHENA}'s 5$^{\prime\prime}$ PSF. Attempts to stack \textit{ATHENA} data in order to probe below the sensitivity limit of its surveys will be hampered by its large PSF, which will cause a large fraction of source extraction regions to overlap with one another, biasing the stacked signal. \par In comparison, \textit{Lynx}, an X-ray observatory currently being developed in a mission concept study \citep{gaskin17}, would be ideally suited to study both AGN and XRB populations at high redshift, and to connect their properties to their host galaxies. \textit{Lynx} is planned to have sub-arcsecond resolution over a 20$^{\prime}\times20^{\prime}$ field of view. This sub-arcsecond resolution would enable the unambiguous association of X-ray sources with their host galaxies down to much lower flux limits than \textit{ATHENA}. \par A square degree survey reaching sensitivity limits of $10^{-18}-10^{-19}$ erg s$^{-1}$ cm$^{-2}$ is being considered as part of the \textit{Lynx} concept study. Based on the CDF-S number-count distribution \citep{luo17}, such a survey would detect about 200,000 sources, $>60$\% of which are likely be normal galaxies, reaching luminosity limits of $L_X\sim10^{41}$ erg s$^{-1}$ out to $z\sim3$ and $L_X\sim10^{42}$ erg s$^{-1}$ out to $z\sim6$. About half of these sources (including $\approx$60,000 galaxies) would have $>25$ net counts, sufficient for measuring meaningful hardness ratios, and the brightest 25,000 sources (including about $\approx$12,000 galaxies) would have $>100$ net counts, sufficient for simple spectral fitting. Thus, \textit{Lynx} would be a revolutionary improvement over current studies of low-luminosity AGN and XRBs at high redshift, allowing the individual detection, basic spectral characterization, and unique multi-wavelength association of sources that are currently only X-ray detectable on average through stacking analysis. | 18 | 8 | 1808.00970 |
1808 | 1808.00377_arXiv.txt | \vspace{1cm} \centerline{\bf ABSTRACT}\vspace{2mm} The accelerated cosmic expansion could be due to dark energy within general relativity (GR), or modified gravity. It is of interest to differentiate between them, by using both the expansion history and the growth history. In the literature, it was proposed that the growth index $\gamma$ is useful to distinguish these two scenarios. In this work, we consider the non-parametric reconstruction of the growth index $\gamma$ as a function of redshift $z$ from the latest observational data as of July 2018 via Gaussian Processes. We find that $f(R)$ theories and dark energy models within GR (especially $\Lambda$CDM) are inconsistent with the results in the moderate redshift range far beyond $3\sigma$ confidence level. A modified gravity scenario different from $f(R)$ theories is favored. However, these results can also be due to other non-trivial possibilities, in which dark energy models within GR (especially $\Lambda$CDM) and $f(R)$ theories might still survive. In all cases, our results suggest that new physics is required. | \label{sec1} Since the discovery of the accelerated expansion of our universe in 1998~\cite{Riess:1998cb,Perlmutter:1998np}, the real cause of this mysterious phenomenon is still unclear so far. As is well known, two main types of scenarios are extensively considered in the literature to this end. The first one is to introduce an unknown component with negative pressure (dark energy) in the framework of general relativity (GR). On the contrary, the second one explains the accelerated expansion by using a modification to GR (modified gravity), without invoking dark energy. We refer to e.g.~\cite{Ishak:2018his,Amendola:2016saw,Joyce:2014kja, Clifton:2011jh} for comprehensive reviews. Until now, both scenarios are competent to interpret the accelerated cosmic expansion. Therefore, it is of interest to differentiate between them. Since they cannot be distinguished by using the expansion history solely, it is necessary to consider the growth history in addition (see e.g.~\cite{Linder:2005in, Linder:2007hg,Wei:2008ig,Wei:2008vw} and references therein). In fact, if the models of dark energy and modified gravity share a same expansion history, they might have different growth histories. Typically, the growth history is characterized by the linear matter density contrast $\delta(z)\equiv\delta\rho_m/\rho_m$ as a function of redshift $z$. It is convenient to introduce the growth rate $f\equiv d\ln\delta/d\ln a$, where $a=(1+z)^{-1}$ is the scale factor. Many years ago, a good approximation $f=\Omega_m^\gamma$ has been first proposed in~\cite{Peebles1980,Lahav:1991wc} within GR, where $\gamma$ is the growth index, and $\Omega_m$ is the fractional density of pressureless matter. In the beginning, $f=\Omega_m^\gamma$ was used only at the present time ($z=0$), and it was not valid for any redshift. Since~\cite{Wang:1998gt} it was applied to anything beyond matter, curvature, and a cosmological constant. Finally, not until~\cite{Lue:2004rj} was it applied to gravity other than GR, and then in~\cite{Linder:2005in} generalized to modified gravity, varying equation of state, and an integral relation for growth. Nowadays, the general form $f(z)=\Omega_m(z)^{\gamma(z)}$ has been extensively used in the literature. In e.g.~\cite{Linder:2005in,Linder:2007hg}, it was proposed that the growth index $\gamma$ is useful to distinguish the scenarios of dark energy and modified gravity. In GR, $\gamma=6/11\simeq 0.545$ for $\Lambda$CDM model~\cite{Linder:2005in,Linder:2007hg} (which is approximately independent of redshift), while $\gamma\simeq 0.55$ for many dark energy models~\cite{Linder:2005in}. In the cases of modified gravity, $\gamma\simeq 0.68$ for Dvali-Gabadadze-Porrati (DGP) model ($\gamma=11/16$ is its high redshift asymptotic value)~\cite{Linder:2007hg,Wei:2008ig}, while $\gamma\simeq 0.42$ for most of viable $f(R)$ theories ($\gamma\,\lsim\, 0.557$ certainly for almost all viable $f(R)$ theories, and $\gamma$ decreases when redshift increases)~\cite{Gannouji:2008wt,Tsujikawa:2009ku,Shafieloo:2012ms, Tsujikawa:2010zza}. Since their $\gamma(z)$ lie in a narrow range around the above values respectively, one might differentiate between them. In the literature, the growth indices for some particular models have been constrained by using the observational data, but only the present value $\gamma_0$ and the derivative $\gamma^\prime_0$ were considered usually. Of course, it is better to study the growth index in a model-independent way. In the literature, a common choice is to consider the model-independent parameterizations for $\gamma(z)$, but a particular function form should be given {\it a~prior}. On the contrary, it is worth noting that the goal function could be directly reconstructed from the input data by using some non-parametric methods, such as principal component analysis, and Gaussian processes, without assuming a particular function form. Here, we consider the non-parametric reconstruction of the growth index $\gamma(z)$ as a function of redshift~$z$ via Gaussian processes~\cite{Rasmussen:2006,Seikel:2012uu}, by using the latest observational data. In Sec.~\ref{sec2}, we briefly describe the methodology. In Secs.~\ref{sec3} and \ref{sec4}, the results and the conclusions are given, respectively. We find that $f(R)$ theories, and dark energy models within GR (especially $\Lambda$CDM), are inconsistent with the results in the moderate redshift range, far beyond $3\sigma$ confidence level (C.L.). A modified gravity scenario different from $f(R)$ theories is favored. However, there might be other possibilities for these results, and we will discuss this issue briefly in Sec.~\ref{sec4}. In all cases, our results suggest that new physics is required. \begin{center} \begin{figure}[tb] \centering \vspace{-6mm} % \includegraphics[width=0.85\textwidth]{fs8se.eps} \caption{\label{fig1} The reconstructed $f\sigma_8$, $\delta/\delta_0$, $\delta^\prime/\delta_0$ and $f$ as functions of redshift $z$, by using Gaussian processes with the squared exponential covariance function. The mean and $1\sigma$, $2\sigma$ uncertainties are indicated by the blue solid lines and the shaded regions, respectively. The observational $f\sigma_{8,\,obs}$ data with error bars are also plotted in the top-left panel. See the text for details.} \end{figure} \end{center} \vspace{-12mm} % | \label{sec4} The accelerated cosmic expansion could be due to dark energy within GR, or modified gravity. It is of interest to differentiate between them, by using both the expansion history and the growth history. In the literature, it was proposed that the growth index $\gamma$ is useful to distinguish these two scenarios. In this work, we consider the non-parametric reconstruction of the growth index $\gamma$ as a function of redshift~$z$ from the latest observational data as of July 2018 via Gaussian Processes. Interestingly, we find that $f(R)$ theories and dark energy models within GR (especially $\Lambda$CDM) are inconsistent with the results in the moderate redshift range far beyond $3\sigma$ C.L., due to the arched structure in the reconstructed $\gamma(z)$. A modified gravity scenario different from $f(R)$ theories is favored. Obviously, this result is unusual, and new physics is required. However, it does not mean that dark energy models within GR (especially $\Lambda$CDM) and $f(R)$ theories certainly end. First, one can doubt the observational data used in this work. For instance, the 63 observational $f\sigma_{8,\,obs}$ data compiled in~\cite{Kazantzidis:2018rnb} might be correlated, and contain duplicated points from the same surveys, while the corrections from the choices of the fiducial cosmology should be taken into account. So, this $f\sigma_{8,\,obs}$ sample might require a re-analysis, as preliminarily considered in~\cite{Nesseris:2017vor,Kazantzidis:2018rnb}. Second, the reliability of Gaussian Processes at high redshift might be questionable. Therefore, we should test the growth index $\gamma$ by using another independent method, and cross-check the corresponding results with the ones from Gaussian Processes. In fact, our relevant work will appear in a separate paper~\cite{Yin:2019rgm}. Finally, in the present work, cold dark matter is implicitly assumed, as in Eq.~(\ref{eq5}). If there is a non-zero interaction between dark energy and dark matter, Eq.~(\ref{eq5}) should be changed to \be{eq11} \Omega_m(z)\equiv\frac{8\pi G\rho_m}{3H^2}= \frac{\Omega_{m0}(1+z)^{3+\xi}}{E^2(z)}\,, \ee where $\xi$ characterizes the deviation from uncoupled cold dark matter (note that $\xi$ can be time-dependent in general). Another non-trivial possibility assumes that dark matter is not cold. In fact, for warm dark matter, its equation-of-state parameter $w_m\not=0$, and hence $\Omega_m(z)$ takes a form similar to Eq.~(\ref{eq11}). In both non-trivial cases, the conclusions should be changed, and dark energy models within GR (especially $\Lambda$CDM) and $f(R)$ theories might still survive (see~\cite{Wei:2008vw, Wei:2013rea} for deeper discussions). The new physics in these non-trivial cases lies in the non-zero interaction between dark energy and dark matter, or the induction of warm dark matter. They deserve further investigations. After all, we would like to mention several technical details. One might note that in Figs.~\ref{fig3}$\,\sim\,$\ref{fig6} the reconstructed $\Omega_m$ becomes larger than $1$ at high redshift $z\,\gsim\, 2$, but this is not unphysical in fact. Yes, in GR, the Friedmann equation $H^2=8\pi G(\rho_m+\rho_X)/3$ is unchanged, and hence $0\leq\Omega_m\leq 1$ by definition. However, this is right only for dark energy models in GR. On the contrary, in modified gravity, the Friedmann equation should be modified, and the modification to GR can be equivalent to an effective ``energy component''. So, the effective $\rho_X$ can be negative at high redshift, and hence $\Omega_m>1$ is possible. Since the effective ``energy component'' is not real matter (it is actually the modification to GR, namely a geometric effect indeed), this is allowed by physics. In fact, noting that $0\leq\Omega_m\leq 1$ must be held for dark energy models in GR, our reconstructed $\Omega_m>1$ at high redshift $z\,\gsim\, 2$ in Figs.~\ref{fig3}$\,\sim\,$\ref{fig6} can be regarded as an extra evidence supporting modified gravity against dark energy models in GR. It is known that in modified gravity, e.g. $f(R)$ theories, the growth rate $f=f(z,k)$ is also spatially scale-dependent in general (see e.g.~\cite{Gannouji:2008wt,Tsujikawa:2009ku,Jennings:2012pt}), where the comoving wavenumber $k$ denotes the scale~\cite{Tsujikawa:2007gd}. In principle, the scale-dependence should be taken into account (we thank the referee for pointing out this issue). However, let us have a closer look. In~\cite{Gannouji:2008wt}, they found that $f$ and hence $\gamma$ is scale-independent at redshift $z\,\lsim\, 0.5$ (see their Fig.~2 and the text below Eq.~(4.17) or Eq.~(58) in the arXiv version). At $z=0$, they found $\gamma_0\simeq 0.41$ independent of the scales $k$. At higher redshift, they have a small difference $\Delta\gamma\,\lsim\, 0.04$ between various scales. $\gamma$ becomes smaller as redshift $z$ increases, so that $\gamma\,\lsim\, 0.41$ at higher redshift. In~\cite{Tsujikawa:2009ku}, the results are quite similar. They found that the dispersion of $\gamma$ with respect to the scale $k$ is very small (see their Fig.~2 and the text in Sec.~IV.B), namely $\gamma$ is nearly scale-independent at redshift $z\,\lsim\, 1$ (especially the dispersion of $\gamma$ is nearly absent for scales $k\geq 0.033h\,{\rm Mpc}^{-1}$). Again, $\gamma$ becomes smaller as redshift $z$ increases. On the other hand, from their Figs.~1, 3, 4, 6, 7, one can see that $\gamma_0=\gamma(z=0)$ is smaller than $\sim 0.557$ for various model parameters of viable $f(R)$ theories, and this is independent of the scales $k$. Keeping the above results of~\cite{Gannouji:2008wt,Tsujikawa:2009ku} in mind, let us turn back to our Fig.~\ref{fig9}. First, in all cases of Fig.~\ref{fig9}, $\gamma_0\,\lsim\, 0.5$ at $z=0$ is clearly inconsistent with our reconstructed $\gamma(z)$ beyond $3\sigma$ C.L. As mentioned above, it is found in~\cite{Gannouji:2008wt,Tsujikawa:2009ku} that $\gamma_0=\gamma(z=0)$ is smaller than $\sim 0.557$ for various model parameters of viable $f(R)$ theories, and this is independent of the scales $k$. So, our results of $\gamma_0$ is a bad news to most of these viable $f(R)$ theories, although it is not so decisive. Second, in all cases of Fig.~\ref{fig9}, $\gamma<0.56$ in the moderate redshift range $0.1\,\lsim\, z\,\lsim\, 0.7$ is also clearly inconsistent with our reconstructed $\gamma(z)$ far beyond $3\sigma$ C.L. Actually, in all cases of Fig.~\ref{fig9}, even $\gamma\,\lsim\, 0.6$ is still inconsistent with our reconstructed $\gamma(z)$ beyond $3\sigma$ C.L. in a relatively narrower moderate redshift range. As mentioned above, it is found in~\cite{Gannouji:2008wt,Tsujikawa:2009ku} that $\gamma$ becomes smaller as redshift $z$ increases. Together with the fact that $\gamma_0=\gamma(z=0)$ is smaller than $\sim 0.557$ for various model parameters of viable $f(R)$ theories, it is easy to see that $\gamma\,\lsim\, 0.557<0.56$ in the moderate redshift range $0.1\,\lsim\, z\,\lsim\, 0.7$. This is also independent of the scales $k$. Note that we can further relax $0.56$ to the larger $0.6$ as mentioned above. Therefore, we can still say that most of viable $f(R)$ theories with various model parameters are inconsistent with our reconstructed $\gamma(z)$ in the moderate redshift range beyond $3\sigma$ C.L., and this conclusion is nearly scale-independent actually. Finally, even in the worst case that our results are not applicable to $f(R)$ theories due to the scale-dependence, our other conclusion that dark energy models in GR (especially $\Lambda$CDM) are inconsistent with our reconstructed $\gamma(z)$ in the moderate redshift range far beyond $3\sigma$ C.L. is still valid. New physics is still required. | 18 | 8 | 1808.00377 |
1808 | 1808.02372_arXiv.txt | We present analysis of the normalised 21-cm bispectrum from fully-numerical simulations of intergalactic-medium heating by stellar sources and high-mass X-ray binaries (\textit{HMXB}) during the cosmic dawn. Lyman-$\alpha$ coupling is assumed to be saturated, we therefore probe the nature of non-Gaussianities produced by X-ray heating processes. We find the evolution of the normalised bispectrum to be very different from that of the power spectrum. It exhibits a turnover whose peak moves from large to small scales with decreasing redshift, and corresponds to the typical separation of emission regions. This characteristic scale reduces as more and more regions move into emission with time. Ultimately, small-scale fluctuations within heated regions come to dominate the normalised bispectrum, which at the end of the simulation is almost entirely driven by fluctuations in the density field. To establish how generic the qualitative evolution of the normalised bispectrum we see in the stellar + \textit{HMXB} simulation is, we examine several other simulations - two fully-numerical simulations that include QSO sources, and two with contrasting source properties produced with the semi-numerical simulation \cmfast. We find the qualitative evolution of the normalised bispectrum during X-ray heating to be generic, unless the sources of X-rays are, as with QSOs, less numerous and so exhibit more distinct isolated heated profiles. Assuming mitigation of foreground and instrumental effects are ultimately effective, we find that we should be sensitive to the normalised bispectrum during the epoch of heating, so long as the spin temperature has not saturated by $z\approx 19$. | \label{sec:intro} One of the priorities of modern astrophysics is to try and understand the first stars and galaxies, as well as their subsequent evolution. The formation of luminous sources drastically changed the properties of the Universe. For example, radiation from such sources ionized the hydrogen and helium in the Inter-Galactic Medium (IGM), ultimately causing the Universe to transition from being largely neutral to almost entirely ionized. This phase transition is generally referred to as the \textit{Epoch of Reionization (EoR)}. Remnants of stars, such as black holes and neutron stars, will also produce X-rays which importantly will heat the neutral IGM. Simulations suggest that the IGM transitioned from adiabatically cooling with the background cosmological expansion, to become universally heated. This transition is often referred to as the \textit{Epoch of Heating (EoH)} (\citealt{Loeb2013} provide a comprehensive overview of both the EoR and EoH). The details of sources during the EoH are uncertain, there is indication that dominant sources of X-rays will be high-mass X-ray binaries (HMXBs) and Active-Galactic Nuclei (AGN), with the hot interstellar-medium contributing to the soft end of the X-ray spectrum \citep{Mineo2012}. It is not currently known how much each will ultimately contribute at high-$z$. AGN are the dominant contributor to the X-ray budget at lower redshift, but their abundance is seen to rapidly reduce beyond $z=3$ \citealt{Fan2001, Lehmer2016}, although, mini-quasars could still be a major contributor at high redshifts \citep{Madau2004, Volonteri2009}. However, it is likely that HMXBs will be the main contributor based on the fact that in low-redshift galaxies (in the absense of AGN) they dominate the X-ray production \citep{Fabbiano2006}, and that their abundance (in contrast to AGN) is seen to increase with redshift \citep{Gilfanov2004, Mirabel2011, Mineo2012a}. Simulations also suggest that the very first generation of Population III stars predominantly formed in binary, or multiple systems \citep{Turk2009, Stacy2010}. In order to establish which of these scenarios is true (or indeed if other heating sources might have contributed), we need observational constraints. It is the hope that high-$z$ observations of the 21-cm line of neutral hydrogen will provide a wealth of information about the EoH (as well as the EoR). The CMB will interact with any neutral hydrogen in its path to us, and by looking at fluctuations in the CMB at the frequencies associated with the 21-cm interaction at different redshifts, we can (in principle) make 21-cm maps and learn about the evolution in the properties of neutral hydrogen with time. The observable for the 21-cm line is the offset of the brightness temperature\footnote{ Intensity $I_\nu$ is usually described in terms of a brightness temperature $T_{\rm b}$, defined such that $I_\nu=B(T_{\rm b})$, where $B(T)$ is the Planck black-body spectrum - well approximated by the Rayleigh-Jeans formula at the frequencies relevant to reionization studies.} ($\delta T_{\mathrm{b}}$) relative to that of the CMB ($T_{\mathrm{cmb}}$) \citep{Field1958, Field1959a, Madau1997}, \begin{equation} \begin{split} \delta T_{\rm b}=&\frac{T_{\rm s}-T_{\textsc{cmb}}}{1+z}(1-e^{-\tau_{\nu_0}})\,,\\ \approx&\,27\,\frac{T_{\rm s}-T_{\textsc{cmb}}}{T_{\rm s}}\,x_{\textsc{hi}}(1+\delta)\left[\frac{H(z)/(1+z)}{\d v_{\rm r}/\d r}\right]\\ &\times \left(\frac{1+z}{10}\frac{0.15}{\Omega_{\rm m}h^2}\right)^{1/2}\left(\frac{\Omega_{\rm b}h^2}{0.023}\right) \rm mK \,.\\ \label{eqn:brightTemp} \end{split} \end{equation} \noindent This depends on the cosmological parameters: the Hubble parameter $H(z)=100\,h$, and the matter ($\Omega_{\rm m}$) and baryon ($\Omega_{\rm b}$) density parameters (where $\Omega_i=\rho_i/\rho_{\rm c}$ and $\rho_{\rm c}$ is the critical density required for flat universe). For the analysis performed in this paper we will adopt a $\Lambda$CDM with $\sigma_8=0.80$, $h=0.70$, $\Omega_{\rm m}=0.27$, $\Omega_{\Lambda}=0.73$, $\Omega_{\rm b}=0.044$ and $n_{\rm s}=0.96$. These values are consistent with the values adopted by the simulations of \citealt{Ross2016} analysed in this work and WMAP 7 \citep{Komatsu2010}. Note that unless otherwise stated the analysis in this paper is done on the mean-subtracted brightness temperature, i.e. $\delta T_{\rm b} - \langle \delta T_{\rm b} \rangle$. More important to our discussion here is the dependence of the brightness temperature on density $\delta$, the neutral fraction $x_{\textsc{hi}}$ (which together measure the amount of neutral hydrogen gas present and so provide sensitivity to the EoR), and the spin temperature $T_{\rm s}$ (which measures the relative distributions of electrons over the two levels associated with the 21cm transition). Stars produce copious amounts of Lyman-$\alpha$ radiation, which is incredibly efficient at coupling $T_{\rm s}$ to the thermal temperature of the gas $T_{\rm k}$. Once Lyman-$\alpha$ coupling is complete, the spin temperature provides a probe of the thermal history of the Universe. However, the spin temperature will saturate as $T_{\rm s} \gg T_{\textsc{cmb}}$ and so the brightness temperature can lose sensitivity to fluctuations in the gas temperature if it gets very high. The first generation of 21-cm radio interferometer, such as LOFAR\footnote{The LOw Frequency ARray \url{http://www.lofar.org/}}, MWA\footnote{The Murchison Wide-field Array \url{http://www.mwatelescope.org/}} and PAPER\footnote{The Precision Array to Probe Epoch of Reionization \url{http://eor.berkeley.edu/}}, have been taking data for several years now, and we are at last starting to see these instruments place some upper-bounds on the 21-cm power-spectrum, e.g. \citealt{Paciga2011, Dillon2013, Ali2015, Pober2015, Beardsley2016c} and \citealt{Patil2017b}. There is also indication from the global experiment EDGES\footnote{The Experiment to Detect the Global EoR Signature \url{http://loco.lab.asu.edu/edges/}} (which is a single antenna experiment observing the mean evolution of the 21-cm signal, rather than attempting to constrain 21-cm fluctuations across the sky) that some form of coupling followed by heating is occurring in the redshift range $15<z<21$ \citep{Bowman2018a}. However, the inferred cosmological signal is far more extreme than expected, and exhibits an unexpected flat evolution over a large range of redshifts. If true, new physics beyond our standard models is required to explain this signal \citep{Bowman2018a}. Given then the challenging nature of the observation (strong foregrounds and ionospheric effects, both of which are observed with a beam that changes with frequency, must be mitigated), confirmation from an independent experiment is needed before we can be confident of the result. \citet{Hills2018} also find that the EDGES fit requires extremely unphysical foreground and ionospheric parameters, casting doubt on the EDGES result. It is therefore important that we do not put all our eggs in the exotic-physics basket and continue in parallel, as we do in this paper, to consider models consistent with our current fiducial astrophysical framework. The current generation of radio interferometers will not be able to observe the EoH over the EDGES redshift range (although it is still hoped that one or more may make a statistical detection of the EoR, and MWA could in principle provide statistical constraints of the EoH at $z<16$). It is expected that the next generation such as HERA\footnote{The Hydrogen Epoch of Reionization Array \url{http://reionization.org/}} and the SKA\footnote{The Square Kilometre Array \url{http://www.skatelescope.org/}} will allow us to observe the EoH. It has been seen from simulations that the signal will be highly non-Gaussian during both the EoH and the EoR \citep{Iliev2006, Mellema2006, Watkinson2014, Watkinson2015, Watkinson2015a, Shimabukuro2016a, Majumdar2017}. As such, it is important that we look to statistics other than the power spectrum, which can only fully describe a Gaussian field. This paper studies the bispectrum, which is sensitive to non-Gaussianities in a map, as measured from the fully numerical EoH simulations of \citet{Ross2016} and \citet{Ross2018}. We focus on their \textit{X-ray + Stellar} simulation, as low-redshift observations indicate that HMXBs are most likely to be the dominant X-ray source out of all the observed sources; we will refer to this simulation as \textit{HMXB} in the remains of the paper. We will also compare with simulations that include some level of contribution from X-rays generated by AGN (or QSO); throughout, we will refer to these as the \textit{HMXB + QSO} and \textit{QSO} simulations \citep{Ross2018}. In Section \ref{sec:sims} we review the numerical N-body + ray tracing simulations that we analyse here. In Section \ref{sec:interp} we discuss the interpretation of the bispectrum. In Section \ref{sec:scale} we define the \textit{normalised bispectrum}, a version of bispectrum, which has been normalised so as to remove the amplitude component. Note that we discuss other common normalisation options in Appendix \ref{app:normdicuss}. In Section \ref{sec:scale} we also present our findings that the normalised bispectrum from the \textit{HMXB} simulation exhibits a turn-over at high redshifts, the scale associated with which corresponds to the typical separation of emission regions. In Section \ref{sec:other_sims}, we will consider how consistent this qualitative evolution of the normalised bispectrum is across other simulations. We consider a totally different type of simulation by studying the normalised bispectrum from the semi-numerical simulation \cmfast as well as the \textit{HMXB + QSO} and \textit{QSO} simulations. We find that the qualitative evolution is the same for all but the \textit{QSO} simulation. This simulation differs in that its heated profiles are more distinct, driven by isolated sources and so imprint a second and stronger turnover corresponding to the typical size of heated regions. In Section \ref{sec:detect} we show that if foregrounds can be mitigated, the bispectrum should be detectable over the redshift range that the simulations we consider predict the EoH occurred. Finally, we conclude this work in Section \ref{sec:conc}. | \label{sec:conc} In this paper we have presented analysis of the 21-cm normalised bispectrum from fully-numerical simulations of the epoch of heating, assuming that the only source of X-rays is HMXBs. In the associated appendix we have also shown that our choice of bispectrum normalisation is the best option for analysing 21-cm data. We have found that if HMXB-like X-ray sources drive heating, then the equilateral bispectrum will be strongest in amplitude compared to other configurations and will exhibit a turnover that shifts from large to small scales with reducing redshift. We find that the scale at which this turnover peaks is correlated with the typical separation of emission regions. It is clear from our analysis that the bispectrum is driven by a complex interplay between the shape and size of heated profiles and their distribution. Cross-terms between the density field and spin temperature dominate at early times reflecting this complex interplay. As X-rays heat the cooler regions of the maps, small-scale sub-structure in the heated regions start to dominate the 21-cm bispectrum, introducing more power on smaller scales than on large. Ultimately, by the end of the simulation, fluctuations in the density field totally dominate the 21-cm bispectrum. We consider how generic the qualitative evolution of the bispectrum is by analysing two contrasting semi-numerical simulations. We observe very similar qualitative behaviour as in the numerical simulation in which HMXBs dominate the evolution. We also consider how the bispectrum is changed if QSOs are included into the numerical simulation, providing a second source of X-rays. At early times the presence of QSOs produces a stronger equilateral bispectrum, but still exhibits a turnover that shifts to smaller scales with decreasing redshift. By the mid phases of the heating process its normalised bispectrum is indistinguishable from that of the HMXB simulation. By analysing a third numerical simulation in which only QSOs provide X-ray radiation, we show that the bispectrum will look quite different than it would if HMXBs (or a similarly wide-spread source of X-rays) drive heating. At early times clustering of sources introduces a large-scale turnover feature. This drops in amplitude as the contrast between the most hot and the most cold regions decrease and is replaced by a turnover that is driven by the typical size of the heated profiles surrounding the heating sources. We consider the observability of the bispecrum with phase-1 of SKA-LOW and find that, assuming foregrounds and instrumental effects are effectively mitigated, we should be able to detect the bispectrum during the Epoch of Heating at $k<0.6$ Mpc$^{-1}$. Measuring the bispectrum should therefore provide a major boost to the information available from the power spectrum alone. Further work is required to get a better handle on the effect of sample variance and other complications to observing statistics such as the bispectrum; for example calibration and foreground removal residuals, and beam effects. | 18 | 8 | 1808.02372 |
1808 | 1808.01284_arXiv.txt | Observations of prestellar cores in star-forming filaments show two distinct morphologies. While molecular line measurements often show broad cores, submillimeter continuum observations predominantly display pinched cores compared to the bulk of the filament gas. In order to explain how different morphologies arise, we use the gravitational instability model where prestellar cores form by growing density perturbations. The radial extent at each position is set by the local line-mass. We show that the ratio of core radius to filament radius is determined by the initial line-mass of the filament. Additionally, the core morphology is independent of perturbation length scale and inclination, which makes it an ideal diagnostic for observations. Filaments with a line-mass of less than half its critical value should form broad cores, whereas filaments with more than half its critical line-mass value should form pinched cores. For filaments embedded in a constant background pressure, the dominant perturbation growth times significantly differ for low and high line-mass filaments. Therefore, we predict that only one population of cores is present if all filaments within a region begin with similar initial perturbations. | \label{sec:introduction} It has long been proposed that core formation in filaments is tied to some kind of fragmentation process \citep{schneider1979, larson1985}. This connection has only been reinforced by observations of the \textit{Herschel} Space Observatory \citep{andre2010, koenyves2010, menshchikov2010, ward-thompson2010, arzoumanian2011, arzoumanian2013, kirk2013, andre2014}, which show that dense cores are contained in an ubiquitous filamentary structure in molecular clouds. As cores are the birth-site of stars \citep{benson1989, klessen1998, mckee2007}, it is essential to understand the process of core formation in order to develop a coherent model for stellar formation. Different models of core formation have been proposed, e.g. by the dissipation of turbulence \citep{padoan2001, klessen2005} or by collapse of density enhancements due to intersecting filaments, so called "hubs" \citep{myers2009}. The complexity of core formation has increased with the observations of fibres \citep{hacar2013, tafalla2015}, trans- and subsonic velocity coherent substructures in filaments, again opening the possibility that cores form by subsonic motions due to gravitational instabilities, potentially modified by magnetic fields either hindering core formation due to magnetic pressure \citep{nagasawa1987, gehman1996b, fiege2000} or facilitating core formation in a magnetically stabilized filament by ambipolar diffusion \citep{shu1987, hosseinirad2017}.\\ A possible indicator to validate this model is the comparison of observed cores with the analytical predictions of overdensities forming by gravitational instabilities. High dynamic range observations in the submillimeter continuum, for instance in the Taurus region, show very thin cores compared to the filament radius \citep{marsh2014}. Contrarily, molecular line observations, which often only trace the dense gas, have mainly revealed cores which are broader than the filament \citep{hacar2011, hacar2013, tafalla2015}. Thus, the interpretation of core radius is complex and core morphology obviously depends on the tracer of observation.\\ Numerical predictions by \citet{nagasawa1987} showed that there are two regimes of the perturbation. One for low line-mass filaments, called deformation instability or "sausage" instability, where the forming cores bulge out and one for high line-mass, named compressional instability, where cores form by compression and thus pinch in. Both morphologies exist in simulations throughout the literature \citep{gehman1996a, gehman1996b, inutsuka1997, fiege2000}. However, in order to determine the morphology of cores it is important to not only predict the radius evolution of the core itself, but also the radius evolution of the material making up the rest of the filament. For a growing perturbation, both evolve simultaneously. We expand on the picture by \citet{nagasawa1987} and show an analytical prediction for the evolution of the radius ratio. | \label{sec:discussion} The gravitational instability model has several shortcomings. The main assumption is that filaments are very idealized cylindrical entities where the mean initial line-mass does not vary much along its length. Moreover, the filament profile requires a certain timescale to adjust to density changes. If the local line-mass varies faster than the radius can adjust, a broad core could be embedded in a filament with a line-mass larger than half the critical value. There are two processes which can lead to a major change in local line-mass on a short timescale. On the one hand, mass accretion increases the overall line-mass. Observed rates are estimated to be on the order of $10-100 \text{ M}_\odot \text{ pc}^{-1} \text{ Myr}^{-1}$ \citep{palmeirim2013}. On the other hand, a filament will longitudinally contract due to self-gravity. In addition, the rapid formation of two cores at the ends of the filament seems to be a typical outcome of the edge-effect \citep{burkert2004}. A different equation of state or additional physical contribute to the radial stability and can change the morphology of cores. Observed radial density profiles are better matched by polytropic indices lower than one \citep{toci2015}. As long as there is a maximum radius in dependence of the line-mass we still expect a dichotomy in morphology but with the division not necessarily at half the critical line-mass. Observationally, it is important to not only include the dense gas in order to reliably measure both filament and core radius. As the density of the outer filament gas is lower than the core gas, the filament radius has to be determined with a tracer of low gas density. If only the dense gas is observed, e.g. \NH, even cores which are nominally pinched can appear broader than the dense gas in the rest of in filament. Moreover, projection effects can reduce the length of a filament and thus increase the apparent line-mass by a substantial factor. This effect is limited by the fact that higher inclined filaments will not resemble a filamentary structure. Additionally, more cores are observed which are thinner than the average widths of star-forming filaments \citep{palmeirim2013, marsh2014, roy2014}, indicating that most filaments have high line-masses. Nevertheless, higher number statics on the local ratio of core-to-filament radius are desirable in order to estimate line-masses. \noindent All in all, our model allows for the following predictions: \begin{itemize} \item The morphology of cores embedded in filaments is set by the initial line-mass. Filaments with an initial line-mass below half the critical value will develop broad cores. Filaments with an initial line-mass above half the critical value will develop pinched cores. \item For filaments which are embedded in the same constant background pressure, the perturbation growth times for low and high line-masses are drastically different. If all filaments start with similar perturbation strengths we expect only one population of cores to be present, only pinched cores at early times and broad cores at late times. \item Using the FWHM to determine the radius underestimates the extent of high density regions of the filament and thus underestimates the ratio of core to filament radius. \item The phase where the radius of pinched cores is significantly different from the overall filament radius is very short and indicates an imminent collapse due to loss of hydrostatic equilibrium. \end{itemize} | 18 | 8 | 1808.01284 |
1808 | 1808.06147_arXiv.txt | We studied dynamical balances in magnetorotational instability (MRI) turbulence with net vertical field in the shearing box model of disks. Analyzing the turbulence dynamics in Fourier (${\bf k}$-)space, we identified three types of active modes that define turbulence characteristics. These modes have lengths similar to the box size, i.e., lie in the small wavenumber region in Fourier space labeled \emph{the vital area} and are: \emph{(i)} the channel mode -- uniform in the disk plane with the smallest vertical wavenumber, \emph{(ii)} the zonal flow mode -- azimuthally and vertically uniform with the smallest radial wavenumber and \emph{(iii)} \emph{the rest modes}. The rest modes comprise those harmonics in the vital area whose energies reach more than $50 \%$ of the maximum spectral energy. The rest modes individually are not so significant compared to the channel and zonal flow modes, however, the combined action of their multitude is dominant over these two modes. These three mode types are governed by interplay of the linear and nonlinear processes, leading to their interdependent dynamics. The linear processes consist in disk flow nonmodality-modified classical MRI with a net vertical field. The main nonlinear process is transfer of modes over wavevector angles in Fourier space -- \textit{the transverse cascade}. The channel mode exhibits episodic bursts supplied by linear MRI growth, while the nonlinear processes mostly oppose this, draining the channel energy and redistributing it to the rest modes. As for the zonal flow, it does not have a linear source and is fed by nonlinear interactions of the rest modes. | \label{sec:Introduction} Anisotropic turbulence offers a means of enhanced transport of angular momentum in astrophysical disks \citep{Shakura_Sunyaev73,Lynden-Bell_Pringle74}, whereas isotropic turbulence is unable to ensure such a transport. It is not surprising that from the 1980s, research in astrophysical disks focused on identifying sources of turbulence and understanding its statistical characteristics. The turning point was the beginning of the 1990s, when a linear instability mediated by a weak vertical magnetic field in differentially rotating conducting fluids \citep{Velikhov59,Chandrasekhar60} -- subsequently named as the magnetorotational instability (MRI) -- was rediscovered for sufficiently ionized astrophysical disks by \citet{Balbus_Hawley91}. MRI is a robust dynamical instability that leads to the exponential growth of axisymmetric perturbations, gives rise to and steadily supplies with energy magnetohydrodynamic (MHD) turbulence, as was demonstrated in earlier numerical simulations shortly after the significance of linear MRI in disks had been realized \citep[e.g.,][]{Hawley_Balbus91,Hawley_Balbus92,Hawley_etal95,Brandenburg_etal95,Balbus_Hawley98}. The generic nonnormality (non-self-adjointness) of a strongly sheared Keplerian flow of astrophysical disks provides an additional important linear mechanism of energy supply to the turbulence - transient, or nonmodal growth of perturbations \citep{Lominadze_etal88, Chagelishvili_etal03,Yecko04,Afshordi_etal05,Tevzadze_etal08, Shtemler_etal11,Salhi_etal12,Pessah_Chan12,Mamatsashvili_etal13, Zhuravlev_Razdoburdin14,Squire_Bhattacharjee14,Razdoburdin_Zhuravlev17}. Although the nonmodal growth is generally less powerful than the classical (exponentially growing) MRI, it is nevertheless capable of driving quite robust MHD turbulence when the latter is absent, for instance, in disks threaded by a purely azimuthal/toroidal magnetic field \citep[e.g.,][]{Hawley_etal95,Fromang_Nelson06, Simon_Hawley09,Guan_etal09,Guan_Gammie11,Flock_etal12a,Nauman_Blackman14, Meheut_etal15,Gogichaishvili_etal17}. In such spectrally, or modally stable magnetized disk flows lacking exponentially growing MRI modes, nonlinear processes are vital for turbulence sustenance. In this situation, the canonical -- direct and inverse -- cascade processes in classical Kolmogorov or Iroshnikov-Kraichnan phenomenologies are not capable of sustaining the turbulence -- this role is taken over by a new kind of the cascade process, so-called \emph{the nonlinear transverse cascade} \citep{Horton_etal10,Mamatsashvili_etal14}. What is the physical reason for the emergence of the transverse cascade in shear flows and its specific nature? The thing is that the anisotropy of the nonnormality/shear-induced linear nonmodal dynamics entails the anisotropy of the nonlinear processes -- transverse, or angular redistribution of perturbation harmonics in Fourier (wavenumber) space. In spectrally stable shear flows, in which the only mechanism for the energy supply to perturbations is the linear transient growth process, the nonlinear transverse cascade, by transferring the harmonics over wavevector angles (i.e., changing the orientation of their wavevectors), continually replenishes those areas in Fourier space where they can experience transient amplification \citep{Mamatsashvili_etal14,Mamatsashvili_etal16}. In this way, the transverse cascade guarantees a long-time sustenance of the turbulence. In the absence of such a feedback, non-axisymmetric modes that undergo in this case the most effective nonmodal transient growth, get eventually sheared away and decay. In the context of disks, this process was studied in detail in our recent paper \citet{Gogichaishvili_etal17} (Paper I) for the above-mentioned case of the Keplerian shear flow with a net azimuthal magnetic field by combining direct simulations of the turbulence and, based on them, the subsequent analysis of the dynamical processes in Fourier space. In this paper, we consider a Keplerian disk threaded by a net vertical/poloidal magnetic field, where there exists classical MRI with an exponential growth of axisymmetric modes \citep{Balbus_Hawley91,Balbus_Hawley98,Wardle99, Pessah_etal06,Lesur_Longaretti07,Pessah_Chan08,Longaretti_Lesur10, Latter_etal15,Shakura_Postnov15}. Because of this, there is no deficit in energy supply and the role of nonlinearity in the sustenance of perturbations is not as vital as in the case of azimuthal field. However, at the same time, MRI also owes its existence to the shear of disk flow and therefore is inevitably subject to nonmodal effects \citep{Mamatsashvili_etal13,Squire_Bhattacharjee14}. As demonstrated in this paper, this has two important consequences. First, the nonmodally-modified, or for short nonmodal MRI growth of both axisymmetric and non-axisymmetric modes during finite times are in fact more important in the energy supply process of turbulence, because the main time scales involved are of the order of dynamical/orbital time \citep{Walker_etal16}, than the modal (exponential) growth of axisymmetric modes prevalent at large times \citep[see also][]{Squire_Bhattacharjee14}. Second, underlying nonlinear processes are necessarily anisotropic in Fourier space as a result of the inherent anisotropy of linear nonmodal dynamics due to the shear. This anisotropy first of all gives rise to the nonlinear transverse cascade and one of the main goals of the present paper is to vividly demonstrate its importance in forming the overall dynamical picture of net vertical field MRI-turbulence. Recently, \citet{Murphy_Pessah15} investigated the saturation of MRI in disks with a net vertical field and the properties of the ensuing MHD turbulence both in physical and Fourier space. Their study was mainly devoted to characterizing the anisotropic nature of this turbulence. They also pointed out a general lack of analysis of the anisotropy in the existing literature on MRI-turbulence: ``Although there have been many studies of the linear phase of the MRI and its nonlinear evolution, only a fraction have explored the mechanism responsible for its saturation in detail, and none have focused explicitly on the evolution of the degree of anisotropy exhibited by the magnetized flow as it evolves from the linear regime of the instability to the ensuing turbulent state''. The main reason for overlooking the anisotropic nature of MRI-driven turbulence in most of previous works that focused on its spectral dynamics \citep[e.g.,][]{Fromang_Papaloizou07,Simon_etal09,Davis_etal10,Lesur_Longaretti11} was a somewhat misleading mathematical treatment, specifically, spherical shell-averaging procedure in Fourier space \citep[borrowed from forced MHD turbulence studies without shear flow, see e.g.,][]{Verma04,Alexakis_etal07}, which had been employed to extract statistical information about the properties of MRI-turbulence. Obviously, the use of the shell-averaging technique, which in fact smears out the transverse cascade, is, strictly speaking, justified for isotropic turbulence, but by no means for shear flow turbulence and therefore for, its special case, MRI-turbulence ``nourished'' in a sheared environment of disk flow. Thus, the shell-averaging is not an optimal tool for analyzing spectra as well as dynamical processes in Fourier space that underlie MRI-driven turbulence and are far from isotropic due to the shear \citep[see also][Paper I]{Hawley_etal95,Nauman_Blackman14,Lesur_Longaretti11,Murphy_Pessah15}. This fact also calls into question those investigations based on the shell-averaging that aim to identify a power-law character and associated slopes of turbulent energy spectrum, because the anisotropy of the energy spectrum itself, a direct consequence of the transverse cascade, is wiped out in these cases. This is perhaps the reason why the kinetic and magnetic energy spectra do not generally exhibit a well-defined power-law behavior in MRI-turbulence in disk flows with a net vertical field \citep{Simon_etal09,Lesur_Longaretti11,Meheut_etal15,Walker_etal16}. In view of the above, in this work we focus on the dynamics and balances in MRI-turbulence with a net vertical field. We adopt the local shearing box model of the disk with constant vertical thermal stratification. The analysis is performed in three-dimensional (3D) Fourier space in full, i.e., without doing the averaging of spectral quantities over spherical shells of given wavenumber magnitude $k = |{\bf k}|$. This allows us to capture the spectral anisotropy of the MRI-turbulence due to the shear and the resulting nonlinear angular redistribution of perturbation modes in Fourier space, i.e., transverse cascade, thereby getting a deeper understanding of spectral and statistical properties of the turbulence. In previous relevant studies on a net vertical field MRI \citep[e.g.,][]{Goodman_Xu94,Hawley_etal95,Sano_Inutsuka01,Lesur_Longaretti07, Bodo_etal08,Latter_etal09,Latter_etal10,Simon_etal09,Pessah_Goodman09, Longaretti_Lesur10,Pessah10,Bai_Stone14,Murphy_Pessah15}, this redistribution (scatter) of modes over wavevector orientations (angles) in Fourier space was attributed to secondary, or \emph{parasitic} instabilities. Specifically, the most unstable, exponentially growing axisymmetric MRI modes (channels) are subject to secondary instabilities of non-axisymmetric modes with growth rates proportional to the amplitude of these channel solutions. In this way, the parasitic instabilities redistribute the energy from the primary axisymmetric channel modes to non-axisymmetric parasitic ones, halting the exponential growth of the former and leading to the saturation of MRI. Several approximations are made in this description: 1. the large amplitude channel mode is a time-independent background on which the small-amplitude parasitic modes feed and 2. the effects of the imposed vertical field, the Coriolis force, and the basic Keplerian shear are all usually neglected. These assumptions clearly simplify the analysis of the excitation and dynamics of the non-axisymmetric parasitic modes, but, more importantly, because of neglecting the basic flow shear, omit independent from the primary MRI modes source of their support - the nonmodal growth of non-axisymmetric modes, which, as discussed above, is an inevitable linear process in shear (disk) flows. To keep an analysis general and self-consistent, together with the nonmodal growth process, we employ the concept of the nonlinear transverse cascade in the present problem of MRI with a net vertical magnetic field that naturally encompasses the secondary instabilities too. This unifying framework enables us to correctly describe the interaction between the channel and non-axisymmetric ``parasitic'' modes when the above assumptions break down -- the amplitudes of these two mode types become comparable, so that it is no longer possible to clearly distinguish between the channel as a primary background and parasites as small perturbations on top of that. This is the case in the developed turbulent state of vertical field MRI, where channels undergo recurrent amplifications (bursts) and decays \citep[e.g.,][]{Sano_Inutsuka01,Lesur_Longaretti07,Bodo_etal08,Simon_etal09,Murphy_Pessah15}. This decay phase is usually attributed to the linear non-axisymmetric parasitic instabilities, however, in the fully developed turbulent state it is more likely governed by nonlinearity. As is shown in this paper, the transverse cascade accounts for the transfer of energy from the axisymmetric channel modes to a broad spectrum of non-axisymmetric ones (referred to as the rest modes here) as well as to the axisymmetric zonal flow mode, which appears to commonly accompany MRI-turbulence \citep{Johansen_etal09,Simon_etal12,Bai_Stone14}. The nonlinear transverse cascade may not be the vital source of the energy supply to turbulence in the presence of purely exponentially growing MRI, but still it shapes the dynamics, sets the saturation level and determines the overall ``design'' of the turbulence the ``building blocks'' of which are these three types of perturbation modes analyzed in this paper. In the spirit of our recent works \citep[][Paper I]{Horton_etal10, Mamatsashvili_etal14,Mamatsashvili_etal16}, here we investigate in detail the roles of underlying different anisotropic linear and nonlinear dynamical processes shaping the net vertical field MRI-turbulence. Namely, we first perform numerical simulations of the turbulence and then, using the simulation data, explicitly calculate individual linear and nonlinear terms and explore their action in Fourier space. The underlying physics, active modes and dynamical balances of the net vertical field MRI-turbulence in disks are, however, entirely different from those of the net azimuthal field one studied in Paper I, primarily because energy sources for the turbulence in these two field configurations differ in essence: in the first case, the turbulence is mostly supplied by the (nonmodally-modified) exponentially growing MRI, while in the second case just by transient growth of non-axisymmetric modes. The present study can be regarded as a generalization of the related works by \citet{Simon_etal09,Lesur_Longaretti11}, where the dynamics of vertical field MRI-turbulence -- the spectra of energy, injection and nonlinear transfers -- were analyzed in Fourier space, however, using a restrictive approach of shell-averaging, which misses out the shear-induced anisotropy of the turbulence and hence the interaction of axisymmetric channel and non-axisymmetric modes. It also extends the study of \citet{Murphy_Pessah15}, who focused on the anisotropy of vertical field MRI-turbulence in both physical and Fourier space and examined, in particular, anisotropic spectra of the magnetic energy and Maxwell stress during the linear growth stage of the channel solutions and after saturation, but not the action of the various linear and nonlinear terms governing their evolution. The paper is organized as follows. The physical model and main equations in Fourier space is given in Section \ref{sec:Basicequations}. The linear nonmodal growth of MRI is analyzed in Section \ref{sec:Optimal}. Simulations of the ensuing MRI-turbulence and its general characteristics, such as time-development, energy spectra and the classification of dynamically active modes are given in Section \ref{sec:DNS}. In this section we also give the main analysis of the individual dynamics of the active modes and their interdependence in Fourier space that underlie the dynamics of the turbulence. Summary and discussions are given in Section \ref{sec:Conclusion}. | \label{sec:Conclusion} In this paper, we investigated the dynamical balances underlying MRI-driven turbulence in Keplerian disks with a nonzero net vertical magnetic field and vertically uniform thermal stratification using shearing box simulations. Focusing on the analysis of the turbulence dynamics in Fourier (${\bf k}$-) space, we identified three key types of modes -- \emph{the channel, the zonal flow mode and the rest modes} -- that are the main ``players'' in the turbulence dynamics. We described the dynamics of these modes separately and then their interdependence, which sets the properties of the nonzero net vertical field MRI-turbulence. The processes of linear origin are defined primarily by nonmodal, rather than modal, growth of MRI due to disk flow nonnormality/shear. This is because the dynamical time of the turbulence is of the order of the orbital/shear time during which the nonmodal effects are important. In the turbulent state, higher values of the stresses and magnetic energy fall just on those active modes that exhibit high nonmodal MRI growth and \emph{not} on the modally (exponentially) most unstable modes. In other words, the properties of the turbulence are determined mostly by nonmodal physics of MRI rather than by the modal one. From all the active modes, the one that exhibits the maximum nonmodal growth is the channel mode, which is horizontally uniform with the largest vertical scale in the domain. As for the nonlinear processes, it can be confidently stated that the decisive agent in forming and maintaining the statistical characteristics of the net vertical field MRI-turbulence is the transverse cascade -- nonlinear redistribution of modes in Fourier space that changes the orientation (angle) of their wavevectors -- arising from the presence of the shear and hence being a generic phenomenon in shear flows. Specifically, the nonlinear transverse cascade redistributes the energy of the channel mode to the rest modes and, subsequently, the energy of the rest modes to the zonal flow mode. (One has to note that the rest modes receive energy not only due to the nonlinear transfers, but they themselves also undergo nonmodal transient growth, however, less than the channel mode does). The combined action of these linear and nonlinear processes leads to the channel mode exhibiting recurrent bursts of the energy and stresses. The nonlinear transfer of its energy to the rest modes causes the decline of the channel mode after each burst and subsequent abrupt increase of the energy of the rest modes. This, in turn, induces similar, burst-like evolution of the integral characteristics of the turbulence -- the total volume-averaged energies, stresses and transport. The rest modes here, playing the main role in draining the channel mode, were referred to as parasitic modes in previous studies of net vertical field MRI. However, there is an important distinction. These parasitic modes are often assumed to be small compared with the channel mode and are treated as linear perturbations imposed on the latter \citep{Goodman_Xu94,Pessah_Goodman09,Latter_etal09,Latter_etal10,Pessah10}. Besides, the effect of shear and hence the nonmodal (transient) physics are neglected with respect to parasitic modes. In the turbulent state, however, the rest modes can reach energies comparable to the channel mode, as has been demonstrated in this paper, so one can no longer separate the channel as a primary background and parasites as small perturbations on top of that. As a result, the complex interaction between these two mode types belongs to the domain of nonlinearity. Our general/unifying approach -- the analysis of the dynamics in 3D Fourier space -- allows us to self-consistently characterize this interaction of the modes. In this case, the transverse cascade defines the spectrum of the rest modes, which contribute to the decline of the channel mode after each burst and subsequently acquire energy in this process. As we found in this study, the net vertical field MRI-turbulence is robust and, in addition, multifarious -- determined by the interdependent/interlaced dynamics of three qualitatively different modes. Consequently, in order to properly quantify the relative contribution of each of these modes in the turbulence characteristics, one has to capture the main aspects of their dynamics in numerical simulations. First of all, this concerns the selection of relevant sizes (aspect ratio) of the simulation box (which is actually arbitrary in the shearing box framework) and resolutions, so that the discrete modes in the selected box densely enough cover the vital area in $(k_x,k_y)$-plane and maximally comprise effectively growing (optimal) modes (see Figure \ref{fig:optimalgrowth}). Besides, one should also avoid artificial/numerical anisotropyzation of nonlinear processes. As was shown in Paper I, the anisotropy of the simulation box in $(k_x,k_y)$-plane introduces artificial anisotropy of nonlinear processes and somewhat ``deforms'' the overall dynamical picture of MRI-turbulence in Fourier space. In the present case, this artificial deformation could result in a change of the relative importance of the above-classified modes in the overall dynamics, for instance, could reduce the effectiveness of the channel mode compared to other modes and hence weaken its manifestation in the turbulence dynamics \citep[see e.g.,][]{Bodo_etal08}. In particular, changing the role of the large-scale channel mode likely affects the generation of the mean azimuthal magnetic field, or the dynamo action, since this field is directly associated with this mode. Susceptibility of the latter to specific factors of the dynamics is clearly shown in Figure \ref{fig:Bya_zt_strat_nostrat}: although the vertical stratification makes only negligible contribution to the turbulence energy, it remarkably enhances the generation of the mean azimuthal magnetic field and regularizes its spatio-temporal variation. The latter process is significant and represents the subject of a special detailed investigation. \subsection{On types of anisotropy} Apart from the analyzed in our paper shear-induced anisotropy of the nonlinear processes, a net nonzero background magnetic field itself, by definition, gives rise to anisotropy of the nonlinear dynamics to its parallel and perpendicular directions. However, these two types of anisotropy, having different origin, essentially differ from each other. Anisotropy of nonlinear cascade processes due to magnetic field (directed along $z$-axis) in plasmas with static equilibrium is analyzed in \cite{Goldreich_Sridhar95}. The source of energy is assumed to be isotropic in both physical and Fourier spaces with some outer scale of turbulence, $L$, corresponding to the smallest wavenumber, $k_0\sim 1/L$. Nonlinear cascade processes take place in the inertial range, at perpendicular and parallel to the field wavenumbers $k_{\perp}, k_z\gtrsim k_0$, which is in fact the region of activity of these processes only. Overall, the regions of the energy supply and nonlinear cascade in Fourier space are separated from each other. As a result, the nonlinear processes do not affect the turbulence energy supply -- they only transfer energy mainly to small perpendicular (to the magnetic field) wavelength, i.e., to large $k_{\perp}$ rather than along $k_z$, thereby forming anisotropic spectrum of the turbulence in $(k_{\perp}, k_z)$-plane. This nonlinear cascade mostly to large $k_{\perp }$, besides increasing the magnitude $k=\sqrt{k_{\perp}^2 + k_z^2}$ of the total wavevector, ${\bf k}=(k_{\perp},k_z)$ (direct cascade), by definition, implies also a change of its orientation that results in some angular redistribution of harmonics (i.e., transverse cascade). At $k_{\perp}/k_z \sim 1$, the direct and transverse cascades are comparable. As the cascade proceeds, the ratio $k_{\perp }/ k_z$ increases, leading ${\bf k}$ to change primarily in magnitude rather than in orientation and hence the direct cascade becomes dominant. In any case, the angular redistribution (transverse cascade) in this case is of secondary importance -- does not affect the turbulence energy supply. A completely different situation arises in the presence of the disk flow shear. Linear (nonmodal) energy supply processes of perturbations/turbulence in shear flows are strongly anisotropic in Fourier space and generally occur over a broad range of wave numbers without leaving a free room (i.e., inertial range) for the action of nonlinear processes only. Due to the inherent anisotropy of the linear dynamics due to the shear, the nonlinear processes in this range of wavenumbers become strongly anisotropic, i.e., the transverse cascade is the dominant nonlinear process. Overall, dynamical picture of the turbulence is formed as a result of the interplay of the linear and nonlinear (transverse cascade) processes. The area of the most intensive interplay of the linear and nonlinear processes we call the vital area of the turbulence. Of course, the nonlinear processes leads to energy exchange between different modes, redistributing perturbation energy in Fourier space while leaving the total energy unchanged. However, the nonlinear transverse cascade in shear flows repopulates the (transiently) growing harmonics and, in this way, indirectly contributes to the energy supply to turbulence. As we have shown above, in the present problem of net vertical field MRI-turbulence in Keplerian shear flow, it significantly affects the dynamical ``design'' of the turbulence and determines the interplay of its ``building blocks'' -- channel, zonal flow and the rest modes. This is clearly illustrated by Figures \ref{fig:nonlinearterms_peak}-\ref{fig:nonlinearterms_average}, showing resulting anisotropic dynamics (transfers) in $(k_x,k_y)$-plane perpendicular to the field. By contrast, in the absence of basic flow velocity shear, nonlinear cascades in this plane are isotropic in Goldreich \& Sridhar theory. | 18 | 8 | 1808.06147 |
1808 | 1808.04232_arXiv.txt | We present visible-light and ultraviolet ($UV$) observations of the supernova PTF\,12glz. The SN was discovered and monitored in near-$UV$ and R bands as part of a joint \textit{GALEX} and Palomar Transient Factory campaign. It is among the most energetic Type IIn supernovae observed to date ($\approx10^{51}$ erg). If the radiated energy mainly came from the thermalization of the shock kinetic energy, we show that PTF\,12glz was surrounded by $\sim1\,\rm{M}_{\odot}$ of circumstellar material (CSM) prior to its explosive death. PTF\,12glz shows a puzzling peculiarity: at early times, while the freely expanding ejecta are presumably masked by the optically thick CSM, the radius of the blackbody that best fits the observations grows at $\approx8000$\,km\,s$^{-1}$. Such a velocity is characteristic of fast moving ejecta rather than optically thick CSM. This phase of radial expansion takes place before any spectroscopic signature of expanding ejecta appears in the spectrum and while both the spectroscopic data and the bolometric luminosity seem to indicate that the CSM is optically thick. We propose a geometrical solution to this puzzle, involving an aspherical structure of the CSM around PTF\,12glz. By modeling radiative diffusion through a slab of CSM, we show that an aspherical geometry of the CSM can result in a growing effective radius. This simple model also allows us to recover the decreasing blackbody temperature of PTF\,12glz. \texttt{SLAB-Diffusion}, the code we wrote to model the radiative diffusion of photons through a slab of CSM and evaluate the observed radius and temperature, is made available on-line. | Type IIn supernovae (SNe) are characterized by prominent and narrow-to-intermediate width Balmer emission lines in their spectra \citep{Schlegel1990, Filippenko1997, Smith2014,Gal-Yam2016}. Rather than a signature of the explosion itself, this spectral specificity is presumably the result of the photoionization of a dense, Hydrogen-rich, circumstellar medium (CSM) which is ejected from the SN progenitor prior to the explosion. The Type IIn class is not a well-defined category of objects, as many SNe show the characteristic narrow Balmer lines in their spectra, sometime during their evolution. These lines are the signature of an external physical phenomenon highly dependent on the surrounding environment, rather than of any intrinsic property of the explosion. Depending on the spatial distribution and physical properties of the CSM, these lines may persist for days (``flash spectroscopy'', \citealt{Gal-Yam2014,Khazov2016,Yaron2017}), weeks (e.g., SN\,1998s, \citealt{Li1998, Fassia2000, Fassia2001}; SN\,2005gl, \citealt{Gal-Yam2007}; SN\,2010mc, \citealt{Ofek2013}), or years (e.g., SN\,1988Z, \citealt{Danziger1991,Stathakis1991,Turatto1993,VanDyk1993,Chugai1994,Fabian1996,Aretxaga1999,Williams2002,Schlegel2006,Smith2017}; 2010\,jl, \citealt{Patat2011, Stoll2011, Gall2014,Ofek2014}). In the last decades, the physical picture governing SN IIn explosions and the wider family of ``interacting'' SNe - SNe whose radiation can be partially or completely accounted for by the ejecta crashing into a dense surrounding medium - has become clearer (see e.g., \citealt{Chevalier1982}, \citealt{Chugai1994}, \citealt{Chugai2004}, \citealt{Ofek2010}, \citealt{Chevalier2011}, \citealt{Ginzburg2014}, \citealt{Moriya2014}). In recent years, there is growing evidence that, in the majority of cases, the high-density CSM originates from explosive phenomena taking place in the months to years prior to the SN explosion. One piece of evidence supporting this conclusion is the direct detection of the so-called precursors (luminous outbursts) in the months to years prior to the SN explosion (e.g., \citealt{Foley2007, Pastorello2007, Fraser2013, Ofek2013, Ofek2014c, Ofek2016, Elias-Rosa2016, Thone2017}). Several theoretical mechanisms have been suggested to explain extreme mass-loss episodes in the final stages of stellar evolution (e.g., \citealt{Woosley2007,Quataert2012,Chevalier2012b,Soker2016}). While in normal core-collapse SNe, the radiation-mediated shock breaks out upon reaching the stellar surface, producing a strong blast in the $UV$ and X-rays \citep{Nakar2010,Rabinak2011}, in the case of SNe IIn the ejecta may crash into the optically thick CSM. The radiation-dominated and radiation-mediated shock runs into the CSM surrounding the star and goes on propagating into it as long as $\tau \gtrapprox c/v_{sh}$, where $\tau$ is the optical depth from the shock to the edge of the wind, $v_{sh}$ is the shock velocity, and $c$ is the speed of light (e.g., \citealt{Ofek2010}). When $\tau \sim c/v_{sh}$, (this condition is verified when the timescale for photons to diffuse from the shocked region to the photosphere becomes comparable to the dynamical timescale of the shock), the shock breaks out: photons diffuse ahead of the shock faster than the ejecta and radiation can escape ahead of the shock \citep{Weaver1976}. After the shock breakout, in the presence of massive CSM above the shock, the radiation-dominated shock transforms into a collisionless shock \citep{Katz2011,Murase2011,Murase2014}. The collisionless shock slows down the ejecta and converts its kinetic energy into hard X-ray photons \citep{Katz2011,Murase2011,Murase2014}. If the optical depth of the CSM above the shock is high enough, the X-rays generated in the collisionless shock are converted into $UV$ and visible radiation (e.g., \citealt{Chevalier2012, Svirski2012}). Without a sufficient optical depth though, the bulk of the X-ray photons will not convert into optical photons. As far as a spectral signature is concerned, the common picture explaining SNe IIn observations is as follows. As long as the CSM is optically thick, the photosphere which emits the continuum is located in the unshocked CSM, masking the observer's view of the shock. The radiation from the shock propagates upstream and photoionizes the slowly moving CSM, resulting in relatively narrow Balmer recombination emission lines in the SN spectrum. As the shock reaches the optically thin medium, broader components can appear in the spectrum - maybe arising from the shocked zone forming at the contact discontinuity between the decelerated ejecta and the shocked CSM \citep{Chugai2004}. Alternatively, if the CSM is optically thin, the lines may be generated in inner regions (e.g., \citealt{Chevalier1994}) Observing SNe IIn at wavelengths where the collisionless shock radiates most - namely $UV$ and X-rays - has the potential to unveil precious information about the explosion mechanism and the CSM properties (e.g., \citealt{Ofek2013a}). In particular, it may provide a much better estimate of the bolometric luminosity of the event. In this paper, we present and analyse the $UV$ and visible-light observations of PTF\,12glz, a SN IIn observed in a joint campaign by \textit{GALEX} and the Palomar Transient Factory (PTF) and detected in the $UV$. PTF\,12glz is one of the six SNe discovered during this campaign \citep{Ganot2016}. The survey was carried out as a proof-of-concept for the $ULTRASAT$ mission \citep{Sagiv2014}. Observations of SNe IIn are usually analyzed within the framework of spherically symmetric models of CSM. However, resolved images of stars undergoing considerable mass loss (e.g., $\eta$ Carinae; \citealt{Davidson1997, Davidson2012}), as well as polarimetry observations \citep{Leonard2000, Hoffman2008, Wang2008, Reilly2017} suggest that asphericity should be taken into account for more realistic modeling. Asphericity of the CSM has recently been invoked to interpret the spectrocopic and spectropolarimetric observations of the Type IIn SN SN2012ab \citep{Bilinski2017}. In this paper, we show that the light curve of PTF\,12glz may be interpreted as evidence for aspherical CSM. We present the aforementioned observations of PTF\,12glz in \S 2. In \S 3, we present the analysis of these observations and the puzzling inconsistency between the spectroscopic and photometric observations. In \S 4, we model the radiative diffusion of photons through an aspherical slab and propose a solution to this puzzle. We then summarize our main results in \S 5. In the Appendix, we make available \texttt{SLAB-Diffusion}, a computer code for modeling radiation through a slab of CSM. | \label{sec:discussion} We presented the observations of the supernova PTF\,12glz by the \textit{GALEX} space telescope and ground-based PTF. Radioactive decay is not sufficient to explain the decay of the light curve of PTF\,12glz and therefore other physical mechanisms must be involved. One possible - yet difficult to verify - scenario is that an internal engine powers the light curve. Another possible scenario - the standard explanation invoked in the case of Type IIn SNe, is that the light curve is powered by interaction between the ejecta and the CSM surrounding the SN. In the case of PTF\,12glz, the spectroscopic analysis is consistent with the following picture: at early times (two first spectra) both the ejecta and the shock are initially masked by a thick, slowly moving, photoionized CSM. At later times (two last spectra), the ejecta have emerged through - at least some of - the optically thick layers and have reached CSM layers that are optically thin enough to expose the ejecta. CSM interaction may still play a role at late times, e.g., by heating the ejecta from the inside, and contributes to slowing the light curve decay. The evolution of $r_{BB}$ - the radius of the deepest transparent emitting layer - seems to contradict this picture. At early times, i.e., at the very time when the opaque CSM seemingly obstructs our view of any growing structure, $r_{BB}$ grows by an order of magnitude, at a speed of $\sim8000$\,km\,s$^{-1}$. In addition to being inconsistent with the spectroscopic analysis, this is also in contradiction - to our knowledge - with all previous observations of either a constant or stalling blackbody radius in SNe IIn (as detailed in \S~\ref{sec:peculiar_r}). If the bulk of the radiation from PTF\,12glz does come from interaction, the explanation for the growing blackbody radius may be geometrical. The question then is whether any peculiar structure of the CSM around the progenitor can reproduce the observations. In this work, we considered a simple aspherical structure of CSM: a slab. We modeled the radiation from an explosion embedded in a slab of CSM by numerically solving the radiative diffusion equation in a slab with different density profiles: $\rho=Const.$, $\rho\propto |z|^{-1}$ and a wind density profile $\rho\propto z^{-2}$. Although this model is simplistic, it allows recovery of the peculiar growth of the blackbody radius $r_{BB}$ observed in the case of PTF\,12glz, as well as the decrease of its blackbody temperature $T_{BB}$. This configuration is not a unique geometrical solution and additional observations, e.g., of the polarization around PTF\,12glz would have been necessary to make it less speculative. As new wide-field transient surveys such as the Zwicky Transient Facility (e.g., \citealt{Bellm2015,Laher2017}) are deployed, many more interacting SNe will be observed and quickly followed up with multiple-band observations. These may also be the brightest sources for the {\it ULTRASAT} $UV$ satellite mission \citep{Sagiv2014}. Some of these interacting SNe may exhibit the same peculiarities as PTF\,12glz. The methodology proposed in this paper offers a framework to analyze them. It could be elaborated upon, to model more complex aspherical geometries, e.g., $\eta$ Carinae-like shapes of the CSM, and give more quantitative predictions of the observables. | 18 | 8 | 1808.04232 |
1808 | 1808.10037_arXiv.txt | The Simons Observatory (SO) will make precision temperature and polarization measurements of the cosmic microwave background (CMB) using a series of telescopes which will cover angular scales between one arcminute and tens of degrees and sample frequencies between 27 and 270 GHz. Here we present the current design of the large aperture telescope receiver (LATR), a 2.4\,m diameter cryostat that will be mounted on the SO 6\,m telescope and will be the largest CMB receiver to date. The cryostat size was chosen to take advantage of the large focal plane area having high Strehl ratios, which is inherent to the Cross-Dragone telescope design. The LATR will be able to accommodate thirteen optics tubes, each having a 36\,cm diameter aperture and illuminating several thousand transition-edge sensor (TES) bolometers. This set of equipment will provide an opportunity to make measurements with unparalleled sensitivity. However, the size and complexity of the LATR also pose numerous technical challenges. In the following paper, we present the design of the LATR and include how we address these challenges. The solutions we develop in the process of designing the LATR will be informative for the general CMB community, and for future CMB experiments like CMB-S4. | \label{sec:intro} % The cosmic microwave background (CMB) has become one of the most powerful probes of the early universe. Measurements of temperature anisotropies on the level of ten parts per million have brought cosmology into a precision era, and have placed tight constraints on the fundamental properties of the universe\cite{Samtleben2007}. Beyond temperature anisotropies, CMB polarization anisotropies not only enrich our understanding of our cosmological model, but could potentially provide clues to the very beginning of the universe via the detection (or non-detection) of primordial gravitational waves. A number of experiments have made and are continuing to refine measurements of the polarization anisotropy. However, these experiments are typically dedicated to a relatively restricted range of angular scales, e.g., large angular scales (tens of degrees) or high resolution/small angular scales (on the order of one arcminute). To provide a complete picture of cosmology, both large and small angular scales are important. Ideally these measurements would be made from the same observing site so that the widest range of angular scales can be probed, at multiple frequencies, on the same regions of the sky. This is the goal of the Simons Observatory (SO). \begin{figure}[H] \begin{center} \begin{tabular}{c} \includegraphics[width = 0.8\linewidth]{telescope.PNG} \end{tabular} \end{center} \caption{A cross-section showing how the LATR couples to the SO 6\,m telescope. In the orientation shown, light enters the telescope from the top. It then reflects off the 6\,m primary (M1) and 6\,m secondary (M2) before being directed to the LATR receiver. The receiver always operates in the horizontal orientation shown, and is not coupled to the telescope elevation structure. Therefore, a separate mechanism is used to rotate the receiver about its long axis as the telescope elevation structure moves M1 and M2 in rotation.} \label{fig:telescope} \end{figure} SO will field a 6\,m Cross-Dragone large aperture telescope (LAT) coupled to the large aperture telescope receiver (LATR, as shown in Fig.~\ref{fig:telescope}). The LAT is designed to have a large FOV\cite{Niemack2016,Parshley2018} capable of supporting a cryostat with up to 19 LATR-like optics tubes. To reduce the development risk of such large cryostat, the LATR is designed to accommodate only up to 13 optics tubes. During the initial deployment, we plan to deploy 7 optics tubes with 3 detector wafers in each for a total of approximately 35,000 detectors, primarily at 90/150\,GHz. It should be noted that each optics tube could be upgraded to support 4.5 wafers for a $\sim$50\% increase in the number of detectors per optics tube. With this upgrade and the deployment of 19 optics tubes, the LAT could support roughly 145,000 detectors at 90/150 GHz. SO will also deploy an array of half-meter small aperture telescopes (SATs) coupled to an additional 30,000 detectors\cite{Galitzki2018}. The unique combination of telescopes in a single CMB observatory, which will be located in Chile\textsc{\char13}s Atacama Desert at an altitude of 5190\,m, will allow us to sample a wide range of angular scales over a common survey area. In this paper, we present the current design of the LATR. In Sec.~\ref{sec:mech}, we will present the mechanical design of each temperature stage and the design philosophy behind each element. In Sec.~\ref{sec:cryo_design}, we will discuss a few key cryogenic components, including cryo-coolers, heat pipes, heat switches, thermometers, and infrared (IR) blocking filters\cite{Tucker2006,Ade2006} that are being considered for the LATR. We will also cover simulation and testing plans for these components. Finally, in Sec.~\ref{sec:detector}, we will briefly discuss the detectors and readout strategy. The challenges we faced when designing the LATR are unique in that we are building the largest and most complex CMB receiver to date. The solutions to these challenges will provide a critical stepping stone for the next generation CMB experiments, in particular, CMB-S4\cite{Abazajian2016,Abitbol2017}. | We have completed the major component design of the 2.4\,m receiver that will operate on the 6\,m millimeter-wave Simons Observatory LAT. The LATR will contain approximately 35,000 detectors (in the first seven of the planned thirteen optics tubes) operating at 100\,mK and sensitive to six frequency bands between 27 and 270\,GHz. With this set of equipment, the LATR will measure fundamental cosmological parameters of our universe with unparalleled sensitivity, find high redshift clusters through the Sunyaev-Zel'dovich effect, constrain properties of neutrinos, and seek signatures of dark matter through gravitational lensing. Manufacture of this exciting, yet challenging, instrument is expected to commence this calendar year. | 18 | 8 | 1808.10037 |
1808 | 1808.09476_arXiv.txt | We present detailed multi-wavelength observations of GRB\,161219B at $z=0.1475$, spanning the radio to X-ray regimes, and the first ALMA light curve of a GRB afterglow. The cm- and mm-band observations before $8.5$\,d require emission in excess of that produced by the afterglow forward shock (FS). These data are consistent with radiation from a refreshed reverse shock (RS) produced by the injection of energy into the FS, signatures of which are also present in the X-ray and optical light curves. We infer a constant-density circumburst environment with an extremely low density, $\dens\approx 3\times10^{-4}$\,\pcc, and show that this is a characteristic of all strong RS detections to date. The VLA observations exhibit unexpected rapid variability on $\sim$ minute timescales, indicative of strong interstellar scintillation. The X-ray, ALMA, and VLA observations together constrain the jet break time, $\tjet\approx32$\,d, yielding a wide jet opening angle of $\thetajet\approx13^{\circ}$, implying beaming corrected $\gamma$-ray and kinetic energies of $\Egamma\approx4.9\times10^{48}$\,erg and $\EK\approx1.3\times10^{50}$\,erg, respectively. Comparing the RS and FS emission, we show that the ejecta are only weakly magnetized, with relative magnetization, $\RB\approx1$, compared to the FS. These direct, multi-frequency measurements of a refreshed RS spanning the optical to radio bands highlight the impact of radio and millimeter data in probing the production and nature of GRB jets. | Long-duration $\gamma$-ray bursts (GRBs) have thus far been almost exclusively discovered through their prompt $\gamma$-ray emission, which unequivocally arises from relativistic outflows at high Lorentz factors, $\Gamma\gtrsim10^2$ \citep{kp91,feh93,wl95,bh95,bh97,ls01}. These outflows are understood to be produced by a nascent, compact central engine, such as a magnetar or accreting black hole, formed in the collapsing core of a dying massive star \citep{wb06,pir05,mgt+11}. The internal shock model proposed to explain the $\gamma$-ray emission invokes collisions between shells with a wide distribution of Lorentz factors ejected by the engine \citep{rm92,kps97,kp00a}. Understanding the distribution of ejecta energy as a function of their Lorentz factor is therefore a critical probe of the nature of the central engine, its energy source, and the energy extraction mechanism \citep{woo93,mw99,ami+00,npk01,zwm03,tmn08}. While monitoring the $\gamma$-ray sky remains an excellent means for detecting GRBs, a detailed description of the energetics of their jets and their progenitor environments is only possible through a study of the long-lasting X-ray to radio afterglow, generated when ejecta interact with their circumburst environment setting up the forward shock, and producing synchrotron radiation \citep[FS;][]{spn98}. Theoretical modeling of detailed multi-wavelength observations in the synchrotron framework yields the energy of the explosion, the degree of jet collimation, the density of the surrounding medium, and the mass loss history of the progenitor star, as well as information about the microphysical processes responsible for relativistic particle acceleration \cite{sph99,cl00,gs02}. Whereas GRB afterglows have traditionally been modeled as arising from jets with a uniform bulk Lorentz factor, radially structured ejecta profiles with energy spanning a range of Lorentz factors are gaining traction as viable models for the observed deviations of X-ray and optical light curves from the synchrotron model\footnote{Alternate explanations include circumburst density enhancements, structured jets, viewing angle effects, varying microphysical parameters, and gravitational microlensing \citep{zfd+06,nkg+06,pmg+06,tiyn06,eg06,gkp06,jyfw07,sd07,kwhc10,dm15,uz14}. } \citep{np03,bhpf02,bgj04,hcg06,jbg06,mmk+08,mgk+09,tsg+12,vmp+13}. Ejecta released later, or at lower Lorentz factors than the initial impulsive shell responsible for the prompt emission, catch up with the contact discontinuity during the afterglow phase and inject energy into the FS \citep{rm98,sm00}. Energy injection through massive ejecta may explain late-time plateaus, re-brightening events, slow decays, and unexpected breaks observed in the X-ray and optical light curves of some afterglows \citep{kp00,zm02,gnp03,pmg+06,mhm+07,gvs+07,mgg+10,hdpm+12,llt+12,gkn+13,pvw13,nef+14,dpko+15,bm17}, and forms a distinct class of models from late-time central engine activity, which has been invoked to explain some rapid X-ray and optical flares \citep{brf+05,ikz05,gngc09,nggc10,mgc+10,mbbd+11,llt+12}. The process of energy transfer between the ejecta and the circumburst medium is expected to be mediated by a reverse shock (RS) propagating into the ejecta during the injection period. This RS is similar to the one expected from the deceleration of the ejecta by the circumburst environment as observed in exquisite detail in the afterglow of GRB\,130427A \citep{lbz+13,pcc+14,vdhpdb+14}; however, an RS supported by energy injection is expected to continue propagating into the ejecta during the entire injection period \citep{zm02}. If injection takes place in the form of a violent shell collision, the resulting strong RS is expected to exhibit a detectable observational signature in the form of an optical flash or radio flare \citep{abb+99,sp99a,kfs+99,sr02,zm02,kz03,bsfk03,sr03,clf04}. In the case of gentle, or continuous energy injection, the RS is long-lasting, and its flux remains proportional to that of the FS during the entire injection period, $\fnumr\propto\Gamma\fnumf$ \citep{sm00,zm02,pk04,gdm07,lcj17}. Thus, it may be possible to detect reverse shocks arising from energy injection in cases both of violent collisions and of interactions at high enough ejecta Lorentz factor. Strong RS signatures are also excellent probes of the magnetization of the jets ($\sigma_{\rm B}$), since high $\sigma_{\rm B}$ effectively increases the sound speed\footnote{In magnetized media, information travels at the speed of the fast magnetosonic wave.}, thereby suppressing shock formation \citep{gma08}. Our previous observations of GRB\,140304A at $z\approx5.3$ yielded the first multi-frequency, multi-epoch detection of a RS from a violent shell collision, lending credence to the multi-shell model \citep{lbm+17}. However, the high redshift of this event impacted the quality of data, limiting the strength of the inference feasible. In an analysis of four GRB afterglows exhibiting late-time optical and X-ray re-brightening events, we constrained the distribution of ejecta energy as a function of Lorentz factor \citep{lbm+15}. In one case, our observations were incompatible with RS radiation from the injection, suggesting collisions in at least some instances may be gentle processes; for the remaining three cases, the observations lacked the requisite temporal sampling and frequency coverage to conclusively rule out an injection RS. The reason may partly stem from the fact that the RS emission peaks in the mm-band for typical shock parameters, and no facilities in this observing window had the requisite sensitivity \citep{duplm+12}. However, the advent of the Atacama Large Millimeter/submillimeter Array (ALMA) now allows us to track the mm-band evolution of afterglows to a sensitivity $\sim 30$--$100\,\mu$Jy for the first time, re-energizing the search for refreshed reverse shocks. Here we report detailed radio through X-ray observations of GRB\,161219B at $z=0.1475$, and present the first ALMA light curve of a GRB afterglow. The cm-band SEDs at $\lesssim 8.5$\,d exhibit unusual spectral features, which we discuss in detail in a separate work (Alexander et el., in prep; henceforth ALB18). Through multi-wavelength modeling of the X-ray, optical, and late radio data, we constrain the parameters of the FS powering the afterglow emission. The resulting model over-predicts the early X-ray emission, which can be explained by an episode of energy injection culminating at $\approx0.25$\,d. We interpret the early optical and radio observations as arising from a reverse shock launched by the same injection event. By tying the RS and FS parameters together, we show that the ejecta were not strongly magnetized. We employ standard cosmological parameters of $\Omega_m=0.31$, $\Omega_{\lambda}=0.69$, and $H_0=68$\,km\,s$^{-1}$\,Mpc$^{-1}$. All magnitudes are in the AB system and not corrected for Galactic extinction\footnote{Galactic extinction correction based on \cite{sf11} is built into our modeling software \citep{lbt+14}.}, all uncertainties are at 1$\sigma$, and all times are relative to the \Swift\ trigger time and in the observer frame, unless otherwise indicated. \begin{deluxetable}{lc} \tabletypesize{\footnotesize} \tablecolumns{2} \tablecaption{XRT Spectral Analysis for GRB\,161219B\label{tab:xrtspect}} \tablehead{ \colhead{Parameter} & \colhead{Value} } \startdata $T_{\rm start}$ (s) & $1.1\times10^{2}$ \\ $T_{\rm end}$ (s) & $1.1\times10^{7}$ \\ $N_{\rm H, gal}$ ($10^{20}$ \pcmsq) & 3.06 \\ $N_{\rm H, int}$ ($10^{21}$ \pcmsq) & $2.2\pm0.1$\\ Photon index, $\Gamma_{\rm X}$ & $1.86\pm0.03$\\ Flux$^{\dag}$ (observed) & $(1.86\pm0.05)\times10^{-12}$ \\ Flux$^{\dag}$ (unabsorbed) & $(2.41\pm0.06)\times10^{-12\ddag}$ \\ C statistic (dof) & 684 (699) \enddata \tablecomments{${}^\dag$ erg\,\pcmsq\,s$^{-1}$ (0.3--10\,keV); $^{\ddag}$ assuming the same fractional uncertainty as for the absorbed flux.} \end{deluxetable} | We have presented detailed multi-wavelength observations of GRB\,161219B, SN\,2016jca, and their host galaxy, including the first ALMA light curve of a GRB afterglow, and the first direct detection of an energy injection RS. Through simultaneous multi-frequency modeling, we constrain the properties of the afterglow, supernova, and host, and determine that the GRB occurred in an extremely low density environment, $\dens\approx3\times10^{-4}$\,\pcc. The data constrain the beaming angle of the relativistic outflow, allowing us to derive the degree of ejecta collimation ($\thetajet\approx13^{\circ}$) and to correct the $\gamma$-ray and kinetic energy for beaming, $E_{\gamma}\approx4.9\times10^{48}$\,erg and $\EK\approx1.3\times10^{50}$\,erg. The prompt efficiency is low, $\etarad\approx4\%$. The early radio and optical data require an additional emission component, which we interpret as synchrotron radiation arising from a refreshed reverse shock, powered by injection of energy into the forward shock through slow-moving ejecta. The combined model explains the X-ray to radio light curves over 8 orders of magnitude in frequency and 5 orders of magnitude in time. We measure a low ejecta magnetization, and our observations provide another confirmation for the internal shock model of GRB prompt emission. The supernova component is fainter and evolves faster than SN\,1998bw, while the stellar mass of the host galaxy is comparable to that of GRB hosts at $z\lesssim1$. We conclude that detailed multi-frequency radio observations and early optical detections are key to constraining refreshed reverse shocks in GRBs, and may yield crucial insight into the production and nature of GRB jets. | 18 | 8 | 1808.09476 |
1808 | 1808.09195_arXiv.txt | The Earth is impacted by 35-40 metre-scale objects every year. These meteoroids are the low mass end of impactors that can do damage on the ground. Despite this they are very poorly surveyed and characterised, too infrequent for ground based fireball observation efforts, and too small to be efficiently detected by NEO telescopic surveys whilst still in interplanetary space. We want to evaluate the suitability of different instruments for characterising metre-scale impactors and where they come from. We use data collected over the first 3 years of operation of the continent-scale Desert Fireball Network, and compare results with other published results as well as orbital sensors. We find that although the orbital sensors have the advantage of using the entire planet as collecting area, there are several serious problems with the accuracy of the data, notably the reported velocity vector, which is key to getting an accurate pre-impact orbit and calculating meteorite fall positions. We also outline dynamic range issues that fireball networks face when observing large meteoroid entries. | The Earth is impacted by 35-40 metre-scale objects every year \citep{2002Natur.420..294B,2006M&PS...41..607B}. These large meteoroids are at the low mass end of potentially damage-causing impacting asteroids like Chelyabinsk \citep{2013Natur.503..238B}. The study of the atmospheric behaviour, physical nature, numbers, and dynamical origin of these objects is therefore important in order to assess the hazard they pose, and prepare an appropriate response should an asteroid be detected and determined to be on a collision course with earth. \subsection{How frequently do these impacts happen?} One of the ways the size frequency distribution (SFD) of metre-scale has been surveyed is by using the so-called US Government (USG) sensors\footnote{\url{https://cneos.jpl.nasa.gov/fireballs/} accessed November 22, 2017}, which are able to detect flashes all around the world, day and night, measure flash energy, and sometimes derive velocities and airburst heights. As outlined by \citet{2013Natur.503..238B}, there might be subtleties in the SFD, namely a larger number of 10-50\,m objects. Indeed the 1-100\,m size range is largely unobserved, with objects too small for telescopes and too infrequent for impact monitoring systems to get representative surveys. So far, there have been 3 cases of asteroids detected before atmospheric impact. These are asteroids 2008 TC3 \citep{2009Natur.458..485J, 2017Icar..294..218F}, 2014\,AA \citep{2016Icar..274..327F}, and 2018\,LA, all discovered by the Catalina Sky Survey only hours before impact. As large deep surveyors like LSST \citep{2008SerAJ.176....1I} come online these types of detections are going to become more common, and predicting the consequences of these impacts is going to be desirable. While the impact location of 2008\,TC3 was well constrained to sub kilometre precision thanks to a very large number ($\simeq$900) of astrometric measurements, the prediction for 2014\,AA was much more uncertain and covered a large area of the Atlantic ocean, as only a total of 7 astrometric positions were available. The impact location of 2018\,LA was very uncertain, until 2 extra observation by the Asteroid Terrestrial-impact Last Alert System (ATLAS) increased the observation arc length from 1.3\,hours to 3.7 hours, which narrowed down the impact location to South Africa. The number of astrometric observations and the length of the observation arc are therefore a critical factors to precisely determining the impact point. Well coordinated, large follow-up networks of telescopes can provide large numbers of such observations and will aid in future impact predictions \citep{2016DPS....4840506L}. \subsection{How dangerous are these impacts?} The damage from an impact depends not only on dynamical parameters, but also on: size, rock type, structure, strength ($s$) and density ($\rho$). To illustrate this, we can use the equations of \citet{2005M&PS...40..817C} to simulate the outcome of the impact of a 2\,m object, with an entry angle of 18\degr, a velocity of 19\,$\mbox{km s}^{-1}$ at the top of the atmosphere (same entry angle and velocity as Chelyabinsk), and various bulk strengths and densities corresponding to different classes of objects (from \citet{1993Natur.361...40C}): \begin{itemize} \item a weak cometary body ($s=10^5$\,Pa, $\rho=1000\,\mbox{kg m}^{-3}$) will break up at a high altitude (60\,km), causing no significant direct damage because the predicted 0.18\,kT TNT of energy released cannot be transferred efficiently to the ground due to the thin atmosphere (1\,kT\,TNT = $4.184\times10^{12}$\,J). \item a chondritic body ($s=10^7$\,Pa, $\rho=3500\,\mbox{kg m}^{-3}$) is likely going to airburst at relatively low altitudes (the model predicts an airburst at 27\,km), releasing around 0.44\,kT TNT of energy that can be propagated more efficiently by the denser atmosphere. \item an iron ($s=10^8$\,Pa, $\rho=7900\,\mbox{kg m}^{-3}$) monolith will reach the surface at hypersonic velocity (3.8\,$\mbox{km s}^{-1}$), causing important but very localised damage, as it only yields $10^{-1}$\,kT TNT. \end{itemize} This is a simplistic example, but it shows how much the response to an imminent asteroid impact depends on both physical and dynamical characteristics of the impactor. Several observation techniques can be levied while the asteroid is still in interplanetary space: \begin{itemize} \item Multi-band photometry in Vis-NIR: size and rotation period, and lower constraint on cohesive strength as a consequence. \item Spectroscopy: likely composition. \item Astrometric observations: pre-encounter orbit, and predictions about the impact geometry, velocity, and location. \item Radar observations: size, shape, rotation period, presence of satellites. \end{itemize} While the size and impacting velocity are well constrained factors using astrometric observations, determining the rock type and structure from remote sensing instruments is more challenging. To some extent spectroscopy can provide insights on the mineralogy of the impactor, but this technique requires a good knowledge of how asteroid spectral types match meteorite types. Another approach is the work of \citet{2014ApJ...786..148M,2014ApJ...789L..22M} on small (metre-scale) asteroids for which spectroscopic work is generally impractical. They used a thermophysical model combined with an orbital model that takes non-gravitational forces into accounts. This model derives physical parameters (likely surface composition, size) by combining both astrometric observations and Near-Infrared photometry. In order to be reliable on large scales, these techniques have to be qualified with direct sample analysis. This active area of research can be tackled in two ways: either direct sample return missions (like Stardust, Hayabusa, Hayabusa 2, OSIRIS-REx), or from a large number of meteorite recoveries with associated orbits that can link to asteroid families: the aim of ground-based efforts like the Desert Fireball Network. The Desert Fireball Network (DFN) is a fireball camera network currently operating in the the Australian outback, designed for the detection and recovery of meteorite falls with associated orbits. Currently 52 observatories are deployed. On January 2, 2015, a particularly bright fireball was observed over South Australia, large enough to be simultaneously detected by the US government (USG) sensors, and by the DFN, which had just started science operation 2 months before. Another similarly bright event, also observed by both the DFN and the USG sensors, happened on June 30, 2017 over South Australia. Over the 3 million km$^2$ that the DFN covers in Australia, the observation of a metre-scale impactor is only expected to happen once every 4-5 years \citep{2002Natur.420..294B}, and once every 8-10 years during night time when most dedicated fireball networks operate (without considering clear sky conditions). The observation of two such events during the first 3 years of operation of the DFN, although outside the nominal collecting area, is somewhat lucky with respect to the size frequency distribution numbers of \citet{2002Natur.420..294B}. These two superbolides are described here and add to the small list of metre-scale impactors that have precisely determined trajectories: \begin{itemize} \item 13 events compiled and discussed by \citet{2016Icar..266...96B}. \item the "Romanian" bolide \citep{2017P&SS..143..147B}. \item the Dishchii'bikoh meteorite, for which initial trajectory details have been reported by \citet{2018arXiv180105072P}. \item the meteorite fall near Crawford Bay in British Columbia (Canada), for which initial trajectory details have been reported by \citet{2018LPI....49.3006H}. \end{itemize} \subsection{Where do they come from?} The current state of the art for source region model for Near-Earth Objects (NEO) is detailed by \citet{2018Icar..312..181G}. They report a significant size dependence of NEO origins, which had not been investigated by earlier similar works \citep{2002Icar..156..399B,2004Icar..170..259B,2012Icar..217..355G}. Their work covers the absolute magnitude range $17<H<25$ (corresponds to diameter $1200>D>30$\,m with an S-type albedo of 0.2), providing little insight on the the metre-size region ($H=32$). Several outstanding issues show that it is not possible to simply interpolate the characteristics of the population of typical macroscopic meteorite dropper meteoroids (decimetre-scale) and the kilometre-scale well surveyed by telescopes. For instance, LL chondrites make up 8\% of meteorite falls, but it is generally thought that $1/3^{rd}$ of observable near-earth small body space is made up of LL compatible asteroids \citep{2008Natur.454..858V}. \citet{2016Natur.530..303G} shows that an unmodelled destructive effect prevents small bodies from stably populating the low perihelion region, further outlying the need to consider body size in the dynamical models. \citet{2016Icar..266...96B} are the first to perform a source region analysis on metre-class NEO bodies, using the \citet{2002Icar..156..399B} model on USG events. Considering the small number statistics they get intermediate source regions proportion that are comparable to previous works on kilometre-size NEO population \citep{2002Icar..156..399B,2004Icar..170..259B,2012Icar..217..355G}. However they also argue for a Halley-type comet (HTC) source region, comparable in importance to the Jupiter-family comets (JFC) source. This source has not been identified previously in NEO works, because of a near-complete lack of such objects in asteroid databases. Their argument is based on three fireball events in the USG dataset that have a Tisserand parameter with Jupiter, $T_J<2$: identified as \textit{20150102-133919}, \textit{20150107-010559}, and \textit{20150311-061859}, not associated to a meteor shower. Because the first two of these events have independently estimated trajectories, an issue that we are interested in is determining if this surprising outcome could be the results of limitations of USG data. This work aims to compile independent information not just for these cases, but for several other metre-scale bodies, to determine the reliability of USG data in general, for population study, orbit determination, as well as undertaking meteorite searches based on these data. We also evaluate the suitability of hardware currently deployed by fireball networks to observe these particularly bright events. | \begin{enumerate} \item USG sensors data are generally unreliable for orbit calculations. The new metre-scale impactors source region of \citet{2016Icar..266...96B} (Halley-type comet orbits) is based on 3 particular USG meteoroid orbits. We have shown that 2 of these are erroneous, seriously questioning the existence of this source region. \item Size frequency distribution work relies on determining rough sizes and having a good knowledge of the probing time-area. The USG seem to achieve both with reasonably good precision. This confirms the sound basis of the work done by \citet{2002Natur.420..294B} and \citet{2013Natur.503..238B}. \item Basic impactor physical properties (size and strength) can be well constrained with USG data. This validates the conclusions of \citet{2016Icar..266...96B} that relate to physical properties of objects. % \item Based on how often the derived trajectories are wrong, it would be naive to invest large amounts of resources to undertake meteorite searching using USG data. \end{enumerate} We also note that ground based fireball networks must find solutions to increase the dynamic range of their observations, in order to get sound observation data when metre-scale objects impact the atmosphere. | 18 | 8 | 1808.09195 |
1808 | 1808.00557_arXiv.txt | The Habitable-zone Planet Finder (HPF) is a highly stabilized fiber fed precision radial velocity (RV) spectrograph working in the Near Infrared (NIR): 810 -- 1280 nm . In this paper we present an overview of the preparation of the optical fibers for HPF. The entire fiber train from the telescope focus down to the cryostat is detailed. We also discuss the fiber polishing, splicing and its integration into the instrument using a fused silica puck. HPF was designed to be able to operate in two modes, High Resolution (HR- the only mode mode currently commissioned) and High Efficiency (HE). We discuss these fiber heads and the procedure we adopted to attach the slit on to the HR fibers. | \label{sec:intro} Optical fibers were introduced in astronomy in the late 1970s and recognized for their multiplexing and scrambling ability\cite{k._serkowski_fabry-perot_1979}. They were proposed as a means of mitigating the guiding errors \cite{black_assessment_1980-1}. Other drivers for using Optical fibers were that it allowed the instrument to be decoupled from the telescope focus and instead to be located in the basement of the observatory. This enabled a constant gravity vector on the instrument and also allows for environmental control\cite{barden_evaluation_1981,heacox_application_1986}. Around the same time, spectroscopic methods were discussed to find extra - solar planets\cite{black_assessment_1980}. The first exoplanet around a hydrogen - burning star was discovered in 1995 using the ELODIE spectrometer \cite{mayor_jupiter-mass_1995}. ELODIE was one of the first precision radial velocity (PRV) spectrometers to use an optical fiber feed \cite{baranne_elodie:_1996}. It used a dual fiber feed for simultaneous science and calibration light. This method enabled it to reach an unprecedented RV precision of $\sim$ 13 m/s. This was taken to the next level in 2003 by High Accuracy Radial velocity Planetary Searcher (HARPS)\cite{mayor_setting_2003}. The instrument was housed in a cryostat maintained under vacuum, which helped it achieve exquisite environmental stability. It was one of the first instruments to approach $\sim$ 1 m/s in its RV precision. Since then, it has continued to remain the gold standard in RV precision, and has even demonstrated sub m/s precision\cite{lovis_exoplanet_2006}. However, the next generation of optical PRV are aiming for $\sim$ 10 cm/s precision, as that is the magnitude of the Doppler reflex motion of a Sun like star due to a true Earth analogue. Similarly, near-infrared (NIR) instruments are seeking to find an Earth like planet in the Habitable zone\cite{kasting_habitable_1993} around M dwarfs. Since these stars are much smaller and cooler than the Sun, such planets produce a reflex motion of the order of 1 m/s in the NIR. HPF is one of these new NIR instruments. It spans 810 to 1280 nm and uses an R4 echelle grating in a white pupil design\cite{mahadevan_habitable-zone_2012}. Deployed at the 10 m class Hobby-Eberly Telescope (HET) in October 2017, it started shared risk science in May 2018. It uses a Teledyne Hawaii-2RG (H2RG) detector with a 1.7 $\mu$m cut off, and the optical bench and optics are cooled down to 180K and kept stable at that temperature using a sophisticated environmental stability system\cite{stefansson_versatile_2016}. | 18 | 8 | 1808.00557 |
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1808 | 1808.07188_arXiv.txt | The radio-X-ray slope in the fundamental plane of radio-loud active galactic nuclei (AGNs) is found to be steeper compared with that of radio-quiet AGNs in previous works. In this work, we reinvestigate the fundamental plane in radio-loud AGNs by compiling a sample of 13 low-excitation radio galaxies (LERG) from the 3CR radio galaxies, for the reason that the accretion mode in LERG is believed to be a radiatively inefficient accretion flow. All the sources in our sample possess the data available at both the 5 GHz core radio luminosity detected by VLA/VLBI/VLBA and the core X-ray luminosity detected by Chandra/XMM-Newton. Surprisingly, we find the slope in the fundamental plane ($\log L_{\rm R}=0.52 \log L_{\rm X}+ 0.84 \log M_{\rm BH} + 10.84$) of LERG is well consistent with that reported by \citet{m2003}. However, the normalization is found to be shifted by about 0.7 dex, which can be due to the difference on magnetic field strength in different objects. A shallower slope of $L_{\rm R}-L_{\rm X}$ relation ($L_{\rm R}\sim L_{\rm X}^{0.63}$) is also given by our sample, which demonstrates that the X-ray emission in LERG may come from accretion disc instead of jets as suggested by previous works. | Accretion process is generally accepted to be the central engine of black hole astrophysical systems \citep{f2002}, e.g., black hole X-ray binaries (BHBs) and AGNs. In black hole systems with relativistic jets, a tight correlation between the X-ray and radio emissions ($L_{\rm R}\propto L_{\rm X}^{0.7}$) was reported by numerous works (e.g., \citealt{c2003,g2003,m2003,f2004}), where the X-ray and radio emissions are believed to come from accretion disc and jet, respectively. Coupled with black hole mass, a so-called fundamental plane (${\rm log}L_{\rm R}=0.6{\rm log}L_{\rm X}+0.78{\rm log}M_{\rm BH}+7.33$) was developed by \citet[][hereinafter, M03]{m2003} to manifest the activity of black hole and further explored by lots of following works (e.g., \citealt{f2004,m2006,l2008,d2015,n2016}). While the fundamental plane is prevailing in recent years, there are still some noises remaining controversial. At first, the secular quasi-simultaneous observations on radio and X-ray fluxs in BHBs illustrated that their evolution can deviate obviously from the original M03 fundamental plane, which are known as 'outliers' \citep{x2007,c2011,c2013}. These outlier tracks possess a much steeper slope ($L_{\rm R}\propto L_{\rm X}^{1.4}$, see, e.g., \citealt{c2011}) compared with the original fundamental plane in radio-X-ray plane. The different radio-X-ray slopes may be originated from different accretion mode (the slopes of 0.7 and 1.4 correspond to radiatively inefficient and radiatively efficient accretion flows, respectively, \citealt{c2011, c2014}), or from the change of viscosity parameter $\alpha$ in a hot accretion flow \citep{x2016}. Secondly, it has long been suggested that the X-ray emission will be dominated by jet in the quiescent state of BHBs when the X-ray luminosity decrease to a critical value ($L_{\rm x,crit}/L_{\rm Edd}\sim 10^{-6}$, see \citealt{f2003,y2005}), resulting in a much steeper radio-X-ray slope (e.g., \citealt{w2007,p2013,r2014}). However, there are some works claiming that, even in quiescent state, the radio-X-ray correlation is still complied with the original M03 fundamental plane (\citealt{g2014,d2015}, but also see \citealt{x2017}). The last point, which is also the focus of this work, is the slope between radio and X-ray in radio-loud AGNs seems appear to be much steeper too compared with the radio-quiet AGNs \citep{w2006,l2008,d2011}, possibly due to the domination of strong jet emissions on radio and/or X-ray bands. Radio galaxies (RG) can be divided into two classes according to the large-scale jet morphology traditionally, i.e., edge-darkened FRI and edge-brightened FRII \citep{f1974}. The difference between FRI and FRII can be originated from the interaction of jets with different power and their ambient mediums \citep{b1995,t2016}, and FRII usually possess higher jet power than FRI. Another important classification of RG is based on their optical spectroscopic information, where the RG with weak and strong emission lines are classified as LERG and high-excitation RG (HERG), respectively (e.g., \citealt{h1979,h2009}). HERG tend to have higher radio luminosity, similar with FRII. However, there isn't an one-to-one match between these two classifications. Both FRI and FRII can comprise LERG and HERG (e.g., see \citealt{l1994}). From the observations of LERG, a radiatively inefficient accretion flow (RIAF) should be present due to their lack of AGN symbols, such as corona and torus, while HERG are believed to be powered by a radiatively efficient cold accretion disk \citep{h2007}. Similarly, low-luminosity AGNs are also found to be different in various aspects with bright AGNs (see, e.g., \citealt{h2008,g2009,s2011,x2011,l2017}). Therefore, we can naturally anticipate that the original M03 fundamental plane will change for HERG because of the transition of accretion mode. As a result, in order to investigate the radio-X-ray slope in radio-loud AGNs, we must discriminate LERG and HERG at first. We notice that in some previous works where a steeper slope was reported, the authors didn't distinguish LERG and HERG from radio-loud AGNs (e.g., \citealt{l2008}). In this work, we reinvestigate the fundamental plane of radio-loud AGNs by constructing a sample satisfying the following conditions: 1), since the accretion modes of LERG and HERG are different, all the sources in the sample should be LERG in order to ensure the accretion flow is RIAF. 2), radio emission mainly comes from jet in radio-loud AGNs, where both the core and lobe can play important roles. To avoid the influence of surrounding mediums, we adopt the source being core dominated only. | In this work, we compile a sample of LERG from the 3CR radio galaxies with optical spectroscopic information \citep{b2009}. After excluding the sources with excitation index $\rm {EI} > 0.95$, a sample of 13 LERG is found to contain both the data of core radio luminosity at 5 GHz detected by VLA/VLBI/VLBA and core X-ray luminosity detected by Chandra/XMM-Newton. Surprisingly, We discover a similar radio-X-ray slope with that of M03 fundamental plane, which suggests that the low-luminosity radio-loud AGNs (LERG) still follow the original M03 fundamental plane. We notice that \citet{d2011} had investigate the fundamental plane in a LERG sample either. They discovered a steeper radio-X-ray slope and advised the X-ray emissions in LERG may be originated from jets, though their X-ray luminosity $L_{\rm X}$ are larger than $L_{\rm crit}$ (see below). The reason for this inconformity may be that their sample also comprised some steep spectrum LERG except for the core dominated flat spectrum LERG. Furthermore, we find the normalization of our sample is larger than that in M03 by about 0.7 dex, though this deviation is still within the range of their error bars. The possible reason for this movement can be the variant magnetic field strength in different objects. If we consider the parameter $\beta$ (the ratio of gas pressure to magnetic pressure) isn't a constant in different objects, the radio flux from jet can be roughly written as $L_{\rm R}\propto \dot{m}^{1.4} \beta^{-1}$ ($L_{\rm R}\propto \dot{m}^{1.4}$ when $\beta$ is constant, see \citealt{h2003}). The X-ray flux can be revised as $L_{\rm X}\propto \dot{m}^{2} \beta^{a}$ ($a>0$) for the same way, because the X-ray flux increase when the magnetic field strength decrease (see \citealt{m1997}). Therefore, the revised $L_{\rm R}$-$L_{\rm X}$ relation can be given as: \begin{equation} \log L_{\rm R}\propto 0.7 \log L_{X}-(1+0.7a)\log \beta. \label{mag} \end{equation} All the objects in our sample are radio-loud, which means higher magnetic field strength and then smaller $\beta$. According to equation (\ref{mag}), we can naturally anticipate a higher normalization for radio-loud objects. Indeed, the large-scale magnetic field is easy to be magnified in a RIAF due to their high radial velocity \citep{l1994,c2011b,l2014} and can strongly affect the activity of black hole. Except for the obvious augment in radio emission, the X-ray emission in a RIAF is a complicated function of the magnetic field strength based on the theoretical research \citep{n1994,n1995,m1997,b2013,b2016,s2015}. Furthermore, large-scale magnetic field can change the value of viscosity parameter $\alpha$ according to the recent magneto hydrodynamical (MHD) simulations \citep{bs2013,s2016}, which can further decrease the X-ray emission of RIAF (e.g., \citealt{n1994,l2017}). These points will be explored in future works. The slope of $L_{\rm R}-L_{\rm X}$ correlation in LERG is also found to be consistent with other black hole systems (e.g., \citealt{c2003,g2003,m2003,f2004}), but much shallower than that found in FRI samples (e.g., \citealt{d2015}). In theory, \citet{y2005} suggested that there is a critical X-ray luminosity ($L_{\rm crit} = L_{\rm {X,2-10 keV}} \sim 10^{-6} L_{\rm Edd}$) to diagnose the origin of X-ray in low-luminosity AGNs. When $L_{\rm X} > L_{\rm crit}$, the X-ray from accretion disc will exceed that from jet, resulting on a shallower slope between the relation of $L_{\rm X}$ and $L_{\rm R}$. The Eddington ratio of ionizing luminosity $L_{\rm ion}/L_{\rm Edd}$ in our sample, which is a significant fraction of the bolometric luminosity \citep{w1999}, is all larger than $L_{\rm crit}$. Therefore, the shallower slope of $L_{\rm X}-L_{\rm R}$ correlation in Fig. \ref{f3} demonstrates that the X-ray emission should come from a RIAF in LERG. Our results indicate that, considering the core emissions of radio and X-ray, the radio-loud AGNs still comply with the physics of truncated accretion disc-jet model (e.g., \citealt{y2014}), which had been successfully applied in low-luminosity AGNs. | 18 | 8 | 1808.07188 |
1808 | 1808.02959_arXiv.txt | An intergalactic magnetic field stronger than $3\times10^{-13}$~G would explain the lack of a bright, extended degree-scale, GeV-energy inverse Compton component in the gamma-ray spectra of TeV-blazars. A robustly predicted consequence of the presence of such a field is the existence of degree-scale GeV-energy gamma-ray halos -- gamma-ray bow ties -- about TeV-bright active galactic nuclei, corresponding to more than half of all radio galaxies. However, the emitting regions of these halos are confined to and aligned with the direction of the relativistic jets associated with gamma-ray sources. Based on the orientation of radio jets, we align and stack corresponding degree-scale gamma-ray images of isolated Fanaroff-Riley class I and II objects and exclude the existence of these halos at overwhelming confidence, limiting the intergalactic field strength to $<10^{-15}$~G for large-scale fields and progressively larger in the diffusive regime when the correlation length of the field becomes small in comparison to 1~Mpc. When combined with prior limits on the strength of the intergalactic magnetic field, this excludes a purely magnetic explanation for the absence of halos. Thus, it requires the existence of novel physical processes that preempt the creation of halos, e.g., the presence of beam-plasma instabilities in the intergalactic medium or a drastic cutoff of the very high energy spectrum of these sources. | Very high energy gamma rays (VHEGRs, above 100~GeV) emitted from active galactic nuclei (AGNs) annihilate on the intergalactic infrared background after propagating over cosmological distances \citep{Gould+66,Stec-deJa-Sala:92,Ahar_etal:06,Fermi_EBL2012}. This results in a population of ultrarelativistic electron-positron pairs (with Lorentz factors of $10^6$), streaming primarily through the intergalactic medium (IGM) in cosmic voids \citep{Gould+66}. The fate of these pairs remains unclear, depending on the competition between nonlinear saturation of virulent plasma beam instabilities and inverse Compton (IC) cooling via the cosmic microwave background \citep{PaperI,Schlickeiser:2013,Chang:2014}. Should the latter dominate, it will effectively reprocess the original VHEGR emission of AGNs to lower energies, 1-100~GeV, generating an IC halo. At these energies, the {\em Fermi Space Telescope} provides a high-resolution (68\% inclusion region of the point spread function, PSF, of $\circp{0}{6}$), high-sensitivity map of the entire sky. The vast majority of observed gamma-ray bright AGNs are blazars, AGNs with jets that are directed at us \citep{3LAC}. This identification indicates that the gamma-ray emission is strongly beamed toward us \citep{Push_etal:09} and aligned with the underlying AGN jet. This anisotropy in the VHEGR emission has already been used to argue for lower limits on the strength of a putative intergalactic magnetic field (IGMF) threading cosmic voids. For a handful of known VHEGR sources, the absorbed VHEGR flux and corresponding IC halo flux have been estimated, and compact ($<\circp{0}{6}$), forward-beamed IC halo components are clearly excluded by {\em Fermi} observations at high significance. One explanation for this disparity is the presence of a strong IGMF, that is, $\gtrsim3\times10^{-16}$~G, that deflects the pairs out of the line of sight prior to their IC emission \citep{Nero-Semi:09,Nero-Vovk:10,Tayl-Vovk-Nero:11,Taka_etal:11,Vovk+12}. However, a robust prediction of this picture is the presence of extended, degree-scale IC halos about gamma-ray sources, corresponding to the IC emission missing from the line of sight \citep{1994ApJ...423L...5A,2009PhRvD..80b3010E,BTI}. The IC halo emission itself is typically exceedingly dim and therefore undetectable for an individual source \citep{1994ApJ...423L...5A,2009PhRvD..80b3010E,BTI}. Many attempts have been made to stack images from known gamma-ray sources and identify extended gamma-ray excesses, though these have met with little success, due, in part, to uncertainties in the PSF \citep{Ando:2010,FLAT-stack:2013}. Currently, the lack of any significant extension of the gamma-ray emission about known {\em Fermi} blazars has placed a stronger limit of $>3\times10^{-13}$~G assuming that the gamma-ray jet lifetimes exceed 10~Myr \citep{FermiIGMF:18}. While shorter jet lifetimes can reduce this limit, lifetimes significantly smaller than those inferred from radio observations of blazars would be inconsistent with the large fraction of nearby blazars observed by {\em Fermi} \citep{3LAC}. Recently, we have proposed exploiting the expected anisotropy of the IC halos to circumvent systematic uncertainties \citep{BTII}, and we have used this to exclude the presence of a large-scale, uniform IGMF at more than 4$\sigma$ \citep{BTIII}. This method relies on the bi-lobed anisotropy that results either from the fact that the electrons and positrons produced by the VHEGR annihilation on the intergalactic infrared background are deflected in opposite directions by a uniform IGMF or from the structure of the initial VHEGR jet in combination with a small-scale tangled IGMF \citep{2010ApJ...719L.130N,2015JCAP...09..065L,BTI,BTII,BTIII,2017JCAP...05..005D}. The structure in the image both increases the surface brightness of the pair halos and distinguishes them from the confounding systematics arising from the instrument response, subthreshold background sources, and diffuse Galactic contributions. The degree of anisotropy and the gamma-ray flux are strong functions of the jet orientation. For blazars, AGNs whose jets are directed at us, this presents a weak anisotropy. In contrast, for oblique jets (i.e., AGNs with jets more than $30^\circ$ off of the line of sight), the expected IC halo structure is striking \citep{BTI}. Such oblique jet sources are not, however, observed to be gamma-ray bright as their intrinsic gamma-ray emission is beamed away from us, and they are therefore visible primarily via their radio emission. Therefore, {\em if we can identify oblique gamma-ray jets and properly orient the images,} any excesses due to IC halos would be detectable with high significance. Here we employ the unified AGN paradigm to identify, align, and stack the oblique counterparts to the gamma-ray-bright blazars observed by {\em Fermi}. The fact that the parent population of radio-loud objects and misaligned blazar counterparts can be identified is supported by the same clustering properties of {\em Fermi} blazars (BL Lacertae and flat-spectrum radio quasars, or FSRQs) and radio-loud AGNs \citep{Allevato2014}. To do this, we utilize existing catalogs of radio jet sources identified in the Very Large Array, Faint Images of the Radio Sky at Twenty-centimeters (VLA FIRST) survey \citep{FIRST_CAT}, from which we obtain 20~cm images that show the radio jet orientation. Given both the locations and orientations of the oblique jets, we then stack the corresponding {\em Fermi} Large Area Telescope (LAT) observations after aligning them. This procedure was followed separately for Fanaroff-Riley class I and II (FR I and II) objects, presumably corresponding to BL Lacs and FSRQs, respectively. We then compare these with the anticipated IC halo signals, associated with their respective gamma-ray AGN object classes. In neither set of comparisons is any evidence for IC halos found. Based on our simulated stacked IC halos, we are able to exclude their existence by more than 6$\sigma$. We explore a variety of potential systematic uncertainties and find that none can adequately explain this nondetection. Therefore, we interpret this in terms of either a novel spectral cutoff between 100~GeV and 1~TeV in gamma-ray-bright AGN (although we provide a number of observational and theoretical arguments, which make this a very unlikely possibility) or as a result of an additional dissipative process that preempts the IC halo formation after the absorption of VHEGRs and the production of the relativistic pair population. This paper is organized as follows. In Section~\ref{sec:method}, the method of selecting and aligning sources with radio jets is described. In Section~\ref{sec:IC_expected}, details of how the expected anisotropic IC halo signal is computed are given. In Section~\ref{sec:results}, the expected halo signal from our stacking procedure is given together with a discussion of various potential systematic uncertainties. In Section~\ref{sec:discussion}, we offer an interpretation of the absence of any evidence of the IC halo signal. We summarize our conclusions in Section~\ref{sec:conclusion}. | \label{sec:conclusion} We have identified, aligned, and stacked the gamma-ray images of the oblique radio analogs of the gamma-ray-bright blazars. This was done independently for FR I and II objects, which we compared to the expectations from simulated IC halos from stacked BL Lacs and FSRQs, respectively. Based on this, we can conservatively exclude the existence of IC halos at more than 6$\sigma$. The apparent absence of IC halos cannot be explained via identified systemic uncertainties in the analysis. Alone this requires an IGMF less than $1\times10^{-15} \eta^{-1} (\lambda_B)$~G. Combined with prior constraints from {\em Fermi} blazar observations, which limit the IGMF to be greater than $3\times10^{-13}$~G (and progressively larger in the diffusive regime when the correlation length of the IGMF becomes smaller than the IC cooling time) for VHEGR jet lifetimes consistent with those assumed here, this precludes any interpretation of this nondetection within the context of an IGMF. That is, there is no IGMF that would simultaneously explain all nondetections. This substantially complicates efforts to probe the IGMF with gamma-ray observations. This suggests either (1) a novel process that suppresses the VHEGR emission dramatically between 100~GeV and 1~TeV, for which there is currently little empirical or theoretical support, or (2) a mechanism by which the IC halos are preempted after the relativistic pairs are generated by VHEGR absorption on the infrared background. Either of these solutions requires fundamental revisions to our understanding of the origin and/or impact of the VHEGR emission of AGNs. Modifications to the emission necessarily imply additional, possibly redshift-dependent, processes within the gamma-ray emission regions of AGNs, placing a severe constraint on their nature and location. Additional cooling processes in the IGM reprocess the VHEGR luminosity of AGNs into forms other than GeV halos, which are essentially decoupled from all other components in the IGM. Beam-plasma instabilities may serve this role \citep{PaperI,Chang:2014}, draining energy from the relativistic pairs into plasma waves and, ultimately, heat. If these instabilities are efficient, they would serve as a mechanism that satisfies (2) and, thus, invalidates the current IGMF limits from the nonobservation of GeV halos around TeV blazars. Additionally, beam-plasma instabilities deposit this energy into heat in the IGM, dominating the energy budget in cosmic voids, where it can raise the IGM temperature by up to an order of magnitude at $z=0$ \citep{PaperII,PaperIV,Lamberts:2015}, substantially modifying the Lyman-$\alpha$ forest at late times ($z\lesssim 2$), and possibly suppressing late-time star formation in galaxies, especially dwarfs \citep{PaperIII}. | 18 | 8 | 1808.02959 |
1808 | 1808.05112_arXiv.txt | { We found a possible $\sim$ 1 hour quasi-periodic oscillation (QPO) in a $\sim$ 55 ks X-ray observation of the narrow-line Seyfert 1 galaxy MCG--06--30--15 made with the {\it XMM-Newton} EPIC/pn detector in the energy range 0.3 -- 10 keV. We identify a total modulation of $\sim$ 16\% in the light curve and find a $\simeq$ 3670~s quasi-period using Lomb-Scargle periodogram (LSP) and weighted wavelet Z-transform (WWZ) techniques. Our analyses of eight light curves of MCG--06--30--15, indicated the possible presence of an oscillation during one of them. The LSP indicates a statistically significant ($\simeq$ 3$\sigma$) QPO detection. A WWZ analysis shows that the signal at this possible roughly 3670~s period is present, and rather persistent, throughout the observation; however, a signal around 8735~s is more persistent. We briefly discuss models that can produce X-ray QPOs with such periods in narrow line Seyfert 1 galaxies, as both other claimed QPO detections in this class of AGN had very similar periods. } | Detections of quasi-periodic oscillations (QPOs) are very rare in active galactic nuclei (AGN), but are fairly common in both black hole (BH) and neutron star binaries in the Milky way and nearby galaxies (Remillard \& McClintock 2006). Over the last decade there have been several claims of QPO detections on diverse timescales ranging from a few tens of minutes to hours to days and even years, using $\gamma-$ray, X-ray, optical and radio monitoring data of various classes of AGN (Gierli{\'n}ski et al.\ 2008; Espaillat et al.\ 2008; Gupta et al.\ 2009; Lachowicz et al.\ 2009; Lin et al.\ 2013; King et al. 2013; Fan et al. 2014; Sandrinelli et al.\ 2014, 2016a, 2016b; Graham et al.\ 2015; Ackermann et al.\ 2015; Pan et al.\ 2016; Bhatta et al.\ 2016; Bhatta 2017; Li et al. 2017; Xiong et al. 2017; Zhang et al. 2017a, 2017b, 2018; Hong et al. 2018; and references therein). Since any such QPOs are almost certainly transient, verification is difficult and claims are strengthened if more than one technique provides consistent results. The first significant QPO reported for a narrow-line Seyfert 1 (NLSy1) galaxy was in a light curve of RE J1034$+$396 (Gierli{\'n}ski et al.\ 2008); this had a roughly 1 hour timescale (3730 s) in the X-ray band and was found in {\it XMM-Newton} data. The second QPO detection in a NLSy1 galaxy was more recently made by Pan et al.\ (2016) in 1H 0707$-$495; it again had a period of $\sim$ 1 hour (3800 s) and was seen in a {\it XMM-Newton} observation. We report a third probable QPO detection in a NLSy1 for MCG--06--30--15 ($z = 0.00775$), once again using {\it XMM-Newton} data. MCG--06--30--15 is well known, particularly for its broad iron K$\alpha$ line; this line that provides strong evidence for the presence of a supermassive black hole (SMBH; e.g.,\ Tanaka et al.\ 1995), which is probably spinning rapidly (e.g.,\ Iwasawa et al.\ 1996). Because its H$\beta$ line width (FWHM: full width at half maxima) of $1933\pm82$ km s$^{-1}$ is $< 2000$ km s$^{-1}$ (Hu et al.\ 2016), this SMBH is classified as a NLSy1. Interestingly, the possible QPO we report on in this work also shows a period of $\sim$ 1 hour. Convincing detection and careful characterizations of additional QPOs in these and other NLSy1 galaxies may shed new light on the physical processes occurring in these sources and can yield information on their SMBH masses and spins (e.g.,\ Abramowicz \& Klu{\'z}niak 2001; Zhou et al.\ 2015; Pan et al.\ 2016). The search for QPOs in light curves of AGN is very important: their presence can provide strong support for the common nature of the accretion process onto BHs ranging from a few solar masses up to the SMBHs present in quasars (e.g.,\ Abramowicz \& Klu{\'z}niak 2001; Remillard \& McClintock 2006; Zhou et al.\ 2015). Plausible models that might explain QPOs in AGN in different wavebands and on diverse timescales have been put forward (e.g.,\ Gierli{\'n}ski et al.\ 2008; Gupta et al.\ 2009; Lachowicz et al.\ 2009; Pan et al.\ 2016; Bhatta 2017, and references therein). However, only once we have at least a handful of good cases of QPOs in different subclasses of AGN and on diverse timescales will it be possible for us to gain a thorough understanding of this phenomenon. \begin{figure} \centering \includegraphics[scale=0.48]{fig1.eps} \\ \caption{XMM-Newton EPIC/pn image of the NLSy1 MCG--06--30--15 and the selected background region, denoted by a circle.} \end{figure} \begin{figure*} \centering \includegraphics[scale=0.90]{fig2.eps} \vspace*{-3.5in} \caption{Light curve of NLS1 MCG--06--30--15 in 0.3 -- 10 keV observed with {\it XMM-Newton} EPIC/pn, with a running average over 5 points given by the continuous red curve.} \end{figure*} In Section 2 of this Letter we briefly describe the X-ray data and how we reduced it. In Section 3 we present a description of the QPO search techniques we employed and the results of those analyses. A discussion and our conclusions are given in Section 4. | The quasi-periods obtained from the LSP and WWZ analyses, which search for the presence of sinusoidal components, agree to within the resolutions of these techniques. Together, they indicate that during this particular observation the NLSy1 MCG--06--30--15 exhibited a QPO at a frequency of $ \simeq 2.73 \times 10^{-4}$ Hz, or at a central period of $\simeq 3670$~s. The LSP approach is the most common frequency-based technique and can provide good estimates of the strength of a signal if the underlying power spectrum can be sensibly modeled, but it does not take into account any possible evolution of a periodic or quasi-periodic signal with time. The time-frequency approach using wavelets has the advantage of quantifying the persistence of such signals in the data. Because we found no indication of a QPO in the other seven extensive light curves for this source that we examined, one may consider that the formal statistical confidence in any long-lived QPO is reduced to $\sim$ 90\%. However, we believe that it is more appropriate to interpret this result as being consistent with the general finding that detectable QPOs from AGN are not long-lived, but rather are transient in nature (Gierli{\'n}ski et al.\ 2008; Pan et al.\ 2016), and thus this object may have a QPO duty cycle of $\sim$12\%. We note that our signal appears to be present throughout the full observation of MCG--06--30--15; this is in contrast to the QPOs detected in the other two NLSy1s, RE J1034$+$396 (Gierli{\'n}ski et al.\ 2008; Czerny et al.\ 2010) and 1H 0707$-$495 (Pan et al.\ 2016), in which the QPOs were only significant in portions of the data trains. Recently the SMBH in MCG--06--30--15 has had its mass estimated in two independent direct ways using the reverberation mapping technique for the H$\beta$ line. Hu et al.\ (2016) found a preferred value of $3.26^{+1.59}_{-1.40} \times 10^6 M_{\odot}$ but Bentz et al.\ (2016) obtained $1.6 \pm 0.4 \times 10^6 M_{\odot}$ using different reverberation campaigns and somewhat different analysis techniques. The time lags these two groups measure at multiple epochs are reasonably consistent, at 6.38 d and 5.33 d, respectively, as are their values of the H$\beta$ FWHM, but Hu et al.\ (2016) used the FWHM and Bentz et al.\ (2016) used the dispersion, $\sigma$. These choices, along with different choices for the value of the scaling factor that accounts for the kinematics and geometry of the broad-line region gas, produce the differences in their best mass values. These values, nevertheless, are still consistent within the errors. These papers estimated Eddington ratios of $\sim 0.12$ (Hu et al.\ 2016) and $\sim 0.04$ (Bentz et al.\ 2016), which are substantial, but lower than that for most NLSy1s (Hu et al.\ 2016). Most of the peculiar properties of NLSy1s can be understood if they are Seyferts with relatively low SMBH masses that are accreting at relatively high rates. High-frequency QPOs (HFQPOs) for BH X-ray binaries are in the range 40--450 Hz and because their frequencies are constant despite large changes in luminosity it has been thought that they are related to the innermost portions of the accretion disks and hence may provide measures of the masses and spins of their BHs (e.g.,\ Abramowicz \& Klu{\'z}niak 2001; Abramowicz et al.\ 2004; Remillard \& McClintock 2006). A plot of HFQPO frequencies against BH mass produces the expected tight inverse relationship for stellar mass BHs and this has recently been extended from Galactic BH QPOs through possible intermediate mass BHs up to SMBHs (Zhou et al.\ 2015, Pan et al.\ 2016, and references therein). The frequency of this probable QPO in MCG--06--30--15 is nearly identical to those found in RE J1034$+$396 and 1H 0707$-$495 at $\simeq 2.7 \times 10^{-4}$Hz. The best mass estimates of this QPO, at $3.26^{+1.59}_{-1.40} \times 10^6 M_{\odot}$ or $1.6 \pm 0.4 \times 10^6 M_{\odot}$, are also very close to the others, which are $4^{+3}_ {-1.5} \times 10^6 M_{\odot}$ for RE J1034$+$396 (Zhou et al.\ 2010) and $5.2(\pm0.5{\rm dex}) \times 10^6 M_{\odot}$ for 1H 0707$-$495 (Pan et al.\ 2016). Hence, our new measurement appears to help confirm the inverse linear dependence of QPO frequency and BH mass and would lie right next to the other AGN points on Fig.\ 4 of Pan et al.\ (2016). \begin{figure*} \centering \includegraphics[angle=90, scale=0.63]{fig4.eps} \caption {Weighted wavelet z-transform of the light curve presented in Fig. 1. The left panel shows the distribution of color-scaled WWZ power (with red most intense) in the time-period plane, and the right panel shows the time-averaged WWZ power (blue curve) as a function of period. The dotted black curve represents $97\%$ global significance.} \end{figure*} This tight bunching of frequencies (and SMBH masses) for AGN QPOs may seem surprising but selection effects probably strongly favor QPO frequencies and SMBH masses relation. The higher the SMBH mass the longer the expected possible QPO period and the longer any observation would need to be to have a chance of detecting a QPO. Of course only Type 1 AGN afford us the opportunity of seeing the region close to the BH from which QPOs emerge. Those BHs are likely to be several times more massive in normal Sy1s than in the NLSy1s in which QPOs have been seen. Most QSOs have much more massive SMBHs, and while they are also much more luminous, they are also much further away. Therefore count rates are no greater and the required continuous observation times needed to perform good QPO searches on much longer timescales are essentially impossible to obtain. The presence of 3:2 ratios for the frequencies of many HFQPOs in X-ray binaries indicates that the physical mechanism producing these HFQPOs involves a resonance phenomenon of some sort (e.g.,\ Abramowicz \& Klu{\'z}niak 2001) and the scaled similarity between those stellar mass systems and these NLSy1 AGNs supports the idea that resonances are important for AGNs as well. Although no pairs of QPOs at that 3:2 ratio (or any other, for that matter) have yet been detected for any AGN and the discovery of such would be extraordinarily important, the resonance hypothesis certainly is reasonable. An estimated Eddington ratio of 0.12 is significantly higher than that of the great majority of AGN. This large Eddington ratio and the likelihood that this is another HFQPO supports the claim by Pan et al.\ (2016) that HFQPOs are associated with BHs accreting at very high rates as is the case for BH X-ray binaries. If a high accretion rate is indeed a prerequisite for engendering a HFQPO and the few detected in AGNs are indeed of the same type as in X-ray binaries, then many otherwise viable models (e.g., review by Belloni \& Stella 2014) are disfavored (Pan et al.\ 2016). A wavelet analysis of the apparent variations of the QPO frequency for the NLS1 RE J1034$+$396 by Czerny et al.\ (2010) indicated that an increase in QPO frequency was accompanied by an increase in X-ray flux. The wavelet approach to the current observation of MCG--06--30--15 does not provide additional evidence for such a trend, as the QPO frequency around 3670 s appears to be very stable. According to Pan et al.\ (2016), only models invoking p-modes trapped in the innermost part of an accretion disk (e.g., Li et al.\ 2003, Hor{\'a}k \& Lai 2013, and references therein) or those involving magnetized disks subject to accretion-ejection instabilities (Tagger \& Pellat 1999) naturally tie HFQPOs to high accretion rate situations. We thank the anonymous referee for constructive comments and suggestions. ACG thanks Prof.\ D.\ Banerjee for providing a computer code for wavelet analysis. ACG is partially supported by the CAS President's International Fellowship Initiative (PIFI), Grant No. 2016VMB073. AT acknowledges support from the China Scholarship Council (CSC), Grant No.\ 2016GXZR89. PJW is grateful for hospitality at SHAO while this paper was written. MFG is supported by the National Science Foundation of China (Grant Nos.\ 11473054 and U1531245) and by the Science and Technology Commission of Shanghai Municipality (Grant No.\ 14ZR1447100). CB was supported by the National Natural Science Foundation of China (Grant No.\ U1531117), Fudan University (Grant No.\ IDH1512060), and the Alexander von Humboldt Foundation. LCH was supported by the National Key R\&D Program of China (2016YFA0400702) and the National Science Foundation of China (11473002, 11721303). This research is based on observations obtained with {\it XMM-Newton}, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. WWZ software is available at URL: https://www.aavso.org/software-directory. | 18 | 8 | 1808.05112 |
1808 | 1808.05438_arXiv.txt | \noindent We present the first results of a pilot program to conduct an ALMA Band~10 spectral line survey of the high-mass star-forming region NGC\,6334I. The observations were taken in exceptional weather conditions (0.19~mm precipitable water) with typical system temperatures $T_{\rm{sys}}$~$<$950~K at $\sim$890~GHz. A bright, bipolar north-south outflow is seen in HDO and CS emission, driven by the embedded massive protostar MM1B. This has allowed, for the first time, a direct comparison of the thermal water in this outflow to the location of water maser emission from prior 22~GHz VLA observations. The maser locations are shown to correspond to the sites along the outflow cavity walls where high velocity gas impacts the surrounding material. We also compare our new observations to prior \emph{Herschel} HIFI spectral line survey data of this field, detecting an order of magnitude more spectral lines (695 vs 65) in the ALMA data. We focus on the strong detections of the complex organic molecule glycolaldehyde (\ce{HC(O)CH2OH}) in the ALMA data that is not detected in the heavily beam-diluted HIFI spectra. Finally, we stress the need for dedicated THz laboratory spectroscopy to support and exploit future high-frequency molecular line observations with ALMA. | \label{intro} \begin{figure*}[b!] \centering \includegraphics[width=0.95\textwidth]{{ngc6334_b10_HDO_4panel}.pdf} \caption{ALMA Band~10 images of the peak (a, b) and integrated intensity (c, d) from $-13.5$ to $+2.5$ \kms\/ for two transitions of HDO with 0.35~mm dust continuum contours overlaid. The 0.35~mm dust continuum contour levels are 50, 75, and 150~K on the Planck temperature scale. Magenta $\times$ symbols show the location of the spectra extracted for MM1 (Fig.~\ref{hifi_comp}). There is substantial absorption of the ground-state transition of HDO toward the MM1 continuum in the integrated intensity maps. This absorption dominates the integrated intensity map, leaving no net emission signal, whereas the peak intensity maps show the locations of this subsumed HDO emission. Primary beam correction has been applied with a cutoff at 0.25 of the FWHM. The synthesized beam of $0\farcs23\times 0\farcs16$ (PA$=39\arcdeg$) is shown in the lower right of each panel.} \label{map} \end{figure*} Observations with the Atacama Large Millimeter Array (ALMA) in Bands 3--7 (84--373~GHz) have proven to be exceptional tools for the detection of new molecular species in the interstellar medium (ISM) and the study of their chemical history and interaction with their physical environment. As a few examples, \citet{Belloche:2014jd} reported the first detection of a branched carbon-chain molecule in the ISM, iso-propyl cyanide (\ce{C3H7CN}), in Band~3 observations of Sgr\,B2. Later, \citet{McGuire:2017gy} detected methoxymethanol (\ce{CH3OCH2OH}) in surprisingly high abundance toward NGC\,6334I in Band 6~and~7 observations, while \citet{Fayolle:2017bg} identified methyl chloride (\ce{CH3Cl}) for the first time using Band~7 observations of IRAS\,16293-2422. These observations, among many others, demonstrate the power of ALMA for studies of our molecular universe in the 1--3~mm wavelength range. Astrochemical observations at higher frequencies, in ALMA Bands 9 (602--720~GHz) and 10 (787--950~GHz) offer complementary benefits to the lower-frequency data, yet few molecular line surveys have been conducted at these frequencies. Here, we explore two advantages of high-frequency spectral line observations. First, the fundamental or first few lowest transitions of many small molecules of interest fall into this range. For example, the HDO $1_{1,1} - 0_{0,0}$ fundamental transition occurs at 893.6~GHz \citep{Messer:1984hh}, providing one of the best opportunities to obtain ground-based measurements of thermal water \citep[see, e.g.,][]{Comito2003}. Second, the transitions of most complex organic molecules (COMs) that fall within this frequency range are typically much higher in energy, providing a robust constraint on excitation conditions within a source. For example, the strongest transitions of glycolaldehyde (\ce{HC(O)CH2OH}), the simplest sugar-related molecule, in Band~6 have upper state energies $E_u$~$\sim$60--200~K in the ground vibrational state. As a result, analysis of the emission of this molecule can be biased toward lower-excitation conditions, although this can be mitigated through the observation of vibrationally excited states in some cases \citep{Jorgensen:2012dw}. Complementary observations at higher frequencies, however, provide access to higher-energy lines to rigorously constrain these excitation temperatures -- the strongest transitions of \ce{HC(O)CH2OH} in Band~10 have $E_u$~=~530-630~K -- and can provide needed confirmatory transitions to secure a lower-frequency detection (see, e.g., \citealt{Jorgensen:2012dw}). These higher-energy transitions also provide selective access to the warmest molecular gas in a source, which prior studies have shown can have substantially different chemistry from the population probed by the lower-energy transitions accessible at lower frequencies \citep{Crockett:2014er}. There is, however, a relative lack of direct laboratory measurements of molecular spectra above $\sim$600~GHz, meaning that many identifications are made from extrapolated quantum mechanical fits. For some species this is a reasonably accurate process, but, as will be shown later, the richness of the Band 10 spectra underscores the need for dedicated high-frequency laboratory work. Finally, at these frequencies it is reasonable to expect that increased dust optical depth effects might ``hide" the deepest, most compact regions of hot molecular cores, and the bright continuum might drive many molecular transitions into absorption. To explore the utility of ALMA Band~9/10 observations, we proposed for a full Band~9 survey, and a pilot Band~10 survey, of the high-mass star-forming region NGC\,6334I in Cycle~5. NGC\,6334I was chosen as the target for three reasons. First, it is an exceptionally molecular line-rich source \citep{McGuire:2017gy} with a relatively small heliocentric distance of 1.3~kpc \citep{Reid:2014km,Chibueze2014}. Second, it has previously been targeted by single-dish observations in overlapping frequency ranges by \citet{Zernickel:2012hx} using the Heterodyne Instrument for the Far-Infrared (HIFI; \citealt{deGraauw:2010gy}) on the \emph{Herschel} Space Observatory \citep{Pilbratt:2010en}. Third, it displays a complex spatial structure consisting of a substantial number of embedded sources and outflows, and several chemically distinct regions separated by only $\sim$2000 au ($\sim$1$\farcs5$; \citealt{Brogan:2016cy,McGuire:2017gy,Bogelund:2018uy}; Figure~\ref{map}). Here, we present a first look at ALMA Band~10 observations of any line-rich source, and discuss the results in the context of both probing favorable transitions of light molecules, and in examining the high-excitation lines of complex organic species. The spectra at these high frequencies are as line-rich as those in the millimeter regime. Contrary to initial expectations, observations of high-mass star-forming regions like NGC\,6334I at ALMA Band 10 do not appear to be substantially hampered by dust opacity, and are in fact generally better suited than previous single-dish facilities such as Herschel. These first-look observations demonstrate the power and versatility of high-frequency observations with ALMA. | We have presented a first look at ALMA Band~10 spectral line survey toward a line-rich source -- the high-mass star-forming region NGC\,6334I -- obtained in exceptional weather conditions. The resulting map shows a bright, bi-polar north-south outflow from the central massive protostar MM1b as traced by both HDO and CS emission. A comparison to archival \emph{Herschel} HIFI data of the source shows the power of spatially resolving underlying substructure with a beam size well-matched to the source, resulting in the unambiguous identification of \ce{CH(O)CH2OH}. A wealth of additional transitions suggest the presence of additional complex molecules that can be identified once high resolution laboratory data are available. | 18 | 8 | 1808.05438 |
1808 | 1808.00501_arXiv.txt | Spectral features in the observed spectra of exoplanets depend on the composition of their atmospheres. A good knowledge of the main atmospheric processes that drive the chemical distribution is therefore essential to interpret exoplanetary spectra. An atmosphere reaches chemical equilibrium if the rates of the forward and backward chemical reactions converge to the same value. However, there are atmospheric processes, such as atmospheric transport, that destabilize this equilibrium. In this work we study the changes in composition driven by a 3D wind field in WASP-43b using our Global Circulation Model, \texttt{THOR}. Our model uses validated temperature- and pressure-dependent chemical timescales that allow us to explore the disequilibrium chemistry of CO, CO$_2$, H$_2$O and CH$_4$. In WASP-43b the formation of the equatorial jet has an important impact in the chemical distribution of the different species across the atmosphere. At low latitudes the chemistry is longitudinally quenched, except for CO$_2$ at solar abundances. The polar vortexes have a distinct chemical distribution since these are regions with lower temperature and atmospheric mixing. Vertical and latitudinal mixing have a secondary impact in the chemical transport. We determine graphically the effect of disequilibrium on observed emission spectra. Our results do not show any significant differences in the emission spectra between the equilibrium and disequilibrium solutions for C/O = 0.5. However, if C/O is increased to 2.0, differences in the spectra due to the disequilibrium chemistry of CH$_4$ become non-negligible. In some spectral ranges the emission spectra can have more than 15$\%$ departures from the equilibrium solution. | \subsection{Background} WASP-43b is a hot Jupiter planet with twice the mass and roughly the same size as Jupiter (\citealt{2011Hellier}). With a semi-major axis of approximately 0.0153 AU, WASP-43b orbits around WASP-43 in just 19.2 hours (\citealt{2012Gillon}). WASP-43 is a K7 star with a mass of 0.73 M$_{\odot}$. The atmosphere of WASP-43b has an equilibrium temperature close to 1440 K (e.g., \citealt{2014Blecic}) and its properties have been studied in previous works. Some of the highlights are the non-detection of thermal inversion on the dayside of the planet (e.g., \citealt{2012Gillon}; \citealt{2013Wang}; \citealt{2014Blecic}; \citealt{2014Stevenson}) and physical constraints on water abundances from \cite{2014Kreidberg}. In this work, we assume WASP-43b to be tidally locked due to its proximity to the parent star, which induces large tidal stresses on the planet (e.g., \citealt{1996Guillot} and \citealt{2014Showman}). The main planet bulk parameters are shown in Table \ref{tab:model}. \cite{2015Kataria} was the first study to explore WASP-43b with a Global Circulation Model (GCM) and to find a reasonable fit of the model dayside emission spectra to the \textit{Hubble Space Telescope }(HST) data from Stevenson et al. 2014. However, in the nightside, the results from \cite{2015Kataria} under-predict the emission fluxes (the nightside is too bright compared to the measured HST data). In \cite{2017Mendoncaa}, we showed that a permanent gray cloud structure in the nightside of WASP-43b is consistent with HST \citep{2014Stevenson} and {\it Spitzer} data (\citealt{2017Mendoncaa}). \cite{2017Mendoncaa} also refers to the possibility that the phase-curve at 4.5 $\mu$m is better fitted by a GCM-generated phase curve that is enhanced in carbon dioxide relative to its chemical equilibrium. Chemical disequilibrium processes occur in the atmosphere mainly due to atmospheric transport and photochemistry processes. However, at high temperatures in the dayside of hot Jupiter planets, the effects of photochemistry become relatively minor compared to thermochemical processes (e.g., \citealt{2011Moses}, \citealt{2012Kopparapu}, \citealt{2014Moses}). Atmospheric dynamics becomes then the dominant process leading to chemical disequilibrium in hot Jupiter atmospheres. This process can be understood in terms of timescales: disequilibrium happens when the fast transport process (by mean motion and eddies) has a shorter timescale ($\tau_{dyn}$) than the chemical timescale ($\tau_{che}$), and the chemical processes are not fast enough to push the chemical constituents back to chemical equilibrium. The chemical timescale refers to the mean characteristic time for production or destruction of a chemical species by local chemical processes. Different scenarios are possible from the competition between dynamical and chemical processes: \begin{itemize} \item $\tau_{dyn}$ $\ll$ $\tau_{che}$. The distribution of the chemical species is dominated by the dynamical transport compared to the localized sources and sinks. In this situation, the dispersive effects of dynamical transport causes the chemical species to be well mixed in the atmosphere \item $\tau_{dyn}$ $\gg$ $\tau_{che}$. Large variability is expected in the distribution of the chemical species, which is kept close to chemical equilibrium (depends solely on local temperature, pressure and elemental abundance). \item $\tau_{dyn}$ and $\tau_{che}$ are the same order of magnitude. This scenario is more difficult to analyse than the two cases above because both chemistry and dynamics have an important role determining the distribution of the chemical species. \end{itemize} The thermochemical equilibrium is maintained in the hot deep regions where $\tau_{che}$ is typically shorter than $\tau_{dyn}$, but at lower pressures the two timescales start to become comparable, and the departures from equilibrium in the atmospheres become larger. When $\tau_{che}$ becomes larger than $\tau_{dyn}$ we say that the chemical component became ``quenched''. A description of the quenching effect was first described in \cite{1977Prinn} to explain the distribution of CO in Jupiter's troposphere, which is associated with a rapid upward mixing from the deeper atmosphere. % \subsection{Motivation} \label{subsec:Motiv} A good understanding of the chemical distribution across the atmosphere is essential to improve our 3D simulations and interpretation of exoplanet atmospheric spectra. Advances on this topic will lead us into a deeper knowledge on how different atmospheric processes control the climate in the planets. Theoretical and observational studies of hot Jovian atmospheres (e.g., \citealt{2009Showman}, \citealt{2009knutson}) show that these atmospheres are very dynamically active with equatorial jets of a few km/s speed, which indicates that a simple thermochemical equilibrium calculation based on local pressure and temperature is not enough to determine its chemical composition across the atmosphere. Studies with different levels of complexity have explored the problem of chemical disequilibrium driven by atmospheric transport. One-dimensional models, for example, include complex kinetic codes with hundreds to thousands of reactions (e.g., \citealt{2012Kopparapu}; \citealt{2011Moses}; \citealt{2015Venot}; \citealt{2014Hu}; \citealt{2016Rimmer}; \citealt{2017Tsai}), but very simplified representations of atmospheric transport (e.g., eddy diffusion parameterisation). To interpret the spectra from observational results it is necessary that we start looking into the impact of three-dimensional atmospheric transport processes, however, the chemical kinetic schemes used in the simple one-dimensional models are too computational expensive to be coupled directly with three-dimensional models (e.g., GCMs). Self-consistent three-dimensional simulations require the computation of the chemical kinetic code for each grid box (usually thousands of boxes covering the entire model domain), over a long time integration until the steady state solution is reached. This latter property is very important in numerical simulations to avoid the final conclusions being influenced by the initial conditions. Despite the complexity, several works have broken through the three-dimensional problem using approaches of intermediate complexity. This hierarchy in model complexity is very important to guide us through the learning process of complex atmospheric mechanisms. Examples are the work of \cite{2012Agundez} and \cite{2014Agundez}. In these two works the complex chemical kinetics model was coupled with a simplified model to represent the dynamics based on a constant solid-body rotation. The prescribed circulation mimicked the broad equatorial jet predicted by theoretical GCMs (e.g., \citealt{2009Showman}; \citealt{2015Heng}). The simplification in the model made it possible to explore conceptually how the atmospheric transport (along the longitude) impacts the chemical distribution across the atmosphere. Another approach was the pioneering work of \cite{2005Cooper, 2006Cooper}, who used a GCM but with a simplified chemical kinetics model. In this work, the complex chemical network is replaced by a simple parameterization called the chemical relaxation method, which consists of a single Newtonian equation where the chemical species are forced towards an equilibrium value. The strength of the source/sink term in the equation depends on a chemical timescale. \cite{2005Cooper, 2006Cooper}, studied the molecular distribution of CO across the HD 209458b planet atmosphere. Their study showed that the vertical winds have an important role in the distribution of CO (vertical quenching). Their chemical relaxation scheme is based on the assumption that a single rate-limiting reaction can describe the interconversion between CO and CH$_4$ over a broad range of pressure and temperature. In \cite{2018Drummond} the simple relaxation method developed in \cite{2006Cooper} was coupled with a nongray radiation scheme. The results show that in HD 209458b the horizontal transport also has an important impact in the final chemical distribution of methane. \cite{2018Tsai} extended and validated the model of \cite{2006Cooper}, and developed a pathway analysis tool that identifies the rate-limiting reaction as a function of temperature, pressure and chemical abundances. \cite{2006Cooper} used essentially a shorter chemical timescale of CO, which was determined by a single rate-limiting reaction. \cite{2006Cooper} and \cite{2018Drummond} only calculate CO using the relaxation method and relate CH$_4$ through mass balance, assuming that all the carbons are locked in either CH$_4$ and CO. We emphasize that this mass balance relation is only valid when the system is in or close to chemical equilibrium and hence not applicable to the relaxation method (see \citealt{2018Tsai}). Our method makes it possible to study chemical disequilibrium of several chemical species using a GCM at very low computational cost. In this work, we apply the formulation developed in \cite{2018Tsai} to study the chemical distribution of CO, CH$_4$, CO$_2$ and H$_2$O on WASP-43b, with our GCM, \texttt{THOR} (\citealt{2016Mendoncab}). This implementation is part of a model hierarchy (such as the others mentioned above) devoted to studying disequilibrium chemistry in the atmosphere of hot Jupiters. The model presented still does not include a self-consistent representation of the chemistry in the radiative transfer code. The representation of the radiation in the GCM simulations are also simplified (gray radiative transfer code) as we explain below. However, our model has an intermediate level of complexity necessary to help us gain physical intuition on the chemical disequilibrium problem. Changes in the atmospheric composition due to atmospheric transport have an impact on the observed atmospheric spectra. Coupling the new relaxation method with \texttt{THOR} allow us to have better knowledge of the distribution of the different chemical species, which is a powerful tool in the interpretation of atmospheric spectra. We also explore and analyse the distribution of different chemical species, however, further detailed analyses of the main dynamical mechanisms transporting these components by, for example, eddies, wave breaking, diffusion, and the impact of different spatial resolutions is beyond the scope of this paper. \subsection{Structure of the present study} In the next section \ref{sec:model}, we explain the theoretical tools, including the GCM and chemical relaxation method used to explore the 3D chemical distribution in WASP-43b. In section \ref{sec:basesimu}, we analyse the results for multiple scenarios integrated with the GCM: with and without thermal inversion in the dayside. In section \ref{sec:des_che}, we study the chemical distribution across the atmosphere for the different species and compare the physical timescales involved to determine the main process controlling the chemical distribution. In section \ref{sec:phcrv_spec}, we compare the emission spectra from equilibrium and disequilibrium chemistry. Final concluding remarks and future prospects are discussed in section \ref{sec:conclu}. | \label{sec:conclu} Consistent modelling of chemistry in exoplanets is still poorly explored. In the current work, we study the disequilibrium chemistry driven by atmospheric circulation using a 3D atmospheric model, which includes a complex representation of the atmospheric flow (solves the deep non-hydrostatic equations) and simplified representations of radiation and chemistry. We have chosen to study the planet WASP-43b and explore its main properties associated with the chemical distribution across the atmosphere. Our GCM solves the deep non-hydrostatic Euler equations and coupled with the radiative transfer code we explored two different scenarios: with and without thermal inversion in the dayside of the upper atmosphere. The main differences between these two simulations are the weaker jet and larger day-night contrast in temperature for the case with thermal inversion. Our new formulation for the chemical timescale developed in \cite{2018Tsai}, allowed us to extend our disequilibrium chemistry study using a GCM to include more molecules than previous works due to the low cost of the parameterization. In the present work we study the chemical distribution of CO$_2$, CO, CH$_4$ and H$_2$O across the atmosphere of WASP-43b. The distribution of the chemical species depends on the strength of the atmospheric mixing and the efficiency of the chemical reactions. Different molecules have different chemical timescales, which result in different patterns of chemical distribution across the atmosphere. In general, the gas molecules are quenched longitudinally in regions near the equator, except for CO$_2$ when $C/O$ is 0.5. CO$_2$ has a shorter chemical timescale than the other molecules for the experiments studied, which resulted in a larger day-night contrast for the abundances of CO$_2$. In contrast with the regions at low latitude the polar vortexes have low levels of dynamical mixing and lower temperatures, which results in a distinct chemical distribution near and inside the polar vortexes. However, if the chemical timescales become very long the chemical species become well mixed over all latitudes and longitudes. The dynamical mixing in the latitudinal and vertical directions is much weaker than in the longitudinal direction, but in the nightside it is more efficient than the chemical production/destruction rates. The main mechanism driving the chemical space distribution in this work for WASP-43b is different than previous works, e.g., \cite{2006Cooper} and \cite{2018Drummond} for HD 209458b. In \cite{2006Cooper} and \cite{2018Drummond} the vertical mixing is the dominant mechanism, however, \cite{2018Drummond} found that the horizontal transport also has an important impact in the final chemical distribution of methane. In our work, we find the transport in the zonal direction (horizontal quenching) to be the dominant process. The disparity is mainly associated with the differences in the dynamical timescales, which can be associated with the different model parameters used: \cite{2006Cooper} and \cite{2018Drummond} explored HD 209458b with a clear atmosphere, and in our work, we explored WASP-43b with a cloudy nightside, which has stronger surface gravity ($\sim$5$\times$ stronger) and faster rotation rate ($\sim$4.4$\times$ faster) than HD 209458b. These results highlight the diversity of solutions and the need to consider 3D dynamical effects when studying disequilibrium chemistry in hot Jupiter planets. The differences in the chemical distribution due to the atmospheric transport have an impact in the emission spectra of the planet. The largest differences exist for the atmosphere at C/O = 2, mainly due to the disequilibrium of CH$_4$. The CH$_4$ molecule has an important impact in the spectral regions between 1.6 to 1.8 $\mu$m, 2.2 to 2.6 $\mu$m and 3 to 4 $\mu$m. The differences due to CH$_4$ become clearly visible in the dayside of the planet. For C/O = 0.5 the atmosphere is dominated by the IR opacity of H$_2$O that is close to equilibrium. Our model is currently being updated with a more sophisticated radiative transfer code capable of computing the feedback due to chemical disequilibrium, and cloud formation. The new \texttt{THOR} will be essential to analyse and interpret the data from \textit{James Webb Space Telescope}, which will have a better spectral resolution and wavelength coverage than current instruments to deal with the problem of disequilibrium chemistry in exoplanet atmospheres. | 18 | 8 | 1808.00501 |
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