subfolder
stringclasses 367
values | filename
stringlengths 13
25
| abstract
stringlengths 1
39.9k
| introduction
stringlengths 0
316k
| conclusions
stringlengths 0
229k
| year
int64 0
99
| month
int64 1
12
| arxiv_id
stringlengths 8
25
|
|---|---|---|---|---|---|---|---|
1404 | 1404.3398_arXiv.txt | We categorically point out why the analysis of Ref.~\cite{Antusch:2014cpa} is incorrect. Here we explicitly show why the sub-Planckian field excursion of the inflaton field can yield large observable tensor-to-scalar ratio, which satisfies both Planck and BICEP constraints. | 14 | 4 | 1404.3398 |
||
1404 | 1404.0022_arXiv.txt | \noindent Motivated by the gamma-ray excess observed from the region surrounding the Galactic Center, we explore particle dark matter models that could potentially account for the spectrum and normalization of this signal. Taking a model-independent approach, we consider an exhaustive list of tree-level diagrams for dark matter annihilation, and determine which could account for the observed gamma-ray emission while simultaneously predicting a thermal relic abundance equal to the measured cosmological dark matter density. We identify a wide variety of models that can meet these criteria without conflicting with existing constraints from direct detection experiments or the Large Hadron Collider (LHC). The prospects for detection in near future dark matter experiments and/or the upcoming 14 TeV LHC appear quite promising. | Over the past several years, a gamma-ray excess from the region surrounding the Galactic Center has been identified in the data of the Fermi Gamma-Ray Space Telescope, with features similar to that expected from annihilating dark matter (DM) particles~\cite{Goodenough:2009gk,Hooper:2010mq,Hooper:2011ti,Abazajian:2012pn,Hooper:2013rwa,Gordon:2013vta,Huang:2013pda,Abazajian:2014fta,Daylan:2014rsa}. Unlike many of the other potential DM signals that have been reported~\cite{Finkbeiner:2004us,Hooper:2007kb,Adriani:2008zr,Weniger:2012tx,Su:2012ft,Bernabei:2008yi,Bernabei:2010mq,Aalseth:2010vx,Aalseth:2011wp,Angloher:2011uu,Agnese:2013rvf,Chang:2008aa}, however, DM interpretations of this gamma-ray excess have become increasingly compelling as the signal has become better measured and characterized. Recent analysis has shown this excess to be robust and highly statistically significant, exhibiting a spectrum and angular distribution that is in good agreement with that expected from the annihilations of $\sim$30 GeV DM particles~\cite{Daylan:2014rsa}. Assuming a DM profile with a local density of 0.3 GeV/cm$^3$, the overall normalization of the signal requires that the DM annihilates with a cross section of $\sigma v \simeq (1.7-2.3)\times 10^{-26}$ cm$^3$/s~\cite{Daylan:2014rsa}, remarkably similar to the value anticipated for a thermal relic~\cite{Steigman:2012nb}. And unlike other astrophysical observations which have received attention as possible detections of DM (the cosmic-ray positron excess, for example~\cite{Hooper:2008kg}), no plausible astrophysical interpretation for the gamma-ray excess has been proposed.\footnote{Although a population of several thousand millisecond pulsars has been discussed as a possible origin of the observed gamma-ray excess~\cite{Hooper:2010mq,Abazajian:2010zy,Hooper:2011ti,Abazajian:2012pn,Hooper:2013rwa,Gordon:2013vta}, the more recent determination that this signal extends to beyond at least $10^{\circ}$ from the Galactic Center~\cite{Daylan:2014rsa,Hooper:2013rwa} strongly disfavors this interpretation~\cite{Hooper:2013nhl}.} In this paper, we attempt to identify the varieties of DM models that could be responsible for the observed gamma-ray excess. Taking a model-independent and bottom-up approach, we construct an exhaustive list of tree-level diagrams for DM annihilation into Standard Model (SM) fermions (see also Ref.~\cite{DiFranzo:2013vra}). By considering tree-level diagrams, instead of effective operators~\cite{Kumar:2013iva,Goodman:2010yf,Bai:2010hh,Goodman:2010ku,Fox:2011fx,Beltran:2010ww,Beltran:2008xg}, we avoid a number of potentially important pitfalls~\cite{Busoni:2013lha,Goodman:2011jq,Buchmueller:2013dya,Shoemaker:2011vi}. For instance, while resonances can be important in determining the annihilation cross section and relic density of the DM, these effects are ``integrated out" in the effective operator approach. By studying the set of tree-level diagrams with all possible combinations of charge- and flavor-conserving renormalizable dimension-four, and super-renormalizable dimension-three, operators compatible with Lorentz invariance, we are able to take a holistic and general view of the types of DM models that could potentially produce the gamma-ray excess observed from the region of the Galactic Center.\footnote{For previous studies which have considered DM models for the Galactic Center excess from an effective field theory perspective, see Refs.~\cite{hooperon,Huang:2013apa}.} For any given model, we impose the following requirements: \begin{enumerate} \item In order to generate the observed spectral shape of the gamma-ray excess, we require that the DM consists of either a $\sim$35 GeV particle that annihilates mostly to $b\bar{b}$ or a $\sim$25 GeV particle that annihilates approximately democratically to SM fermions~\cite{Daylan:2014rsa}. \label{enum:0} \item To accommodate the observed intensity of the gamma-ray excess, we require that the DM annihilates in the low-velocity limit with a cross section of $\langle \sigma v \rangle = (0.77-3.23) \times 10^{-26}$ cm$^3$/s or $\langle \sigma v \rangle = (0.63-2.40) \times 10^{-26}$ cm$^3$/s for the two cases described in Criterion~\ref{enum:0}, respectively~\cite{Daylan:2014rsa}. These ranges take into account the uncertainty in the local DM density~\cite{Iocco:2011jz}. The necessary cross sections are doubled in the case that the DM is not self-conjugate. \label{enum:1} \item We require that the thermal relic density of the DM satisfies $\Omega_{\rm DM} = 0.268^{+0.013}_{-0.010}$, in accordance with measurements from WMAP and Planck~\cite{Ade:2013zuv}. \label{enum:2} \item We require that the elastic scattering cross sections of the DM with nuclei are consistent with the constraints from LUX~\cite{Akerib:2013tjd} and other direct detection experiments.\label{enum:3} \item We require that no constraints from the LHC or other accelerator experiments are violated.~\label{enum:4} \end{enumerate} Criteria~\ref{enum:1} and~\ref{enum:2} roughly correspond to the requirement that the DM is a thermal relic whose annihilations proceed largely through $s$-wave processes. Criterion~\ref{enum:3} roughly requires that any coherent (spin-independent) DM scattering with nuclei must be suppressed, such as by powers of momentum or relative velocity. In evaluating Criterion~\ref{enum:4}, we consider mono-jet and mono-$W/Z$ constraints from the LHC, mono-$b$ projections, as well as accelerator constraints on various classes of particles that might mediate the interactions of the DM. For the DM and its mediator, we consider any combination of spin-0, -1/2, and -1 particles with interactions of the following general forms: \alg{ \label{schematic} \cL_s \supset& \pL {\rm \overline{DM}~DM~mediator} \pR + \pL {\rm \overline{SM}~SM~mediator} \pR, \\ \cL_t \supset& \pL {\rm \overline{DM}~SM~mediator} \pR + \pL {\rm \overline{SM}~DM~mediator} \pR. } These refer to $s$-channel or $t$-~and $u$-channel annihilation diagrams, respectively. We will continue to use this terminology even when talking about elastic scattering processes for which the Feynman diagrams are oriented differently. For the purpose of avoiding ambiguities regarding the labels for the DM and mediating particles, we adopt the conventions shown in Table~\ref{EFT}. We constrain the interactions of the mediator with DM and SM fermions only by the requirement that Lorentz invariance is respected at every vertex. We then consider all allowed combinations of scalar ($1$), pseudoscalar ($\gamma^5$), vector ($\gamma^{\mu}$), and axial ($\gamma^{\mu} \gamma^5$) interactions. We do not attempt to construct an ultraviolet completion for any model, leaving such exploration for future work. \begin{table}[t] \centering \begin{tabular}{| c || c | c | c |} \hline & Scalar & Fermion & Vector \\ \hline \hline DM & $\phi$ & $\chi$ & $X^\mu$ \\ \hline Mediator & $A$ & $\psi$ & $V^\mu$ \\ \hline SM (fermions) & $-$ & $f$ & $-$ \\ \hline \end{tabular} \caption{The particle notation used throughout this study.} \label{EFT} \end{table} The remainder of this article is structured as follows. In Secs.~\ref{fermionsection} and~\ref{scalarsection}, we consider fermionic (spin-$1/2$) and bosonic (spin-0 or spin-1) DM, respectively, annihilating through $s$-channel Feynman diagrams. In each case, we determine which combination of spins and interaction types can satisfy the five criteria described in this section. In Sec.~\ref{tchannel}, we consider cases in which the DM annihilates through the $t$-channel exchange of a colored and charged mediator. In Sec.~\ref{collider}, we discuss constraints from collider experiments on the mass and couplings of the particles that mediate the DM's interactions. In Sec.~\ref{directD}, we discuss the prospects for operating and upcoming direct detection experiments. In Sec.~\ref{conclusions} we summarize our results and conclusions. This paper contains an extensive set of appendices which include, among other information, the full expressions for the DM annihilation and elastic scattering cross sections used in this study. | \label{conclusions} In this study, we have taken a ``simplified model'' approach to determine which classes of dark matter models are capable of producing the gamma-ray excess observed from the region surrounding the Galactic Center. In doing so, we have identified 16 different models that can generate the observed excess without exceeding any of the constraints from direct detection experiments or from colliders (see Table~\ref{goodmodels}). These 16 models can be divided into the following three groups: \begin{itemize} \item{Models in which the dark matter (which could be spin-$0$, $1/2$, or $1$) annihilates through the exchange of a spin-$0$ particle with pseudoscalar interactions. Such a mediator could potentially be observed in future searches for heavy neutral Higgs bosons at the LHC.} \item{Models in which the dark matter is a fermion that annihilates through the exchange of a spin-$1$ particle with axial couplings to standard model fermions, or with vector couplings to third generation standard model fermions. Assuming perturbative couplings, LHC constraints from dijet searches require that the mass of the mediator be {\it less} than $\sim$1 TeV.} \item{Models in which the dark matter annihilates into b-quark pairs through the $t$-channel exchange of a colored and charged particle. Constraints from sbottom searches at the LHC restrict the mediator mass be greater than $\sim$600 GeV. Both LUX and the LHC should be able to conclusively test this class of models in the near future.} \end{itemize} Upon reviewing this list of possibilities, it is clear that a wide range of simple dark matter models could be responsible for the Galactic Center's gamma-ray excess without running afoul of existing constraints. Moreover, the prospects for detecting the dark matter in these scenarios at either direct detection experiments or at the LHC appear to be quite promising. Of the 16 viable models identified in our study, LUX and XENON1T are expected to be sensitive to 7. Only 3 of these 16 models predict an elastic scattering cross section that will remain beyond the reach of future direct detection experiments due to the irreducible neutrino floor. Mono-jet searches, sbottom searches, and searches for heavy Higgs bosons at the LHC will further restrict the range of model parameters that remains viable. With 13-14 TeV data from the LHC, it will be possible to conclusively test several of the scenarios presented here. Many of the results presented in this study nicely illustrate the complementarity between indirect, direct, and collider searches for dark matter. Although future astrophysical observations (such as gamma-ray searches for dark matter annihilating in dwarf galaxies~\cite{Ackermann:2013yva} or future cosmic-ray anti-proton measurements~\cite{Cirelli:2013hv,Fornengo:2013xda}) may provide additional support for a dark matter interpretation of the Galactic Center gamma-ray excess, indirect detection signals alone are expected to determine little more than the mass and annihilation cross section of the particles that make up the dark matter, leaving many questions unanswered. Information from a combination of direct detection experiments and colliders will be needed if one is to identify the underlying interactions and particle content of the dark sector. \bigskip | 14 | 4 | 1404.0022 |
1404 | 1404.7744_arXiv.txt | {The external regions of galaxy clusters may be under strong influence of the dark energy, discovered by observations of the SN Ia at redshift $z<1$.}{The presence of the dark energy in the gravitational equilibrium equation, with the Einstein $\Lambda$ term, contrasts the gravity, making the equilibrium configuration more extended in radius.}{In this paper we derive and solve the kinetic equation for an equilibrium configuration in presence of the dark energy, by considering the Newtonian regime, being the observed velocities of the galaxies inside a cluster largely smaller than the light velocity.}{The presence of the dark energy in the gravitational equilibrium equation leads to wide regions in the $W_0$- $\rho_\Lambda$ diagram where the equilibrium solutions are not permitted, due to the prevalence of the effects of the dark energy on the gravity.}{} | It was shown by \cite{2001Chernin,2008Chernin} that outer parts of galaxy clusters (GC) may be under strong influence of the dark energy (DE), discovered by observations of SN Ia at redshift $z\le 1$ \citep{1998Riess,1999Perlmutter}, and in the spectrum of fluctuations of CMB radiation \citep[see e.g.][]{2003Spergel,2004Tegmark}. Equilibrium solutions for polytropic configurations in presence of DE have been obtained in papers of \cite{2006Balaguera,2007Balaguera} and \cite{2012Merafina}. Here we derive a Boltzmann-Vlasov kinetic equation in presence of DE, in Newtonian gravity, and obtain its solutions. These solutions generalize the ones, obtained by \cite{1993Bisnovatyi,1998Bisnovatyi} for the kinetic equation without DE. Here we consider the problem in Newtonian approximation, because the observed chaotic velocities of galaxies inside a cluster are much less than the light velocity $v_{gc}\ll c$. The general relativistic solution in presence of DE could be applied for equilibrium configurations of point masses from some exotic particles, interacting only gravitationally. On early stages of the universe expansion, before and during the inflation stage, these particles may form gravitationally bound configurations which collapse during the inflation, when antigravity decreases. As a result of this collapse such hypothetical objects may be transformed into primordial black holes appearing after the end of inflation. The relativistic kinetic equation, and its solutions in presence of DE will be considered elsewhere. | We have calculated the equilibrium configurations of Newtonian clusters with a truncated Maxwellian distribution function, in presence of DE. All these clusters have a structural equilibrium, being satisfied the condition $R<R_{\Lambda}$, and result dynamically stable. By considering the relaxation time, we obtain a value larger than the age of the Universe and, therefore, we can conclude that thermodynamical instabilities are not relevant in current evolution of galaxy clusters. On the other hand, the evaluation of parameters characterizing this kind of clusters suggests that these systems are collisionless. In any case, the critical point of the onset of thermodynamic instability lies far from the first maximum mass, at larger values of $W_0$ in the curve with $\hat{\rho}_{\Lambda}=0$ of Fig.\,\ref{fig4} \citep{2006bkm}, as well in curves with $\hat{\rho}_{\Lambda}\neq 0$, the most part of the equilibrium configurations results thermodynamically stable. Presently the density distribution inside galaxy clusters is described by several phenomenological functions, some if which follow from numerical simulations \citep[see][]{2013chernin}. Qualitatively the truncated Maxwellian distribution, considered here is similar to the non-singular density distribution suggested by \cite{2013chernin}. It may be used for the more detailed study of the density and velocity distributions on the periphery of rich clusters, where the influence of DE is important, and their comparison with observations. The galaxies in the outer parts of the clusters are not numerous, and they have smaller masses and luminosities in presence of even weaker relaxation. Therefore, only largest telescopes should be used for a search of galaxies at the cluster peripheries. The most sensitive X-ray telescopes are needed for detection of the hot gas in the outer parts of the clusters, and its possible outflow in presence of DE, considered by \cite{2013bkm}. | 14 | 4 | 1404.7744 |
1404 | 1404.2944_arXiv.txt | Spatially homogeneous universes can be described in (loop) quantum gravity as condensates of elementary excitations of space. Their treatment is easiest in the second-quantised group field theory formalism which allows the adaptation of techniques from the description of Bose--Einstein condensates in condensed matter physics. Dynamical equations for the states can be derived directly from the underlying quantum gravity dynamics. The analogue of the Gross--Pitaevskii equation defines an anisotropic quantum cosmology model, in which the condensate wavefunction becomes a quantum cosmology wavefunction on minisuperspace. To illustrate this general formalism, we give a mapping of the gauge-invariant geometric data for a tetrahedron to a minisuperspace of homogeneous anisotropic 3-metrics. We then study an example for which we give the resulting quantum cosmology model in the general anisotropic case and derive the general analytical solution for isotropic universes. We discuss the interpretation of these solutions. We suggest that the WKB approximation used in previous studies, corresponding to semiclassical fundamental degrees of freedom of quantum geometry, should be replaced by a notion of semiclassicality that refers to large-scale observables instead. | There is by now a variety of approaches to the problem of quantum gravity which are actively pursued \cite{danielebook}. Research into any of these directions generally addresses one of two basic aims. The first is to show that a proposed theory of quantum gravity is in itself consistent and that its objects are mathematically well-defined, computable, and can be translated into observable quantities, so that the theory can, at least in principle, be confronted with experiment. The second aim is a derivation of the phenomenology of the theory, which usually requires taking a `low-energy' or `semiclassical' regime, in which the theory should at least be consistent with present observational constraints on deviations from the predictions of general relativity and the standard model of particle physics. It is then often claimed that any genuine quantum-gravitational effect, going beyond separate predictions of general relativity or the standard model, would be intrinsically unobservable, since the Planck scale is many orders of magnitude above the energy scales probed in particle accelerators or hypothetical experiments. However, while it is indeed difficult to come up with present-day experiments that probe Planck-scale physics (for some efforts in this direction, see \cite{giovanni}), the very early universe provides a natural laboratory in which quantum gravity effects can be expected to play a role. Inflation, the standard paradigm for the physics of the very early universe, has been spectacularly corroborated in the recent observations made by Planck \cite{planck} and BICEP2 \cite{bicep}. However, despite its phenomenological success, there are several theoretical issues that remain open: the inflaton and its potential are not part of the standard model, and have to be added by hand. While inflation provides a picture in which the physics at the Big Bang singularity is not observationally relevant today, as its imprint has been stretched outside the causal horizon during the accelerated expansion, theorems such as \cite{theorem} show that inflationary spacetimes have a past singularity, so that there is still a need for a more complete theory. Eternal inflation seems to have drastic and contentious theoretical consequences \cite{eternal}. Observationally, the BICEP2 results seem to imply a violation of the Lyth bound \cite{lyth}: the inflaton field presumably varies over super-Planckian scales during inflation. All of this motivates the study of quantum-gravitational models with regard to their predictions for cosmology. The spacetimes relevant for cosmology are to a very good approximation spatially homogeneous. One can use this fact and perform a symmetry reduction of the classical theory (general relativity coupled to a scalar field or other matter) assuming spatial homogeneity, followed by a `quantisation' of the reduced system. Inhomogeneities are usually added perturbatively. This leads to models of quantum cosmology \cite{admrefs} which can be studied on their own, without the need for a manageable full theory of quantum gravity. While this approach can be pursued with profit to some extent, and is claimed to make potentially observable predictions \cite{kiefercmb}, there is no unambiguous interpretation of calculations that supposedly result from truncation of an unknown underlying theory. For instance, since one is generally ignorant about the physical inner product in full quantum gravity, the predictive power of computing wavefunctions is not clear. Loop quantum gravity (LQG) has some of the structures one would expect in a full theory of quantum gravity: kinematical states corresponding to functionals of the Ashtekar--Barbero connection can be rigorously defined, and geometric observables such as areas and volumes are well-defined as operators, typically with discrete spectrum \cite{LQGbook}. Using the LQG formalism in quantising symmetry-reduced gravity leads to loop quantum cosmology (LQC) \cite{LivRev}. Because of the structures of LQG, LQC allows a rigorous analysis of issues that could not be addressed within the Wheeler--DeWitt quantisation of conventional quantum cosmology, such as a definition of the physical inner product. Recently, LQC has made contact with CMB (cosmic microwave back-ground) observations, as the usual inflationary scenario is now discussed in LQC \cite{abhayagullo}. One missing ingredient in the formalism of LQC is its embedding into the full setting of LQG. Just as in conventional quantum cosmology, one has performed a symmetry reduction before quantisation, and truncated almost all degrees of freedom present in the full Hilbert space of LQG. A different approach aiming at a more complete picture would be to work within the full Hilbert space, identify states that can represent macroscopic, (approximately) spatially homogeneous universes, and extract information about their dynamics. Clearly, this last step will involve many approximations, but since these are approximations for equations of the full theory, one has some control about the error made. Already the identification of suitable states that represent cosmological spacetimes is challenging in a theory like LQG: because of the notion of background independence built into the definition of the theory, the most natural notion of vacuum state is the `no space' state, which has zero expectation value for geometric observables (areas, volumes, etc). Elementary excitations over this vacuum are usually interpreted as distributional geometries, and a macroscopic nondegenerate configuration is unlikely to be found as a small perturbation of this vacuum. A new approach towards addressing the issue of how to describe cosmologically relevant universes in (loop) quantum gravity was recently proposed in \cite{PRL,JHEP} \footnote{For some alternative approaches towards the same problem, see \cite{uther}.}. This proposal uses the group field theory (GFT) formalism, itself a {\em second quantisation} formulation of the kinematics and dynamics of LQG \cite{2ndq}: one has a Fock space of LQG spin network vertices (or tetrahedra, as building blocks of a simplicial complex), annihilated and created by the field operator $\hat\varphi$ and its Hermitian conjugate $\hat\varphi^{\dagger}$, respectively. The advantage of using this reformulation is that field-theoretic techniques are available, as a GFT is a standard quantum field theory on a curved (group) manifold ({\em not} to be interpreted as spacetime). In particular, one can define {\em coherent} or {\em squeezed} states for the GFT field, analogous to states used in the physics of Bose--Einstein condensates or in quantum optics; these represent {\em quantum gravity condensates}. They describe a large number of degrees of freedom of quantum geometry in the same microscopic quantum state, which is the analogue of homogeneity for a differentiable metric geometry. This idea was made explicit in \cite{JHEP}: after embedding a condensate of tetrahedra into a smooth manifold representing a spatial hypersurface, one shows that the spatial metric (in a fixed frame) reconstructed from the quantum state is compatible with spatial homogeneity. As the number of tetrahedra is taken to infinity, a continuum homogeneous metric can be approximated to a better and better degree. At this stage, the condensate states defined in this way are kinematical. They are gauge-invariant (locally Lorentz invariant) by construction, and represent geometric data invariant under (active) spatial diffeomorphisms, but they do not satisfy any dynamical equations corresponding to a Hamiltonian constraint in geometrodynamics. The strategy followed in \cite{PRL, JHEP} for extracting information about the dynamics of these states is the use of Schwinger--Dyson equations of a given GFT model. These give constraints on the $n$-point functions of the theory evaluated in a given condensate state (approximating a non-perturbative vacuum), which can be translated into differential equations for the `condensate wavefunction' used in the definition of the state. Again, this is analogous to condensate states in many-body quantum physics, where such an expectation value gives, in the simplest case, the Gross--Pitaevskii equation for the condensate wavefunction. The truncation of the infinite tower of such equations to the simplest ones is part of the approximations made. As argued in \cite{PRL, JHEP}, the effective dynamical equations thus obtained can be viewed as defining a {\em quantum cosmology model}, with the condensate wavefunction interpreted as a quantum cosmology wavefunction. This provides a general procedure for deriving an effective cosmological dynamics directly from the underlying theory of quantum gravity. In a specific example, it was shown how a particular quantum cosmology equation of this type, in a semiclassical WKB limit and for isotropic universes, reduces to the classical Friedmann equation of homogeneous, isotropic universes in general relativity. The purpose of this paper, apart from reviewing the formalism introduced in detail in \cite{JHEP}, is to analyse more carefully the quantum cosmological models derived from quantum gravity condensate states in GFT. In particular, the formalism identifies the gauge-invariant configuration space of a tetrahedron with the minisuperspace of homogeneous (generally anisotropic) geometries. We will justify this interpretation and propose a convenient set of variables for the gauge-invariant geometric data, which can be mapped to the variables of a general anisotropic Bianchi model (for which the metric is not diagonalised and has six components). We will then revisit the example that led to the Friedmann equation in \cite{PRL, JHEP} and study it directly as a quantum cosmology equation, without a WKB limit. The Friedmann equation arising in a WKB limit in \cite{PRL} appeared to have no solutions, as there was a mismatch between the curvature of the gravitational connection, assumed to be small on the scale of the tetrahedra, and the spatial curvature term which was large on the same scale. Here we find simple solutions to the full quantum equation, corresponding to isotropic universes. They can only satisfy the condition of rapid oscillation of the WKB approximation for large positive values of the coupling $\mu$ in the GFT model. For $\mu<0$, states are sharply peaked on small values for the curvature, describing a condensate of near-flat building blocks, but these do not oscillate. This supports the view that rather than requiring semiclassical behaviour at the Planck scale, semiclassicality should be imposed only on large-scale observables. | Condensate states in group field theory can be used to derive effective quantum cosmology models directly from a proposal for the dynamics of a quantum theory of discrete geometries. We have illustrated the interpretation of the configuration space of gauge-invariant geometric data of a tetrahedron, the domain of the condensate wavefunction, as a minisuperspace of spatially homogeneous 3-metrics. The approach taken here is very different from the more conventional one of quantising only classical degrees of freedom that remain after imposing a symmetry. It makes assumptions about the approximate form of a fully dynamical quantum gravity state, similar to the assumptions one makes when treating interacting quantum systems in condensed matter physics. The validity of these assumptions can be verified; for instance, one can compute fluctuations around the mean field given by the condensate wavefunction $\sigma(g_I)$ and see whether they remain small. From the classical interpretation of the geometric data associated to these states, one expects that they are good approximations as long as the curvature remains small on the scale of the tetrahedra \cite{JHEP}. Again, this assumption can be verified by analysing the effective quantum cosmology equations, and in \cite{PRL, JHEP} there seemed to be a tension as the WKB approximation indicated that the curvature was peaked at large (presumably Planckian) values. This was our main motivation for revisiting the model of \cite{PRL, JHEP}. We found that a more consistently derived WKB approximation has a second solution corresponding to flat universes. Then, rather than assuming semiclassical properties for the condensate wavefunction, we gave simple isotropic solutions to the full quantum cosmology equation. These solutions depend on the `mass parameter' $\mu$. For negative $\mu$, they violate the assumptions of the WKB approximation, but are peaked on small curvature $p\ll 1$ and thus consistent with the expectations of the classical picture, as well as with the Friedmann equation $p\approx 0$. For positive $\mu$ states show a wider distribution of curvature. The effective Friedmann equation (\ref{fried}), as discussed in \cite{PRL, JHEP}, came out of the WKB approximation of (\ref{effeq}) with (\ref{lapl}). Once it is accepted that a WKB-type condensate wavefunction may not give a physically relevant approximation to the dynamics, one might ask whether (\ref{lapl}) still corresponds to an interesting model of quantum cosmology. The explicit examples given in the paper show that this depends strongly on the value of the `mass parameter' $\mu$ in the fundamental theory; this parameter dropped out in the WKB limit. For negative $\mu$ solutions are strongly peaked near $p=0$, and (\ref{lapl}) implements the Friedmann equation $p=0$ describing a pure vacuum, spatially flat universe. In any case, (\ref{lapl}) remains a useful example to consider, because it is simple enough for explicit solutions to be constructed, so that the physical interpretation of GFT condensates and their cosmology can be discussed. Further work will be required to conclusively answer whether the model can reproduce some features of general relativity. In the reduction to isotropic states, the model we have studied is fully constrained: there is only one degree of freedom (essentially given by the scale factor) and one constraint. One could add anisotropies or include matter degrees of freedom into the model. For instance, a massless scalar field can be introduced \cite{JHEP} by taking \ben \hat\mathcal{K}=\sum_{I=1}^4 \Delta_{g_I}+\tau\,\frac{\partial^2}{\partial\phi^2}+\mu \label{knew} \een for an extended GFT model with a field on $G^4\times\bR$ where $\bR$ parametrises the scalar field. Adopting this prescription and decomposing $\sigma(p,\phi)=\sum_{\omega}\sigma_\omega(p)e^{\im\omega\phi}$, general isotropic solutions would be superpositions of the solutions given above with $\hat\mu=\frac{1}{4}(\mu-\tau\omega^2)$, and one could try to construct wavepackets similar to \cite{kiefer}. Requiring these to be composed out of rapidly oscillating modes would require choosing $\tau<0$ and/or a restriction on the values of $\omega$, depending on the value of $\mu$. The physical meaning of these conditions and of the choice (\ref{knew}) from the viewpoint of quantum gravity is however rather unclear. We conclude that the criterion of {\em semiclassicality} for condensate states describing quantum cosmology has to be phrased more carefully to justify results such as an effective Friedmann equation (\ref{fried}) that can support the potential usefulness of the choice (\ref{lapl}) for quantum cosmology. It is only for large-scale observables, such as the total volume (of the universe), that semiclassical behaviour is required. The condensate wavefunction itself captures the properties of what presumably describes a highly quantum-mechanical many-particle state of Planck-scale objects. It carries much more information than a usual quantum cosmology wavefunction, {\em e.g.} about correlations between different quanta or about the scaling of geometric observables with the particle number. Using this information will be necessary for adding inhomogeneities \cite{JHEP}, and for potentially making contact with CMB observations. All of this motivates further systematic studies of the quantum cosmology of (loop) quantum gravity condensates. | 14 | 4 | 1404.2944 |
1404 | 1404.5755_arXiv.txt | We present a new determination of the solar fluorine abundance together with abundance measurements of fluorine in two Galactic open clusters. We analyzed a sunspot spectrum, observed by L. Wallace and W. Livingston with the FTS at the McMath/Pierce Solar Telescope situated on Kitt Peak and spectra of four giants in the old cluster M~67 ($\sim$4.5 Gyr) and three giants in the young cluster NGC~6404 ($\sim$0.5 Gyr), obtained with the CRIRES spectrograph at VLT. Fluorine was measured through synthesis of the available HF lines. We adopted the recent set of experimental molecular parameters of HF delivered by the HITRAN database, and found a new solar fluorine abundance of $A(F) = 4.40\pm 0.25$, in good agreement with the M~67 average fluorine abundance of $A(F) = 4.49\pm 0.20$. The new solar abundance is in a very good agreement with the meteoritic value. The used modern spectrosynthesis tools, the agreement with the meteoritic value and with the results in open cluster M67, known to be a solar analogue, make our solar determination very robust. At the same time, the fluorine measurement in the above-mentioned open clusters is the first step in the understanding of its evolution during the last $\sim$10 Gyr in the Galactic disk. In order to develop this project, a larger sample of open clusters is required, so that it would allow us to trace the evolution of fluorine as a function of time and, in turn, to better understand its origin. | The origin and the evolution of fluorine in the Galaxy are still nowadays a matter of debate. The available observational constraints, coupled with stellar nucleosynthesis models, have not yet clarified which stellar mass ranges and in which evolutionary stages are the mainly responsible for the fluorine production. Therefore, further and new observational evidence is needed, to understand where fluorine is produced and its implications on the stellar nucleosynthesis and Galactic chemical evolution. The state of the art proposes three means of fluorine production: neutrino spallation on $^{20}$Ne in gravitational supernovae (SNII; Woosley \& Haxton 1988), hydrostatic nucleosynthesis in the He-burning core of heavily mass-losing Wolf-Rayet (WR) stars (Meynet \& Arnould 2000), and hydrostatic nucleosynthesis in the He-rich intershell of thermally pulsing (TP) AGB stars (Busso et al. 1999). It is still unknown which of the three above sources is the main contributor for fluorine. The above scenario is based on several observational studies that, during the last decade, have been addressed the problem of the fluorine origin and evolution. Fluorine determinations were carried out in different environments: namely: i) in the Large Magellanic Cloud (LMC) \citep{cun03}; ii) the globular cluster M~4 \citep{smi05} and M 22 \citep{dor13}; iii) the Milky Way Bulge \citep{cun08}; iv) pre-main sequence stars of the Orion nebula cluster \citep{cun05} and dwarf stars of the solar neighbourhood \citep{rec12}; v) Galactic and extragalactic asymptotic giant branch (AGB) stars \citep{abi09,abi10,abi11,utt08}; vi) in one hot post-AGB star \citep{wer05}; vii) in C-Rich low-metallicity stars \citep{luc11}; viii) in planetary nebulae \citep{zha05}; ix) in the interstellar medium surrounding Type II supernovae \citep{fed05}. These recent studies enlarged and in some cases reanalyzed the sample of stars presented in \citet{jor92}.% Almost all the above fluorine analyses have been developed using spectral features of the HF molecule (mostly the R 9 line at $\lambda_{vacuum}$=2336.47 nm). A \textit{theoretical} list of HF molecular parameters (e.g. log\textit{gf}, E$_{low}$), provided by R.H.\ Tipping (see e.g.\ Abia et al.\ 2009), was in general adopted, together with an old solar abundance derived by \citet{hal69}, A(F)$_{\odot}$=4.56. Very recently, a new list of \textit{experimental} molecular parameters for the HF molecule has been delivered by the HITRAN database \citep[see][for details on this database]{rot13}. Therefore, we started a new analysis of the fluorine abundance based on these new data. More in detail: i) we reanalyzed the solar fluorine abundance as observed in sunspot spectra with modern techniques: sunspot specific spectrosynthesis simulations and atmospheric model. ii) We collected, for the first time, spectra of giant star members of two Galactic open clusters (OCs), M~67 and NGC~6404, in the infrared region, where HF lines were detected and analyzed. The solar fluorine abundance is used as a zero-point for all the other dedicated studies, hence a redetermination in light of the above new molecular parameters and of more recent analysis techniques was needed \citep[see also][]{asp09}. On the other hand, this work is the starting point of a new project which consists in the fluorine determination in several open clusters, with different ages and Galactocentric distances (R$_{GC}$). In fact, this investigation offers the opportunity of measuring fluorine evolution during the last $\sim$10 Gyr as a function of time. This can be done since the age estimate of OCs can be performed with a smaller uncertainty than for field stars. In turn, the knowledge of the fluorine evolution provides also a further constraint to understand which kind of stars (low mass or more massive stars) is mostly responsible to its production. At the same time, the analysis of the M~67 cluster is also a good test of the solar fluorine determination, since this cluster shows a solar-like abundance distribution and its age, metallicity and R$_{GC}$ resemble those of the Sun, so that it can be considered as a \textit{solar analogue} cluster. In Section 2 we describe the observations and the analysis for the fluorine determination in the Sun, while Section 3 is focused on fluorine in open clusters. Section 4 shows our results, while in Section 5 and 6 we discuss results, giving our final conclusions. | \begin{itemize} \item We derived a new solar fluorine abundance in the spectral atlas of \citet{wal01} of a medium strong sunspot umbra. We used experimental molecular data from the HITRAN database for the HF lines and modern spectrosynthesis tools, taking into account the magnetic field of the sunspot: our result is A(F)$_{\odot}$=4.40$\pm$0.25. \item We collected new spectra in the infrared region, with the CRIRES spectrograph, for 7 giant stars of two open clusters: M~67 and NGC~6404. \item We derived fluorine abundances for the observed stars using: i) stellar parameters derived with photometric calibrations, and ii) the synthesis of their spectra. Uncertainties due to CN blends and stellar parameters were evaluated. The total error was estimated to be $\pm$0.20 dex. The abundance in M 67 is in a very good agreement with our new solar estimate, while fluorine in the younger OC NGC 6404 is $\sim$0.1 dex higher than the value in M 67. Looking at [F/Fe] ratios, we found solar values in the two analyzed OCs. \item Future studies of fluorine in several other open clusters with different ages and located at different R$_{GC}$, will allow us to trace its evolution in different zones of the Galaxy. It will also show the relevance of the AGB contribution to the synthesis of fluorine, improving our understanding of the its origin. \end{itemize} | 14 | 4 | 1404.5755 |
1404 | 1404.2205_arXiv.txt | {We aim at unveiling the observational imprint of physical mechanisms that govern planetary formation in young, multiple systems. In particular, we investigate the impact of tidal truncation on the inner circumstellar disks. We observed the emblematic system GG~Tau at high-angular resolution: a hierarchical quadruple system composed of low-mass T~Tauri binary stars surrounded by a well-studied, massive circumbinary disk in Keplerian rotation. We used the near-IR 4-telescope combiner PIONIER on the VLTI and sparse-aperture-masking techniques on VLT/NaCo to probe this proto-planetary system at sub-au scales. We report the discovery of a significant closure-phase signal in $H$ and $K_s$ bands that can be reproduced with an additional low-mass companion orbiting GG~Tau~Ab, at a (projected) separation $\rho = 31.7\pm 0.2$\,mas (4.4\,au) and $PA = 219.6 \pm 0.3^\circ$. This finding offers a simple explanation for several key questions in this system, including the missing-stellar-mass problem and the asymmetry of continuum emission from the inner dust disks observed at millimeter wavelengths. Composed of now five co-eval stars with $0.02 \le M_{\star} \le 0.7 $\,M$_{\odot}$, the quintuple system GG~Tau has become an ideal test case to constrain stellar evolution models at young ages (few $10^6$\,yr). } | Planet formation is a common process that can occur in different environments. While the first decade of planet searches has preferentially focused on single, solar-like host stars, it has been more recently shown that a large proportion of extrasolar giant planets are born in binary systems \citep{udry07}. The recent diskoveries of transiting circumbinary planets in close binary systems \citep[Kepler 16, 34,35, ][]{doyle11, welsh12}, as well as the planet candidate directly imaged around the young low-mass binary 2MASS J0103 \citep{delorme13} have proven that planets can also appear in a circumbinary disk, despite the strong dynamical mechanisms that shape the disk and can rapidly clear out its inner region. Stars in young binary systems are expected to be surrounded by two inner disks, located inside the Roche lobes and an outer circumbinary ring or disk outside the outer Lindblad resonances \citep[e.g., ][]{artymowicz94}. Persistent signs of accretion in binary systems, as well as direct imaging of residual gas in the inner region, demonstrate that gas and dust can flow from the outer reservoir through this gravitationally unstable zone to nurture inner circumstellar disks (where planet formation may also occur), which otherwise would not survive. Understanding how the inner disks are replenished is also important in the general context of planetary system formation, since binary stars provide a scaled-up version of a proto-planet environment in a circumstellar (CS) disk. Finally, mutiple systems can provide essential clues in testing stellar evolution models, as they provide a set of co-eval stars with different masses at a common distance \citep[e.g,][]{white99} In the past two decades, the young hierarchical quadruple system GG~Tau, composed of two low-mass binary systems, has been subject of many detailed studies. With its relatively massive (0.15\,M$_{\odot}$) and bright outer ring, \object{GG~Tau~A} is one of the best known nearby (140\,pc) T~Tauri binaries, with a 0.26\arcsec\ separation ($36$\,au on the sky plane). The circumbinary disk ($R_{\rm in} =180$\,au) has been observed in thermal dust emission \citep{dutrey94, guilloteau99,pietu11} and in scattered light \citep{roddier96,silber00}, and is in Keplerian rotation \citep{dutrey94}. The scattered-light images proved that the gravitationally unstable zone is not empty of dust. Indirect evidence for gas flow from the ring towards the inner system(s) has been found from $^{12}$CO J=2-1 gas image \citep{guilloteau01} and from near-IR H$_2$ transitions \citep{beck12}. The warm H$_2$ gas may be heated by shocks, as material from the circumbinary ring is accreted onto material close to the stars. The existence of inner CS disks is independently attested by mm excess emission on GG~Tau~Aa \citep{pietu11}, strong H$_\alpha$ accretion signature separately detected around Aa and Ab, [OI] line detection around Ab \citep{white99,hartigan03}, and 10$\mu$m silicate feature from hot grains in both Aa and Ab environments \citep{skemer11}. We have recently undertaken a very high-spatial resolution observing program of GG~Tau~A, from UV to mm wavelengths. In this letter, we report near-IR VLT interferometric observations of the inner region of the GG~Tau~A system, where we detect a new component and direct evidence for resolved circumstellar dust emission. \begin{figure*} \centering \includegraphics[width=5.5cm,angle=0]{Fig_chi2map_Aa-Ab1_3s_H6+7_medium.pdf} \includegraphics[width=5.5cm,angle=0]{Fig_chi2map_Aa-Ab2_3s_H6+7_medium.pdf} \includegraphics[width=5.35cm,angle=0]{Fig_chi2map_Ab_2punct_CP-2012red2_large.pdf} \caption{Chi2 maps for: NACO-SAM $H$-band closure-phase data (6+7 Dec. 2012): location of GG~Tau Ab1 (a) and Ab2 (b) around the primary star GG~Tau Aa, and for PIONIER (c) $H$-band closure-phase data for Ab2 location around Ab1, with VLTI $(u,v)$ sampling in the inset. } \label{fig_chi2maps} \end{figure*} | The emblematic, young binary system GG~Tau A (0.26\arcsec separation) has been successfully observed with interferometric and AO\,/\,sparse-aperture-masking techniques (SAM) at the Very Large Telescope. We found that: \begin{enumerate} \item The secondary GG~Tau Ab is itself a close binary, with a projected separation of $0.032\arcsec$ (or 4.5\,au) and $PA=220^\circ$ (end 2012). It is consistent with a M3V (Ab2) and M2V (Ab1) low-mass binary. This finding solves the discrepancy between the dynamical stellar mass derived from CO gas kinematics \citep{guilloteau99} and the most recent spectral-type estimate of Aa and Ab \citep{hartigan03}. Based on a tentative {\it a posteriori } identification in archival (2003) VLT/NACO images, its orbital period is estimated to $P_{\rm Ab1-Ab2} \sim16$\,yr, a value consistent with the period derived from the binary mass and separation. \item All stars in this triple system present significant IR excesses, confirming the presence of circumstellar material. Around GG Tau Aa, the NIR emission is partly resolved at 1.65\micron\ and the derived geometrical ring radius is typical of proto-planetary disks around low-luminosity stars. For GG Tau~Ab, due to tidal truncation, a (deprojected) separation of $\sim5.1$\,au sets a strong constraint on the maximum radial extent of any circumstellar disk surrounding Ab1 and/or Ab2 (\Rout\,$\simless 2$\,au). The binary nature of this system also provides a simple explanation to the intriguing non detection of mm continuum emission at the location of Ab. \item With five coeval low-mass stars, this young multiple system becomes an ideal test case to constrain evolutionary models, provided that future astrometric studies will refine the stars physical parameters of the Ab system. \end{enumerate} | 14 | 4 | 1404.2205 |
1404 | 1404.3804_arXiv.txt | We show the possibility that the observational results of the primordial gravitational waves from Planck and BICEP2 for the tensor-to-scalar ratio $r$ can be reconciled when an early dark energy was included. This early dark energy behaves like a radiation component at very early epoch. This is equivalent to induce additional number of effective neutrino species: $\Delta N_{eff}=[\frac{7}{8}(\frac{4}{11})^{4/3}]^{-1}\rho_{de}(a)/\rho_{\gamma}(a)$, where $\rho_{\gamma}(a)$ is the photon energy density and the numerical factors arise from converting to effective neutrino species. And $\rho_{de}(a)$ is the energy density of early dark energy. Combining the Planck temperature data, the WMAP9 polarization data, and the baryon acoustic oscillation data with and without BICEP2 data, we find that in this early dark energy model the tension between the observations from Planck and BICEP2 was relived at $2\sigma$ regions. But it cannot be removed completely due to the small ratio of early dark energy constrained by the other cosmic observations. As a byproduct, the tension between observed values of Hubble parameter from Planck and the direct measurement of the Hubble constant was removed in this early dark energy model. | The Background Imaging of Cosmic Extragalactic Polarization (BICEP2) experiment \cite{ref:BICEP21,ref:BICEP22} has detected the B-modes of polarization in the cosmic microwave background, where the tensor-to-scalar ratio $r=0.20^{+0.07}_{-0.05}$ with $r=0$ disfavored at $7.0\sigma$ of the lensed-$\Lambda$CDM model was found. However, combining with WMAP9 polarization data, ACT and SPT, Planck group reported a much smaller tensor-to-scalar ratio, compared to that from BICEP2 $r<0.11$ at $95\%$ C.L. in the $\Lambda$CDM+$r$ model \cite{ref:Planck2013tensor}. It apparently shows the tension between these two observations. To relieve this tension, some possible extensions to the $\Lambda$CDM+$r$ model have been discussed such as the ruining spectral index \cite{ref:BICEP21} introduced by BICEP2 (see also in Refs. \cite{ref:runMa,ref:runCzery}), the additional relativistic degrees of freedom beyond the three active neutrinos and photons \cite{ref:Zhang,ref:Hu,ref:Li,ref:Anchordoqui,ref:Ko}, the suppression of the adiabatic perturbations on large scales \cite{ref:Contaldi,ref:Miranda,ref:Abazajian,ref:Hazra}, the blue tilted tensor \cite{ref:Gerbino,ref:Ashoorioon,ref:Wu2014}, and the isocurvature mode \cite{ref:Kawasaki1,ref:Kawasaki2} and so on. In this paper, we take a somewhat related approach to investigate the possibility of relaxing the tension by considering early dark energy which mimics radiation at early epoch. Therefore, effective neutrino species \cite{ref:edeLinder} \begin{equation} \Delta N_{eff}=\left[\frac{7}{8}\left(\frac{4}{11}\right)^{4/3}\right]^{-1}\rho_{de}(a)/\rho_{\gamma}(a)\approx4.4032\frac{\Omega_{de}(a)}{\Omega_{\gamma}(a)}, \end{equation} were induced, where $\rho_{\gamma}(a)$ and $\rho_{de}(a)$ are the photon energy density and dark energy density respectively. If dark energy contributes to the energy budge about $\Omega^{e}_{de}=0.10$ at early time, there is equivalently a half effective neutrino, when one assumes two main components: radiation+early dark energy. This effective neutrino would modify the sound horizon at recombination, then the inferred distance from both CMB and BAO would be changed. Therefore, it is helpful to relax the tension between that with local measurement of the Hubble parameter. We investigate the effects on the tensor-to-scalar ratio and tension relaxation in a quantitative way by performing Markov chain Monte Carlo (MCMC) analysis by using Planck, BAO and BICEP2 data. This paper is structured as follows. In Section \ref{sec:ede}, we give a brief review of an early dark energy model. The data sets and constrained results are given in Section \ref{sec:results}. Section \ref{sec:conclusion} is the conclusion. | \label{sec:conclusion} In this brief paper, we consider the possible relaxation of the tension between the observed values of the tensor-to-scalar ratio $r$ from the Planck and BICEP2 by a special kind of early dark energy model which is characterized by its equation of state. The interesting property is that this early dark energy behaves like radiation early epoch and like dark energy at late time. And a phase transition between them happens around the reionization epoch. Therefore, at early time an effective neutrino species characterized by the ratio of early dark energy component are contributed to the energy budge. And this extra radiation component is helpful to relive the tension between the observations from Planck and BICEP2 at $2\sigma$ region as shown is this paper. However, the tension cannot be removed completely due to the small values of early dark energy contained by other cosmic observations. As a byproduct, the early dark energy model can remove the tension between observed values of Hubble parameter from Planck and the direct measurement of the Hubble constant. In this paper, we did not study the details of the sequences of phase transition of the early dark energy and the physics behind them. It deserves to be studied in the future. We expect this work can shed lights on the discovery of the nature of inflation and dark energy. | 14 | 4 | 1404.3804 |
1404 | 1404.1801_arXiv.txt | { We report a detection of the baryon acoustic oscillation (BAO) feature in the flux-correlation function of the \Lya forest of high-redshift quasars with a statistical significance of five standard deviations. The study uses 137,562 quasars in the redshift range $2.1\le z \le 3.5$ from the Data Release 11 (DR11) of the Baryon Oscillation Spectroscopic Survey (BOSS) of SDSS-III. This sample contains three times the number of quasars used in previous studies. The measured position of the BAO peak determines the angular distance, $\dA(z=2.34)$ and expansion rate, $H(z=2.34)$, both on a scale set by the sound horizon at the drag epoch, $r_d$. We find $\dA/r_d=11.28\pm0.65(1\sigma)^{+2.8}_{-1.2}(2\sigma)$ and $\dH/r_d=9.18\pm0.28(1\sigma)\pm0.6(2\sigma)$ where $\dH=c/H$. The optimal combination, $\sim\dH^{0.7}\dA^{0.3}/r_d$ is determined with a precision of $\sim2\%$. For the value $r_d=147.4~{\rm Mpc}$, consistent with the cosmic microwave background power spectrum measured by Planck, we find $\dA(z=2.34)=1662\pm96(1\sigma)~{\rm Mpc}$ and $H(z=2.34)=222\pm7(1\sigma)~{\rm km\,s^{-1}Mpc^{-1}}$. Tests with mock catalogs and variations of our analysis procedure have revealed no systematic uncertainties comparable to our statistical errors. Our results agree with the previously reported BAO measurement at the same redshift using the quasar-\Lya forest cross-correlation. The autocorrelation and cross-correlation approaches are complementary because of the quite different impact of redshift-space distortion on the two measurements. The combined constraints from the two correlation functions imply values of $\dA/r_d$ that are 7\% lower and 7\% higher for $\dH/r_d$ than the predictions of a flat \lcdm cosmological model with the best-fit Planck parameters. With our estimated statistical errors, the significance of this discrepancy is $\approx 2.5\sigma$. } | \label{introsec} Observation of the peak in the matter correlation function due to baryon acoustic oscillations (BAO) in the pre-recombination epoch is now an established tool to constrain cosmological models. The BAO peak at a redshift $z$ appears at an angular separation $\Delta\theta=r_d/[(1+z)\dA(z)]$ and at a redshift separation $\Delta z=r_d/\dH(z)$, where $\dA$ and $\dH=c/H$ are the angular and Hubble distances, and $r_d$ is the sound horizon at the drag epoch\footnote{ We follow the convention of \citet{anderson13}, $r_d=r_s(z_d)$, where $r_s$ is the sound horizon and $z_d$ is the drag redshift (baryon decoupling from photons), to be distinguished from $z_*$ (the redshift corresponding to unity optical depth for CMB photons). Earlier publications on BAO generally denoted $r_d$ simply as $r_s$. For models with cold dark matter, baryons and three light neutrino species, $r_d$ can be evaluated with Eq. (55) of \citet{anderson13}, which agrees with the CAMB-derived value to better than 0.1 per cent. }. Measurement of the peak position at any redshift thus constrains the combinations of cosmological parameters that determine $\dH/r_d$ and $\dA/r_d$. The BAO peak has been observed primarily in the galaxy-galaxy correlation function obtained in redshift surveys. The small statistical significance of the first studies gave only constraints on $\dv/r_d$ where $\dv$ is the combination $\dv=[(1+z)\dA]^{2/3}[z\dH]^{1/3}$, which determines the peak position for the galaxy correlation function when averaged over directions with respect to the line of sight. The first measurements were at $z\sim 0.3$ by the SDSS \citep{sdss1bao} and 2dFGRS \citep{2dfbao} with results from the combined data set presented by \citet{percival10}. A refined analysis using reconstruction \citep{eisenstein07,padmanabhan09} to improve the precision $\dv/r_d$ was presented by \citet{padmanabhan12} and \citet{mehta12}. Other measurements of $\dv/r_d$ were made at $z\sim0.1$ by the 6dFGRS \citep{beutler11}, at $(0.4<z<0.8)$ by WiggleZ \citep{blake11}, and, using galaxy clusters, at $z\sim0.3$ by \citet{veropalumbo13}. The Baryon Oscillation Spectroscopic Survey (BOSS; \citealt{bossoverview}) of SDSS-III \citep{eisenstein11} has presented measurements of $\dv/r_d$ at $z\sim 0.57$ and $z\sim0.32$ \citep{anderson12}. A measurement at $z\sim0.54$ of of $\dA/r_d$ using BOSS photometric data was made by \citet{seo12}. The first combined constraints on $\dH/r_d$ and $\dA/r_d$ were obtained using the $z\sim0.3$ SDSS data by \citet{chuang12} and \citet{xu12}. Recently, BOSS has provided precise constraints on $\dH/r_d$ and $\dA/r_d$ at $z=0.57$ \citep{anderson13,kazin13}. At higher redshifts, the BAO feature can be observed using absorption in the \Lya forest to trace mass, as suggested by \citet{macdonald03}, \citet{white03} and \citet{mcdoeisen07}. After the observation of the predicted large-scale correlations in early BOSS data by \citet{slosar11}, a BAO peak in the \Lya forest correlation function was measured by BOSS in the SDSS data release DR9 \citep{busca13,slosar13,kirkby13}. The peak in the quasar-\Lya forest cross-correlation function was detected in the larger data sets of DR11 \citep{font13}. The DR10 data are now public \citep{dr10ref}, and the DR11 data will be made public simultaneously with the final SDSS-III data release (DR12) in late 2014. This paper presents a new measurement of the \Lya forest autocorrelation function and uses it to study BAO at $z=2.34$. It is based on the methods used by \citet{busca13}, but introduces several improvements in the analysis. First, and most important, is a tripling of the number of quasars by using the DR11 catalog of 158,401 quasars in the redshift range 2.1~$\leq$~$z_q$~$\leq$~3.5. Second, to further increase the statistical power we used a slightly expanded forest range as well as quasars that have damped \Lya troughs in the forest. Finally, the \citet{busca13} analysis was based on a decomposition of the correlation function into monopole and quadrupole components. Here, we fit the full correlation $\xi(\rperp,\rpar)$ as a function of separations perpendicular, $\rperp$, and parallel, $\rpar$, to the line of sight. This more complete treatment is made possible by a more careful determination of the covariance matrix than was used by \citet{busca13}. Our analysis uses a fiducial cosmological model in two places. First, flux pixel pairs separated in angle and wavelength are assigned a co-moving separation (in $\hMpc$) using the $\dA(z)$ and $\dH(z)$ calculated with the adopted parameters. Second, to determine the observed peak position, we compare our measured correlation function with a correlation function generated using CAMB \citep{cambref} as described in \citet{kirkby13}. We adopt the same (flat) \lcdm model as used in \citet{busca13}, \citet{slosar13}, and \citet{font13}; with the parameters given in Table~\ref{modeltable}. The fiducial model has values of $\dA/r_d$ and $\dH/r_d$ at $z=2.34$ that differ by about 1\% from the values given by the models favored by CMB data \citep{planck13,calabrese13} given in the second and third columns of Table~\ref{modeltable}. This paper is organized as follows: Section \ref{samplesec} describes the DR11 data used in this analysis. Section \ref{mocksec} gives a brief description of the mock spectra used to test the analysis procedure; a more detailed description is given in \citet{bautista14}. Section \ref{xisec} presents our method of estimating the correlation function $\xi(\rperp,\rpar)$ and its associated covariance matrix. In Sect. \ref{fitssec} we fit the data to derive the BAO peak position parameters, $\dA(z=2.34)/r_d$ and $\dH(z=2.34)/r_d$. Section \ref{systsec} investigates possible systematic errors in the measurement. In Sect. \ref{cosmosec} we compare our measured peak position with that measured by the Quasar-\Lya-forest cross-correlation \citep{font13} and study \lcdm models that are consistent with these results. Section \ref{conclusionssec} concludes the paper. \begin{table} \begin{center} \caption{Parameters of the fiducial flat \lcdm cosmological model used for this analysis, the flat \lcdm model derived from Planck and low-$\ell$ WMAP polarization data, `Planck + WP'' \citep{planck13}, and a flat \lcdm model derived from the WMAP, ACT, and SPT data \citep{calabrese13}. The models are defined by the cold dark matter, baryon, and massive neutrinos densities, the Hubble constant, and the number of light neutrino species. The sound horizon at the drag epoch, $r_d$ is calculated using CAMB (which can be approximated with Eq. (55) of \citet{anderson13} to a precision of 0.1\%). } \begin{tabular}{l c c c} & fiducial & Planck &WMAP9 \\ & & + WP & +ACT+SPT \\ \hline \hline \noalign{\smallskip} $\om h^2$ & 0.1323 & 0.14305 & 0.1347 \\ $=\oc h^2$ & 0.1090 & 0.12038 & 0.1122 \\ $\;+\ob h^2$ & 0.0227 & 0.022032 & 0.02252 \\ $\;+\on h^2$ & 0.0006 & 0.0006 & 0 \\ $h $ & 0.7 & 0.6704 & 0.714 \\ $N_\nu$ & 3 & 3 & 3 \\ \hline \noalign{\smallskip} $\om$ &0.27 & 0.3183 & 0.265 \\ $r_d$ (Mpc) & 149.7 & 147.4 & 149.1 \\ & (104.80$~h^{-1}$) & (98.79$~h^{-1}$) & (106.4$~h^{-1}$) \\ $\dA(2.34)/r_d$ & 11.59 & 11.76 & 11.47 \\ $\dH(2.34)/r_d$ & 8.708 & 8.570 & 8.648 \\ \end{tabular} \end{center} \label{modeltable} \end{table}% | \label{conclusionssec} The \Lya correlation data presented in this study constrain $\dH/r_d$ and $\dA/r_d$ at $z\sim2.34$.\footnote{ The baofit software used in this paper is publicly available at http://github.com/deepzot/baofit/. The measured cross-correlation function and its covariance matrix, and the instructions to reproduce the main BAO results presented in this paper, can be downloaded from http://darkmatter.ps.uci.edu/baofit/, together with the likelihood surfaces used to generate the contours in Fig. \ref{crossautofig}. } The 3.0\% precision on $\dH/r_d$ and 5.8\% precision on $\dA/r_d$ obtained here improve on the precision of previous measurements: 8\% on $\dH/r_d$ \citep{busca13}, and 3.4\% on $\dH/r_d$ and 7.2\% on $\dA/r_d$ \citep{slosar13}. The increasing precision of the three studies is primarily due to their increasing statistical power, rather than to methodological improvements. The 2\% precision on the optimal combination $\dH^{0.7}\dA^{0.3}/r_d$ can be compared with the 1\% precision for $D_V(z=0.57)/r_d$ obtained by \citet{anderson13}. The derived values of $\dH/r_d$ and $\dA/r_d$ obtained here with the \Lya\ autocorrelation are similar to those inferred from the Quasar-Ly$\alpha$-forest cross-correlation \citep{font13}, as shown in Fig. \ref{crossautofig}. At the two-standard-deviation level, the two techniques are separately compatible with the Planck+WP and fiducial models of Table~\ref{modeltable}. However, the combined constraints are inconsistent with the Planck+WP \lcdm model at $\approx 2.5\sigma$ significance, given our estimated statistical uncertainties. The tests presented in earlier sections suggest that our statistical error estimates are accurate and that systematic uncertainties associated with our modeling and analysis procedures are smaller than these statistical errors. We are in the process of addressing what we consider to be the main weaknesses of our analysis. The artifacts in the spectrophotometric calibration due, for example, to Balmer lines, will be eliminated. More sophisticated continuum modeling making use of spectral features will allow us to verify that unsuspected correlated continua in neighboring quasars are not introducing artifacts in the autocorrelation function. Finally, we are producing realistic mock catalogs with quasar positions correlated with \lya~absorption features in the corresponding forests. Such mocks would allow us to verify the statistical errors for the cross-correlation measurement and to search for unsuspected correlations between the cross- and autocorrelation function measurements. All of these improvements in the analysis procedure will be used for publications using the higher statistical power of the upcoming DR12 data release. The cosmological implications of our results will be investigated in much greater depth in a forthcoming paper (The BOSS collaboration, in preparation), where we combine the Ly$\alpha$-forest BAO with the BOSS galaxy BAO results at lower redshift and with CMB and supernova data, which enables interesting constraints on a variety of theoretical models. | 14 | 4 | 1404.1801 |
1404 | 1404.1927.txt | We propose different scenarios where a keV dark matter annihilates to produce a monochromatic signal. The process is generated through the exchange of a light scalar of mass of order 300 keV - 50 MeV coupling to photon through loops or higher dimensional operators. For natural values of the couplings and scales, the model can generate a gamma-ray line which can fit with the recently identified 3.5 keV X-ray line. %{\color{red}Surprisingly, %the cross section needed to fit the observations, $\sigma v \simeq 10^{-33}~\mrm{cm^3 s^{-1}}$ is of the %order %predicted if the dark matter abundance is generated through freeze-in mechanism.} | \label{sec:introduction} \setcounter{equation}{0} %%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 14 | 4 | 1404.1927 |
|
1404 | 1404.2613_arXiv.txt | We investigate the star formation-feedback cycle in cosmological galaxy formation simulations, focusing on progenitors of Milky Way (MW)-sized galaxies. We find that in order to reproduce key properties of the MW progenitors, such as semi-empirically derived star formation histories and the shape of rotation curves, our implementation of star formation and stellar feedback requires 1) a combination of local early momentum feedback via radiation pressure and stellar winds and subsequent efficient supernovae feedback, and 2) efficacy of feedback that results in {\it self-regulation} of the global star formation rate on kiloparsec scales. We show that such feedback-driven self-regulation is achieved globally for a {\it local} star formation efficiency per free fall time of $\epsilon_{\rm ff}\approx 10\%$. Although this value is larger that the $\epsilon_{\rm ff}\sim 1\%$ value usually inferred from the Kennicutt-Schmidt (KS) relation, we show that it is consistent with direct observational estimates of $\epsilon_{\rm ff}$ in molecular clouds. Moreover, we show that simulations with local efficiency of $\epsilon_{\rm ff}\approx 10\%$ reproduce the global observed KS relation. Such simulations also reproduce the cosmic star formation history of the Milky Way sized galaxies and satisfy a number of other observational constraints. Conversely, we find that simulations that a priori assume an inefficient mode of star formation, instead of achieving it via stellar feedback regulation, fail to produce sufficiently vigorous outflows and do not reproduce observations. This illustrates the importance of understanding the complex interplay between star formation and feedback and the detailed processes that contribute to the feedback-regulated formation of galaxies. | \label{sect:introduction} The basic scenario of hierarchical galaxy formation \citep[][]{WhiteRees78,FallEfstathiou80} has been greatly elaborated and put on a firm footing within the Cold Dark Matter paradigm during the last three decades. Although the $\Lambda$ Cold Dark Matter ($\Lambda$CDM) model has proven broadly successful in explaining and predicting a variety of observations, such as the Cosmic Microwave Background temperature anisotropies \citep[e.g][]{komatsu_etal11,hinshaw_etal13,Planckparam2013}, the evolution of cluster abundance \citep{vikhlinin_etal09b}, and large scale distribution of matter in the Universe \citep{conroy_etal06,Springel2006}, many aspects of the theory of galaxy formation are not yet fully understood \citep[see, e.g.,][for a recent review]{Silk2012}. One of the most pressing problems in galaxy formation modelling is understanding why galaxies forming at the centers of dark matter halos are so inefficient in converting their baryons into stars. A number of different methods, such as dark matter halo abundance matching \citep{conroy_wechsler09,Guo2010}, satellite kinematics \citep{klypin_prada09,More2011}, and weak lensing \citep{Mandelbaum2006} \citep[see][for a comprehensive discussion]{kravtsov_etal14} point towards \emph{peak} stellar to dark matter mass fractions of $M_\star/M_{\rm h}\approx 3-5\,\%$ on average for $L_\star$ galaxies \citep[e.g.,][]{kravtsov_etal14}, far below the cosmological baryon fraction $\Omega_{\rm b}/\Omega_{\rm m}\approx 16\%$ \citep{Planckparam2013}. The low galactic baryon fractions are believed to be due to galactic winds driven by stellar feedback at the faint end of the stellar mass function \citep{DekelSilk86,Efstathiou00} and by the active galactic nuclei (AGN) and the bright end \citep{SilkRees1998,Benson2003}. Modeling these processes in fully cosmological hydrodynamical simulations has proven to be a daunting task due to the multi-scale nature of galaxy formation, where properties of the intergalactic distribution of baryons, on scales $\gtrsim 100\kpc$, are affected by star formation and feedback processes on scales of individual star clusters ($\lesssim 1\pc$). Although a formal spatial resolution of $\sim 10-100\pc$, comparable to the scale of massive giant molecular clouds (GMCs), is not uncommon in modern cosmological galaxy formation simulations \citep[e.g.][]{kravtsov03,GnedinKravtsov2010,Agertz09b,Hopkins2014}, the relevant star formation and feedback processes remain ``subgrid''. In particular, substantial differences in resulting galaxies may arise when different implementations and paramterizations of these processes are used in simulations, even when the same initial conditions are used \citep{Governato10,Aquila}. Implementations of stellar feedback in galaxy formation simulations have been explored in many studies over the last two decades \citep[e.g.][]{Katz92,NavarroWhite93, Katz1996, ThackerCouchman2001, Stinson06, Governato07,Scannapieco08,Colin2010,Agertz2011,AvilaReese2011,Guedes2011,Aquila,Hopkins2011,Brook2012,Stinson2013,Agertz2013,Ceverino2013,Roskar2013,Booth2013,Christensen2014}. Nevertheless, we still do not have a full understanding of what processes matter most for suppressing star formation and driving galactic winds over the vast range of observed galaxy masses. Recent studies \citep[][]{Leitner2012,Weinmann2012, Behroozi2013,Moster2013} have shown that not only is galaxy formation an inefficient process, but also that star formation in progenitors of most galaxies ($L\lesssim L_\star$) is significantly suppressed during the first 3 Gyr of cosmic evolution. \cite{vandokkum2013} recently reached a similar conclusion by matching cumulative co-moving number densities in the 3D-HST and CANDELS Treasury surveys, demonstrating that $\sim 90\%$ of the stellar mass in Milky Way mass galaxies formed after $z\sim2.5$. Much effort has gone into reproducing the $z=0$ $M_\star-M_{\rm h}$ relation over a large range of galaxy masses in simulations with efficient feedback \citep[e.g.,][]{Munshi2013}. At the same time, predicting its evolution, and hence reproducing the significant suppression of star formation necessary at $z\gtrsim2$ has proven more difficult. \cite{Brook2012} and \cite{Stinson2013} discussed the importance of ``early feedback''\footnote{Feedback that operates at times before the first SNe explosions, i.e. $t\lesssim 4\Myr$, for a coeval stellar population.} in their SPH simulations, here modeled by assuming that $10\%$ of the bolometric luminosity radiated by young stars get converted into thermal energy. This large energy injection resulted in star formation histories consistent with the data of \cite{Moster2013}. Similar results were found by \cite{Aumer2013} who considered a momentum based model of radiation pressure, although with the value of the infrared optical depth of $\tau_{\rm IR}\sim 25$, larger than in the models by \cite{Hopkins2011} and \cite{Agertz2013}. \cite{Hopkins2014}, \cite{Ceverino2013}, and \cite{TrujilloGomez2013} also found that radiative feedback, both due to photoionizaiton and radiation pressure, could play an important role in low mass galaxies (here progenitors of galaxies with $M_{\rm vir}(z=0)\lesssim 10^{12}\Msol$) at high redshifts, even for more moderate values of photon trapping by dust. While suppression of star formation in simulations of galaxy formation via strong stellar feedback has been widely explored in the recent literature, freedoms in the way in which star formation in the interstellar medium (ISM) is modeled has received less attention. Recent work by \citet[][see also \citealt{GnedinKravtsov2010,GnedinKravtsov11,Kuhlen2012,Christensen2014}]{Gnedin09} demonstrated how a star formation model based on the local abundance of $\HH$ could explain the observed steepening for $\Sigma_{\rm gas}<100\Msol\pc^{-2}$ in the Kennicutt-Schmidt (KS) relation for $z\approx 3-4$ Damped Lyman-$\alpha$ systems and Lyman Break Galaxies (LBGs). \cite{Governato10} found that a high threshold for star formation, in conjunction with higher resolution and strong feedback, can lead to more correlated feedback events and a more realistic halo baryon fraction. These results illustrate that it is paramount to explore how parameters of the star formation recipe and feedback implementation affect basic properties of galaxies forming in a given halo. In this paper we present results of a systematic study of such dependencies using high resolution, cosmological simulations of the Milky Way (MW) sized progenitors that include our new model for stellar feedback described in \cite{Agertz2013}. We specifically explore how the interplay between various modes of star formation and feedback models affect galactic characteristics at $z\gtrsim 1$. The paper is organized as follows: In \S\,\ref{sect:method} we outline our numerical method as well as star formation and feedback models. In \S \ref{sec:sfeff} we discuss empirical constraints on the efficiency of star formation in molecular clouds -- one of the most important parameters in our implementation of the star formation -- feedback cycle, and show that observations often indicate an efficiency in massive star forming clouds considerably larger than implied by the global normalization of the Kennicutt-Schmidt relation. \S\,\ref{sect:IC} describes the initial conditions and the simulation suite. In \S\ref{sect:results} we present our suite of cosmological simulations and demonstrate how two different models of star formation and feedback can match several observational properties of galaxies, including the star formation history, the total stellar mass expected from abundance matching, average gas metallicity and the rotational velocity. In \S\,\ref{sect:degeneracy} we discuss how the degeneracy between the two parameterizations can be broken, and show how only the simulation with efficient stellar feedback together with a high local efficiency of star formation can reproduce all the observed properties of galaxies. Finally, we discuss our results and conclusions in \S\,\ref{sect:discussion} and \ref{sect:conclusions}. | \label{sect:conclusions} In this paper we have presented a suite of high resolution cosmological zoom-in simulations of galaxy formation, focusing on the formation of a Milky Way-sized galaxy with a halo mass of $M_{\rm 200}\approx 10^{12}\Msol$ at $z=0$. We have focused on exploring how variations in the modeled star formation and feedback physics affect galaxy evolution and how properties of the simulated MW progenitors compare to modern high redshift ($z\gtrsim1$) estimates of global characteristics, such as star formation histories, the mass-metallicity relation, the Kennicutt-Schmidt relation and the stellar mass - halo mass relation. Our simulations adopt the feedback model presented recently by \cite{Agertz2013}, which accounts for energy and momentum injection via radiation pressure, stellar winds and supernovae type II and Ia. Furthermore, star formation is modeled using the local density of molecular gas \citep{kmt09,Gnedin09,Kuhlen2012}. Perhaps the central result of our study is that in our implementation feedback becomes efficient in suppressing star formation and driving outflows only if the {\it local} star formation efficiency per free fall time is sufficiently large, $\epsilon_{\rm ff}\approx 10\%$ for the density field resolved in the current simulations. Such large efficiency allows for a high degree of temporal and spatial correlation of energy and momentum injection. We show (in \S~\ref{sec:sfeff}) that such value of the local efficiency is consistent with observational estimates in giant molecular clouds. We confirm that in the models with efficient feedback, the star formation efficiency measured on global, kiloparsec scales self-regulates to the low value inferred from observations. The rest of our results can be summarized as follows. \begin{itemize} \item At the peak spatial resolution of our simulations, $\Delta x\sim75$ pc, simulated galaxy relations are sensitive not only to the details of stellar feedback processes and their parameters, but also to the underlying star formation model and the adopted efficiency of star formation. This highlights the fact that it is important to model carefully the entire star formation and feedback cycle. \item If the adopted $\epsilon_{\rm ff}$ is low (here $\sim 1\%$, relevant for the currently resolved gas density field), hence treating the observed inefficiency of galactic star formation as a model input rather than a prediction from the star formation--feedback cycle, the strength of feedback must be artificially boosted in order to regulate galaxy masses via galactic outflows. We show that although this can lead to a successful match to the semi-empirical stellar mass--halo mass relation, such simulations may be in tension with the normalization of the Kennicutt-Schmidt relation. Furthermore, in agreement with other recent studies \citep{Agertz2011,Roskar2013}, we also find that simply boosting feedback with a low $\epsilon_{\rm ff}$, to match global relations, prevents the formation of a well-defined gaseous disk, even at relatively low redshifts ($z\lesssim 1$). The morphology of the gaseous and stellar galactic disks may therefore serve as one of the key additional constraints on the parameters of the star formation--feedback loop. \item Our simulations indicate a complex interplay between the parameters of star formation and stellar feedback. If the star formation efficiency is sufficiently large to allow for feedback self-regulation, removing key feedback sources, such as radiation pressure or efficient thermal feedback, moves the galaxy off observed scaling relations, but in a complex manner. \item Encouragingly, we find that our fiducial model provides a good match to all considered observables at different redshifts: semi-empirically derived star formation histories, the stellar mass-gas metallicity relation and its evolution, the Kennicutt-Schmidt relation, the $M_\star-M_{\rm halo}$ relation and its evolution, as well as the flat shape of rotation curves and galaxy morphology. In particular, we show that our fiducial simulation, with feedback sufficient to drive vigorous galactic winds at high-$z$, is sufficiently gentle to allow for a young thin stellar disk to form by $z=0$. The disk has a flat rotation curve, with gas and stellar velocity dispersions consistent with observations of the Milky Way's at the solar circle. \end{itemize} Our results are encouraging, as they show that a comprehensive model that satisfies a number of non-trivial observational constraints and tests is feasible. In this work we have mostly discussed the $z\gtrsim 1$ results for our simulations, as the majority of them were stopped at high redshifts due to the computational expense. A significant fraction of stars in the $z=0$ thin disk is expected to form after $z\sim 1$ \citep{vandokkum2013}, which also appears to be the case in our fiducial model where we see the formation of well-defined thin stellar disk as soon as the turbulent gas rich disk enters an epoch of ``disk settling'' \citep{Kassin2012} at $z\lesssim 1$ (see Figure \ref{fig:map2}, \ref{fig:z0disk} and related text). Building upon the exploratory study presented here, we will in future work (Agertz \& Kravtsov, in prep) study and contrast galaxy sizes and morphologies at $z=0$ for a subset of the simulated galaxies. | 14 | 4 | 1404.2613 |
1404 | 1404.7200_arXiv.txt | We present the first extragalactic detections of several CH rotational transitions in the far-infrared (FIR) in four nearby galaxies: NGC 1068, Arp 220, M 82 and NGC 253 using the \textit{Herschel Space Observatory}. The CH lines in all four galaxies are a factor of 2 - 4 brighter than the adjacent HCN and \hcop\ J = 6-5 lines (also detected in the same spectra). In the star formation dominated galaxies, M 82, NGC 253 and Arp 220, the CH/CO abundance ratio is low ($\sim 10^{-5}$), implying that the CH is primarily arising in diffuse and translucent gas where the chemistry is driven by UV radiation as found in the Milky Way ISM. In NGC 1068, which has a luminous AGN, the CH/CO ratio is an order of magnitude higher suggesting that CH formation is driven by an X-ray dominated region. Our XDR models show that both the CH and CO abundances in NGC 1068 can be explained by an XDR-driven chemistry for gas densities and molecular hydrogen column densities that are well constrained by the CO observations. We conclude that the CH/CO ratio may a good indicator of the presence of AGN in galaxies. We also discuss the feasibility of detecting CH in intermediate- to high-$z$ galaxies with ALMA. | The methylidyne radical CH has been studied extensively at visible wavelengths through its electronic transitions in diffuse Galactic gas \citep{federman97, sheffer08}. From these observations, CH was found to be a powerful tracer of the molecular hydrogen in diffuse and translucent gas. Because CH is a light molecule its ground state rotational transitions lie at submm/FIR wavelengths and are impossible to observe from the ground. \textit{Herschel} made the first observations of the rotational transitions of CH arising in the far-infrared/sub-millimeter regime in Galactic star forming regions \citep{gerin10b, qin10, bruderer10, naylor10b}. In this work we present the first extragalactic detections of CH in four prototypical galaxies dominated by starbursts or AGN: NGC 1068, Arp 220, M 82 and NGC 253. CH can be present in high- or low-density gas depending on the formation scenario. If the abundance of ionized carbon is substantial, CH formation is believed to be initiated by the radiative association of C$^+$ with vibrationally excited molecular hydrogen, \htwo, in the outer layers of photon-dominated regions (PDRs), where the chemistry is dominated by UV radiation. The chemical network forming CH \citep[described in][]{black73} involves the following reactions: \begin{eqnarray}\label{chreact} C^{+} + H_2 \rightarrow CH_{2}^{+} + h\nu \\ CH^{+}_{2} + H_2 \rightarrow CH_{3}^{+} + H \nonumber \\ CH^{+}_{2} + e^{-} \rightarrow CH + H \nonumber \\ CH^{+}_{3} + e^{-} \rightarrow CH + H_2 \nonumber \end{eqnarray} CH can be formed in high-density gas via reactions described in Eq.\,[\ref{chreact}] or it can also be produced during \chp\ synthesis in lower density material ($\sim 50$ \cmthree) from MHD shocks \citep{draine86, pineau86}. In Galactic star forming regions, CH is about a factor of 1 -- 3 more abundant than \chp\ \citep{godard12}. The most efficient reaction forming \chp\ is $C^{+} + H_2 \rightarrow CH^{+} + H$, which has a high endothermic barrier of 4640 K. This reaction can form \chp\ in a dense and highly illuminated PDR but is inefficient in the cold diffuse ISM. The recent investigations by \citet[][also see \citet{falgarone10}]{godard12, godard09} suggest that \chp\ can be produced in the diffuse ISM via kinetic energy from turbulent dissipation. The \chp\ molecule is also rapidly destroyed at high densities. By comparing CH/\chp\ ratios with other galaxies we can determine whether these molecules are tracing similar environments and densities as the Milky Way (MW) or if their production is influenced by strong starbursts and AGN. The CH energy level diagram shown in Figure \ref{chel} is taken from \citet{stacey87}. The rotational lines of CH in the submm/FIR have a characteristic doublet pattern due to lambda doubling: the spin-orbit interaction of the unpaired $\pi$ electron splits the rotational levels (N) into two ladders depending on the relative orientation between the electron's spin and orbital angular momentum vectors. The rotational levels, J, in the individual ladders split into $\Lambda$-doublet states (denoted by + or $-$) from the relative orientations of the electron's orbital momentum axis and the molecular rotation axis. The magnetic hyperfine interaction further splits the $\Lambda$-doublet states. The major rotational transitions at 560 \mic, 203 \mic, 180 \mic, and 149 \mic\ (highlighted by the red arrows in Figure \ref{chel}) are accessible with \textit{Herschel}. The 560 \mic\ transition is the easiest to detect by both the SPIRE Fourier Transform Spectrometer (FTS) and Heterodyne Instrument for the Far Infrared (HIFI) because of their high sensitivity at this wavelength. This transition has six hyperfine components grouped near 532.8 GHz ($1+ \rightarrow 1-, 2+ \rightarrow 1-, \mathrm{and} \, 1+ \rightarrow 0-$) and 536.8 GHz ($2- \rightarrow 1+, 1- \rightarrow 1+, \mathrm{and} \, 1- \rightarrow 0+$), very close in frequency to the HCN and \hcop\ J = 6-5 lines and therefore can only be resolved by the higher resolution of HIFI. In this work we present sub-millimeter observations of these CH lines and additionally \chp\ lines from \textit{Herschel} in four prototypical galaxies: Arp 220 (Starburst/AGN), NGC1068 (AGN Seyfert-2), M82 (Starburst/ no AGN) and NGC 253 (starburst/ no AGN). These observations are presented in Section 2. The properties of molecular gas derived from CH are compared with CO observations to investigate where CH is arising in these galaxies and if its formation in these starburst/AGN dominated galaxies differs from the Milky Way (sections 3 and 4). The feasibility of detecting CH in intermediate- to high-redshift galaxies with ALMA is discussed in section 5. | 14 | 4 | 1404.7200 |
|
1404 | 1404.3938_arXiv.txt | The structure of the photospheric magnetic field during solar flares is examined using {\it echelle} spectropolarimetric observations. The study is based on several Fe~{\sc I} and Cr~I lines observed at locations corresponding to brightest H$\alpha$ emission during thermal phase of flares. The analysis is performed by comparing magnetic field values deduced from lines with different magnetic sensitivities, as well as by examining the fine structure of $I\pm V$ Stokes profiles splitting. It is shown that the field has at least two components, with stronger unresolved flux tubes embedded in weaker ambient field. Based on a two-component magnetic field model, we compare observed and synthetic line profiles and show that the field strength in small-scale flux tubes is about $2-3$~kG. Furthermore, we find that the small-scale flux tubes are associated with flare emission, which may have implications for flare phenomenology. | \label{intro} There is observational evidence that the photospheric magnetic field is very inhomogeneous at small scales ({\it i.e.} $\approx 100$ km) (see, {\it e.g.} \opencite{sola93} for a review). Early magnetographic observations showed that the magnetic field strengths evaluated using spectral lines with similar characteristics but different magnetic sensitivity ({\it i.e.} different Lande factor [$g$]) can vary by up to a factor of $2.5$ \cite{host72,sten73}. This effect was interpreted as an indicator of unresolved multi-component structure with intense magnetic flux tubes embedded in a non-magnetic atmosphere or an atmosphere with weaker ambient field (Figure~\ref{f-sketch}). It is believed that these small-scale magnetic flux tubes may account for nearly $90\,\%$ of the photospheric magnetic flux outside sunspots \cite{frst72}. \begin{figure} % \centerline{\includegraphics[width=1.0\textwidth,clip=]{f01.pdf}} \caption{Left panel: two-component field with stronger unresolved magnetic flux tubes with field $B_\mathrm{sub}$ embedded into weaker ambient field $B_\mathrm{back}$, as shown in the left panel. Apart from different magnetic field strength, the two photospheric components may have different intensities and filling factors [$\mathcal{F}$], different widths (represented by their effective temperatures) and different line-of-sight (LOS) velocities. Right panels: $I+V$ components corresponding to the background (blue dot-dashed line) and unresolved flux tubes (red dotted line) (C1 and C2, respectively), resulting $I+V$ and $I-V$ profiles (solid and dashed black lines, respectively). Lower panel shows corresponding Stokes-$V$ profile.} \label{f-sketch} \end{figure} Most existing instruments in the optical range can directly resolve features with sizes of $\approx 1$ Mm. Hence, here by ``unresolved'' we mean spatial scales $\lesssim 1$~Mm. The most direct measurements of unresolved field structure were carried out using speckle interferometry in the line Fe~{\sc I}~5250.2~$\ang$ \cite{keva92,kell92} and revealed magnetic elements with a field strength of a few kG and sizes of $100\,-\,200$~km, which, perhaps, can be considered as an upper limit on the diameters of small-scale magnetic flux tubes. Also, \citext{lin95} observed full Stokes profiles in magneto-sensitive infrared Fe~{\sc I} lines 15648~$\ang$ and 15652~$\ang$ and found two types of small-scale magnetic elements: stronger elements with the field of $1.4$~kG and sizes $\approx 10^2\,-\,10^3$~km and weaker ones with the field strength of $\approx 500$~G and sizes about $70$~km. The two types of magnetic elements were attributed to network and inter-network flux tubes. It was also concluded that the inter-network magnetic field elements have rather short lifetimes of about few hours. There are a number of indirect measurements of unresolved magnetic field structure characteristics. Estimations of horizontal sizes of intense magnetic flux tubes vary significantly from tens of kilometers \cite{wieh78,lots89} to hundreds of kilometers \cite{sanc98}. Comparison of the effective field values obtained using spectral lines with different magnetic sensitivity shows that the magnetic field in such flux tubes is about $1.0\,-\,3.0$~kG, although there are some indications that it could be substantially higher \cite{race05,lozi09}. The main reason behind such large discrepancies in estimations is that even the two-component model, which is used to fit the observational data, has about ten free parameters, such as magnetic field strengths, field inclinations, and surface brightness for both components, along with the filling factor and other parameters. Hence, diagnostics of the small-scale magnetic field require some realistic assumptions about the field structure to reduce the number of free parameters. For instance, it might be safe to assume that the intense, unresolved flux tubes within a sampled area are almost identical \cite{ulre09} and the key free parameters are the strength and vertical gradient of magnetic field in these flux tubes, and the filling factor. The fine structure of the magnetic field in solar flares is understood even less than that in the quiet photosphere. There is strong evidence of unresolved magnetic field in flares, but their diagnostics is a challenging task, because of the associated temperature and velocity inhomogeneities at small scales. Furthermore, it is not always possible to distinguish between horizontal and vertical inhomogeneities at sub-telescopic scales. \citext{lolo94} observed full Stokes profiles of several Fe~{\sc I} lines in order to study the structure of magnetic field in the 2B solar flare of 16~June~1989. It was found that the small-scale field strength was between $1.0$ and $1.5$~kG, and it substantially changed during the flare. It was also found that the filling factor decreased with time. More recently, the new generation of solar space observatories along with advanced ground-based instruments have provided more evidence for fine structure of the magnetic field in solar flares. Thus, full-Stokes-imaging spectropolarimetry of a C-class flare with the Interferometric Bidimensional Spectropolarimeter (IBIS) shows that Stokes profiles are highly irregular, indicating the presence of unresolved multi-component magnetic and velocity fields \cite{klei12}. The resolution of the instrument (up to 0.33~arcsec, or 240~km) provides the upper limit for the sizes of these unresolved magnetic elements. \citext{fise12} used the data from the Synoptic Optical Long-term Investigations of the Sun (SOLIS) vector-spectrograph in order to investigate the evolution of magnetic field in an X-class flare and also found that the Stokes profiles demonstrate complex, highly asymmetric structure that may be explained by a multi-component velocity field or by substantial perturbations of the spectral line profile due to heating. In addition, they show that a small patch of the photosphere, co-spatial with hard X-ray footpoints observed by Ramaty High Energy Solar Spectroscopic Imager (RHESSI), exhibits an unresolved fine structure. A quantitative estimation of a typical cross-section of small-scale magnetic elements can be made based on the analysis by \citext{anro12}. They carried out observations of the coronal rain using the Crisp Imaging Spectro-polarimeter (CRISP) at the Swedish Solar Telescope and concluded that the coronal rain consists of elements with a typical width of $\approx$310~km. The structure and temporal variations of the small-scale magnetic field during flares may be related to the fast evolution of magnetic field in the corona and, therefore, the small-scale field structure in active regions and especially during solar flares, deserves more attention. In the present work, we aim to study unresolved structure of magnetic field at the photosphere during solar flares using two different approaches. The first approach is based on the analysis of the relationship between magnetic field strengths measured using different spectral lines and their Lande factors [$g$], which is similar to the method applied in magnetographic observations. The second approach is based on the analysis of fine structure of $I\pm V$ Stokes-profile splitting, which is possible only in observations with relatively high spectral resolution. | \label{concl} Our observational data demonstrates that, similarly to the quiet photopshere \cite{sten73,sola93}, the magnetic field at the photospheric level in flares is very inhomogeneous at unresolved scales. The observational picture is more complicated than in the quiet photosphere and in plages due to the presence of emission and fast plasma flows along the line-of-sight. The comparison of the magnetic field values deduced from spectral lines with different Lande factor $g$ shows that the effective field strength $B_\mathrm{eff}$ increases with $g$ (see Section ~\ref{observ-lr}), in contrast with what is normally observed in the quiet photosphere. Analysis of the synthetic $I\pm V$ Stokes profiles for two-component fields shows that this is possible in two cases: when the strong magnetic field has its polarity opposite to the polarity of the weaker ambient field or when the spectral components corresponding to the unresolved field (C2) show emission. The presence of very strong field of opposite polarity should be associated with high current densities. Hence, the second possibility, with emission from intense magnetic flux tubes, looks more realistic. Furthermore, the second option seems to be more viable as emission peaks are often observed in metallic lines in moderate and bright flares. The fine structure of $I\pm V$ profiles observed in flares has been studied using bisector splitting functions ($\Delta \lambda_\mathrm{H}(\Delta \lambda_\mathrm{c})$) (Section ~\ref{observ-bs}). It is known that the centre of a Fraunhofer line is formed predominantly in cooler regions, while the wings correspond to higher temperatures and are formed slightly deeper. Hence, the average trend of $\Delta \lambda_\mathrm{H}(\Delta \lambda_\mathrm{c})$ can be considered as the result of corellation between the magnetic field and the temperature. Hence, it is very unlikely, that the decrease of $\Delta \lambda_\mathrm{H}$ with $\Delta \lambda_\mathrm{c}$ is the result of vertical magnetic field gradients, as it would imply magnetic field increase with height. Horizontal inhomogeneity of the magnetic field seems to be rather realistic explanation for the observed trends of $\Delta \lambda_\mathrm{H}(\Delta \lambda_\mathrm{c})$, with higher magnetic field corresponding to lower temperatures and turbulent velocities. Analysis of the bisector splitting in synthetic $I\pm V$ Stokes profiles shows that the localized extrema, similar to those seen in Figure~\ref{f-bsall}, can be explained by the presence of narrow emission or absorption components of the spectral line with substantial shift $\Delta \lambda_\mathrm{H}$. Thus, based on synthetic profiles, it is possible to estimate the magnitude of the field, that would result in peaks on the bisector splitting profiles. The peak at $\Delta \lambda_\mathrm{c}\approx 75$~m$\ang$ in the case of the line similar to Fe~{\sc I}~$5233$~$\ang$ can be caused by field strength $\approx 2.5-3.0$~kG. However, these estimations are very sensitive to the width of the main component of the spectral line, and should be used with caution. In general, our preliminary results confirm previous conclusions that the photospheric magnetic field in flares has at least two components. Furthermore, our study shows that in flares thermodynamic conditions in the intense magnetic flux tubes are very different from conditions outside. The most interesting feature that has been revealed by this study, is an apparent link between the strong unresolved field and emission in flares. This effect can be easily seen in observed profiles: the Zeeman split of emission peaks in line cores is normally bigger than that of the absorption component of a spectral line. This finding could be quite important for the flare phenomenology. There are several possible explanations for emission observed in photospheric lines: atoms can be excited due to non-thermal particle precipitation, heating by propagating waves, or by conduction. In any case, it is very likely that the observed emission is directly related to energy release in the corona. Hence, the revealed connection between the emission in metallic lines and stronger field component may indicate that the intense unresolved photospheric magnetic elements are topologically connected to the coronal field, while weak ambient photospheric magnetic fluxes only form the low-level magnetic canopy (see, {\it e.g.}, \opencite{sole99}). This study gives a rough estimate of the backgound and strong photospheric magnetic field components. Obviously, more work needs to be done before these can be evaluated more reliably. Observationally, higher precision may be achieved by using combination of different methods, for instance, by using the data obtained from Zeeman and Hanle measurements. This, however, would not help to answer the question about the size of the small-scale magnetic elements. In order to address this problem, direct observations with higher spatial resolution are needed and these are likely to be possible with future missions. Current instruments provide reasonably high spatial resolution of $\approx 200-300$ km, but this is still not sufficient for reliable investigation of magnetic field fine structure. Thus, high-resolution magnetic field maps of flares 12 and 13 obtained with Helioseismic and Magnetic Imager (HMI) onboard Solar Dynamic Observatory (SDO) show that the field is still very inhomogeneous even on a scale of one pixel. Alternatively, observations in other spectral bands may provide an opportunity to resolve very small scales. For example, future solar observations with Atacama Large Millimeter/submillimeter Array (ALMA) would be able to provide polarimetric data in sub-THz range with mili-arcsec resolution, possibly giving a unique insight into solar magnetic field structure \cite{loue09}. Additionally, local helioseismology may provide an indirect estimations, as $p$-mode scattering and absorption can be sensitive to the magnetic flux tube sizes (see, {\it e.g.}, \opencite{choe98}; \opencite{goja08}; \opencite{jago08}; \opencite{fele12}). Finally, more theoretical work concerning the radiative transfer in strongly inhomogeneous magnetic field is needed in order to explain adequately the observed fine structure of Stokes profiles. \begin{acks} The authors thank Philippa Browning for useful comments from which the article strongly benefited. MG is supported by the Science and Technology Facilities Council (UK). \end{acks} | 14 | 4 | 1404.3938 |
1404 | 1404.1083_arXiv.txt | According to the photo-heating model of the intergalactic medium (IGM), He II reionization is expected to affect its thermal evolution. Evidence for additional energy injection into the IGM has been found at $3\lesssim z\lesssim4$, though the evidence for the subsequent fall-off below $z\sim2.8$ is weaker and depends on the slope of the temperature--density relation, $\gamma$. Here we present, for the first time, an extension of the IGM temperature measurements down to the atmospheric cut-off of the H I Lyman-$\alpha$ forest at $z\simeq1.5$. Applying the curvature method on a sample of 60 UVES spectra we investigated the thermal history of the IGM at $z<3$ with precision comparable to the higher redshift results. We find that the temperature of the cosmic gas traced by the Ly-$\alpha$ forest [$T(\bar{\Delta})]$ increases for increasing overdensity from $T(\bar{\Delta})\sim 22670$ K to 33740 K in the redshift range $z\sim2.8-1.6$. Under the assumption of two reasonable values for $\gamma$, the temperature at the mean density ($T_{0}$) shows a tendency to flatten at $z\lesssim 2.8$. In the case of $\gamma\sim1.5$, our results are consistent with previous ones which indicate a falling $T_{0}$ for redshifts $z\lesssim2.8$. Finally, our $T(\bar{\Delta})$ values show reasonable agreement with moderate blazar heating models. | Starting from a very hot plasma made of electrons and protons after the Big Bang, to the gas that now fills the space between galaxies, the Intergalactic medium (IGM) has been one of the main ``recorders'' of the different phases of evolution of the Universe. Changes in the thermodynamic state and chemical composition of this gas reflect the conditions for the formation and the evolution of the structures that we can observe today. In particular, the IGM thermal history can be an important source of information about reionizing processes that injected vast amounts of energy into this gas on relatively short timescales. For this reason considerable efforts have been made to find any ``footprint" of either H I ($z <6$) or He II ($z<3$) reionization. Because the ionisation potential of He II (from He II to He III) is 54.4 eV and fully ionized helium recombines more than 5 times faster than ionized hydrogen, this second reionization event should have begun later, after the reionization of hydrogen and He I ($11\gtrsim z \gtrsim 6$, \citealt{LL11}; \citealt{Fan06}; \citealt{Becker01}) when quasars started to dominate the UV background (\citealt{MiraldaescudeQ00}). Theoretically, the much harder photons from quasars would have been able to fully ionize He II around redshifts $3\lesssim z \lesssim 4.5$ but these estimates change depending on assumptions about the abundance of quasars (QSOs) and the hardness of their spectra (\citealt{Meiksin05}). While the direct observation, through the detection of the ``Gunn--Peterson effect'', recently suggests the end of the He II reionization at $z\sim 2.7$ (e.g. \citealt{Shull10};\citealt{Worseck11}; \citealt{Syphers13}; \citealt{Syphers14}; ), any current constraint on the physics of this phenomenon is limited by the cosmic variance among the small sample of ``clean'' lines of sight, those along which the He II Lyman-$\alpha$ transition is not blocked by higher-redshift HI Lyman limit absorption. For this reason indirect methods have been developed to obtain a detailed characterization of the He II reionization. It is predicted that the IGM should be reheated by photo-ionization heating during He II reionization and, because its cooling time is long, the low density gas retains some useful memory of when and how it was reionized. In fact, at the mean density of the IGM the characteristic signature of reionization is expected to be a peak: a gradual heating followed by cooling due to adiabatic expansion (e.g. \citealt{McQuinn09}). In the last decade, the search for this feature and the study of the thermal history of the IGM as a function of redshift have been the objectives of different efforts, not only to verify the theoretical prediction and constrain the timing of He II reionization, but also to obtain information on the nature of the ionizing sources and on the physics of the related ionizing mechanisms. To obtain measurements of the temperature of the IGM, studying the absorption features of the HI Lyman-$\alpha$ forest has proven to be a useful method so far. The widths of these lines are sensitive to thermal broadening but are also affected by Hubble broadening and peculiar velocities. Cosmological simulations are therefore required to characterize the large scale structure and bulk motion of the IGM (\citealt{Meiksin10}), before the gas temperature can be determined. Previous efforts to extract information on the thermal state of the cosmic gas from the Lyman-$\alpha$ forest can be divided into two main approches: the study of individual absorption features and the quantification of the absorption structures with a global statistical analysis of the entire forest. The first method consists of decomposing the spectra into a set of Voigt profiles. \citet{Schaye00} found evidence using this technique for an increase in the IGM temperature consistent with He II reionization at $z\simeq3$. In contrast, \citet{McDonald01} found a constant temperature over $z\sim2-4$. A characterization of the flux probability distribution (PDF) based on pixel statistics has also been used to analyse the forest and extract information from the comparison with theoretical models (\citealt{Bolton08}; \citealt{Calura12}). However, the PDF is sensitive to a range of systematic effects, including the placement of the unabsorbed continuum. A further approach is to use wavelet analysis to characterize the Ly$\alpha$ line-widths distribution in terms of discrete wavelets. \citet{Theuns02} and \citet{Lidz10} found evidence using this technique for He II reionization completing near $z\sim3.4$, but with large statistical uncertainties. In the recent work of \citet{Garzilli12}, the PDF and the wavelet decomposition methods have been compared and tested on Lyman-$\alpha$ spectra at low redshift. While the results are in formal agreement with previous measurements, the uncertainties are still large and there is a mild tension between the two analyses. Recently, \citet{Becker11} developed a statistical approach based on the flux curvature. This work constrained the temperature over $2\lesssim z\lesssim 4.8$ of a ``optimal" or ``characteristic" overdensity, which evolves with redshift. The error bars were considerably reduced compared to previous studies, partially at the expense of determining the temperature at a single density only, rather than attempting to constrain the temperature--density relation. Some evidence was found for a gradual reheating of the IGM over $3 \lesssim z \lesssim 4$ but with no clear evidence for a temperature peak. Given these uncertainties, the mark of the He II reionization still needs a clear confirmation. Nevertheless, the curvature method is promising because it is relatively robust to continuum placement errors: the curvature of the flux is sensitive to the shape of the absorption lines and not strongly dependent on the flux normalization. Furthermore, because it incorporates the temperature information from the entire Lyman-$\alpha$ forest, this statistic has the advantage of using more of the available information, as opposed to the line-fitting method which relies on selecting lines that are dominated by thermal broadening. An injection of substantial amounts of thermal energy is predicted to result in both an increase in the IGM temperature and a change in the temperature--density (T--$\rho$) relation. The detailed study of this process has to take into consideration the effects of the IGM inhomogeneities driven by the diffusion and percolation of the ionized bubbles around single sources, and currently constitutes an important object of investigation through hydrodynamical simulations (e.g. \citealt{Compostella13}). In the simplest scenario, for gas at overdensities $\Delta\lesssim 10$ ($\Delta=\rho / \bar{\rho}$, where $\bar{\rho}$ is the mean density of the IGM), the temperature is related to the density with a power-law relation of the form: \begin{equation} T(\Delta)=T_{0}\Delta^{\gamma -1} , \end{equation} where $T_{0}$ is the temperature at the mean density (\citealt{HuiGnedin1997}; \citealt{Valageas02}). The evolution of the parameters $T_{0}$ and $\gamma$ as a function of redshift then describes the thermal history of the IGM. A balance between photo-heating and cooling due to adiabatic expansion of the Universe will asymptotically produce a power law with $\gamma=1.6$ (\citealt{HuiGnedin1997}). During the reionization the slope is expected to flatten temporarily before evolving back to the asymptotic value. Possible evidence for this flattening at $z\simeq3$, seems to be consistent with He II reionization occurring around this time (e.g., \citealt{Ricotti00}; \citealt{Schaye00}). Some analyses of the flux PDF have indicated that the T--$\rho$ relation may even become inverted (e.g., \citealt{Becker07}; \citealt{Bolton08}; \citealt{Viel09}; \citealt{Calura12}; \citealt{Garzilli12}). However, the observational uncertainties in this measurement are considerable (see discussion in \citealt{Bolton13}). A possible explanation was suggested by considering radiative transfer effects (\citealt{Bolton08}). Although it appears difficult to produce this result considering only He II photo-heating by quasars (\citealt{McQuinn09};\citealt{Bolton09}), a new idea of volumetric heating from blazar TeV emission predicts an inverted temperature--density relation at low redshift and at low densities. According to these models, heating by blazar $\gamma$-ray emission would start to dominate at $z\simeq3$, obscuring the ``footprint" of He II reionization (\citealt{Chang12}; \citealt{Puchwein12}). Even if in the most recent analysis, with the line-fitting method (\citealt{Rudie13}; \citealt{Bolton13}), the inversion in the temperature--density relation has not been confirmed, a general lack of knowledge about the behaviour of the T--$\rho$ relation at low redshift ($z<3$) still emerges, accompanied with no clear evidence for the He II reionization peak. A further investigation of the temperature evolution in this redshifts regime therefore assumes some importance in order to find constraints for the physics of the He II reionization and the temperature--density relation of the IGM. The purpose of this work is to apply the curvature method to obtain new, robust temperature measurements at redshift $z<3$, extending the previous results, for the first time, down to the optical limit for the Lyman-$\alpha $ forest at $z\simeq1.5$. By pushing the measurement down to such a low redshift, we attempt to better constrain the thermal history in this regime, comparing the results with the theoretical predictions for the different heating processes. We infer temperature measurements by computing the curvature on a new set of quasar spectra at high resolution obtained from the archive of the UVES spectrograph on the VLT. Synthetic spectra, obtained from hydrodynamical simulations used in the analysis of \citet{Becker11} and extended down to the new redshift regime are used for the comparison with the observational data. Similar to \citet{Becker11}, we constrain the temperature of the IGM at a characteristic overdensity, $\bar{\Delta}$, traced by the Lyman-$\alpha$ forest, which evolves with redshift. We do not attempt to constrain the T--$\rho$ relation, but we use fiducial values of the parameter $\gamma $ in Eq. 1 to present results for the temperature at the mean density, $T_{0}$. This paper is organised as follows. In Section 2 we present the observational data sample obtained from the VLT archive, while the simulations used to interpret the measurements are introduced in Section 3. In Section 4 the curvature method and our analysis procedure are summarized. In Section 5 we present the data analysis and we discuss the strategies applied to reduce the systematic uncertainties. The calibration and the analysis of the simulations is described in Section 6. The results are presented in Section 7 for the temperature at the characteristic overdensities and the temperature at the mean density for different values of $\gamma$. We discuss the comparison with theoretical models in Section 8, and conclude in Section 9. | In this work we have utilized a sample of 60 VLT/UVES quasar spectra to make a new measurement of the IGM temperature evolution at low redshift, $1.5\la z\la2.8$, with the curvature method applied to the H{\sc \,i} Lyman-$\alpha$ forest. For the first time we have pushed the measurements to the lowest optically-accessible redshifts, $z\sim1.5$. Our new measurements of the temperature at the characteristic overdensities traced by the Ly-$\alpha$ forest, $T(\bar{\Delta})$, are consistent with the previous results of Becker et al.~(2011) in the overlapping redshift range, $2.0 < z < 2.6$, despite the datasets being completely independent. They show the same increasing trend for $T(\bar{\Delta})$ towards lower redshifts while, in the newly-probed redshift interval $1.5\la z\la2.0$, the evolution of $T(\bar{\Delta})$ is broadly consistent with the extrapolated trend at higher redshifts. The translation of the $T(\bar{\Delta})$ measurements into values of temperature at the mean density, $T_{0}$, depends on the slope of the temperature--density relation, $\gamma$, which we do not constrain in this work. However, for reasonable, roughly constant, assumptions of this parameter, we do observe some evidence for a change in the slope of the temperature evolution for redshifts $z\lesssim2.8$, with indications of at least a flattening, and possibly a reversal, of the increasing temperature towards lower redshifts seen in our results and those of \citet{Becker11} for $2.8\lesssim z\lesssim 4$. In particular, for the minimum $T_{0}$ case, with $\gamma\sim 1.5$, the extension towards lower redshifts provided by this work adds to existing evidence for a decrease in the IGM temperature from $z\sim2.8$ down to the lowest redshifts probed here, $z\sim1.5$. This could be interpreted as the footprint of the completion of the reheating process connected with the He II reionization. Following the additional hypothesis that our low redshift temperature measurements are already tracing the thermal asymptote, the cooling of $T_{0}$ inferred at $z\lesssim2.8$ (assuming $\gamma\sim 1.5$) may suggest that the UV background has changed, hardening during the He II reionization epoch. However, the expectation for the evolution of $T_{0}$ following HeII reionization will depend on the evolution in $\gamma$ and on details of the reionization model. We also compared our T($\bar{\Delta}$) measurements with the expectations for the models of \citet{Puchwein12} with and without blazar heating contributions. To allow a fair comparison with our observed values, the model predictions were computed at the corresponding (redshift-dependent) characteristic overdensities ($\bar{\Delta}$). Our observational results seem to be in reasonable agreement with a moderate blazar heating scenario. However, to definitely confirm or rule out any specific thermal history it is necessary to obtain new, model-independent measurements of the temperature at the mean density. With the IGM curvature now constrained from $z\sim4.8$ down to $z\sim1.5$, the main observational priority now is clearly to tightly constrain the slope of the temperature--density relation, $\gamma$, and its evolution over the redshift range $1.5\la z\la 4$. This is vital in order to fix the absolute values of the temperature at the mean density and to comprehensively rule out or confirm any particular heating scenarios. Finally, we note that, even though our new measurements have extended down to $z\sim1.5$, there is still a dearth of quasar spectra with high enough S/N in the 3000--3300\,\AA\ spectral range to provide curvature information in our lowest redshift bin, $1.5<z<1.7$. We have searched the archives of both the VLT/UVES and Keck/HIRES instruments for new spectra to contribute to this bin. However, the few additional spectra that we identified had relatively low S/N and, when included in our analysis, contributed negligibly to the final temperature constraints. Therefore, new observations of UV-bright quasars with emission redshifts $1.5\la z_{\rm em}\la1.9$ are required to improve the temperature constraint at $1.5<z<1.7$ to a similar precision as those we have presented at $z>1.7$. | 14 | 4 | 1404.1083 |
1404 | 1404.3565_arXiv.txt | We use large-scale lattice simulations to compute the rate of baryon number violating processes (the sphaleron rate), the Higgs field expectation value, and the critical temperature in the Standard Model across the electroweak phase transition temperature. While there is no true phase transition between the high-temperature symmetric phase and the low-temperature broken phase, the cross-over is sharply defined at $T_c = (159\pm 1)$\,GeV. The sphaleron rate in the symmetric phase ($T> T_c$) is $\Gamma/T^4 = (18\pm 3)\alpha_W^5$, and in the broken phase in the physically interesting temperature range $130\mbox{\,GeV} < T < T_c$ it can be parametrized as $\log(\Gamma/T^4) = (0.83\pm 0.01)T/{\rm GeV} - (147.7\pm 1.9)$. The freeze-out temperature in the early Universe, where the Hubble rate wins over the baryon number violation rate, is $T_* = (131.7\pm 2.3)$\,GeV. These values, beyond being intrinsic properties of the Standard Model, are relevant for e.g. low-scale leptogenesis scenarios. \end {abstract} | 14 | 4 | 1404.3565 |
||
1404 | 1404.6947_arXiv.txt | We present astrometric analysis of archival data of water masers in the star-forming region Sharpless~269 (S269) IRS~2w, observed with the VLBI Exploration of Radio Astrometry. An annual parallax of one of the bright maser features in this region was previously reported to be $0.189\pm0.008$~milliarcsecond (mas) using part of the same archival data as we used. However, we found that this maser feature is not the best to represent the annual parallax to S269~IRS~2w because the morphology is remarkably elongated in the east--west direction. For this study we have selected another maser feature showing simpler morphology. This makes the new annual parallax estimate more credible. Our newly obtained annual parallax is $0.247 \pm 0.034$~mas, corresponding to $4.05^{+0.65}_{-0.49}$~kpc. This value is well consistent with the 3.7--3.8~kpc obtained using the kinematic distance estimates and photometric distance modulus. We considered two hypotheses for the water maser spatial distribution, a bipolar outflow and an expanding ring, in a kinematic model fitting analysis with a radially expanding flow. At this stage, any conclusions about the systemic proper motion could not be drawn from the kinematic analysis. Alternatively, we evaluated the mean proper motion to be ($0.39 \pm 0.92$,~$-1.27 \pm 0.90$)~mas~yr$^{-1}$ eastward and northward, respectively, from the obtained proper motions of the detected water maser features. The newly obtained annual parallax and mean proper motion give the peculiar motion of S269~IRS~2w to be ($U_{\mathrm{s}}$,~$V_{\mathrm{s}}$,~$W_{\mathrm{s}}$) of ($8 \pm 6$, $-21 \pm 17$, $1 \pm 18$)~km~s$^{-1}$. | \label{sec:01} Sharpless~269 (hereafter, S269) is one of several H~$_{\mathrm{II}}$ regions in the outer Galaxy. Astrophysical masers in quantum transitions of water molecules have been detected in one of the compact infrared sources, IRS~2w \citep{Lo1973, Genzel1977}. The water masers were monitored from 2004 to 2006 with the VLBI Exploration of Radio Astrometry (VERA) of the National Astronomical Observatory of Japan (NAOJ) to obtain the annual parallax. A sharp-peaked spectrum at $V_{\mathrm{LSR}}$ of 19.7~km~s$^{-1}$ was observed at the position of S269~IRS~2w, and the annual parallax of the emission was reported to be $0.189 \pm 0.008$~mas using model fitting for the eastward sinusoidal motion, corresponding to $5.28^{+0.24}_{-0.22}$~kpc \citep[][hereafter H2007]{Honma2007}. However, this value is considerably higher than distance estimates obtained using other methods. For example, kinematic distance estimates based on the radial velocity of CO molecules provided value of 3.7~kpc \citep[see,][] {Wouterloot1989}. The latest kinematic distance estimated by \cite{Xu2009} yielded 3.7~kpc using the previously adopted values of galactic rotation parameters and 3.0~kpc using the values improved by \cite{Reid2009}. This estimate was based on the radial velocity of the CS(2--1) molecular radio line which traces dense cores associated with H~$_{\mathrm{II}}$ regions \citep{Bronfman1996}. Although kinematic distance estimates are uncertain in general, they considerably exceed values of the parallax measurements, especially in the outer Galaxy, and S269 was the only significant exclusion from this rule \citep{Reid2009}. Another independent method of the distance estimate based on a distance modulus of the luminous star from the S269 stellar cluster gives 3.8~kpc \citep{Moffat1979}. Therefore, H2007's value is 40\%--80\% larger than all the above distance estimates, and thus there exists an unusually large discrepancy. Another aspect which we have to carefully consider for S269 is its three-dimensional (3D) motion in the Milky Way. H2007 calculated the 3D motion from the absolute proper motion and radial velocity of the maser emission. They suggested that S269 has a very small peculiar motion with respect to a Milky Way flat rotation curve. On the other hand, Miyoshi~et~al. (2012, hereafter M2012) showed, from the same data as H2007, that the water maser emissions of S269~IRS~2w are widely distributed in space and radial velocity, and have various proper motions. A systemic proper motion of S269 can be different from the absolute proper motion of the single maser emission because it may be a part of a complicated internal motion such as an outflow with a typical speed of a few tens of km~s$^{-1}$ from a massive protostar. It is worthwhile to reanalyze the archival data to intensively inspect whether S269 is really in line with the Galactic rotation by imaging a wide area of the maser emitting region. In order to revisit the distance to S269 and its 3D motion, we extensively analyzed the VERA archival data, a part of which was already published by H2007 and M2012, including those from the follow-up observations. We describe the VERA monitoring program of S269~IRS~2w in Section~{\ref{sec:02}}. The data reduction procedure is presented in Section~{\ref{sec:03}. Our astrometric analysis results are described in Section~{\ref{sec:04}}. We discuss the systemic proper motion and the 3D motion of S269~IRS~2w in Section~{\ref{sec:05}}, and summarize this study in Section~{\ref{sec:06}}. We often refer results of a specific epoch observation (2005 March 14, or ``epoch C") to compare our data reduction process with that previously reported by H2007. In this paper, we adopted a line-of-sight systemic velocity of $17.7$~km~s$^{-1}$ and its standard deviation of $3.6$~km~s$^{-1}$ determined from CO molecular line observations \citep{Carpenter1990} with the local standard of rest (LSR) defined by \cite{Kerr1986}. Note that we discuss the 3D motion of S269 in Section~\ref{sec:05} with respect to the solar motion reported by \cite{Schonrich2010}. Hereafter we define a maser ``spot" as an emission in a single velocity channel and a maser ``feature" as a group of spots observed in at least two consecutive velocity channels at a coincident or at very closely located positions \cite[e.g.,][]{Imai2002}. | \label{sec:06} We presented the data reduction results of the VERA archival data of S269~IRS~2w, part of which was already analyzed by H2007 and M2012. The water masers ranging in LSR velocity from 4 to 21~km~s$^{-1}$ spatially distributed along the northeast--southwest direction. We obtained the annual parallax of $0.247 \pm 0.034$~mas for a maser feature with simple morphology residing in this star-forming region, corresponding to $4.05^{+0.65}_{-0.49}$~kpc. In order to consider the systemic proper motion of S269~IRS~2w we presented two major hypotheses regarding the dynamical center position: within the aligned maser structure and beside it. We found that at present these two cases cannot be distinguished by reasoning based on both statistical and physical plausibility. Because we have not had clear evidence on the certain location of the source of maser excitation, we found it impossible to draw any conclusion about the systemic proper motion. Instead, using the mean motion of the water maser proper motion, ($0.39 \pm 0.92$,~$-1.27 \pm 0.90$) mas~yr$^{-1}$, the annual parallax of $0.247 \pm 0.034$~mas, and radial velocity of the S269 molecular cloud, the peculiar motion of S269~IRS~2w is estimated to be ($U_{\mathrm{s}}$,~$V_{\mathrm{s}}$,~$W_{\mathrm{s}}$) of ($8 \pm 6$, $-21 \pm 17$, $1 \pm 18$)~km~s$^{-1}$. | 14 | 4 | 1404.6947 |
1404 | 1404.3369_arXiv.txt | {The spectra of the extended narrow-line regions (ENLRs) of Seyfert 2 galaxies probe the physics of the central active galaxy nucleus (AGN), since they encode the energy distribution of the ionising photons, the radiative flux and radiation pressure, nuclear chemical abundances and the mechanical energy input of the (unseen) central AGN.} {We aim to constrain the chemical abundance in the interstellar medium of the ENLR by measuring the abundance gradient in the circum-nuclear \ion{H}{ii} regions to determine the nuclear chemical abundances, and to use these to in turn determine the EUV spectral energy distribution for comparison with theoretical models.} {We have used the Wide Field Spectrograph (WiFeS) on the ANU 2.3m telescope at Siding Spring to observe the nearby, nearly face-on, Seyfert 2 galaxy, NGC~5427. We have obtained integral field spectroscopy of both the nuclear regions and the \ion{H}{ii} regions in the spiral arms. The observed spectra have been modelled using the MAPPINGS IV photoionisation code, both to derive the chemical abundances in the \ion{H}{ii} regions and the Seyfert nucleus, and to constrain the EUV spectral energy distribution of the AGN illuminating the ENLR.} {We find a very high nuclear abundance, 3.0 times solar, with clear evidence of a nuclear enhancement of N and He, possibly caused by massive star formation in the extended ($\sim 100$pc) central disk structure. The circum-nuclear narrow-line region spectrum is fit by a radiation pressure dominated photoionisation model model with an input EUV spectrum from a Black Hole with mass $5\times10^7 M_{\odot}$ radiating at $\sim 0.1$ of its Eddington luminosity. The bolometric luminosity is closely constrained to be $\log L_{\mathrm bol.} = 44.3\pm 0.1$ erg~s$^{-1}$. The EUV spectrum characterised by a soft accretion disk and a harder component extending to above 15keV. The ENLR region is extended in the NW-SE direction. The line ratio variation in circum-nuclear spaxels can be understood as the result of mixing \ion{H}{ii} regions with an ENLR having a radius-invariant spectrum.} {} | Seyfert Galaxies, like their more luminous cousins the quasars, contain at their nucleus an accretion disk around a supermassive black hole. As a result of mass accretion into the black hole, their nuclei produce both copious amounts of EUV photons which in turn excite extended narrow-line regions (ENLRs). In addition, energetic bipolar jets of relativistic plasma are produced, which may also shock and energise the ENLR. According to the standard unified model of AGN \citep{Antonucci:90apj,Antonucci:93araa} and its extensions \citep{Dopita:97pasa}, the Seyfert~1 galaxies are seen pole-on relative to the accretion disk, and these display very broad permitted lines originating in rapidly moving gas close to the central engine. In the Seyfert~2 galaxies, the thick accretion disk obscures the central engine, and an ENLR -- often confined within an ``ionisation cone" -- is observed. Within this general paradigm for the formation of the ENLR, many issues remain to be resolved: \begin{enumerate} \item {Do Seyfert~2s differ in spatial extent and cone opening angle from the Seyfert~1s \citep{Clarke:98apj,Schmitt:03apjs,Schmitt:03apj}?} \item {What is the relative energy flux in the EUV continuum compared with the jets, and are the radio-loud objects (which clearly have relativistic jets) more kinematically disturbed \citep{Bicknell:98apj,Evans:99apj, Wilson:99apj}?} \item {What mechanisms control the non-thermal EUV spectrum, and can these be constrained by observation \citep{Bland-Hawthorn:97apss,Allen:99apj,Done:12mnras,Jin:12a,Jin:12b,Jin:12c}}? \item{What is the ionisation parameter in the ENLR, how well can this be constrained by the coronal lines \citep{MullerSanchez:11apj} or by other line ratio diagnostics?} \item{Are all ENLRs dominated by radiation pressure acting on dust \citep{Dopita:02apj,Groves:04apjs, Groves:04apj}?} \item {What is the chemical abundance distribution in Seyfert nuclei, how well can this be constrained by observations of the surrounding \ion{H}{ii} regions \citep{Evans:87apj}, and how is this correlated with the host galaxy mass? } \item {To what extent is Seyfert activity triggered by tidal interactions between galaxies, and what is the role of mergers in feeding the AGN, and producing nuclear starbursts?} \end{enumerate} To address and understand these issues, especially items 3--6 (above), we have undertaken the {Siding Spring Southern Seyfert Spectroscopic Snapshot Survey} (S7). The S7 survey is an integral field survey in the optical of over 100 southern Seyfert galaxies. The survey uses the Wide Field Spectrograph (WiFeS) mounted on the Nasmyth focus of the ANU 2.3m telescope \citep{Dopita:2010aa}, and will be described in detail in another paper. However, here we present detailed WiFeS observations of the Seyfert~2 galaxy NGC~5427 with the particular objective of providing answers to the questions 3 -- 6 (above) for this galaxy. NGC~5427 is a giant Sc-type spiral located at a distance of $\sim 40$~Mpc (corresponding to a spatial scale of 177~pc~arcsec$^{-1}$). It contains a Seyfert 2 nucleus \citep{Veron-Cetty:2006aa}. Together with NGC~5426, it forms the interacting pair Arp~271. The interaction between NGC~5427 and NGC~5426 at 20~kpc separation has resulted in a tidal bridge between both galaxies, which is seen in the UV and in H$\alpha$ \citep{Evans:1996aa, Smith:2010aa}. The gas kinematics in the interacting pair has been studied using H$\alpha$ kinematical maps obtained with Fabry Perot systems \citep{Fuentes-Carrera:2004aa, Hernandez:2008aa, Font:2011aa}. \cite{Font:2011aa} find evidence for a transfer of hydrogen gas from NGC~5426 to NGC~5427 as well as for a galactic wind in NGC~5427. \cite{Fuentes-Carrera:2004aa} present a detailed study of the orbital configuration of the interaction between NGC~5427 and NGC~5426, providing an estimate of the total dynamical mass of both galaxies within D$_{25}$/2 of $6.72$-$11.2\times 10^{10}\ M_\odot$ for NGC~5426 and $4.5$-$7.5\times 10^{10}\ M_\odot$ for NGC~5427, suggesting that NGC~5426 is the more massive galaxy of the pair within the optical radius. The host galaxy of NGC~5427 has been classified as SA(s)c pec by \cite{de-Vaucouleurs:1991aa} and is seen nearly face-on. Based on ellipse fitting in the $B$-band, \cite{Marinova:2007aa} report an inclination of $i= 38\deg$ and a position angle of $PA = 11\deg$, although the tidal interactions may make these values rather uncertain. The galaxy stellar mass is estimated to be $M_\star = 4.61\times 10^{10}\ M_\odot$ \citep{Weinzirl:2009aa}, based on the $M_\star$-$(B-V)$ relation from \cite{Bell:2003aa}. Using an SED model for the global UV, optical, and far-infrared (far-IR) emission of NGC~5427, \cite{Misiriotis:2004aa} derive a star-formation rate of $\log \mathrm{SFR} = 0.95\ M_\odot\ \mathrm{yr^{-1}}$, a dust mass of $\log M_\mathrm{d} = 7.83\ M_\odot$ and report a gas mass of $\log M_\mathrm{g} = 10.38\ M_\odot$. NGC~5427 has been variably classified as barred or un-barred \citep{Fuentes-Carrera:2004aa, Marinova:2007aa, Weinzirl:2009aa, Comeron:2010aa}. It has a nuclear ring or pseudo-ring \cite[][and references therein]{Comeron:2010aa} as well as 4-5 nuclear dust spirals in the inner 700~pc \citep{Martini:2003aa}. The circum-nuclear ring of \ion{H}{ii} regions is at a distance of about 1~kpc from the nucleus in H$\alpha$ images \citep{Evans:1996aa, Gonzalez-Delgado:1997aa}. NGC~5427 hosts a $\sim 10^7\ M_\odot$ super-massive black hole. By modelling the width of ionised gas emission lines in HST spectra, \cite{Beifiori:2009aa} derive upper limits for the black-hole mass of NGC~5427 of $8.1\times 10^7\ M_\odot$ or $2\times 10^7\ M_\odot$ for a Keplerian disk model assuming inclinations of $i=33\deg$ and $i=81\deg$, respectively. If the accretion disk inclination was the same as the disk inclination of $i=38\deg$ \citep{Marinova:2007aa}, then the BH mass would be $\sim 7\times 10^7\ M_\odot$. \cite{Woo:2002aa} estimate a black hole mass of $2.45\times 10^6\ M_\odot$, using the $M_\mathrm{BH}$-$\sigma_\star$ relation from \cite{Tremaine:2002aa} and a stellar velocity dispersion of $\sigma_\star = 74\ \mathrm{km\ s^{-1}}$. However, this is an indirect method, and therefore subject to greater uncertainty. They also report a bolometric luminosity of $\log (L_\mathrm{bol}) = 44.12\ \mathrm{erg\ s^{-1}}$, estimated by flux integration over the measured (UV to far-IR) spectral energy distribution (SED). If all this luminosity were due to the black hole itself, it would correspond to an Eddington ratio of $\sim 0.4$ with their black hole mass estimate. However, the \cite{Woo:2002aa} estimate is likely to be an underestimate of the true bolometric luminosity when the EUV contribution is taken into account. In this paper, we derive a model for the EUV SED and luminosity of the accreting supermassive black hole in NGC~5427 based upon a study of the gas excitation and physical conditions in the ENLR surrounding the AGN in NGC~5427. This model is constrained by chemical abundance measurements derived from \ion{H}{ii} region spectra in the host galaxy. The paper is organised as follows. In Section \ref{obs} we present details of our observations and the data reduction techniques employed, in Section \ref{results} we present emission maps, observed line ratios, and derive the total abundances estimated from the optical spectra of the \ion{H}{ii} regions. In Section \ref{nuc} we describe our photoionisation analysis technique, which we use to constrain the chemical abundances, pressure and ionisation parameter in the NLR, and to estimate the shape of the EUV spectrum. | In this paper we have presented high-quality IFU observations of the Seyfert galaxy NGC~5427. From these we have been able to measure the abundance gradient accurately and so constrain the chemical abundances of the Seyfert 2 nucleus. Using these abundances, we have built a detailed photoionisation model for the nuclear region. We have identified the sensitivity of the various observed line ratios to the input parameters of the model; the intensity of the hard X-ray component relative to the accretion disk component of the EUV/X-ray SED. We found no evidence for an intermediate-energy Comptonised component in this object. We discovered that, in order to obtain an excellent fit to the measured emission-line ratios, the He abundance needed to be enhanced above that found from the \ion{H}{ii} regions by +0.16~dex, and the N abundance by +0.3~dex. This implies that massive fast-rotating stars are being formed within the accretion flow near the black hole and chemically polluting both the extended ($\sim 100$pc) outer disk structure and the ENLR. Such stars could potentially provide a source of viscosity in the disk, enhancing the accretion rate between 100 and 10~pc (approximately), and possibly providing a solution to the ``angular momentum problem''. The nuclear luminosity is loosely constrained by the models; $\log L_{\mathrm bol.} = 44.3\pm 0.1$ erg~s$^{-1}$, of which about 25\% is able to ionise hydrogen. The far-IR to UV flux estimated from the model, $\log L = 44.17 \pm 0.10$ erg~s$^{-1}$, agrees closely with what has been estimated by direct integration of the observed fluxes over these wavebands ($\log L = 44.12$ erg~s$^{-1}$, \citet{Woo:2002aa}). The best fit model has a Black Hole with a mass of $5\times10^7 M_{\odot}$ radiating at $\sim 0.1$ of its Eddington luminosity. Finally, we have established that the spectra observed for individual spaxels in the ENLR of NGC~5427 out to a radius of $\sim 2$~kpc can be understood as a mixing between the ENLR and background star formation activity giving rise to normal \ion{H}{ii} region-like spectra. The ENLR is extended more in the NW-SE direction, roughly perpendicularly to the nuclear ring of star formation. In this sense it forms a classical ``ionisation cone". The methodology we have established in this paper is generic and applicable to other Seyfert galaxies and their ENLR. It establishes the utility of using the \ion{H}{ii} regions to constrain the chemical abundances in the nucleus, and providing strict limits on the form of the ionising spectrum of the central engine. | 14 | 4 | 1404.3369 |
1404 | 1404.6250_arXiv.txt | We present the results of an X-ray analysis of 80 galaxy clusters selected in the 2500 deg$^2$ South Pole Telescope survey and observed with the \emph{Chandra X-ray Observatory}. We divide the full sample into subsamples of $\sim$20 clusters based on redshift and central density, performing a joint X-ray spectral fit to all clusters in a subsample simultaneously, assuming self-similarity of the temperature profile. This approach allows us to constrain the shape of the temperature profile over $0 < r < 1.5R_{500}$, which would be impossible on a per-cluster basis, since the observations of individual clusters have, on average, 2000 X-ray counts. The results presented here represent the first constraints on the evolution of the average temperature profile from $z = 0$ to $z = 1.2$. We find that high-$z$ ($0.6 < z < 1.2$) clusters are slightly ($\sim$30\%) cooler both in the inner ($r<0.1R_{500}$) and outer ($r>R_{500}$) regions than their low-$z$ ($0.3 < z<0.6$) counterparts. Combining the average temperature profile with measured gas density profiles from our earlier work, we infer the average pressure and entropy profiles for each subsample. Confirming earlier results from this data set, we find an absence of strong cool cores at high $z$, manifested in this analysis as a significantly lower observed pressure in the central $0.1R_{500}$ of the high-$z$ cool-core subset of clusters compared to the low-$z$ cool-core subset. Overall, our observed pressure profiles agree well with earlier lower-redshift measurements, suggesting minimal redshift evolution in the pressure profile outside of the core. We find no measurable redshift evolution in the entropy profile at $r\lesssim0.7R_{500}$ -- this may reflect a long-standing balance between cooling and feedback over long timescales and large physical scales. We observe a slight flattening of the entropy profile at $r\gtrsim R_{500}$ in our high-$z$ subsample. This flattening is consistent with a temperature bias due to the enhanced ($\sim$3$\times$) rate at which group-mass ($\sim$2\,keV) halos, which would go undetected at our survey depth, are accreting onto the cluster at $z\sim1$. This work demonstrates a powerful method for inferring spatially-resolved cluster properties in the case where individual cluster signal-to-noise is low, but the number of observed clusters is high. | \setcounter{footnote}{0} Galaxy clusters, despite what the name implies, consist primarily of matter that is not in galaxies. A typical cluster is well-modeled by a central dark matter halo ($\sim$85\% by mass) and a diffuse, optically-thin plasma ($\sim$15\% by mass). The response of this hot ($\gtrsim10^7$~K) plasma, known as the intracluster medium (ICM), to the evolving gravitational potential is one of our best probes of the current state and evolution of galaxy clusters. X-ray imaging and spectroscopy of the ICM allow estimates of the cluster mass profile via the spectroscopic temperature and gas density \citep[e.g.,][]{forman82,nevalainen00,sanderson03,arnaud05,kravtsov06, vikhlinin06a, arnaud07}, the enrichment history of the cluster via the ICM metallicity \citep[e.g.,][]{deyoung78,matteucci86,deplaa07,bregman10,bulbul12}, the cooling history via the cooling time or entropy \citep[e.g.,][]{white97, peres98, cavagnolo08, mcdonald13b}, the feedback history via the presence of X-ray bubbles \citep[e.g.,][]{rafferty06, mcnamara07,rafferty08,hlavacek12}, and the current dynamical state and merger history of the cluster via the X-ray morphology \citep[e.g.,][]{jones79,mohr95,roettiger96,schuecker01,jeltema05,nurgaliev13}. While there is much diversity in the ICM from cluster to cluster, it is valuable to determine if there are broad similarities in clusters of a given mass and redshift. The construction of a ``Universal'' pressure profile, for example, can allow comparisons to simulated galaxy clusters, as well as provide a functional form for matched-filtering algorithms, such as those that are used to select galaxy clusters using the Sunyaev-Zel'dovich \citep[SZ;][]{sunyaev72} effect. Much effort has been made to quantify the average temperature \citep[e.g.,][]{loken02, vikhlinin06a, pratt07, leccardi08a, baldi12}, entropy \citep[e.g.,][]{voit05, piffaretti05,cavagnolo09, pratt10}, and pressure \citep[e.g.][]{arnaud10, sun11,bonamente12,planck13} profiles for low-redshift galaxy groups and clusters based on both X-ray and SZ selection. In all cases, the average profiles have a substantial amount of scatter at $r\lesssim0.2R_{500}$, due to the presence (or lack) of a cool, dense core \citep[e.g.,][]{vikhlinin06a,cavagnolo09,arnaud10}, but collapse onto the self-similar expectation at larger radii. This suggests that non-gravitational processes (e.g., cooling, AGN feedback) are important in the central region of the cluster while gravity is the dominant force in the outer region. While the aforementioned studies have made significant progress in quantifying the average temperature, entropy, and pressure profiles of galaxy groups and clusters, they have focused almost entirely on low-redshift ($z\lesssim0.2$) systems. This is in part due to the relative ease with which one can measure the temperature profile in nearby systems, but also due to the fact that, until recently, large, well-selected samples of galaxy clusters at high redshift did not exist. This has changed, with the recent success of large SZ surveys from the Atacama Cosmology Telescope \citep[ACT;][]{act11,act13}, \emph{Planck} \citep{planck11,planck13b}, and the South Pole Telescope \citep[SPT;][]{vanderlinde10,reichardt13}. These surveys have discovered hundreds of new galaxy clusters at $z>0.3$, allowing the study of galaxy cluster evolution for the first time out to $z>1$ using large, homogeneous data sets. In this paper, we present a joint-fit spectroscopic analysis of 80 SPT-selected galaxy clusters in the SPT-XVP sample \citep[][Benson \etal in prep]{mcdonald13b}. Utilizing uniform-depth X-ray observations of these clusters we can, for the first time, constrain the redshift evolution of the average ICM temperature, entropy, and pressure profiles. We present the details of this analysis in \S2, including the resulting projected and deprojected temperature profiles in \S2.2 and \S2.3, respectively. In \S3 we infer the average pressure (\S3.1) and entropy (\S3.2) profiles. In \S4 we discuss the implications of the observed evolution, specifically in the inner $\sim$100\,kpc and outskirts ($r\gtrsim R_{500}$) of the mean pressure and entropy profiles, and assess any potential biases in our analysis. Finally, we summarize these results in \S5 before suggesting future applications of these data. Throughout this work, we assume H$_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\Omega_M$ = 0.27, and $\Omega_{\Lambda}$ = 0.73. | Below, we discuss the results of \S2 and \S3, comparing these to previous observations and simulations in order to aid in their interpretation. In addition, we investigate potential systematic errors and/or biases which may conspire to influence our conclusions. \subsection{Comparison to Simulations} In Figure \ref{fig:univP+sims} we compare our average pressure profiles to those from simulations, as presented in \cite{battaglia12} and Dolag \etal (in prep). For the latter simulations, pressure profiles are computed and presented in Liu \etal (in prep). For each simulation, we show three redshift slices, similar to our low- and high-$z$ subsamples as well as a $z\sim0$ subsample for comparison to P13, and have made mass cuts similar to the SPT 2500 deg$^2$ survey selection function. All samples have been normalized by $\left<f(M)\right>$ (see \S3.1). We do not show comparisons for CC and NCC subsamples here, since i) we do not have cuspiness measurements for the simulated clusters, and ii) it is unlikely that the simulated clusters will span the full range of properties from cool core to non-cool core clusters. In general, simulations struggle to get the complicated balance between cooling and feedback right in the cores of clusters ($r<0.1R_{500}$), but perform well outside of the core where gravity dominates. As shown in Figure \ref{fig:univP+sims}, our measured pressure profiles and the pressure profiles from both sets of simulations agree reasonably well at $r > 0.1R_{500}$ for all redshift slices. At $r<0.1R_{500}$, the difference between the simulations and the data becomes worse with increasing redshift. The simulated clusters appear to have massive cool cores in place already at $z\sim1$, while the observed clusters are becoming more centrally concentrated over the past $\sim$8\,Gyr \citep{mcdonald13b}. At large radii, the best-fit profile is consistent with Dolag \etal (in prep), and slightly steeper than that predicted by \cite{battaglia12}, but all profiles are consistent at the 1$\sigma$ level with the data. We stress that any steepening of the pressure profile may be artificial, indicative of a bias due to clumping of the ICM at higher redshift, a point we will address below. We can also compare our observations of an unevolving entropy core (Figure \ref{fig:univK}) to simulations, this time by \cite{gaspari12} who focus on the delicate balance between AGN feedback and cooling in the cores of simulated galaxy clusters. These simulations demonstrate that, while at the very center ($<$10\,kpc) of the cluster the entropy can fluctuate significantly (factors of $\sim$2--3) on short (Myr) timescales, the entropy at $\gtrsim$20 kpc is relatively stable of $\sim$5\,Gyr timescales. These simulations, which reproduce realistic condensation rates of cool gas from the ICM, suggest that a gentle, nearly-continuous injection of mechanical energy from the central AGN is sufficient both to offset the majority of the cooling (preventing the cooling catastrophe) and to effectively ``freeze'' the entropy profile in place. Overall, the agreement between observations and simulations is encouraging. The primary difference between the two occurs at $r\lesssim0.1R_{500}$, with excess pressure in the simulated cores. As these are the radii where the complicated interplay between ICM cooling, bulk ICM motions, and AGN feedback is most important, it is perhaps unsurprising that the deviations between data and model are most severe in this regime. \subsection{Cluster Outskirts: Halo Accretion?} In recent years, a number of different studies have observed a flattening of the entropy profile for a number of different galaxy clusters at the virial radius \citep[e.g.,][]{bautz09,walker13,reiprich13,urban14}. This flattening, while not observed in all clusters \citep[e.g.,][]{eckert13}, has been attributed to clumping in the intracluster medium \citep[see e.g.,][]{simionescu11, nagai11, urban14}. If a substantial fraction of the ICM beyond the virial radius is in small, overdense clumps, the measured electron density ($n_e$) over a large annulus will be biased high, due to surface brightness being proportional to $n_e^2$. These clumps are thought to be the halos of infalling galaxies or small groups. Due to their low mass, they ought to be cooler than the ambient ICM, which could also lead to the measured temperature being biased low. Given that we measure, on average, lower temperatures (and entropies) at large radii ($r\gtrsim R_{500}$) in high-$z$ clusters, it is worth discussing whether this result could be driven by clumping and, specifically, how massive these clumps could be. If we assume that an extended source with $<$20 X-ray counts would go undetected against the diffuse cluster emission, we can estimate a limiting X-ray luminosity at $z=0.8$ of $L_X\sim2\times10^{43}$ erg s$^{-1}$, corresponding to a halo mass of $M_{500} \sim 8\times10^{13}$ M$_{\odot}$ and temperature of $\sim$2 keV \citep{vikhlinin09a}. Thus, it is quite possible that the measured temperature in the outskirts of clusters at $z>0.6$ is biased low due to our inability to detect and mask group-sized halos which are in the process of accreting onto the massive cluster. The entropy flattening that we measure in Figure \ref{fig:univK} is driven primarily by the evolution in the temperature profile (Figures \ref{fig:univT}--\ref{fig:univT_model}), with only a small, insignificant evolution measured in the outer part of the gas density profile \citep{mcdonald13b}. This makes sense, if the temperature profile is in fact biased by infalling $>10^{13}$ M$_{\odot}$ groups at $\sim$R$_{500}$. Figure \ref{fig:cartoon} illustrates this scenario, showing the density and temperature profiles for a typical SPT-selected cluster (M$_{500} = 6\times10^{14}$ M$_{\odot}$, $kT_{500} = 6.5$ keV), and an infalling group-sized system (M$_{500} = 6\times10^{13}$ M$_{\odot}$, $kT_{500} = 1.5$ keV). For simplicity, we assume that the infalling group is isothermal and constant density, with $\rho_g = M_{g,500}/\frac{4}{3}\pi R_{500}^3$, where both M$_{g,500}$ and $R_{500}$ can be derived from the group mass, assuming a gas fraction of 0.12. This simple test shows that, at $r\gtrsim 1.7R_{500}$, group-sized halos will significantly bias the measured density high, while at $r\lesssim 1.7R_{500}$ they will bias the measured temperature low. At $\sim$R$_{500}$, where we measure a flattening of the entropy profile, the density of the infalling group and the ambient ICM are roughly equal, with a factor of $\sim$3 difference in temperature. This temperature contrast would result in an artificial steepening of the temperature profile, as we observe in Figures \ref{fig:univT}--\ref{fig:univT_model}). Following \citep{vikhlinin06b}, we estimate that the group-sized halos would need to contribute $\sim$30--40\% of the total X-ray counts in the outer annuli to bias the temperature low by the observed 40\%, with the exact fraction depending on the relative temperature of the cluster and group. \begin{figure}[htb] \centering \includegraphics[width=0.49\textwidth]{cartoon.eps} \caption{Idealized depiction of a group-sized (M$_{500} = 6\times10^{13}$ M$_{\odot}$; blue lines) halo falling into a massive (M$_{500} = 6\times10^{14}$ M$_{\odot}$; red lines) galaxy cluster. The infalling group is assumed to be isothermal and constant density, with the density equal to $\rho_g = 0.12$M$_{500}/\frac{4}{3}\pi R_{500}^3$ and temperature taken from the M--T$_X$ relation \citep{vikhlinin09a}. This figure demonstrates that, as a group-sized halo falls into a massive cluster, it will first significantly bias the density high at $r\gtrsim 1.7R_{500}$ (right of dashed vertical line), and then bias the temperature low at $r\lesssim 1.7R_{500}$ (left of dashed vertical line). The latter effect may be driving the steep temperature profile (Figures \ref{fig:univT}--\ref{fig:univT_model}) and entropy flattening (Figure \ref{fig:univK}) that we observe in high-$z$ clusters. } \label{fig:cartoon} \end{figure} Simulations suggest that at $z\gtrsim1$ there is significantly more massive substructure in the outskirts of galaxy clusters. For example, \cite{tillson11} find that the accretion rate onto massive halos evolves by a factor of $\sim$3.5 from $z\sim1.5$ to $z\sim0$, while \cite{fakhouri10} find that 10$^{14}$ M$_{\odot}$ halos are accreting 10$^{13}$ M$_{\odot}$ subhalos at a rate $\sim$3 times higher at $z\sim1$ than at $z\sim0$. These results suggest that the entropy flattening which we measure (Figure \ref{fig:univK}) is consistent with a temperature bias due to our inability to detect (and mask) large substructures in the outskirts of SPT-selected clusters. We stress that this ``superclumping'' is qualitatively different than the ``clumping'' inferred in nearby clusters \citep[e.g.,][]{simionescu11,nagai11,urban14}, which is commonly interpreted as large numbers of small subhalos raining onto clusters at the virial radius, where group-sized halos would be detected and masked. \subsection{Cool Core Evolution} In an earlier analysis of this dataset \citep{mcdonald13b}, we saw evidence for evolution in the central gas density of cool cores over the past 8~Gyr but no evidence that the minimum ICM entropy in the central $\sim$10~kpc had evolved since $z \sim 1$, maintaining a floor at $\sim$10 keV cm$^2$. Now, with a more rigorous joint-fit analysis to constrain the central temperature, providing a more accurate estimate of the central entropy, we revisit this result. From Figure \ref{fig:univK}, we see no measureable evolution in the central entropy bin ($0<r<0.04R_{500}$), from $K/K_{500} = 0.095_{-0.02}^{+0.02}$ at low-$z$ to $K/K_{500} = 0.102_{-0.01}^{+0.02}$ at low-$z$. Indeed, the average cool core entropy profile shows no evidence for evolution interior to $r<0.7R_{500}$ since $z\sim1$ (Figure \ref{fig:Kevol}). In the absence of feedback or redistribution of entropy, one would expect the average entropy to drop rapidly in the cores of these clusters, on Gyr or shorter timescales. Given the 5~Gyr spanned by this sample, and the consistency with the $z\sim0$ work by \cite{pratt10}, we can argue that some form of feedback is precisely offsetting cooling between $z\sim1$ and $z\sim0$. Specifically, as the central gas density increases, the core temperature also increases. This trend is contrary to what one would expect from simple hydrostatic equilibrium in a dark matter-dominated halo, but is consistent with the expectation for adiabatic compression of the gas. A likely culprit for this heat injection is radio-mode feedback \citep[e.g.,][]{churazov01, fabian12,mcnamara12}, which has been shown to be operating steadily over similar timescales \citep{hlavacek12}. Indeed, \cite{gaspari11} demonstrate that the immediate result of a burst of AGN feedback is to increase the core temperature of the gas, while leave the large-scale ($r\gtrsim0.1R_{500}$) distribution of temperatures unchanged. We finish by stressing that this work and that of \cite{mcdonald13b} refer to the entropy in the inner $\sim$40~kpc as the ``central entropy''. This annulus, which contains all of the lowest entropy gas falling onto the central cluster galaxy, is limited in size by our relatively shallow exposures. Indeed, \cite{panagoulia13} show that with improved angular resolution the entropy continues to drop toward the central AGN. Our discussion of an ``entropy floor'' is always referring to a fixed radius, within which the mean entropy is not evolving. \subsection{Systematic Biases/Uncertainties} Below we briefly address three potential issues with our data analysis: whether the low signal-to-noise in cluster outskirts is driving the entropy flattening, whether joint spectral fitting yields the same results as averaging individual fits, and whether the average temperature profile is mass-dependent. \subsubsection{X-ray Spectrum Signal-to-Noise} While our observing program was designed to obtain 2000 X-ray counts per cluster, a variety of effects conspired to create the scatter in the observed number of counts per cluster (see Figure \ref{fig:sample}). These factors include uncertainties in the $\xi$--L$_X$ relation, uncertainties on early redshift measurements, and the presence or lack of a cool core. Here, we investigate how strongly the measured average entropy profile depends on the S/N of the included observations. \begin{figure}[h!] \centering \includegraphics[width=0.49\textwidth,trim=0cm 1cm 0cm 0cm]{univK_bycounts.eps} \caption{Joint-fit entropy profile for both the low- and high-$z$ subsamples (see also Figure \ref{fig:univK}). The red and blue points correspond to the joint-fit profiles for low- and high-S/N subsamples, respectively, as described in \S4.4.1. We find that, at large radii, the flattening of the entropy profile correlates with both increasing redshift and decreasing S/N. The most significant flattening is present in the high-$z$, low-S/N subsample, which contains 7 of the 8 clusters at $z>1$ and all four $z>1.1$ clusters. Given that the low-$z$ and high-$z$ low-S/N subsamples have similar S/N but different degrees of flattening, we propose that the observed flattening is driven primarily by increasing redshift. } \label{fig:univK_bycounts} \end{figure} In Figure \ref{fig:univK_bycounts} we have divided the low-$z$ and high-$z$ subsamples by the S/N in the three outermost bins ($r>0.75R_{500}$), specifically so that we can test whether the observed entropy flattening is a function of S/N. For the low-S/N subsamples, there are a total of $\sim$1000 X-ray counts in each of the three outermost bins and $\sim$2700 counts per radial bin over the full radial range, compared to $\sim$2800 (outer) and $\sim$4600 (full radial range) per bin for the high-S/N subsamples. For the low-$z$ clusters, the measured entropy profile appears to be independent of the S/N -- the difference of a factor of $\sim$2 in the total number of X-ray counts used in the spectral modeling does not appear to have a significant affect on the resulting entropy profile. For the high-$z$ clusters, the low- and high-S/N profiles are identical at $r < 0.6 R_{500}$, with more flattening at larger radii in the low-S/N clusters. Since the low-S/N clusters also tend to be higher redshift (the high-$z$, low-S/N subsample contains 7 of the 8 clusters at $z>1$ and all 4 clusters at $z>1.1$), it is not clear which effect is most responsible for the flattening. In general, there is a trend of more flattening going to both higher redshift and lower S/N. We do not expect a significant bias from low cluster counts, due to our background modeling on an observation-by-observation basis (\S 2.2), but we can not rule out this possibility. Given that the low-$z$, low-S/N clusters have equally low S/N to the high-$z$, low-S/N clusters, we suggest that the flattening is more significantly driven by redshift evolution. \subsubsection{Joint-Fitting Versus Profile Averaging} To test whether our joint-fitting technique is introducing a systematic bias, we compute individual temperature profiles for our low-$z$ subsample (Figure \ref{fig:allfits}). Given that each annulus has on the order of $\sim100$ X-ray counts, these individual fits are poorly constrained. However by averaging $\sim$40 profiles (unweighted), we can constrain the average temperature profile for this subsample. For comparison, we show the results of our joint-fit analysis for the same clusters. We find that the joint-fit method and the averaging method yield consistent results. Since the uncertainty on the joint-fit analysis is really the scatter in the mean for a number of realizations (black points), we have shown the standard error on the mean (standard deviation divided by $\sqrt{N}$) in the average profile (red points) in order to make a fair comparison. This simple test confirms that our method of joint-fitting multiple spectra is largely unbiased with respect to the true average profile. Naively, one might expect a joint-fit analysis to be biased towards the highest signal-to-noise spectra, since each cluster is essentially weighted by its total X-ray counts, while each cluster is weighted equally in the averaging method. However, this test shows that any bias that would be imparted by joint-fitting spectra of varying signal-to-noise is offset by randomly drawing and fitting subsamples of spectra. \subsubsection{Mass Bias} \cite{vikhlinin06a} show that, for a sample of relaxed, low-redshift clusters, low-mass systems tend to have higher central temperatures than their high-mass counterparts. We explore this idea in Figure \ref{fig:univT_Mg} by dividing our high-$z$ subsample by total gas mass, $M_{g,500}$, rather than by cuspiness (the following results hold for the low-$z$ subsample as well). This figure confirms that the temperature profiles of galaxy clusters are not self-similar at $r\lesssim0.3R_{500}$. We find that low mass systems have temperatures $\sim$20--30\% higher in their cores, consistent with work by \cite{vikhlinin06a} which covered a larger mass range. At $r>0.3R_{500}$ there appear to be no deviations from self-similarity, suggesting that non-gravitational processes are most likely driving the differences in the core. This figure demonstrates how important a well-selected sample is for such a joint-fit analysis to be successful and yield results representative of the true population. We expect that, given the similar mass distribution of our low- and high-$z$ subsamples (see Figure \ref{fig:sample}), this mass bias is not driving any of the trends discussed in \S3. \begin{figure}[htb] \centering \includegraphics[width=0.49\textwidth]{univT_allfits.eps} \caption{This figure demonstrates the similarity in the average temperature profile (red) and the ``joint-fit'' profile (black; see \S2.2). Individual cluster profiles are shown as red dashes, while the average of these profiles is shown as thick red points. The uncertainty shown for the average profile is the standard error on the mean (standard deviation divided by $\sqrt{N}$) to allow a better comparison to the joint-fit uncertainties, which are measuring the scatter in the mean temperature for a large number of realizations. The joint-fit result, which is fully consistent with the average profile, is shown in black. This figure demonstrates that our joint-fit analysis is not strongly affected by combining spectra of varying signal-to-noise.} \label{fig:allfits} \end{figure} \begin{figure}[h!] \centering \includegraphics[width=0.49\textwidth]{univT_Mg.eps} \caption{Average temperature profiles for high-$z$ clusters. We show the combined fits in grey, low-mass systems in blue, and high-mass systems in red. This figure demonstrates that the deviation from self-similarity interior to $0.3R_{500}$, consistent with earlier work by \cite{vikhlinin06a}, is present out to $z\sim1$. Beyond $0.3R_{500}$, there is no evidence for a mass bias.} \label{fig:univT_Mg} \end{figure} It is also possible that our use of the $Y_{X,500}$--M$_{500}$ relation to infer R$_{500}$ could impart a bias in these results, if the assumed evolution on this relation is incorrect. To investigate this potential bias, we re-extracted spectra using R$_{500}$ estimates based on the M$_{gas}$--M$_{500}$ relation, and repeated the analysis described in \S2.2. The resulting temperature profiles were consistent with what we have presented here, suggesting that our assumed evolution on the $Y_{X,500}$--M$_{500}$ relation is appropriate out to $z\sim1.2$. | 14 | 4 | 1404.6250 |
1404 | 1404.6066_arXiv.txt | We present the ultra-deep Subaru narrowband imaging survey for \lya\ emitters (LAEs) at $\zi=7.3$ in SXDS and COSMOS fields ($\sim 0.5$ deg$^2$) with a total integration time of 106 hours. Exploiting our new sharp bandwidth filter, $\textit{NB101}$, installed on Suprime-Cam, we have reached $L( \mathrm{Ly}\alpha ) = 2.4 \times 10^{42} \ \mathrm{erg} \ \mathrm{s}^{-1}$ ($5\sigma$) for $\zi=7.3$ LAEs, about 4 times deeper than previous Subaru $\zi \gtrsim 7$ studies, which allows us to reliably investigate the evolution of the \lya\ luminosity function (LF), for the first time, down to the luminosity limit same as those of Subaru $z=3.1-6.6$ LAE samples. Surprisingly, we only find three and four LAEs in SXDS and COSMOS fields, respectively, while one expects a total of $\sim 65$ LAEs by our survey in the case of no \lya\ LF evolution from $\zi = 6.6$ to $7.3$. We identify a decrease of the \lya\ LF from $\zi =6.6$ to $7.3$ at the $> 90\%$ confidence level from our $\zi = 7.3$ \lya\ LF with the best-fit Schechter parameters of $L^{*}_{\mathrm{Ly}\alpha} = 2.7^{+8.0}_{-1.2} \times 10^{42} \ \mathrm{erg} \ \mathrm{s}^{-1}$ and $\phi^{*} = 3.7^{+17.6}_{-3.3} \times 10^{-4} \ \mathrm{Mpc}^{-3}$ for a fixed $\alpha = -1.5$. Moreover, the evolution of the \lya\ LF is clearly accelerated at $\zi > 6.6$ beyond the measurement uncertainties including cosmic variance. Because no such accelerated evolution of the UV-continuum LF or the cosmic star-formation rate (SFR) is found at $z\sim 7$, but suggested only at $z>8$ \citep{2013ApJ...773...75O, 2014arXiv1403.4295B}, this accelerated \lya\ LF evolution is explained by physical mechanisms different from a pure SFR decrease but related to the \lya\ production and escape in the process of cosmic reionization. Because a simple accelerating increase of IGM neutral hydrogen absorbing \lya\ would not reconcile with Thomson scattering optical depth measurements from \textit{WMAP} and \textit{Planck}, our findings may support new physical pictures suggested by recent theoretical studies, such as the existence of {\sc Hi} clumpy clouds within cosmic ionized bubbles selectively absorbing \lya\ and the large ionizing photon escape fraction of galaxies making weak \lya\ emission. | 14 | 4 | 1404.6066 |
||
1404 | 1404.3705_arXiv.txt | Magnetars are a special type of neutron stars, considered to have extreme {\it dipole} magnetic fields reaching $\sim 10^{11}$ T. The magnetar 4U 0142+61, one of prototypes of this class, was studied in broadband X-rays (0.5--70 keV) with the {\it Suzaku} observatory. In hard X-rays (15--40 keV), its 8.69 sec pulsations suffered slow phase modulations by $\pm 0.7$ sec, with a period of $\sim 15$ hours. When this effect is interpreted as free precession of the neutron star, the object is inferred to deviate from spherical symmetry by $\sim 1.6 \times 10^{-4}$ in its moments of inertia. This deformation, when ascribed to magnetic pressure, suggests a strong {\it toroidal} magnetic field, $\sim10^{12}$ T, residing inside the object. This provides one of the first observational approaches towards toroidal magnetic fields of magnetars. | 14 | 4 | 1404.3705 |
||
1404 | 1404.5969.txt | We use direct $N$-body calculations to study the evolution of the unusually extended outer halo globular cluster Palomar 4 (Pal~4) over its entire lifetime in order to reproduce its observed mass, half-light radius, velocity dispersion and mass function slope at different radii. We find that models evolving on circular orbits, and starting from a non-mass segregated, canonical initial mass function (IMF) can reproduce neither Pal 4s overall mass function slope nor the observed amount of mass segregation. Including either primordial mass segregation or initially flattened IMFs does not reproduce the observed amount of mass segregation and mass function flattening simultaneously. Unresolved binaries cannot reconcile this discrepancy either. We find that only models with both a flattened IMF \textit{and} primordial segregation are able to fit the observations. The initial (i.e. after gas expulsion) mass and half-mass radius of Pal~4 in this case are about 57000 M${\odot}$ and 10 pc, respectively. This configuration is more extended than most globular clusters we observe, showing that the conditions under which Pal~4 formed must have been significantly different from that of the majority of globular clusters. We discuss possible scenarios for such an unusual configuration of Pal~4 in its early years. | Globular clusters are ideal astrophysical systems whose long-term evolution is determined by several internal and external processes, like mass loss due to stellar evolution and the energy-equipartition processes as well as tidal removal of star. In this regard, numerous numerical investigations have been carried out to understand their dynamical evolution (e.g.~\citealt{Giersz11}; see also the textbook by Heggie \& Hut 2003). However, only within the last few years, with the introduction of graphics processing unit (GPU)-accelerated $N$-body codes such as \textsc{nbody6} \citep{Aarseth03, Nitadori12} it has become feasible to compute the dynamical evolution of massive star clusters over their entire lifetimes on a star-by-star basis. In Paper I of this series \citep{Zonoozi11}, we presented the first direct $N$-body simulation of a Milky Way globular cluster over a Hubble time. For this project we chose the outer-halo globular cluster Palomar~14 (Pal~14), due to its relatively low mass and its large half-mass radius. In the paper, we presented a comprehensive set of $N$-body computations of Pal~14's evolution over its entire lifetime and compared the results to the observed mass, half-light radius, flattened stellar mass function and velocity dispersion of Pal~14, which have been presented by Jordi et al. (2009). We showed that dynamical mass segregation alone cannot explain the mass function flattening in the cluster centre when starting from a canonical Kroupa initial mass function (IMF), and that a very high degree of primordial mass segregation would be necessary to explain this discrepancy. We concluded that such initial conditions for Pal~14 might be obtained by a violent early gas-expulsion phase from an embedded cluster born with mass segregation and a canonical IMF for low-mass stars, a thesis supported later by an independent study of an ensemble of globular clusters \citep{Marks08}. Here we aim at modelling the globular cluster Palomar~4 (Pal~4), which is similar to Pal~14 but has more complete observational data. Recently, Frank et al. (2012) presented an extensive observational study of Pal~4, revealing a flattened stellar mass function and significant mass segregation throughout the cluster. This additional knowledge of Pal~4 puts much stronger constraints on its current dynamical state. Star clusters can undergo significant changes not only at birth but also during the course of their dynamical evolutions. It is therefore essential to specify to what extent the present-day properties of a globular cluster, e.g.~their degree of mass segregation, are imprinted by early evolution and the formation processes, and to what extent they are the outcome of long-term dynamical evolution. There are certain distinct mechanisms that can cause mass segregation. Dynamical mass segregation is the process by which the more massive stars of a gravitationally bound system sink towards the central regions, while lighter stars move further away from the centre. This process is a consequence of evolution towards energy equipartition driven by two-body encounters and is usually associated with the long-term evolution of clusters through the two-body relaxation process. However, a number of observational studies (e.g.~\citealt{Hillenbrand97, Fischer98, Hillenbrand98, de Grijs02, Sirianni02, Gouliermis04, Stolte06, Sabbi08, Allison09, Gouliermis09}) have found evidence of mass segregation in clusters with ages shorter than the time needed to produce the observed segregation via two-body relaxation (see also~\citealt{de Grijs10}). It has been suggested that the observed mass segregation in young clusters could be primordial -- imprinted by the early star-formation process \citep{Bonnell97, Bonnell01, Bonnell98, Klessen01, Bonnell06}. Such mass segregation could be due to the higher accretion rate of proto-stars in high-density regions of molecular clouds fragmenting into clumps. If individual clumps are mass segregated, it has been shown by McMillan, Vesperini, \& Portegies Zwart (2007), that such primordial mass segregation would not be erased in the violent-relaxation phase during which clumps merge. The final system would preserve the mass segregation of the original clumps (see also \citealt{Fellhauer09, Moeckel09}). But even if such clumps are not initially segregated, the internal segregation time-scale can be shorter than the time needed for the clumps to merge. Hence, they will segregate internally via two-body relaxation and preserve this segregation after they have merged (\citealt{McMillan07}). Alternatively, Bastian et al. (2008) found observational evidence for a strong expansion in the first 20 Myr of the evolution of six young M51 clusters and pointed out that this expansion could also lead to a rapid variation in the cluster relaxation time, thus, using the present-day relaxation time might lead to an underestimation of the possible role played by two-body relaxation in generating mass segregation in the early phases of a cluster's dynamical evolution. Regardless of the mechanism producing mass segregation, the presence of primordial (or early) mass segregation significantly affects the global dynamical evolution of star clusters (\citealt{Gurkan04, Baumgardt08a}). For example, the early mass loss due to stellar evolution of high-mass stars has a stronger impact on initially segregated clusters than on unsegregated clusters (\citealt{Vesperini09b}). The degree of primordial or early mass segregation is therefore a crucial parameter in the modelling of globular clusters. Another important quantity that has to be taken into account in the modelling of star clusters is the IMF. Its shape strongly influences the dynamical evolution of star clusters. The canonical IMF as observed in young star clusters in the Milky Way is often expressed as a two-part power-law function ($\frac{dN}{dm}\propto m^{-\alpha}$) with near Salpeter-like slope above $0.5\,\mbox{M}{\odot}$ (i.e., $\alpha=2.3$; \citealt{Salpeter55}), and a shallower slope of $\alpha=1.3$ for stars in the mass range $0.08-0.5\,\mbox{M}{\odot}$ \citep{Kroupa01, Kroupa08, Kroupa13}. The mass function of stars in clusters evolves through stellar evolution and through dynamical evolution, i.e. via preferential loss of low-mass stars \citep{Vesperini97,Baumgardt03}. This effect should be more pronounced in concentrated clusters, since the two-body relaxation times-cale is shorter for such systems. However, based on a data set of observed mass functions of a sample of globular clusters, De Marchi et al. (2007) found that all high concentration clusters have steep mass functions (i.e., larger $\alpha$), while low concentration ones have a smaller $\alpha$, although the opposite is expected. This effect is not well understood yet. Marks et al. (2008) suggested that the `De Marchi relation' is due to early gas expulsion. They showed that for initially mass-segregated clusters mostly low-mass stars are lost due to gas expulsion, which yields a shallower slope in the low-mass range in clusters with low concentration. Moreover, the mass functions of some outer-halo globular clusters also show a flattening at comparatively high stellar masses (i.e., the range $0.55\leq m/M{\odot}\leq0.85$; \citealt{Jordi09, Frank12}). Some of the ideas that have been proposed to explain a shallowness of the slope at the high-mass end, again, include primordial mass segregation of stars in the cluster (e.g., \citealt{Vesperini97, Kroupa02, Mouri02}). It remains to be shown if such scenarios can really reproduce the observational findings. Direct $N$-body simulations offer the possibility to test these scenarios. In this paper we perform a set of direct $N$-body simulations of Pal~4 to determine its most likely initial conditions in terms of total mass, initial half-mass radius, stellar mass function and primordial mass segregation. We furthermore investigate the effect of unresolved binaries on the observed mass function of this cluster. The paper is organized as follows. In Section~\ref{Sec:Observational data} we describe the observational data of Pal~4, including the velocity dispersion, the mass function and the total stellar mass to which we later compare our simulations. In Section~\ref{Sec:Description of the models} we describe the set-up of the $N$-body models. This is followed by a presentation of the results of simulations for different scenarios in Section~\ref{Sec:Results}. A discussion and conclusions are presented in Section~\ref{Sec:Conclusions}. | \label{Sec:Conclusions} This paper is the second study in which we model the dynamical evolution of a Galactic globular cluster over its entire lifetime by direct $N$-body simulations on a star-by star basis. While we focused on Pal~14 in Paper I \citep{Zonoozi11}, we here investigate the diffuse outer halo globular cluster Pal~4 using the $N$-body code \textsc{nbody6} \citep{Aarseth03}. Recent observational work on Pal 4 \citep{Frank12} has shown that the global mass function slope in the mass range 0.55-0.85 M${\odot}$ is $\alpha=1.4\pm0.25$, i.e.~significantly shallower than a canonical mass function slope of about 2.3 \citep{Kroupa01}. Similar results have been found for a number of Milky Way globular clusters (see, e.g., \citealt{De Marchi07,Jordi09,Paust10,Frank12,Hamren13}). Interestingly, \cite{Frank12} also found that the slope of the mass function steepens with radius from a slope of $\alpha \leq 1$ inside about $1.3 r_h$ to $\alpha \geq 2.3$ at the largest observed radii, indicating the presence of mass segregation in Pal~4 and therefore constraining numerical models much more than our previous target Pal~14 could do. A preferential loss of low-mass stars due to two-body relaxation would be a natural explanation for the observed mass function depletion \citep{Baumgardt03}. However, for diffuse outer halo clusters such as Pal~4 and Pal~14 (i.e., a low mass together with a large half-mass radius), the present-day two-body relaxation time is of the order of a Hubble time. Therefore, relaxation should be inefficient in these clusters and the observations should be an indication for primordial mass segregation. Alternatively the cluster could have been more compact in the past such that relaxation was more important at that time. To test these scenarios, we have tried to find the best possible evolutionary model for Pal~4 by running a set of models with varying initial half-mass radii and total masses, until we got an adequate fit to the observed structural parameters. While it is relatively straightforward to find initial models which reproduce the observed structural parameters of Pal~4, i.e. half-light radius, total mass and velocity dispersion, it is very difficult if not impossible to reproduce its global mass function and degree of mass segregation. Because the models have to start with a comparatively low mass and large half-mass radius of about 55000 M${\odot}$ and 10 pc, low-mass star depletion and mass segregation are very ineffective in these clusters. We showed that models evolving on circular orbits, starting with a Kroupa IMF, and without primordial mass segregation do not produce enough depletion in the slope of the mass function. In addition, these models do not develop enough mass segregation within the cluster lifetime to match the observations. It should be noted that the current conclusions are based on the assumption of a circular orbit for Pal~4. The orbit of Pal 4 is unknown however. In the case of an eccentric orbit the Galactic field changes with time, which could significantly affect the dynamical evolution of Pal 4. We also find that the present-day global mass function slope of Pal 4 cannot be reproduced in models starting with a canonical but primordially segregated IMF, not even by using very high degrees of primordial segregation. Models starting with a flattened IMF reach enough depletion in the global mass function to be compatible with the observations. However, the radial variation of the mass function slope is significantly better reproduced when we include both a flattened IMF \textit{and} primordial mass segregation. This is similar to our findings from Paper I \citep{Zonoozi11}, where we concluded that Pal~14 must have undergone one of two scenarios: \begin{enumerate} \item the observed mass function may be a result of dynamical evolution starting from a canonical Kroupa IMF with a high degree of primordial mass segregation;\\ \item the observed mass function may be the result of an already established non-canonical IMF depleted in low-mass stars, which might have been obtained during a violent early phase of gas-expulsion of an embedded cluster with primordial mass segregation \citep{Marks10}. \end{enumerate} Now, for Pal~4 we can exclude the first scenario as we have observations covering larger parts of the cluster and hence have more precise knowledge of the present-day global mass function and the degree of mass segregation. This leaves us with the assumption that the peculiar mass function and the cluster's unusual extent have been imprinted on Pal~4 during its very early lifetime. The inferred initial half-light radius of about 10 pc is significantly larger than the present-day half-light radii of most globular clusters, which are narrowly distributed around 3 pc \citep{Jordan05}. This could be a footprint of the weaker external tidal force of Pal~4's host galaxy during its formation. The Galactic tidal field, which we here model as being static, has evolved significantly since Pal~4's birth and might have been much weaker 11 Gyr ago. The cluster might have also been born in a now detached/disrupted dwarf galaxy. Alternatively, since star clusters lose more mass during pericentric passages on eccentric orbits, and undergo stronger expansion due to the weaker tidal fields at larger Galactic radii \citep{Madrid12}, an eccentric cluster orbit might have had an important influence on Pal~4's evolution, as it could have had a much smaller initial size and significantly higher mass. We leave this scenario for an upcoming paper to be investigated in detail (K\"upper et al., in preparation). | 14 | 4 | 1404.5969 |
1404 | 1404.1536_arXiv.txt | \setcounter{equation}{0} Our understanding on the origin of the Universe has advanced considerably in recent years through interactions between experiments and theories. We have a large number and variety of ongoing and upcoming experiments that are mapping the entire observable universe. One of the most important achievements of these experiments is to produce, one way or the other, different and complimentary maps of the distributions of large scale structures, including various spectra and objects, in our Universe. These maps are the gold mines to advance our knowledge in cosmology. All these large scale structures originated from some tiny fluctuations in the very early universe, the primordial perturbations. One of the most beautiful ideas in modern cosmology is that these perturbations are seeded by quantum fluctuations of fields present in an early epoch responsible for the Big Bang. By studying properties of these maps, we learn properties of this epoch, as well as fundamental physics in conditions that are inaccessible for experiments on Earth. In the past two decades, the data from CMB and Large Scale Structures (LSS) strongly support the inflationary paradigm \cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi,Starobinsky:1980te,Sato:1980yn} as the leading candidate for this primordial epoch. The simplest inflationary models predict the primordial perturbations to be superhorizon, approximately scale-invariant, adiabatic and Gaussian \cite{Mukhanov:1981xt,Hawking:1982cz,Starobinsky:1982ee,Guth:1982ec,Bardeen:1983qw}. All of these have been verified to some extent by the results from the Wilkinson Microwave Anisotropy Probe (WMAP) \cite{Komatsu:2010fb} and the Planck satellite \cite{Ade:2013zuv,Ade:2013uln}. The properties of these perturbations are summarized quantitatively by two of the six parameters in the Standard Model of Cosmology, the $\Lambda$CDM model. On the other hand, other possibilities have also been speculated as alternative theories to the inflationary scenario. From the perspective of theoretical model building, none of them has been as successful as inflation. See Ref.~\cite{Brandenberger:2012zb,Ijjas:2013vea,Guth:2013sya,Linde:2014nna} for the current status. Nonetheless, models may be improved or become complicated to fit the data. This is possible because there are only two parameters in the Standard Model that are relevant to the primordial epoch, leaving rooms for theoretical freedoms. Therefore an equally important approach in cosmology is to search for beyond-Standard-Model signals in data that can be used to distinguish different scenarios. Phenomenologically one can distinguish four different kinds of primordial epochs, classified by the time dependence of the scale factor $a(t)\sim t^p$: the fast-expanding or fast-contracting scenarios, and the slowly-expanding or slowly-contracting scenarios. (The contracting scenarios require a bounce to match the Big Bang.) Each of them has a different fingerprint index in terms of the parameter $p$ \cite{Chen:2011zf,Chen:2011tu}. The acceleratedly-expanding scenario, namely inflation, has $|p|>1$; the fast-contracting scenario \cite{Wands:1998yp,Finelli:2001sr} has $p\sim \CO(1)<1$; the slowly-expanding scenario has $-1 \ll p <0$; and the slowly-contracting scenario \cite{Khoury:2001wf} has $0<p\ll 1$.\footnote{The case $p\sim \CO(-1) > -1$ ($-\infty<t<0$) is also acceleratedly expanding.} For $p>1$, $t$ runs from $0$ to $+\infty$; for all other $p$, $t$ runs from $-\infty$ to $0$. The choices of $t$ are based on the requirement that the quantum fluctuations in this epoch exit the horizon, so that they can give rise to the acoustic oscillations in the CMB after reentry during the Big Bang. The primordial perturbations, which are seeded by quantum fluctuations in these epochs, consist of scalar and tensor modes. While the scalar mode determines the density perturbations at the beginning of the Big Bang as the source of the large scale structures, the tensor mode corresponds to the gravitational quantum fluctuations and records the magnitude of the Hubble parameter during the epoch. Therefore the tensor mode serves as a good discriminator between the scenarios with fast-evolving scale factor and those with slowly-evolving scale factor. In particular, if the tensor mode origin of the recent CMB B-mode detection by the BICEP2 experiment \cite{Ade:2014xna} is confirmed, both the slowly-expanding and slowly-contracting scenarios will be ruled out. Nonetheless, phenomenologically, the tensor mode does not distinguish the inflation from the fast-contracting scenarios. For example, both the inflation and the matter contraction can give rise to observable tensor mode with approximately scale-invariant spectra \cite{Wands:1998yp,Finelli:2001sr}. In this letter, we consider a different type of observables. A main reason the degeneracy of scenarios could exist is that the observables we mentioned so far (namely the approximately scale-invariant scalar and tensor modes) are all convoluted consequences of the scale factor $a(t)$, the defining property of different scenarios. A direct measurement of $a(t)$ would provide an independent and direct evidence for a scenario, as was done for the late-time accelerated universe using the Standard Candles \cite{Perlmutter:1997zf,Riess:1998cb}. This turns out to be possible: oscillating massive fields in the primordial epoch can serve for this purpose as the Standard Clocks \cite{Chen:2011zf,Chen:2011tu,Chen:2012ja}. The massive field oscillates with a frequency that can be thought of as ticks of a clock. This Standard Clock imprints its ticks as a special type of features in the primordial perturbations, thereby letting some imprints in the CMB angular power spectra, the non-Gaussianities and the distribution of large scale structures. The patterns of these ticks are a direct record of $a(t)$ of the primordial universe. In this Letter, after summarising the main results of the theoretical proposal of the Standard Clock, we compare its key predictions with the Planck 2013 residual data. A full-scale comparison will be the subject of the next paper \cite{ToAppear}. Here we focus on one interesting candidate emerging from this comparison, although it is still not statistically significant. Motivated by this candidate we construct explicit Standard Clock models and compute the full power spectrum. This is a completion of the above key predictions, under the same number of model parameters. We again see encouraging signs after this prediction is compared with the Planck data. | Following the theoretical proposal of Standard Clock, we have started to compare its predictions with the Planck data. We have constructed and worked out a Standard Clock model with full predictions on the power spectrum. We have presented an interesting candidate in the Planck data, characteristic of the Standard Clock in the inflationary scenario. Although this candidate is not yet statistically significant, we use this to motivate detailed theoretical model-building, and use it as an example to illustrate how Standard Clock appears in CMB and how they can be further tested by future data. Such a Standard Clock candidate is an extensive feature covering nearly the entire range of scales probed by Planck. It would potentially add at least four more parameters to the $\Lambda$CDM Standard Model. It has highly correlated predictions in the polarization map, non-Gaussianities and other large scale structure maps over the same wide range of scales, and so can be tested with further analyses and future data. The Standard Clock signal contains important information on the early universe, and in particular can be used to directly measure the time-dependence of the scale factor of the primordial universe. If any of these candidates is verified, such a signal would provide an independent and direct evidence for the inflationary paradigm. \begin{figure}[H] \centering \includegraphics[width=0.85\textwidth]{numerical_smallfield.PDF} \includegraphics[width=0.90\textwidth]{FullSC_smallfield_TT.PDF} \caption{A Standard Clock full numerical prediction using a small field example ({\it top}) and its comparison with the Planck residuals using the template (\ref{template}) ($k_r=0.106~{\rm Mpc}^{-1}$, $A=0.07$) and CAMB \cite{Lewis:1999bs} ({\it bottom}).} \label{Fig:FullModel_smallfield} \end{figure} \begin{figure}[H] \centering \includegraphics[width=0.90\textwidth]{FullSC_smallfield_TE.PDF} \caption{Predictions on EE and TE residuals correlated with Fig.~\ref{Fig:FullModel_smallfield}.} \label{Fig:smallfield_TE} \end{figure} \medskip | 14 | 4 | 1404.1536 |
|
1404 | 1404.2960.txt | {Small planets, 1--4x the size of Earth, are extremely common around Sun-like stars, and surprisingly so, as they are missing in our solar system. Recent detections have yielded enough information about this class of exoplanets to begin characterizing their occurrence rates, orbits, masses, densities, and internal structures. The {\em Kepler} mission finds the smallest planets to be most common, as 26\% of Sun-like stars have small, 1-2 \rearth planets with orbital periods under 100 days, and 11\% have 1--2 \rearth planets that receive 1-4x the incident stellar flux that warms our Earth. These Earth-size planets are sprinkled uniformly with orbital distance (logarithmically) out to 0.4 AU, and probably beyond. Mass measurements for 33 transiting planets of 1--4 \rearth show that the smallest of them, $R < 1.5$ \rearthe, have the density expected for rocky planets. Their {\it densities increase with increasing radius}, likely caused by gravitational compression. Including solar system planets yields a relation: $\rho = 2.32 + 3.19 R/R_{\oplus}$ [\gcc]. Larger planets, in the radius range 1.5--4.0 \rearthe, have {\it densities that decline} with increasing radius, revealing increasing amounts of low-density material (H and He or ices) in an envelope surrounding a rocky core, befitting the appellation ``mini-Neptunes.'' Planets of $\sim$ 1.5 \rearth have the highest densities, averaging near 10 \gcc. The gas giant planets occur preferentially around stars that are rich in heavy elements, while rocky planets occur around stars having a range of heavy element abundances. One explanation is that the fast formation of rocky cores in protoplanetary disks enriched in heavy elements permits the gravitational accumulation of gas before it vanishes, forming giant planets. But models of the formation of 1--4 \rearth planets remain uncertain. Defining habitable zones remains difficult, without benefit of either detections of life elsewhere or an understanding of life's biochemical origins.} | Among the nearly 4000 planets known around other stars, the most common are 1--4x the size of Earth. A quarter of Sun-like stars have such planets orbiting within half an Earth's orbital distance of them, and more surely orbit farther out. Measurements of density show that the smallest planets are mostly rocky while the bigger ones have rocky cores fluffed out with hydrogen and helium gas, and likely water, befitting the term ``mini-Neptunes.'' The division between these two regimes is near 1.5 \rearthe. Considering exoplanet hospitality, 11\% of Sun-like stars have a planet of 1--2x the size of Earth that receives between 1.0--4.0x the incident stellar light that our Earth enjoys. However, we remain ignorant of the origins of, and existence of, exobiology, leaving the location of the habitable zone uncertain. \bigskip %% The first letter of the article should be drop cap: \dropcap{} %\dropcap{I}n this article we study the evolution of ''almost-sharp'' fronts %% Enter the text of your article beginning here and ending before %% \begin{acknowledgements} %% Section head commands for your reference: %% | 14 | 4 | 1404.2960 |
|
1404 | 1404.7126_arXiv.txt | We present an analysis of five Suzaku observations of the bright Seyfert1 galaxy IC 4329A. The broad energy band and high signal-to-noise ratio of the data give new constraints on the iron K$\alpha$ line profile and its relationship with the Compton hump at higher energies. The Fe K bandpass is dominated by a narrow core (EW=57$_{-3}^{+3}$~eV) at 6.4 keV consistent with neutral material. Using a physically-motivated model, our analysis also reveals the presence of a broad Iron K$\alpha$ line (EW=124$_{-11}^{+11}$~eV), most likely produced in the inner part of the accretion disk and blurred by general relativistic effects. This component is not immediately evident from the individual spectra, but is clearly present in the stacked residuals of all five observations, and has high statistical significance. This highlights the difficulty in identifying broad iron lines in AGN, even in data with very high signal-to-noise ratio, as they are difficult to disentangle from the continuum. The data are consistent with the narrow and broad iron line components tracking the Compton Hump, but do not provide clear evidence that this is the case. An additional narrow Fe~{\sc xxvi} emission line at 6.94 keV is also seen, suggesting the presence of ionized material relatively distant from the central region. There is also a hint of variability, so the precise origin of this line remains unclear. | The X-ray spectra of the Active Galactic Nuclei (AGN) are dominated by a simple power law with a photon index of $\Gamma\sim1.9$ (e.g. \citealt{nandra+94}, \citealt{Piconcelli+05}). This emission is thought to originate from inverse Comptonization of soft photons by a hot corona near the central black hole (e.g., \citealt{Shapiro+76}; \citealt{Sunyaev+80}; \citealt{Haardt+93}; \citealt{Haardt+94}). The optically thick and cold disk around the black hole (BH) reprocesses the X-rays producing the so-called Compton reflection Hump peaking near 20-30 keV (e.g., \citealt{Pounds+90}; \citealt{George+91}). Other signatures of this reflection are the absorption and emission lines produced by photoelectric absorption and fluorescence (\citealt{Matt+97}). The most prominent is the Fe K$\alpha$ emission line seen at 6.4 keV (\citealt{nandra+94}). Observations with high spectral resolution, performed with \textit{XMM-Newton} and \textit{Chandra}, have revealed that a relatively narrow core to the Fe K$\alpha$ line is common in type 1 AGN, with FWHM's of several thousand km s$^{-1}$ (e.g., \citealt{Yaqoob+04}; \citealt{Nandra+06}). The narrow component is thought to be produced in Compton-thick material distant from the black hole, such as the molecular torus (e.g. \citealt{Krolik+87}). The reflection can also arise in the inner parts of the accretion disk, producing a relativistically broadened component (\citealt{Tanaka+95}), modified by gravitational redshift and relativistic Doppler effects (e.g., \citealt{Fabian+89}; \citealt{Fabian+02}). The analysis of the Iron K$\alpha$ line components is a key probe of the innermost region of the AGN (\citealt{Fabian+09}). Broad iron lines are expected, and observed, to be a widespread feature in bright AGN, with the observed fraction between $\sim30\%$ and $\sim80\%$ among nearby Seyfert galaxies (\citealt{Nandra+07}; \citealt{Calle+10}). Key outstanding issues are why some AGN apparently lack a disk line component (\citealt{Bhayani+11}) and whether or not the Fe K$\alpha$ emission is correlated with the associated hard X-ray reflection continuum, as it should if the line arises from optically thick material. IC 4329A ($z=0.01605$; \citealt{Willmer+91}) is one of the brightest Seyfert 1 AGN ($F\sim2\times10^{-10}$ erg s$^{-1}$ cm$^{-2}$), embedded in a nearly edge-on host galaxy. The nature of the iron K-complex in this source, and in particular the presence or otherwise of broad component to the iron line, has been controversial. Analyzing simultaneous \textit{ASCA} and \textit{RXTE} observations, \citet{Done+00} detected a moderately broadened (FWHM=43,000$\pm$11,000 km s$^{-1}$) Fe K$\alpha$ line with an equivalent width of EW=180$\pm$50 keV peaking at $\sim$6.4 keV. While the line is significantly broadened, it is not as expected from an accretion disk that extends down to the last stable orbit around a black hole. Similar results had previously been found based on \textit{ASCA} data by \citet{Nandra+97}. An \textit{XMM-Newton} observation reported by \citet{Gondoin+01} revealed a narrow core for the iron line (EW=43$\pm$1 eV) originating in mostly neutral material. \citet{Nandra+07} analyzed two XMM-Newton observations, finding one with purely narrow emission and another moderately broadened like the \textit{ASCA} case. A higher resolution view of the iron line complex in IC 4329A was provided using a 60 ks \textit{Chandra}-HETGS observation by \citet{McKernan+04}, who detected a narrow core for the 6.4 keV line together with an additional emission line near 6.9 keV. This double-peaked feature could be reproduced by several different models, included dual Gaussians, dual disk lines or a single disk line. Finally we note that \citet{Markowitz+06}, \citet{Nandra+07} and \citet{Tombesi+10} have all reported the possible presence of a blue shifted absorption line in an XMM-Newton observation of IC 4329A, which may be the signature of a very fast outflow. The aim of this paper is to study the nature of the iron line and reflection component in this Seyfert galaxy using data from the \textit{Suzaku} satellite. With its large effective area and broad bandpass this provides an ideal opportunity to determine the nature of the iron K$\alpha$ emission and its relationship to the Compton Hump above 10 keV. | We have analyzed 5 Suzaku observations in order to determine the nature of the Fe K$\alpha$ emission and its relation to the Compton Hump in the Seyfert 1 galaxy IC 4329A. The analysis of the iron line shows a core with an energy that is consistent with an origin in neutral reflection form distant material. Once this narrow component of the line is fitted, broad residuals in the Fe K band appear, but these are only clearly evident when all datasets are combined (Fig. \ref{1gaussXIS}): they are not clearly visible in the single data/model ratio of each observation (Fig. \ref{power}). These broad residuals are well reproduced by a model appropriate to an origin in the inner accretion disk, where the line and the continuum are blurred by relativistic effects. This is the first clear demonstration of a relativistic reflection component in this source, with previous reports suggesting either a narrow line only, or moderate broadening, the latter having been interpreted as evidence for a truncated accretion disk. The apparent lack of relativistic signatures are in the X-ray spectra of some nearby AGN -- including IC 4329A -- is a major outstanding issue for the standard accretion disc paradigm for feeding of AGN. If a relatively cool, optically thick accretion disk is present around the black hole and X-ray emission impinges on the disk, this component is hard to avoid. Disk truncation to large radii seems implausible because of the very high radiative efficiency of these systems, and the fact that most of the energy dissipation is expected to occur in the central regions. \citet{Bhayani+11} have suggested in general, and specifically for this source, that strong relativistic effects, disk ionization and/or high disk inclination can explain the apparent lack of relativistic signatures due to the difficulty of disentangling very broad features from the continuum. While these effects may be operating in IC 4329A the spectra are well fit by a relatively simple, normal relativistic, reflection component from a neutral disk in a Schwarzschild geometry seen at a relatively low inclination ($35\deg$). The difficulty in recognizing the relativistic component in IC 4329A stems from its relative weakness compared to the expectations for a semi-infinite flat disk geometry illuminated by a point source ($R_{B}\sim 0.3$ compared to $R\sim 1$). It has been noted, however, that the latter geometry is that which produces the very strongest reflection (e.g. \citealt{Murphy+09}), and any significant deviation from this results in lower reflection fractions. General relativistic effects close to the black hole can also result in reflection that is either stronger, or weaker, than this expectation (\citealt{Miniutti+03}). Either or both effects could be at play in IC 4329A. Our work supports the contention of \citet{Calle+10} and \citet{Nandra+07} that very high signal-to-noise ratio is required to disentangle broad iron line components even in nearby AGN. Despite IC 4329A being the second brightest Seyfert galaxy in the 2-10 keV bands, it was necessary to sum all the data/model ratios and reach a total exposure time of $\sim$130 ks in order to clearly detect the important residual at the energy of the Iron line. Another aspect of our current work is the examination of the relationship between the iron line strength and the Compton Hump. Being features of the same reflected spectrum, they should vary together from spectrum-to-spectrum. Fits with the phenomenological Gaussian line model have revealed no clear relationship between these components in our fits. On the other hand, the physically motivated reflection model, in which the line is forced to follow the reflection continuum, provides a much better fit to the data, suggesting that such a correlation might in fact hold. As a further test, we allowed the iron abundance to be free in the reflection model. While largely unphysical (one would not expect the iron abundance to vary on these short timescales), the effect of this is to decouple the emission line from the reflection continuum, while retaining the requirement for both to be present, and accounting for the fact that there is evidence in the spectra for both a blurred and distant reflection component. Applying such a model we find consistency in the iron abundance for all spectra, with tentative evidence that it is sub-solar ($A_{Fe}\sim 0.5$). If this is the case, a low iron abundance could be another factor contributing to the difficulty of detecting the broad emission line in this source. Regardless, the consistency of the iron abundance between the five observations shows that the spectra are consistent with the iron line tracking the reflection continuum, as expected in the standard model. We also detected a narrow emission at 6.94 keV the 99.7$\%$ significance in two observations, in agreement with the work of \citet{McKernan+04}. This feature does not appear to be the blue wing of a disk line, so is more likely to be associated with highly ionized gas, and identified with the Fe~{\sc xxvi} emission line, albeit with a small redshift. | 14 | 4 | 1404.7126 |
1404 | 1404.5315_arXiv.txt | Using the deepest available \textit{Chandra} observations of NGC 4649 we find strong evidences of cavities, ripples and ring like structures in the hot interstellar medium (ISM) that appear to be morphologically related with the central radio emission. {These structures show no significant temperature variations in correspondence with higher pressure regions (\(0.5\mbox{ kpc}<r<3\mbox{ kpc}\)).} On the same spatial scale, a discrepancy between the mass profiles obtained from stellar dynamic and \textit{Chandra} data represents the telltale evidence of a significant non-thermal pressure component in this hot gas, which is related to the radio jet and lobes. On larger scale we find agreement between the mass profile obtained form \textit{Chandra} data and planetary nebulae and globular cluster dynamics. The nucleus of NGC 4649 appears to be extremely radiatively inefficient, with highly sub-Bondi accretion flow. Consistently with this finding, the jet power evaluated from the observed X-ray cavities implies that a small fraction of the accretion power calculated for the Bondi mass accretion rate emerges as kinetic energy. Comparing the jet power to radio and nuclear X-ray luminosity the observed cavities show similar behavior to those of other giant elliptical galaxies. | Evidence of the interaction of Active Galactic Nuclei (AGN) with the surrounding hot gas in nearby galaxies and clusters has been observed as morphological disturbances in the X-ray halos in the form of ripples and cavities \citep[e.g.][]{2000MNRAS.318L..65F,2003MNRAS.344L..43F,2005ApJ...635..894F,2006MNRAS.366..417F}. The AGN-induced disturbances have also been observed in the hot interstellar medium (ISM) in the halos of a number of normal elliptical galaxies \citep[e.g.][]{2007ApJ...668..150D}, and are interpreted as a consequence of the thermal X-ray emitting gas being displaced by the AGN jets. NGC 4649, also known as M60, is a nearby\footnote{We assume a distance to NGC 4649 of (\(16\mbox{ Mpc}\)). At this distance \(1''\) corresponds to \(77\mbox{ pc}\).} X-ray-bright giant elliptical galaxy located in a group at the eastern edge of the Virgo cluster. Its companion, located at \(\sim 2.5'\) to the northwest, is the spiral galaxy NGC 4647. NGC 4649 harbors a faint nuclear radio source \citep{2002AJ....124..675C}. Although earlier \textit{Chandra} data revealed a relaxed X-ray morphology close to hydrostatic equilibrium - first reported by \citet{1995ApJ...452..522B} - there has been a debate on the presence of inhomogeneities correlated with the nuclear radio source. {Finger-like structures in the inner \(\sim 5\) kpc of the diffuse X-ray emission from NGC 4649 have been reported by \citet{2004ApJ...600..729R,2006ApJ...636..200R} in their study of \textit{Chandra} and \textit{XMM-Newton} data. These structures are compared by the authors with those predicted by hydrodynamical simulations of cooling flows in elliptical galaxies \citep{1998MNRAS.301..343K}, that is, brighter, cooler inflowing gas surrounded by fainter, hotter outflowing jets. However, the authors found no significant temperature variations across the observed structures.} \citet{2008MNRAS.383..923S}, using {the same \(\sim 37\mbox{ ksec}\) \textit{Chandra} observation of \citet{2004ApJ...600..729R} analysis}, found morphological disturbances in the X-ray emitting gas, and interpreted them as the result of interaction with the central AGN. Instead a subsequent analysis of deeper \textit{Chandra} observations by \citet{2008ApJ...683..161H} showed a generally undisturbed X-ray morphology, consistent with that expected from a hot ISM in hydrostatic equilibrium. Later, the analysis by \citet{2010MNRAS.404..180D} of \textit{Chandra} observations shallower that those of \citeauthor{2008ApJ...683..161H} - but deeper than \citeauthor{2008MNRAS.383..923S} - revealed (again) disturbances and cavities in the ISM connected with the radio emission. Thanks to the relatively small distance and large supermassive black hole (SMBH) mass (\(\sim\) few \({10}^9\,M_{\astrosun}\)) of NGC 4649, \textit{Chandra} resolves radii close to the Bondi accretion radius, \(r_{acc}\approx 100-200\mbox{ pc}\), at which the gravitational binding energy of a gas element becomes larger than its thermal energy \citep{1952MNRAS.112..195B}. Thus, NGC 4649 represents an ideal case to investigate the following questions: what is the mass accretion rate? What fraction of the accretion power is in the observed nuclear luminosity, and what in the observed jet power? Is NGC 4649 consistent with the previously found correlations between the Bondi mass accretion rate \(\dot M_{B}\), the power associated with the observed cavities \(P_{cav}\), and radio luminosity \citep{2010ApJ...720.1066C,2013MNRAS.432..530R}? Due to its low radio power, and its non-being a central dominant galaxy, NGC 4649 is also an ideal case to investigate whether these correlations, mostly found for radio-bright central dominant galaxies in groups or clusters, work equally well in more ``normal'' elliptical galaxies. In this paper we revisit the properties of the hot ISM of NGC 4649, using much deeper \textit{Chandra} data with respect to previous studies and updated atomic databases. We find strong evidence of cavities, ripples and ring like structures that appear to be morphologically related with the central radio emission. In addition, we find that the hot halo is subject to an additional non-thermal pressure term as already reported in previous studies \citep[e.g.][]{2009ApJ...705.1672B,2010MNRAS.409.1362D,2011MNRAS.415.1244D,2013MNRAS.430.1516H}. We show that the non-thermal pressure is found on the same scale spatial as the disturbances of the halo and it is spatially correlated with the hot gas pressure and the minimum pressure derived from the radio data. The nucleus of NGC 4649 appears to be extremely radiatively inefficient, with highly sub-Bondi accretion flow, releasing a very small fraction of the accretion power in form of kinetic energy in the surrounding halo. The paper is organized as follows: Section \ref{sec:data} describes the data sets used in this work, the reduction procedures, and the image and spectral analysis. In Section \ref{sec:discussion} we discuss our results, and Section \ref{sec:summary} is dedicated to our conclusions. | \label{sec:summary} We investigated the presence of AGN feedback in the ISM of the giant elliptical NGC 4649 by using a total of 280 ks \textit{Chandra} observations This source has been studied several times in different wavelength, and in particular in the X-rays \citep[e.g.][]{2010MNRAS.404.1165C,2010MNRAS.409.1362D,2012ApJ...757..121L}. \citet{2008MNRAS.383..923S} and \citet{2008ApJ...683..161H}, making use of \textit{Chandra} observations, studied the properties of the ISM in NGC 4649 using the unsharp-mask technique looking for morphological disturbances pointing to deviations from the hydrostatic equilibrium condition suggested by the generally relaxed X-ray morphology. Interestingly, while the former authors using a shallow \(\sim 37\mbox{ ksec}\) observation found evidences of structures and cavities in the ISM that they interpreted as connected with the central, faint radio source, the latter authors using deeper \(\sim 81 \mbox{ ksec}\) data did not found any evidence of such disturbances. A subsequent analysis by \citet{2010MNRAS.404..180D} showed disturbances and cavities in the ISM as residuals of the X-ray surface brightness from a spherical \(\beta\) model, connected with the radio emission. Using much deeper \textit{Chandra} data with a total exposure \(\sim 280\mbox{ ksec}\), we used the latter approach to investigate the morphological distribution of the ISM in NGC 4649. We studied the deviation of the X-ray surface brightness from an elliptical \(\beta\) model, which is expected to describe the hot gas distribution in relaxed galaxies. The residuals of this fitting procedure, presented in Figure \ref{fig:sign10}, show significant cavities, ripples and ring like structures on the inner \(0.5\div 3\mbox{ kpc}\) scale. This is at variance with the \(\lesssim 2\sigma\) significance of the cavities reported by \citet{2008MNRAS.383..923S} as evaluated by \citet{2008ApJ...683..161H}. The deeper data considered here revealed these structures with high significance; moreover the structures appear to be morphologically related with the central radio emission, with cavities lying in correspondence with the extended radio lobes and regions of enhanced emission situated on the side of them and, on larger scale, taking the form of ring like ripples which seems reminiscent of the structures observed in NGC 1275 \citep{2006MNRAS.366..417F}. In common with this source, we found no significant temperature variations in correspondence with higher pressure regions. So, if radio ejecta driven shocks are responsible for the observed ISM morphology, the observed structures may be isothermal waves whose energy is dissipated by viscosity, with thermal conduction and sound waves effectively distributing the energy from the radio source. Evidences of deviations from the hydrostatic equilibrium are also provided by the mass profiles presented in Figure \ref{fig:masses} (left panel). A significant non-thermal pressure is found on the same scale of the residual structures, where it reaches \(\sim 30\%\) of the observed gas pressure. In addition, the excess gas pressure and non-thermal pressure profiles appear to be strongly correlated, indicating the radio ejecta as the likely origin for this additional pressure component. At smaller scales, similarly to a few other early type galaxies harboring low power radio sources, NGC 4649 shows increased temperatures in the inner \(0.5\mbox{ kpc}\) region. The nucleus of NGC 4649 appears to be extremely sub-Eddington, with the accretion flow emitting less than predicted from the fuel observed to be available, even when allowing for a RIAF with angular momentum at the outer radius of the accretion flow. Also the jet power evaluated from the observed X-ray cavities appears to be much smaller than that predicted for elliptical galaxies from the Bondi accretion power \(P_B\). If the mass accretion rate accounting for the observed nuclear X-ray luminosity is adopted - which requires, in addition to a low radiative efficiency, a significant reduction of the accretion rate with respect to the Bondi value, due, e.g., to outflows/convective motions - then the corresponding accretion power \(P\) is \(\sim 10\) times larger than the observed kinetic {power}. When comparing the jet power to radio and nuclear X-ray luminosity, on the other hand, the observed cavities show similar behavior to those of other giant elliptical galaxies. | 14 | 4 | 1404.5315 |
1404 | 1404.2159_arXiv.txt | The arrival of solid particles from outside our solar system would present us with an invaluable source of scientific information. Attempts to detect such interstellar particles among the meteors observed in Earth's atmosphere have almost exclusively assumed that those particles moving above the Solar System's escape speed -- particles on orbits hyperbolic with respect to the Sun-- were precisely the extrasolar particles being searched for. Here we show that hyperbolic particles can be generated entirely within the Solar System by gravitational scattering of interplanetary dust and meteoroids by the planets. These particles have necessarily short lifetimes as they quickly escape our star system; nonetheless some may arrive at Earth at speeds comparable to those expected of interstellar meteoroids. Some of these are associated with the encounter of planets with the debris streams of individual comets: Comet C/1995 O1 Hale-Bopp's 1996 pre-perihelion encounter with Jupiter could have scattered particles that would have reached our planet with velocities of almost 1 km~s$^{-1}$ above the hyperbolic velocity at Earth; however, such encounters are relatively rare. The rates of occurrence of hyperbolically-scattered sporadic meteors are also quite low. Only one of every $\sim10^{4}$ optical meteors observed at Earth is expected to be such a locally generated hyperbolic and its heliocentric velocity is typically only a hundred meters per second above the heliocentric escape velocity at Earth's orbit. The majority of such gravitationally-scattered hyperbolics originate at Mercury, though Venus and Mars also contribute. Mercury and Venus are predicted to generate weak 'hyperbolic meteor showers': the restrictive geometry of scattering to our planet means that a radiant near the Sun from which hyperbolic meteors arrive at Earth should recur with the planet's synodic period. However, though planetary scattering can produce meteoroids with speeds comparable to interstellar meteors and at fluxes near current upper limits for such events, the majority of this locally-generated component of hyperbolic meteoroids is just above the heliocentric escape velocity and should be easily distinguishable from true interstellar meteoroids. | The first measurement of a meteor velocity may have been due to \cite{elk00}. He used a bicycle wheel as the basis for a rotating shutter that would interrupt the meteor's image on a photograph, and the segmented image was used to determine the meteor's velocity. The idea of using photographs to measure meteor velocity goes back further, at least to \cite{lan60} but was not initially widely-used. The difficulty with photographic observations was its limited sensitivity in its early days: a hundred hours of observation might be required for a single successful result \eg \cite{lov54}. Though photography goes back to the early 1800's, as late as 1932 \cite{shaopiboo32} noted that ``several hundred meteors are visible to the unaided eye to one that can be photographed''. Naked-eye visual observations provided the first substantial number of meteor velocity measurements, along with the first indication of meteors that might be from outside our Solar System. Von Nei{\ss}l and Hoffmeister's Fireball Catalogue \citep{vonhof25} contains visually determined orbits of fireballs. Many of their entries are hyperbolic with respect to the Sun, that is, their velocities are so large that they cannot be gravitationally bound to our solar system. At the Earth's orbit, the parabolic or escape velocity with respect to the Sun is about 42 km~s$^{-1}$, and 79\% of \cite{vonhof25}'s orbits exceed this value, some ranging up to 99 km~s$^{-1}$ (as quoted in \cite{lov54}). A simple interpretation of hyperbolic meteors was that, since they were not bound to our Solar System, they must be from outside it and thus represent material originating elsewhere in our Galaxy. However, not all researchers agreed that substantial numbers of meteors had hyperbolic velocities, attributing them rather to measurement error. The discussion of the reality of hyperbolic meteors centred on the sporadic meteors: many showers were already accepted to be on bound orbits near those of their parent comets and thus part of streams of particles originating within our planetary system. A vigorous debate as to the existence of hyperbolic meteors spanned the next few decades. \cite{fis28} and \cite{wat39} concluded that since the hyperbolic meteors of \cite{vonhof25} largely coincided in time with the major showers and in space with the ecliptic plane, that they were unlikely to be of true interstellar origin; they instead concluded that a systematic over-estimation of the velocities, which were at the time still measured by observers using the naked eye, was more likely. Others argued conversely that, if some showers were in fact interstellar in nature, meteor showers and hyperbolic meteors might be expected to coincide in some cases. The problem was considered sufficiently important that the Harvard Observatory organized the Arizona Expedition to resolve the question \citep{shaopiboo32}. Ernst {\"O}pik led a campaign that erected two 'meteor houses' in Arizona where observers would record meteor data -- still taken visually by human observers -- in an organized fashion. The houses, really small protective shelters for the observers, had windows with built-in reticules to aid in positional measurements. The campaign also made use of the clever 'double-pendulum' or 'rocking mirror' technique whereby the meteors' motion would by translated into a pseudo-cycloidal motion, the number of cusps/loops of which could be used to facilitate trail length and speed measurements \citep{shaopiboo32,mcfash10}. The initial results of this work \citep{opi40} reported 57.2\% of meteors as being hyperbolic, with heliocentric velocities in rare cases reaching over 280 km~s$^{-1}$. The expedition leader was aware of the potential pitfalls: ``As in all kinds of visual observations of meteors in which the observer has finally to rely upon his memory, considerable accidental and systematic errors are involved in the observed velocities too; in a statistical discussion of velocities such as given below the data must be freed, in the first place, from the influence of these errors. Only after that can the bearing of the statistical data upon cosmic problems be investigated'' \citep{opi40}. Despite the Arizona results, some astronomers remained sceptical that sporadic meteors were interstellar. \cite{por43,por44}, working from other visual observations, concluded that meteors are not hyperbolic in any great numbers, and emphasized the need for a careful statistical analysis of a sample with known errors. The interstellar hypothesis received a serious blow when the first photographic meteor studies \citep{whi40} identified the Taurid stream, once conjectured to be an interstellar stream, and found it to be bound to the Sun and associated with Comet Encke. Though {\"O}pik continued to stand by the Arizona expedition's findings of a high fraction of hyperbolic sporadic meteors \citep{opi50}, new photographic programs began finding the interstellar fraction of meteors to be quite small. The Harvard Super-Schmidt photographic program \citep{jacwhi61} detected very few hyperbolic orbits. Radar meteor observations from Jodrell Bank (\cite{almdavlov51,almdavlov52,almdavlov53,cle52}; see \cite{gun05} for a review) and from Ottawa \citep{mck49,mck51} showed little or no evidence for interstellar velocities. \cite{opi69} eventually conceded that there was a failure in the basic assumptions underlying the rocking mirror technique, due partly to height differences between sporadic and shower meteors, and partly due to 'psychological' differences in their perception by observers. Though the hyperbolic component is now recognized to be small at visual meteor sizes ($\gas$~1~mm), they have been detected convincingly in interplanetary space at smaller sizes. Dust detectors aboard the Ulysses, Galileo and Helios spacecraft \cite[]{gruzoobag93,fridorgei99,krulanalt07} have detected very small ($10^{-18}-10^{-13}$~kg) grains moving at speeds above the local solar system escape velocity and parallel to the local flow of interstellar gas. This result provides perhaps the first generally-accepted detection of interstellar meteoroids. However, these particles are too small to be detected as meteors at the Earth: sizes $\gas10^{-10}$~kg may be required for this. Meteor radars are typically more sensitive that the human eye or photographs and can detect particles much smaller than those seen in early surveys of sporadic meteors. The Advanced Meteor Orbit Radar (AMOR) reported that a few percent of meteors with sizes of $\sim50\mu$m or masses$\sim10^{-10}$~kg were hyperbolic, many of which were proposed to be particles ejected from the $\beta$ Pictoris dust disk \citep{bag99,bag00,baggal01}. Radar hyperbolics have also seen by \cite{janmeimat01} at Arecibo though these were determined to be meteoroids accelerated by solar radiation pressure and originating within our own Solar System. Some Arecibo radar meteor detections have been interpreted as true interstellar meteoroids \citep{meijanmat02a,meijanmat02b}, and thus the possibility remains that substantial numbers of interstellar meteoroids reach the Earth at sizes that do not produce naked eye meteors. However, not all radar studies show evidence for hyperbolic meteors: radar observations at the Canadian Meteor Orbit Radar (CMOR) do not contain appreciable hyperbolics: less than 0.0008\% \citep{werbro04}. Harvard super-Schmidt photographic observations \citep{mcrpos61}, photographic meteor observations \citep{babkra67}, TV meteor observations \citep{jonsar85}, and image-intensified video meteor studies \citep{hawwoo97} all show a minority fraction of hyperbolic orbits. Between 1\% and 22\% of of meteors observed at the Earth by various surveys, optical and radar-based, have shown a hyperbolic component according to reviews by \cite{hawclowoo99} and \cite{bagmarclo07}. It remains unclear whether these represent true hyperbolics or result from experimental uncertainties. \cite{muswerbro12} present image-intensified optical results of a small number (22 of 1739) possible interstellar meteoroids but ultimately attribute these to measurement errors. Work by Hajdukov\'a and her colleagues \cite[]{hajpau02,hajhaj06,hajpau07,haj08,haj11} has shown that -- even with modern photographic and video techniques -- in many cases hyperbolic meteors only appear so as the result of measurement errors. The true population of interstellar meteoroids within our Solar System remains unknown. The most recent theoretical work on the expected component of true interstellar meteoroids in our Solar System is \cite{murweicap04}, but the question must ultimately be answered by measurement. However, even given an unequivocal measurement of hyperbolic velocity for a meteor, the question remains: did the particle originate outside our solar system? Given that processes within our Solar System might produce particles with high velocities and that could ``contaminate'' our sample of interstellar meteors, we need to understand the population of hyperbolic meteors produced internally to our planetary system in order to tease the two apart. Here we address the question of whether or not hyperbolic meteoroids could be produced within our own Solar System, in particular by the gravitational slingshot effect. It has long been recognized that planetary scattering must produce some hyperbolic meteors whose origin is contained wholly within our Solar System \citep{lov54,opi69}. Not that gravitational scattering is the only mechanism by which they might be produced. The ejection processes of comets may inject meteoroids directly onto hyperbolic orbits. Small ($\las 1\mu$m) cometary particles may feel enough radiation pressure to be unbound from the Sun independent of their ejection velocity from their parent (the so-called beta meteoroids). Collisional or rotational breakup of meteoroids and subsequent radiation pressure modification of their orbits are also thought to contribute to this population \citep{zoober75}. The magnetic fields of the planets Jupiter and Saturn are also able to accelerate electrically charged grains to escape speeds, as was measured by the Ulysses and Cassini spacecraft \citep{gruzoobag93,kemsrahor05,flakruham11}. Most these mechanisms only produce hyperbolic meteoroids if the particles are small enough that radiation pressure can play a role in accelerating them out of our Solar System, and are not effective for larger particles. Particles small enough to be ejected via radiation pressure are not easily detected by Earth-based meteor sensors, and so reports of high-speed meteors are unlikely to be of this origin. Direct cometary ejection may also produce larger particles on hyperbolic orbits without the action of radiation forces; however, meteoroids of this type are likely to be part of a freshly deposited meteoroid stream and unlikely to be mistaken for interstellar meteoroids. Gravitational scattering is perhaps the only source of meteoroids larger than a few microns in size which could be readily confused with an interstellar influx of particles. Here we examine the properties of gravitationally scattered meteoroids produced within our Solar System. Though in some sense ``contaminants'' of the interstellar meteoroid sample, their intrinsic properties are of interest as well for they come to us directly from the vicinity of the planets and may thus carry important information about these bodies and their environments. The techniques of meteor velocity measurement continue to improve and it is only a question of time before a substantial number of reliable hyperbolic meteors is measured: an understanding of the flux of such meteors produced within our Solar System is needed before the true interstellar component can be separated from the local one. | The number of hyperbolic meteors observed at Earth that might be produced by gravitational scattering by the planets is calculated: overall the numbers are small, certainly compared to the background rate of bound meteors. According to \cite{cambra11}, the sporadic flux of video meteors at Earth is $0.18\pm0.04$~meteoroids~km$^{-2}$~hr$^{-1}$ while for the interstellar flux \cite{muswerbro12} give an upper limit of $2 \times 10^{-4}$~meteoroids~km$^{-2}$~hr$^{-1}$ at optical sizes, so current limits from optical systems put the flux of interstellars at 1 in 1000 at most. Our work here predicts 1 meteor in $10^4$ at Earth will be a hyperbolic originating in a scattering event at Mercury. As a result, the contamination of a sample of interstellar meteors in this size range is expected to be at least at the 10\% level and possibly higher and so needs to be addressed. The problem is mitigated by the fact that the contamination is primarily at low $\vex$, much smaller than those expected of true interstellars. At radar sizes, CMOR \citep{werbro04} sees perhaps 1 in $10^5$ meteors as hyperbolic. Our work here still predicts 1 in $10^4$ internally-generated hyperbolics at these sizes, but again their excess velocities are too low to constitute a significant source of confusion in current samples. We conclude that hyperbolic particles can in principle be generated wholly within our Solar System and at speeds which rivals those expected of interstellar meteors. However, the properties of the meteoroid environment of our planetary system are not conducive scattering larger numbers of them onto Earth-intercepting orbits. Though selecting a sample of presumed-interstellar meteors solely on the basis of their heliocentric velocity is likely to produce a substantially contaminated sample, the internally-generated hyperbolics are relatively easy to account for, as their excess velocities at the Earth are expected to be only about 100~m~s$^{-1}$ in most cases. Higher $\vex$ may occur in exceptional cases or when a planet encounters a freshly-deposited debris stream from a comet. However, in all such cases a sufficiently precise velocity determination followed by a careful examination of the pre-atmospheric trajectory can determine whether such a scattering event occurred. Thus while the search for interstellar meteors is complicated by planetary scattering, continuing improvements in detection methods mean that the phenomenon is not likely to prove a substantial obstacle to the study of interstellar meteoroids. | 14 | 4 | 1404.2159 |
1404 | 1404.2645_arXiv.txt | { The late-B magnetic chemically peculiar star \vir\ is one of the fastest rotators among the intermediate-mass stars with strong fossil magnetic fields. It shows a prominent rotational modulation of the spectral energy distribution and absorption line profiles due to chemical spots and exhibits a unique strongly beamed variable radio emission. } { Little is known about the magnetic field topology of \vir. In this study we aim to derive, for the first time, detailed maps of the magnetic field distribution over the surface of this star. } { We use high-resolution spectropolarimetric observations covering the entire rotational period. These data are interpreted using a multi-line technique of least-squares deconvolution (LSD) and a new Zeeman Doppler imaging code based on detailed polarised radiative transfer modelling of the Stokes $I$ and $V$ LSD profiles. This new magnetic inversion approach relies on the spectrum synthesis calculations over the full wavelength range covered by observations and does not assume that the LSD profiles behave as a single spectral line with mean parameters. } { We present magnetic and chemical abundance maps derived from the Si and Fe lines. Mean polarisation profiles of both elements reveal a significant departure of the magnetic field topology of \vir\ from the commonly assumed axisymmetric dipolar configuration. The field of \vir\ is dipolar-like, but clearly non-axisymmetric, showing a large difference of the field strength between the regions of opposite polarity. The main relative abundance depletion features in both Si and Fe maps coincide with the weak-field region in the magnetic map. } { Detailed information on the distorted dipolar magnetic field topology of \vir\ provided by our study is essential for understanding chemical spot formation, radio emission, and rotational period variation of this star. } | \label{intro} The bright late-B star \vir\ (HR\,5313, HD\,124224, HIP\,69389) is one of the best known intermediate-mass magnetic chemically peculiar (CP) stars. This object shows abnormally strong lines of Si, weak lines of He, and a notable spectroscopic \citep{deutsch:1952} and photometric \citep{hardie:1958} variability with a period of $\sim$\,0.5~d. These periodic changes arise due to a combination of stellar rotation and an inhomogeneous distribution of chemical elements over the stellar surface. It is generally believed that high-contrast chemical spots are produced by an anisotropic atomic diffusion in the presence of strong, globally-organised magnetic field \citep{michaud:1981,alecian:1981,alecian:2010}, although a detailed diffusion theory capable of explaining the surface structure of individual magnetic CP stars is still lacking. Being one of the most rapidly rotating magnetic CP stars, \vir\ represents an optimal target for the reconstruction of the surface chemical spot distributions with the Doppler imaging (DI) method. Using this technique, maps of He, Mg, Si, Cr, and Fe were obtained for \vir\ by different authors \citep{goncharskii:1983,hiesberger:1995,hatzes:1997,kuschnig:1999}. The study by \citet{krticka:2012} showed that these uneven chemical distributions provide a successful quantitative explanation for the observed spectrophotometric behaviour of \vir\ in the ultraviolet and optical wavelength regions. The fast rotation and stability of the photometric light curves of \vir\ enable a very precise period determination and allows one to study subtle long-term variations of stellar rotation. \vir\ is one of a few stars for which rotational period changes have been securely detected. Remarkably, this star appears to show period glitches \citep{pyper:1998,pyper:2013} or possibly cyclic period variation \citep{mikulasek:2011} rather than a constantly increasing period as expected from magnetic spin down models \citep{ud-doula:2009}. The origin of these period changes is currently unknown. In addition to the photometric and spectroscopic variability typical of magnetic CP stars, \vir\ exhibits a unique radio emission signature, not seen in any other star. \citet{trigilio:2000} reported two sharp 100 per cent circularly polarised radio pulses occurring at particular rotational phases. This emission was subsequently studied in different radio bands \citep{trigilio:2008,trigilio:2011,ravi:2010,lo:2012} and interpreted as an electron cyclotron maser emission originating in the stellar magnetosphere above one of the magnetic poles \citep[see][and references therein]{lo:2012}. The sharp radio emission pulses can also be used for a very precise measurement of the stellar rotational period. However, it is not clear if the phase shifts observed in the radio are related to intrinsic variation of the stellar rotation or to instabilities in the region where the radio emission is produced \citep{ravi:2010,pyper:2013}. Evidently, the magnetic field is a key parameter for understanding different variability phenomena observed in \vir. The strength and geometry of the surface magnetic field influences the photospheric distribution of chemical elements, while the structure of the magnetosphere at a distance of 2--3 stellar radii is the main ingredient of any possible non-thermal radio emission process. The presence of a kG-strength magnetic field in \vir\ was established by \citet{landstreet:1977} and \citet{borra:1980} using a Balmer line magnetometer. Their mean longitudinal magnetic field (\bz) measurements revealed a smooth, roughly sinusoidal variation with an amplitude of $\approx$\,600~G and reversing sign. Several oblique dipolar models were derived using these \bz\ data \citep{borra:1980,trigilio:2000}, resulting in estimates of the polar field strength $B_{\rm d}$\,=\,3.0--4.5~kG and magnetic obliquity $\beta$\,=\,70--90\degr, depending on the assumed inclination angle $i$. On the other hand, \citet{hatzes:1997} and \citet{glagolevskij:2002} argued that the magnetic field geometry of \vir\ is better described by an offset dipole or by a combination of dipolar and quadrupolar components of comparable strengths. Observations of the mean longitudinal magnetic field available in the literature are insufficient to definitively distinguish between these alternative magnetic field models. The goal of our investigation is to obtain a much more precise and detailed picture of the magnetic field topology of \vir\ by using new high-resolution circular polarisation observations. To interpret these data we developed a new version of the Zeeman Doppler imaging (ZDI) methodology, which for the first time combines detailed polarised spectrum synthesis calculations based on realistic model atmospheres with a multi-line approach widely used in stellar spectropolarimetry. The rest of this paper is organised as follows. Sect.~\ref{methods} presents spectropolarimetric observations and outlines corresponding analysis methodology, including calculation of the mean polarisation profiles and their interpretation in terms of the magnetic and chemical stellar surface maps. In Sect.~\ref{results} we present revised stellar parameters, discuss new longitudinal field measurements, and derive magnetic field topology and chemical abundance distributions with the help of ZDI inversions. The results of our investigation are summarised and discussed in the context of other recent studies in Sect.~\ref{discussion}. | \label{discussion} \subsection{Comparison with previous DI studies} Previously conventional abundance DI mapping was carried out for \vir\ by \citet{hatzes:1997} and \citet{kuschnig:1999}. The first study presented a Si equivalent width map derived from a single line; the second study obtained abundance maps for five different elements, including Si and Fe for which multiple lines were used. Both analyses gave a high weight to the strongest \ion{Si}{ii} lines. Consequently, their results may be distorted by an interplay of horizontal and vertical abundance inhomogeneities and therefore are less representative than our chemical spot maps derived from the average line profiles. In addition, previous DI studies completely neglected magnetic field and ignored local modifications of the atmospheric structure. Despite these differences, main features in our Si and Fe abundance maps agree reasonably well with the distributions obtained by \citet{kuschnig:1999}. Taking into account the 0.5 phase offset between the ephemeris adopted in the two studies, we both find an underabundance area around phase 0.0, with Fe underabundance zone trailing slightly behind the one for Si. Also, similar to our study, the maps by \citet{kuschnig:1999} exhibit arch-like features around phase 0.5--0.6 when the Si and Fe abundance reaches maximum. Kuschnig et al. reported a surface abundance variation from +2.5 to $-1.5$ for Si and from +1.3 to $-0.7$ for Fe. These ranges are somewhat larger compared to our study, which may be related to a lower inclination angle adopted by these authors and/or their neglect of the continuum intensity variation due to spots. \subsection{Verification of the Si abundance map} In this study we attempted to derive chemical abundance distributions from LSD profiles, which represent an average over many spectral features with different characteristics. Since this approach is used here for the first time, it is of interest to check how the resulting chemical spot maps reproduce the Stokes $I$ and $V$ profiles of individual spectral lines. In the particular case of \vir\ such an assessment is best carried out for Si since there are many reasonably unblended lines of this element in the stellar spectrum. Using the results of magnetic and chemical abundance mapping with the Si LSD profiles (see Figs.~\ref{fig:fld_si} and \ref{fig:abn}), we computed about a dozen of \ion{Si}{ii} spectral lines and compared their predicted profiles with observations. Figure~\ref{fig:lines} illustrates this comparison for five representative spectral features. \begin{figure*}[!th] \centering \fifps{15cm}{23472_f14-eps-converted-to.pdf} \caption{Comparison of the observed Stokes $I$ and $V$ profiles (thin solid line) of individual \ion{Si}{ii} lines with the theoretical spectra (thick solid lines) predicted by the magnetic field model and Si abundance distribution derived from silicon LSD profiles (see Figs.~\ref{fig:fld_si} and \ref{fig:abn}). The oscillator strengths of the high-excitation Si blends at $\lambda$ 4673 and 5868~\AA\ have been increased by 0.5--1.5~dex relative to the values given in VALD to match the observations. The spectra predicted with original oscillator strengths are shown with dotted lines. The spectra corresponding to different rotation phases (indicated to the right of the Stokes $V$ panel) are offset vertically. The bars at the lower left corner of each panel indicate the horizontal (1~\AA) and vertical (1\% of $I_{\rm c}$) scales.} \label{fig:lines} \end{figure*} We find that the LSD Si abundance map reproduces fairly well weak and intermediate strength \ion{Si}{ii} lines with the excitation potential of the lower level $E_{\rm i}\approx10$~eV (e.g. \ion{Si}{ii} 5041, 5056~\AA). However, a number of discrepancies is found for other lines. For example, the higher excitation \ion{Si}{ii} lines, such as 4276 and 5268~\AA\ ($E_{\rm i}=13$--15~eV), appear substantially stronger in observations than in our spectrum synthesis. We had to increase their oscillator strengths by 0.5--1.5~dex relative to the values given in VALD to match the observed profiles. On the other hand, very strong silicon lines such as \ion{Si}{ii} 6347 and 6371~\AA\ (the latter not shown in Fig.~\ref{fig:lines}) exhibit a somewhat different profile shape, especially in the phase interval 0.23--0.42 when a distinct flattening of the line cores is observed. This line to line scatter is almost certainly caused by an inhomogeneous vertical distribution of Si in the atmosphere of \vir. If the Si abundance decreases with height, as found for magnetic CP stars of similar \teff\ \citep{glagolevskii:2005,ryabchikova:2006} and predicted by theoretical radiative diffusion models \citep{leblanc:2009,alecian:2010}, one might expect anomalously strong high-excitation lines (formed in deep atmospheric layers) and profile distortions in strong lower excitation features (sensitive to a large range of heights, including high atmospheric layers). This qualitative scenario describes reasonably well our observations of \vir. This star appears to be an interesting target for studying how the vertical stratification of Si changes across the stellar surface. In conclusion, our LSD-based Si abundance map provides a satisfactory description of the intensity and circular polarisation spectra of typical individual \ion{Si}{ii}. It does not, however, explain the profiles of certain features likely to be affected by chemical stratification. \subsection{Magnetosphere and radio emission of \vir} \vir\ shows a complex and interesting behaviour in the radio domain. Similar to some other magnetic chemically peculiar stars it shows a smoothly varying, rotationally modulated gyrosynchrotron emission \citep{leone:1996}. In addition, unlike any other known radio-loud stellar source, it exhibits two 100 per cent circularly polarised radio pulses every rotation period \citep{trigilio:2000}. These radio pulses are attributed to the electron cyclotron maser (ECM) mechanism, which produces emission directed at a certain narrow ranges of angles to the magnetic field lines. This emission is believed to be produced in the stellar magnetosphere at a distance of about two stellar radii in the vicinity of one of the magnetic poles. Different modifications of the ECM model have been examined for \vir\ \citep[see][]{lo:2012}, but it is still not entirely clear which one is the most appropriate. Our work has provided several important constraints for the nature of the magnetosphere of \vir\ and its relation to the unique coherent radio emission of this star. First, our magnetic observations are essentially concurrent with the recent radio studies. This allows us to establish the phase relation between the longitudinal magnetic field curve and the times of arrival of the radio pulses (see Sect.~\ref{longit}) with a much higher certainty than in the previous studies, which had to rely on the \bz\ measurements by \citet{borra:1980} acquired several decades ago. Along these lines, we report new precise phase offsets between the radio pulses at 1.4~GHz and the phases of the zero longitudinal magnetic field. Furthermore, we found for \vir\ non-negligible deviations from a dipolar magnetic field topology. All previous radio emission models considered purely dipolar field configurations, occasionally speculating about the presence of non-dipolar field components to explain the asymmetry of the radio pulses \citep[e.g.][]{trigilio:2000}. Our ZDI maps characterised these non-dipolar field components and provided detailed information about the surface magnetic field topology that can be used for realistic modelling of the stellar magnetosphere and the radio emission processes. Although such analysis is beyond the scope of our study, it is already clear that our magnetic maps are generally compatible with the asymmetry of the beamed radio emission. Evidently, the negative field region is more compact, is geometrically much better defined and has a stronger field compared to the positive field zone. Therefore, association of the radio pulses with only the negative surface field extremum comes as no surprise given our ZDI results. Magnetospheric models invoked to explain the radio emission of \vir\ estimated that magnetic field channels the mass loss out to the Alfv\'en radius of $\sim$\,15\,$R$ \citep{trigilio:2004}. This is much larger than the Keplerian radius, thus providing conditions for accumulation of material in corotating clouds or in a warped disk \citep{townsend:2005a}. For \vir\ several studies suggested the presence of a torus of relatively cold, dense material surrounding the star at the magnetic equatorial plane \citep{trigilio:2004,leto:2006}. \citet{lo:2012} argued that refraction from this cold torus represents the most promising version of the ECM mechanism to explain the frequency dependence and shapes of the radio pulse profiles. The circumstellar material in the torus may produce characteristic spectroscopic variation in the hydrogen Balmer lines, similar to the behaviour commonly observed in hotter early-B magnetic stars \citep{townsend:2005,petit:2013}. We searched for such variations using the Narval spectra of \vir. The H$\beta$--$\delta$ lines exhibit smooth changes in their Stark wings due to a non-uniform He distribution over the stellar surface \citep{kuschnig:1999,shulyak:2004}. The maximum line width is found at phase $\approx$\,0.10; the minimum occurs at phase $\approx$\,0.65. In addition, H$\alpha$ shows subtle variation, primarily on the red side of the line core, similar to the spectroscopic changes described earlier by \citet{shore:2004}. Relative to the mean profile the line core alternates between a lack of absorption at phase $\approx$\,0.24 and an excess at phase $\approx$\,0.75. This variability in H$\alpha$ is, again, likely a result of He abundance spots rather than cold, dense gas trapped in the stellar magnetosphere. Thus, our observations provide no direct evidence for the existence of circumstellar material in the vicinity of \vir. \subsection{Conclusions} In this study we developed a new Zeeman Doppler imaging procedure, combining for the first time the LSD multi-line technique with the detailed magnetic spectrum synthesis using realistic model atmospheres. This new magnetic inversion methodology was successfully applied to a set of high-quality circular polarisation observations of the well-known magnetic chemically peculiar star \vir. The main conclusion of our investigation can be summarised as follows. \begin{itemize} \item The new ZDI methodology enables simultaneous reconstruction of the stellar magnetic field topology and distribution of chemical abundance (or temperature for cool active stars) from the LSD Stokes profiles without making simplifying approximations typical of previous ZDI studies. In particular, the new ZDI method does not assume that a response of the LSD profiles to the variation of magnetic field, chemical abundance or temperature is equivalent to that of a single spectral line. \item The revised fundamental parameters of \vir\ suggest that this is a young magnetic star observed from an intermediate inclination. It is a relatively rapid rotator, presumably due to its youth, but its equatorial rotational velocity is not high enough to significantly distort the spherical shape of the stellar surface. \item The magnetic field topology of \vir\ is generally dipolar-like, but is characterised by a strong field modulus asymmetry between the two poles and by a major deviation from an axisymmetric configuration. This field geometry is mainly poloidal and nearly all field energy is contained in the harmonic modes with the angular degree $\ell$\,=\,1--3. \item Our calculations prove that neither the purely dipolar magnetic field model nor a more complex dipole plus quadrupole field topology derived from the longitudinal field curve of \vir\ can describe the observed circular polarisation signatures in spectral lines. \item The chemical abundance maps of Si and Fe are dominated by a large underabundance region where the field is positive, predominantly radial, and weak. In contrast, the spots of element overabundance occur at the stronger negative and positive field regions where a significant horizontal field component is also present. \item The magnetic and chemical abundance maps obtained from the Si LSD profiles match variability of the Stokes $I$ and $V$ spectra in individual \ion{Si}{ii} lines. However, observations show some line to line scatter in the Si profile shapes and intensities, which we tentatively attribute to a vertical stratification of silicon in the atmosphere of \vir. \item The magnetic field maps derived in our study explain the asymmetry of the pulsed radio emission of \vir\ and provide detailed constraints for future magnetospheric modelling. \end{itemize} | 14 | 4 | 1404.2645 |
1404 | 1404.6470_arXiv.txt | { We simultaneously analyze the spectral line energy distributions (SLEDs) of CO and \Hmol\ of six local luminous infrared (IR) Seyfert galaxies. For the CO SLEDs, we used new \Herschel\slash SPIRE FTS data (from $J=4-3$ to $J=13-12$) and ground-based observations for the lower-$J$ CO transitions. The \Hmol\ SLEDs were constructed using archival mid-IR \Spitzer\slash IRS and near-IR VLT\slash SINFONI data for the rotational and ro-vibrational \Hmol\ transitions, respectively. In total, the SLEDs contain 26 transitions with upper level energies between 5 and 15\,000\,K. A single, constant density, model ($n_{\rm H_2}\sim 10^{4.5-6}$\,cm$^{-3}$) with a broken power-law temperature distribution reproduces well both the CO and \Hmol\ SLEDs. The power-law indices are $\beta_1\sim 1-3$ for warm molecular gas ($20$\,K$<T<100$\,K) and $\beta_2\sim 4-5$ for hot molecular gas ($T>100$\,K). We show that the steeper temperature distribution (higher $\beta$) for hot molecular gas can be explained by shocks and photodissociation region (PDR) models, however, the exact $\beta$ values are not reproduced by PDR or shock models alone and a combination of both is needed. We find that the three major mergers among our targets have shallower temperature distributions for warm molecular gas than the other three spiral galaxies. This can be explained by a higher relative contribution of shock excitation, with respect to PDR excitation, for the warm molecular gas in these mergers. For only one of the mergers, IRASF~05189--2524, the shallower \Hmol\ temperature distribution differs from that of the spiral galaxies. The presence of a bright active galactic nucleus in this source might explain the warmer molecular gas observed. } | Molecular gas is an important phase of the interstellar medium (ISM). This phase contains a significant fraction of the total mass, and stars form in it. But the study of molecular gas presents some complications. First, the lower energy levels of \Hmol, the main component of the ISM phase, have energies $>500$\,K, thus in cold molecular gas ($T< 100$\,K) most of the \Hmol\ is in the fundamental state and no \Hmol\ emission lines are produced. And second, only the near infrared (IR) ro-vibrational \Hmol\ transitions, with $E_{\rm up}>6000$\,K, are observable from ground telescopes, so only very high-temperature ($T > 1500$\,K) molecular gas can be detected. To overcome the first caveat, other abundant molecules with observable transitions in the millimeter range (like CO, HCN, etc.) are used as tracers of molecular gas. In particular, the lowest rotational transitions of CO, the second most abundant molecule, are commonly used to study the molecular gas content of galaxies. However, these low-$J$ CO transitions mainly originate in the coldest molecular gas. Thus, ground observations are limited to the study of either the warmest or the coldest molecular gas. Just recently, thanks to IR and sub-millimeter space observatories like the \textit{Infrared Space Observatory} (\textit{ISO}; \citealt{Kessler1996}), the \textit{Spitzer Space Telescope} \citep{Werner2004}, and the \Herschel\ Space Observatory \citep{Pilbratt2010Herschel}, the rotational \Hmol\ transitions as well as the intermediate-$J$ CO transitions became accessible for a large number of local galaxies (e.g., \citealt{Rigopoulou02,Roussel07,vanderWerf2010,Pereira2013}). Therefore, now for the first time, it is possible to obtain a complete snapshot of molecular gas emission and study its physical properties (temperature, density, column density, etc.) and the excitation mechanisms (ultraviolet (UV) radiation, shocks, and X-ray and cosmic rays). In this work, we present new data obtained by the Fourier transform spectrometer (FTS) module of the Spectral and Photometric Imaging Receiver (SPIRE) instrument on-board \Herschel\ \citep{Griffin2010SPIRE,Naylor2010,Swinyard2010} for six local active luminous IR galaxies. These \FTS\ data cover the 210--670\micron\ (450--1440\,GHz) spectral range, so the mid-$J$ CO lines ($J=4-3$ to $J=13-12$) are observed. We completed the CO spectral line energy distributions (SLEDs) with ground-based observations of the three lowest $J$ CO transitions. In addition, we complemented the CO SLEDs with the \Hmol\ SLEDs obtained from near- and mid-IR observations of these galaxies. We used the available mid-IR spectroscopy obtained by the \Spitzer\ IR spectrograph (IRS; \citealt{HouckIRS}) to measure the lowest rotational \Hmol\ transitions (e.g., \citealt{Wu2009, Pereira2010}), and near-IR integral field spectroscopy obtained by the Spectrograph for INtegral Field Observations in the Near-Infrared (SINFONI; \citealt{Eisenhauer2003}) on the Very Large Telescope (VLT) for the ro-vibrational \Hmol\ transitions. For the first time, we have performed a radiation transfer analysis of the whole set of molecular lines together (i.e., CO rotational and \Hmol\ rotational and ro-vibrational) in local IR bright galaxies. In total, the compiled CO and \Hmol\ SLEDs contain information for 26 transitions with upper level energies between 5 and 15\,000\,K, thus the emission from most of the molecular gas is included. The paper is organized as follows. In Section \ref{s:observations}, we present the sample and the data reduction. Sections \ref{s:rad_models} and \ref{s:sled_fit} describe the radiative transfer models used, and the fitting of the SLEDs. The cold-to-warm molecular gas ratio and the heating mechanisms are discussed in Sections \ref{s:cold-to-warm_ratio} and \ref{s:heating}, respectively. We summarize the main results in Section \ref{s:conclusions}. | \label{s:conclusions} We have studied the integrated CO and \Hmol\ emission of six local IR bright galaxies using non-LTE models. Assuming a broken power-law distribution for the molecular gas temperatures, our model reproduces both the CO SLED (from $J_{\rm up}=1$ to $J_{\rm up}=13$) and the \Hmol\ SLED ($J_{\rm up}\leq 7$ for the lowest three vibrational levels) in our sample of galaxies. The main findings of this work are summarized in the following: \begin{enumerate} \item With a single power-law temperature distribution it is not possible to fit simultaneously the CO and \Hmol\ SLEDs. The \Hmol\ SLEDs have a much steeper power-law index ($\beta_2\sim 4-5$) than the CO SLED ($\beta_1\sim 1-3$). This is the expected behavior for the temperature distributions in PDRs and shocks. \item We found that for most of the galaxies, the models with $n_{\rm H_2}=10^6$\,cm$^{-3}$ provide the best fit to the observed data, thus the majority of the transitions are close to LTE. The minimum acceptable density for the warm gas is $\sim n_{\rm H_2}=10^{4 \pm 0.5}$\,cm$^{-3}$, lower densities would imply CO abundances higher than the atomic C abundance. Likewise, we obtained that the minimum CO abundance in the warm gas is $x_{\rm CO}\sim 10^{-6}-10^{-5}$ assuming LTE conditions. \item The column densities of the warm molecular gas ($T>20$\,K) represents between 10 and 100\,\% of the molecular gas traced by the \CO1 transition. \item We used PDR and shock models to determine the excitation mechanism of the molecular gas. Our models show that the temperature distributions are steeper for PDRs than for shocks. We also found that the temperature distribution of the warmest gas ($T>100$\,K) emitting in \Hmol\ is steeper than that of coldest gas ($T>30$\,K), which produces the mid-$J$ CO emission for both PDR and shocks models. This is because of the different main coolant of the warm and cold molecular gas (\Hmol, and CO, respectively). \item Neither PDR nor shocks alone can explain the derived temperature distributions. A combination of both is needed to reproduce the observed SLEDs. In this case, the lower-$J$ CO transitions would be dominated by PDR emission, whereas the higher-$J$ CO transitions would be dominated by shocks. \item The three major mergers among our targets (NGC~34, IRASF~05189--2524, and UGC~05101) have shallower temperature distributions for CO than the other three spirals. This suggests that the relative contribution of shocks to the heating of warm molecular gas ($T<100$\,K) in these major mergers is higher than in the other three spirals. \item For only one of the mergers, IRASF~05189--2524, the shallower \Hmol\ temperature distribution (hot molecular gas) suggests that the relative importance of shocks is high. This galaxy also has a bright AGN that dominates the bolometric luminosity, which can contribute to the molecular gas heating. For the other two mergers, the \Hmol\ temperature distribution is similar to that of the spiral galaxies. Therefore the shocks producing the extra contribution to the CO emission in these mergers are not able to heat the molecular gas to temperatures higher than 100\,K, which would be necessary to see the \Hmol\ emission. \end{enumerate} | 14 | 4 | 1404.6470 |
1404 | 1404.4196_arXiv.txt | Debris disks are thought to be sculptured by neighboring planets. The same is true for the Edgeworth-Kuiper debris disk, yet no direct observational evidence for signatures of giant planets in the Kuiper belt dust distribution has been found so far. Here we model the dust distribution in the outer solar system to reproduce the dust impact rates onto the dust detector onboard the New Horizons spacecraft measured so far and to predict the rates during the Neptune orbit traverse. To this end, we take a realistic distribution of transneptunian objects to launch a sufficient number of dust grains of different sizes and follow their orbits by including radiation pressure, Poynting-Robertson and stellar wind drag, as well as the perturbations of four giant planets. In a subsequent statistical analysis, we calculate number densities and lifetimes of the dust grains in order to simulate a collisional cascade. In contrast to the previous work, our model not only considers collisional elimination of particles, but also includes production of finer debris. We find that particles captured in the $3$:$2$ resonance with Neptune build clumps that are not removed by collisions, because the depleting effect of collisions is counteracted by production of smaller fragments. Our model successfully reproduces the dust impact rates measured by New Horizons out to $\approx{}23\AU$ and predicts an increase of the impact rate of about a factor of two or three around the Neptune orbit crossing. This result is robust with respect to the variation of the vaguely known number of dust-producing scattered disk objects, collisional outcomes, and the dust properties. | The symbiosis between debris disks and planets is multifaceted. Not only do both represent natural outcomes of the planetesimal and planet accretion processes, it has also long been realized that planets should sculpt the disks by their gravity, and thus the observed structure in debris disks can be used as a tracer of planets. In one case~--- $\beta$~Pic~--- the observed disk structure has been undoubtedly attributed to interactions with a directly imaged planet \citep{Lagrange-et-al-2009,Lagrange-et-al-2010,Lagrange-et-al-2012a}. In a few other cases, such as Fomalhaut \citep{Kalas-et-al-2008,Kalas-et-al-2013}, HR~8799 \citep{Su-et-al-2009}, and HD 95086 \citep{Rameau-et-al-2013a}, the relation between the disks and directly imaged planets has yet to be understood. Many more systems have been found where resolved debris disks reveal various types of structure, possibly driven by gravity of unseen planets orbiting in the disks' inner cavities, exemplified by the clumpy disk of $\epsilon$~Eri \citep{Greaves-et-al-2005} or spirals in the HD~141569 disk \citep{Wyatt-2005b}. However, some of these features are yet to be confirmed observationally, for instance requiring disambiguation with possible background objects. Also, models of planet-induced structure in the disk suffer from many uncertainties, especially those arising from poorly known dust properties and from difficulties of including collisions into the models. Besides, alternative explanations of the observed disk structure that do not require the presence of planets are possible \citep[e.g.][]{Artymowicz-Clampin-1997,Grigorieva-et-al-2007,Debes-et-al-2009}. It is natural and tempting to look at our solar system. Here, the orbits and masses of planets are precisely known. Largely known are also the populations of parent bodies, maintaining its debris disk, which include asteroids, comets, and Edgeworth-Kuiper belt (EKB) objects (EKBOs). Less well known are the properties of the dust cloud that these parent bodies replenish and maintain. Yet a dust ring encompassing the Earth orbit has been discovered and its structure successfully ascribed to resonant interactions between dust grains and the Earth \citep{Dermott-et-al-1994}. In the outer solar system, it has been predicted that Neptune should create two dusty clumps just exterior to its orbits, one ahead and one behind the planet slightly outside its orbit \citep{Liou-Zook-1999}. The simulated images of those clumps from the \citet{Liou-Zook-1999} paper have been used in dozens of subsequent papers as a vivid prototype of what can also be expected in extrasolar debris disks. However, the predicted EKB clumps are still lacking direct observational confirmation. Remote observations of thermal emission are hampered by an extremely low dust density in the EKB debris disk \citep{Vitense-et-al-2012}, which causes the foreground emission of the zodiacal cloud to outshine the thermal flux from the EKB dust. Luckily, the New Horizons spacecraft is now on its way to the outer solar system. Onboard, beside several other instruments, is the Venetia Burney Student Dust Counter (SDC), capable of measuring impacts of grains between $10^{-12}\g< m < 10^{-9}\g$ in mass \citep{Horanyi-et-al-2008}. New Horizons will traverse the trailing clump around early 2015 and might see a dust enhancement there. This paper addresses the question of whether the SDC has chances to detect the dust clump, proving its existence, and whether the expected data can be used to constrain the models. So far, there have been several attempts to model the EKB dust clumps \citep[e.g.,][]{Liou-Zook-1999,MoroMartin-Malhotra-2002,MoroMartin-Malhotra-2003,Kuchner-Stark-2010,Poppe-et-al-2010,Han-et-al-2011}. However, most of these studies ignored possible effects of grain-grain collisions. \citet{Kuchner-Stark-2010} did include them with the aid of their novel ``collisional grooming'' algorithm. Nevertheless, the collisions in their model only acted to eliminate the colliders; it is not surprizing therefore that the collisions in their model tends to erase the clumps. Including collisional {\em production} of fine debris might counteract this, putting dust enhancements back into play; we check this in this paper. Besides, none of the previous studies provided detailed predictions for the dust impact rates for the SDC detector aboard New Horizons along its trajectory; we do this here. We also try to use a more realistic ``true'' distribution of the EKB objects, acting as parent bodies for the dust \citep{Vitense-et-al-2010}, and to include perturbations of all four giant planets. Section~\ref{sec:dust_production} describes the procedure and results of the modeling in a collisionless approximation. Section 3 does the same with the collisions included. Conclusions are made in section 4 and discussed in section~5. | \label{sec:conclusion} In this paper we developed a model to reproduce and predict dust impact rates onto the dust detector of New Horizons. To this end, we used the debiased EKBO populations from \citet{Vitense-et-al-2010} and launched $27000$ dust grains of nine sizes. The particle sizes were chosen in such a way as to cover the entire mass sensitivity range of the dust detector aboard New Horizons. We then integrated their orbits by including the gravity of the four giant planets, stellar wind and the Poynting-Robertson effect, and recorded the position and velocities every orbit of Neptune. The resulting $2 \times 10^9$ records were used to draw a collisionless dust density map. We then complemented this collisionless model with a post-processing algorithm to roughly simulate a collisional cascade. In the collisionless approximation, our model fully reproduces the results obtained previously with similar methods. The most salient feature of all collisionless models, pioneered by the \citet{Liou-Zook-1999} study, is a pair of resonant clumps ahead and behind the Neptune location. These clumps are dominated by grains that are large enough to be efficiently captured in resonances with Neptune. Including collisions has a potential of changing the number density map for small (submicrometer-sized) grains dramatically. Since it is these grains that are predominantly detected by the New Horizons dust detector, the issue is important. \citet{Kuchner-Stark-2010} made the first attempt to include collisions in the models of azimuthal structure in the EKB region and showed that the collisonal elimination of grains should essentially wash out the clumps. However, they only considered collisional loss of grains. Our model refines theirs by adding the collisional replenishment of small grains. This tends to increase the amount of dust in the clumps almost back to the level predicted by the collisionless models. We have calculated the dust impact fluxes onto the New Horizons dust detector to find that it should be able to detect an increase of particle flux when it enters the region of the clump ahead the position of Neptune. The dust flux should increase by a factor of two or three. By varying the assumed material composition, mechanical strength of dust and the exponent of the crushing law, we checked that this result is rather robust, making our predictions quite certain. At the same time, this implies that New Horizons will probably not be able to provide additional contraints on the dust properties. The same applies to uncertainties stemming from a poorly known distribution of the dust parent bodies, especially of the scattered population of the Kuiper belt objects. Including or excluding this population from the simulations does not affect the predicted impact rates considerably. | 14 | 4 | 1404.4196 |
1404 | 1404.3639_arXiv.txt | In a recent work, Baldi et al. highlighted the issue of cosmic degeneracies, consisting in the fact that the standard statistics of the large-scale structure might not be sufficient to conclusively test cosmological models beyond $\Lambda $CDM when multiple extensions of the standard scenario coexist in nature. In particular, it was shown that the characteristic features of an $f(R)$ Modified Gravity theory and of massive neutrinos with an appreciable total mass $\Sigma _{i}m_{\nu _{i}}$ are suppressed in most of the basic large-scale structure observables for a specific combination of the main parameters of the two non-standard models. In the present work, we explore the possibility that the mean specific size of the supercluster spines -- which was recently proposed as a non-standard statistics by Shim and Lee to probe gravity at large scales -- can help to break this cosmic degeneracy. By analyzing the halo samples from N-body simulations featuring various combinations of $f(R)$ and $\Sigma _{i}m_{\nu _{i}}$ we find that -- at the present epoch -- the value of $\Sigma _{i}m_{\nu _{i}}$ required to maximally suppress the effects of $f(R)$ gravity on the specific sizes of the superclusters spines is different from that found for the other standard statistics. Furthermore, it is also shown that at higher redshifts ($z\ge 0.3$) the deviations of the mean specific sizes of the supercluster spines for all of the four considered combinations from its value for the standard $\Lambda$CDM case are statistically significant. | The clustering of the large-scale structures (LSS, hereafter) and its evolution with redshift encode information about the statistical properties of the initial conditions of the universe and the laws governing the gravitational instability processes that determine the growth of primordial density perturbations. The basic LSS statistics, such as the two point-correlation function of the density field, the mass function of galaxy clusters and the halo bias, are frequently employed to quantify the main properties of the LSS clustering, thereby providing a handle on the underlying cosmological model. Such basic statistics have played a vital role in establishing the concordance $\Lambda $CDM scenario by allowing to place tight constraints on its basic cosmological parameters. Furthermore, the standard LSS statistics are routinely employed to test not only the key cosmological parameters of the $\Lambda$CDM cosmology \citep[e.g.,][and references therein]{addison-etal13,didio-etal14} but also the viability of alternative cosmological models, such as e.g. the Warm Dark Matter model \citep[WDM, see e.g.][]{SM11,viel-etal12}, the coupled Dark Energy scenario \citep[cDE, see e.g.][]{moresco-etal13}, and various Modified Gravity theories \citep[MG, see e.g.,][]{SR08,SK09,stril-etal10,lombriser-etal12,abebe-etal13}. The main motivation behind most of these alternative models was to overcome the observational and theoretical shortcomings of the $\Lambda$CDM cosmology. For instance, the free-streaming effect of the WDM particles has been invoked as a possible solution to the tension between the $\Lambda$CDM predictions and astrophysical observations on (sub-)galactic scales \citep[e.g.,][and references therein]{menci-etal12}, even though recent constraints from the Lyman-alpha forest seem to exclude the WDM particle mass range required to address the tension \citep[][]{viel-etal13}. Similarly, the dark sector interactions that characterise the cDE scenario can alleviate the fine-tuning problems of the cosmological constant \citep[e.g.,][]{wetterich95,amendola00,amendola04}, while possible large-scale deviations of the laws of gravity from their standard GR form can accommodate the present acceleration of the universe without requiring dark energy with negative pressure \citep[for a comprehensive review, see][]{mg_review12}. However, some cautionary remarks have been recently raised on the rosy prospects for the basic statistics of LSS as a powerful discriminator of alternative cosmologies. \citet{wei-etal13} theoretically proved that the WDM, cDE and MG models are hard to be differentiated from one another just by tracing the expansion and growth history with LSS observations. Additionally, it was shown that the presence of a cosmological background of massive neutrinos might significantly suppress the main observational footprints of various cDE and MG models \citep[e.g.,][]{lavacca-etal09,motohashi-etal13,he13}. Such early predictions based on linear observables have been recently confirmed and extended to the non-linear regime of structure formation by \citet{baldi-etal14}, who studied the simultaneous effect on structure formation of a $f(R)$ MG model and of massive neutrinos (the ``$f(R)+\nu$" model, hereafter) by means of large $N$-body simulations. The $f(R)$ gravity is one of the viable and most widely investigated MG theories \citep[see e.g.,][and references therein]{lombriser14} where the Ricci scalar $R$ is replaced by a function $f(R)$ in the Einstein-Hilbert action. With suitable choices of such function, the model can be tuned to match the same expansion history of the standard $\Lambda $CDM cosmology \citep[][]{HS07}. Nonetheless, the derivative of the $f$ function $df/dR$ -- which represents an additional degree of freedom called ``the scalaron" -- is expected to mediate a {\it fifth-force} \citep[see][for a review]{DT10,SF10}, implying that the evolution of density perturbations will be different as compared to $\Lambda $CDM even for an identical expansion history. The viability of the model is ensured by the Chameleon screening mechanism \citep[e.g.][]{KW04,brax-etal08} that allows to recover the behaviour of standard GR in overdense regions of the Universe. In their recent work, \citet{baldi-etal14} noted that for a given $f(R)$ model there is a specific value of the total neutrino mass $\Sigma_{i}m_{\nu _{i}}$ that cancels the effects of the Modified Gravity in most of the standard basic LSS observables, thereby yielding -- besides an identical expansion history -- also the same LSS statistics as $\Lambda$CDM at the present level of observational accuracy. In other words, the free-streaming effect of massive neutrinos \citep{LP06} effectively cancels out that of the fifth-force of the $f(R)$ gravity on the large scale structure. Calling it a ``cosmic degeneracy", \citet{baldi-etal14} regarded this result as an indication of the fundamental limitation of the basic statistics of LSS as a test of alternative cosmologies, concluding that some novel independent statistics or some other independent constraints (as e.g. a laboratory determination of the neutrino mass) are necessary to break the degeneracy. Meanwhile, \citet{SL13} have recently developed a new diagnostic based on the filamentary cosmic web for testing gravity at large scales. Such new approach has shown that the filamentary pattern of the cosmic web is significantly affected by deviations from the standard gravitational behaviour, either in terms of MG or cDE models. More specifically, considering the supercluster spines (i.e. the main stems of the superclusters) as the richest filamentary structures in the cosmic web, \citet{SL13} determined their specific sizes to quantify the degree of the straightness of the superclusters, finding that the specific size distributions of the supercluster spines substantially differ among various cDE models \citep[see again][]{SL13}, implying it being indeed a good indicator of cDE. Similarly, in a subsequent paper \citet{shim-etal14} also investigated the effect of $f(R)$ gravity on the specific sizes of the supercluster spines and showed that the evolution trends of the specific size distributions of the superclusters differ between the cDE and the $f(R)$ gravity models, which indicated that this new diagnostic is in principle capable of breaking the degeneracy between the two models. In the light of the works of \citet{shim-etal14} and \citet{baldi-etal14}, in the present study we aim to explore whether the the degree of the straightness of the superclusters quantified by the specific size of the supercluster spines can be also useful for breaking the cosmic degeneracy between the $\Lambda$CDM and the $f(R)+\nu$ models. In section \ref{sec:data} we will describe how the supercluster samples are obtained from the simulation datasets obtained from the work of \citet{baldi-etal14}. In section \ref{sec:result} we will show how the specific sizes of the supercluster spines are affected by the simultaneous effects of the $f(R)$ gravity and the massive neutrinos. In section \ref{sec:con} we discuss the implications of our results as well as possible prospects for future further improvements. | \label{sec:con} This work has been inspired by two recent publications. On one hand, the work of \citet{shim-etal14} demonstrated that the specific size of the supercluster spine is a powerful diagnostic not only for detecting any $f(R)$ modification of gravity but also for distinguishing the effects of $f(R)$ gravity from that of the presence of interacting Dark Energy. On the other hand, the recent work of \citet{baldi-etal14} raised for the first time a cosmic degeneracy problem consisting in the fact that the basic statistics of LSS (such as e.g. the cluster mass function, the two-point correlation function of the density field and the halo bias) might inherently fail to distinguish a suitable combination of a $f(R)$ MG model and of a specific value of the total mass of neutrinos from the standard $\Lambda$CDM cosmology, since the free streaming effect of massive neutrinos effectively suppresses the extra clustering effect associated with the scalar fifth-force of $f(R)$ gravity. Using the numerical data from the N-body simulations carried out by \citet{baldi-etal14} and the techniques developed by \citet{shim-etal14}, we have determined the mean specific sizes of the supercluster spines at three different redshifts ($z=0,\ 0.3,\ 0.6$) for an $f(R)$ gravity cosmology with different values of the total neutrino mass. Our investigation showed that the neutrino mass required to maximally suppress the effects of the MG fifth force on the specific size distribution of the supercluster spines at $z=0$ is different from the value that determines the maximum degeneracy for the basic statistics of LSS. Furthermore, we have also found that at higher redshifts all the considered combined models show larger differences in the mean specific size of the supercluster spine from the standard GR+$\Lambda$CDM cosmology. Therefore, we conclude that the evolution of the specific size of the supercluster spine can in principle play a crucial role in breaking the cosmic degeneracy highlighted by \citet{baldi-etal14}. Although such conclusion emerges clearly from the present study, a more thorough investigation of the power of the supercluster spines in breaking the MG-massive neutrinos degeneracy will be required in order to quantify precisely the gain of discriminating power obtained by combining this new statistics with the standard LSS probes. First of all, the datasets employed for our analysis were obtained from a suite of intermediate-resolution $N$-body simulations. Since the identification of the superclusters -- especially at high redshifts -- is likely to be affected by the resolution of the original $N$-body realization, a new set of halo catalogs from higher resolution simulations is required to examine more carefully how the mean specific sizes of the supercluster spines evolve with redshifts. Secondly, in the current work we have considered only the case of a rather ``extreme" (although still possibly consistent with standard LSS data for a reasonable value of the neutrino mass) $f(R)$ scenario, namely $f_{R0}=-10^{-4}$. Therefore, it will be necessary to examine whether the specific size of the supercluster spine is capable of breaking the cosmic degeneracy even for the more challenging case of less severe deviations from the standard GR gravity. Second of all, our conclusion is based not on a realistic observational error but only on the statistical significance of the differences in the mean specific size of the superclusters spine among the models. It will be essential to estimate an observational confidence region around the fiducial model for the specific size of the supercluster, as \citet{baldi-etal14} did for the LSS statistics. Another direction in which some improvements are needed is the development of an analytic model for the mean specific size of the supercluster spine. One of the reasons why the basic LSS statistics have been so widely employed as a test for cosmology is that they can be analytically (or semi-analytically) predicted. On the contrary, the mean specific sizes of the supercluster spines have so far been only numerically determined, without being guided by any analytic prescription, which is an obvious drawback that has to be overcome for its practical application in the future. We leave these three main extensions of our analysis for future work. | 14 | 4 | 1404.3639 |
1404 | 1404.1120_arXiv.txt | We report here a series of observations of the interstellar scintillation (ISS) of the double pulsar J0737$-$3039 over the course of 18 months. As in earlier work \citep{col05} the basic phenomenon is the variation in the ISS caused by the changing transverse velocities of each pulsar, the ionized interstellar medium (IISM), and the Earth. The transverse velocity of the binary system can be determined both by VLBI and timing observations. The orbital velocity and inclination is almost completely determined from timing observations, but the direction of the orbital angular momentum is not known. Since the Earth's velocity is known, and can be compared with the orbital velocity by its effect on the timescale of the ISS, we can determine the orientation $\Omega$ of the pulsar orbit with respect to equatorial coordinates ($\Omega = 65\pm2\deg$). We also resolve the ambiguity ($i= 88.7$ or $91.3\deg$) in the inclination of the orbit deduced from the measured Shapiro delay by our estimate $i=88.1\pm0.5\deg$. This relies on analysis of the ISS over both frequency and time and provides a model for the location, anisotropy, turbulence level and transverse phase gradient of the IISM. We find that the IISM can be well-modeled during each observation, typically of a few orbital periods, but its turbulence level and mean velocity vary significantly over the 18 months. | The double pulsar binary system J0737$-$3039 is in a highly relativistic orbit with significant eccentricity \citep{lyne04}. It is an eclipsing binary that is a wonderful laboratory for studies of general relativity \citep{kramer06}. Detailed measurements of the eclipses of A have been used to probe the magnetosphere of the B neutron star \citep{mcl04b, lyutikov} and provide a measurement of geodetic precession \citep{breton08}. The changes in the pulse profiles have been used to explore the precession of the emission beams, the evolution of the orbital system and the dynamics of B's supernova \citep{stairs06, ferdman, per10, per12}. It is also a fine system for study of the interstellar plasma (IISM) because the scattering is dominated by a compact region, the velocities are well determined, and the plasma turbulence can be probed on two neighbouring lines of sight. \citet{ransom04} first reported how the interstellar intensity scintillation (ISS) of the A pulsar exhibits dramatic modulation in timescale over its orbital period (2.45 hr). Following the method proposed by \citet{lyne84} and developed by \citet{ord02} for PSR J1141$-$6545, the authors estimated a rather high center of mass velocity for the double pulsar. Subsequent analysis of the same data showed that the scattering must be anisotropic and inclusion of this effect in the analysis greatly reduced the implied center of mass velocity \citep{col05} -- hereafter Paper 1. In this paper we found correlation between the ISS of pulsars A and B near the time of A's eclipse by B. From the correlation we concluded that the orbital inclination angle was considerably closer to 90$^\circ$ than had been expected on the basis of the original measurements of the Shapiro delay \citep{lyne04}. The main purpose of the observations reported here was to make use of the Earth's orbital velocity to improve, calibrate, and align the earlier scintillation analyses. This would allow us to correct the center of mass velocity for the motion of the Earth, to orient the binary orbit with respect to the celestial reference frame, and to locate the distance of the scattering region. The additional observations were also expected to improve the estimates of the inclination of the orbit, the anisotropy of the IISM and the spatial spectrum of the electron density of the IISM. However, two factors made the original plan of analysis impossible. First the phase at which emission from B is easily detectable drifted away from the time of A's eclipse during the course of the 2004-5 observations. This made measurements of the correlation between the A and B pulsars much less consistent and reliable than had been expected. Second we observed that, although the turbulence in the IISM is homogeneous over several binary orbits, it is not stationary over a year, nor is the velocity of the IISM constant over the year. This phenomenon was also observed by Ord [private communication] when his group attempted to observe the effect of the Earth's orbit on scintillations of PSR J1141$-$6545. Subsequent pulse timing measurements have determined the Shapiro delay and proper motion with greater precision \citep{kramer06}. There have also been long baseline interferometry measurements of parallax and proper motion \citep{del09}. We have altered our analysis of the time scale variations to take advantage of these observations, and we have modeled the entire two dimensional time-frequency correlation function of the scintillations, rather than simply modeled its time scale. These changes have made it (just) possible to obtain a consistent interpretation. This provides the distance to the scattering region and the orientation of the pulsar orbit in celestial coordinates. It also provides an inclination estimate that is consistent with the Shapiro delay. We now realize that the scattering is homogeneous over several binary orbits because the proper motion of the pulsar is low and the binary orbit remains entirely within the ``scattering disc''. Since the measured intensity is a summation over waves that have traveled through all possible paths through the scattering disc, it is quite homogeneous over that area, even if the underlying turbulence is not. However, from month to month, as we repeated the observations, different realizations of the plasma turbulence occupied the scattering disc. The level of turbulence was clearly non-stationary on this time scale. | \label{sec:conclusions} Although we were not able to achieve all the original objectives of this experiment, we have been able to determine the orientation of the orbital plane of the system ($\Omega = 65\pm2\deg$ see Table \ref{tab:inc}) and the inclination of the orbit ($i=88.1\pm0.5\deg$). The inclination of the orbit is consistent with earlier measurements of the Shapiro delay, if the smaller of the two possible Shapiro delay solutions ($88.7\deg$) is taken. Knowing the orientation of the system relative to its proper motion will allow immediate progress on other scientific problems. The first of these is the nature of the supernova that created the B pulsar. Several attempts have been made to determine the magnitude of the kick to the nascent B neutron star (e.g., \citet{piran, willems,stairs06}); these involve tracing the path of the binary system back to possible birthplaces in the Galactic Plane, evolving the orbital eccentricity and semi-major axis to the appropriate age, and determining which set of kick magnitudes and directions is compatible with the known properties of the system. Most of these studies point to a small kick and therefore likely very little tilt of the post-supernova orbital plane relative to the pre-supernova one; this is reinforced by the finding that the spin axis of the A pulsar is close to aligned with the orbital angular momentum \citep{ferdman}. Previous work on the evolution of this system did not use a constraint on $\Omega$, as the anisotropy of the IISM prevented a robust determination of the angle at the time \citep{stairs06} Recent modeling \citep{wong} has found that kicks in the plane of the pre-supernova orbit are preferred over polar kicks; the rather large angle (~60$^{\circ}$) found here between the Line of Nodes and the proper motion may be at odds with this result. Revised modeling incorporating the constraint on $\Omega$ will be presented elsewhere. The determination of $\Omega$, plus the resolution of the sign ambiguity in $\cos i$, permits a strengthening of the double-pulsar test of preferred-frame effects in semi-conservative theories of gravity \citep{wex}. Such effects would produce periodic changes in the longitude of periastron and the eccentricity of the system, in a manner that depends on the orientation of the system relative to the coordinates of the preferred frame. Therefore, the knowledge of the orientation will allow a more precise limit to be set on the parameters of this theory. We model the dominant interstellar scattering plasma as a thin layer located at a distance from the pulsar of $73\pm1$\% of the distance to the Earth. The success of this ``thin screen'' model emphasizes the highly localized distribution of scattering plasma along the line of sight. We also measured its velocity and scattering parameters. The velocity is about 40 km s$^{-1}$ with respect to the Sun, with variations of about 10 km s$^{-1}$, that are comparable with expected Alfv\'{e}n speeds. The level of turbulence varied by a factor of two on a time scale of months, much greater than the statistical variation expected due to refractive effects in a homogeneous Kolmogorov random process. At 5 epochs we measured its anisotropy (axial ratio $1.2 - 1.7$) and phase gradient (due to a transverse gradient in the electron column density). There is some significant variability between epochs in both of these parameters; however the variations are at the level one might expect in different refractive realizations of a homogeneous Kolmogorov random process concentrated in a thin layer. Our results add to the evidence suggesting that the IISM model of homogeneous isotropic Kolmogorov turbulence is no longer adequate. There is accumulating evidence for: anisotropy and intermittency in the turbulence on sub AU scales \citep{ric02,dt03,tunstov2013,Hill,brisken} and for persistent phase gradients \citep{keith}. Evidently this default model of turbulence in the IISM will need to be modified. Apart from the light this throws on the interstellar plasma, turbulence in the IISM is a problem for accurate pulsar timing which is limited in precision by dispersion and scattering \citep{keith}. | 14 | 4 | 1404.1120 |
1404 | 1404.2123_arXiv.txt | We study the stellar haloes of galaxies out to 70-100 kpc as a function of stellar mass and galaxy type by stacking aligned $r$ and $g$ band images from a sample of 45508 galaxies from SDSS DR9 in the redshift range $0.06\,\le\,z\,\le\,0.1$ and in the mass range $10^{10.0} \msun < M_{*} < 10^{11.4} \msun$r. We derive surface brightness profiles to a depth of almost $\mu_r \sim 32 \,\mathrm{mag\,arcsec}^{-2}$. We find that the ellipticity of the stellar halo is a function of galaxy stellar mass and that the haloes of high concentration ($C > 2.6$) galaxies are more elliptical than those of low concentration ($C < 2.6$) galaxies. The $g$-$r$ colour profile of high concentration galaxies reveals that the $g$-$r$ colour of the stellar population in the stellar halo is bluer than in the main galaxy, and the colour of the stellar halo is redder for higher mass galaxies. We further demonstrate that the full two-dimensional surface intensity distribution of our galaxy stacks can only be fit through multi-component S\'{e}rsic models. Double-S\'{e}rsic profiles adequately model the average surface brightness distributions of high concentration galaxies, while triple-S\'{e}rsic profiles are often needed to model the surface brightness distributions of low concentration galaxies. Using the fraction of light in the outer component of the models as a proxy for the fraction of accreted stellar light, we show that this fraction is a function of stellar mass and galaxy type. For high concentration galaxies, the fraction of accreted stellar light rises from $30\%$ to $70\%$ for galaxies in the stellar mass range from $10^{10.0} \msun$ to $10^{11.4} \msun$. The fraction of accreted light is much smaller in low concentration systems, increasing from $2\%$ to $25\%$ over the same mass range. This work provides important constraints for the theoretical understanding of the formation of stellar haloes of galaxies. | Traditionally, galaxies have been studied through their surface brightness profiles \citep{Hubble,deVauc}. This has not only revealed a wealth of information about their different morphologies but also hints about their formation processes. De Vaucouleurs (1948) first characterised the surface brightness profiles of giant elliptical galaxies as a simple $\log I(R) \propto R^{1/4}$ law, which was later also found to fit the bulges of disk galaxies. On the other hand, the disks of spiral galaxies have been traditionally fit with exponential profiles \citep{Freeman}. \cite{Sersic} showed that all these profiles are specific cases of a more general $\log I(R) \propto R^{1/n}$ function, which fits the surface brightness profile of a large number of galaxies from disks to spheroidals, dwarfs, ellipticals and bulges. The shape of the surface brightness profile provides valuable clues about the way in which different galaxies formed. As deeper and more resolved surface brightness data became available, deviations from these simple laws became clearly evident, indicating that galaxy formation was a more complex process than previously believed \citep{Kormendy}. This discovery motivated the use of multiple components to model the surface brightness profiles of galaxies \citep{Kormendy77, Simard, Lackner} Bulge-disk decompositions helped distinguish pseudo-bulges ($n\sim 1$) from classical bulges ($n \sim 4$). Pseudo-bulges are dense central components of disk galaxies that are flattened and rotationally supported and believed to be built out of disk gas. Classical bulges lie on the fundamental plane linking galaxy size, luminosity and velocity dispersion \citep{Bender}. With the advent of deeper imaging (through Hubble Space Telescope and medium-sized, ground-based telescopes), it has become possible to detect additional fainter stellar structures around both normal galaxies and brightest cluster galaxies \citep{Schweizer_a, Malin, Schweizer_b} Today, stellar haloes of galaxies have been observed and confirmed not only in clusters as intracluster light (ICL), but also in a large variety of field galaxies from early-type to late-type spirals. This is consistent with the idea that the faint stellar halo is built up from the debris of smaller galaxies and satellites that are tidally disrupted (e.g. \citealt{Bullock} and \citealt{Cooper10}). In the Milky Way and in other nearby disk galaxies, the stellar halo and other tidal features have been directly detected through star counts \citep{Bell2,Ibata,Monachesi}. Observing the stellar halo through star counts is limited to the Local Universe. The integrated light from deep imaging has enabled studies of the stellar haloes of more distant elliptical and spiral galaxies (see e.g. \citealt{mihos}, \citealt{delgado}, \citealt{Tal2} and \citealt{dokkum}). By using modest aperture telescopes \citep{delgado} with innovative telescope design optimised for low surface brightness emission \citep{dokkum}, one can reduce the systematic errors related to flat fielding and the complex point spread function (PSFs) of the telescope and reach much deeper limiting depths of $\mu_g \sim 32 \,\mathrm{mag\,arcsec}^{-2}$. Alternatively, stacking the images of a large number of similar galaxies (e.g. \citealt{Zibetti04}, \citealt{Zibetti05}, \citealt{Tal} and \citealt{Cooper}) enables one to study the average stellar haloes of statistical samples of more distant galaxies. The disadvantage is that information on detailed structure is lost. \cite{Zibetti04} used stacking techniques to study the stellar haloes of edge-on disk galaxies; \cite{Tal} studied the stellar haloes of luminous red galaxies out to $z \sim 0.34$. Theoretical models \citep{Cooper,Purcell,Oser,Lackner2} predict not only large variations between individual stellar haloes of galaxies, but also systematic variations in the average properties of stellar haloes as a function of certain galaxy parameters (for example, halo mass, stellar mass, galaxy bulge-to-disk ratio, etc). In order to constrain theoretical models for the formation of stellar haloes, it is important to study the average properties of the surface brightness profiles of galaxies as a function of these galaxy parameters. In this paper, we stack a large number of galaxy images and study them as a function of stellar mass and galaxy type (late-type or early-type). The SDSS imaging data set is well-suited to study the faint stellar haloes of galaxies \citep{Zibetti04,Zibetti05,Tal}. The systematics of stacking many SDSS images to produce a very deep image have been well understood and quantified. This is important because studying low-surface brightness structures is highly dependent on a proper estimation and removal of the sky background. We pay particular attention to the residual sky background obtained after stacking the sky-subtracted images from SDSS DR9. We then model the surface brightness profile of the stacked galaxy including the stellar halo through multi-component fits. We then parametrise the contribution of the stellar halo by deriving the fraction of light in the outer component of the galaxy. In Section \ref{sec:selection}, we describe how we select and prepare our galaxy images for stacking. In Section \ref{sec:stacking}, we describe in detail the stacking procedure, our error analysis, PSF analysis and the methodology we employ to derive the ellipticity, surface brightness and the colour profiles for each galaxy stack. In Section \ref{sec:analysis}, we present the surface brightness and colour as a function of the stellar mass of the galaxy and of galaxy type. In Section \ref{sec:modelling}, we fit models to these surface brightness profiles and determine the fraction of light in the outer faint stellar component. In Section \ref{sec:summary}, we summarise and in Section \ref{sec:discussion}, we discuss our results in light of our theoretical understanding of the formation of stellar haloes of galaxies. Throughout this paper, we assume a flat $\Lambda$CDM cosmology, $\Omega_{\mathrm{m}}=0.25$, $\Omega_{\mathrm{\Lambda}}=0.75$ and Hubble parameter $h=0.73$. | \label{sec:summary} In this work, we have shown that stacking $g$ and $r$ band mosaics of similar galaxies allows us to derive reliable surface brightness profiles upto a depth of $\mu_r \sim 32 \,\mathrm{mag\,arcsec}^{-2}$. We study surface brightness, ellipticity and $g$-$r$ colour profiles as a function of stellar mass and galaxy type. We perform fits to the stacked images using multi-component S\'{e}rsic models. This enables us to estimate the fraction of the stellar light/mass in the outermost component, which we hypothesize to be built up from accreted stellar material, and to set constraints on theories for the formation of stellar haloes through hierarchical merging. The main results of this paper can be summarized as follows. \begin{enumerate} \item The fraction of accreted stellar material increases with stellar mass. At fixed mass, the fraction of accreted material is higher in early-type than in late-type galaxies. \item The stellar haloes of high concentration galaxies ($C>2.6$) tend to be more elliptical than those of low concentration galaxies ($C<2.6$). The ellipticity of the outer stellar halo increases strongly with stellar mass for high concentration galaxies, and more weakly with stellar mass for low concentration galaxies. \item Because we stack galaxies that are nearly face-on, we are only able to probe the colour of the outer accreted component in high concentration galaxies. In these systems, the $g$-$r$ colour of the outer halo light is bluer than than the centre of the galaxy and is an increasing function of stellar mass. \item We find that a single-S\'{e}rsic profile cannot fit the entire two-dimensional surface brightness distribution of any of our stacked images . Multi-component models are needed to model the excess light in the outer parts of the galaxy, especially between $\mu_r \sim 28 - 32 \,\mathrm{mag\,arcsec}^{-2}$, and to account for the radial dependence of the ellipticity of the light distribution. \item Double-S\'{e}rsic profiles adequately model the surface brightness distributions of high concentration galaxies ($C>2.6$), while triple-S\'{e}rsic profiles are often needed to model the surface brightness profile of low concentration galaxies ($C<2.6$). \item Using the fraction of light in the outer component of our models as a measure of the fraction of the total stellar mass composed of accreted stellar material, we find that this fraction is an increasing function of stellar mass. At fixed stellar mass, it is also a function of concentration. For high concentration galaxies, the fraction of accreted stellar light rises from $30\%$ to $70\%$, while for low concentration galaxies the fraction of stellar light rises from $2\%$ to $25\%$ for galaxies in the stellar mass range $10^{10.0} \msun$ to $10^{11.4} \msun$. \end{enumerate} We have attempted to characterise the stellar halo of galaxies through modelling their surface brightness. It is the depth, the large dynamic range and the two-dimensional shape information (ellipticity) of our surface brightness profiles which enables us to recognise deviations from a single component profile and to model successfully the stellar halo of our galaxy stacks out to 100 kpc with two or three components. An important outcome is that a single S\'{e}rsic component cannot fit the surface brightness profiles of high concentration galaxies over a large dynamic range in radius and surface brightness, but can only fit the inner parts of galaxies. The inability of a single S\'{e}rsic to fit the two-dimensional surface brightness profile of galaxies has also been confirmed by the studies of \cite{Bernardi}, \cite{Simard} and \cite{Lackner}. Multi-component models are needed to model the full two-dimensional surface brightness profiles of galaxies. We have demonstrated that it is both the average shape of the surface brightness profile and the radial variation in ellipticity of the light in a galaxy stacks that constrain such models. For high concentration galaxies, the effective radius of the outer component is twice as large as the effective radius of the inner component. For low concentration galaxies, the effective radius of the outer component is much larger than the inner components. For high concentration galaxies, the luminosity of the outer component is a significant fraction of the total luminosity of the galaxy and ranges from $30\%$ to $70\%$. It also dominates over a large radial range of the galaxy. On the other hand, in low concentration galaxies, the outer component occupies a smaller fraction (from $2\%$ to $25\%$) and is only dominant at radii larger than $20-30$ kpc. In both cases, the fraction of light in the outer component increases with stellar mass (see the red line in the top plots of Figure \ref{fig:light-fraction_lc} and Figure \ref{fig:light-fraction_uc}). We propose in this work that the fraction of light in the outer component provides a measure of the amount of accreted stellar light in the galaxy. While a direct one-to-one correspondence between the fraction of light in the outer component and the fraction of accreted stellar light cannot be directly proven, the trends in the fraction of light in the outer component agree qualitatively with the trends of the accreted light fraction as a function of mass and galaxy-type in the particle-tagging models of \citep{Cooper}. Interestingly, the rate of increase of accreted stellar mass increases dramatically above $M_{*} \sim 10^{10.6} \msun$. Interestingly, this corresponds to the stellar mass where galaxies transition from blue/star-forming to red/passive systems \citep{Kauffmann03b}. A significant jump in the accreted mass fraction may be most simply explained by \emph{in-situ} growth of the galaxy being terminated by feedback processes, such as energy injection from relativistic jets produced by black holes in massive galaxies \citep{Croton}. In the two stage model of massive galaxy formation proposed by \cite{Oser}, an early, rapid \emph{in-situ} star formation period is followed by a late merger-dominated period. In the later phase, galaxies tend to grow predominately through minor mergers. We note that the particle tagging models of \cite{Cooper} are directly tied to semi-analytic models that include AGN feedback prescription, and thus also include quenching of \emph{in-situ} growth of galaxies through cooling and star formation. In future work, we intend to undertake a detailed comparison with these models. Measuring the ellipticity of the outer stellar halo of galaxy also provides us with hints about the formation processes for the stellar halo. A high ellipticity is likely to imply that satellite systems are preferentially accreted along the major axis of the main galaxy \citep{Tal}. The variance in the outer stellar halo profile between different galaxies can be predicted from our surface brightness profiles. This variance results from the fact that similar galaxies can have stellar haloes with very different masses, sizes and shapes. The physical origin of this variance as predicted by the $\Lambda$CDM models, is that galaxies of the same mass have had a range of merger histories, resulting in different accreted stellar mass fractions. This has also been clearly demonstrated using particle-tagging techniques on the Aquarius haloes \citep{Cooper10}, which show very large halo-to-halo differences. We also note that the integrated surface brightness of the galaxy, including the stellar halo, includes considerably more light that measured by the SDSS \texttt{model} and \texttt{cModel} magnitudes. For example, for high concentration galaxies in the stellar mass range $10^{11.0} \msun < M_{*} < 10^{11.4} \msun$, there is about 50\% more light contained in the stellar halo at surface brightnesses greater than $\mu_r \sim 24.5 \,\mathrm{mag\,arcsec}^{-2}$. This implies that there is considerably more stellar material in the galaxy that one might infer from the SDSS photometry. The stellar masses defined by the MPA-JHU catalogue and used in this work are only used to define the stellar mass bins, and are systematically less than the true stellar mass of the galaxy. This will also be the subject of future work. | 14 | 4 | 1404.2123 |
1404 | 1404.2912_arXiv.txt | If the recent measurement of B-mode polarization by BICEP2 is due to primordial gravitational waves, it implies that inflation was driven by energy densities at the GUT scale $M_{GUT} \sim 2\times 10^{16} GeV$. This favors single-field chaotic inflation models. These models require transplanckian excursions of the inflaton, forcing one to address the UV completion of the theory. We use a benchmark 4d effective field theory of axion-4-form inflation to argue that inflation driven by a quadratic potential (with small corrections) is well motivated in the context of high-scale string theory models; that it presents an interesting incitement for string model building; and the dynamics of the UV completion can have observable consequences. \vskip .2cm \noindent{BRX-TH-676} | The detection of B-mode polarization by the BICEP2 experiment \cite{Ade:2014xna,Ade:2014gua}\ gives tantalizing evidence that quantum gravity may be directly relevant for observational cosmology. If the future checks confirm that the BICEP2 B-modes are generated by primordial gravitational waves, this will give support to the idea that these waves are induced by quantum fluctuations of the graviton. Furthermore, simple chaotic inflation models such as $V = \frac{1}{2} m^2 \phi^2$ \cite{Linde:1983gd} are in excellent current agreement with the data. To generate the observed density fluctuations in the CMB and large scale structure, chaotic inflation models occur at energy densities $\la M_{GUT}^4 \sim (2\times 10^{16}\ GeV)^4$. Any models generating observable primordial gravitational radiation require transplanckian excursions of the inflaton in field space of order $\Delta \phi \sim 10 m_{pl}$ \cite{Lyth:1996im, Efstathiou:2005tq}\ over the course of inflation.\footnote{We use the reduced Planck mass $m_{pl} = 2.4 \times 10^{18}\ GeV$; we take $M_{GUT}$ to be the value suggested by supersymetric coupling unification \cite{Dimopoulos:1981yj,Amaldi:1991cn,Langacker:1992rq}.} Therefore all such models are sensitive to the UV completion at the Planck scale, requiring a good understanding of quantum gravity. The purpose of this paper is to consider constraints on the UV completion of models dominated by a (possibly distorted) quadratic potential, and present them as an incitement for string model builders. Chaotic inflation is stable against perturbative quantum corrections, as exemplified as early as \cite{Smolin:1979ca,Linde:1987yb}, in response to the concerns about irrelevant operator contributions raised in \cite{Enqvist:1986vd}. This in fact follows if one protects the theory with (approximate) shift symmetries, so that all couplings to the inflaton are either very weak or via the derivatives of the inflaton. However, nonperturbative quantum-gravitational effects are expected to break such global symmetries, inducing ${\cal O}(1)$ coefficients for all Planck-suppressed operators. When $\phi$ ranges over super-Planckian distances, these operators may spoil the small curvature of the inflaton potential required for slow roll inflation. Another concern is that in a string compactification, the degrees of freedom can shift substantially as fields move over super-Planckian distances; this can also manifest itself in dangerous Planck-suppressed operators. A good UV-complete realization must control these operators.\footnote{For a more complete review and discussion of the issues in this paragraph, see \cite{Kaloper:2011jz}.} \begin{figure} \includegraphics[scale=.75]{monopot.pdf} \caption{Energies as a function of $\phi$, for the potential $V = \frac12 (\mu \phi + q)^2$. The picture repeats itself each time one shifts $\phi \to \phi + e^2/\mu \equiv \phi + f$.} \end{figure} Another technically natural high-scale model takes as the inflaton a periodic pseudoscalar (aka an axion) for which the potential is generated by instantons \cite{Freese:1990rb,ArkaniHamed:2003wu,ArkaniHamed:2003mz}: the topology of field space protects the shift symmetry of the inflaton. The usual potential is $V \sim \Lambda^4 \cos(\phi/f)$. However, high-scale slow roll inflation requires axion decay constants $f > \phi> m_{pl}$. Gravitational instantons, such as wormholes \cite{Kallosh:1995hi}, and string theoretic instantons \cite{Banks:2003sx}\ typically have actions of order $S \sim (\frac{m_{pl}}{f})^p$; thus higher-order instantons can spoil slow-roll inflation \cite{Banks:2003sx,ArkaniHamed:2006dz}. Axion monodromy models control the Planck-suppressed operators by combining chaotic and natural inflation. In this scenario, the inflaton is a compact scalar with periodicity $f < m_{pl}$. The presence of fluxes or branes ``unwraps" the inflaton configuration space \cite{Silverstein:2008sg,McAllister:2008hb,Kaloper:2008fb,Kaloper:2011jz,Dubovsky:2011tu}, leading to a multivalued potential as in Figure 1. In what follows we will focus on the axion-4-form models of \cite{Kaloper:2008fb,Kaloper:2011jz}, which we dub ``natural chaotic inflation", as a benchmark theory from which to discuss the general phenomenon of axion monodromy. We start with the Lagrangian density \be {\cal L} = \frac12 m_{pl}^2 R - \frac{1}{2} (\p\phi)^2 - \frac{1}{48} F_{(4)}^2 + \frac{\mu}{24} \phi {}^*F_{(4)}\, , \label{eq:axionfour} \ee where $F_{(4)} = dA_{(3)}$ is a four-form field strength, with $A_{(3)}$ a three-form gauge field. The canonical momentum $p_{A_{123}} = {}^*F_{(4)} - \mu\phi$ is quantized in units of the electric charge $e^2$; using this, one may write the Hamiltonian as \be H = \frac{1}{2} \left(p_\phi\right)^2 + \frac{1}{2} \left(\nabla \phi\right)^2 + \frac{1}{2}\left(n e^2 + \mu \phi\right)^2 \, .\label{eq:afham} \ee If $\phi \to \phi + f$, we must have $\mu f = e^2$ for consistency. The quantum number $n$ can be shifted by the nucleation of membranes, leading to the multivalued potential for $\phi$ shown in Figure 1. At tree level, if membrane nucleation is suppressed, one has a model of chaotic inflation with a quadratic potential. The fundamental periodicity of the scalar implies that corrections to $V = \frac{1}{2} \mu^2 \phi^2$ take the form $(V/M_{uv}^4)^n$, which can be small \cite{Kaloper:2008fb,Kaloper:2011jz}. A danger to the model remains: couplings such as $\mu$ can depend on moduli with Planck-suppressed couplings. If the moduli masses are of the order or smaller than the Hubble scale $H$, they potentially destabilize inflation. The axion monodromy models of \cite{Silverstein:2008sg,McAllister:2008hb} turn these bugs into features, by providing a UV completion in which the overall effect of higher-order corrections and couplings to moduli is to flatten the tree-level potential \cite{Silverstein:2008sg,McAllister:2008hb,Dong:2010in}. The constructions in those papers give potentials significantly flatter than quadratic while still being large-field models, with significant tensor-to-scalar ratio $r$. These may be consistent with the BICEP2 result, given the current statistical significance of the results. Further, additional UV complete models realizing the monodromy mechanism starting from steeper potentials with larger $r$ are under development \cite{monopowers}. Here we will focus on whether a quadratic potential generated by integrating the 4-form starting with (\ref{eq:axionfour}) can be realized with only small corrections, using effective field theory as a guide. Our motivation for selecting this particular class of models is the relative simplicity of the dynamics which follows; yet we will find that there is a range of interesting potential signatures. We will discuss the relevant dynamical scales arising in various string theory scenarios, to argue that existing models consistent with grand unification are at the edge of being viable in this regard. If viable, they can give a finite probability for nonperturbative transitions to occur in early epochs of inflation \cite{Kaloper:2011jz}, and can give observable corrections to the tree-level quadratic potential. In particular, we will show that the natural consequence of a single UV scale $M_{uv} \gtrsim M_{GUT}$ close to the GUT scale yields potentially observable corrections to the scalar and tensor spectrum. We present this as motivation for further work on high-scale string models. | The observation of primordial gravitational waves by BICEP2 has ushered a new era in cosmology. The results are pointing very strongly to the large field slow roll models of inflation \cite{Linde:1983gd} as the dynamics that shaped the universe (although a lively discourse on this is still ongoing, \cite{Collins:2014yua,Contaldi:2014zua,Kehagias:2014wza,Dent:2014rga,Freese:2014nla,Czerny:2014qqa,Lyth:2014yya,DiBari:2014oja,Creminelli:2014oaa,Ashoorioon:2014nta,Ho:2014xza}). This presents a constructive incitement to model building, in particular in string theory, because large field models are sensitive to the details of the UV completion in which they should be embedded. In this work we have addressed the UV sensitivity of the specific class of models described in \cite{Kaloper:2008fb,Kaloper:2011jz}, which lead to quadratic potentials and are in excellent agreement with the current cosmological data\footnote{After the appearance of the first version of this work, several papers have appeared with very similar constructions of effective large field models of inflation from string theory \cite{Marchesano:2014mla,Hebecker:2014eua,Blumenhagen:2014gta,Harigaya:2014eta,Grimm:2014vva,Ibanez:2014kia}.}. Since such models operate at the GUT scale, and string theory consistent with grand unification have fundamental scales close to the GUT scale, the UV physics may leave imprints in the sky in the form of small corrections to the leading order CMB observables. These signatures can at least be used to constrain the UV physics. Even more interestingly, future observational tests could use these signatures as benchmarks to search for, in order to further test large field inflation. We feel the discussion here only reinforces the conclusion of \cite{Kaloper:2011jz}: that this class of chaotic inflation models lives at the edge of respectability, which is often the most interesting place to be. \vskip.5cm {\bf Acknowledgments}: We would like to thank Guido D'Amico, Matt Kleban, Andrei Linde, Fernando Quevedo, Eva Silverstein, Martin Sloth, Lorenzo Sorbo and Alexander Westphal for very useful discussions. AL would like to thank the Aspen Center for Physics for its hospitality during the ``New Perspectives on Thermalization" winter conference, at the onset of this work. NK is supported by the DOE Grant DE-FG03-91ER40674. AL is supported by the DOE grant DE-SC0009987. | 14 | 4 | 1404.2912 |
1404 | 1404.5703_arXiv.txt | Here we present the analysis of multifrequency data gathered for the Fanaroff-Riley type-II (FR\,II) radio galaxy \rg, hosted in the brightest cluster galaxy in the center of \cl. The galaxy harbors one of the most massive black holes known to date and our analysis of the acquired optical data reveals that this black hole is only weakly active, with a mass accretion rate $\dot{M}_{\rm acc} \sim 2 \times 10^{-4} \, \dot{M}_{\rm Edd} \sim 0.02 \, M_{\odot}$\,yr$^{-1}$. Based on analysis of new \cxo\ and \xmm\ X-ray observations and archival radio data, and assuming the well-established model for the evolution of FR II radio galaxies, we derive the preferred range for the jet kinetic luminosity $L_{\rm j} \sim (1-6) \times 10^{-3}\, L_{\rm Edd} \sim (0.5-3) \times 10^{45}$\,\Lu. This is above the values implied by various scaling relations proposed for radio sources in galaxy clusters, being instead very close to the maximum jet power allowed for the given accretion rate. We also constrain the radio source lifetime as $\tau_{\rm j} \sim 40-70$\,Myrs, meaning the total amount of deposited jet energy $E_{\rm tot} \sim (2-8) \times 10^{60}$\,ergs. We argue that approximately half of this energy goes into shock-heating of the surrounding thermal gas, and the remaining $50\%$ is deposited into the internal energy of the jet cavity. The detailed analysis of the X-ray data provides indication for the presence of a bow-shock driven by the expanding radio lobes into the \cl\ cluster environment. We derive the corresponding shock Mach number in the range $\mathcal{M}_{sh} \sim 2-4$, which is one of the highest claimed for clusters or groups of galaxies. This, together with the recently growing evidence that powerful FR\,II radio galaxies may not be uncommon in the centers of clusters at higher redshifts, supports the idea that jet-induced shock heating may indeed play an important role in shaping the properties of clusters, galaxy groups, and galaxies in formation. In this context, we speculate on a possible bias against detecting stronger jet-driven shocks in poorer environments, resulting from an inefficient electron heating at the shock front, combined with a relatively long electron-ion temperature equilibration timescale. | \label{intro} One of the most widely debated problems in modern astrophysics is related to heating of the hot gaseous environment in galaxy clusters, which is required by the observed temperature and entropy profiles of the intracluster medium (ICM). This problem is manifested by the apparent lack of observed low-temperature gas ($< T_{\rm vir}/3$, where $T_{\rm vir}$ is the virial temperature) falling at very high rates ($\sim 100-1000\,M_{\odot}$\,yr$^{-1}$) onto the centers of the so-called ``cooling-flow'' systems \citep[see][and references therein]{fab94,pet06}. Although the debate regarding the exact physical processes at work is in this context still ongoing, mechanical heating by relativistic jets and lobes expanding from active galactic nuclei (AGN) of giant ellipticals located in cluster centers (hereafter `brightest cluster galaxies', BCGs), is considered as one of the most promising scenarios \citep[see, e.g., the topical reviews by][and references therein]{mcn07,mcn12}. This idea is supported by the finding that BCGs are indeed typically radio-loud \citep[e.g.,][]{bur90,bes05,crof07}, and that the jets and lobes produced in those systems are, on average, capable of suppressing the dramatic ICM cooling as far as the total available energy release is considered \citep[e.g.,][but see the discussion below]{bir04,raf06,dun08}. The impact of large-scale jets and lobes on their environment is not restricted to quenching of cooling flows in clusters. Radio-loud AGN are now understood to play a crucial role also in altering the gas properties in galaxy groups (where, in fact, most of the galaxies in the Universe reside) and isolated massive ellipticals \citep[see][for a review]{mat03,sun12}, or even to suppress star-formation processes in the interstellar medium (ISM) of systems in formation (particularly at high redshifts), influencing in this way the co-evolution of galaxies and central supermassive black holes (SMBHs) in a complex feedback loop \citep{bow06,cat06,crot06}. Hence, understanding the energetics of AGN jets and lobes, as well as how exactly they interact with the ambient medium, is crucial for understanding the structure formation in the Universe in general. The first models of the ICM heating by AGN jets involved energy dissipation at strong shocks driven in the ambient medium by the expanding jet cocoons, i.e., lobes \citep{cla97,hei98,kai99}. Such shocks were naturally expected in the framework of widely accepted evolutionary scenarios developed for `classical doubles' \citep{sch74,beg89,kai97}. However, while the cavities in the X-ray emitting cluster gas at the positions of radio lobes of BCG-hosted AGN were being discovered already in \ros\ data (see \citealt{boe93} for Perseus\,A, \citealt{car94} for Cygnus\,A, or \citealt{hua98} for Abell\,4059), indicating clearly that the expanding jets do displace the ICM out of the cluster central regions, no radiative signatures for the expected shocks heating up the X-ray emitting gas around the boundaries of the radio lobes were found even in the first high-resolution \cxo\ observations of systems like Hydra\,A \citep{mcn00}, Perseus\,A \citep{fab00}, Abell\,2052 \citep{bla01}, RBS\,797 \citep{sch01}, or Abell\,4059 \citep{hei02}. This absence of evidence for strong shocks at the lobes' edges led to the conclusion that the expected supersonic expansion of jet cocoons takes place only in the earliest phases of the source evolution (basically only when the radio structure is still confined within the ISM of the host), and is quickly replaced by a transonic expansion as soon as the lobes reach the ICM scale \citep{rey01}. This transonic expansion of the lobes, combined with the anticipated duration of the jet lifetime/jet duty cycle of the order of $\sim 10-100$\,Myr, which is much shorter than the cluster lifetime, led next to the idea that, after switching off the jet activity, the lobes detach from the central AGN and rise buoyantly in the stratified cluster atmosphere in the form of low-density `bubbles.' The rising bubbles uplift the cool gas from the center, converting their enthalpy into the gas kinetic energy, which may next be thermalized \emph{somehow} in the bubbles' wake \citep{chu01,chu02}. This scenario seemed supported by the discovery of `ghost cavities' \citep[e.g.,][]{fab00}, and also by numerical simulations of the jet evolution in rich cluster environment (e.g., \citealt{bru02,rey02,rus04}; although one should note that the stability of such structures against various types of magneto-hydrodynamic instabilities, which depends crucially on the hardly known ICM magnetic field configuration, is an open issue, e.g., \citealt{jon05,die08,one09}). Later, however, the evidence for the presence of \emph{weak} shocks driven by expanding lobes in the gaseous atmospheres of clusters, galaxy groups, and isolated systems, started to emerge in deep and very deep \cxo\ and \xmm\ exposures (see \S\,\ref{sh} below and references therein), and placed increased attention back again on the possibility of shock heating of the ICM \citep{voi05}. Yet highly over-pressured jet cocoons expanding supersonically are expected in luminous classical doubles, i.e. Fanaroff-Riley type II sources (FR\,IIs), rather than in low-power Fanaroff-Riley type I radio galaxies (FR\,Is) typically found in the centers of rich clusters \citep[see in this context][]{bru07,guo10,per11,har13}. In fact, a number of observational studies indicated that FR\,IIs at low redshifts avoid dense cluster environment, and that there is a clear difference in richnesses of clusters harboring FR\,I and FR\,IIs central radio galaxies \emph{on average}, albeit with a substantial dispersion \citep[e.g.,][]{lon79,pre88,all93,wan96,zir97,har02,sle08,win11}. There are however strong indications that this may not hold in the higher-$z$ Universe, where luminous FR\,IIs are found also in richer systems, equivalent to Abell class I or higher \citep{yat89,hil91,sie05,bel07,ant12}. According to \citet{fan74}, radio galaxies with total spectral luminosity density exceeding $P_{\rm 1.4\,GHz} \sim 10^{25}$\,W\,Hz$^{-1}$ are characterized almost exclusively by a `classical-double', edge-brightened morphology, consisting of collimated jets terminating in well-localized bright hotspots, and surrounded by prominent lobes. On the other hand, the overwhelming majority of sources with 1.4\,GHz spectral power below this critical value possess edge-darkened large-scale radio structures with less-collimated jets lacking any clear termination points, but instead extend in a plume-like fashion to larger distances, or form amorphous bubbles of radio-emitting plasma. These two types of objects are referred to as the aforementioned FR\,II and FR\,I type radio galaxies, respectively. \citet{led96} later argued that the FR\,II/FR\,I division depends not only on the total radio power of a source, but also on the properties of the galactic hosts, and that the borderline radio luminosity in particular scales with the host optical luminosity as $\propto L_{\rm opt}^{1.8}$. Later investigations showed however that the \citeauthor{led96} scaling is not an absolute separation, as there is a substantial scatter of FR\,IIs around the proposed borderline \citep[e.g.,][]{lin10,koz11,win11}. Still, even only a rough positive dependance of the jet radiative luminosity on the host luminosity is a strong indication that the difference in large-scale radio morphology of a source is due to a combination of the jet kinetic power (scaling somehow with the lobes' radio luminosity) and the properties of the ambient medium (on the galactic scale), rather than due to the jet power alone \citep[see in this context also][]{gop00,kai07,kaw09}. Namely, for a given pressure and density of the ambient medium, the jet has to be powerful enough to pierce through the surrounding gas and to form a strong termination shock (observed as a hotspot) at the tip of an outflow, through which the jet plasma passes by, back-flowing into an over-pressured cocoon/lobe surrounding the jet; the lobe expands next sideways with supersonic velocity driving weaker bow-shock all along its edges. We note that the FR\,II and FR\,I populations seem to differ also in the AGN emission line properties whereby sources with prominent emission lines (`high-excitation radio galaxies'; HERGs) are typically associated with FR\,II-type radio structures, while majority of FR\,I-type sources possess nuclei with no strong emission lines \citep[`low-excitation radio galaxies'; LERGs; see][]{lai94}. Again, the division is not a strict one, since many FR\,IIs are classified as LERGs, and some FR\,Is are found to have HERG-like nuclei \citep[e.g.,][]{har07,but10,bes12,gen13,ine13}. Despite this caveat, the established correspondence is in general consistent with the idea that the production of powerful jets --- like the ones found in FR\,IIs --- requires higher accretion rates (evidenced by prominent AGN emission lines), with the commonly anticipated borderline at the accretion luminosities in Eddington units, $L_{\rm acc}/L_{\rm Edd} > 0.01$ \citep[see][]{ghi01}. This issue is then particularly interesting in the context of BCG-hosted jets, which are often assumed to be powered only by a hot rarefied gas accreting at a limited rate \citep{all06,mer07,mcn11,rus13}, even though large amounts of cold gas capable of an efficient fueling of central SMBHs are now being routinely found in many giant ellipticals at the centers of evolved clusters \citep[see, e.g.,][]{don11}. Studying classical doubles in rich cluster environment is therefore important for several reasons, and in particular for understanding (i) the AGN activity and jet duty cycle in the evolved systems, (ii) the evolution of relativistic outflows and their interactions with the ambient medium, as well as (iii) the mechanical heating of the ICM by the expanding AGN jets and lobes. Thus motivated, here we present the analysis of multifrequency data for the FR\,II source \rg\ located in the center of the \cl\ cluster. The paper is organized as follows. In \S\,\ref{data} we describe the target selection, and present the analysis of the arc-second resolution radio maps from the Very Large Array archive, along with the newly acquired optical spectra from the William Herschel Telescope, as well as \cxo\ and \xmm\ X-ray data. In \S\,\ref{results} we discuss the main results of the data analysis for the \rg/\cl\ system. In \S\,\ref{discuss} we summarize our findings in a broader context of the `radio-mode' feedback operating in galaxy clusters. We assumed a standard cosmology with $H_0=71$\,km\,s$^{-1}$\,Mpc$^{-1}$, $\Omega_m=0.27$, and $\Omega_{\Lambda}=0.73$, so that the redshift of the target $z = 0.0363$ corresponds to the luminosity distance of 158\,Mpc and the conversion scale of $0.71$\,kpc/\as. | \label{discuss} In this paper we analyze in detail the gathered radio (\vla), optical (\wht), and X-ray (\xmm\ and \cxo) data for the radio galaxy \rg\ located in the center of the \cl\ cluster. We find that \rg\ appears under-luminous ($L_{\rm 1.4\,GHz} < 10^{41}$\,\Lu) considering its large-scale FR\,II morphology, as well as the starlight luminosity of its host, thus challenging the strictness of the Fanaroff-Riley and Ledlow-Owen divisions in classifying extragalactic radio sources. The host of \rg\ is the brightest cluster galaxy in \cl, harboring one of the most massive black holes known to date, $M_{\rm BH} \simeq 4 \times 10^9 \, M_{\odot}$. Our analysis of the obtained \wht\ data reveals that this black hole is only weakly active, with the corresponding LINER-type nuclear luminosity of $L_{\rm nuc} \sim 4 \times 10^{-5} \, L_{\rm Edd} \sim 2 \times 10^{43}$\,\Lu. Following the results of numerical simulations by \citet{sch07}, we assume an efficient electron heating in the accretion disk in the system; this leads to a relatively high ($\sim 1\%$) radiative efficiency of the disk despite its low accretion rate, estimated here as $\dot{M}_{\rm acc} \sim 2 \times 10^{-4} \, \dot{M}_{\rm Edd} \sim 0.02 \, M_{\odot}$\,yr$^{-1}$ with a corresponding accretion luminosity $L_{\rm acc} \sim 2 \times 10^{-3}\,L_{\rm Edd} \sim 10^{45}$\,\Lu. The derived $\dot{M}_{\rm acc}$ is likely lower than the Bondi value, which we could only restrict to a relatively wide range ($0.03-2.5 \, M_{\odot}$\,yr$^{-1}$\,$\ll \dot{M}_{\rm Bon} \ll 2.5 \, M_{\odot}$\,yr$^{-1}$) because the accretion radius is not resolved in our \cxo\ observation. Based on the collected radio and X-ray data for \rg, along with the well-established model for the evolution of FR\,II radio galaxies, we derive the preferred range for the jet kinetic luminosity $L_{\rm j} \simeq (1-6) \times 10^{-3}\, L_{\rm Edd} \sim (0.5-3) \times 10^{45}$\,\Lu. This range is significantly above those implied by various scaling relations proposed for radio sources in galaxy clusters, being instead very close to the maximum jet power allowed for the given accretion rate ($\sim 3 \, \dot{M}_{\rm acc} c^2$) thus implying a very high jet production efficiency in a low-accretion rate AGN. This is in agreement with recent results of numerical simulations of jets launched by spinning SMBHs \citep{mck12}. We also constrain the source lifetime as $\tau_{\rm j} \simeq 40-70$\,Myrs, meaning the total amount of deposited jet energy is $E_{\rm tot} \simeq 2\,L_{\rm j} / \tau_{\rm j} \sim (2-8) \times 10^{60}$\,ergs. We argue that about half of this energy goes into shock-heating of the surrounding thermal gas, while about $50\%$ is deposited into internal energy of the jet cavity (while a negligible fraction is radiated away), in agreement with expectations from recent long-term simulations of relativistic jet evolution \citep{per11}. The detailed analysis of the obtained X-ray data for the \rg/\cl\ system reveals some evidence for the presence of a bow-shock driven by the expanding radio lobes into the cluster environment. We derive the shock Mach number in the range $\mathcal{M}_{sh} \simeq 2-4$, which is one of the highest claimed for clusters or groups of galaxies. This, together with the recent growing evidence that powerful FR\,II radio galaxies may not be uncommon in the centers of clusters at higher redshifts, strongly supports the idea that jet-induced shock heating may indeed play an important and possibly dominant role in shaping the properties of clusters, galaxy groups, and massive ellipticals in formation. Such a role may be underestimated in studies focused on low-power radio galaxies typically found in the centers of low-redshift systems. At the same time, we note that this `mechanical' feedback in more distant sources may be difficult to access observationally due to several various reasons, including limited sensitivities and resolutions of the available X-ray detectors, but also the fact that the collective plasma processes responsible for electron heating at the shock front may become increasingly inefficient for stronger shocks, as in fact observed in the Solar System. This latter effect, when combined with a relatively long electron-ion Coulomb equilibration timescale, may result in stronger shocks driven by the expanding jets in a lower-density plasma appearing `weaker'. | 14 | 4 | 1404.5703 |
1404 | 1404.5968_arXiv.txt | As an Oort Cloud object with a record small perihelion distance of 2.7~{\Rsun} and discovered more~than a year before its encounter with the Sun, comet C/2012~S1 is a subject of considerable scientific interest. Its activity along the orbit's inbound leg evolved through a series of cycles. Two remarkable events preserved in SOHO's and/or STEREO's near-perihelion images of its tail were an early massive production of gravel at heliocentric distances of up to $\sim$100~AU(!), evidently by the annealing of amorphous water ice on and near the nucleus' surface; and, about a week before perihelion, a rapid series of powerful explosions, from the comet's interior, of water vapor with dust at extremely high rates, causing precipitous fragmentation of the nucleus, shattering it into a vast number of sublimating boulders, and ending up, a few days later, with a major, sudden drop in gas emission. The disintegration of the comet was completed by about 3.5 hours before perihelion, at a heliocentric distance of 5.2~{\Rsun}, when C/2012 S1 ceased to exist. The orbital motion in this period of time was subjected to progressively increasing outgassing-driven perturbations. A comprehensive orbital analysis results in successfully fitting the comet's observed motion from 2011 to $\sim$7 hours before perihelion. | There are several reasons for an unusually intense scientific interest in comet C/2012 S1. Perhaps the most compelling one is its record small perihelion distance, merely 2.7 solar radii (1 solar radius = 1 {\Rsun} = 0.0046548 AU), among known dynamically new comets, i.e., those arriving from the Oort Cloud. This perihelion distance beats the previous record, held by comet C/1962 C1 (Seki-Lines) by more than 4 \Rsun. Also beneficial are the early discovery of C/2012 S1 by Nevski \& Novichonok (2012) and subsequent detections of pre-discovery images of the comet when it was as far as 9.4~AU from the Sun. The early discovery allowed a comprehensive monitoring of the comet's activity on its way to perihelion. Anticipated with particular interest was the comet's behavior near perihelion and chances of its survival. An optimistic side of the controversy was argued e.g.\ by Knight \& Welsh (2013), while the most skeptical view was Ferr\'{\i}n's (2013, 2014), who nearly two months before perihelion predicted the comet's impending demise. Much effort in the present investigation is expended to examine extensive evidence on the comet's physical state as a function of time, especially in the last weeks before perihelion when the brightness, the coma and tail morphology, and the orbital motion were subject to rapid and profound changes. | 14 | 4 | 1404.5968 |
|
1404 | 1404.2854_arXiv.txt | The Fisher Information Matrix formalism \citep{Fisher} is extended to cases where the data is divided into two parts ($\bX,\Y$), where the expectation value of $\Y$ depends on $\bX$ according to some theoretical model, and $\bX$ and $\Y$ both have errors with arbitrary covariance. In the simplest case, ($\bX,\Y$) represent data pairs of abscissa and ordinate, in which case the analysis deals with the case of data pairs with errors in both coordinates, but $\bX$ can be {\em any} measured quantities on which $\Y$ depends. The analysis applies for arbitrary covariance, provided all errors are gaussian, and provided the errors in $\bX$ are small, both in comparison with the scale over which the expected signal $\Y$ changes, and with the width of the prior distribution. This generalises the Fisher Matrix approach, which normally only considers errors in the `ordinate' $\Y$. In this work, we include errors in $\bX$ by marginalising over latent variables, effectively employing a Bayesian hierarchical model, and deriving the Fisher Matrix for this more general case. The methods here also extend to likelihood surfaces which are not gaussian in the parameter space, and so techniques such as DALI (Derivative Approximation for Likelihoods) can be generalised straightforwardly to include arbitrary gaussian data error covariances. For simple mock data and theoretical models, we compare to Markov Chain Monte Carlo experiments, illustrating the method with cosmological supernova data. We also include the new method in the Fisher4Cast software. | The Fisher Information Matrix or simply Fisher Matrix has become one of the most widely used statistical tools for forecasting the errors in parameter estimation problems. It provides lower limits on the variances \change{(through the Cram\' er-Rao inequality)}, and the expected covariances of estimates of model parameters from maximum likelihood, or maximum posterior, techniques, for a given experimental design. \change{If we further assume gaussianity in two respects: that the data are jointly gaussian-distributed, and that the posterior for the parameters is gaussian, then the Fisher matrix determines the full expected posterior}. For data pairs $\{X_i,Y_i\}$ with no errors in $X$, the problem was solved many years ago \citep{Fisher}. The main value of the Fisher matrix technique is in being able to obtain error forecasts without any data, real or simulated, and is generally much faster than computing full posterior distributions with simulations \citep{2012ApJ...749...72A, Bassett:2009wr}. It is however only a first step, as it assumes the posteriors are well described by multivariate gaussian distributions, and this may not hold \citep[e.g.,][]{Wolz}, when more sophisticated analysis may be required, but it is still a very valuable tool for experimental design. Furthermore, more sophisticated forecasts for likelihood surfaces which are non-gaussian in the parameter space now exist \citep{Sellentin}. From the initial derivations of the Fisher Matrix in the cosmological context \citep{VS96,Tegmark:1997}, we have arrived today at very mature applications and implementations \citep[e.g.,][]{Bassett:2009wr, Coe:2009ti, ARefregier:2011ho}. The Fisher Matrix has been useful in proposals and projections for surveys, such as for the Cosmic Microwave Background \citep{1997astro.ph..7265T}, spectroscopic galaxy surveys \citep{BigBoss}, the Dark Energy Survey \citep{2005astro.ph.10346T}, large-scale structure \citep{2009PhRvD..79f3009C}, and in the broader discussion of the investigation of Dark Energy \citep{2006astro.ph..9591A} and estimation of neutrino masses with the future European Space Agency Euclid mission \citep{Kitching2008}. For the purposes of review and later reference in this work, we summarise the basic Fisher Matrix formalism. We begin with the likelihood of a set of data, ${\mathbf d}$ given (or conditional upon) a set of model parameters, represented by a vector $\btheta$: $p({\mathbf d}|{\btheta})$. In the simplest case, $\mathbf d$ represents only the ordinates, $\Y$. Later in the paper, we will take it to be the union of the ordinates and any other measured quantities on which $\Y$ depends, such as abscissa values, and which may be subject to error. In practice what is typically required is the posterior distribution of $\btheta$, given the data $\data$. Assuming an uninformative prior on the parameters, $p(\btheta)=$ constant, Bayes' Theorem implies $p({\btheta}|{\mathbf d})\propto p({\mathbf d}|{\btheta})=L$, the likelihood. The log-likelihood is then Taylor-expanded about its maximum. The first term is a constant, irrelevant for the discussion of parameter constraint forecasts; the second term is the first derivative, which vanishes at the point of maximum likelihood; the third term is the Hessian (curvature matrix) of the likelihood, and is the term whose ensemble average (over the data) gives the Fisher Matrix: \begin{equation} F_{\alpha\beta} = -\left\langle \frac{\partial^2 \ln L}{\partial \theta_\alpha \partial \theta_\beta} \right\rangle, \end{equation} where $\alpha$ and $\beta$ label the parameters. For the case of a gaussian likelihood, this is analytically computable, and can depend only on the expectation values of the data, $\mu(\btheta)\equiv \langle {\bf d}(\btheta)\rangle$, and the covariance, $\C(\btheta)\equiv\langle ({\bf d}-\bmu)^T({\bf d}-\bmu)\rangle$. This results in the following form for the Fisher Matrix \citep{Tegmark:1997}. \begin{equation} \F_{\alpha\beta} = \frac{1}{2}{\rm{Tr}}\left[\C^{-1}\C_{,\alpha}\C^{-1}\C_{,\beta} + \C^{-1}(\bmu_{,\alpha}\bmu_{,\beta}^T+\bmu_{,\beta}\bmu_{,\alpha}^T)\right]. \label{eqn:originalFM} \end{equation} An early example of dealing with errors in both variables was straight-line fitting, where both the statistics and astronomy communities used either {\em ad hoc} choices for the axis, or ultimately arbitrary combinations e.g., the bisector or the average of the one-dimensional fits on either axis. The evolution to two-dimensional or joint-distribution fitting was accompanied by a slow transition to the Bayesian perspective \citep{Gull:1989uy}. New tools for fitting data in the presence of two-dimensional errors have been developed and used to extract improved cosmological constraints from supernovae populations \citep{2011MNRAS.418.2308M}. Here, we develop the application of two-dimensional errors in the predictive Fisher Matrix formalism itself, \change{but the formalism can treat more general cases where the signal depends on arbitrary extra parameters} . For pedagogical discussions of straight-line fitting and Bayesian approaches to fitting, see for example \cite{Hogg, DAgostini:2005we,Kelly}. The remainder of the paper is organized as follows: \S\ref{sec:formalism} describes the formal derivation of the generalized Fisher matrix for the case of dependence of $\Y$ on an arbitrary set of gaussian-distributed variables $\bX$; \S \ref{sec:examples} describes an application of this formalism to a particular experiment, with tests on simulated data. We present conclusions in \S\ref{sec:conclusion}. \change{For the reader who is interested only in the application of the result, this is effected by simply replacing the covariance matrix $\C$ in equation (\ref{eqn:originalFM}) by the matrix $\R$ computed in equation (\ref{Rmatrix})}. | \label{sec:conclusion} In this paper we have considered the Fisher Information Matrix where some subset of the data ($\Y$) depends via a theoretical model $\langle \Y\rangle = \bmu(\bX,\btheta)$ on some other set of measured variables ($\bX$), and a set of model parameters $\btheta$ whose posterior distribution is desired. $\bX$ and $\Y$ are assumed to have gaussian errors which can have arbitrary covariance. This includes as a subset the case of ($X,Y$) data pairs with errors in both coordinates, with correlations between one independent variable and a different dependent variable, but the analysis is more general, and $\bX$ can included any other measured quantities. The main result, equation (\ref{Fnew}), is similar to the standard Fisher matrix, but with the covariance matrix replaced by a more complicated matrix (\ref{Rmatrix}) derived from the expanded covariance matrix of all variables, and the partial derivatives of the expected signals with respect to the dependent variables. The result is valid for situations where two conditions hold: the first is that a Taylor expansion of the expected signal to linear order is valid across the gaussian error of the independent variables; the second is that the errors in the independent variables are small compared with the width of the prior distribution. At the price of some complexity, we present a perturbative correction when the latter condition does not hold. In the case when the errors are uncorrelated between data pairs, the result reduces to the result one obtains from propagation of errors, where the variance of the dependent variable is increased from $\sigma_{\rm Y}^2$ to $\sigma_{\rm Y}^2+(\partial \mu/\partial x)^2\sigma_{\rm X}^2$. Since we compute the likelihood itself, it may be used to evaluate the expected likelihood surface when it is not gaussian in the parameter space, straightforwardly generalizing the DALI technique of \cite{Sellentin}. Finally, the generalised Fisher Matrix will be implemented in the Fisher4Cast software, available at http://www.mathworks.com/matlabcentral/fileexchange/20008-fisher-matrix-toolbox-fisher4cast. \noindent{\bf Acknowledgments}\\ We are grateful to the organisers of the Cape Town International Cosmology School, where this work started as a student project, to Roberto Trotta, Daniel Mortlock and Andrew Jaffe for useful discussions, and to the anonymous referee for very helpful comments and suggestions. \appendix | 14 | 4 | 1404.2854 |
1404 | 1404.2253_arXiv.txt | The DRIFT-IId dark matter detector is a m$^3$-scale low-pressure TPC with directional sensitivity to WIMP-induced nuclear recoils. Its primary backgrounds were due to alpha decays from contamination on the central cathode. Efforts to reduce these backgrounds led to replacing the $20$~\si{\micro \meter} wire central cathode with one constructed from $0.9$~\si{\micro \meter} aluminized mylar, which is almost totally transparent to alpha particles. Detailed modeling of the nature and origin of the remaining backgrounds led to an in-situ, ppt-sensitive assay of alpha decay backgrounds from the central cathode. This led to further improvements in the thin-film cathode resulting in over 2 orders of magnitude reduction in backgrounds compared to the wire cathode. Finally, the addition of O$_2$ to CS$_2$ gas was found to produce multiple species of electronegative charge carriers, providing a method to determine the absolute position of nuclear recoils and reject all known remaining backgrounds while retaining a high efficiency for nuclear recoil detection. | \label{sec:Intro} Weakly Interacting Massive Particles (WIMPs) are an attractive dark matter candidate. % It is thought that these non-baryonic particles form halos around galaxies including our own. As the baryonic portion of the galaxy rotates, our solar system passes through this halo at 220~km/s, providing an effective wind of WIMPs in the laboratory frame which comes from the constellation Cygnus \citep{Spergel1988}. Recoils induced from this WIMP wind exhibit two distinct signatures that can be used to distinguish between a true dark matter signal and background events. First, the motion of the Earth around the Sun at 30~km/s causes the apparent speed of the dark matter particles seen from Earth to vary over the course of the year and affects the rate of WIMP interactions. This results in a maximum rate in June and a minimum rate in December; this is the annual modulation. Second, the rotation of the Earth on its axis causes the incoming direction of the dark matter particles to vary over the course of the day in the laboratory frame; this is the sidereal modulation. The Directional Recoil Information From Tracks (DRIFT) experiment aims to provide an unambiguous detection of dark matter by observing this signature. | 14 | 4 | 1404.2253 |
|
1404 | 1404.7462_arXiv.txt | We present precision radial velocity (RV) datasets from Keck-HIRES and from Lick Observatory's new Automated Planet Finder Telescope and Levy Spectrometer on Mt. Hamilton that reveal a multiple-planet system orbiting the nearby, slightly evolved, K-type star HD~141399. Our 91 observations over 10.5 years suggest the presence of four planets with orbital periods of 94.35, 202.08, 1070.35, and 3717.35 days and minimum masses of 0.46, 1.36, 1.22, and 0.69 $M_J$ respectively. The orbital eccentricities of the three inner planets are small, and the phase curves are well sampled. The inner two planets lie just outside the 2:1 resonance, suggesting that the system may have experienced dissipative evolution during the protoplanetary disk phase. The fourth companion is a Jupiter-like planet with a Jupiter-like orbital period. Its orbital eccentricity is consistent with zero, but more data will be required for an accurate eccentricity determination. | The detection and characterization of extrasolar planets constitutes one of the high water marks for the entire field of astronomy and astrophysics. Slightly more than two decades ago, our own solar system was the only planetary system known, whereas at conservative last count, over a thousand extrasolar planets have now been confirmed, and thousands of additional high-quality planetary candidates are awaiting follow-up. The quest for Earth-mass and Earth-sized worlds has been effectively fulfilled, both by the Kepler Mission \citep{Borucki12} and by the Doppler velocity surveys \citep{Dumusque12}. The extraordinary string of successes has increasingly left the study of extrasolar planets standing at a crossroads. It has been clear for a number of years that the discovery of new planets merely for the sake of discovery has essentially lost its luster. Doppler velocity resources are increasingly being employed in the service of detecting very low mass planets orbiting very nearby stars \citep{Vogt10, Vogt12, Pepe11}, as well as in efforts to follow up on interesting transiting planetary candidates, especially those discovered by Kepler \citep{Howard13, Pepe13}. As a consequence, the steady stream of ``ordinary'' Jovian-mass planets with periods ranging from 100 to 1000 days orbiting nearby stars, as evidenced by Figure \ref{fig:massYrDisc}, has slowed to a relative trickle. The present article continues an extensive series of papers that have described the detection of the planetary systems that have emerged from the long-running Doppler Velocity monitoring efforts carried out with the Keck Telescope. Precision Doppler velocities were first obtained in quantity with Keck I's HIRES Spectrometer \citep{Vogt94} in 1996, following the installation of an iodine cell that can be inserted into the beam of collected starlight \citep{Butler96}. In the intervening years, a database of over 34,000 precision velocity measurements have been collected at Keck from a total of 1300 (mostly nearby) FGKM stars. The median time baseline over which velocities have been obtained for a given star in this catalog is $\tau_{d}$ = 2,563 days, and 197 stars have at least 50 high-precision radial velocity observations. Well over 100 planets have been detected with Keck, including ``firsts'', such as Gliese 436~b \citep{Butler04}, the first Neptune-mass extrasolar planet, Gliese 876~d, the first super-Earth with a secure mass determination \citep{Rivera05}, and HD~209458 \citep{Henry00}, the first known transiting planet. The target star considered in this paper, HD 141399, is a relatively nearby (36 pc distant) slightly evolved, slightly metal-rich K-type star located high in the northern sky ($46^\circ$ declination). Although it is quite bright, with V=7.2, its overall mediocrity has ensured that it has remained generally obscure, even in astronomical circles. A standard Simbad search, for example, turns up only four noncommittal mentions of the star in the literature between 1850 and 2013. Yet because HD 141399 is bright, and because it is chromospherically inactive, it has been on the Keck Radial Velocity program for over a decade. Its first iodine spectrum dates to July 2003. In recent months, it has also been repeatedly observed with the new Automated Planet Finder Telescope (APF) at Lick Observatory \citep{Vogt13}. In this paper, we report that our set of 91 velocities for HD~141399 (including 77 measurements obtained at Keck and 14 new measurements obtained at the APF) indicate that the star is accompanied by an unusual subsystem of three giant planets with near-circular orbits and with periods ranging from 94.0 days to 1070.0 days. The size of the annulus around this star spanned by these planetary orbits is associated with the zone of the terrestrial planets in our own solar system. Our radial velocities, furthermore, indicate the presence of a nearly Jupiter-mass planet at a Jupiter-like distance from the star with an orbital period of roughly a decade. Our time baseline of observations does not yet allow a definitive eccentricity determiniation for this outer planet. HD~141399's planetary system rises above a minimal threshold of interest as a consequence both of the proximity of its inner two planets to the 2:1 mean motion resonance, as well as the fact that it may well harbor an (apparently) rare near-twin of Jupiter. The plan for this paper is as follows: in \S 2 we review our Doppler velocity pipeline, with an emphasis on the new APF telescope. In \S 3, we give a brief overview of the properties of HD~141399. In \S 4, we discuss our four-planet model to explain the observed Doppler velocity variations exhibited by the star. In \S 5 we discuss the dynamical properties of the multiple-planet system that we have detected. In \S 6 we present our photometric observations of HD~141399 and in \S 7 we place our results into a larger context and conclude. \begin{figure} \plotone{f1.pdf} \caption{Diagram showing the masses of extrasolar planets versus their discovery year. Black points are extrasolar planets discovered by radial velocity observations with or without transits. Red points are extrasolar planets initially discovered by transits. The number of planets discovered by RV monitoring and having $M > 1.0\,M_J$ and $100 < P < 1000$ days has decreased in recent years, from an average of 18 such planets per year in 2007-2011, to 15 in 2012 and 7 in 2013. Data are from www.exoplanets.org and exoplanetarchive.ipac.caltech.edu, accessed 03/03/2013.} \label{fig:massYrDisc} \end{figure} | \begin{figure} \plotone{f12.pdf} \caption{{\it Green circles:} 634 planets securely detected by the radial velocity method (either with or without photometric transits). {\it Red circles:} The regular satellites of the Jovian planets in the Solar System. {\it Gray circles:} 1501 Kepler candidates and objects of interest in which {\it multiple} transiting candidate planets are associated with a single primary. Radii for these candidate planets, as reported in \citep{Batalha13}, are converted to masses assuming $M/M_{\oplus}=(R/R_{\oplus})^{2.06}$ \citep{Lissauer11}, which is obtained by fitting the masses and radii of the solar system planets bounded in mass by Venus and Saturn. Data are from www.exoplanets.org and exoplanetarchive.ipac.caltech.edu, accessed 03/03/2013. \label{fig:period_mass_ratio}} \end{figure} \begin{figure} \plotone{f13.pdf} \caption{Eccentricity-Period diagram for extrasolar planets detected via Doppler Velocity measurements. HD 141399 b, c, d, and e are indicated as open circles, and the positions of the Solar System Planets are indicated as well. Symbol size is proportional to $(M\sin(i))^{1/3}$. The color scale represents angular separation between the longitude of periastron and the line of sight to the Earth. Magenta coloring indicates that the periastron line is close to the line of sight to Earth, and signals a higher probability of transit. Planets with circular orbits or unmeasured eccentricities are indicated in gray.} \label{fig:period_ecc} \end{figure} HD~141399's retinue of planets joins a relatively select group of systems that orbit bright parent stars and which also contain multiple Jovian-mass planets observed over multiple orbital periods. As with other members of this particular class, the orbital parameters can be determined with a high degree of precision. A number of such configurations, with prime examples being provided by Upsilon Andromedae \citep{Butler99}, 55 Cancri \citep{Marcy02}, and Gliese 876 \citep{Marcy98}, were discovered during the first decade of the high-precision Doppler velocity surveys, but the discovery rate of such systems has fallen off in recent years, as ease of detection has been supplanted by intrinsic scarcity. Indeed, we expect that relatively few additional systems with such unambiguously determined orbits and large planetary radial-velocity half-amplitudes remain to be discovered orbiting $V<8$ primaries. This system -- with three Jovian-mass planets lying at stellocentric distances normally associated with the terrestrial planets of our own solar system -- is fundamentally alien. As shown in Figure \ref{fig:period_mass_ratio} and Figure \ref{fig:period_ecc}, HD~141399's companions do, however, lie within the outskirts of the now well-demarcated distribution of giant planets that have been found over the years by the radial velocity surveys. It appears that $\sim10$\% of the F, G, and K dwarf stars in the solar neighborhood \citep{Cumming08} harbor such planets, although in many cases they are (i) more massive than Jupiter, (ii) of somewhat longer period than HD~141399 b and c, and (iii) single with large orbital eccentricity. It is of substantial interest to know whether planets such as the companions to HD~141399 formed {\it in situ}, or whether they accreted the bulk of their mass further out in the protoplanetary disk and subsequently suffered Type II migration and attendant orbital decay \citep{2011exop.book..347L}. In this regard, the proximity of planets b and c to the 2:1 mean motion resonance may provide an important clue. A history of quiescent inward migration would suggest that these two planets should have been captured into resonance. Their current configuration, however, can be placed outside of the resonance with a very high degree of confidence, due to the high precision to which the orbital eccentricities have been determined. The aggregate of candidate systems with multiple transiting planets observed by Kepler show no overall preference for configurations of planets lying in low order mean motion resonance. The Kepler systems do, however, show a mild preference for configurations in which the period ratios are a few percent larger than the nominal resonant value \citep{Lissauer11}. The HD 141399 system conforms to this particular pattern. Recently, \citet{2013AJ....145....1B} and \citet{2012ApJ...756L..11L} have shown that when two planets in the vicinity of a low-order resonance interact gravitationally in the presence of dissipation, the initial orbital separation increases as orbital energy is converted to heat. Initially near-resonant pairs are driven toward orbits that are both more circular and separated by an increased distance that scales with the total integrated dissipation experienced. \citet{2012ApJ...756L..11L} suggest that the observed overdensity of near-resonant pairs can arise if tidal dissipation is unexpectedly efficient, with $Q\sim10$. This explanation seems unlikely for HD 141399 b and c, which are likely gas giants, and which likely have tidal quality factors that are orders of magnitude away from the required value. \citet{2013AJ....145....1B} argue that the dissipative mechanism is provided by interaction with the surrounding protoplanetary disk. This mechanism would appear to be more viable in this case, although the hydrodynamical details are somewhat vague and remain to be worked out. In conclusion, HD~141399 harbors a fairly unusual system in which three (and likely a fourth) Jovian-mass planets lie on low-eccentricity orbits with periods that conform to one's naive expectation for terrestrial planets. Confirmation of this system was significantly aided by velocity measurements from the APF telescope. The quality of the measurements that the APF is obtaining show it is functioning as intended, and that it is producing Doppler measurements that conform to the current state-of-the-art. | 14 | 4 | 1404.7462 |
1404 | 1404.7712_arXiv.txt | Complementing our Washington photometric studies on Galactic open clusters (OCs), we now focus on four poorly studied OCs located in the first and fourth Galactic quadrants, namely BH\,84, NGC\,5381, BH\,211 and Czernik\,37. We have obtained CCD photometry in the Washington system $C$ and $T_1$ passbands down to $T_1$ $\sim$ 18.5 magnitudes for these four clusters. Their positions and sizes were determined using the stellar density radial profiles. We derived reddening, distance, age and metallicity of the clusters from extracted $(C-T_1,T_1)$ color-magnitude diagrams (CMDs), using theoretical isochrones computed for the Washington system. There are no previous photometric data in the optical band for BH\,84, NGC\, 5381 and BH\,211. The CMDs of the observed clusters show relatively well defined main sequences, except for Czernik\,37, wherein significant differential reddening seems to be present. The red giant clump is clearly seen only in BH\,211. For this cluster, we estimated the age in (1000$^{+260}_{-200}$) Myr, assuming a metallicity of $Z$ = 0.019. BH\,84 was found to be much older than it was previously believed, while NGC\,5381 happened to be much younger than previously reported. The heliocentric distances to these clusters are found to range between 1.4 and 3.4 kpc. BH\,84 appears to be located at the solar galactocentric distance, while NGC\,5381, BH\,211 and Czernik\,37 are situated inside the solar ring. | \label{} Galactic open clusters (OCs) have long been considered excellent targets not only to probe the Galactic disc properties \citep*{l82,f95,pietal06,bietal06} but also to trace its chemical evolution (see, e.g., Chen et al., 2003 and references therein). Because it is relatively simple to estimate ages, distances and metallicities of OCs fairly accurately, their basic parameters constitute excellent tracers to the structure and chemical evolution of the Galactic disc. The proximity of most OCs to the Galactic plane, however, usually restricts this analysis to the most populous ones and/or to those located within a few kpc from the Sun (Bonatto et al., 2006). Although there are at present estimates of a total of about $25\times10^3$ OCs in the Milky Way \citep{poetal10}, there is not yet an estimation of fundamental parameters such as reddening, distance and age for nearly 30\% of the $\sim$ 2200 catalogued Galactic OCs \citep*{dietal02, bietal03,duetal03}. The present work is part of a current project of photometric observation of Galactic OCs in the Washington system that is being developed at the Observatorio Astron\'omico de la Universidad Nacional de C\'ordoba (Argentina). This project aims at determining the fundamental parameters or at refining the quality of observationally determined properties for some unstudied or poorly studied OCs, located in different regions of the Milky Way. Washington photometry has proved to be a valuable tool to determine the fundamental parameters of OCs since information on cluster membership, reddening, distance, age and metallicity is obtained through the analysis of the $(C-T_1,T_1)$ color-magnitude diagram (CMD). We have already reported results based on Washington CCD $CT_1$ photometric data on several young (e.g., Piatti et al. 2003a), intermediate-age (e.g., Clari\'a et al., 2007) and old Galactic OCs (e.g., Piatti et al. 2004). These studies have contributed not only to the individual characterization of these stellar systems but also to the global understanding of some properties of the Galactic disc (e.g., Parisi et al. 2005). In this study we provide new high-quality photometric CCD data obtained with the Washington system $C$ and $T_1$ passbands down to $T_1$ $\sim$ 18.5 magnitudes in the fields of four faint, poorly studied OCs, namely BH\,84, NGC\,5381, BH\,211 and Czernik\,37. The equatorial and Galactic coordinates of the cluster centers taken from the WEBDA Open Cluster Database \citep{me05} are listed in Table 1, together with the angular sizes given by \citet{ah03}. The selected clusters are located in the first and fourth Galactic quadrants (280$^\circ$ $<$ {\it l} $<$ 3$^\circ$) near the Galactic plane ($\mid$b$\mid$ $\leq$ 3$^\circ$). As far as we know, no previous photometric data in the optical band exist for BH\,84, NGC\,5381 and BH\,211. The four selected clusters have been examined by \citet{buetal11} and \citet{ta08,ta11} using Two-Micron All-Sky Survey (2MASS) data. Some preliminary results about BH\,84, NGC\,5381 and BH\,211 are presented in \citet{maetal13}. A brief description of these objects as well as a summary of previous results for the fields under investigacion is given below: {\it BH\,84}. First recognized as an OC by \citet{vh75}, this object (IAU designation C0959-579) seems to be a detached, relatively poor and faint OC in the Carina constellation (Fig. 1). It shows the typical morphology of a Trumpler class II-1p cluster, which is characterized by a slight concentration of member stars of similar brightness bf and relatively small population. The only observational data for this object are those given in the 2MASS catalog and discussed by \citet{buetal11}. These authors derived a reddening $E(B-V)$ = 0.60 and suggest that BH\,84 is a young cluster ($\sim$ 18 Myr), located at a heliocentric distance $d$ = (2.92 $\pm$ 0.19) kpc. {\it NGC\,5381}. This is a cluster in Centaurus, also designated as BH\,156 \citep{vh75}. \citet{ah03} refer to this object as belonging to Trumpler class II-2m, i.e., a moderately rich, detached cluster with little central concentration and medium-range bright stars (Fig. 1). According to these authors, NGC\,5381 has a comparatively large angular diameter of 11$'$. A search for variable stars in the cluster field was carried out by \citet{pietal97}. Using 2MASS data, \citet{ta11} suggests that NGC\,5381, slightly reddened by $E(B-V)$ = 0.06, is an intermediate-age cluster ($\sim$ 1.6 Gyr) located at 1.2 kpc from the Sun. {\it BH\,211}. This object (C1658-410) appears to be somewhat elongated in the East-West direction (Fig. 1). BH\,211 is a detached, moderately poor and relatively faint group of stars, first recognized as an OC in Scorpius by \citet{vh75}. It is a small-sized OC situated very near the Galactic center direction, practically on the Galactic plane (Table 1). The only observational data-set for this object is the one given in the 2MASS catalog and discussed by \citet{buetal11} who found the following results: $E(B-V)$ = 0.48, $d$ = (1.38 $\pm$ 0.09) kpc and $\sim$ 1.6 Gyr. {\it Czernik\,37}. Also known as BH\,253 \citep{vh75}, this is a relatively faint cluster (C1750-273) first recognized in Sagittarius by \citet{cz66}. As indicated by its Trumpler class (II-1m), it shows a slight central concentration but can be identified by its relatively dense population compared to that of the field stars (Fig. 1). This cluster is projected on to the central bulge of the Galaxy, only 2$^\circ$ from the Galactic center direction. \citet{caetal05} presented CCD $BVI$ photometry in the field of Czernik\,37. Although they conclude that this may be a sparse but real cluster superimposed on the Galactic bulge population, they do not provide its physical parameters. Using 2MASS data, \citet{ta08} derived a heliocentric distance of 1.7 kpc, $E(B-V)$ = 1.03 and an age of 0.6 Gyr. The layout of this paper is as follows. Section 2 provides details on our observations and the data reduction procedure. In Section 3 we determine the cluster centers and the stellar density radial profiles. Section 4 deals with the determination of cluster fundamental parameters through the fitting of theoretical isochrones. A brief description of the results, including a comparison with previous findings, is presented in Section 5, while the final conclusions are summarized in Section 6. | We have presented new CCD Washington $CT_1$ photometry in the field of four Galactic OCs projected onto the two inner quadrants of the Galactic plane. These data were obtained with the main purpose of estimating the cluster fundamental astrophysical parameters. We performed a star count analysis of the cluster fields to assess the clusters' reality as over-densities of stars with respect to the field and estimated the cluster radii. We determined the center of the clusters by finding the maximum surface number density of the stars in each cluster. New equatorial coordinates for the 2000.0 epoch are now provided. We outlined possible solutions for cluster fundamental parameters by matching theoretical isochrones, which reasonably reproduce the main cluster features in the $(C-T_1,T_1)$ CMDs. In all cases, the best fits were obtained using solar metallicity isochrones. BH\,211 was found to be the oldest object of our sample with an age of around 1.0 Gyr. Czernik\,37, the most heavily reddened cluster of the sample, with a mean colour excess $E(B-V)$ = 1.47, is very likely affected by differential reddening. BH\,84 is located in the fourth Galactic quadrant just before the tangent to the Carina branch of the Carina-Sagittarius spiral arm. This cluster turned out to be much older than previously believed. Conversely, NGC\,5381 was found to be much younger than previously reported. It appears to have a relatively small but conspicuous nucleus and a low-density extended corona. We estimated the angular core and corona radii as $\sim$ 2.2' and $\sim$ 5.7', respectively. The derived fundamental properties for the studied clusters are listed in Table 9. Previous estimates of cluster parameters are listed in Table 10, for easy comparison with the present results. Since two of the studied clusters, BH\,211 and Czernik\,37, are in the VISTA Variables in the Via Lactea (VVV) survey \citep{mietal10}, additional observational information about these two objects can be found in this database. \begin{table} \caption{Basic parameters of the four open clusters} \vspace{0.1cm} \begin{tabular}{@{}|llcccc|lc|}\hline \multicolumn{6}{|c|}{WEBDA} &\multicolumn{2}{c|}{This study} \\ Cluster & $\alpha$$_{\rm 2000}$ & $\delta$$_{\rm 2000}$ & {\it l} & $b$ & Diam. & $\alpha$$_{\rm 2000}$ & $\delta$$_{\rm 2000}$ \\ & (h m s) & ($^\circ$ ' ") & ($^\circ$) & ($^\circ$) & (') & (h m s) & ($^\circ$ ' ") \\ \hline \hline BH\,84 & 10 01 19 & -58 13 00 & 280.06 & -2.42 & 4.5 & 10 01 19 & -58 13 33 \\ NGC\,5381 & 14 00 41 & -59 35 12 & 311.60 & 2.11 & 11.0 & 14 00 41 & -59 35 20 \\ BH\,211 & 17 02 11 & -41 06 00 & 344.97 & 0.46 & 4.0 & 17 02 10 & -41 05 57 \\ Czernik\,37 & 17 53 16 & -27 22 00 & 2.22 & -0.64 & 3.0 & 17 53 17 & -27 22 36 \\ \hline \end{tabular} \end{table} \begin{center} \begin{longtable}{|ccccc|} \caption{Observation log of observed clusters} \\ \hline Cluster & Date & Filter & Exposure & Airmass \\ & & & (sec) & (") \\ \hline \hline \endfirsthead \multicolumn{5}{c} {\tablename\ \thetable\ -- \textit{continued}} \\ \hline Cluster & Date & Filter & Exposure & Airmass \\ & & & (sec) & (") \\ \hline \hline \endhead \hline % \endfoot \hline \endlastfoot BH\,84 & May 9, 2008 & $C$ & 30 & 1.13 \\ & & $C$ & 45 & 1.13 \\ & & $C$ & 300 & 1.13 \\ & & $C$ & 450 & 1.13 \\ & & $R$ & 5 & 1.13 \\ & & $R$ & 7 & 1.12 \\ & & $R$ & 30 & 1.12 \\ & & $R$ & 45 & 1.12 \\ \hline NGC\,5381 & May 11, 2008 & $C$ & 90 & 1.15 \\ & & $C$ & 120 & 1.15 \\ & & $C$ & 600 & 1.15 \\ & & $C$ & 30 & 1.16 \\ & & $R$ & 3 & 1.16 \\ & & $R$ & 120 & 1.16 \\ & & $R$ & 120 & 1.16 \\ \hline BH\,211 & & $C$ & 30 & 1.02 \\ & & $C$ & 45 & 1.02 \\ & & $C$ & 300 & 1.02 \\ & & $C$ & 450 & 1.02 \\ & & $R$ & 5 & 1.02 \\ & & $R$ & 7 & 1.02 \\ & & $R$ & 30 & 1.02 \\ & & $R$ & 45 & 1.02 \\ \hline Czernik\,37 & May 10, 2008 & $C$ & 30 & 1.00 \\ & & $C$ & 45 & 1.00 \\ & & $C$ & 300 & 1.00 \\ & & $C$ & 450 & 1.00 \\ & & $R$ & 30 & 1.00 \\ & & $R$ & 45 & 1.00 \\ & & $R$ & 5 & 1.00 \\ & & $R$ & 7 & 1.00 \\ \end{longtable} \end{center} \begin{table}[h] \begin{center} \caption{Standard system mean calibration coefficients} \vspace{0.1cm} \begin{tabular}{@{}|cc|}\hline $C$ & $T_1$ \\ \hline \hline $a_1$ = 3.61 $\pm$ 0.03 & $b_1$ = 3.04 $\pm$ 0.02 \\ $a_2$ = 0.56 $\pm$ 0.01 & $b_2$ = 0.33 $\pm$ 0.02 \\ $a_3$ = -0.19 $\pm$ 0.01 & $b_3$ = -0.03 $\pm$ 0.01 \\ \hline \end{tabular} \end{center} \end{table} \begin{table} \caption[]{CCD $CT_1$ data of stars in the field of BH\,84} \vspace{0.1cm} \begin{tabular}{|cccccccc|} \hline Star & $X$\hspace{0.2cm} & $Y$\hspace{0.2cm} & $T_{1}$ & $\sigma$$T_1$ & $C-T_1$ & $\sigma$$(C-T_1)$ & n \\ & (pixel) & (pixel) & (mag) & (mag) & (mag) & (mag) & \\ \hline 495 & 1870.105 & 584.378 & 17.590 & 0.094 & 3.237 & 0.071 & 1 \\ 496 & 926.421 & 585.257 & 15.910 & 0.018 & 1.955 & 0.015 & 2 \\ 497 & 1568.217 & 585.395 & 13.705 & 0.010 & 1.653 & 0.009 & 1 \\ - & - & - & - & - & - & - & - \\ - & - & - & - & - & - & - & - \\ \hline \end{tabular} \end{table} \begin{table} \caption[]{CCD $CT_1$ data of stars in the field of NGC\,5381} \vspace{0.1cm} \begin{tabular}{|cccccccc|} \hline Star & $X$\hspace{0.2cm} & $Y$\hspace{0.2cm} & $T_{1}$ & $\sigma$$T_1$ & $C-T_1$ & $\sigma$$(C-T_1)$ & n \\ & (pixel) & (pixel) & (mag) & (mag) & (mag) & (mag) & \\ \hline 499 & 918.737 & 415.590 & 13.625 & 0.005 & 1.874 & 0.005 & 2 \\ 500 & 501.666 & 418.198 & 14.553 & 0.005 & 1.225 & 0.005 & 1 \\ 501 & 1000.223& 422.144 & 16.461 & 0.022 & 1.801 & 0.017 & 2 \\ - & - & - & - & - & - & - & - \\ - & - & - & - & - & - & - & - \\ \hline \end{tabular} \end{table} \begin{table} \caption[]{CCD $CT_1$ data of stars in the field of BH\,211} \vspace{0.1cm} \begin{tabular}{|cccccccc|} \hline Star & $X$\hspace{0.2cm} & $Y$\hspace{0.2cm} & $T_{1}$ & $\sigma$$T_1$ & $C-T_1$ & $\sigma$$(C-T_1)$ & n \\ & (pixel) & (pixel) & (mag) & (mag) & (mag) & (mag) & \\ \hline 496 & 1255.054 & 1053.165 & 16.858 & 0.043 & 2.626 & 0.034 & 2 \\ 497 & 1137.926 & 2039.064 & 16.899 & 0.052 & 2.597 & 0.040 & 1 \\ 498 & 983.109 & 1055.830 & 15.363 & 0.016 & 2.018 & 0.014 & 1 \\ - & - & - & - & - & - & - & - \\ - & - & - & - & - & - & - & - \\ \hline \end{tabular} \end{table} \begin{table} \caption[]{CCD $CT_1$ data of stars in the field of Czernik\,37} \vspace{0.1cm} \begin{tabular}{|cccccccc|} \hline Star & $X$\hspace{0.2cm} & $Y$\hspace{0.2cm} & $T_{1}$ & $\sigma$$T_1$ & $C-T_1$ & $\sigma$$(C-T_1)$ & n \\ & (pixel) & (pixel) & (mag) & (mag) & (mag) & (mag) & \\ \hline 500 & 1480.077 & 796.860 & 15.898 & 0.038 & 3.126 & 0.033 & 1 \\ 501 & 961.983 & 801.728 & 15.552 & 0.021 & 2.819 & 0.018 & 2 \\ 502 & 1114.297 & 805.020 & 15.474 & 0.041 & 2.794 & 0.03 & 1 \\ - & - & - & - & - & - & - & - \\ - & - & - & - & - & - & - & - \\ \hline \end{tabular} \end{table} \begin{table} \caption{Cluster sizes} \vspace{0.1cm} \begin{tabular}{|l|cc|cc|cc|} \hline Cluster & \multicolumn{2}{|c|}{$r_{FWHM}$} & \multicolumn{2}{c|}{$r_{clean}$} & \multicolumn{2}{c|}{$r_{cl}$} \\ & (px) & (pc) & (px) & (pc) & (pix) & (pc) \\ \hline BH\,84 & 150 & 1.0 & 200 & 1.3 & 550 & 3.6 \\ NGC\,5381 & 150 & 0.8 & 280 & 1.4 & 1200 & 6.1 \\ BH\,211 & 150 & 0.4 & 250 & 0.7 & 580 & 1.6 \\ Czernick\,37 & 220 & 0.6 & 320 & 0.9 & 550 & 1.5 \\ \hline \end{tabular} \end{table} \begin{table} \caption{Fundamental properties of the observed clusters} \vspace{0.1cm} \begin{tabular}{|lcccccccc|} \hline Cluster & $E(B-V)$ & $d$ & Age & [Fe/H] & $X$ & $Y$ & $Z$ & $R_{GC}$ \\ & (mag) & (kpc) & (Myr) & (dex) & (kpc) & (kpc) & (kpc) & (kpc) \\ \hline \hline BH\,84 & 0.63 $\pm$ 0.05 & 3.37 $\pm$ 0.48 & 560$^{+150}_{-110}$ & 0.0 & 7.91 & -3.32 & -0.14 & 8.58 \\ NGC\,5381 & 0.46 $\pm$ 0.04 & 2.63 $\pm$ 0.40 & 250$^{+65}_{-50}$ & 0.0 & 6.76 & -1.97 & 0.10 & 7.04 \\ BH\,211 & 0.61 $\pm$ 0.05 & 1.44 $\pm$ 0.21 & 1000$^{+260}_{-200}$ & 0.0 & 7.11 & -0.37 & 0.01 & 7.12 \\ Czernik\,37 & 1.47 $\pm$ 0.25 & 1.44 $\pm$ 0.86 & 250$^{+100}_{-65}$ & 0.0 & 7.06 & 0.06 & -0.02 & 7.06 \\ \hline \end{tabular} \end{table} \begin{table} \caption{Previous stimates of cluster parameters} \vspace{0.1cm} \begin{tabular}{|lcccc|} \hline Cluster & $E(B-V)$ & $d$ & Age & Reference \\ & (mag) & (kpc) & (Myr) & \# \\ \hline \hline BH\,84 & 0.60 & 2.92 $\pm$ 0.19 & 18 & 1 \\ NGC\,5381 & 0.06 & 1.17 $\pm$ 0.05 & 1600 & 2 \\ BH\,211 & 0.48 & 1.38 $\pm$ 0.09 & 1600 & 1 \\ Czernik\,37 & 1.03 & 1.73 $\pm$ 0.08 & 600 & 3 \\ \hline \end{tabular} References: (1) Bukowiecki et al.(2011); (2) Tadross (2011); (3) Tadross (2008) \end{table} \begin{figure} \begin{center} \includegraphics*[width=12cm]{Fig1.ps} \end{center} \caption{Schematic finding charts of the stars observed in BH\,84 (top left), NGC\,5381 (top right), BH\,211 (bottom left) and Czernik\,37 (bottom right). North is up and East is to the left. The sizes of the plotting symbols are proportional to the $T_1$ brightness of the stars. Two circles $r_{clean}$ and $r_{cl}$ wide are shown around the cluster centers ({\it crosses}).} \label{fig1} \end{figure} \begin{figure} \begin{center} \includegraphics*[width=12cm]{Fig2.ps} \end{center} \caption{$T_1$ magnitude and C-T$_1$ color photometric errors as a function of $T_1$ for stars measured in the field of NGC\,5381.} \label{fig2} \end{figure} \begin{figure} \begin{center} \includegraphics*[width=12cm]{Fig3.ps} \end{center} \caption{($C-T_1$,$T_1)$ CMDs for stars observed in the field of BH\,84, NGC\,5381, BH\,211 and Czernik\,37.} \label{fig3} \end{figure} \begin{figure} \begin{center} \includegraphics*[width=12cm]{Fig4.ps} \end{center} \caption{Cluster stellar density radial profiles normalized to a circular area of 50 pixel radius. The radius at the FWHM (r$_{FWHM}$) and the adopted cluster radius (r$_{cl}$) are indicated by green vertical lines. The red horizontal lines represent the measured background levels.} \label{fig4} \end{figure} \begin{figure} \begin{center} \includegraphics*[width=12cm]{Fig5.ps} \end{center} \caption{CMDs for stars observed in different extracted circular regions around BH\,84 center as indicated in each panel.} \label{fig5} \end{figure} \begin{figure} \begin{center} \includegraphics*[width=12cm]{Fig6.ps} \end{center} \caption{CMDs for stars observed in different extracted circular regions around NGC\,5381 center as indicated in each panel.} \label{fig6} \end{figure} \begin{figure} \begin{center} \includegraphics*[width=12cm]{Fig7.ps} \end{center} \caption{CMDs for stars observed in different extracted circular regions around BH\,211 center as indicated in each panel.} \label{fig7} \end{figure} \begin{figure} \begin{center} \includegraphics*[width=12cm]{Fig8.ps} \end{center} \caption{CMDs for stars observed in different extracted circular regions around Czernik\,37 center as indicated in each panel.} \label{fig8} \end{figure} \begin{figure} \begin{center} \includegraphics*[width=12cm]{Fig9.ps} \end{center} \caption{r $<$ r$_{clean}$ (C-$T_1$,$T_1$) CMDs for stars in: BH\,84 (top left), NGC\,5381 (top right), BH\,211 (bottom left) and Czernik\,37 (bottom right). The ZAMS and the adopted isochrones from Girardi et al. (2002) are overplotted with solid lines. The isochrones associated to the cluster age errors are indicated by dashed lines, for comparison purposes.} \end{figure} | 14 | 4 | 1404.7712 |
1404 | 1404.0132_arXiv.txt | {} {To investigate the evolution of plasma properties and Stokes parameters in photospheric magnetic bright points using 3D magneto-hydrodynamical simulations and radiative diagnostics of solar granulation.} {Simulated time-dependent radiation parameters and plasma properties were investigated throughout the evolution of a bright point. Synthetic Stokes profiles for the FeI $630.25~\mathrm{nm}$ line were calculated, which allowed the evolution of the Stokes-$I$ line strength and Stokes-$V$ area and amplitude asymmetries to also be investigated.} {Our results are consistent with theoretical predictions and published observations describing convective collapse, and confirm this as the bright point formation process. Through degradation of the simulated data to match the spatial resolution of SOT, we show that high spatial resolution is crucial for the detection of changing spectro-polarimetric signatures throughout a magnetic bright point's lifetime. We also show that the signature downflow associated with the convective collapse process is reduced towards zero as the radiation intensity in the bright point peaks, due to the magnetic forces present restricting the flow of material in the flux tube.} {} | \begin{figure*} \includegraphics[width=\textwidth]{fig1.pdf} \caption{Simulated temporal evolution of plasma properties during the formation and disappearance of a MBP. \emph{Top left:} Normalised G-band intensity. \emph{Top right:} Evolution of total pressure (solid), gas pressure (dotted), magnetic pressure (dashed) and mass density (dash-dot) at log($\tau_{500\mathrm{nm}}$)=0. \emph{Bottom left:} Evolution of modulus of the magnetic field (solid) and line-of-sight velocity (dashed) at log($\tau_{500\mathrm{nm}}$)=0, where downward (red-shifted) velocity is positive. \emph{Bottom right:} Plasma temperature evolution at log($\tau_{500\mathrm{nm}}$)=0 (solid), log($\tau_{500\mathrm{nm}}$)=-1 (dashed) and log($\tau_{500\mathrm{nm}}$)=-2 (dotted). Markers along the x-axis indicate times at which images and profiles were taken for Figure~\ref{stokes1}.} \label{t-0} \end{figure*} The dominant pattern of the solar photosphere, outside sunspots, is granulation. Granular flows remove magnetic flux from the granules and advect it into the dark intergranular lanes, where the magnetic field accumulates until it reaches a limit ($\sim$500G) given by the equipartition of energy \citep{parker2,bellot1}. However, further intensification occurs by a process called convective collapse. The magnetic field within the intergranular lanes suppresses horizontal convective motions which carry heat to the downflow regions, causing the gas to become cooler and denser, and accelerates the downflow. Due to this, the flux tube becomes partially evacuated. To restore balance between the tube and its surroundings, the tube walls collapse inwards, until stable, and compress the magnetic field lines, thereby intensifying the magnetic field with strengths often in excess of $1~\mathrm{kG}$ \citep{spruit1, takeuchi1, schussler1}. The enhanced magnetic field decreases the density within the flux tube which, in turn, lowers the surface of optical depth unity to layers at higher temperatures. It also leads to increased heating from the granular walls and produces the bright appearance of magnetic bright points (MBPs) \citep{keller1, voegler2, steiner1}. MBPs evolve on timescales smaller than that of granulation \citep{mehltretter1}, with lifetimes of less than $120~\mathrm{s}$ for almost all MBPs \citep{abramenko1, keys1}. Their diameters range from $\sim120-600~\mathrm{km}$ \citep{bovelet1} with \citet{wiehr1} reporting a predominant diameter of 160$\pm$20 km. Semi-automatic routines for the detection of MBPs in solar observations were recently developed, and statistical studies of their area distribution \citep{crockett1}, transverse velocity distribution \citep{keys1}, were undertaken based on MHD numerical modelling with MuRAM \citep{voegler1} and high-cadence observations with the Rapid Oscillations in Solar Atmosphere (ROSA; \citet{jess2}) instrument at the Dunn Solar Telescope (USA). A recent study by \citet{shelyagav} demonstrated the presence of Alfv{\'e}n waves in the intergranular magnetic field concentrations, previously attributed to photospheric vortices \citep{shelyagvort}. These waves carry Poynting flux \citep{shelyagpf}, which potentially is able to propagate and supply the energy to higher layers of the solar atmosphere. While the MBPs exhibit rotational behaviour in the photospheric intergranular lanes, as was observationally demonstrated by \citet{bonetbp}, they do not show ``Alfv{\'e}n-type" transverse oscillatory motions since they are formed deeper than the region where Alfv{\'e}n waves are detected in simulations. \citet{berger1} reported on the dynamics of small-scale magnetic elements. However, the first observations which allowed the simultaneous study of plasma properties and magnetic properties simultaneously were undertaken by \citet{nagata1}. They used the Solar Optical Telescope (SOT; \citet{tsuneta1,suematsu1}) onboard Hinode to investigate the evolution of an MBP and confirmed the convective collapse process as the formation mechanism. \citet{narayan1} used observations from the Crisp Imaging SpectroPolarimeter (CRISP; \citet{scharmer2}) installed at the Swedish 1-m Solar Telescope (SST; \citet{scharmer1}), and also made use of spectro-polarimetric data to detect downflows associated with small magnetic features. While plenty of observations, simulations and analyses of MBPs already exist, there is a need to provide details on the processes of formation and evolution of MBPs in connection to the time-dependent physical and radiative properties of the background solar plasma. In this paper we use radiative MHD simulations of solar magneto-convection to study the temporal evolution of MBPs. Section 2 provides a brief description of the simulations used in this work. Physical parameters such as magnetic field strength, density, line-of-sight velocity, as well as observable quantities such as Stokes parameters are presented in Section 3. A comparison between simulations and published observational findings is also made. Concluding remarks are given in Section 4. | In this paper, we investigated the plasma properties and spectro-polarimetric parameters of radiation during the evolution of a photospheric magnetic bright point using numerical simulations and detailed radiative diagnostics of the simulated photospheric magneto-convection models. We emphasise that the intergranular lanes are magnetic regions with the MBP being just a feature in a more continuous magnetic field appearing as a result of radiative transfer effects. Our approach in this paper has been to follow, in time, the area where we see the appearance and disappearance of a MBP and study the flows and magnetic fields in and around that area. We define the MBP feature as an enhancement of the intensity, and magnetic field strength and investigate that location. Certain aspects of the evolution of the MBP plasma parameters and Stokes profiles in simulation have been analysed such as the magnetic field, line-of-sight velocity, pressure, density and temperature along with the full Stokes vector. The evolution of these properties are shown to be consistent with observations of convective collapse \citep{nagata1}. We observe a strong downflow evacuating the flux tube followed by an enhancement in G-band intensity as optical depth unity occurs lower in the atmosphere. The density was observed to decrease as a result of the evacuation, causing the flux tube walls to collapse inwards compressing the field lines and intensifying the magnetic field strength . \citet{nagata1}, \citet{fischer} and \citet{shimizu} observed strong downflows associated with MBPs to precede the magnetic field intensification whereas, \citet{narayan1} found the downflows to accompany the magnetic field increase. Our results agree with the former alongside the theory of convective collapse which expects the downflow to precede the magnetic field intensification. We surmise that the results from \citet{narayan1} are probably due to lower temporal resolution. As the downflow tends towards zero, the evacuation stops thereby starting to increase the density in the tube and causing the magnetic field to begin to reduce as well. This decrease in downflow was noted to be at the MBP centre and caused by strong magnetic pressure that builds up in the flux tube restricting material flow. The redshift of the Stokes profiles confirm the downflow that occurs during the photospheric MBP formation process, and the Zeeman splitting observed in the Stokes-$I$ profiles indicates the strong magnetic field that builds up throughout its lifetime. However, these signatures of convective collapse are yet to be observed. The increasing amplitude of Stokes-$V$ profiles throughout the process has been observed by \citet{viticchie1} and \citet{shimizu} but the reduction of the amplitude after the process is, as yet, only recorded by \citet{grossmann1} using 2D numerical simulations. The comparison between the Stokes profiles and the plasma properties in Figure~\ref{t-0} allows us to comment on the continued increase in Stokes-$V$ amplitude after the MBP has peaked. Our results confirm convective collapse as the process involved in the MBP formation. However, the high resolution provided by the simulations allow much more information and structure to be detected and analysed, such as the Zeeman splitting. The spatial resolution of current instrumentation also misses some asymmetry structure in the Stokes-$V$ profiles which provides important information on the line-of-sight velocity and magnetic field gradients. The velocity maps show no downflow in the MBP centre as it peaks, which has not yet been observed. Therefore, high spatial and temporal resolution observations combined with vector magnetograms are essential for understanding the small-scale photospheric structures and dynamics. The Advanced Technology Solar Telescope (ATST; \citet{rimmele}) will detect and track myriads of individual flux concentrations down to a size of 25 km and provide insights on the energy hidden in the dark internetwork field of the quiescent Sun. | 14 | 4 | 1404.0132 |
1404 | 1404.6522_arXiv.txt | { In this paper, we examine observational constraints on the power law cosmology; essentially dependent on two parameters $H_0$ (Hubble constant) and $q$ (deceleration parameter). We investigate the constraints on these parameters using the latest 28 points of H(z) data and 580 points of Union2.1 compilation data and, compare the results with the results of $\Lambda$CDM. We also forecast constraints using a simulated data set for the future JDEM, supernovae survey. Our studies give better insight into power law cosmology than the earlier done analysis by Kumar [arXiv:1109.6924] indicating it tuning well with Union2.1 compilation data but not with H(z) data. However, the constraints obtained on $<H_0>$ and $<q>$ i.e. $H_0$ average and $q$ average using the simulated data set for the future JDEM, supernovae survey are found to be inconsistent with the values obtained from the H(z) and Union2.1 compilation data. We also perform the statefinder analysis and find that the power-law cosmological models approach the standard $\Lambda$CDM model as $q\rightarrow -1$. Finally, we observe that although the power law cosmology explains several prominent features of evolution of the Universe, it fails in details.} \date{\today} \keywords {inflation, dark energy theory } \begin{document} | \label{sect1}The Standard Cosmological Model (SM) of Universe {\it a la} $\Lambda$CDM complemented by the inflationary phase is remarkably a successful theory, although, the cosmological constant problem still remains to be one of the major unsolved problems \cite{sami} of our times. It is therefore reasonable to examine the alternative cosmological models to explain the observed Universe. Power-law cosmology is one of the interesting alternatives to deal with some usual problems (age, flatness and horizon problems etc.) associated with the standard model. In such a model, the cosmological evolution is explained by the geometrical scale factor $a(t) \propto t^\beta$ with $\beta$ as a positive constant. The power law evolution with $\beta\geq1$ has been discussed at length in a series of articles in distinct contexts \cite{lohiya,batra1,batra2,geh1,geh2,dev1,dev2,sethi2,zhu}; phantom power-law cosmology is discussed in reference \cite{kae}. The motivation for such a scenario comes from a number of considerations. For example, power-law cosmology does not face the horizon problem \cite{sethi2}, as well as the flatness problem. Another remarkable feature of these models is that they easily accommodate high redshift objects and hence reduce the age problem. These models also deal with the fine tuning problem, in an attempt to dynamically solve the cosmological constant problem \cite{mann,allen,dol,ford,wein}. A power law evolution of the cosmological scale factor with $\beta \approx 1$ is an excellent fit to a host of cosmological observations. Any model supporting such a coasting presents itself as a falsifiable model as far as classical cosmological tests are concerned as it exhibits distinguishable and verifiable features. Classical cosmological tests also support such kind of evolution, such as the galaxy number counts as a function of redshift and the data on angular diameter distance as a function of redshift \cite{kolb}. However, these tests are not considered as reliable tests of a viable model since these are marred by evolutionary effects (e.g. mergers). Now, SNe Ia (reliable standard candles), and hubble test have become more reliable to that of a precision measurement. Cosmological parameters prove to be the backbone of any of the cosmological models, therefore it becomes important to obtain a concise range or more specifically, the estimated values of such parameters using available observational data, so that the said model can explain the present evolution of Universe more precisely. In this series of cosmological parameters we observe that Hubble constant ($H_{0}$) and deceleration parameter ($q$) are very important in describing the current nature of the Universe. $H_0$ explains the current expansion rate of the Universe whereas $q$ describes the nature of the expansion rate. In last few years, various attempts have been done to evaluate the value of $H_{0}$. Freedman et al. \cite{free} evaluated a value of $H_{0}=72 \pm 8$ km/s/Mpc. Suyu et al.\cite{suy} evaluated $H_{0}$ as $69.7_{-5.0}^{+4.9}$ km/s/Mpc. WMAP7 evaluated the value of $H_{0}=71.0\pm 2.5$ km/s/Mpc (with WMAP alone), and $H_{0} = 70.4_{-1.4}^{+1.3}$ km/s/Mpc (with Gaussian priors ) \cite{jar}. Numerous other evaluates of $H_{0}$ are $73.8 \pm 2.4$ km/s/Mpc \cite{rie}, $67.0\pm 3.2$ km/s/Mpc \cite{beu}. Most recent PLANCK evaluate of the Hubble constant gives a value of $H_{0}$ = $67.3\pm 1.2$ km/s/Mpc \cite{Ade:2013zuv}. Along with the above mentioned evaluates of $H_0$, several other authors, \cite{dev1}, \cite{sethi2,zhu,bgum,sur,gum} obtained the constraints on cosmological parameters including $H_0$, $q$ and $\beta$ for open, closed and flat power law cosmology. Numerical results for flat power-law cosmology have been described in Table \ref{tabparm}. In a recent paper, Kumar \cite{sur} has investigated observational constraints on the power-law cosmological parameters using H(z) and SN Ia data and discussed various features of power-law cosmology. In the present work, we are investigating the scenario similar to an analysis done in reference \cite{sur} for flat power law cosmology. We compare the results of the model under consideration with the results obtained from $\Lambda$CDM and with that of Kumar \cite{sur}. We use the most recent observational datasets such as 28 points of H(z) data \cite{Farooq:2013hq} and Union2.1 compilation (SN) data \cite{Suzuki:2011hu} (taking into account the full covariance matrix). Here, we also forecast constraints using a simulated data set for the future JDEM, supernovae survey \cite{hol,ald} and also employ Statefinder analysis of the results obtained. | Precision cosmological observations offer the possibility of uncovering essential properties of the Universe. Here, we have investigated power-law cosmology $a(t) \, \propto t^{\beta}$, which has some prominent features, making it unique when compared to other models of the Universe. For example, for $\beta\geq1$, it addresses to the horizon, flatness and age problems \cite{kolb, mann, allen} and all these features provide viability to the power-law cosmology to dynamically solve the cosmological constant problem. In the work presented here, we used the most recent observational data sets from H(z) and SNe Ia observations and obtained the constraints on the two crucial cosmological parameters $H_0$ and $q$ and compared our results with reference \cite{sur}. We also have forecasted these constraints with simulated data for large future surveys like JDEM. Statistically, this model may be preferred over other models as we have to fit only two parameters. Numerical results obtained have been concluded in the Tables \ref{tabparm}, \ref{tabjdem} and \ref{tabstat}. In this work, we observed that though $q$ is negative in the constraints from both H(z) and SNe Ia observations respectively but with bad $\chi_{\delta}^2$ in case of H(z) data, thus we can say that latter explains the present cosmic acceleration more efficiently than that of former in the context of power law cosmology. With the latest H(z) data, we found that obtained best fit value of $H_0$ for power law is outside more than 2$\sigma$ confidence level from the value of $\Lambda$CDM and also, we noticed the large discrepancy between the values of their $\chi^2_{\delta}$. In contrast, with SN data, we found that the value of $H_0$ for power law agrees with $\Lambda$CDM within 1$\sigma$ confidence level and also, its $\chi^2_{\delta}$ is approximately equal to the value of $\chi^2_{\delta}$ for $\Lambda$CDM. Therefore, we conclude that power law model fits well with SN data but not with H(z) data. Contour plots for both the models with H(z) data and SN data have been shown in figures \ref{hsn} and \ref{SNcont} respectively. Best fitted behaviour for power law model with data error bars have been shown in figure \ref{erbar}. On comparing our results with reference \cite{sur}, we observe that in the new analysis, best fit value of $q$ with H(z) data is drastically different from the constraint with SN data but it was not so significant in reference \cite{sur}. This discrepancy has been observed because in the current analysis, power law model does not fit well with latest H(z) data due to significant difference between the value of its $\chi^2_{\delta}$ and $\Lambda$CDM. Also we observe that new SN analysis gives larger error bars on both $q$ and $H_0$ than that of reference \cite{sur} because of taking into account the full covariance matrix. Corresponding variations in the values of statefinders have also been observed which have been summarized in Table \ref{tabstat}. More explicitly, we see the differences in our study and of reference \cite{sur} as: in latter one it had been shown that the power law model fits well with both H(z) and SNe Ia observations but in our case we observed that it fits well only with SN data having larger error bars on both of the parameters and on the contrary it fails to fit with latest H(z) data shifting best fit value of $q$ significantly with bad $\chi_{\delta}^2$. Thus, we can say that our study explains merits and demerits of power law model in explaining the evolution of Universe in a more clear and sophisticated manner than that of reference \cite{sur}. The statefinder diagnostic carried out shows that power law cosmological model will finally approach the $\Lambda$CDM model as shown in figure \ref{rsq}. From the results mentioned in Table \ref{tabjdem}, one can also conclude that future surveys like JDEM demands an accelerated expansion of the Universe but with smaller values of Hubble constant within the framework of power law cosmology. From the above discussed results, it can be concluded that though power law cosmology has several prominent features but still it fails to explain redshift based transition of the Universe from deceleration to acceleration, because here we do not have redshift or time dependent deceleration parameter $q$. Thus, in nutshell it can clearly be said that despite having numerous remarkable features, the power law cosmology does not fit well in dealing with all cosmological challenges. | 14 | 4 | 1404.6522 |
1404 | 1404.1134_arXiv.txt | We present the results of the search for an astrophysical gravitational-wave stochastic background during the second Einstein Telescope mock data and science challenge. Assuming that the loudest sources can be detected individually and removed from the data, we show that the residual background can be recovered with an accuracy of $1\%$ with the standard cross-correlation statistic, after correction of a systematic bias due to the non-isotropy of the sources. | Introduction} The first generation of gravitational-wave (GW) detectors such as LIGO or Virgo (2002-2013) were able to reach their design sensitivities, demonstrating the feasibility of the experiment. With the second generation, Advanced LIGO \cite{aLIGO} and Advanced Virgo \cite{AdVIRGO}, expected to start collecting data in 2015, we will enter the era of the first GW detections. With a sensitivity about 10 times better than that of initial LIGO/Virgo, we expect the detection of a few or a few tens of compact binary coalescences (CBC) a year. With the third generation european antenna Einstein Telescope (ET) \cite{ET,ETdesign} planned to be operational in $\sim 2025$, GW astronomy will definitely take a step further, with the possibility of detecting a large number of sources (up to $10^4-10^5$ CBC a year) from a large range of processes, such as core collapses to neutron stars or black holes \cite{2005PhRvD..72h4001B,2006PhRvD..73j4024S,2009MNRAS.398..293M,2010MNRAS.409L.132Z}, rotating neutron stars \cite{2001A&A...376..381R,2012PhRvD..86j4007R} including magnetars \cite{2006A&A...447....1R,2011MNRAS.410.2123H,2011MNRAS.411.2549M,2013PhRvD..87d2002W}, phase transition \cite{2009GReGr..41.1389D} or initial instabilities in young neutron stars \cite{1999MNRAS.303..258F,2011ApJ...729...59Z,2004MNRAS.351.1237H,2011ApJ...729...59Z} or compact binary mergers \cite{2011ApJ...739...86Z,2011PhRvD..84h4004R,2011PhRvD..84l4037M,2012PhRvD..85j4024W,2013MNRAS.431..882Z} (see \cite{2011RAA....11..369R} and references therein). Besides the emission produced by the coalescence of the nearest binary systems, the superposition of a large number of unresolved sources at high redshifts will produce a background of gravitational waves \cite{2011ApJ...739...86Z,2011PhRvD..84h4004R,2011PhRvD..84l4037M,2012PhRvD..85j4024W,2013MNRAS.431..882Z} that may dominate over the cosmological background in the range $10-1000$ Hz where terrestrial detectors are the most sensitive. The detection of the cosmological background would provide very important constraints on the first instant of the Universe, up to the limits of the Planck era and the Big Bang, while the detection of the astrophysical background would provide crucial information about the star formation history, the mass range of neutron star or black hole progenitors and the rate of compact binary mergers. The issue with ET will not be the detection but rather the estimation of the parameters and the interpretation of the results in term of astronomy, cosmology and fundamental physics. In order to get prepared and test our ability to extract valuable information from the data, we initiated a series of mock data and science challenges, with increasing degree of complexity. For the first ET mock data and science challenge (ET MDSC1) \cite{2012PhRvD..86l2001R}, we produced one month of simulated data containing simulated gaussian colored noise and the GW signal from a simulated population of double neutron stars in the redshift range $z =0-6$. Using a modified version of the LIGO/Virgo data analysis pipeline IHope \cite{2010PhRvD..82j2001A}, we were able to recover the intrinsic chirp mass and total mass distributions with an error of less than 1\% and 5\% respectively. We also analyzed the data with the standard isotropic cross-correlation (CC) statistic and measured the amplitude of the background with an accuracy better than 5\%\footnote{Unlike initial LISA which was designed to be a single detector with three arms in a triangle configuration, ET will consist of three nested detectors (six independent arms in total) \cite{ET,ETdesign}, and one can use cross-correlation methods to extract GW stochastic backgrounds from the instrumental noise.}. Finally, one of our main result was to verify the existence of a null stream canceling the GW signal and giving a very precise estimate of the noise power spectral density (PSD). By subtracting the null stream from the data, we showed that we could recover the typical shape of the PSD of the GW signal. After the success of the first challenge, we extended our data generation package and produced three new sets of data. The first one (ET MDSC2-a) contains all types of stellar compact binary coalescences, composed of two neutron stars (NS-NS), a black hole and a neutron star (BH-NS) or two black holes (BH-BH). The second data set (ET MDSC2-b) contains the population of CBC too faint to be detected individually and which creates a residual GW background. We assume here individual detections can be successfully subtracted from the data, as it has been done with success for the population of white dwarf binaries in the context of the LISA Mock Data Challenge, using Markov Chain Monte Carlo techniques \cite{2008CQGra..25r4026B}.The third data set (ET MDSC2-c) contains the same population of CBC as ET MDSC2-a, two supernovae and two f-modes, and also a population of intermediate-mass black hole binary coalescences and intermediate mass ratio inspiral, which could be quite numerous at low frequencies but whose existence has not been confirmed yet \cite{2011GReGr..43..485G}. In this paper we use the standard cross correlation statistic which is known to be optimal in the case of a Gaussian, isotropic stochastic background to search for the residual GW background in the second data set ET MDSC2-b. This analysis complement the search for individual CBC (Meacher et al., in preparation), as the majority of the sources contributing to the residual background are at redshift above the detection range. The paper will be organized as follow. In section II we present the CBC population model and summarize briefly the simulation procedure. In section III we discuss the spectral properties of the GW signal in the first and second data sets. In section IV we present the results of the analysis. Section V contains a conclusion and suggestions for further research. | Conclusion} In this paper we reported on the analysis of the second Einstein Telescope mock data and science challenge, searching for the residual GW background resulting from the superposition of all the CBC sources that are too faint to be detected individually. We used the standard cross correlation statistic which is known to be optimal in the case of a Gaussian, isotropic stochastic background. Confirming the results of the ET MDSC1 and the recent work of Meacher et al., we obtained that the non continuity or non gaussianity of the background \cite{2011PhRvD..84h4004R,2012PhRvD..85j4024W,2012arXiv1205.4621K,2012PhRvD..86l2001R} does not significantly affect the analysis (what's important is the total number of sources and not whether they overlap or not). But because of the GW selection effect that favored the detection of the best oriented and located sources, especially at larger redshift, the assumption of an isotropic stochastic background is not verified and the estimate given by the standard cross correlation statistic presents a systematic bias in the case of the residual background. Deriving a correction for the overlap reduction function we obtained a point estimate that agrees with the expected value with a precision $<1\%$. The detection of the residual background would have very important consequences in cosmology and astrophysics as it would probe the high redshift population, complementing individual detections at smaller redshift. The residual background from CBC may dominate in the frequency band of ET. In future ET MDSC, we will investigate how one can use the non Gaussian signature to separate this background or foreground and recover the cosmological background. | 14 | 4 | 1404.1134 |
1404 | 1404.4182_arXiv.txt | {The $v$=1 and $v$=2 \juc\ (43 GHz), and $v$=1 \jdu\ (86 GHz) SiO masers are intense in Asymptotic Giant Branch (AGB) stars and have been mapped using Very Long Baseline Interferometry (VLBI) showing ring-like distributions. Those of the $v$=1, $v$=2 \juc\ masers are similar, but the spots are rarely coincident, while the $v$=1 \jdu\ maser arises from a well separated region farther out. These relative locations can be explained by models tools that include the overlap of two IR lines of SiO and H$_2$O. The $v$=3 \juc\ line is not directly affected by any line overlap and its spot structure and position, relative to the other lines, is a good test to the standard pumping models. } {The aim of this project are to gain insight into the properties and the general theoretical considerations of the different SiO masers that can help to understand them.} {We present single-dish and simultaneous VLBI observations of the $v$=1, $v$=2, and $v$=3 \juc\ maser transitions of $^{28}$SiO in several AGB stars. The results are compared to the predictions of radiative models of SiO masers that both include and not include the effect of IR line overlap. } {The spatial distribution of the SiO maser emission in the $v$=3 \juc\ transition from AGB stars is systematically composed of a series of spots that occupy a ring-like structure (as often found in SiO masers). The overall ring structure is extremely similar to that found in the other 43 GHz transitions and is very different from the structure of the $v$=1 \jdu\ maser. The positions of the individual spots of the different 43 GHz lines are, however, very rarely coincident, which in general is separated by about 0.3 AU (between 1 and 5 mas). These results are very difficult to reconcile with standard pumping models, which predict that the masers of rotational transitions within a given vibrational state require very similar excitation conditions (since the levels are placed practically at the same energy from the ground), while the transitions of different vibrational states (which are separated by an energy of 1800 K) should appear in different positions. However, models including line overlap tend to predict $v$=1, $v$=2, $v$=3 \juc\ population inversion to occur under very similar conditions, while the requirements for $v$=1 \jdu\ appear clearly different, and are compatible with the observational results. } {} | Many Asymptotic Giant Branch (AGB) stars have been mapped in SiO maser emission in the \juc\ $v$=1 and $v$=2 lines\footnote{In this paper, $v$=1, $v$=2, $v$=3 refers to masers at the $v$=1, $v$=2, or $v$=3 states.} using Very Long Baseline Interferometry (VLBI) \citep[][etc]{diam94, desmurs00, cot06}. The maser emission is found to form a ring of spots at a few stellar radii from the center of the star. In general, both distributions are similar, although the spots are very rarely coincident, and the $v$=2 ring is slightly closer to the star than the $v$=1 ring \citep[see e.g.][]{desmurs00}. The similar distributions of the $v$=1, $v$=2 \juc\ transitions were first interpreted as favoring collisional pumping, because the radiative mechanism discriminates the location of the two masers more. On the contrary, the lack of true coincidence was used to argue in favor of radiative pumping, which leads to the well-known, long-lasting discrepancy in the interpretation of the $v$=1, $v$=2 \juc\ maps in terms of pumping mechanisms; see detailed discussion in \cite{desmurs00}. Our understanding of this topic changed dramatically when the first comparisons between the $v$=1 \juc\ and $v$=1 \jdu\ maser distributions were performed \citep{soria04,soria05,soria07}. In contrast to predictions from both models (radiative and collisional), the $v$=1 \jdu\ maser spots are systematically found to occupy a ring with a significantly larger radius (by about 30\%) than that of $v$=1 \juc, where both spot distributions being completely unrelated. \cite{soria04} explained these unexpected results by invoking line overlap between the ro-vibrational transitions $v$=2--1 \juc\ of SiO and $v_2$=0--1 $J_{K_a,K_c}$=$12_{7,5}$--$11_{6,6}$ of H$_2$O. According to \cite{soria04}, this phenomenon, which was first proposed by \cite{olof81} to explain the anomalous weakness of the $v$=2 \jdu\ SiO maser, would also introduce a strong coupling of the $v$=1 and $v$=2 \juc\ lines, explaining their similar distributions. In the simplest theoretical interpretation \citep[not including line overlap, see][]{buj81,buj94,loc92,hum02}, the $v$=3 \juc\ emission requires completely different excitation conditions than the other less excited lines, since the $\Delta v$=1 energy separation is very high, $\sim$1800~K. The $v$=3 \juc\ spatial distribution should in principle be different from the $v$=1, $v$=2 \juc\ ones and, of course, the $v$=1 \jdu\ maser, and placed in a still more inner ring. However, we have seen that line overlap strongly affects the $v$=1 and $v$=2 \juc\ maser pumping. In particular, this phenomenon changes the conditions required to pump both lines, which now tend to require higher densities. In this paper, we compare the $v$=1, $v$=2, $v$=3 \juc\ maser distribution and analyzed the results in the framework of the pumping models. | We have observed the VLBA four AGB stars, R Leo, TX Cam, U Her, and IK Tau, and obtained reliable maps of \juc\ SiO masers in the first three vibrationally excited states ($v$=1, $v$=2, and $v$=3) toward the four sources. We find that the spatial brightness distribution of the $v$=3 maser does not show significant differences with respect to those of the $v$=1 and $v$=2 lines. The $v$=3 maser emission is distributed in a ring-like pattern and is coincident or slightly inner than those of $v$=1, $v$=2 \juc. Despite our initial expectation, this agrees with model predictions and can be easily explained by the range of physical conditions that give rise to the $v$=1, $v$=2, $v$=3 maser lines, which are predicted when the effects of the overlaping of two IR lines of SiO and H$_2$O are taken into account (Sect.\ 3). When line overlap is not taken into account, the observed distributions of the SiO maser lines cannot be explained by current models invoking either collisional or radiative maser pumping. The excitation conditions are very different for lines within the different vibrationally excited states ($v$=1, $v$=2, and $v$=3), which are separated by an energy equivalent to almost 1800 K. On the other hand, the conditions required to excite the $v$=1 \juc\ and \jdu\ lines are almost the same, since the energy levels are separated by a few degrees. Since standard inversion schemes do not significantly discriminate the low-$J$ levels, we expect, clearly different spot distributions for the masers in different $v$-states and very similar distributions for masers in the same state under both collisional or radiative pumping mechanisms. However, the overlap of the above mentioned IR lines significantly affects the pumping of the $v$=1, $v$=2 \juc: our calculations show that these two masers are then strongly coupled and require higher excitation conditions, which is similar to those of the $v$=3 lines (which are not significantly affected by the line overlap); see details in Sect.\ 3. Therefore, the $v$=1, $v$=2, $v$=3 \juc\ lines require quite similar excitation conditions and should appear in practically the same circumstellar shells due to the overlap effects. However, the $v$=1 \jdu\ maser, which is practically not affected by the considered pair of IR lines, requires lower excitation condition, and should appear in outer shells. Under the physical conditions adopted in our models we stress that the pumping of the SiO masers is mainly radiative and, indeed, that the effects of line overlap tend to be more important in the radiative pumping regime than for the collisional one. However, we cannot rule out that collisional models, includes the effects of line overlap, could also explain the relative spatial distributions of the different maser lines. The predictions presented here are, therefore, compatible with the existing maps of these four maser lines, notably for $J$=1--0 presented in this work and for our previous $v$=1 $J$=2--1 maps \citep{soria04,soria05,soria07}. Even the small (but systematic) differences found between the radii at which the three $J$=1--0 maser emissions appear in the maps are qualitatively compatible with the model predictions. | 14 | 4 | 1404.4182 |
1404 | 1404.1558_arXiv.txt | In this brief report, we try to constrain general parameterized forms of scalar and tensor mode power spectra, $P_{s}(k)\equiv A_s(k/k_0)^{n_s-1+\frac{1}{2}\alpha_s\ln(k/k_0)}$ and $P_{t}(k)\equiv A_t(k/k_0)^{n_t+\frac{1}{2}\alpha_t\ln(k/k_0)}$ by the recently released BICEP2 data set plus {\it Planck} 2013, WMAP9 and BAO. We loosen the inflationary consistence relations, and take $A_s$, $n_s$, $A_t$ and $n_t$ as free model parameters, via the Markov chain Monte Carlo method, the interested model parameter space was investigated, we obtained marginalized $68\%$ limits on the interested parameters are: $n_s=0.96339_{-0.00554}^{+0.00560}$, $n_t=1.70490_{-0.56979}^{+0.56104}$, ${\rm{ln}}(10^{10} A_s)=3.08682_{-0.02614}^{+0.02353}$ and ${\rm{ln}}(10^{10} A_t)=3.98376_{-0.54885}^{+0.86045}$. The ratio of the amplitude at the scale $k=0.002 \text{Mpc} ^{-1}$ is $r=0.01655_{-0.01655}^{+0.00011}$ which is consistent with the {\it Planck} 2013 result. | The BICEP2 experiment \cite{ref:BICEP21,ref:BICEP22} has detected the B-modes of polarization in the cosmic microwave background. And this observed B-modes power spectrum gives the constraint to the tensor-to-scalar ratio with $r=0.20^{+0.07}_{-0.05}$ at the $1\sigma$ level of the lensed-$\Lambda$CDM model \cite{ref:BICEP21,ref:BICEP22}. And the tensor spectral tilt $n_t$ can be obtained, when the first oder consistency relation, $n_t=-r/8$, was respected. Also, relaxing this consistency relation by taking $n_t$ as a free model parameter \cite{ref:Huang2014}, $r_{0.002}=0.21^{+0.04}_{-0.10}$ and $n_t=-0.06^{+0.25}_{-0.23}$ were obtained by using BICEP2 only. By combining {\it Planck}, WMAP9 and BAO data, it was already found that a blue tilt is slightly favored, but it is still well consistent with flat or red tilt \cite{ref:Wu2014}. However, one can go further by taking generalized parameterized forms of scalar and tensor mode power spectra as \begin{eqnarray} P_{s}(k)&\equiv& A_s(k/k_0)^{n_s-1+\frac{1}{2}\alpha_s\ln(k/k_0)},\label{eq:ps}\\ P_{t}(k)&\equiv& A_t(k/k_0)^{n_t+\frac{1}{2}\alpha_t\ln(k/k_0)},\label{eq:pt} \end{eqnarray} without assuming any idea about inflation, in other words throwing away the consistence relations, just considering the possible deviation from the scale invariant power spectra, i.e. the Harrison-Zel'dovich-Peebles spectra. And how to interpreter it is another issue. Of course, one can relate it to the so-called inflation, where the consistence relations should be respected. And in this way, one can test the viability of inflation models. But, one can also explain it through the bounce expansion. Here $n_s-1$ and $n_t$ are tilts of power spectrum of scalar and tensor modes, $k_0=0.05 \text{Mpc} ^{-1}$ is the pivot scale, $\alpha_s=d n_s/d\ln k$ and $\alpha_t=d n_t/d\ln k$ are the running of the scalar and tensor spectral tilts. The primordial tensor-to-scalar ratio is defined by $r\equiv A_t/A_s$ at different pivot scale, here, they are $r_{0.05}$ defined at $k_0=0.05 \text{Mpc} ^{-1}$ and $r_{0.002}$ defined at $k_0=0.002 \text{Mpc} ^{-1}$. In this paper, without any other specification, $r_{0.002}$ will be donated by $r$. And we also denote $A_t/A_s$ as the amplitude ratio of the tensor and scalar mode power spectrum at $k\equiv k_0$, i.e. the scale independent tensor-to-scalar ratio. In our calculation, adiabatic initial conditions were assumed in this paper. Actually, if one wants to relate the parameterized primordial power spectra, the following relations, the so-called consistency relation should be respected \cite{ref:Mukhanov1999} \begin{equation} r=-8c_s n_t,\label{eq:consnt} \end{equation} By taking these parameters, $n_t$ and $A_t$ as free ones, one can test these consistency relation by the recently released BICEP2 data. Here, we are mainly focusing on the model parameters which are related to the primordial power spectra. Therefore, in this brief paper, by combing the following data sets, we report the constrained results on the interested parameters: (i) The newly released BICEP2 CMB B-mode data \cite{ref:BICEP21,ref:BICEP22}. (ii) The full information of CMB which include the recently released {\it Planck} data sets which include the high-l TT likelihood ({\it CAMSpec}) up to a maximum multipole number of $l_{max}=2500$ from $l=50$, the low-l TT likelihood ({\it lowl}) up to $l=49$ and the low-l TE, EE, BB likelihood up to $l=32$ from WMAP9, the data sets are available on line \cite{ref:Planckdata}. (iii) For the BAO data points as 'standard ruler', we use the measured ratio of $D_V/r_s$, where $r_s$ is the co-moving sound horizon scale at the recombination epoch, $D_V$ is the 'volume distance' which is defined as \begin{equation} D_V(z)=[(1+z)^2D^2_A(z)cz/H(z)]^{1/3}, \end{equation} where $D_A$ is the angular diameter distance. The BAO data include $D_V(0.106) = 456\pm 27$ [Mpc] from 6dF Galaxy Redshift Survey \cite{ref:BAO6dF}; $D_V(0.35)/r_s = 8.88\pm 0.17$ from SDSS DR7 data \cite{ref:BAOsdssdr7}; $D_V(0.57)/r_s = 13.62\pm 0.22$ from BOSS DR9 data \cite{ref:sdssdr9}. Here the BAO measurements from WiggleZ are not included, as they come from the same galaxy sample as $P(k)$ measurement. We will present the method and obtained results in the next section \ref{sec:results}. Section \ref{sec:conclusion} is the conclusion. | \label{sec:conclusion} In this brief paper, we loosen the inflation consistency relation constraint, and take the spectral tilts $n_s$, $n_t$, $A_s$ and $A_t$ as free model parameters. Combining the recently released BICEP2 data, {\it Planck} 2013, WMAP9 and BAO via the MCMC method, the model parameter space was scanned. We found that $n_s=0.96339_{-0.00554}^{+0.00560}$, $n_t=1.70490_{-0.56979}^{+0.56104}$, ${\rm{ln}}(10^{10} A_s)=3.08682_{-0.02614}^{+0.02353}$ and ${\rm{ln}}(10^{10} A_t)=3.98376_{-0.54885}^{+0.86045}$. The ratio of the amplitude at the scale $k=0.002 \text{Mpc} ^{-1}$ is $r=0.01655_{-0.01655}^{+0.00011}$ which is consistent with the {\it Planck} 2013 result. And a blue tensor tilt is favored at $1\sigma$ C.L.. And $n_t$ is positive above $2\sigma$ C.L.. It implies the broken of consistency relation $r=-8c_s n_t$ at $2\sigma$ C.L., when the speed of sound $c_s>0$ is respected. | 14 | 4 | 1404.1558 |
1404 | 1404.3288_arXiv.txt | Particle acceleration is one of the most significant features that are ubiquitous among space and cosmic plasmas. It is most prominent during flares in the case of the Sun, with which huge amount of electromagnetic radiation and high-energy particles are expelled into the interplanetary space through acceleration of plasma particles in the corona. Though it has been well understood that energies of flares are supplied by the mechanism called magnetic reconnection based on the observations in X-rays and EUV with space telescopes, where and how in the flaring magnetic field plasmas are accelerated has remained unknown due to the low plasma density in the flaring corona. We here report the first observational identification of the energetic non-thermal electrons around the point of the ongoing magnetic reconnection (X-point); with the location of the X-point identified by soft X-ray imagery and the localized presence of non-thermal electrons identified from imaging-spectroscopic data at two microwave frequencies. Considering the existence of the reconnection outflows that carries both plasma particles and magnetic fields out from the X-point, our identified non-thermal microwave emissions around the X-point indicate that the electrons are accelerated around the reconnection X-point. Additionally, the plasma around the X-point was also thermally heated up to 10~MK. The estimated reconnection rate of this event is $\sim$~0.017. | \label{sec:Introduction} Particle acceleration is ubiquitously observed among space and cosmic plasmas, including those around the Earth's magneto-tail \citep[{\it e.g.},][]{oie02}, in supernova remnants \citep[{\it e.g.},][]{koy95}, and even in distant galaxies \citep[{\it e.g.},][]{mus77}. For our Sun, acceleration of plasma particles is most prominent in flares, with which vast amounts of electromagnetic radiation from accelerated particles \citep[{\it e.g.},][]{sak96, yok02}, as well as the particles themselves \citep[{\it e.g.},][]{lin96}, are expelled into interplanetary space. Observations from space in the past few decades have established that it is magnetic reconnection during flares that converts magnetic energy in the corona into the kinetic and thermal energies of the plasma particles \citep[{\it e.g.},][]{tsu92, mas94, yok01, har11, ima13, su13}. It is thus expected that particle acceleration in flares should also be closely related to the magnetic reconnection process through which as much as half of the liberated energy is converted into acceleration of particles \citep{lin76}. There have been a number of theoretical studies made so far that have considered various portions of the reconnecting magnetic structure as the site of electron acceleration in solar flares. These include, inside the closed magnetic loop formed by reconnection \citep[{\it e.g.},][]{fle08}, in the magneto-hydrodynamic fast-shock structure expected to be formed above the closed loop \citep[{\it e.g.},][]{tsu98}, in the magnetic cusp region where field lines are contracting downward \citep[{\it e.g.},][]{som97}, and around the X-point (or in the current sheet) \citep[{\it e.g.},][]{lit96,pri06,dra06,oka10}. Meanwhile, some observations have addressed possible site of the electron acceleration. \citet{asc96} studied timing relationship among hard X-ray emissions in different energies and quantitatively estimated locations of electron acceleration region as above the flaring loops. \citet{sak98} argued that electron acceleration should take place at the reconnection site which is a common place among the two different magnetic field configurations identified from foot point motions of hard X-ray sources during the impulsive phases. These observations gave us the supposition that the particle acceleration would occur above the flaring loops. However, there have not been any observations that pinpoint where and how in the flaring magnetic field plasmas are accelerated. This is partly due to the fact that magnetic reconnection in flares takes place in the high corona where plasma density is low, and hard X-ray fluxes from energetic non-thermal electrons, which are emitted by Bremsstrahlung (interaction with ambient coronal plasma), are too weak to be imaged by existing modulation-collimator type hard X-ray telescopes whose dynamic range is not high, especially for the reconnection region. Meanwhile, there is a possibility that microwaves emitted from the energetic non-thermal electrons by gyro-syncrotoron mechanism (interaction with coronal magnetic fields) can be detected even in the low density corona including the reconnection region. In this paper, we report the first identification of energetic non-thermal electrons around the reconnection X-point in a flare; with the location of the X-point identified by soft X-ray imagery and the localized presence of non-thermal electrons identified from imaging-spectroscopic data at two microwave frequencies. This is a strong evidence of the electrons being accelerated around the X-point, providing an observational clue toward understanding the mechanism of electron acceleration in solar flares. | \label{sec:Discussion} We have shown that there is microwave emission from non-thermal electrons most clearly seen around the X-point. This indicates that there are energetic electrons present around the X-point during the course of the magnetic reconnection. Meanwhile, the soft X-ray observations of the flare were in good agreement, in a morphological sense, with the canonical reconnection picture (Figure~\ref{fig:data}~(a)) from which we expect, in addition to the inflow, bi-directional outflow of plasma particles and magnetic fields away from the X-point towards the bright soft X-ray loop and in the opposite direction \citep{shi11}. Assuming the presence of the expected reconnection outflow, it would be reasonable to conclude that the energetic non-thermal electrons are supplied from, namely, accelerated at, the region around the X-point rather than assuming that they travel from a region lower in altitude than the X-point against the counter-streaming downward outflow, or that they come from a region higher in altitude beyond the X-point against the counter-streaming upward outflow from the X-point. The observed decreasing trend in the column emission measure (Figure~\ref{fig:sxr}~(c)), {\it i.e.}, a decrease in density of soft-X-ray-emitting thermal electrons, would be due to the ambient electrons being swept away from the X-point by the reconnection outflow and/or that they departed from thermal equilibrium due to collisions with accelerated electrons \citep{kru10}. Furthermore, as seen in Movie~4, the negative alpha signal traveled from around the X-point to the foot point of the flare loop along the northern outer edge of the flare loop around the timing of the plasmoid ejection. This gives further support that the accelerated electrons manifesting themselves as the negative alpha signal are originated from around the X-point, although the traveling negative alpha signal does not immediately correspond to the individual energetic electrons. Hence, we argue that electron acceleration around the X-point is the initial (first-stage) acceleration in a flare that immediately follows the onset of reconnection. We further take a closer look into the region around the X-point. The spatial extent of the negative alpha index in the lateral direction (parallel to the solar surface) is possibly caused by the inclined configuration of the flaring loop to the line-of-sight as reported by \citet{nar06}. Meanwhile, the length of the localized negative alpha index, {\it i.e.}, the region where the energetic non-thermal electrons exists, has a spatial extent of $\sim$~20,000 km along the direction of the height, and is located at the top-most of the cusp structure, for the period of ongoing reconnection (see Figure~\ref{fig:data}~(l)). Considering this spatial extent along the direction of the height, the possible scenarios of the electron acceleration around the X-point are follows: (i) An X-point alone can produce energetic electrons with the inductive reconnection electric field \citep[{\it e.g.},][]{pri06}. Such energetic electrons can easily escape from the X-point, and emit microwaves from a spatially extended region as seen in our data. (ii) Multi-island coalescence along a single current sheet should also be considered as an electron acceleration scenario \citep[{\it e.g.},][]{oka10}. (iii) Fragmented islands scenario with multiple current sheets \citep[{\it e.g.},][]{dra06, shi01} is also possible, since, considering the spatial resolution of our microwave data ($\sim$ 10$\arcsec$), the multiple current sheets cannot be resolved, and the non-thermal signals in microwaves should be observed as only one source as seen in the presented data. We note that the hot thermal electrons ($>$ 10~MK) spatially coexists with the energetic non-thermal electrons (Figure~\ref{fig:data}~(e)--(g)) around the X-point. Since the coronal temperature below the X-point are cooler than around the X-point, we can say that the hot plasmas around the X-point is not heated by the thermal conduction from the plasmas located below the X-point. These high temperature plasmas are directly created by the released magnetic energy that would also accelerate the electrons, and/or are the result of the thermalization of the accelerated electrons. In our flare, we also find microwave non-thermal sources other than around the X-point (see the negative alpha index distribution in Figure~\ref{fig:data}, especially in Figure~\ref{fig:data}~(m)). One is located at the foot points of the flaring loop. This negative alpha index is also due to the non-thermal emission, not by the gyro-resonance emission, since the sunspots were occulted by the solar disk in this flare. The other non-thermal source is distributed above the top of the flaring loop extending from the X-point region. These two kinds of observed microwave non-thermal sources became noticeable around the timing of the plasmoid ejection (see Figure~\ref{fig:data}~(m)). This is consistent with the timing where intensity enhancement in the hard X-rays was observed around the timing of plasmoid ejections \citep[{\it e.g.},][]{ohy98}. Hence, these two kinds of the sources may have close relationship with the well known hard X-ray non-thermal sources, {\it e.g.}, double-foot-point sources \citep[{\it e.g.},][]{sak96} and above-the-loop-top source \citep[{\it e.g.},][]{mas94}. Considering the spatially-smooth distribution of the negative alpha index extending from the X-point region (larger negative alpha index $\sim -1$) to above-the-loop-top region (smaller negative alpha index $\sim 0$) at the timing when the microwave non-thermal sources became noticeable (see Figure~\ref{fig:data}~(m)), it is reasonable to expect that no major acceleration site other than around the X-point exists, although some additional sub-accelerations might occur in other locations. The hard X-ray above-the-loop-top source \citep[{\it e.g.},][]{mas94} may appear as a lower part of the spatial distribution of non-thermal electrons from the X-point region to above the loop top manifesting themselves in microwaves (with negative alpha index), and is detected due to the denser coronal plasmas near the flaring loop with which Bremsstrahlung hard X-rays are efficiently emitted. Meanwhile, the reconnection rate can be estimated to be $\sim$~0.017 from the data set of following parameters considering the line-of-sight effect \citep{nar06}: the inflow velocity ($\sim$~16.3 km s$^{-1}$) derived from Figure~\ref{fig:microwave}~(a), the foot point expanding velocity ($\sim$~3.3 km s$^{-1}$) of the flare loop from X-ray data, the averaged magnetic field strength at the sunspot ($\sim$~98 Gauss) from the magnetogram, and the coronal column emission measure ($\sim 10^{28.4}$ cm$^{-5}$ that corresponds to the number density of $10^{9.3}$ cm$^{-3}$ for the line-of-sight depth of 60,000 km) from the soft X-ray data with the filter ratio method. This reconnection rate supports the Petschek type reconnection against Sweet-Parker type \citep{pri82}. The solar corona provides a unique opportunity for investigating the acceleration process associated with magnetic reconnection, not only in that the entire view of the reconnecting magnetic structure can be obtained by imagery, but also because the spatio-temporal evolution of any sources with certain spectral features in the dynamically-evolving reconnecting structure can be traced by imaging-spectroscopic approach. With the present result being the first step, progress in theoretical investigations as well as future observations such as from Atacama Large Millimeter/submillimeter Array (ALMA) \citep{bas02} with very high spatial resolution, mirror-focusing hard X-ray \citep[FOXSI;][]{kru13} and soft X-ray \citep{sak12} telescopes with photon-counting capabilities should help us solve the long-standing questions about particle acceleration in solar flares. | 14 | 4 | 1404.3288 |
1404 | 1404.0148_arXiv.txt | Dark matter is one of the most important scientific goals for neutrino telescopes. These instruments have particular advantages with respect to other experimental approaches. Compared to direct searches, the sensitivity of neutrino telescopes to probe the spin-dependent cross section of WIMP-proton is unsurpassed. On the other hand, neutrino telescopes can look for dark matter in the Sun, so a potential signal would be a strong indication of dark matter, contrary to the case of other indirect searches like gammas or cosmic rays, where more conventional astrophysical interpretations are very hard to rule out. We present here the results of a binned search for neutralino annihilation in the Sun using data gathered by the ANTARES neutrino telescope during 2007-2008. These result include limits on the neutrino and muon flux and on the spin-dependent and spin-independent cross section of the WIMP-proton scattering. | \label{intro} Dark matter existence has been soundly proofed by different experimental evidence, including the observations from Planck~\cite{planck}, the results on the Big Bang Nucleosynthesis~\cite{jedamzik}, the rotation curves of galaxies~\cite{rubin} and the studies of highly red-shifted Ia supernovae~\cite{kowalski}. These results show that the only about 30\% of the content of the Universe is matter and about 70\% is dark energy. Moreover, 85\% of the matter is non-barionic. Explanations for the nature of this non-barionic component have to be outside the Standard Model. The basic conditions required to a particle dark matter candidate are to have interaction cross section of the order of that of the weak interaction and to be massive and stable. Particles like neutrinos, which fulfill these requirement for a good dark matter candidate, are not viable as a dominant component, since they are relativistic and cannot explain the large-scale structure of the Universe. Given these contiditons, a generic familly of particles fulfilling these conditions are callend WIMPs (Weakly Interacting Massive Particles). The most popular model which provides WIMP candidates is Supersymmetry (SUSY). In particular, in this analysis we have looked for neutralinos, which in many of the possible scenarios is the lightest SUSY particle. Its stability is preserved by the conservation of the R-parity. The results have been analyzed with respect to two implementanions of the SUSY framework: CMSSM~\cite{cmssm} and MSSM-7~\cite{mssm7}. The analysis presented here is a binned search for neutrinos produced after the neutralino annihilations in the Sun direction using 2007-2008 data of the ANTARES neutrino telescope, since neutralinos would accumulate in massive objects like the Sun and their annihilation would produce high energy neutrinos~\cite{paper}. One of the advantages of this kind of searches, compared to other indirect searches like looking for gammas in the Galactic Center or excesses of positrons is that a potential signal would be a very robust indication of dark matter, since no other astrophysical explanations are expected. The structure of this paper is as follows. The ANTARES detector is introduced in Section~\ref{antares}. The estimations for signal and background are described in Section~\ref{simulation}. Section~\ref{cuts} explains the optimization procedure. Finally, the results are presented in Section~\ref{results} and the conclusions are summarized in Section~\ref{conclusions}. | \label{conclusions} The ANTARES data corresponding to 2007-2008 have been used to search for an excess of high energy neutrinos in the Sun's direction, which could indicate annihiliation of dark matter particles like neutralinos. The analysis is a binned search that has shown no excess with respect to the expectations from background. Upper limits both in the neutrino and muon flux have been set. Assuming that equilibrium between capture and annihilation has been reached in the Sun, these limits can be translated into limits in the spin dependent and spin independent cross section of the WIMP-proton scattering. The results for spin dependent cross section are particularly competitive with respect to direct search experiments. A comparison with the parameter space allowed by the CMSSM and MSSM-7 models has been shown. | 14 | 4 | 1404.0148 |
1404 | 1404.5301_arXiv.txt | We explore the distribution of cool ($\sim10^4\,$K) gas around galaxies and its dependence on galaxy properties. By cross-correlating about 50,000 \mgii\ absorbers with millions of sources from the SDSS (optical), WISE (IR), and GALEX (UV) surveys we effectively extract about 2,000 galaxy-absorber pairs at $z\sim0.5$ and probe relations between absorption strength and galaxy type, impact parameter and azimuthal angle. We find that cool gas traced by \mgii\ absorbers exists around both star-forming and passive galaxies with a similar incidence rate on scales greater than {100}\,kpc but each galaxy type exhibits a different behavior on smaller scales: \mgii\ equivalent width does not correlate with the presence of passive galaxies whereas stronger \mgii\ absorbers tend to be found in the vicinity of star-forming galaxies. This effect is preferentially seen along the minor axis of these galaxies, suggesting that some of the gas is associated with outflowing material. In contrast, the distribution of cool gas around passive galaxies is consistent with being isotropic on the same scales. We quantify the average excess \mgii\ equivalent width $\langle \delta W_{0}^{\rm Mg\,II}\rangle$ as a function of galaxy properties and find $\langle\delta W_{0}^{\rm Mg\,II}\rangle \propto SFR^{1.2}$, $sSFR^{0.5}$ and $M_\ast^{0.4}$ for star-forming galaxies. This work demonstrates that the dichotomy between star-forming and passive galaxies is reflected in the circumgalactic medium traced by low-ionized gas. We also measure the covering fraction of \mgii\ absorption and find it to be about 2-10 times higher for star-forming galaxies than passive ones within 50\,kpc. We estimate the amount of neutral gas in the halo of $\langle \log M_{\ast}/{\rm M_\odot} \rangle \sim10.8$ galaxies to be a few $\times\, 10^9\,{\rm M_\odot}$ for both types of galaxies. Finally, we find that correlations between absorbers and sources detected in the UV and IR lead to physical trends consistent with those measured in the optical. | The circum-galactic medium (CGM), tracing baryons surrounding galaxies within their own dark matter halos, plays an important role in galaxy formation and evolution. This interface between galaxies and the inter-galactic medium (IGM) is known to harbor gas flows as accretion and/or outflows but its properties and dependence on galaxy properties are poorly known. For several decades, metal absorption lines imprinted in the spectra of background sources have been the main tool to probe the gas distribution with the CGM. Studies have made use of galaxy-absorber pairs, for which the galaxy has been spectroscopically confirmed to be close to the redshift of the absorber. Various absorption lines can be used for such analyses. Among them the \mgiidoublet\ absorption lines, which trace cool gas (T $\sim10^{4}$ K), have been extensively used due to their strength and visibility from the ground: $0.3<z<2.5$ in the optical range. From the first detection of a galaxy-\mgii\ absorber pair by \citet{Bergeron1986}, to the collection of a sample of about 50 systems by \cite{Steidel1994} to the latest extension to about 200 pairs by \citet{NielsenMgIIcatalogI}, numerous authors have attempted to find correlations between galaxy and absorber properties but the large scatter typically observed among such quantities has mostly led to non-detections \citep[e.g.][]{Kacprzak2011}. Another line of investigation has been based on statistical analyses of large surveys \citep[e.g.][]{Zibetti2007,MenardOII,BordoloiMgII}. While some interesting trends have emerged from these measurements, the connection between galaxies and the distribution of cool gas in the CGM is, after more than two decades of studies, still poorly constrained. To provide some guidance to the theoretical understanding of the CGM, additional observational constraints that can possibly reveal the physical connections between absorber and galaxy properties are needed. In this paper, we statistically extract about 2000 galaxy-absorber pairs using data from the Sloan Digital Sky Survey (SDSS; \citealt{YorkSDSS}). These pairs can then be used to measure correlations between absorber and galaxy properties such as star formation rate (SFR), stellar mass, and azimuthal angle as a function of impact parameter. We also perform our analysis to data from the GALEX and WISE surveys and provide additional support to our findings using UV and IR data. The outline of the paper is as follows: we present the datasets and statistical estimators in Section 2 and the results of galaxy-absorber correlations in Section 3. We discuss our findings in Section 4 and summarize our results in Section 5. In this analysis, we adopt flat $\Lambda$CDM cosmology with $h=0.7$ and $\Omega_{M}=0.3$. Throughout this work we use AB magnitude system and unless stated otherwise scales are in physical units. \begin{figure*} \begin{center} \includegraphics[viewport=20 350 900 920,scale=0.45]{140717_DF_with_median_redshift_value.pdf} \caption{\emph{Top}: SDSS, WISE, and GALEX mean number of galaxies per absorber as a function of $z_{\rm abs}$ from $0.4<z_{\rm abs}<2.8$. We search galaxies with impact parameter from 10 to 200 kpc in SDSS and 50 kpc to 200 kpc in WISE and GALEX. Three surveys all detect galaxies associated with \mgii\ absorbers up to $z\sim1$. From redshift 0.4 to 0.6, SDSS detects 70\% of galaxies associated with \mgii\ absorbers, WISE detects 30\% and GALEX detects 4\%. \emph{Bottom}: The redshift distributions of DR7 and DR9 \mgii\ absorbers.} \label{fig:3surveys} \end{center} \end{figure*} | We have shown that cool gas traced by \mgii\ absorption is found around both star-forming galaxies and passive galaxies, with a similar incidence rate at impact parameters greater than about 50 kpc. In contrast, at smaller impact parameters we find that the strength of \mgii\ absorption depends on the level of star formation of the central galaxy. Our results are consistent with other types of observational constraints on the galaxy-absorber connection: \citet{Zibetti2007} used stacked SDSS images to measure the mean flux excess correlated with the presence of \mgii\ absorbers and found stronger \mgii\ absorbers to be preferentially associated with bluer emission. \citet{BordoloiMgII} used stacked galaxy spectra to measure the mean \mgii\ absorption induced by the presence of galaxies along the line-of-sight and showed that the mean \mgii\ absorption strength around star-forming galaxies is higher than around passive galaxies. Our analysis has, in addition, allowed us to show that while \mgii\ absorption is commonly found around passive galaxies, no correlation between equivalent width and galaxy properties can be detected. This is in contrast to the relation observed between star-forming galaxies and absorbers where the mean equivalent width depends on the color and/or star formation rate of the galaxies. The basic dichotomy between star-forming and passive galaxies is therefore reflected in the properties of the cool gas in the CGM. We also find that strong absorbers tend to be found along the minor axis of star-forming galaxies at small impact parameters, indicating that some of the gas detected in absorption is likely associated with outflows from galaxies. This is consistent with what was reported by \citet{BordoloiMgII}, \citet{BoucheMgII} and \citet{KacprzakAZ}, and is similar to the measurements of \caii\ absorption by \citet{ZhuCaII}. Our results also show that at impact parameters greater than 50\,kpc, the covering fraction of strong \mgii\ absorbers declines with scale proportionally to the overall mass distribution for both star-forming and passive galaxies. Here we note that \citet{ZhuMgII} found a similar result by measuring the \emph{total} \mgii\ absorption around SDSS luminous red galaxies. By measuring the mean absorption correlated with the presence of such galaxies, these authors showed that the average \mgii-to-dark matter ratio is roughly scale-independent. Strong \mgii\ absorbers therefore trace the overall \mgii\ absorption field on large scales. \subsection{The connection to baryons in galaxy halos} It is interesting to point out the contrast between the ubiquitous presence of cool, low-ionized gas traced by \mgii\ around both star-forming and passive galaxies and the distribution of highly ionized gas traced by \ovi\ absorption which currently is only detected in the halo of star-forming galaxies \citep{TumlinsonOVI}. However, despite such a difference, the two gas distributions share a common property: excess absorption is detected when the galaxy specific star formation rate $sSFR\gtrsim 10^{-11}\,{\rm yr}^{-1}$. This threshold is found to be the same for both phases. It appears to be a characteristic sSFR value for determining the gaseous properties of the CGM. Similarly, we detect excess absorption when the galaxy star formation rate ${\rm SFR}\gtrsim 1\,\rm M_\odot/yr$ and note that this threshold is similar to the one for which blueshifted \mgii\ self-absorption is seen in galaxies \citep[e.g.][]{Weiner2009,Rubin2013, Bordoloioutflow}. The interpretation of such studies is usually limited by the lack of information on the spatial scales over which the gas flow is occurring. The similarity of our results suggests that the two lines of investigation are probing the same material. If so, the gas seen as blueshifted self-absorption could extend up to scales greater than a few kpc, as typically assumed by these authors when inferring mass outflow rates from such measurements. Having characterized the covering fraction of \mgii\ absorbers as a function of impact parameter, we can attempt to estimate the typical amount of cool HI gas traced by \mgii\ absorbers residing in the CGM of $z\sim0.5$ galaxies. To do so we can use the empirical relation between \mgii\ equivalent width and median HI column density measured by \citet{2006ApJ...636..610R} and quantified by \citet{2009MNRAS.393..808M}: $\langle N_{\rm HI} \rangle(W_0^{\rm Mg\,II})$ (their Equation 5). Using this relation we can write \begin{eqnarray} M_{\rm HI}^{\rm CGM}(>W_0^{\rm Mg\,II}) &\sim& 2\pi \int_{20\,{\rm kpc}}^{150\,{\rm kpc}} \langle N_{\rm HI} \rangle f_{c}(r_{p})\,r_{p}\,{\rm d}r_{p}\, \nonumber\\ \, \end{eqnarray} where $f_{c}(r_{p})$ is the \mgii\ covering fraction and $r_{p}$ is the projected distance. The inner impact parameter limit (20 kpc) is selected based on the range for which we have robust covering fraction measurements, while the outer impact parameter limit (150 kpc) corresponds to the maximum impact parameter of galaxy-absorber pairs for which COS-Halos team searched. To estimate this quantity numerically we consider the cool HI gas traced by $W^{\rm MgII}_{0}> 1\rm\,\AA$ ($\langle W^{\rm Mg\,II}_{0}\rangle\simeq1.6 \rm\, \AA$), corresponding to the covering fraction shown in Figure~\ref{fig:covering_fraction_cumulative}. For star-forming galaxies with $\langle \log M_\ast/M_\odot \rangle\sim 10.6$ we find \begin{eqnarray} \log M_{\rm HI}^{\rm CGM}/{\rm M_\odot} &\sim& {9.5}. \end{eqnarray} For passive galaxies with $\langle \log M_\ast/M_\odot \rangle\sim 10.9$ we find \begin{eqnarray} \log M_{\rm HI}^{\rm CGM}/{\rm M_\odot} &\sim& {9.2}. \end{eqnarray} The above numbers correspond only to HI gas traced by \mgii\ absorbers stronger than $1\,{\rm \AA}$. The estimate can be considered as a lower limit for the total amount of HI found in galaxy halos. These values imply that the ratio $M_{\rm HI}^{\rm CGM}/M_\ast$ is about four times lower around passive galaxies than around star-forming ones. Recently, \citet{Werk2014} used HST-COS observations to probe the gaseous distribution around $z\sim 0.2$ $L\sim L_\ast$ galaxies. Using photo-ionization models they inferred that such galaxies are surrounded by cool ($T \sim 10^4\,$K) gas amounting to at least $\log M_{\rm H}/M_\odot = 10.4$ within about 150\,kpc. Most (99\%) of this gas is found to be ionized, implying that the neutral component probed by the lines-of-sight of the COS-Halos survey (excluding a few damped systems) amounts to $\langle \log M_{\rm HI}/M_\odot \rangle\sim 8.4$. This is a factor 5-10 lower than the amount of HI probed by \mgii\ absorbers presented above. We note that given the low covering fraction for strong \mgii\ absorbers on such scales: $f_c(100\,{\rm kpc})\sim 0.05$, most of randomly-selected lines-of-sight are not expected to intercept such clouds. Interestingly, the above numbers indicate that most of the neutral hydrogen probed by metal absorbers is located in strong \mgii\ absorbers, despite their low covering fraction. An additional neutral gas contribution can be associated with pristine or low metallicity gas. Such clouds would not give rise to strong \mgii\ absorbers. We also note that the spatial dependence of the covering fraction of strong \mgii\ absorbers (derived in Section~\ref{covering}) can be converted into a minimum value for the \emph{mean} absorption equivalent width as a function of impact parameter. By selecting absorbers with $W_0^{\rm Mg\,II}>1.0$\,\AA\ with $1.6$\,\AA\ mean absorption and focusing on red galaxies, we find that \begin{equation} \langle W_0^{\rm Mg\,II} \rangle(r_p) \sim 0.05\,\left( \frac{r_p}{100\,{\rm kpc}} \right)^{-1.1} {\rm \AA} \;. \end{equation} This dependence can be compared to the results obtained by \citet{ZhuMgII}. These authors measured the mean absorption around galaxies averaged over all lines-of-sight. By comparing the two sets of results, we find that absorbers with $W_0^{\rm Mg\,II}>1\,$\AA\ contribute to about half of the total absorption signal. By selecting systems with $W_0^{\rm Mg\,II}>0.4\,$\AA\ we find this contribution to be comparable to the mean absorption level. This indicates that, on large scales around galaxies, strong \mgii\ absorbers dominate the total \mgii\ absorption budget. Our results also allow us to reveal that the covering fraction for cool gas around galaxies appears to change as a function of redshift. At $z\sim 0.5$ we found $f_c \propto r_p^{-1}$ for \mgii\ absorbers at impact parameters greater than about 50\,kpc around both star-forming and passive galaxies. As mentioned above, this relation steepens on smaller scales around star-forming galaxies. Using a statistical analysis, \cite{2010ApJ...717..289S} inferred the radial dependence of the covering fraction of cool gas around star-forming (Lyman break) galaxies at $2\lesssim z \lesssim 3$. They derived a significantly shallower radial dependence: $f_c \propto r_p^{-\gamma}$ with $0.2\lesssim \gamma \lesssim 0.6$, depending on the transition. The comparison of the two analyses therefore shows that the radial dependence of metals in cool gas appears to steepen from $z\simeq2-3$ down to $z\sim0.5$. Finally, if we assume that the neutral gas traced by \mgii\ absorbers is not only enriched in metals in the gas phase but also dusty, it implies a CGM dust mass of $M_{\rm dust}^{\rm CGM}\sim 4\times 10^{7}\,{\rm M_\odot}$ {using the global dust-to-gas ratio in \mgii\ clouds \citep[$\sim1/50$;][]{Menarddust}. This value is consistent with the findings of \citet{2010MNRAS.405.1025M} and \citet{Peek2014} who statistically mapped out the distribution of dust in galaxy halos using reddening measurements and inferred its total mass. It also implies that most of the circum-galactic dust is associated with \mgii\ absorbers. | 14 | 4 | 1404.5301 |
1404 | 1404.7132_arXiv.txt | We compare theoretical dust yields for stars with mass $\rm 1\,M_{\odot} \le \rm m_{star} \le \rm 8 \, M_{\odot}$, and metallicities $\rm 0.001 \le \rm Z \le \rm 0.008$ with observed dust production rates (DPR) by carbon-rich and oxygen-rich Asymptotic Giant Branch (C-AGB and O-AGB) stars in the Large and Small Magellanic Clouds (LMC, SMC). The measured DPR of C-AGB in the LMC are reproduced only if the mass loss from AGB stars is very efficient during the carbon-star stage. The same yields over-predict the observed DPR in the SMC, suggesting a stronger metallicity dependence of the mass-loss rates during the carbon-star stage. DPR of O-AGB stars suggest that rapid silicate dust enrichment occurs due to efficient hot-bottom-burning if $\rm m_{star} \ge \rm 3 \, M_{\odot}$ and $\rm Z \ge 0.001$. When compared to the most recent observations, our models support a stellar origin for the existing dust mass, if no significant destruction in the ISM occurs, with a contribution from AGB stars of $70 \%$ in the LMC and $15 \%$ in the SMC. | The best way to compare the total dust input from evolved stars in a galaxy to the total dust budget is to detect the entire population of dusty stars at infrared (IR) wavelengths and estimate the dust-injection rate of each. These global measurements are not possible in our Galaxy, due to source confusion in the Galactic plane, but have been attempted on nearby external galaxies, such as the Small and Large Magellanic Clouds (hereafter SMC and LMC, respectively). Extragalactic studies have also the advantage that distances to sources within a galaxy can be assumed to be the same. The importance of the Magellanic Clouds as laboratories of dust enrichment by stellar sources has been thoroughly discussed in the literature (Matsuura et al. 2009, 2013; Srinivasan et al. 2009; Boyer et al. 2012; Riebel et al. 2012). A growing body of observational data has been made available to the community by means of dedicated large photometric surveys. Among the others, the Magellanic Clouds Photometric Survey (MCPS, Zaritsky et al. 2004), the Two Micron All Sky Survey (2MASS, Skrutskie et al. 2006), Surveying the Agents of a Galaxy's Evolution Survey (SAGE) with the {\it Spitzer} Space Telescope (SAGE-LMC, Meixner et al. 2006; SAGE-SMC, Gordon et al. 2011), and {\it HERschel} Inventory of The Agents of Galaxy Evolution (HERITAGE, Meixner et al. 2010, 2013) have provided catalogues of point sources as well as high-resolution maps of the emission by the warm and cold dust components in the interstellar medium (ISM). In addition, a wealth of complementary data allows to recontruct the recent and past star formation histories of the galaxies (see, among others, Harris \& Zaritsky 2004, 2009; Bolatto et al. 2011; Skibba et al. 2012; Weisz et al. 2013; Cignoni et al. 2013), their metal enrichment histories (see e.g. Carrera et al. 2008a, 2008b; Piatti 2012; Piatti \& Geisler 2013), and their present-day global gas, stellar and dust content (see Meixner et al. 2013 for a recent collection of observational data). Hence, these two galaxies represent an excellent astrophysical laboratory to investigate the life-cycle of dust in the ISM, providing a fundamental benchmark to theoretical models. Knowledge of the amount and composition of dust formed by stars of intermediate mass ($\rm 1 M_{\odot} \le \rm m_{star} \le \rm 8 M_{\odot}$) and in the ejecta of core-collapse supernovae ($\rm m_{star} > 8 M_{\odot}$, where $\rm m_{star}$ is the stellar mass at zero-age main sequence) represents the first step towards the understanding of dust enrichment from the most distant galaxies to the Local Universe. The relative importance of these stellar sources of dust depends on the mass-dependent dust yields, on the stellar initial mass function (IMF), as well as on the star formation history (SFH) of each galaxy (Valiante et al. 2009, 2011). Contrary to previous claims, dust production at high redshifts, $\rm z > 6$, is not necessarily driven by massive stars, as stars of intermediate mass on their Asymptotic Giant Branch (AGB) can dominate dust production on a timescale which ranges between 150 and 500 Myr (Valiante et al. 2009). Local observations of dust emission in the circumstellar shells of evolved stars or in recent supernovae and supernova remnants have the potential to significantly reduce uncertainties associated to theoretical dust yields. In this paper, our main aim is to compare dust production rates by AGB stars calculated by means of theoretical models to observations of evolved stars in the Magellanic Clouds. In particular, we consider a new grid of dust yields for different stellar masses and metallicities (Ventura et al. 2012a, 2012b, Di Criscienzo et al. 2013; Ventura et al. 2014) which is based on models calculated with a full integration of the stellar structure, following the evolution from the pre-main sequence phase using the code ATON (Ventura et al. 1998; Ventura \& D'Antona 2009). This represents an important difference with respect to previous studies by Ferrarotti \& Gail, whose dust yields are based on stellar properties computed from synthetic\footnote{In this context, by synthetic models we refer to models in which the evolution is described with analytical relations derived by fitting the results of full evolutionary models.} models (Ferrarotti \& Gail 2001, 2002, 2006; Zhukovska, Gail \& Trieloff 2008) or with respect to more recent hybrid models where the integration is limited to the envelope structure of the stars (Marigo et al. 2013; Nanni et al. 2013). As a consequence, the mass and metallicity dependence of carbon and silicate dust yields based on ATON stellar models can not be reproduced by dust yields based on synthetic models, due to the different treatment of physical processes, such as the third dredge-up and the hot bottom burning, which alter the surface chemistry of AGB stars (Ventura et al. 2014). By comparing the predictions of different sets of AGB stars dust yields to observations of carbon-rich and oxygen-rich AGB stars in the Magellanic Clouds, we can hope to constrain some of the model uncertainties. In addition, the difference in the gas metallicity of the two galaxies, with $\rm Z_{SMC} = 0.004$ and $\rm Z_{LMC} = 0.008$, allows to explore the complex and poorly understood dependence of AGB stars dust production rates on their progenitors initial metallicity (Ventura et al. 2012b). Finally, we can assess the relative importance of AGB stars and supernovae as dust producers in the two galaxies, comparing the overall contribution of stellar sources to the estimated total dust mass in the ISM. In a recent paper, Zhukovska \& Henning (2013) have done a similar analysis, discussing the dust input from AGB stars in the LMC. For the first time, theoretically calculated dust production rates of AGB stars have been compared to those derived from IR observations of AGB stars for the entire galaxy. In their comparison, they consider synthetic yields by Zhukovska et al. (2008) but discuss also the implications of their models when ATON yields are adopted. They find that while synthetic models lead to carbon and silicate dust production rates in good agreement with observations, ATON (hereafter {\bf old} ATON) models under-predict carbon-dust production rates, favouring silicate dust production, in contrast to the observations. Motivated by these findings, we have recently investigated the dependence of the predicted dust yields on the macrophysics adopted to describe the AGB evolution (Ventura et al. 2014) and a new grid of ATON (hereafter {\bf new} ATON) dust yields has been computed for different initial metallicities, which range between $3 \times 10^{-4}$ to $8 \times 10^{-3}$, including the metallicity of the SMC, $\rm Z = 0.004$, that has not been considered in previous calculations. Here we extend the comparison to this new ATON grid. In addition, we do not limit the analysis to AGB stars in the LMC but we test the models against observations in the SMC. The paper is organized as follows: in Section 2, we briefly summarize the AGB stellar dust yields predicted by different theoretical models; in Section 3 we review the observationally constrained star formation and chemical enrichment histories of the Magellanic Clouds; the associated dust production rates for different sets of dust yields are presented in Section 4 and compared to observational data. The best fit models are then used, in Section 5, to assess the role of AGB stars and supernovae in the global dust budget of the LMC and SMC. Finally, in Section 6 we summarize and discuss our conclusions. \begin{figure*} \includegraphics[width=58mm]{AGBdustcomp01.eps} \includegraphics[width=58mm]{AGBdustcomp04.eps} \includegraphics[width=58mm]{AGBdustcomp08.eps} \caption{Dust yields from AGB stars as a function of their initial mass for three initial metallicities: Z = 0.001, 0.004, and 0.008 (left, central, and right panels, respectively), with separate contributions from carbon dust, silicates, SiC and Iron grains. In each panel, empty squares show the predictions from Zhukovska et al. (2008), filled circles and squares are the old and new ATON yields, respectively (see text).} \label{fig:agbyields} \end{figure*} | In this paper, we have compared theoretical dust yields for AGB stars to observations of dust production rates by carbon-rich and oxygen-rich AGB stars in the Small and Large Magellanic Clouds. Our aim is to test whether current observations have the potential to discriminate among different models and to shed light on the complex dependence of the dust yields on the mass and metallicity of progenitor stars. Using metallicity dependent star formation histories inferred by Harris \& Zaritsky (2004, 2009) based on the MCPS survey, we find that: \begin{itemize} \item Observed dust production rates by carbon-rich AGB stars in the LMC favour theoretical models with strong mass loss, as predicted by the new ATON models with Wachter et al. (2008) mass loss prescription. \item The same yields, however, exceed the dust production rate observed for carbon-rich AGB stars in the SMC by approximately one order of magnitudes. Hence, current data of the SMC seem to favour the old ATON yields, which were based on the less efficient mass loss rate prescription by Bl{\"o}cker et al. (1995). This conclusion is independent of the uncertainties associated to the star formation or the metal enrichment histories of the galaxy and may be an indication of a stronger metallicity dependence of the mass-loss rates during the carbon-star stage. \item Efficient Hot Bottom Burning in ATON models allows stars with $\rm m_{stars} > 3 \, M_{\odot}$ to enrich the interstellar medium with silicate and iron grains already at very low metallicities, $\rm Z \ge 0.001$, at odds with dust yields based on synthetic AGB models. If the dominant stellar populations in the SMC have metallicities $\rm Z \le 0.004$, observations of dust production rates by oxygen-rich stars have the potential to confirm or refute the theoretical predictions. \item The latest analysis by the HERITAGE team (Gordon et al. 2014) leads to integrated dust masses in the LMC and SMC that are a factor 2-4 smaller than previous estimates. When compared to our model predictions, we find that the existing dust mass in the ISM of the MCs can have a stellar origin, even without resorting to extreme yields, unless significant destruction of the newly formed dust in SN reverse shock or in the ISM takes place. \end{itemize} Our study confirms the potential of the Magellanic Clouds as fundamental astrophysical laboratories to test our current understanding of the dust cycle in the interstellar medium. Yet, conclusions based on detailed comparison between models and observations are hampered by uncertainties on the star formation and chemical enrichment history of the galaxies, particularly of the Small Magellanic Cloud. Future observational studies complemented by a more realistic modelling of the two brightest satellites of the Milky Way in a cosmological context (Boylan-Kolchin et al. 2011) are needed to make substantial progress. | 14 | 4 | 1404.7132 |
1404 | 1404.3711_arXiv.txt | We analyze string-theoretic large-field inflation in the regime of spontaneously-broken supergravity with conventional moduli stabilization by fluxes and non-perturbative effects. The main ingredient is a shift-symmetric K\"ahler potential, supplemented by flux-induced shift symmetry breaking in the superpotential. The central technical observation is that all these features are present for D7-brane position moduli in Type IIB orientifolds, allowing for a realization of the axion monodromy proposal in a controlled string theory compactification. On the one hand, in the large complex structure regime the D7-brane position moduli inherit a shift symmetry from their mirror-dual Type IIA Wilson lines. On the other hand, the Type IIB flux superpotential generically breaks this shift symmetry and allows, by appealing to the large flux discretuum, to tune the relevant coefficients to be small. The shift-symmetric direction in D7-brane moduli space can then play the role of the inflaton: While the D7-brane circles a certain trajectory on the Calabi-Yau many times, the corresponding $F$-term energy density grows only very slowly, thanks to the above-mentioned tuning of the flux. Thus, the large-field inflationary trajectory can be realized in a regime where K\"ahler, complex structure and other brane moduli are stabilized in a conventional manner, as we demonstrate using the example of the Large Volume Scenario. | The standard theory of cosmological evolution involves a period of primordial inflation which, in its simplest realization, is driven by the potential energy density of a slowly rolling scalar field, the inflaton $\vp$. This theory of slow-roll inflation is sensitive to higher-dimensional operators, thereby probing its UV completion. Consequently, any such inflationary model needs to be implemented in a UV-complete theory of quantum gravity, such as string theory. Models of slow-roll inflation can be classified according to the distance the inflaton rolls during inflation and are either of the large-field type, $\Delta \vp > M_p$, or of the small-field type, $\Delta \vp <M_p$. While there has been much progress in constructing small-field models in string theory (for a review see \cite{Burgess:2013sla,Baumann:2014nda}), realizing large-field models is notoriously difficult. In field theory, the latter are well studied, the prime candidate being chaotic inflation \cite{Linde:1983gd}. Crucially, in any viable representative of this class of models one needs to control all higher-dimensional operators. This is commonly done by imposing a shift symmetry for the inflaton. This symmetry is broken, e.g.\ by a term $\sim m^2 \vp^2$, with $m \ll 1 $ in units of the reduced Planck mass. The shift symmetry is restored in the limit $m\to 0$ and thus the model is technically natural in field theory. In string theory, however, typical inflaton candidates like D-brane positions \cite{Kachru:2003sx,Dasgupta:2002ew}, Wilson lines \cite{Avgoustidis:2006zp}, and axions generically have a field range which is limited to sub-planckian values. The same is true for K\"ahler moduli \cite{Conlon:2005jm}, except where the inflaton is identified with a breathing mode of the compact space \cite{Cicoli:2008gp}. Overall, realizing large-field models in a UV-complete theory of quantum gravity is challenging. Clearly, there are several proposed ways how one can, despite of the limited field range, construct scenarios in string theory which are effectively of the large-field type. For example, one may consider a large number of axions as in N-flation \cite{Kim:2004rp,Dimopoulos:2005ac,Kallosh:2007cc,Grimm:2007hs,Conlon:2012tz} or similar proposals \cite{Ashoorioon:2009wa, Ashoorioon:2011ki}. However, a recent analysis of an embedding of N-flation in Type IIB string theory shows that the number of axions $N$ has to be as large as $10^5$ \cite{Cicoli:2014sva}. It is questionable if such a large number can be attained. A different interesting proposal is the use of monodromy to break the periodicity and enlarge the field space of an axion \cite{Silverstein:2008sg,McAllister:2008hb,Berg:2009tg,Palti:2014kza}, a mechanism also analyzed in field theory \cite{Kaloper:2008fb,Kaloper:2011jz,Dubovsky:2011tu,Lawrence:2012ua}. These models are plagued by control issues: In the original proposal it is a pair of NS5 and anti-NS5 branes which needs to be embedded in the compact space (see, however, \cite{Palti:2014kza}). Thus, supersymmetry is broken at the string scale and it is unclear whether the description in terms of an effective supersymmetric 4d action with the anti-branes treated as probes is valid \cite{Conlon:2011qp}. In this letter we propose a novel way to realize large-field inflaton in string theory, using the position modulus of a D7-brane as the inflaton. Our model features the appealing mechanisms of a shift symmetry and a monodromy. Thus, in spirit it is similar to the proposals of \cite{Silverstein:2008sg,McAllister:2008hb,Berg:2009tg,Palti:2014kza}, however, with one major advantage: The model does not suffer from the control issues described above, i.e.\ it allows for a description in terms of an effective supergravity lagrangian. Furthermore, a rather explicit discussion of moduli stabilization e.g.\ in the Large Volume Scenario \cite{Balasubramanian:2005zx} is possible. The basic ingredients of our proposal of large field inflation with D7-branes are the following: First, we recall that the K\"ahler potential for a D7-brane modulus features a shift symmetry in the vicinity of the large complex structure point. This structure arises as the mirror dual version of the shift symmetry enjoyed by a Wilson line on a D6-brane in Type IIA string theory at large volume \cite{Kerstan:2011dy,Grimm:2011dx,Hebecker:2012qp,Hebecker:2013lha}. Disk-instantons will break the shift symmetry \cite{Kachru:2000ih}, but these effects are exponentially suppressed by the volume of the disk on the IIA side or, rather, by a complex structure modulus in the Type IIB description. The shift symmetry is crucial to avoid the supergravity $\eta$-problem \cite{Kachru:2003sx}, a mechanism equally important in the small-field cousins \cite{Hebecker:2011hk,Hebecker:2012aw,wip} of the model proposed here. Second, in the absence of fluxes the D7-brane modulus parametrizes a Riemann surface which generically has one-cycles, such that the field space of the modulus is periodic.\footnote{Immediateley after this work appeared the possibility of realising an inflation potential on Riemann surfaces was proposed in \cite{Harigaya:2014eta}.} In fact, all we need is a closed trajectory along the shift-symmetric direction in the D7-brane position moduli space. Fluxes will lead to an appearance of the brane modulus in the superpotential, such that the periodicity will be broken and a monodromy arises.\footnote{Inflation using a monodromy in the field space of a D3-brane was analyzed in \cite{Shlaer:2012by}. However, it is acknowledged in that paper that, since the proposal relies on the existence of non-trivial one-cycles in the compact space, much of the recent progress regarding moduli stabilization is not applicable in that model.} Inflation occurs along the shift-symmetric direction in the D7 moduli space. The situation is illustrated in figure~\ref{fig:Riemann}. \begin{figure} \centering \begin{overpic}[width=0.45\textwidth,tics=10]{img02.pdf} \put (18,26) {O7} \put (23,44) {O7}\put (61,44) {O7}\put (81,27) {O7}\put (63,14) {D7} \end{overpic} \caption{Illustration of the D7-brane position modulus parameter space. Inflation occurs when the D7-brane moves along a one-cycle in the parameter space, which need not necessarily be non-trivial in homology.}\label{fig:Riemann} \end{figure} Displacing the D7-brane from its minimum leads to $F$-terms in the effective action which generically destabilize the potential, i.e.\ they lead to a runaway direction in the K\"ahler moduli space. Therefore, in order to ensure stability of the system during inflation, we have to tune the coefficients of the brane-modulus-dependent terms in the superpotential to small values. This can be viewed as a tuning of complex structure moduli by a suitable choice of fluxes. We assume that the landscape will provide a model with this feature and will not discuss this tuning in any detail in this letter. Rather, given the very limited understanding of large field inflation in string theory, we think it is important to demonstrate that such models can be realized in principle in a controlled string-derived supergravity framework. As a result of working in Type IIB string theory, K\"ahler moduli stabilization can be analyzed very explicitly in our model, e.g.\ in the Large Volume Scenario, and gives non-trivial constraints on the size of the overall volume of the compact space and the coefficients of the brane moduli in the superpotential. An additional motivation for studying large-field inflation in string theory comes from the recent measurement of B-mode polarization \cite{Ade:2014xna} by the BICEP2 collaboration. The measured spectrum was fit in this reference to a spectrum from primordial gravitational waves, generated during an epoch of inflation. The corresponding amplitude of the tensor mode perturbations can be quantified in terms of the tensor-to-scalar ratio, which was quoted as $r = 0.2_{-0.05}^{+0.07}$. Such a large value for $r$ forces the inflaton $\vp$ to traverse a super-planckian field range during inflation \cite{Lyth:1996im,Boubekeur:2005zm}.\footnote{See, however, \cite{Avgoustidis:2008zu, Choudhury:2013iaa, Choudhury:2014kma, Antusch:2014cpa}.} Though the measurement and its attribution to primordial gravitational waves should clearly be confirmed independently, it certainly encourages our analysis of embedding a large-field model of inflation in string theory. Related results \cite{Marchesano:2014mla, Blumenhagen:2014gta, Grimm:2014vva, Ibanez:2014kia} appeared immediately before and after this work. | In this paper, we have outlined a scenario which has the potential to realize large-field inflation in Type IIB string theory, within a controlled 4d supergravity description with conventional moduli stabilization. More specifically, our inflaton is a D7-brane position modulus with shift-symmetric K\"ahler potential. This shift symmetry is inherited from the shift symmetry of a D6-brane Wilson line in the mirror-dual Type IIA model. Furthermore, since this latter shift symmetry requires large volume, we need to be at large complex structure in our Type IIB scenario. Shift-symmetry-breaking corrections to the K\"ahler potential are exponentially suppressed in the large periods of the complex-structure and D7-brane moduli space. Hence, they are relatively easy to control. The inflaton potential is quadratic at leading order. It is induced by the flux-superpotential which also depends on D7-brane positions. The coefficients of the relevant terms in the superpotential depend on complex structure moduli and other D7-brane positions. They can hence be tuned to be small, given a sufficiently large flux discretuum. As a result, the coefficient of the quadratic inflaton potential (i.e. the inflaton mass) can be made small. Clearly, going to a large VEV of the D7-brane position is impossible within the standard D7-brane moduli space, which is rather small. However, the natural periodicity of this space is broken by the flux mentioned above, such that a non-trivial monodromy arises. Thus, large-field inflation arises because a D7-brane circles a closed trajectory in its moduli space many times, thereby slowly growing a significant contribution to the $F$-term potential. Our parametric analysis demonstrates that the above-mentioned flux-tuning allows us to prevent this contribution from destabilizing other moduli. We analyzed a concrete example based on the Large Volume Scenario, where the most dangerous destabilization direction is that of the overall volume. However, a tuning of the coefficients to about $10^{-3}$ of their natural value, combined with an overall superpotential $W_0\sim 1$ and a volume ${\cal V}\sim 10^2$, allows us to escape destabilization. Finally, we have also analyzed naively sub-leading (in the fine-tuned small coefficients) effects which correct the quadratic form of the potential. Very interestingly, it turns out that some of these effects can become considerable in the region of inflation relevant for the presently observed CMB perturbations. Many open questions had to be left for future work. They include an explicit demonstration of the flux-based tuning, a more detailed phenomenology of the inflationary potential, the combination of our D7-brane inflation scenario with other K\"ahler moduli stabilization mechanisms, and the discussion of corrections associated with the uplifting contribution. \vspace*{0.5cm} | 14 | 4 | 1404.3711 |
1404 | 1404.4307_arXiv.txt | We present in detail the scientific objectives in fundamental physics of the Space-Time Explorer and QUantum Equivalence Space Test (STE-QUEST) space mission. STE-QUEST was pre-selected by the European Space Agency together with four other missions for the cosmic vision M3 launch opportunity planned around 2024. It carries out tests of different aspects of the Einstein Equivalence Principle using atomic clocks, matter wave interferometry and long distance time/frequency links, providing fascinating science at the interface between quantum mechanics and gravitation that cannot be achieved, at that level of precision, in ground experiments. We especially emphasize the specific strong interest of performing equivalence principle tests in the quantum regime, \textit{i.e.} using quantum atomic wave interferometry. Although STE-QUEST was finally not selected in early 2014 because of budgetary and technological reasons, its science case was very highly rated. Our aim is to expose that science to a large audience in order to allow future projects and proposals to take advantage of the STE-QUEST experience. | \label{sec:intro} \subsection{Scientific Motivations} \label{sec:scmotiv} Our best knowledge of the physical Universe, at the deepest fundamental level, is based on two theories: Quantum Mechanics (or, more precisely, Quantum Field Theory) and the classical theory of General Relativity. Quantum Field Theory has been extremely successful in providing an understanding of the observed phenomena of atomic, particle, and high energy physics and has allowed a unified description of three of the four fundamental interactions that are known to us: electromagnetic, weak and strong interactions (the fourth one being gravitation). It has led to the Standard Model of particle physics that has been highly successful in interpreting all observed particle phenomena, and has been strongly confirmed with the recent discovery at the LHC of the Higgs (or, more precisely, Brout-Englert-Higgs) boson, which could in fact be viewed as the discovery of a fifth fundamental interaction. Although open questions remain within the Standard Model of particle physics, it is clearly the most compelling model for fundamental interactions at the microscopic level that we have at present. On the other hand, Einstein's theory of General Relativity (GR) is a cornerstone of our current description of the physical world at macroscopic scales. It is used to understand the flow of time in the presence of gravity, the motion of bodies from satellites to galaxy clusters, the propagation of electromagnetic waves in the vicinity of massive bodies, the evolution of stars, and the dynamics of the Universe as a whole. GR brilliantly accounts for all observed phenomena related to gravitation, in particular all observations in the Earth's environment, the Solar system, in relativistic binary pulsars and, beyond that, on galactic and cosmological scales. The assumed validity of GR at cosmological scales, and the fact that non-gravitational interactions are described by the Standard Model of particle physics, together with a hypothesis of homogeneity and isotropy of cosmological solutions of these theories, have led to the ``concordance model'' of cosmology, referred to as the $\Lambda$-CDM (Cold Dark Matter) model, which is in agreement with all present-day observations at large scales, notably the most recent observations of the anisotropies of the cosmic microwave background by the Planck satellite~\cite{Ade2013}. However, important puzzles remain, in particular the necessary introduction of dark energy, described by a cosmological constant $\Lambda$, and of cold dark matter, made of some unknown, yet to be discovered, stable particle. There is a potential conflict on the problem of dark matter between the concordance model of cosmology and the Standard Model of particles. On the one hand, there is strong evidence~\cite{Ade2013} that 26.8 \% of the mass-energy of the Universe is made of non-baryonic dark matter particles, which should certainly be predicted by some extension of the Standard Model of particles. On the other hand, there is no indication of new physics beyond the Standard Model which has been found at the LHC. For instance, the search of supersymmetry at LHC has for the moment failed. Although very successful so far, GR as well as numerous other alternative or more general theories of gravitation are classical theories. As such, they are fundamentally incomplete, because they do not include quantum effects. A theory solving this problem would represent a crucial step towards the unification of all fundamental forces of Nature. Most physicists believe that GR and the Standard Model of particle physics are only low-energy approximations of a more fundamental theory that remains to be discovered. Several concepts have been proposed and are currently under investigation (\textit{e.g.}, string theory, loop quantum gravity, extra spatial dimensions) to bridge this gap and most of them lead to tiny violations of the basic principles of GR. One of the most desirable attributes of that fundamental theory is the unification of the fundamental interactions of Nature, \textit{i.e.} a unified description of gravity and the three other fundamental interactions. There are several attempts at formulating such a theory, but none of them is widely accepted and considered successful. Furthermore, they make very few precise quantitative predictions that could be verified experimentally. One of them is the Hawking radiation of black holes, which is however far from being testable experimentally for stellar-size black holes we observe in astrophysics. Therefore, a fuller understanding of gravity will require observations or experiments able to determine the relationship of gravity with the quantum world. This topic is a prominent field of activity with repercussions covering the complete range of physical phenomena, from particle and nuclear physics to galaxies and the Universe as a whole, including dark matter and dark energy. A central point in this field is that most unification theories have in common a violation at some (\textit{a priori} unknown) level of one of the basic postulates of GR, which can be tested experimentally: the Einstein Equivalence Principle (EEP). Let us emphasize that the Weak Equivalence Principle (WEP) is not a fundamental symmetry of physics, contrary to \textit{e.g.} the principle of local gauge invariance in particle physics. An important challenge is therefore to test with the best possible accuracy the EEP. This is then the main motivation of many experiments in fundamental physics, both on Earth and in space. Precision measurements are at the heart of the scientific method that, since Galileo's time, is being used for unveiling Nature and understanding its fundamental laws. The assumptions and predictions of GR can be challenged by precision experiments on scales ranging from micrometers in the laboratory to the Solar System size, in the latter case using spacecrafts or the orbiting Earth, Moon and planets. The implementation of tests with significantly improved sensitivity obviously requires the use of state-of-the-art technology, and in case of satellite-based experiments the challenge is to make such technology compatible with use in space, \textit{i.e.} extremely robust, reliable, and automatized. \subsection{The Space Mission STE-QUEST} \label{sec:mission} The satellite STE-QUEST (Space-Time Explorer and QUantum Equivalence Space Test) is specifically designed for testing different aspects of the EEP and searching for its violation with high precision. The weak equivalence principle has been verified with high precision using torsion balances on ground~\cite{Schlamminger2008} and the Lunar laser ranging~\cite{Williams2004}. It will be tested in Earth orbit by the CNES satellite $\mu$-SCOPE (Micro-Satellite \`a tra\^in\'ee Compens\'ee pour l'Observation du Principe d'Equivalence) in 2016~\cite{Touboul2001}. On the other hand, the gravitational red-shift, a different aspect of the EEP, was first measured using gamma ray spectroscopy in the laboratory~\cite{Pound1960}, and the most precise test so far was done in space with the GP-A experiment~\cite{Vessot1979}. The ESA mission ACES (Atomic Clock Ensemble in Space) will test the gravitational red-shift with the highly accurate laser-cooled atomic clock PHARAO on the International Space Station (ISS) in 2016~\cite{Cacciapuoti2009}. Atomic clocks and high-performance time and frequency links, atom interferometers and classical accelerometers are today able to measure frequency, time, and distances, and furthermore to track the motion of massive bodies, quantum particles, and light to accuracy levels never reached before. These instruments achieve their ultimate performance in space, where the clean environment and the free-fall conditions become essential for identifying tiny deformations in space-time that might bring the signature of new physics or new fundamental constituents. From this point of view, it is not surprising that fundamental physics pervades all aspects of space science. STE-QUEST was proposed in the fall of 2010 in response to ESA's M3 call in the Cosmic Vision programme (with launch date in the 2022-24 time interval), by a science team under coordination by S.~Schiller and E.M.~Rasel with support from 67 colleagues from Europe and the USA. STE-QUEST is based on the earlier proposals EGE~\cite{Schiller2009} and MWXG~\cite{Ertmer2009}, submitted to ESA's M2 call. ESA performed a ``concurrent design facility'' study of a mission concept similar to EGE, named STE, in 2010. Previously, ESA had also convened a Fundamental Physics Advisory Team which in 2009-10 developed a roadmap on fundamental physics in space. STE-QUEST, together with three other mission proposals, was selected in early 2011 by ESA's advisory structure as one candidate mission. STE-QUEST went through an assessment phase study of the satellite and payload (see the Yellow Book of the mission~\cite{YellowBook}). As a result of the assessment phase and in agreement with the national space agencies, STE-QUEST was removed from the candidate pool in late 2013, before the final selection of a single mission for the M3 slot, because of budgetary and technological reasons. Nevertheless, ESA's advisory committees evaluated the science aspects of STE-QUEST in early 2014 (together with the remaining M3 candidates) and ranked them highly. It is likely that STE-QUEST will recompete for the M4 launch slot. The primary science objectives of STE-QUEST is testing the different aspects of the Einstein Equivalence Principle with quantum sensors. The payload consists of a differential atom interferometer comparing the free propagation of matter waves of different composition under the effect of gravity and a frequency comparison link in the microwave domain for comparing atomic clocks on ground. STE-QUEST performs a direct test of the WEP by comparing the free fall of quantum objects of different composition. The E\"otv\"os ratio between the matter waves of two isotopes of the Rubidium atom is measured in a differential atom interferometer down to the $2 \times 10^{-15}$ uncertainty level. While present limits on WEP tests involving classical objects reach an uncertainty of a few parts in $10^{13}$, measurements performed on quantum objects (matter waves in states which have no classical counterpart, \textit{e.g.} spatio-temporal quantum superpositions) are still at the level of a few parts in $10^7$~\cite{Fray2004, Schlippert2014, Tarallo2014}. From this point of view, STE-QUEST will explore the boundaries between gravitation and quantum mechanics, significantly improving existing measurements and complementing experiments such as $\mu$-SCOPE, designed for a classical WEP test in space to the level $1 \times 10^{-15}$. \begin{figure}[t] \begin{center} \begin{tabular}{c} \includegraphics[width=14cm]{stequest.jpg} \end{tabular} \caption{The STE-QUEST spacecraft in orbit around the Earth. The mission is designed to test the Einstein Equivalence Principle by tracking the free-fall motion of quantum matter waves, by performing gravitational red-shift tests between ground clocks on intercontinental distances and with (optionally) a high-stability and high-accuracy onboard clock, and by performing tests of local Lorentz invariance and CPT symmetry. (OL: optical link; MWL: microwave link.)}\label{fig:stequest} \end{center} \end{figure} \begin{table*}[] \begin{center} \begin{tabular}{|p{4.7cm}||p{11.3cm}|} \hline \centerline{\textbf{Science Investigation}} & \centerline{\textbf{Measurement Requirement}} \\\hline \multicolumn{2}{|l|}{\textbf{Weak Equivalence Principle Tests}} \\ \hline \textit{Universality of propagation of matter-waves} & Test the universality of the free propagation of matter waves to an uncertainty in the E\"otv\"os parameter better than $2\times 10^{-15}$. \\ \hline \multicolumn{2}{|l|}{\textbf{Gravitational Red-shift Tests}}\\ \hline \textit{Sun gravitational red-shift} & Test of the Sun gravitational red-shift effect to a fractional frequency uncertainty of $2\times 10^{-6}$, with an ultimate goal of $5\times 10^{-7}$. \\ \hline \textit{Moon gravitational red-shift} & Test of the Moon gravitational red-shift effect to a fractional frequency uncertainty of $4\times 10^{-4}$, with an ultimate goal of $9\times 10^{-5}$. \\ \hline \textit{Earth gravitational red-shift (optional)\footnote{This scientific investigation can be performed only if the STE-QUEST payload is equipped with a high-stability and high-accuracy atomic clock.}} & Measurement of the Earth gravitational red-shift effect to a fractional frequency uncertainty of $2\times 10^{-7}$. \\ \hline \multicolumn{2}{|l|}{\textbf{Local Lorentz Invariance and CPT Tests}}\\ \hline \textit{LLI and CPT} & Provide significant improvements on the determination of several LLI and CPT parameters of the Lorentz and CPT symmetry violating Standard Model Extension. \\ \hline \end{tabular} \caption{Science investigations \textit{vs.} measurement requirements for topics in fundamental physics that shall be investigated by STE-QUEST.}\label{tab:topics} \end{center} \end{table*} STE-QUEST also tests another complementary aspect of the Einstein Equivalence Principle, one of the most fascinating effects predicted by GR and other metric theories of gravity: the gravitational red-shift or gravitational time dilation effect. As direct consequence of the EEP, time runs (or clocks tick) more slowly near a massive body. This effect can be detected when comparing the time intervals measured by identical clocks placed at different depths in a gravitational field. The microwave link (MWL) of the STE-QUEST satellite allows comparing ground clocks down to the $1 \times 10^{-18}$ uncertainty level. Such measurements, far beyond the capabilities of existing frequency transfer systems, will perform clock red-shift tests in the field of the Sun and the Moon, respectively at the $2 \times 10^{-6}$ and $4 \times 10^{-4}$ uncertainty levels. For comparison, existing measurements of the Sun red-shift effect are at the few \% uncertainty level while, to our knowledge, no such tests have ever been performed in the field of the Moon. An optional (depending on available funding) onboard clock allows additionally a red-shift measurement in the Earth field by taking advantage of the high apogee and high eccentricity of the orbit. The clock under consideration is derived from the PHARAO cold atom \text{Cs} clock to be flown on the ISS~\cite{Cacciapuoti2009}. The version planned for STE-QUEST is designed to reach an uncertainty in the Earth field red-shift test of $2 \times 10^{-7}$, one order of magnitude better than the objective of ACES. The relativistic theory for time and frequency transfer needed for frequency links in space missions such as ACES and STE-QUEST is described in Ref.~\cite{Blanchet2001}. Clock red-shift measurements obtained in the field of the Earth, the Sun or the Moon test the Local Position Invariance (LPI) principle and search for anomalous couplings depending on the composition of the source of the gravitational field. LPI is a constituent of EEP together with WEP and the Local Lorentz Invariance (LLI) principle, see Sec.~\ref{sec:facet}. As we shall discuss in Sec.~\ref{sec:diffEEP}, in generic frameworks modelling a possible violation of EEP, WEP and clock red-shift tests are complementary and need to be pursued with equal vigor as, depending on the model used, either one of the tests can prove significantly more sensitive than the other. Improving the accuracy of these tests will bring significant progress in restricting the parameters space and discriminating between theories seeking to unify quantum mechanics with gravity. The eventual detection of an EEP violation would carry the signature of new fundamental constituents or interactions in the Universe (\textit{e.g.} scalar fields for dark energy, particles for dark matter, fundamental strings, \textit{etc.}). In this case, STE-QUEST tests would have a significant impact not only for fundamental physics research, but also for cosmology and particle physics. The ensemble of fundamental physics science objectives of STE-QUEST is summarized in Table~\ref{tab:topics} and Fig.~\ref{fig:stequest}. STE-QUEST has also important applications in domains other than fundamental physics, in particular in the fields of time and frequency metrology and for geodesy studies. As mentionned, the STE-QUEST high-performance MWL provides the means for connecting atomic clocks on ground in a global network, enabling comparisons down to the $1 \times 10^{-18}$ fractional frequency uncertainty level. Clock comparisons \textit{via} STE-QUEST will contribute to the realization of international atomic time scales (UTC, TAI, \textit{etc.}) and to the improvement of their stability and accuracy. Synchronization of clocks, space-to-ground and ground-to-ground, to better than $50\,\text{ps}$ can be achieved through STE-QUEST for distributing time scales to unprecedented performance levels. Common-view comparisons of ground clocks, primarily used for gravitational red-shift tests in the field of Sun or Moon, also provide direct information on the geopotential differences at the locations of the two ground clocks. STE-QUEST will therefore contribute to establishing a global reference frame for the Earth gravitational potential at the sub-$\text{cm}$ level through local measurements. This method is complementary to current and future satellite gravimetry missions such as CHAMP, GRACE and GOCE as well as to altimetry missions like JASON and Envisat in defining the Global Geodetic Observing System (GGOS). The Table~\ref{tab:othertopics} (relegated in the conclusion section~\ref{sec:conclusion}) summarizes the list of topics other than fundamental physics that shall be investigated by STE-QUEST. The present paper is an adapted version of the fundamental physics science objectives of STE-QUEST extracted from the Yellow Book of STE-QUEST which is available in Ref.~\cite{YellowBook} (see also Ref.~\cite{Aguilera2013}). The Yellow Book also gives an overview of science objectives in other fields (geodesy, time/frequency metrology, reference frames) and details on the mission and payload, which are however beyond the scope of this paper that focuses on the fundamental physics objectives. In Sec.~\ref{sec:EEP} we shall review in more detail the EEP and its different facets. In Sec.~\ref{sec:EEPtoday} we shall discuss the status of EEP in Physics today and particularly in the contexts of cosmology and particle physics. Quantum mechanics and the EEP and the potential interest of quantum tests of the EEP will be analyzed in Sec.~\ref{sec:quantum}. The specific tests of the EEP which will be achieved by STE-QUEST will be presented in Sec.~\ref{sec:STEtest}. The paper ends with the main conclusions in Sec.~\ref{sec:conclusion}. | \label{sec:conclusion} \begin{table*}[t] \begin{center} \begin{tabular}{|p{5.4cm}||p{11cm}|} \hline \centerline{\textbf{Science Investigation}} & \centerline{\textbf{Measurement Requirement}} \\ \hline \multicolumn{2}{|l|}{\textbf{Clock Comparisons and International Atomic Time Scales}} \\ \hline \textit{Common-view comparisons of ground clocks} & Common-view comparison of ground clocks at the $1\times 10^{-18}$ fractional frequency uncertainty level after a few days of integration time with the STE-QUEST microwave link and a few hours by using the optical link. \\ \hline \textit{Space-to-ground time transfer} & Space-to-ground time transfer with accuracy better than $50\,\text{ps}$. \\ \hline Synchronization of ground clocks & Synchronization of clocks on ground to better than $50\,\text{ps}$. \\ \hline \textit{Atomic time scales} & Contribution to the generation of atomic time scales to fractional frequency inaccuracy lower than $1\times 10^{-16}$. \\ \hline \textit{GNSS clocks and time scales (optional)\footnote{This scientific investigation can be performed only if the STE-QUEST payload is equipped with a high-stability and high-accuracy atomic clock.}} & Monitoring of the stability of on-board GPS, GALILEO, and GLONASS clocks. \\ \hline \multicolumn{2}{|l|}{\textbf{Geodesy}}\\ \hline \textit{On-site differential geopotential measurements} & Differential geopotential measurements between two points on the Earth's surface with resolution in the gravitational potential $U$ at the level of $0.15\,\text{m}^2/\text{s}^2$ (equivalent to $1.5\,\text{cm}$ on the geoid height difference). \\ \hline \multicolumn{2}{|l|}{\textbf{Reference Frames}}\\ \hline \textit{Earth terrestrial and celestial reference frame} & Realization and unification of the terrestrial and the celestial reference frame of the Earth. \\ \hline \end{tabular} \caption{Science investigations \textit{vs.} measurement requirements for topics other than fundamental physics that shall be investigated by STE-QUEST.}\label{tab:othertopics} \end{center} \end{table*} We have presented the fundamental physics science objectives of STE-QUEST, which are centered on tests of the three different aspects of the Einstein Equivalence Principle (EEP): the Weak Equivalence Principle (WEP) or Universality of Free Fall, Local Position Invariance (LPI) or Universality of Clock Rates, and Local Lorentz Invariance (LLI) coupled to CPT symmetry. One of the unique strengths of STE-QUEST is that it will test all three aspects of the EEP, using a combination of measurements in space and on the ground (relative acceleration of different atomic isotopes, comparison of distant clocks). Although the three sub-principles are connected by Schiff's conjecture, the actual quantitative merit of the different experiments is model-dependent (see Sec.~\ref{sec:diffEEP}). As a consequence, it is not known \textit{a priori} which test (WEP, LPI, or LLI) is more likely to first detect a violation and the most reasonable approach is to pursue tests of the three sub-principles with equal vigor. This is the baseline of STE-QUEST, which carries out, and improves on, state-of-the-art tests of all three sub-principles. Another unique feature of STE-QUEST is its capability to carry out the WEP test using quantum matter waves in superpositions that have no classical counterpart. Although, we know of no viable model that predicts an EEP violation specific to such quantum matter waves, one should be prepared for surprises as there are numerous open questions at the interface between gravitation and quantum mechanics reviewed in Sec.~\ref{sec:quantum}. Those issues provide good reason for exploring the foundations of general relativity, \textit{i.e.} the EEP, with as diverse objects as possible, like antimatter or quantum degenerate gases and superposition states. Finally let us mention that although the primary science objectives of STE-QUEST are in fundamental physics, the mission will also provide a wealth of legacy science for other fields like time/frequency metrology, reference frames and geodesy. These are summarized in Table~\ref{tab:othertopics}, with more details available in the STE-QUEST Yellow Book~\cite{YellowBook}. To conclude, all of this gives a unique science case for STE-QUEST, making it the first space mission to carry out such a complete test of the foundations of gravitation theory using quantum sensors, and providing additionally some legacy science for other applications. If selected in an upcoming call, it will follow in the wake of precursor missions like ACES and $\mu$-SCOPE (to be launched by ESA and CNES in 2016) and firmly establish Europe's lead in fundamental physics as a space science discipline. | 14 | 4 | 1404.4307 |
1404 | 1404.1378_arXiv.txt | We discuss the rest-frame ultraviolet emission from the starbursting galaxy HFLS3 at a redshift of 6.34. The galaxy was discovered in {\it Herschel}/SPIRE data due to its {\it red} color in the sub-mm wavelengths from 250 to 500 $\mu$m. The apparent instantaneous star-formation rate of HFLS3 inferred from the total far-IR luminosity measured with over 15 photometric data points between 100 and 1000 $\mu$m is 2900 M$_{\odot}$ yr$^{-1}$. Keck/NIRC2 K$_s$-band adaptive optics imaging data showed two potential near-IR counterparts near HFLS3. Previously, the northern galaxy was taken to be in the foreground at $z=2.1$ while the southern galaxy was assumed to HFLS3's near-IR counterpart. The recently acquired {\it Hubble}/WFC3 and ACS imaging data show conclusively that both optically-bright galaxies are in the foreground at $z < 6$. A new lensing model based on the {\it Hubble} imaging data and the mm-wave continuum emission yields a magnification factor of $2.2 \pm 0.3$. The lack of multiple imaging constrains the lensing magnification to be lower than either 2.7 or 3.5 at the 95\% confidence level for the two scenarios, which attribute one or two components to HFLS3 in the source plane. Once accounting for the possibility of gravitational lensing, the instantaneous star-formation rate is 1320 M$_{\sun}$ yr$^{-1}$ with the 95\% confidence lower limit around 830 M$_{\sun}$ yr$^{-1}$. Using models for the rest-frame UV to far-IR spectral energy distribution (SED) we determine the average star-formation rate over the last 100 Myr to be around 660 M$_{\odot}$ yr$^{-1}$. The dust and stellar masses of HFLS3 from the same SED models are at the level of $3 \times 10^8$ M$_{\odot}$ and $\sim 5 \times 10^{10}$ M$_{\odot}$, respectively, with large systematic uncertainties on assumptions related to the SED model. With {\it Hubble}/WFC3 images we also find diffuse near-IR emission about 0.5 arcseconds ($\sim$ 3 kpc) to the South-West of HFLS3 that remains undetected in the ACS imaging data. The emission has a photometric redshift consistent with either $ z\sim 6$ or a dusty galaxy template at $ z\sim 2$. If at the same redshift as HFLS3 the detected diffuse emission could be part of the complex merger system that could be triggering the starburst. Alternatively, it could be part of the foreground structure at $z \sim 2.1$ that is responsible for lensing of HFLS3. | The unexpected discovery of HFLS3 (HerMES J170647.8+584623) at a redshift of $6.3369 \pm 0.0009$ in {\it Herschel} Space Observatory's \citep{Pilbratt10} has led to the possibility that massive starbursting galaxies could be an appreciable contributor to the star-formation rate density of the Universe during the epoch of reionization \citep{Riechers13}. The galaxy was first identified in {\it Herschel} Multi-Tiered Extragalactic Survey (HerMES\footnote{http://hermes.sussex.ac.uk}, \citealt{Oliver12}) as a high-redshift candidate due to its ``red'' color in the SPIRE \citep{Griffin10} data, with $S_{500}/S_{350} \sim 1.45$ and $S_{500} \sim 47 \pm 3$ mJy. The redshift of HFLS3 was secured through the detection of more than 20 individual molecular and atomic lines at far-IR/sub-mm wavelengths with ground-based interferometers. HFLS3 was found to be luminous ($L_{\rm IR}=(3.4 \pm 0.3)\times10^{13}$ L$_{\odot}$), gas-rich ($M_{\rm gas} \sim 10^{11}$ M$_{\odot}$) and dusty ($T_{\rm d}=49\pm2$ K). The instantaneous star-formation rate (SFR) implied by the above total IR luminosity \citep{Kennicutt98} is around 2900 M$_{\odot}$ yr$^{-1}$ for a \citet{Chabrier03} initial mass function. It is also the highest redshift sub-mm galaxy (SMG) known to date, potentially probing the earliest formation epoch of dust in the Universe \citep*[for a recent review of SMGs and dusty star-forming galaxies in general see][]{Casey14}. One complication in interpreting the properties of HFLS3 is that it was found to be $\sim 0.5\arcsec$ to the South of a $z= 2.09$ galaxy (Figure~1), identified by Keck/NIRC2 K-band AO imaging and Keck/LRIS spectroscopy. This suggests some possibility that the flux density of HFLS3 is enhanced by gravitational lensing with a magnification factor, $\mu_{\rm lens}$. Due to the steepness of the SMG number counts and their high redshifts, and the corresponding high magnification bias, sub-mm surveys are known to be highly sensitive to gravitational lensing modifications \citep{Blain96,Perrotta02,Negrello07,Paciga09}. At the bright-end of the number counts at wavelengths longer than 350 $\mu$m, lensed SMGs appear as a power-law distinct from the intrinsic counts (e.g., \citealt{Negrello10,Wardlow13,Vieira13}). At $z > 4$, we expect the lensing fraction to be substantial for current generation surveys, where the flux density limit for the source detection is relatively high. An example of a high efficiency lensing selection at $z > 3$ is the bright SMG sample from the South Pole Telescope at 1.4 mm \citep{Vieira13,Weiss13}. If lensing is a statistically important correction to the flux densities of high-redshift SMGs we expect them to be discovered near foreground galaxies and groups. Such a close association with a foreground galaxy is consistent with the existing indications that a reasonable fraction of the $z > 7$ Lyman-break drop-outs are also magnified by $\mu_{\rm lens}\sim$ few due to their closeness to foreground bright galaxies \citep{Wyithe11}. In the case of HFLS3, a possibility for lensing was expected since the Keck/LRIS spectroscopy showed emission lines corresponding to a foreground galaxy at a $z=2.1$ within one arcsecond of the peak 1.1 mm continuum emission. The high resolution Keck/NIRC2 LGS-AO imaging data in the K$_s$-band showed two galaxies within $1\farcs5$ of HFLS3. In \citet{Riechers13} the northern component was taken to be the $z=2.1$ foreground galaxy, while the southern component, close to the peak 1.1 mm emission, was taken to be the rest-frame optical counterpart, or the least obscured part, of HFLS3. Under such an assumption, deblended NIRC2 and {\it Spitzer}/IRAC photometry suggested a stellar mass of $\sim 3.7 \times 10^{10}$ M$_{\sun}$. Thus, HFLS3 is already a stellar mass-rich galaxy at $z=6.34$, while continuing to form stars at a very high rate of $>2000$ M$_{\odot}$ yr$^{-1}$. The lack of multiple images of HFLS3 in mm-wave interferometric imaging data was inferred to imply that the lensing magnification factor is negligible, with $\mu_{\rm lens}=1.5 \pm 0.7$ associated with lensing by the foreground galaxy to the north of the assumed rest-frame optical counterpart. Due to such a small magnification a lensing correction to the properties of HFLS3 was not included in \citet{Riechers13}. However, the lensing magnification determination is subject to assumptions related to the counterpart identification and the location of foreground galaxies relative to the mm-wave emission. Since the true mass and star-formation rate of HFLS3 are directly related to its cosmic rarity, a potential lensing correction is even more important when addressing whether HFLS3 is a rare source among the SMG sample or if it is a source typical of $z > 4$ SMGs \citep{Daddi09,Coppin10,Capak11,Walter12,Combes12}. Here we report {\it Hubble}/WFC3 and ACS imaging observations of HFLS3 in five filters from optical to near-IR wavelengths. We use these data to study the physical properties of HFLS3 by improving the lensing model and by identifying rest-frame optical/UV emission for a new estimate of the stellar mass of HFLS3. This paper is organized as follows. In the next Section we summarize the observations and the analysis. We discuss the counterpart identification and {\sc Galfit} (Peng et al. 2002) models in Section~3. Our lens models and the magnification factor of HFLS3 are presented in Section~4. In Section~5 we present the modeling of optical to IR SED of foreground galaxies and the UV to far-IR SED of HFLS3. We present a discussion of our key results and the implications for the presence of massive, dusty starbursts galaxies at high redshifts in Section~6 and conclude with a summary in Section~7. For lensing and SED models we assume the best-fit concordance cosmology consistent with WMAP-9 year and Planck data \citep{Hinshaw13,Planck13}. \begin{figure*}[htb] \centerline{ \includegraphics[scale = 0.4]{Fig1a.eps} \includegraphics[scale = 0.4]{Fig1b.eps} } \caption{\label{fig:galfit} {\it Left:} The three-color image using HST/ACS combined F625W and F814W (blue), HST/WFC3-IR combined F160W, F125W and F105W, and Keck/NIRC-2 $K_s$-band LGS-AO (red) images. Note the clear detection of two galaxies close to HFLS3 shown here in terms of the IRAM/PdBI 1.1 mm (rest-frame 158 $\mu$m) emission. The r.m.s. uncertainty in the PdBI A-array configuration data is 180 $\mu$Jy beam$^{-1}$ and the contours are shown in steps of 3$\sigma$ starting at 5$\sigma$. The instrumental beam is shown to the bottom right with FWHM of 0.35$''$ $\times$ 0.23$''$. {\it Right:} The three-color {\sc GALFIT} residual map where we remove models for the HST/ACS-detected galaxies in HST/WFC3. Here we show the combination of ACS/F625W$+$F814W (blue), WFC3/F105W (green) and WFC3/F160W (red). Both G1 and G2 are detected in the combined ACS/F625W and F814W stack, consistent with the scenario that both G1 and G2 are at $z < 6$ and G2 is not the least obscured region, or the rest-frame optical counterpart, of HFLS3, as was previously assumed. We find a marginal detection of rest-UV emission at the location of HFLS3 (labeled R2) and a higher significance diffuse emission $0\farcs5$ to the South-West of HFLS3 (labeled R1). We use WFC3 fluxes and ACS upper limits of R2 for combined SED modeling of HFLS3 with far-IR/sub-mm flux densities. We detemine a photometric redshift for R1 and find it to be consistent with emission from either a galaxy at $z \sim 6$ or a dusty galaxy at $z\sim 2$. } \end{figure*} | \label{sec:discussion} The {\sc Magphys} SED models of HFLS3 described above lead to SFR, dust mass, stellar mass among other properties. As outlined in \citet{Riechers13}, instantaneous SFR, using the FIR luminosity and assuming a \citet{Chabrier03} IMF to scale the \citet{Kennicutt98} relation, is $\sim$ 2900 M$_{\odot}$ yr$^{-1}$. Using the {\sc Magphys} SED model, we find that the apparent SFR, averaged over the last 100 Myr, to be 1450 $\pm$ 100 M$_{\odot}$ yr$^{-1}$. Note that these SFRs must be corrected down by the factor $\mu_{\rm lens}$ to account for lensing magnification. With our preferred best-fit correction factor of $2.2 \pm 0.3$ for the model involving two components to describe 1.1-mm emission from HFLS3, the instantaneous and 100-Myr averaged SFRs are $\sim$ 1300 M$_{\odot}$ yr$^{-1}$ and $\sim 660$ M$_{\odot}$ yr$^{-1}$, respectively. The two are different as the \citet{Kennicutt98} relation assumes a bolometric luminosity of a constant star-formation lasting over 100 Myr emitted in the infrared \citep{Kennicutt98,LeithererHeckman95}. For a constant star-formation, bolometric luminosity after the first 10 Myr evolves relatively slowly as the rate of birth and death of massive stars that dominate the bolometric luminosity reach a steady state. For starbursting galaxies, however, the SFR is likely changing rapidly over the 100 Myr time interval and we may be observing the galaxy at the peak of the SFR. Such a possibility then naturally explains why the instantaneous SFR is a factor of two higher than the SFR averaged over the last 100 Myr. We can also place a strict lower limit on the SFRs using the 95\% confidence level upper limit on lensing magnification. This leads to values of $>$ 780 M$_{\odot}$ yr$^{-1}$ and 390 M$_{\odot}$ yr$^{-1}$ for instantaneous and 100-Myr averaged SFRs, respectively. This revision of the SFR to a lower value is consistent with a similar revision to the SFR of $z =5.3$ SMG AzTEC-3 \citep{Capak11}. While the total IR luminosity implies a SFR of 1800 M$_{\odot}$ yr$^{-1}$ \citep{Capak11,Riechers10}, SED modeling of the fluxes with population synthesis models have shown the SFR, averged over the last 100 Myr, to be as low as 500 M$_{\odot}$ yr$^{-1}$ \citep{Dwek11}. Our SED models also show that the age of the oldest stars in HFLS3 is around 200 Myr, suggesting that HFLS3 started assembling its stars at a redshift of $\sim 8$, during the epoch of reionization. Using the far-IR/sub-mm SED and the standard assumptions used in {\sc Magphys}, and correcting for magnification, the dust mass of HFLS3 is $\sim 3 \times 10^8$ M$_{\odot}$, with a lower limit at $2 \times 10^8$ M$_{\odot}$. The ISM includes two components with dust temperatures of 24 $\pm$2 and 50 $\pm$ 2 K. The best-fit SED model is such that $>$ 90\% of the dust mass is in the warm phase, contrary to low-redshift star-forming galaxies that have a lower ratio. Such a high ratio for HFLS3 establishes that most of the dust is associated with star-bursting clumps and not the diffuse cirrus. The implied dust temperature of the cold phase component is comparable to the CMB temperature at $z=6.3$, suggesting that the extended cirrus of this galaxy may be in radiative equillibrium with the CMB. Using the \citet{Chabrier03} IMF, with parameters derived again from the SED fits using {\sc Magphys}, and with lensing magnification included, we find that HFLS3 has a stellar mass of about $5\times 10^{10}$ M$_{\odot}$. This stellar mass, however, is highly uncertain as it is based on just three detections at the rest-frame UV wavelengths. And in all of these cases, the detections are at the level of 3$\sigma$. Furthermore, we have assumed that the magnification factor derived with 1.1-mm continuum map also applies for the rest-frame UV emission form which the stellar mass is derived. Regardless of these uncertainities, we find that HFLS3 has formed a substantial amount of stellar mass already. Such a high stellar mass is already at the limits allowed by the dynamical mass of HFLS3 reported in \citet{Riechers13}. While the SED-based stellar mass is uncertain by an order of magnitude once all modeling errors are accounted for, the dust mass of HFLS3 with a value $\sim 3 \times 10^8$ M$_{\odot}$ provides an additional constraint on the stellar mass of HFLS3. This comes from models related to the dust formation mechanisms in massive starbursts where core collapse supernovae (CCSNe) are expected to be the origin of the bulk of the elements that formed the dust. The contribution of low mass stars to the refractory elements is negligible in a young galaxy such as HFLS3. Thus, the total number of CCSNe that exploded in the galaxy dictates the maximum dust mass. Following the arguments in Watson et~al. (2014, in preparation), from an observed dust mass, we can infer the minimum number of supernovae that occurred and for a particular initial mass function, the resulting lower bound on the stellar mass. The simplest and most robust way to make such an estimate on the stellar mass is to work from observations. SN\,1987A is close to the mass-integrated mean CCSN mass for most IMFs and is the best-observed CCSN remnant known. Assuming SN\,1987A as a good mass-weighted mean for the dust production, and using the preferred value of a carbonaceous and silicate grain mix of $0.6$ to $0.7$\,M$_{\odot}$ \citep{Matsuura11,Indebetouw14}, we can infer that at least $2\times10^{8}$\,M$_{\odot}$ CCSNe exploded in HFLS3 to account for the dust mass of $3 \times 10^8$ M$_{\odot}$. The stellar-to-dust mass ratio should be around 100 for a Chabrier IMF, and a factor of two larger than this for a Salpeter IMF. The precise value of this ratio depends on how CCSNe produce refractory metals as a function of mass. When considering model uncertainties, the stellar-to-dust mass ratio is within 20\% of the value quoted above, where the dust masses are tied through the observational pivot point provided by the dust mass observed in SN\,1987A. Note that this argument is currently based on the dust mass observed in SN\,1987A as an indication of the refractory element production, rather than claiming that CCSNe necessarily produce all the dust directly. But since the dust mass observed is believed to be close to the maximal dust production for this SN \citep{Indebetouw14}, it is therefore a reasonable reflection of the most dust we could ultimately expect to be produced by the elements synthesised by SN\,1987A. Thus, for $3\times10^{8}$\,M$_{\odot}$ mass of dust in the galaxy, we expect a minimum of $\sim 2\times10^{10}$\,M$_{\odot}$ mass of stars for a Chabrier IMF, and twice this for a Salpeter IMF. This is comparable to the lensing magnification-corrected stellar mass inferred from {\sc Magphys} at $5\times 10^{10}$ M$_{\odot}$, though we note once again that this value has a large uncertainty due to various assumptions and low signal-to-noise ratio of the rest-frame UV measurements. For a SFR averaged over 100 Myr of about 660\,M$_{\odot}$\,yr$^{-1}$, the above arguments imply a characteristic dust production time of at least 40\,Myr, assuming a negligible dust destruction during the same period. This is lower than the suggested lifetime for dust mass assembly in AzTEC-3 of about 200 Myr \citep{Dwek11}. While our current estimates are uncertain, the above argument, however, can be strengthened in the future with more precise measurements of dust and stellar masses to constrain dust production mechanisms at $ z\sim 6$. We also attempted a SED model with far-IR/sub-mm data points combined with rest-UV fluxes from R1, with peak emission 0.5$''$ to the South-West of HFLS3 (Figure~1). This emission is detected in all three WFC3 bands at significances greater tha 6 $\sigma$ in each, although the emission remains undetected in ACS. The emission, however, is blended in IRAC data with the near-IR emission from the two galaxies (Figure~4). The {\sc Magphys} fit was considerably poor as there was no consistent SED that can fit the four orders of magnitude in wavelength from UV to sub-mm in that case with the best-fit case having a reduced $\chi^2$ of greater than 5. This ruled out a scenario in which HFLS3 sub-mm emission is associated with R1. It also rules out an extreme scenario in which our relative astrometry between IRAM/PdBI and {\it Hubble} images are wrong such that the near-IR counterpart to HFLS3 is R1. We find that R1 must be a separate source. The {\sc Hyperz} SED model shown in Figure~5 leads to a photometric redshift for this emission is consistent with a source at $z \sim 6.3$ (Figure~2 bottom panel), though a dusty galaxy SED at $z \sim 2$ is also consistent with this emission. The {\sc Hyperz} fit to the data leads to a stellar mass of $\sim 1.2 \times 10^{10}$ M$_{\odot}$ for R1 if we assume the redshift is $z=6.3$, following the $z \sim 6$ photo-z solution. We have two possibilities for this new source. It could be part of the emission associated with a complex galaxy merger system involving HFLS3, especially if HFLS3 starburst is triggered by a merger as is the case for most $z \sim 2$ bright SMGs. Alternatively, it could be part of the $z \sim 2.1$ foreground structure that is responsible for lensing of HFLS3. If the latter is indeed the case, the region in the foreground of HFLS3 involves a massive galaxy group, but the magnification upper limit of 3.7 we have derived here is unlikely to be revised higher as it accounts for a wide variation of model parameters, including to the total lens mass in the foreground. It is far more likely that R1 is part of the complex merger system associated with HFLS3. | 14 | 4 | 1404.1378 |
1404 | 1404.6244_arXiv.txt | Supergravity models of natural inflation and its generalizations are presented. These models are special examples of the class of supergravity models proposed in \cite{Kallosh:2010ug}, which have a shift symmetric \K\, potential, superpotential linear in goldstino, and stable Minkowski vacua. We present a class of supergravity models with arbitrary potentials modulated by sinusoidal oscillations, similar to the potentials associated with axion monodromy models. We show that one can implement natural inflation in supergravity even in the models of a single axion field with axion parameters $O(1)$. We also discuss the irrational axion landscape in supergravity, which describes a potential with infinite number of stable Minkowski and metastable dS minima. | It appears that one of the popular models of inflation, called natural inflation \cite{Freese:1990rb}, which was proposed 24 years ago, has not yet been generalized to supergravity with stabilization of all moduli. The goal is to find a supergravity model that would lead to the natural inflation potential of the axion field $\phi$ \be V=\Lambda^4 \Bigl(1-\cos{\phi\over f}\Bigr) \label{Natural}\ee with Minkowski minimum at $\phi=0$. The supergravity axion valley models proposed and studied in \cite{Kallosh:2007ig,Kallosh:2007cc}, and used more recently in \cite{Czerny:2014qqa}, almost did the job. They have the following \K\, potential and superpotential \be\label{ax} K= {(T+\bar T)^2\over 2}\, , \qquad W= W_0 + A e^{-aT} + B e^{-bT} \ . \ee The real part of the modulus is stabilized in this model and the imaginary part plays the role of the light axion $\phi$. The resulting potential is almost of the form \rf{Natural}. However, in this class of models the minimum of the potential is in AdS space. Therefore one has to specify an uplifting procedure, which uplifts the AdS minimum to a Minkowski one, or even better, to a de Sitter minimum with a tiny cosmological constant. Various uplifting procedures have been proposed over the years, but some of them cannot be described at the supergravity level, whereas some others may lead to modification of the functional form of the potential upon uplifting. As a result, to the best of our knowledge, explicit supergravity models realizing such an uplifting in a way consistent with moduli stabilization and leading to natural inflation \rf{Natural} are still unavailable. For a recent discussion of the axion inflation models see for instance \cite{Baumann:2014nda} and \cite{Pajer:2013fsa}. The purpose of this note is to present a very simple supergravity model with non-negative potential which upon stabilization of the non-inflaton moduli produces the natural inflation potential \rf{Natural}. It will be achieved in the context of the general class of models \cite{Kallosh:2010ug} describing chaotic inflation in supergravity. This class of models generalized the supergravity realization of the simplest chaotic inflation scenario ${m^{2}\over 2}\phi^{2}$ proposed in \cite{Kawasaki:2000yn}. The class of models developed in \cite{Kallosh:2010ug} has a built-in feature which makes the potential non-negative. The superpotential in these models is linear in the goldstino superfield $S$, whereas the \K\, potential is some function of either $\Phi + \bar \Phi$ or $\Phi + \bar \Phi$, and of $S\bar S$: \be W = S f(\Phi) \ , \qquad K = K ((\Phi\pm \bar \Phi)^2, S\bar S) \ . \ee The \K\, potential $K ((\Phi\pm \bar \Phi)^2, S\bar S)$ does not depend on one of the combinations $(\Phi\mp \bar \Phi)$, which plays the role of the inflaton field in this scenario. If one can stabilize the field $S$ at $S = 0$, then $W = 0$, and the potential becomes manifestly non-negative: \be V= e^K (|DW|^2 -3W^2)|_{S=0}= e^K \partial_S W \partial_{\bar S} \bar W \geq 0 \ . \ee If, in addition, one can ensure that one of the combinations of the fields $(\Phi\pm \bar \Phi)$, which is orthogonal to the inflaton field, vanishes during inflation, then the inflaton potential becomes \be V= |f(\Phi)|^2 \ . \label{general}\ee The required stabilization conditions are rather mild, which allows to have a functional freedom in the choice of the inflaton potential in supergravity \cite{Kallosh:2010ug}. As we will see, this class of models can easily incorporate natural inflation. Moreover, by a simple extension of the supergravity versions of natural inflation, one can find a family of positive definite inflationary potentials of arbitrary shape modulated by sinusoidal oscillations. These potentials are similar to the string theory inflaton potentials associated with axion monodromy \cite{Silverstein:2008sg,Flauger:2009ab}. | 14 | 4 | 1404.6244 |
|
1404 | 1404.4288_arXiv.txt | {We demonstrate with a nonlinear MHD code that {angular momentum can be transported} due to the magnetic instability of toroidal fields under the influence of differential rotation, and that the resulting effective viscosity may be high enough to explain the almost rigid-body rotation observed in radiative stellar cores. Only stationary current-free fields and only those combinations of rotation rates and magnetic field amplitudes { which provide maximal numerical values of the viscosity are considered}. We find that the dimensionless { ratio of the effective over molecular viscosity}, $\nu_{\rm T}/\nu$, linearly grows with the Reynolds number of the rotating fluid multiplied with the square-root of the magnetic Prandtl number -- which is of order unity for the considered red sub-giant KIC 7341231. For the considered interval of magnetic Reynolds numbers -- which is restricted by numerical constraints of the nonlinear MHD code -- there is a remarkable influence of the magnetic Prandtl number on the { relative importance of the} contributions of the Reynolds stress and the Maxwell stress to the total viscosity, which is magnetically dominated { only} for $\rm Pm\gsim 0.5$. We also find { that} the magnetized { plasma behaves as a non-Newtonian fluid}, i.e. the resulting effective viscosity depends on the shear in the rotation law. The decay time of { the differential rotation} thus depends on its shear and becomes { longer and longer} during the spin-down of a stellar core.} | Model calculations for the formation of red giants without turbulent or magnetic angular momentum transport lead to rather steep radial profiles of the angular velocity in the innermost core of the star. Ceillier et al. (2012) report a theoretical (quasi-Keplerian) profile $\Om \propto R^{-q}$ with $q\simeq 1.6$ for the low-mass red giant KIC 7341231. Note that the exponent $q=1$ would describe a quasi-galactic rotation profile with $U_\phi=$ const while $q=2$ represents the rotation law for uniform angular momentum $\Om R^2\approx$ const. The {\em Kepler} data, however, lead to much flatter rotation laws of the observed red giants: the cores of several sub-giants and young red giants seem to rotate only (say) five times faster than the outer convection zone (Deheuvels et al. 2012, Deheuvels et al. 2014). Eggenberger et al. (2012) argue that only an additional viscosity of $3 \times 10^4$ cm$^2$/s may explain the unexpectedly flat internal rotation law of the more massive red giant KIC 8366239. The outward flux of angular momentum { due to this enhanced} viscosity,{ which exceeds} the molecular value by a factor of $\nu_{T}/\nu \approx 500$, suffices to produce the observed spin-down of the inner radiative core. R\"udiger \& Kitchatinov (1996) needed just this viscosity value to produce by Maxwell stress the high degree of uniformity of the internal solar rotation, derived from helioseismologic measurements down to $0.15 R_\odot$. Rotation laws with $q<2$ are hydrodynamically stable. Under the presence of toroidal fields, however, they become unstable against nonaxisymmetric disturbances if the amplitude of the toroidal field is high enough but does not exceed $\Om_{\rm A}\simeq \Om$ with $\Om_{\rm A}=B_\phi/\sqrt{\mu_0 \rho R^2}$ as the Alfv\'{e}n frequency for incompressible fluids. This nonaxisymmetric instability even exists for toroidal fields which are current-free in the considered domain. Because of this force-free character of the magnetic field, the instability has been called the `azimuthal magnetorotational instability' (AMRI). Within a cylindrical setup, AMRI has been studied in the linear approximation, but also by nonlinear simulations (R\"udiger et al. 2014). The consequences of both compressibility and heat transport (see Spruit 2002) cannot be studied with the present model. We know, however, that these influences become negligible for strong fields. One also can show that, with thermodynamics included, the radial components of flow and field are strongly damped, so that the resulting angular momentum transport should be reduced by the `negative buoyancy'. The viscosity values derived in the present paper are thus maximum values. If they are not high enough for $\nu_{\rm T}/\nu\simeq 500$ then the concept of the instability of magnetic fields in the stellar interior is proven as not working. An important basis for realistic numerical simulations is the knowledge of the magnetic Prandtl number \begin{equation} {\rm Pm}=\frac{\nu}{\eta}, \label{Pm} \end{equation} where $\nu$ is the molecular viscosity of the fluid and $\eta$ its magnetic diffusivity. So far, numerical nonlinear simulations are only possible for $\rm Pm$ exceeding (say) 0.01. The magnetic Prandtl number of the plasma inside main-sequence stars, however, is smaller (see Brandenburg \& Subramanian 2005). We have thus first to probe the value of $\rm Pm$ in the radiative interiors of the considered red {\em Kepler} stars. | In the sense of a proof of existence, nonaxisymmetric magnetic instabilities under the influence of differential rotation are { shown} to transport angular momentum in a direction orthogonal to the rotation axis. The toroidal background field is chosen as current-free (outside the axis), so that it cannot decay. The mass density is assumed as uniform, suppressing the influence of the buoyancy. The model is thus not applicable to solar type stars { on the main sequence}, but only to the very hot radiative cores of (sub)giants. It is also shown, by use of a numerical stellar model, that the high temperatures in such cores { result in} microscopic magnetic Prandtl numbers varying between $0.1$ and $10$. { This makes it possible to use} a nonlinear MHD code which only works for such Prandtl numbers. The code solves the nonlinear and nonaxisymmetric MHD differential equations in a cylindric setup and provides the $R$-$\phi$-component of the complete stress tensor (Reynolds stress plus Maxwell stress), and averages the resulting number over the azimuth and the axial coordinate. A spectral element code is used based on the hydrodynamic code of Fournier et al. (2005). The solutions are expanded into Fourier modes in the azimuthal direction. The remaining meridional problems are solved with a Legendre spectral method. If a linear code is used to find the classical eigenvalues for the onset of the marginal instability then all solutions are optimized in the wave number $k$ yielding the smallest possible Reynolds numbers. The boundary conditions at the cylinder walls are assumed to be no-slip and perfectly conducting. The resulting cross correlations thus vanish { at the surfaces of} the two cylinders, { while they have} a maximum { near the} center of the gap between the cylinders (Fig. \ref{fig22}). These { maximum values} are calculated and transformed into viscosity values by means of Eq. (\ref{uRuphi}). The resulting viscosity values are compared { for various} Hartmann numbers in the instability cone but for the same Reynolds number (see Fig. \ref{fig1}). If the same procedure is done for various Reynolds numbers and various magnetic Prandtl numbers, one obtains the results presented in Fig. \ref{fig31} for two different rotation laws. Note that the data for different Reynolds numbers do {\em not} belong to the same Hartmann number. The { figures} do not, therefore, contain evolutionary scenarios. As expected, the consequence of the results given in Fig. \ref{fig31} is that the viscosity for given shear linearly grows with the angular velocity. The most interesting result, however, is the dependence of the effective viscosity on the shear: the steeper the rotation law the higher the viscosity value. Hence, the decay of a nonuniform stellar rotation law is a nonlinear process. The decay time-scale does not remain constant in time as it becomes larger and larger. The non-Newtonian behavior of the magnetized conducting fluid is basically connected with the mechanism of the azimuthal magnetorotational instability which exists only for differential rotation. It { thus leads} automatically to saturation during stellar spin-down process, which is most violent at its beginning and becomes slower { at later times}. As mentioned, the toroidal magnetic field with the current-free radial profile does not dissipate. As we thus fix { the magnetic field} for a given magnetic Prandtl number and consider the dependence of the effective viscosity on the angular velocity, the effective viscosity will grow, passes its maximum close to the line of maximum growth rate shown in Fig. \ref{fig1}, and again sinks to zero at the lowest possible Reynolds number. Figure \ref{fig31} demonstrates that, up to Reynolds numbers of the order of $10^3$, along the line of maximum growth rate there is no saturation of the effective viscosities. The stronger the fields, the higher is the { maximum viscosity which can be achieved}. While the instability domain is only slightly modified for Prandtl numbers smaller than unity, this is not true for the behavior of the second order correlations. The effective viscosity formed by Reynolds stress and Maxwell stress { decreases} for small $\rm Pm$. It varies by one order of magnitude when $\rm Pm$ varies by two orders of magnitude. The ratio of magnetic energy to kinetic energy (taken for maximum viscosity) { also decreases} for small $\rm Pm$ (Fig. \ref{fig32}, top). Consequently, the Maxwell stress in relation to the Reynolds stress also sinks for small magnetic Prandtl numbers (Fig. \ref{fig32}, bottom). The { effectiveness} of the magnetic perturbations in transporting angular momentum { is thus reduced} for small $\rm Pm$, i.e. with cooler temperatures. Our results support the conclusion that magnetic instabilities of toroidal magnetic fields in the presence of differential rotation are a viable mechanism to explain angular momentum redistribution in stellar interiors, especially for sub-giant and young red giant stars. | 14 | 4 | 1404.4288 |
1404 | 1404.2099_arXiv.txt | {Stellar properties and in particular stellar radii of exoplanet host-stars are essential for measuring properties of exoplanets. Therefore it is becoming increasingly important to be able to supply reliable stellar radii fast. Grid-modelling is an obvious choice to do this, but that only offers a low degree of transparency to the non-experts.} {Here we present a new method to obtain stellar properties for stars exhibiting solar-like oscillations in an easy, fast, and transparent way. The method, called Asteroseismology Made Easy (AME), can determine stellar masses, mean-densities, radii, and surface gravities, as well as estimate ages. In this writing we present AME as a visual and powerful tool which could be useful; in particular in the light of the large number of exoplanets being found.} {AME consists of a set of figures from which the stellar parameters are deduced. These figures are made from a grid of stellar evolutionary models that cover masses ranging from $0.7 M_\sun$ to $1.6 M_\sun$ in steps of $0.1 M_\sun$ and metallicities in the interval $-0.3 \ \mathrm{dex} \leq [\mathrm{Fe}/\mathrm{H}] \leq +0.3 \ \mathrm{dex}$ in increments of $0.1 \ \mathrm{dex}$. The stellar evolutionary models are computed using the Modules for Experiments in Stellar Astrophysics (MESA) code with simple input physics.} {We have compared the results from AME with results for three groups of stars; stars with radii determined from interferometry (and measured parallaxes), stars with radii determined from measurements of their parallaxes (and calculated angular diameters), and stars with results based on the modelling of their individual oscillation frequencies. We find that a comparison of the radii from interferometry to those from AME yield a weighted mean of the fractional differences of just $2\%$. This is also the level of deviation that we find when we compare the parallax-based radii to the radii determined from AME.} {The comparison between independently determined stellar parameters and those found using AME show that our method can provide reliable stellar masses, radii, and ages, with median uncertainties in the order of $4\%$, $2\%$, and $25\%$ respectively.} | \label{sec:intro} NASA's \textit{Kepler Mission} \citep{ref:kepler_koch} has provided photometric light curves of high precision which has enabled the detection of solar-like oscillations in more than 500 stars \citep{ref:kepler_chaplin}. Asteroseismic scaling relations are generally used to determine properties of such stars \citep[for a recent review on those issues see][]{ref:ast_sca_chaplin}. If one considers stars in hydrostatic equilibrium, it follows from homology that since the dynamical time scale is proportional to ${\bar{\rho}}^{-1/2}$ (where $\bar{\rho}$ is the mean-density) the oscillation frequencies will be proportional to ${\bar{\rho}}^{1/2}$. In the asymptotic approximation \citep[see e.g. page 218 of][]{ref:ast_book} the frequency spectrum for a star can be written as a function of radial order ($n$) and angular degree ($\ell$) as \begin{equation} \label{eq:asymp_rel} \nu_{n,\ell} = \Delta\nu \cdot \left( n + \ell /2 + \varepsilon \right) - \ell \left( \ell + 1 \right) \cdot D_0 \ . \end{equation} Since the frequencies scale with the square root of the mean-density, the main regularity in the asymptotic relation (expression~(\ref{eq:asymp_rel})), the large separation ($\Delta\nu$), will therefore be proportional to $\sqrt{\bar{\rho}}$. The parameters $D_0$ and $\varepsilon$ will depend on the detailed internal structure of a given star and can e.g. be used to estimate its age. The large separation can be found from the sound-speed integral: \begin{equation} \label{eq:sound_speed_integral} \Delta\nu = \left( 2 \int \limits_0^R \frac{\mathrm{d}r}{c(r)} \right) ^{-1} \ , \end{equation} where $R$ is the stellar radius and $c(r)$ the sound-speed in the stellar interior. We will in this study use this most fundamental asteroseismic parameter $\Delta\nu$ to estimate the stellar mean-density. It has been common to also use the frequency at maximum oscillation power ($\nu_\mathrm{max}$) to measure the gravity of a given star since observationally it has been shown to be proportional to the acoustic cut-off frequency in the atmosphere (which scales with $g/\sqrt{T_\mathrm{eff}}$). While the observational link to the acoustic cut-off frequency is established \citep{ref:stello_2009} we still need to better understand the theoretical background for this relation \citep{ref:belkacem_2011,ref:belkacem_2013}. An issue in relation to the use of $\nu_\mathrm{max}$ is that there is a tight relation between $\nu_\mathrm{max}$ and $\Delta\nu$ \citep[see][]{ref:stello_2009}. This is the reason why we in the present study only consider $\Delta\nu$ and not $\nu_\mathrm{max}$. Obtaining stellar parameters such as mass, radius, and age through detailed modelling (modelling based on individual oscillation frequencies) for just one star showing solar-like pulsations can be a very time-consuming task. As an alternative, one or more of the analysis pipelines described in for instance \citet{ref:pip_seek} (SEEK), \citet{ref:pip_mathur}, \citet{ref:pip_radius} (RADIUS), \citet{ref:pip_yb}, and \citet{ref:pip_ybplus} can be utilised. However, this grid-based modelling does not offer a very transparent procedure and inconsistencies can occur between the mean-density derived using the large separation ($\bar{\rho} \propto \Delta \nu ^2$) and the mean-density found using the grid-determined mass and radius ($\bar{\rho} \propto M/R^3$). This is primarily an effect of mass and radius being determined individually \citep{ref:pip_seek} and the fact that grid-based modelling sometimes yield bimodal parameter distributions. \citet{ref:pip_white} discuss the potential of constraining stellar masses and ages for main-sequence (MS) and sub-giant solar-like oscillators using diagrams based on their pulsation properties, such as the Christensen-Dalsgaard and the $\varepsilon$ diagrams. However, $\varepsilon$ can be hard to constrain \citep{ref:white_epsilon}, and the Christensen-Dalsgaard diagram requires the value of the small separation ($\delta \nu$) to be known and this is usually not available. The reason for this is that to find $\delta\nu$ some modes with $\ell=2$ need to be discernible in the power spectrum, which is often not the case. Therefore we have developed a method to obtain self-consistent mean-densities along with other stellar parameters in an easy and fast way for MS and sub-giant stars showing solar-like oscillations. We consider this method to be transparent because it is easy to follow, from considering the figures, what effect a small change in one of the input parameters would have on the result. By scaling to the Sun we have made our method less sensitive to the stellar models and evolution code used. We call this method AME - Asteroseismology Made Easy. The aim of this paper is to present AME as a powerful and visual tool which does not need the frequency of maximum power to yield basic stellar properties. | \label{sec:conclusion} We have shown that AME as a tool to determine basic stellar parameters works very well. It offers a transparent and easy way to find stellar properties for stars exhibiting solar-like oscillations at a level where the large separation can be detected. The full set of AME figures which are the backbone of the method can be found in the appendix~\ref{om_sec:morefigs}. They cover a range in masses from $0.7 M_\sun$ in steps of $0.1 M_\sun$ to $1.6 M_\sun$, and in metallicity the span is $-0.3 \ \mathrm{dex} \leq [\mathrm{Fe}/\mathrm{H}] \leq +0.3 \ \mathrm{dex}$ in increments of $0.1 \ \mathrm{dex}$. We have used AME on a total of $43$ stars with stellar parameters determined in different ways. We have compared the results from AME to those found in the literature and in some cases to results obtained with the BaSTI grid. The overall agreement is good and the differences can be explained. AME was found to perform as well as the BaSTI grid in terms of agreement with the values quoted in the literature for the $16$ stars with results from detailed modelling. Furthermore, the radii from AME was found to be consistent with both radii from interferometry (and measured parallaxes) and radii determined from parallaxes (and computed angular diameters). In the future, the AME grid will be extended to cover a wider range of metallicities and masses and potentially made denser. We are currently working on a web interface to enable people to use AME without having to plot their stars in the diagrams themselves. | 14 | 4 | 1404.2099 |
1404 | 1404.5446_arXiv.txt | \noindent We consider a decaying magnetic dark matter explaining the X-ray line at $3.55\,{\rm keV}$ shown recently from XMM-Newton observations. We introduce two singlet Majorana fermions that have almost degenerate masses and fermion-portal couplings with a charged scalar of weak scale mass. In our model, an approximate $Z_2$ symmetry gives rise to a tiny transition magnetic moment between the Majorana fermions at one loop. The heavier Majorana fermion becomes a thermal dark matter due to the sizable fermion-portal coupling to the SM charged fermions. We find the parameter space for the masses of dark matter and charged scalar and their couplings, being consistent with both the relic density and the X-ray line. Various phenomenological constraints on the model are also discussed. | Dark matter is a dominant component of the matter density in the universe, playing a crucial role in the structure formation and explaining the flatness of galaxy rotation curves, etc. The evidences for dark matter are without question but there is no understanding of the properties of dark matter such as its mass or coupling to the SM particle, except gravitational interaction. Recently, there was an interesting indication of dark matter \cite{xray} from the stacked X-ray spectrum of galaxies and clusters \cite{xmm}, which shows an unexplained line signal at the energy $3.55\,{\rm keV}$. A sterile neutrino having a very small mixing with the active neutrinos, can explain the X-ray line signal by a small transition magnetic moment but the link to the generation of neutrino masses via see-saw mechanism is unclear due to a small mixing \cite{xray}. The model building issue with keV sterile neutrino was discussed in Ref.~\cite{keVnu}. There have been more candidates suggested for dark matter after the X-ray line was identified \cite{models}. In this work, we consider a decaying dark matter model with two Majorana fermions. Similarly to the sterile neutrino case, the X-ray line can be obtained from the decay of the heavier Majorana fermion, for a sufficiently small magnetic transition dipole moment and a mass difference of $3.55\,{\rm keV}$ between the two Majorana fermion masses. To this purpose, we propose a microscopic model for dark matter containing the fermion-portal couplings with a new charged scalar \cite{ibarra,bai}. We introduce an approximate $Z_2$ symmetry for the long-lived dark matter and a $U(1)_X$ global symmetry for controlling the $Z_2$ breaking to a small amount. As a consequence, a fermion-portal coupling for the heavier Majorana fermion preserves $Z_2$ and is sizable while a tiny coupling for the lighter Majorana fermion is produced after an explicit breaking of $Z_2$. Then, when fermion-portal couplings break CP, a tiny transition magnetic moment between two Majorana fermions is degenerated at one loop \cite{RHneutrino}, even for the charged scalar of weak scale mass. Therefore, the heavier Majorana fermion is sufficiently long-lived for explaining the X-ray line and it can be thermally produced due to the sizable fermion-portal coupling to the SM charged fermion. We study the parameter space for the masses and couplings of dark matter and the charged scalar, that are consistent with both the X-ray line and the dark matter relic density in our model. We also discuss the phenomenological constraints on the model, coming from indirect and direct detection experiments, precision measurements such as the anomalous magnetic moment of muon, and collider experiments. The paper is organized as follows. We begin with the properties of magnetic dark matter explaining the X-ray line and describe a microscopic model for the magnetic dark matter in a simple extension of the SM with a charged scalar. Then, we impose the condition for the relic density in our model and discuss the compatibility with the X-ray line signal. Next, various phenomenological constraints on the model are given. Finally, conclusions are drawn. | We have studied a simple model of dark matter with two singlet Majorana fermions, explaining the X-ray line as the decay product of dark matter through a transition magnetic moment. We showed that a tiny value of the transition magnetic moment is generated at one loop by the interplay between the $CP$ violating fermion-portal couplings of a charged scalar to the two Majorana fermions: one to the heavier state preserves $Z_2$ and the other to the lighter state breaks $Z_2$ by a tiny amount. Therefore, even with a weak-scale charged scalar, the heavier Majorana fermion decays into the lighter one with a sufficiently long lifetime, emitting the X-ray line at $3.55\,{\rm keV}$. Moreover, the decaying Majorana fermion can be a thermal dark matter due the sizable $Z_2$-symmetric coupling to the SM charged fermions. It was shown that the proposed model can satisfy the current bounds from the relic density and various experiments and it can be testable in the upgraded LHC. | 14 | 4 | 1404.5446 |
1404 | 1404.6735_arXiv.txt | {} {We study the effects of moderate deviations from solar abundances upon the atmospheric structure and colours of typical Mira variables.} {We present two model series of dynamical opacity-sampling models of Mira variables which have (1) $\frac{1}{3}$ solar metallicity and (2) ``mild'' S-type C/O abundance ratio ([C/O]=0.9) with typical Zr enhancement (solar +1.0). These series are compared to a previously studied solar-abundance series which has similar fundamental parameters (mass, luminosity, period, radius) that are close to those of $o$~Cet.} {Both series show noticeable effects of abundance upon stratifications and infrared colours but cycle-to-cycle differences mask these effects at most pulsation phases, with the exception of a narrow-water-filter colour near minimum phase.} {} | The density stratification of upper atmospheric layers of Mira variables is determined by outwards travelling shock fronts. These shock fronts are seen in typical emission lines of hydrogen and other atoms (Fox et al. 1984; Richter \& Wood 2004) and lead to a geometrically very extended stellar atmosphere resulting in strong dependence of the Mira diameter observed in different absorption features (e.g. Ireland et al. 2004; Woodruff et al. 2008, 2009; Zhao-Geisler et al. 2012). Models based on sophisticated treatment of radiative transfer have become available in recent years (H{\"o}fner et al. 2003; Ireland et al. 2008 ({\tt CODEX1}), 2011 ({\tt CODEX2})). The {\tt CODEX} model series, which are based upon a self-excited pulsation model for each series, show that differences in position and strength of outward travelling - or sometimes receding - shock fronts in different cycles may lead to noticeable cycle-to-cycle differences of density and temperature in upper layers and to differences of spectral features formed in these layers. Comparison of these models with observations of Miras were published by Woodruff et al. (2009), Wittkowski et al. (2011), Hillen et al. (2012). {\tt CODEX} models have so far been computed for 4 different sets of parameters given by 4 non-pulsating ``parent stars'' (series R52, C50, C81, o54; see Tab.~1). Solar element abundances (Z = 0.02; Grevesse et al. 1996) were adopted for all series. Details of constructing pulsation models and of computing atmospheric temperatures and spectra, based upon an opacity-sampling treatment of absorption coefficients, are given in Ireland et al. (2008, 2011). In this paper, we construct two {\tt CODEX} model series with (i) ``mildly'' sub-solar metallicity and (ii) ``mild'' S-type C/O abundance ratio. In order to look for abundance-dependent differences, we compare stratifications and colours of these models with those of models of the o54 series which has, except for abundances, almost identical parameters. We note that intermediate-period Mira variables are most typically associated with the thick disk (Groenewegen \& Blommaert 2006) which does not have solar abundances and, in particular, has [Ti/Fe] enhancement of +0.2 and [O/Fe] and [C/Fe] enhancements of +0.3 to +0.4 (e.g. Reddy et al. 2006). Assuming a standard [$\alpha$/Fe]=0.0, our chosen metallicity Z=0.02/3 corresponds to [Ti/H]=-0.5, [O/H]=-0.5 and [C/H]=-0.5. Given that Ti, O and C are the heavy elements that most significantly influence the spectra of O-rich Mira variables, the metallicity Z=0.02/3 can thus approximately represent a thin-disk star of [Fe/H]=-0.5 as well as a thick-disk star with typical $\alpha$-enhancement of +0.3 and with [Fe/H]=-0.8. This range is typical of stars kinematically associated with the thick disk (e.g. Adibekyan et al. 2011, Cheng et al. 2012). Of course, significant abundance variations are still expected throughout the intermediate-period Mira variable population. This study attempts to look for simple observational effects of these variations. | We computed two model series which allow the study of the effects of ``mild'' deviations from solar element abundances upon the atmospheric stratification and standard spectral colours of typical ($o$~Cet-like) Mira variables. The x54 model series is the $\frac{1}{3}$ metallicity counterpart of the solar-metallicity o54 series discussed by Ireland et al. (2008, 2011), while the s54 model series is the S-type counterpart of the o54 series. Both model series show noticeable stratification and spectral-colour differences from the o54 models. It turns out, however, that such abundance effects are readily masked by significant cycle-to-cycle differences at most phases. Also, in real stars, one needs to consider in addition the effects caused by modest differences of stellar parameters, so one has to conclude that ``mild'' deviations from solar element abundances will barely be detectable in the structure of the stellar atmosphere and the broadband emitted spectral flux. This confirms the findings of Scholz \& Wood (2004) who infer, based on a series of less elaborated models, that there is no easy way to determine the metallicity of an M-type Mira field star for moderate deviations (factor of the order of 2) from solar metallicity. We suggest that future observational campaigns should both focus on phase-dependent measurements in the W water filter and on high-spectral-resolution bandpasses that include a continuum (e.g. J and H bands) and where the classical method of line analysis can be attempted. | 14 | 4 | 1404.6735 |
1404 | 1404.1265_arXiv.txt | We explore a new type of entropic mechanism for generating density perturbations in a contracting phase in which there are two scalar fields, but only one has a steep negative potential. This first field dominates the energy density and is the source of the ekpyrotic equation of state. The second field has a negligible potential, but its kinetic energy density is coupled to the first field with a non-linear sigma-model type interaction. We show that for any ekpyrotic equation of state it is possible to choose the potential and the kinetic coupling such that exactly scale-invariant (or nearly scale-invariant) entropy perturbations are produced. The corresponding background solutions are stable, and the bispectrum of the entropy perturbations vanishes as no non-Gaussianity is produced during the ekpyrotic phase. Hence, the only contribution to non-Gaussianity comes from the non-linearity of the conversion process during which entropic perturbations are turned into adiabatic ones, resulting in a local non-Gaussianity parameter $f_{\textsc{nl}} \sim 5$. | Recent {\it Planck} satellite measurements \cite{Ade:2013lta,Ade:2013rta,Ade:2013ydc}, together with earlier observations from WMAP, ACT, SPT, and other experiments \cite{Sievers:2013ica}, showed with high precision that the spectrum of primordial density fluctuations is nearly scale-invariant, Gaussian, and adiabatic. The currently best known mechanisms for the generation of cosmological perturbations are inflation \cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi} and ekpyrosis \cite{Khoury:2001wf}. Inflation is a period of accelerated expansion following the big bang, characterized by a large Hubble parameter $H$ and an equation of state $w \approx -1$. Ekpyrosis is a period of ultra-slow contraction preceding the big bang, characterized by a small $H$ and $w > 1$. A distinctive feature of all currently known ekpyrotic models is that, during the ekpyrotic contraction phase, gravitational waves are not amplified. More specifically, the gravitational waves have a blue spectrum, but their quantum state does not become squeezed; hence, they cannot be given a classical interpretation \cite{Tseng:2012qd,Battarra:2013cha}. During the expanding phase, the (classical) scalar curvature perturbations provide a source for gravitational waves at second order in perturbation theory, leading to a small-amplitude gravitational wave background \cite{Baumann:2007zm}. This small-amplitude background is not compatible with a tensor-to-scalar ratio $r \sim 0.2$, as reported by the BICEP2 collaboration \cite{Ade:2014xna}. However, questions have been raised about the BICEP2 claim \cite{Flauger:2014qra} and it will take other ongoing experiments to determine if there really exist any detectable tensor B-modes. Hence, in the meantime, it is reasonable to assume the Planck2013 bound on $r$ and continue studying cyclic/ekpyrotic scenarios, given their conceptual advantages regarding a number of important open issues in early universe cosmology (such as the initial conditions and measure problems). Furthermore, it is conceivable that a detectable gravitational wave spectrum can be created in the context of cyclic/ekpyrotic scenarios, {\it e.g.} by phase transitions, topological defects or during the bounce -- these avenues remain to be explored. Using the field picture, the ekpyrotic phase can be described by a scalar field, $\phi$, rolling down a steep negative potential \begin{equation} V = - V_0 e^{-\sqrt{2\epsilon}\phi} , \end{equation} where $V_0$ is a constant and $\epsilon$ denotes the equation-of-state parameter \begin{equation}\label{es} \epsilon \equiv \frac{3}{2}\left(1 + w \right)\quad \text{with}\quad w \equiv \frac{\rho_S}{p_S}\,, \end{equation} where $w$ is the equation of state, $\rho_S$ the energy density, and $p_S$ the pressure of the smoothing background component. It has been shown \cite{Tseng:2012qd,Battarra:2013cha} that, if there is only a single field in the contracting phase, the (adiabatic) perturbations are not amplified and cannot be the seed of structure in the post-bang universe. The currently best-understood way around this problem is the entropic mechanism, where pre-bang isocurvature fluctuations are generated by adding a second ekpyrotic field, $\phi_2$ \cite{Notari:2002yc,Finelli:2002we,Lehners:2007ac,Buchbinder:2007ad}. These isocurvature modes are then converted into density perturbations which source structure in the post-bang universe. A simple example of an action describing the standard ekpyrotic mechanism is \begin{eqnarray}\label{old_action} S &=& \int d^4 x \sqrt{-g}\frac{R}{2}\nonumber\\ &-& \int d^4 x \sqrt{-g}\bigg(\frac{1}{2}\partial_{\mu}\phi_1\partial^{\mu}\phi_1+V_1 e^{-c_1\phi_1}\bigg) \nonumber\\ &-&\int d^4 x \sqrt{-g}\bigg(\frac{1}{2}\partial_{\mu}\phi_2\partial^{\mu}\phi_2 + V_2 e^{-c_2\phi_2}\bigg), \end{eqnarray} where $V_1, V_2, c_1, c_2$ are constants and the two fields have separate ekpyrotic potentials. (Here and throughout this paper we choose units such that $M_{\text{Pl}}^2 \equiv 1$, where $M_{\text{Pl}}^2 = (8\pi \mathrm{G})^{-1}$ is the reduced Planck mass and $\mathrm{G}$ is Newton's constant.) The background evolution is determined by the linear combination of these potentials, or equivalently, after performing a rotation in field space, by the adiabatic field, $\sigma$, (defined to point tangentially along the background trajectory, with $\dot{\sigma} = (\dot{\phi_1}^2+\dot{\phi_2}^2)^{1/2}$) while the evolution of perturbations is governed by the entropy field, $s$ (which is, by definition, perpendicular to the $\sigma$-field). At the end of the ekpyrotic phase and before the bounce, the background trajectory bends and the isocurvature perturbations are converted into adiabatic ones. However, it is well-known that these ekyprotic solutions for $\phi_1$ and $\phi_2$ are {\it unstable}, in that the $\sigma$ direction runs along a ridge in the potential that is unstable to variations in the $s$ direction (possible consequences in a cyclic context were discussed in \cite{Lehners:2009eg,Lehners:2011ig}). Also, to obtain nearly scale-invariant spectra requires a steep negative potential which results in the generation of non-negligible non-Gaussianity {\it during the ekpyrotic phase} that dominates the non-Gaussianity generated during the conversion of entropic fluctuations to curvature fluctuations after the ekpyrotic phase \cite{Koyama:2007if,Buchbinder:2007at,Lehners:2007wc,Lehners:2008my}. Furthermore, the steepness of the potential and the instability involve additional tuning of parameters and initial conditions such that, from a theoretical point of view, it would be desirable to find an alternative approach that avoids them. In this paper, we explore a new type of entropic mechanism in which there are two scalar fields, as before, but only one has a steep negative potential, $V(\phi)$. This first field, $\phi$, dominates the energy density and is the source of the ekpyrotic equation of state. The second field, $\chi$, has a negligible potential, perhaps precisely zero potential, but its kinetic energy density is multiplied by a function of the first field, $\Omega^2(\phi)$, with a non-linear sigma-model type interaction. This model shows certain similarities with conformal cosmology \cite{Rubakov:2009np} and pseudo-conformal cosmology \cite{Hinterbichler:2011qk}. A specific example of our model was introduced in \cite{Li:2013hga} and \cite{Qiu:2013eoa} where both the potential and the non-trivial kinetic coupling are proportional to $e^{-\lambda\phi}$, where $\lambda$ is a positive constant. This model, which is characterized by a constant equation of state $\epsilon$, admits stable scaling solutions that generate (nearly) scale-invariant spectra and, as shown by \cite{Fertig:2013kwa}, the bispectrum of this model vanishes such that no non-Gaussianity is produced during the ekpyrotic phase. As such, these models fit well within the Planck2013 bounds on non-Gaussianity; hence it is worthwhile studying how general these results are. Here, we show that these results can be extended to an entire class of ekpyrotic models: we show that scale-invariant entropic perturbations can be produced continuously as modes leave the horizon for {\it any time-dependent} ekpyrotic background equation of state. This has the additional advantage of reducing fine-tuning constraints. The corresponding background solutions are stable and the bispectrum of these perturbations vanishes, such that no non-Gaussianity is produced during the ekpyrotic phase. Hence, the only contribution to non-Gaussianity comes from the non-linearity of the conversion process during which entropic perturbations are turned into adiabatic ones. The paper is organized as follows. In Sec.\,2 we introduce a generic action involving two fields, derive the background equations of motions and briefly discuss their properties. In Sec.\,3 we derive the equations of motion at first order in perturbation theory and show that for each background potential, $V(\phi)$, we can define a non-trivial field-space metric such that the spectrum of entropy perturbations, produced by the $\chi$-field, is scale-invariant. We illustrate our finding on a simple class of ekpyrotic models with equation-of-state parameter $\epsilon = \bar{\epsilon}(-\tau)^p$, where $p>0$. In Sec.\,4 we compute the bispectrum of the perturbations and we show that, for models with constant spectral tilt, no non-Gaussianity is generated during the ekpyrotic phase. We conclude in Sec.\,5 by summarizing our results and discussing directions for future research. | In this paper, we explored a new class of two-field ekpyrotic models with a massive ekpyrotic field governing the background evolution and a second field with no or negligible mass and non-canonical kinetic term. The crucial ingredient of our model is the non-trivial coupling of the background field to the kinetic term of the second, massless field, which plays the role of the entropy field governing the perturbations. Remarkably, we have found that for each background equation of state there exists a non-trivial kinetic coupling such that our model admits scale-invariant solutions (or, more generally, constant spectral index solutions) at first order in perturbation theory. At second order, we have found that the bispectrum of these perturbations vanishes, such that no non-Gaussianity is produced during the ekpyrotic phase. Hence, the only contribution to non-Gaussianity comes from the non-linearity of the conversion process during which entropic perturbations are turned into adiabatic ones. This process is model-dependent, but for an efficient conversion mechanism the final bispectrum remains small, with $f_{\textsc{nl}}^{\text{local}} \sim 5$, which is in accord with current cosmic microwave background measurements \cite{Ade:2013ydc}. This analysis leaves many avenues for future work. A natural extension of our analysis is the calculation of the 4-point function and predictions for the trispectrum (thus extending the analysis of \cite{Lehners:2009ja} to non-trivial field space metrics), in particular since forthcoming data releases from the Planck satellite and large-scale structure experiments will be able to constrain the primordial trispectrum increasingly tightly. Throughout our analysis, we worked with a minimal extension of the standard ekpyrotic theory, studying a two-field Lagrangian. It might be worthwhile to see if a multi-field generalization adds to our model in improving cyclic theories. Similarly, it would be interesting to explore the implications of including a non-negligible mass for the entropy field. | 14 | 4 | 1404.1265 |
1404 | 1404.4992_arXiv.txt | Planetary rotation rate is a key parameter in determining atmospheric circulation and hence the spatial pattern of clouds. Since clouds can exert a dominant control on planetary radiation balance, rotation rate could be critical for determining mean planetary climate. Here we investigate this idea using a three-dimensional general circulation model with a sophisticated cloud scheme. We find that slowly rotating planets (like Venus) can maintain an Earth-like climate at nearly twice the stellar flux as rapidly rotating planets (like Earth). This suggests that many exoplanets previously believed to be too hot may actually be habitable, depending on their rotation rate. The explanation for this behavior is that slowly rotating planets have a weak Coriolis force and long daytime illumination, which promotes strong convergence and convection in the substellar region. This produces a large area of optically thick clouds, which greatly increases the planetary albedo. In contrast, on rapidly rotating planets a much narrower belt of clouds form in the deep tropics, leading to a relatively low albedo. A particularly striking example of the importance of rotation rate suggested by our simulations is that a planet with modern Earth's atmosphere, in Venus' orbit, and with modern Venus' (slow) rotation rate would be habitable. This would imply that if Venus went through a runaway greenhouse, it had a higher rotation rate at that time. | It is traditional to define the habitable zone based on whether liquid water can be maintained on a planet's surface, which is primarily controlled by the planet's surface temperature \citep{Kastingetal1993, Kastingetal2014}. Accurate estimates of the stellar flux boundaries of the habitable zone are critical for estimating parameters of astrophysical interest such as the frequency of Earth-like planets \citep[e.g.,][]{Kopparapu2013}. The inner edge of the habitable zone is set by the runaway greenhouse effect, a positive feedback through which an entire ocean can be evaporated into the atmosphere \citep{Nakajimaetal1992}. Our ability to constrain the stellar flux corresponding to the inner edge of the habitable zone has been severely hampered by the inability of 1D radiative-convective models to predict cloud behavior \citep{Selsisetal2007}. Clouds are critical to planetary energy balance. Cloud reflection of solar radiation is responsible for most of the planetary albedo on modern Earth \citep{DonohoeandBattisti2011}, and clouds also significantly increase Earth's greenhouse effect by absorbing terrestrial infrared emission \citep{Harrisonetal1990}. Cloud coverage and location are primarily controlled by large-scale atmospheric circulation, which is determined by a variety of factors including stellar flux, orbital parameters, and rotation rate. As the stellar flux increases, cloud coverage and thickness may increase, potentially leading to a higher albedo and a negative feedback, or decrease, potentially leading to a lower albedo and a positive feedback\footnote{Here we assume that changes in cloud reflection of stellar radiation (cooling) dominate over changes in cloud absorption of planetary infrared radiation (warming), which is the case in the simulations we present below.}. Orbital parameters such as obliquity and eccentricity can both drive large-amplitude seasonal cycles of the atmospheric circulation and surface temperature \citep[e.g.,][]{Ferreiraetal2014}, but they tend to minimally affect the annual-mean climate \citep{WilliamsandPollard2002, WilliamsandPollard2003}. \begin{figure*}[] \begin{center} \vspace{-22mm} \begin{center} \includegraphics[angle=0, width=34pc]{./df1_TS_ALB_4panels_AquaNoIceNoDZA.eps} \end{center} \vspace{-8mm} \caption{This figure demonstrates the dependence of planetary climate on rotation period ($P_{rot}$) for planets orbiting a Sun-like star. (a) and (b): Global-mean surface temperature (TS) and planetary albedo as a function of $P_{rot}$ for a given stellar flux ($S_0$). Black line: $S_0$\,=\,1365\,W\,m$^{-2}$ and the surface heat capacity (D) is equivalent to 50~m of water; blue line: $S_0$\,=\,1365\,W\,m$^{-2}$ and D\,=\,1~m; red line: clouds are switched off, $S_0$\,=\,1150\,W\,m$^{-2}$, and D\,=\,50~m. For $P_{rot}$\,=\,365~days, the planet is in a synchronously rotating state. (c) and (d): Global-mean surface temperature (TS) and planetary albedo as a function of stellar flux ($S_0$) for a given $P_{rot}$ with D\,=\,50~m. The vertical dashed lines denote the stellar flux of early and modern Venus. The upper horizontal axis in (c--d) is the corresponding semi-major axis between a Sun-like star and the planet in AU. In all these simulations the orbital period is 365 days and there is no sea ice.} \label{fig1} \end{center} \end{figure*} Planetary rotation rate determines the strength of the Coriolis force (the apparent force felt due to the rotation of the planet) and the length of day (and night). The Coriolis force is a key parameter in determining the atmospheric circulation \citep[e.g.,][]{Schneider2006, Showmanetal2013}. If the Coriolis force is strong, thermally direct latitudinal circulations (Hadley cells) are constrained to low latitudes, and the atmosphere organizes into banded, roughly longitudinally symmetric regions. If the Coriolis force is weak, horizontal temperature gradients become small throughout the atmosphere and the Hadley cells can extend globally. The length of day, combined with surface and atmospheric thermal inertia, helps determine the surface temperature distribution, which drives atmospheric circulation \citep{Pierrehumbert2010}. For a short day (or large thermal inertia), the surface temperature difference between day and night is small. If the day is long enough that the dayside is much warmer than the nightside, atmospheric circulation is characterized by ascent on the warm dayside and descent on the cold nightside. Recently a number of calculations have been done with 3D general circulation models (GCMs) to assess the effects of atmospheric circulation, subsaturation, and clouds on the inner edge of the habitable zone \citep{Yangetal2013,Leconteetal2013b,WolfandToon2014}. In the extreme case of tidally locked synchronously rotating planets orbiting M-stars, strong atmospheric ascent on the dayside leads to thick dayside cloud coverage and a very high planetary albedo \citep{Yangetal2013}. This can allow a planet to remain habitable at twice the stellar flux 1D model calculations would suggest. In contrast, for rapidly rotating planets with banded atmospheric circulations, cloud behavior remains roughly similar to modern Earth's (which the albedo of 1D models are tuned to) so that the inner edge of the habitable zone in 3D models is similar to that in 1D models \citep{Leconteetal2013b,WolfandToon2014}. Another interesting difference is that the cloud feedback near the runaway greenhouse threshold appears to be negative for tidally locked planets \citep{Yangetal2013} and positive for rapidly rotating planets \citep{Leconteetal2013b,WolfandToon2014}\footnote{\citet{WolfandToon2014} find a cloud feedback that starts negative, then becomes positive near the runaway greenhouse.}. \begin{figure*}[] \vspace{-10mm} \begin{center} \includegraphics[angle=0, width=34pc]{./df5_Clouds_TS_Winds_RELHUM_1vs128_Aqua.eps} \end{center} \vspace{-22mm} \caption{Differences in clouds and atmospheric circulation between rapidly (left) and slowly (right) rotating planets with an orbital period of 365~days. The rotation period is 1~day for the rapidly rotating planet and the stellar flux is 1365~or~1600~W\,m$^{-2}$. The rotation period is 128~days for the slowly rotating planet and the stellar flux is 1365~or~3000~W\,m$^{-2}$. (a--d): Total cloud coverage (\%, shaded) and surface temperature (K, black contours with an interval of 5~K). (e--h): Relative humidity at 450 mbar (\%, shaded) and near-surface winds (m\,s$^{-1}$, vectors). The black dot in the right panels is the transient substellar point, which moves westward around the planet with a period of 197~days. All variables are averaged over 30 days.} \label{fig2} \end{figure*} The purpose of this study is to investigate the effects of a range of planetary rotation rates, between tidally locked and rapidly rotating, and stellar types on cloud behavior and the inner edge of the habitable zone. To do this we use a 3D GCM with a sophisticated cloud scheme that reproduces cloud behavior well in the large range of climates observed on modern Earth. Although we do not push the model significantly outside of this range, it is important to note that cloud modeling is difficult, and other models may yield quantitatively different results. Nevertheless, we focus on results due to robust physical processes that should be qualitatively similar in any 3D model. Our main conclusion is that for all stellar types slowly rotating planets (orbital period $\approx$100 days or more) behave similarly to tidally locked planets and have a high planetary albedo near the inner edge of the habitable zone. The width of the habitable zone is therefore strongly dependent on planetary rotation rate. | This work demonstrates that the inner edge of the habitable zone for slowly rotating planets could be at twice the stellar flux as for rapidly rotating planets. Numerical simulations suggest that the rotation periods of planets at formation could vary between 10 hours and 400 days \citep{MiguelandBrunini2010}, and tidal interactions can further slow planetary rotation \citep{LagoandCazenave1979}. It is therefore probable that a large number of planets rotate slowly enough to have a greatly expanded habitable zone. Additionally, our simulations suggesting that an Earth-like planet with Venus' present orbit and rotation rate would be habitable demonstrate that empirical limits on the habitable zone based on solar system planets \citep[e.g.,][]{Kastingetal2014} may not be as strong constraints as previously believed, depending on factors such as rotation rate and planetary history. Finally, we note that although we expect our results to be qualitatively robust, the details will differ with other models that have different cloud and radiation schemes. We can hope to resolve this issue by comparing GCMs, applying cloud resolving models to the problem \citep[e.g.,][]{Abbot2014}, and eventually observing planets using methods such as those suggested by \citet{Yangetal2013}. | 14 | 4 | 1404.4992 |
1404 | 1404.2446_arXiv.txt | { In the framework of the Virtual Observatory (VO), the German Astrophysical Virtual Observatory (GAVO) developed the registered service TheoSSA (Theoretical Stellar Spectra Access). It provides easy access to stellar spectral energy distributions (SEDs) and is intended to ingest SEDs calculated by any model-atmosphere code, generally for all effective temperature, surface gravities, and elemental compositions. We will establish a database of SEDs of flux standards that are easily accessible via TheoSSA's web interface. } {The OB-type subdwarf \fg is a standard star for flux calibration. State-of-the-art non-local thermodynamic equilibrium (NLTE) stellar-atmosphere models that consider opacities of species up to trans-iron elements will be used to provide a reliable synthetic spectrum to compare with observations. } { In case of \fg, we demonstrate that the model reproduces not only its overall continuum shape from the far-ultraviolet (FUV) to the optical wavelength range but also the numerous metal lines exhibited in its FUV spectrum. } { We present a state-of-the-art spectral analysis of \fg. We determined \Teffw{47\,250 \pm 2000}, \loggw{6.00 \pm 0.20}, and the abundances of He, N, P, S, Ti, V, Cr, Mn, Fe, Co, Ni, Zn, and Ge. Ti, V, Mn, Co, Zn, and Ge were identified for the first time in this star. Upper abundance limits were derived for C, O, Si, Ca, and Sc. } { The TheoSSA database of theoretical SEDs of stellar flux standards guarantees that the flux calibration of astronomical data and cross-calibration between different instruments can be based on models and SEDs calculated with state-of-the-art model-atmosphere codes. } | \label{sect:intro} \fg is a bright \citep[$m_\mathrm{V} = 11.845 \pm 0.010$,][]{kharchenkoroeser2009}, subluminous OB-star (type sdOB, \citealt{heberetal1984b}; type sdO D,\citealt{vennesetal2011}). It is widely used as a spectrophotometric standard star \citep[e.g\@.][]{oke1990,turnsheketal1990,bohlinetal1990}. Since \fg will be used as a reference star for the flux calibration of X-SHOOTER\footnote{\url{http://www.eso.org/sci/facilities/paranal/instruments/xshooter.html}} \citep{vernetetal2011} observations from 3000\,\AA\ to 25\,000\,\AA\ \citep{moehleretal2014}, we decided to reanalyze its spectrum with state-of-the-art model-atmosphere techniques. An early spectral analysis with approximate LTE\footnote{local thermodynamic equilibrium}, line-blanketed hydrogen model atmospheres yielded an effective temperature \Teffw{39\,000} and a surface gravity $\log (g\,/\,\mathrm{cm/s^2}) = 6.5$ \citep{newell1973}. \citet{kudritzki1976} showed that both, the consideration of deviations from LTE (NLTE\footnote{non-local thermodynamic equilibrium}) as well as of opacities of elements heavier than H, have a significant influence on the determination of \Teff and \logg in an analysis of optical spectra (Table\,\ref{tab:kud76}). \citet{heberetal1984b} extended the analysis of \fg to the ultraviolet (UV) wavelength range (IUE\footnote{International Ultraviolet Explorer} observations, $1150\,\mathrm{\AA}\,\la\,\lambda\,\la\,2000\,\mathrm{\AA}$) in addition to high-resolution optical spectra ($4000\,\mathrm{\AA}\,\la\,\lambda\,\la\,5100\,\mathrm{\AA}$) and derived \Teffw{40\,000^{+5000}_{-3000}}, \loggw{5.0 \pm 0.3}, and He/H = 0.03$^{+0.03}_{-0.02}$ (by number) using H+He (with subsequent C+N+Si line-formation calculations) NLTE models. \onltab{ \begin{table}\centering \caption{\Teff and \logg of \fg determined by \citet{kudritzki1976}. He/H gives his models' abundance ratio by number.} \label{tab:kud76} \begin{tabular}{rr@{.}lcrr@{.}lr@{.}l} \hline \multicolumn{3}{c}{LTE} & & \multicolumn{3}{c}{NLTE} & \multicolumn{2}{c}{} \\ \cline{1-3} \cline{5-7} \multicolumn{8}{c}{~} \vspace{-5.5mm} \\ \multicolumn{7}{c}{~} & \multicolumn{2}{c}{He/H} \vspace{-2.5mm} \\ \noalign{\smallskip} \multicolumn{1}{c}{$T_\mathrm{eff}$\,/\,K} & \multicolumn{2}{c}{$\log\,(g\,/\,\mathrm{cm/s^2)}$} & & \multicolumn{1}{c}{$T_\mathrm{eff}$\,/\,K} & \multicolumn{2}{c}{$\log\,(g\,/\,\mathrm{cm/s^2)}$} & \multicolumn{2}{c}{} \\ \noalign{\smallskip} \hline \noalign{\smallskip}42\,600 & \hbox{}\hspace{2mm}6&3 && 44\,600 & \hbox{}\hspace{2mm}5&9 & \hbox{}\hspace{2mm}0&1 \\ 42\,400 & 6&5 && 42\,700 & 6&4 & 1&0 \\ \hline \end{tabular} \end{table} } With the FUSE\footnote{Far Ultraviolet Spectroscopic Explorer} mission, the interstellar deuterium and oxygen column densities toward \fg were measured. \citet{friedmanetal2002} used optical spectra and estimated the atmospheric parameters by comparison with a grid of synthetic NLTE model-atmosphere spectra \citep[using TLUSTY to compute the stellar atmosphere model and SYNSPEC to generate the SED,][just ``TLUSTY'' here after]{hubenylanz1995}, that considered H and He. They achieved \Teffw{42\,300 \pm 1000}, \loggw{5.95 \pm 0.15}, and He/H = $0.011 \pm 0.005$. With the higher \logg \citep[in agreement with][]{kudritzki1976}, their spectroscopic distance of $d =288 \pm 43\,\mathrm{pc}$ agreed with the Hipparcos\footnote{\url{http://www.rssd.esa.int/index.php?project=HIPPARCOS}} parallax distance of $d =179^{+265}_{-67}\,\mathrm{pc}$. In the following, we describe our analysis in detail. In Sect\@.\,\ref{sect:obs}, we give some remarks on the observations. Then, we introduce our models and the considered atomic data (Sect\@.\,\ref{sect:models}) and start with a preliminary analysis (Sect\@.\,\ref{sect:prelim}) of the optical spectrum based on H+He models followed by a highly sophisticated analysis with metal-line blanketed models (Sect\@.\,\ref{sect:analysis}). We summarize our results and conclude in Sect\@.\,\ref{sect:results}. | \label{sect:results} We performed a comprehensive spectral analysis of \fg, based on observations from the FUV to the optical wavelength range. We determined \Teffw{47\,250 \pm 2000} and \loggw{6.00 \pm 0.20}. The ionization equilibria of \ion{He}{i} / \ion{He}{ii}, \ion{N}{iii} / \ion{N}{iv} / \ion{N}{v}, \ion{P}{iv} / \ion{P}{v}, \ion{S}{iv} / \ion{S}{v} / \ion{S}{vi}, \ion{Ti}{iv} / \ion{Ti}{v}, \ion{V}{iv} / \ion{V}{v}, \ion{Cr}{iv} / \ion{Cr}{v} / \ion{Cr}{vi}, \ion{Mn}{v} / \ion{Mn}{vi}, \ion{Fe}{v} / \ion{Fe}{vi}, \ion{Co}{v} / \ion{Co}{vi}, and \ion{Ni}{v} / \ion{Ni}{vi} are well reproduced with these values. The photospheric abundances were determined based on the FUSE and optical observations (Table\,\ref{tab:abund}). Figure\,\ref{fig:abund} shows a comparison of the photospheric abundances patterns of three hot O(B)-type subdwarfs. While the intermediate-mass metals are solar or subsolar in all these stars, the iron-group elements but Fe have strongly super-solar values. An exception is Fe in \object{AA\,Dor} and \object{EC\,11481$-$2303}, that appears to be solar. Neither this Fe peculiarity nor the extremely low C and Si abundances in \fg can be explained. \begin{figure} \resizebox{\hsize}{!}{\includegraphics{23711_f17.eps}} \caption{Comparison of the determined photospheric abundances (arrows indicate upper limits) of the three OB-type subdwarfs \object{AA\,Dor} \citep{klepprauch2011}, \object{EC\,11481$-$2303} \citet{rauchetal2010b}, and \fg. Their \Teff and \logg are shown in the legend.} \label{fig:abund} \end{figure} The position of \fg in the \Teff $-$ \logg plane shows that it is located directly on the He main sequence (Fig.\,\ref{fig:tefflogg}). \fg belongs, like \object{AA\,Dor} or \object{EC\,11481$-$2303}, to the hottest post-EHB\footnote{extended horizontal branch} stars. From a comparison to post-EHB tracks \citep{dormanetal1993}, we can extrapolate a stellar mass of $M = 0.469 \pm 0.001$\,\Msol. With $R = \sqrt{GM/g}$ ($G$ is the gravitational constant), we calculated the stellar radius of $R = 0.114^{+0.030}_{-0.024}\,R_\odot$. \begin{figure} \resizebox{\hsize}{!}{\includegraphics{23711_f18.eps}} \caption{Location of \fg in the \Teff $-$ \logg plane compared to sdBs and sdOBs from \citet{edelmann2003}, \citet[\object{EC\,11481$-$2303}]{rauchetal2010b}, and \citet[\object{AA\,Dor}]{klepprauch2011}. Post-EHB tracks from \citet[][$Y_\mathrm{HB} = 0.288$, labeled with the respective stellar masses in $M_\odot$]{dormanetal1993} are also shown. Their start and kink points are used to illustrate the location of the zero-age and terminal age EHB for this He composition. The He main sequence is taken from \citet{paczynski1971}. } \label{fig:tefflogg} \end{figure} We determined the distance of \fg following the flux calibration of \citet{heberetal1984a} for $\lambda_\mathrm{eff} = 5454\,\mathrm{\AA}$, \begin{equation} d_\mathrm{spec} = 7.11 \times 10^4 \cdot \sqrt{H_\nu\cdot M \cdot 10^{0.4\, m_{\mathrm{V}_0}-\log g}}\,\mathrm{pc}\,, \label{eq:distance} \end{equation} \noindent with $m_\mathrm{V_o} = m_\mathrm{V} - 2.175 c$, $c = 1.475 E_\mathrm{B-V}$, and the Eddington flux $H_\nu = 7.24 \pm 0.37 \times 10^{-4}\, \mathrm{erg/cm^{2}/s/Hz}$ at $\lambda_\mathrm{eff}$ of our final model atmosphere. We used $\ebv =0.027 \pm 0.007$ (Sect.\,\ref{sect:prelim}), $M = 0.469 \pm +0.001$\,\Msol, and $m_\mathrm{V} = 11.847 \pm 0.010$ \citep{kharchenkoroeser2009} and derived a distance of $d_\mathrm{spec}=297^{+62}_{-77}$\,pc and a height below the Galactic plane of $z=255^{+53}_{-66}$\,pc. This distance is about a factor of three larger than the new Hipparcos parallax-measurement reduction \citep[\object{HIP115195}, $\pi = 9.76 \pm 3.44$\,mas]{vanleeuwen2007} of $d_\mathrm{parallax}=102.46^{+55.78}_{-26.69}$\,pc. Interestingly, the older Hipparcos measurement published by \citet[$\pi = 5.59 \pm 3.34$\,mas, $d_\mathrm{parallax}=178.89^{+265.55}_{-66.91}$\,pc]{perrymanetal1997} deviates from this new value by a factor of almost two and would be in agreement with our spectroscopic distance within error limits. The discrepancy between spectroscopic and parallax distances is a significant problem. \logg cannot be higher by about 0.5 to achieve a distance agreement, because the spectral lines in the models appear too broad and too shallow. This apparently is not a problem of our TMAP code, because \citet[$d_\mathrm{spec}=288 \pm 43$\,pc]{friedmanetal2002} used the TLUSTY code and encountered the same problem. Similar discrepancies are reported by \citet[][for \object{LSV+46$\degr$21} with TMAP: $d_\mathrm{spec} = 224^{+46}_{-58}\,\mathrm{pc}$ vs\@. $d_\mathrm{parallax} = 129^{+6}_{-5}\,\mathrm{pc}$]{rauchetal2007} and by \citet[][for \object{BD+28$\degr$4211} with TLUSTY, \Teffw{82\,000}, \loggw{6.2}, and an assumed $M = 0.5$\,\Msol: $d_\mathrm{spec} = 157\,\mathrm{pc}$ (no error estimate given) vs\@. $d_\mathrm{parallax} = 92^{+13}_{-11}\,\mathrm{pc}$]{latouretal2013}. \citet{latouretal2013} mentioned that a relatively high \logg value and/or a low mass may be the solution and since they regard the HIPPARCOS measurement as fully reliable and their TLUSTY results reasonably reliable, the mass of \object{BD+28$\degr$4211} must be much less than the canonical post-EHB mass of about 0.5\,\Msol. For their $d_\mathrm{parallax}/d_\mathrm{spec} = 0.59$, the mass has to be about 0.17\,\Msol. In case of \fg, with $d_\mathrm{parallax}/d_\mathrm{spec} = 0.31$, the mass has to be about 0.10\,\Msol. In both cases, the mass can be higher, if \logg is higher. Thus, since \logg is also the main error source in the spectroscopic distance (Eq.\,\ref{eq:distance}), one might speculate about the applied broadening theory for lines that are used to determine \logg. For the relevant \ion{H}{i} and \ion{He}{ii} lines (linear Stark effect), TMAP as well as TLUSTY use the same data of \citet{tremblaybergeron2009} and \citet{schoeningbutler1989}, respectively. However, all the narrow metal lines (e.g\@. of the iron-group element) in the UV, that are broadened by the quadratic Stark effect, cannot be reproduced at a much higher \logg. To summarize, the distance discrepancy is as yet unexplained. The analysis of the FUV spectrum has shown that the lack of reliably measured wavelengths of lines of the iron-group elements (Ca - Ni) and of elements heavier than Ni hampers the line-identification. Efforts in this field in the near future are highly desirable. The established database of spectrophotometric standard stars in TheoSSA was extended by the OB-type subdwarf \fg. The successfully launched GAIA\footnote{\url{http://www.esa.int/Our_Activities/Space_Science/Gaia }} mission will provide accurate parallax measurements for spectrophotometric standard stars. This will strengthen the importance of a VO-compliant database like TheoSSA that provides easy access to the best synthetic spectra calculated for these stars. | 14 | 4 | 1404.2446 |
1404 | 1404.5861_arXiv.txt | {} {An effort has been made to simulate the expected Gaia Catalogue, including the effect of observational errors. We statistically analyse this simulated Gaia data to better understand what can be obtained from the Gaia astrometric mission. This catalogue is used to investigate the potential yield in astrometric, photometric, and spectroscopic information and the extent and effect of observational errors on the true Gaia Catalogue. This article is a follow-up to Robin et al. (A\&A 543, A100, 2012), where the expected Gaia Catalogue content was reviewed but without the simulation of observational errors.} {We analysed the Gaia Object Generator (GOG) catalogue using the Gaia Analysis Tool (GAT), thereby producing a number of statistics about the catalogue.} {A simulated catalogue of one billion objects is presented, with detailed information on the 523 million individual single stars it contains. Detailed information is provided for the expected errors in parallax, position, proper motion, radial velocity, and the photometry in the four Gaia bands and for the physical parameter determination including temperature, metallicity, and line of sight extinction. } {} | Gaia, a cornerstone ESA mission, launched in December 2013, will produce the fullest 3D galactic census to date, and it is expected to yield a huge advancement in our understanding of the composition, structure, and evolution of the Galaxy \citep{Gaia}. Through Gaia's photometric instruments, object detection up to $G=20$ mag will be possible (see \cite{Gmag} for a definition of $G$ magnitude), including measurements of positions, proper motions, and parallaxes up to micro arcsecond accuracy. The on-board radial velocity spectrometer will provide radial velocity measurements for stars down to a limit of $G_{RVS}=17$ mag. With low-resolution spectra providing information on effective temperature, line of sight extinction, surface gravity, and chemical composition, Gaia will yield a detailed catalogue that contains roughly 1\% of the entire galactic stellar population. Gaia will represent a huge advance on its predecessor, Hipparcos \citep{Hipparcos}, both in terms of massive increases in precision and in the numbers of objects observed. Thanks to accurate observations of large numbers of stars of all kinds, including rare objects, large numbers of Cepheids and other variable stars, and direct parallax measurements for stars in all galactic populations (thin disk, thick disk, halo, and bulge), Gaia data is expected to have a strong impact on luminosity calibration and improvement of the distance scale. This, along with applications to studies of galactic dynamics and evolution and of fields ranging from exoplanets to general relativity and cosmology, Gaia's impact is expected to be significant and far reaching. During its five years of data collection, Gaia is expected to transmit some 150 terabytes of raw data to Earth, leading to production of a catalogue of $10^9$ individual objects. After on-ground processing, the full database is expected to be in the range of one to two petabytes of data. Preparation for acquiring this huge amount of data is essential. Work has begun to model the expected output of Gaia in order to predict the content of the Gaia Catalogue, to facilitate the production of tools required to effectively validate the real data before publication, and to analyse the real data set at the end of the mission. To this end, the Gaia Data Processing and Analysis Consortium (DPAC) has been preparing a set of simulators, including a simulator called the Gaia Object Generator (GOG), which simulates the end-of-mission catalogue, including observational errors. Here a full description of GOG is provided, including the models assumed for the performance of the Gaia satellite and an overview of its simulated end-of-mission catalogue. A selection of statistics from this catalogue is provided to give an idea of the performance and output of Gaia. In Sect. \ref{sec:gaia}, a brief description of the Gaia instrument and an overview of the current simulation effort is given, followed by definitions of the error models assumed for the performance of the Gaia satellite in Sec. \ref{sec:simulations}. In Sec. \ref{sec:methods}, the methods used for searching the simulated catalogue and generating statistics are described. In Sec. \ref{sec:results}, we present the results of the full sky simulation, broken up into sections for each parameter in the catalogue and specific object types of interest. Finally, in Sec. \ref{sec:conclusions} we provide a summery and conclusions. | \label{sec:conclusions} The Gaia Object Generator provides the most complete picture to date of what can be expected from the Gaia astrometric mission. Its simulated catalogue provides useful insight into how various types of objects will be observed and how each of their observables will appear after including observational errors and instrument effects. The simulated catalogue includes directly observed quantities, such as sky position and parallax, as well as derived quantities, such as interstellar extinction and metallicity. Additionally, the full sky simulation described here is useful for gaining an idea of the size and format of the eventual Gaia Catalogue, for preparing tools and hardware for hosting and distribution of the data, and for becomeing familiar with working with such a large and rich dataset. In addition to the stellar simulation described in this paper, there are plans to generate other simulated catalogues of interest, such as open clusters, Magallanic Clouds, supernovae, and other types of extragalactic objects, so that a more complete version of the simulated Gaia Catalogue can be compiled. Here we have focussed on the simulated catalogue from the inbuilt Gaia Universe Model, based on the Besan\c{c}on Galaxy model. However, GOG can alternatively be supplied with an input catalogue generated by the user. This way, simulated data from any other model can be processed with GOG to obtain simulated Gaia observations of specific interest to the individual user. The input can be either synthetic data on a specific star or catalogue, or an entire simulated survey such as those generated using Galaxia \citep{Galaxia}, provided a minimum of input information is supplied (e.g. position, distance, apparent magnitude, and colour). With GOG, the capabilities of the instrument can be explored, and it is possible to gain insight into the expected performance for specific types of objects. While only a subset of the available statistics have been reproduced here, it is possible to obtain the full set of available statistics at request. We are working to make the full simulated catalogue publicly available, so that interested individuals can begin working with data similar to the forthcoming Gaia Catalogue. | 14 | 4 | 1404.5861 |
1404 | 1404.5266_arXiv.txt | {The structure of inner region of protoplanetary disks around young pre-main-sequence stars is still poorly understood. This part of the disk is shaped by various forces that influence dust and gas dynamics, and by dust sublimation, which creates abrupt drops in the dust density. This region also emits strong near-infrared excess that cannot be explained by classical accretion disk models, which suggests the existence of some unusual dust distribution or disk shape. The most prevalent explanation to date is the puffed-up inner disk rim model, where the disk exhibits an optically thin cavity around the star up to the distance of dust sublimation. The critical parameter in this model is the inner disk rim height $z_{\rm max}$ relative to the rim distance from the star $R_{\rm in}$. Observations often require $z_{\rm max}/R_{\rm in}\gtrsim0.2$ to reproduce the near-infrared excess in the spectra. We compile a comprehensive list of processes that can shape the inner disk rim and combine them into a self-consistent model. Two of them, radiation pressure force and the gas velocity profile, have never been applied in this context before. The aim was to find the most plausible theoretical values of $z_{\rm max}/R_{\rm in}$. The results show that this value is $\lesssim$0.13 for Herbig Ae stars, $\lesssim$0.11 for T Tau stars, and $\lesssim$0.10 for young brown dwarfs. This is lower than the observational requirements for Herbig Ae stars. We argue that the same problem exists in T Tau stars as well. We conclude that the puffed-up inner rim model cannot be the sole explanation for the near-infrared excess in young pre-main-sequence stars. } | Protoplanetary disks surrounding young pre-main-sequence stars have been a subject of intense theoretical and observational scrutiny in recent decades \citep{Williams}. Since these disks are expected to host planet formation, particular focus has been on dust properties and dust distribution within the disks. While this led to an advanced understanding of the dust evolution beyond $\sim$1 AU from the star, the inner disk region within $\lesssim$1 AU remained a controversial topic \citep{Millan-Gabet,DullemondARAA}. The main obstacle is the inability of current telescopes to spatially resolve this region, except by the near-infrared interferometry, which in itself depends on modeling assumptions and currently provides valuable but still limiting information on the disk structure \citep{Wolf}. The inner disk rim is observable at near-infrared wavelengths because dust emission in this part of the disk comes from dust temperatures above 1,000K. Even the most resilient dust grains sublimate in these conditions, but the exact sublimation distance from the star depends on the grain size and chemistry and on the local gas pressure. Removal of dust grains creates large steps in the opacity gradient, which complicates modeling the dusty disk structure \citep{KMD09,Vinkovic12}. Modeling efforts have therefore been devoted to approximate methods, motivated by the peculiar near-infrared bump in the spectral energy distribution of many pre-main-sequence stars. This bump was first noticed in Herbig Ae/Be stars by \citet{Hillenbrand}, while \citet{Chiang} showed that ordinary accretion disk models do not produce enough flux to explain the bump. \citet{DDN} introduced the most popular explanation for this phenomenon to date, based on the assumption that the disk gas is optically thin close to the star in the zone where dust grains cannot survive. This allows stellar radiation to heat the whole vertical profile of the inner dusty disk rim, including the disk interior, which is typically much colder than the disk surface. This causes the disk to expand vertically, or "puff-up", which also increases the emitting area of the hot dusty disk rim. The near-infrared bump is then merely a measure of how much the disk expands vertically. That assumption of an optically thin hole was confirmed by near-infrared interferometry \citep{Millan-Gabet01,Millan-Gabet}, which boosted the popularity of the puffed-up disk model. The model was additionally improved \citep[e.g.][]{DD04a,DD04b,Isella05,Tannirkulam07,Thi} and is still the most prevalent description of the inner disk structure \citep{DullemondARAA}. However, over the years, two main concerns about the plausibility of the model have been raised. One concern is that the near-infrared excess predicted by the puffed-up disk model is too low to explain the strongest near-infrared bumps, unless physically unrealistic vertical puffing is invoked, with heights several times above hydrostatic equilibrium \citep{DAlessio,VIJE,Meijer,Schegerer,Acke,Verhoeff10,Verhoeff11,Mulders,McClure2013a}. In addition, as \citet{Mulders} explain in their Appendix C, such strong puffing up casts a shadow over large parts of the outer disk, which would decrease the disk temperature and reduce its mid-infrared emission below the observed levels. Moreover, the claim that an optically thin hole does not contribute to the near-infrared excess and images has been challenged recently by some interferometry observations \citep{Akeson2005,Tannirkulam08,Benisty10,Benisty11} and also theoretically \citep{Vinkovic06,KMD09}. The other concern is that the theoretical approaches to the modeling of the inner disk have been lacking some critical physical elements and therefore cannot correctly describe its structure. Small grains have often been included into the dust opacity of the inner rim, but now we know that the shape of inner disk rim is dictated solely by big grains ($\gtrsim 1\mu m$) because they are more resilient to sublimation than smaller grains. Hence, big grains populate the inner disk surface and shield smaller grains in the disk interior from direct stellar radiation \citep{KMD09,Vinkovic12}. An unexpected consequence of radiative transfer in big grains is that the location of the inner disk rim becomes a nontrivial problem, with the temperature inversion effect producing a local temperature maximum within the dust cloud (see the appendix). Here we focus only on the height and shape of optically thick disk, where we use the simple prescription for the rim radius presented by \citet{Vinkovic06}. Moreover, processes such as vertical gas velocity profile \citep{TakeuchiLin02} or radiation pressure on dust grains \citep{TakeuchiLin03,Vinkovic09} have been lacking from the debate on the puffed-up rim model. In this paper we construct a model that takes into consideration a comprehensive list of physical processes that influence the rim shape. Our goal is to find $z_{\rm max}/R_{\rm in}$ for a wide range of disk parameters and determine constraints on the applicability of the puffed-up inner rim model. | We investigated the inner disk rim height using a self-consistent dust dynamics model based on a long list of physical processes: big dust grains populating the disk surface as the dust grains most resilient to the direct stellar heating, the temperature inversion effect that affects the position of the inner disk radius, accretion luminosity, viscous heating, turbulent diffusion, dust settling due to the gas drag force, vertical hydrostatic equilibrium, radiation pressure force, and the gas velocity profile. The last two processes have never been included in such an analysis before. The objects under consideration were Herbig Ae stars, T Tau stars, and brown dwarfs. A wide range of parameters was explored, resulting in more than 6,000 models of the rim height. The results were compared with the rim height implied by the observations of near-infrared excess from protoplanetary disks. Comparison showed that although we built a comprehensive self-consistent model, the modeled rim heights are insufficient to explain the observed flux of Herbig Ae stars. This confirms some previous studies that reached the same conclusion using models that incorporate a smaller list of physical processes. It also implies that some additional physical processes are influencing dust dynamics in such a way that they increase the overall volume of observed hot optically thin dust. This problem is closely related to the question of the shape (i.e. curvature) of the inner disk rim. We showed that radiation pressure creates a cut-off height above which the dust flows outward. This means that the dusty disk cannot exist above this height unless some other forces counterbalance the radiation pressure force. We argued that recent observations of extreme near-infrared excess in T Tau stars are also indicative of extensions of the puffed-up inner rim model. | 14 | 4 | 1404.5266 |
1404 | 1404.0555.txt | Strong shear flow regions found in astrophysical jets are shown to be important dissipation regions, where the shear flow kinetic energy is converted into electric and magnetic field energy via shear instabilities. The emergence of these self-consistent fields make shear flows significant sites for radiation emission and particle acceleration. We focus on electron-scale instabilities, namely the collisionless, unmagnetized Kelvin-Helmholtz instability (KHI) and a large-scale dc magnetic field generation mechanism on the electron scales. We show that these processes are important candidates to generate magnetic fields in the presence of strong velocity shears, which may naturally originate in energetic matter outburst of active galactic nuclei and gamma-ray bursters. We show that the KHI is robust to density jumps between shearing flows, thus operating in various scenarios with different density contrasts. Multidimensional particle-in-cell (PIC) simulations of the KHI, performed with OSIRIS, reveal the emergence of a strong and large-scale dc magnetic field component, which isÊnot captured by the standard linear fluid theory. This dc component arises from kinetic effects associated with the thermal expansion of electrons of one flow into the other across the shear layer, whilst ions remain unperturbed due to their inertia. The electron expansion forms dc current sheets, which induce a dc magnetic field. Our results indicate that most of the electromagnetic energy developed in the KHI is stored in the dc component, reaching values of equipartition on the order of $10^{-3}$ in the electron time-scale, and persists longer than the proton time-scale. Particle scattering/acceleration in the self generated fields of these shear flow instabilities is also analyzed. | Relativistic jets are found in a wide range of extreme astrophysical scenarios like active galactic nuclei (AGN) and gamma-ray bursts (GRBs) \cite{bridle84,mirabel99}. The energetic outflows of plasma associated with astrophysical jets represent massive sources of free-energy for collisionless plasma instabilities to operate. The onset of plasma instabilities play a central role in dissipating the jet's kinetic energy into electric and magnetic turbulence \cite{gruzinovwaxman99,medvedev99} resulting in particle acceleration to ultra-high energies and nonthermal radiation emission. A deep understanding of these processes and their interplay is challenging, requiring full kinetic simulations to address their highly nonlinear nature. First principle modeling of these processes are, however, computationally intensive due to the wide range of temporal and spatial scales involved. Therefore, full kinetic simulations demand massive computational resources and advanced numerical and visualization techniques. Much attention has been devoted to relativistic shocks, which are thought to be a strong mechanism for particle acceleration. Such shocks arise from the collision and bulk interpenetration of different velocity plasma shells, due to either intermittencies or inhomogeneities of the ejecta. The Weibel \cite{weibel59} and the purely transverse two stream instabilities \cite{silva03} act as the dissipation mechanism in these scenarios, and are critical for shock formation. A vast number of fully kinetic simulations have focused on shock formation settings, where long-lived equipartition magnetic field generation via the Weibel instability has been observed \cite{silva03,fonseca03,frederiksen04,nishikawa05}. A Fermi-like particle acceleration process has also been identified in simulations of long-term evolution of collisionless shocks \cite{spitkovsky08,martins09}. These previous works have only considered shearless flows. However, in addition to bulk plasma collision sites, the transition layers of shear flows have also been probed \cite{gruzinov08} and shown to constitute important dissipation regions \cite{alves12,grismayer13,Liang13}. Increasing evidence has pointed to a general stratified organization of the structure of jets in AGN and GRBs \cite{granot03, rieger04}, where different internal shear layers can occur; rotating inner cores vs. axially moving outer shells, or fast inner cores vs. slower outer shells. Moreover, external shear layers, resulting from the interaction of the jet with the interstellar medium, may also be considered. In these scenarios, collisionless shear instabilities such as the Kelvin-Helmholtz instability (KHI) \cite{dangelo65,gruzinov08,macfadyen09} play a role in the dissipation of the jet kinetic energy into electric and magnetic turbulence \cite{macfadyen09,alves12,Liang13}. In fact, the combined effect of shear flow with collisionless shock formation has not yet been addressed, and may also lead to interesting novel phenomenology since density inhomogeneities generated by shear instabilities can also constitute important scattering sites for particle acceleration. Recent fully kinetic simulations of shear flow settings have probed the self-consistent evolution of the electron-scale KHI, demonstrating that the operation of kinetic effects are responsible for the generation of large-scale, equipartition magnetic fields \cite{alves12,grismayer13}. Nonthermal particle acceleration has also been investigated in hybrid electron-positron-ion shear flows \cite{boettcher,Liang13}, with different pair/ion ratio compositions, showing spectral features similar to those found in GRBs. In laboratory experiments, scenarios where the unmagnetized KHI can be triggered are now being examined both in the collisional \cite{harding09,Hurricane12} and in the collisionless regimes\cite{kuramitsu12} (the latter is explored in this paper). In this work we focus on electron-scale processes triggered by velocity shears, namely the unmagnetized KHI and a dc magnetic field generation mechanism. In Section 2, we develop the linear theory for the cold unmagnetized KHI, and analyze the impact of density contrast between sharp shearing flows. We find the onset of the KHI is robust to density contrasts, allowing for a strong development in various density contrast regimes (inner shears with low density contrasts, and outer shears with high density contrasts). We then extend the analysis to finite shear gradients, where we find that KHI growth rate decreases with increasing shear gradient length. Particle-in-cell (PIC) simulations are performed to verify the theoretical predictions. At late times, PIC simulations reveal the formation of a large-scale, dc magnetic field extending along the entire shear surface between flows, which is not predicted by the linear KHI theory. This dc magnetic field is the dominant feature of the magnetic field structure of the instability at late times. In Section 3, we find that the dc magnetic field results from kinetic effects associated with electron mixing between shearing flows, which is driven by the nonlinear development of the cold KHI. The dc magnetic field generation is discussed and an analytical model is developed that captures the main features of the dc magnetic field evolution and saturation. In Section 4, we analyze the dynamics of the electrons in the self-generated fields. The electrons are scattered in the self consistent electric and magnetic fields generated by the KHI, and are accelerated to high energies. We discuss the particle energy spectra resulting from the development of these shear instabilities, and we investigate the mechanism underlying the acceleration of energetic particles using advanced particle tracking diagnostics. %\input{2_fluid_regime} \newpage | Shear instabilities in plasmas are usually studied within the framework of magnetohydrodynamics where the plasma is considered as a magnetized fluid and where the typical time scale is governed by ion motion. We have shown in this work that electron scale physics leads to a variety of new effects when one considers an initially unmagnetized cold shearing collisionless plasma. The collective electron dynamics can be described in first approximation by using a two-fluid model that allows to take into account electron inertial physics, not captured in MHD models. In this fluid framework, we have presented the derivation of the equations for the linear development of the longitudinal KHI, assuming arbitrary velocity and density profiles. This framework was applied to a special case, where the velocity and density profiles were simple step-functions, allowing analytical solutions to the equations. The model provided new insights into the effect of density-contrasts between shearing flows, namely that the development of the KHI is robust to density jumps, making it ubiquitous in astrophysical settings. We also observed that the unstable modes begin to drift when the density symmetry is broken. In large density-contrast regimes, the KHI dominates over other common astrophysical plasma instabilities such as the Weibel and Two-Stream instabilities. The case of a finite shear profile has also been investigated in detail. A smooth velocity profile (non step-like) induces a phase mixing of the eigenmodes of the system that results into a damping term. The combined effect of the instability due to the shear with the damping term suggests that the maximum growth rate is a decreasing function of the shear gradient length. All of these results obtained in the limit of linearized fluid equations have been accurately verified by 2D PIC simulations. PIC simulation results have also demonstrated the emergence of a large-scale dc magnetic field after the onset of the electron-scale KHI . However this dc field is not captured by the two-fluid KHI theory nor MHD model. We have shown that the emergence of the dc magnetic field is intrinsically associated with electron-ion shear flows. The dc magnetic field is induced through the formation of dc current sheets driven by the expansion of electrons in the shear region due to either a thermal expansion or the development of the cold fluid electron-scale KHI perturbations. We have presented an analytical description of the formation of the dc field in agreement with 1D, 2D, and 3D PIC simulations. The dc magnetic-field saturation on the electron time scale is independent of the type of the expansion and persists beyond ion time scales. Finally, we addressed the particle acceleration physics due to scattering in the self-generated fields of electron-scale instabilities triggered in unmagnetized shear flows. An understanding of how electrons are accelerated is essential if we are to fully interpret observations since it is presumably the radiation from energetic electrons that is most often observed from astrophysical sources. To address this issue, we have tracked the most energetic electrons in our simulations to identify the acceleration mechanism. It was found that the kinetic energy, initially stored in the drift, was mainly redistributed thermally. The bulk of the energy distribution energy of the electrons has a typical temperature comparable with the Lorentz factor of the flow. Nevertheless, the energy distribution also displays a non-thermal tail. The most energetic electrons, that make up the power-law tail of the distribution, are accelerated while surfing close to the speed of light on electric and magnetic field patches, self consistently developed by the electron-scale KHI, which are carried by the flow. This results in an efficient acceleration mechanism where electrons can reach energies on the order of $\gamma_0^4$, where $\gamma_0$ is the initial Lorentz factor of the flow. %\input{Appendix} \ack E. P. Alves and T. Grismayer contributed equally to this work. This work was supported by the European Research Council (ERCÑ2010ÑAdG Grant 267841) and FCT (Portugal) grants SFRH/BD/75558/2010, SFRH/BPD/75462/2010, and PTDC/FIS/111720/2009. We acknowledge PRACE for providing access to resource SuperMUC based in Germany at the Leibniz research center. Simulations were performed at the IST cluster (Lisbon, Portugal) and SuperMUC (Germany). | 14 | 4 | 1404.0555 |
1404 | 1404.0219_arXiv.txt | There is a remarkable correlation between the spin periods of the accreting neutron stars in Be/X-ray binaries (BeXBs) and their orbital periods . Recently \citet{k2011} showed that the distribution of the spin periods contains two distinct subpopulations peaked at $\sim 10$ s and $\sim 200$ s respectively, and suggested that they may be related to two types of supernovae for the formation of the neutron stars, i.e., core-collapse and electron-capture supernovae. Here we propose that the bimodal spin period distribution is likely to be ascribed to different accretion modes of the neutron stars in BeXBs. When the neutron star tends to capture material from the warped, outer part of the Be star disk and experiences giant outbursts, a radiatively-cooling dominated disk is formed around the neutron star, which spins up the neutron star, and is responsible for the short period subpopulation. In BeXBs that are dominated by normal outbursts or persistent, the accretion flow is advection-dominated or quasi-spherical. The spin-up process is accordingly inefficient, leading to longer periods of the neuron stars. The potential relation between the subpopulations and the supernova mechanisms is also discussed. | High-mass X-ray binaries (HMXBs) usually consist of a neutron star (NS) and an optical companion star of mass higher than about $8 M_{\odot}$. According to the spectral characteristics of the optical companions, HMXBs can be further divided into supergiant X-ray binaries (SGXBs) and Be/X-ray binaries (BeXBs) \citep[][for a recent review]{Reig2011}. Most BeXBs are transient systems and present moderately eccentric orbits ($e\gtrsim 0.3$). The NS captures the wind material from its companion, producing X-ray radiation. Meanwhile, the spin of the NS changes with time, and both spin-up and spin-down have been observed when accretion took place \citep{n1989,b97}. \citet{Corbet1984,Corbet1985,Corbet1986} first noticed that different subgroups of HMXBs appear to be located in different regions in the spin perion ($P_{\rm s})$ vs. orbital period ($P_{\rm orb}$) diagram (also called the Corbet diagram). In particular, there exists a positive correlation between $P_{\rm s}$ and $P_{\rm orb}$ for BeXBs, although with a large observed scatter. The relations between $P_{\rm s}$ and $P_{\rm orb}$ in HMXBs are likely to reflect the wind structure and accretion processes in HMXBs \citep[e.g.,][]{swr1986,hr87,wv1989,king1991,lv1996}. It is generally thought that the interaction between the NS magnetic field and the captured material from its binary companion can lead to a so-called equilibrium spin period $P_{\rm eq}$ of the NS \citep{bh91}. However, the derived values of $P_{\rm eq}$ for wind-accreting SGXBs are always lower than the observed ones of $P_{\rm s}$. It was suggested that the present $P_{\rm s}$ distribution may result from the equilibrium spin period when the companion star was still on the main sequence with a much weaker wind, and the wind of a SG is unable to transfer enough angular momentum to move the NS towards a new equilibrium value (e.g. Stella et al. 1986). The situation is more complicated in BeXBs. Be star winds are known to be disklike rather spherically expanding as in SGs, and a Be star can transform to be a B star and vice versa from time to time. The mechanism for this transition is still unknown. The varying Be star wind and the eccentric orbit imply that there does not exist a stable equilibrium spin period. \citet{wv1989} argued that the observational selection effects, that is, BeXBs are more likely to be observed when the NS moves within the dense equatorial disk wind, and the difference between the spin-up and spin-down timescales when the NS accretes within and outside of the disk wind, imply that $P_{\rm s}$ is potentially correlated with the accretion rate during outbursts, and thus $P_{\rm orb}$. Recently \citet{k2011} showed that the $P_{\rm s}-P_{\rm orb}$ correlation in BeXBs becomes more dispersed and a bimodal distribution for both $P_{\rm s}$ and $P_{\rm orb}$ seems to exist. While the bimodality is somewhat marginal in $P_{\rm orb}$, the $P_{\rm s}$ distribution has a clear gap at $\sim 40$ s with two peaks around 10 s and 200 s, respectively. \citet{k2011} proposed that two types of supernovae (SNe) may be responsible for the two subpopulation of BeXBs: the electron-capture supernovae (ECS) usually produce NSs with shorter spin periods, and lower eccentricities, while iron core-collapse supernovae (CCS) are preferred for NSs with longer spin periods and higher eccentricities. The original idea for the ECS candidates in BeXBs stems from a subclass of BeXBs that can be explained by low NS kicks \citep{p2002}. These BeXBs are characterized by long spin periods, persistent low X-ray luminosities ($\sim 10^{34}-10^{35}$ ergs$^{-1}$), wide binary orbits ($P_{\rm orb}> 30$ days), and low eccentricities ($e\lesssim 0.2$). The prototype of them is X Persei \citep[$P_{\rm s}=837$ s,][]{w1976}. Other sources include RX J0146.9$+$6121 \citep[1412 s,][]{h1998}, RX J1037.5$+$5647 (860 s) and RX J0440.9+4431 (202.5 s) \citep{rr1999}. Recently discovered BeXBs SXP 1062 \citep[1062 s,][]{h2012}, 1RXS J225352.8$+$624354 \citep[47 s,][]{e2013}, and SWJ2000.6$+$3210 \citep[890 s,][]{p2013} may also belong to this subclass. \citet{p2004} and \citet{vdh2004} suggested that the ECS mechanism may account for the low kicks in these BeXBs. This is different from the proposal by \citet{k2011} that ECS-BeXBs may have short spin periods, relatively narrow orbits and low eccentricities. A thorough investigation on the origin of the subpopulations of BeXBs requires a population synthesis incorporating stellar and binary evolution, SN explosions, the Be star wind structure, mass transfer processes, and the NS evolution, which is beyond the scope of this paper. Here we focus on the origin of the the $P_{\rm s}$ distribution, which shows a much clearer bimodal feature than the $P_{\rm orb}$ distribution in the current sample. We expect that the orbital periods of BeXBs may be largely dependent on the initial parameters of the progenitor binaries and the SN mechanisms, since tidal interaction is unable to change them effectively in wide orbits, while the NS spin periods are likely to be determined by the accretion processes in BeXBs. In Section 2 we compare the statistical characteristics of the outbursts in the two subpopulations, and present qualitative argument that the bimodal $P_{\rm s}$ distribution can be ascribed to different accretion modes of the NSs. We discuss the possible implications on the SN mechanisms in Section 3, and summarize in Section 4. | We emphasize that the argument in Section 2 is based on the global properties of the BeXB population rather than the individual source characteristics. Actually Fig.~1 shows that there is no strict one to one correspondence between short/long $P_{\rm s}$ and the occurrence of giant/normal outbursts. The transient nature of BeXBs means that it is impossible to monitor all outbursts for each source, so our classification of the outburst behavior is obviously incomplete. Evolution of the Be star disk also influences the characteristics of outbursts and the NS spin evolution. Nevertheless, the features in Figs.~1- 3 strongly suggest that different accretion modes of the NSs may be behind the origin of the subpopulations. The SN mechanisms can be related with the subpopulations through their influence on the initial orbital period, eccentricity, and misalignment of the Be star disk. Population synthesis calculations by \citet{l2009} showed that the ECS channel may be efficient at forming BeXBs, especially in the SMC, in which the population of HMXBs has been found to have relatively wide orbits and low eccentricities. However, they did not compare the characteristics of BeXBs formed through the CCS and ECS channels. To examine the influence of different kinds of SNe on the orbital period distribution, we employ a Monte-Carlo method to simulate the formation of BeXBs. We adopt the binary population synthesis (BPS) code developed by \citet{h2000,h2002} to calculate the evolutions of a large number of the primordial binaries, which is similar to the code {\em StarTrack} used by \citet{l2009}. We consider Solar abundance for the stars, with most of the input parameters (i.e., the distributions of the orbital separation, mass ratio, the initial mass function of the mass of the primary star) same as the standard ones described by \citet{h2002}. Detailed description of the method and calculated results will be presented elsewhere \citep{sl2014}. Some relevant key points in the calculations are listed below. \begin{enumerate} \item We consider both CCS and ECS for the NS formation. For ECS, we adopt the following criterion suggested by \citet{f12}. If the core mass $ M_{\rm c,bagb}$ of the primary star (i.e., the NS's progenitor) at the base of asymptotic giant branch is between $ 1.83 M_{\odot} $ and $ 2.25 M_{\odot} $, the CO core will non-explosively burn into an ONe core, and the core mass is accumulated gradually. If its mass can reach $ M_{\rm esc} = 1.38 M_{\odot} $, the ONe core will collapse due to electron capture into Mg and form a NS. If the mass is less than $ M_{\rm esc} $, it will leave an ONe WD. \item We apply a Maxwellian distribution for the SN kick velocity imparted to the newborn NS, with one dimensional rms velocity $\sigma=265$ kms$^{-1}$ for CCS \citep{h2005} and $50$ kms$^{-1}$ for ECS \citep{p2002}, respectively. \item We define a BeXB to be a binary consisting of a NS and a main-sequence (i.e., core H burning) companion star with mass between 8 $M_{\odot}$ and 20 $M_{\odot}$, which does not fill its Roche-lobe. Since Be stars are rapidly rotating, we also consider the influence of tidal synchronization on the Be stars. Only systems with the synchronization timescale greater than the main-sequence lifetime of the Be star are taken into account. \end{enumerate} The calculated normalized orbital period distributions of BeXBs are plotted in Fig.~5. The red and black curves denote systems formed through CCS and ECS, respectively. In the left panel the secondary star is regarded as a Be star when its rotational velocity is accelerated to $80\%$ of its break-up velocity due to previous mass transfer. In the right panel we assume that a constant fraction of B stars are Be stars. Note that in both cases the orbital periods have similar, wide distributions, with CCS-BeXBs peaked at $P_{\rm orb}\sim 40-50$ days and ECS-BeXBs at $\sim 100$ days. It seems that CCS-BeXBs dominate at $P_{\rm orb}<\sim 20-50$ days, but at longer $P_{\rm orb}$ the numbers of the two classes of objects are comparable. We need to caution that the number and distribution of ECS-BeXBs depend on the adopted mass range of the ECS progenitors \citep{n1984,n1987,p2004,s2007,p2008}, which is not well understood. Nevertheless, a bimodal orbital period distribution from the two types of SN channels seems not to exist. Another factor that can influence the $P_{\rm s}$ distribution and might be related to the SN mechanisms is the NS magnetic field. In the above estimates we assume that all the NSs possess a magnetic field of order $10^{12}$ G. It is not known how NSs born in CCS and ECS differ in their magnetic fields. For BeXBs with measured cyclotron resonance scattering features in their X-ray spectra, the characteristic line energies range from 10 to 55 keV \citep[][and references therein]{p2012}, suggesting comparable field strengths (a few $10^{12}$ G) in the NSs. However, there is indirect evidence that the NS magnetic field strengths may occupy a wide range. For example, from the measured spin-up rate in the 9.28 s Be/X-ray pulsar 2S1553$-$542, \citet{pp2012} derived a relatively low field $B\sim 5\times 10^{11}$ G for the NS. In another case, the spin-down rate measured in the 1062 s Be/X-ray pulsar SXP1062 implies that the NS possesses a magnetic field $B\gtrsim 10^{14}$ G \citep{fl2012}. The large scatter of BeXBs in the Corbet diagram may be partly due to the distribution and evolution of the NS magnetic field. | 14 | 4 | 1404.0219 |
1404 | 1404.1359_arXiv.txt | We use the observed radial profiles of the mass surface densities of total, $\Sigma_g$, and molecular, $\Sigma_{\rm{H2}}$, gas, rotation velocity and star formation rate (SFR) surface density, $\Sigma_{\rm{sfr}}$, of the molecular-rich ($\Sigma_{\rm{H2}}\ge\Sigma_{\rm{HI}}/2$) regions of 16 nearby disk galaxies to test several star formation laws: a ``Kennicutt-Schmidt'' law, $\Sigma_{\rm{sfr}}=A_g\Sigma_{g,2}^{1.5}$; a ``Constant Molecular'' law, $\Sigma_{\rm sfr}=A_{\rm H2}\Sigma_{\rm{H2,2}}$; the turbulence-regulated laws of Krumholz \& McKee (KM05) and Krumholz, McKee \& Tumlinson (KMT09), a ``Gas-$\Omega$'' law, $\Sigma_{\rm{sfr}}=B_\Omega\Sigma_g\Omega$; and a shear-driven ``GMC Collision'' law, $\Sigma_{\rm{sfr}}=B_{\rm{CC}}\Sigma_g\Omega(1-0.7\beta)$, where $\beta\equiv\:d\:{\rm{ln}}\:v_{\rm{circ}}/d\:{\rm{ln}}\:r$. If allowed one free normalization parameter for each galaxy, these laws predict the SFR with rms errors of factors of 1.4 to 1.8. If a single normalization parameter is used by each law for the entire galaxy sample, then rms errors range from factors of 1.5 to 2.1. Although the Constant Molecular law gives the smallest rms errors, the improvement over the KMT, Kennicutt-Schmidt and GMC Collision laws is not especially significant, particularly given the different observational inputs that the laws utilize and the scope of included physics, which ranges from empirical relations to detailed treatment of interstellar medium processes. We next search for systematic variation of star formation law parameters with local and global galactic dynamical properties of disk shear rate (related to $\beta$), rotation speed and presence of a bar. We demonstrate with high significance that higher shear rates enhance star formation efficiency per local orbital time. Such a trend is expected if GMC collisions play an important role in star formation, while an opposite trend would be expected if development of disk gravitational instabilities is the controlling physics. | \label{S:intro} Understanding the rate at which stars form from a given galactic gas inventory is a basic input for models of galaxy evolution. Global and kiloparsec-scale correlations between star formation activity, gas content and galactic dynamical properties have been observed (e.g., Kennicutt \& Evans 2012). However, most star formation is known to occur in highly clustered $\sim 1-10$~pc-scale regions within giant molecular clouds (GMCs) and the physical processes linking these large and small scales, i.e., the ``micro-physics'' of galactic star formation laws, remain uncertain. Tan (2010, hereafter Paper I), analyzed data from Leroy et al. (2008) for the molecular dominated regions of 12 nearby disk galaxies. The predictions of six star formation laws, described below, were tested against the observed radial profiles in the galaxies. In this paper, after summarizing the star formation (SF) laws to be considered (\S\ref{S:laws}) and adopting similar methods as Paper I (\S\ref{S:method}), we have extended this work by: (1) utilizing a modestly expanded sample of 16 galaxies, which are now explicity selected to be relatively large disk galaxies with mean circular velocity $\geq 100\:{\rm km\:s^{-1}}$ (11 galaxies overlap with the sample of Paper I); (2) also considering ``molecular rich'' regions where $\Sigma_{\rm HI}/2<\Sigma_{\rm H2}<\Sigma_{\rm HI}$, in addition to the ``molecular dominated'' regions (the results of relative comparison of the different SF laws in these regions are presented in \S\ref{S:lawtest}); (3) searching for correlations of SF law parameters with galactic dynamical properties (\S\ref{S:properties}), i.e., galactic disk shear (rotation curve gradient), rotation speed, and presence of a bar. We conclude in \S\ref{S:conclusion}. | \label{S:conclusion} We have tested six star formation laws against the resolved profiles of 16 molecular dominated and molecular rich regions of nearby, massive disk galaxies. There is a range from about a factor of 1.4 to 1.8 dispersion in the residuals of the best-fits when allowing each galaxy one free parameter to normalize the star formation laws, rising to 1.5 to 2.1 when a single global parameter is fit to the sample for each law. Since the different laws involve different inputs, which can have varying levels of observational uncertainties and varying degrees to which they connect to fundmental galactic physical properties, the relative ordering of the laws is not of primary importance (formally, the Constant Molecular law does best in having the smallest residuals; see Table 3). More interesting is the comparison of laws within similar classes. Thus the turbulence-regulated model of KMT09 is seen to be a clear improvement over the KM05 model. The GMC Collision model improves over the Gas-$\Omega$ model. The reason for this latter effect is the predicted decrease in star formation efficiency per orbital time with decreasing shear rate (increasing $\beta$) in the disk due to a reduced rate of shear-driven GMC-GMC collisions (Gammie et al. 1991; Tan 2000; Tasker \& Tan 2009), which is elucidated in Figure~\ref{fig:beta}. We estimate that the significance of this trend over the range $0<\beta<0.5$ is at least at the $4\sigma$ level. Such a trend is the opposite of that expected if development of gravitational instabilities (e.g., leading to GMC formation) from the diffuse interstellar medium is the rate limiting step for star formation activity. Confirmation of this result with a larger sample of galaxies, together with more careful investigation of potential systematic uncertainties, such as galactic radial gradients in normalization of star formation rate indicators and CO to $\rm H_2$ conversion factors (although this latter does not appear to be a major effect; Sandstrom et al. 2013), is desirous. More tentatively, we have found evidence that the presence of a bar boosts star formation efficiency per orbit. This could potentially be due to the influence of the bar on the strength of spiral arms (or more general axisymmetric structure; Kendall et al. 2011) in the larger-scale star-forming disks of the galaxies, although the influence of spiral arms on star formation activity in NGC 628, NGC 5194 and NGC 6946 has been found to be small ($\lesssim 10\%$) (Foyle et al. 2010). More detailed study of the influence of spiral arms (including potential inducement by the presence of bars) and their effect on star formation efficiency per orbit in a larger sample of galaxies is needed. Also worthwhile is further theoretical work on the influence of bars and spiral arms on the global GMC collision rate and its link to star formation, i.e., compared to that in more axisymmetric, flocculent galaxies. | 14 | 4 | 1404.1359 |
1404 | 1404.4056_arXiv.txt | Currently there are two competing scenarios to explain the origin of the stellar population in globular clusters (GCs). The main difference between them is whether or not multiple events of star formation took place within GCs. In this paper we present the star formation history (SFH) of Cluster 1, a massive young cluster in NGC 34 $(\sim10^7\mbox{ M}_\odot)$. We use \dinbas, a spectrum fitting algorithm, to retrieve the SFH and find that Cluster 1 is consistent with a single stellar population of solar metallicity with an age of $100\pm30$ Myr and a mass of $1.9\pm0.4\times10^7\mbox{ M}_\odot$. These results are in conflict with the expectations/predictions of the scenarios that invoke extended or multiple episodes within 30--100 Myr of the initial star-formation burst in young massive clusters. | \label{sec:intro} The classical notion of globular clusters (GCs) being simple stellar populations (SSPs) has been challenged by the presence of chemical anomalies and multiple sequences in the colour--magnitude diagrams (CMD) of GCs. The chemical anomalies are present only in light elements (namely C, N, O, Na and Al - e.g. \citealt{Carretta:2009p2165}) and are generally found only within GCs and not in the field population (e.g. \citealt{Martell:2011p2166}). To date, only a handful of clusters have been found with significant Fe spreads among their stellar populations (e.g. $\omega$ Centauri, Terzan 5, M52, M22 and NGC1851 - e.g. \citealt{2014arXiv1401.4323M} and references therein). Additionally, significant spreads in He abundance within GCs have been proposed to explain multiple main sequences and turn-offs, as well as the shape of the horizontal branch, in colour--magnitude diagrams of some GCs (e.g. \citealt{Milone:2012p2160}). Most models that attempt to explain the chemical anomalies and CMD morphology observed in GCs assume that these features are the product of multiple generations of stars. The basic idea is that a second generation of stars is created from the chemically processed ejecta of some very precise kinds of stars from the first generation (\emph{polluter} stars). Stars that have been suggested to be \emph{polluters} include: Asymptotic Giant Branch (AGB) stars (e.g. \citealt{Dercole:2008p2154}), fast rotating massive stars (also known as spin-stars e.g. \citealt{Decressin:2009p2155}), and massive stars in interacting binary systems \citep{DeMink:2009p2156}. All these multiple-population scenarios do well reproducing many of the observed anomalies mentioned before, and they predict that star clusters forming today should undergo a second generation of star formation. If spin-stars or massive interacting binaries are the source of the enriched material, then the second generation is expected to form within $\sim10$~Myr of the first generation. Alternatively, if AGB stars are the source, a difference of $30-200$~Myr between the 1st and 2nd generation is expected (e.g. \citealt{Conroy:2011p1997}). Recently, an alternative scenario has been proposed that does not invoke multiple star formation events within massive clusters. In this scenario, \cite{Bastian:2013p2152} suggest that the chemically enriched material is ejected by spin stars or high mass interacting binaries, and is accreted onto circumstellar disks of pre-main-sequence low mass stars of the same generation. Additionally it has been suggested that the observed extended main sequence turn-offs (eMSTO) and ``dual red clumps'' observed in intermediate age (1--2 Gyr) Small and Large Magellanic Cloud (LMC) clusters may be the product of extended (200--500 Myr) star formation events (e.g. \citealt{Mackey2007,Goudfrooij:2009p2234,Goudfrooij:2011p2248,Goudfrooij:2011p2254,Milone:2009p2291,Rubele:2013p2304}). Some studies propose a common evolution of these intermediate age clusters with GCs (e.g. \citealt{Conroy:2011p1997}). On the other hand, there are some claims for the opposite, for example \cite{Mucciarelli:2008p2339} indicated that the eMSTO of intermediate-age clusters were not related to the multiple stellar populations seen in globular clusters, due to the lack of abundance spreads between the stars of the younger clusters. Alternatively, different mechanisms have been put forward to explain such anomalies in intermediate age clusters e.g. stellar rotation (e.g. \citealt{Bastian:2009p2488,Yang:2013p2467}) or interacting binaries (e.g. \citealt{Yang2011}). \cite{Bastian:2013p2152} argue that if such extended (or multiple) star formation events took place in these intermediate age clusters and GCs, it would be expected that younger ($<500$ Myr) massive clusters should be currently forming stars. To test this, \cite{Bastian:2013p2199} studied the CMD of two young (180--280 Myr) massive ($\sim10^5$ M$_\odot$) clusters in the LMC and assessed an upper limit of 35 Myr for the possible age spread in these clusters. Also \cite{Bastian:2013p2022} presented a catalog containing more than 100 young (10--1000 Myr) massive ($10^4$--$10^8$ M$_\odot$) clusters where they do not find evidence of any ongoing star formation within the clusters, and concluded that any extended star formation within clusters lasting for hundreds of Myr are ruled out at high significance (unless strong stellar initial mass function -IMF- variations are invoked). Their study was sensitive to $\sim2$\% of the current cluster mass being formed today. If such extended ($200-500$~Myr) star formation events were common, the authors estimate that roughly 50\% of their sample should have shown evidence for ongoing star-formation. In this work, we approach the problems of the origin of multiple populations in GCs and eMSTO/``dual red clumps'' detected in intermediate age clusters by analysing the integrated spectrum of a young massive star cluster, looking for evidence for multiple events of star formation within this cluster. The cluster we chose for this initial study is young ($\sim$150~Myr), is found in the wet-merger galaxy NGC 34 and does not show any evidence for ongoing star-formation, based on the lack of optical emission lines in its spectrum (e.g., \citealt{Schweizer:2007p2018}; \citealt{Bastian:2013p2022}). This young globular cluster (Cluster 1, hereafter) has an estimated mass of about 15--20$\times10^6\mbox{ M}_\odot$ \citep{Schweizer:2007p2018}, which is 3--4 times more massive than that of $\omega$ Centauri, the most massive GC in the Galaxy. The fact that Cluster 1 is so massive and young makes it rather suitable to probe both families of formation scenarios, given that it can easily retain the ejecta of the polluter stars of the first generation, and we should be able to find evidence of a second generation of stars if a secondary burst has already taken place in the cluster. The paper is organised as follows: In \S \ref{sec:data} we present the optical spectrum of Cluster 1 and in \S \ref{sec:dinbas} we introduce the fitting method and models used in the SFH analysis. The degeneracies and uncertainties in the fits are discussed in \S \ref{sec:uncertainties}, and we discuss our results and present our conclusions in \S \ref{sec:discussion} and \S\ref{sec:conclusions}, respectively. | \label{sec:conclusions} By fitting the normalised spectrum of Cluster 1 in NGC 34 with model SSP spectra, we have determined an age of $100\pm30$ Myr for the cluster and estimated a mass of $1.9\pm0.4\times10^7\mbox{ M}_\odot$, based on published photometry and SSP models for this age. We do not find evidence for multiple star formation episodes, and we can confidently rule out the presence of a 2nd generation of stars for ages outside the range from 70 to 130 Myr with mass ratios between the second and first generation greater than 0.1. These results are consistent with GC formation scenarios where multiple generations of stars are separated by $<30$ Myr in age (e.g. \citealt{Decressin:2009p2155,DeMink:2009p2156}) or scenarios that do not invoke multiple star forming events (\citealt{Bastian:2013p2152}). Our results do not support any GC formation scenarios that involve multiple generations of stars separated by $>30$ Myr in age. However, it is still possible that a secondary burst might happen in the future (i.e. with an age difference between the first and second generation of stars that is greater than 100 Myr). To improve our understanding of how GCs form, further spectroscopic studies of young massive clusters covering a wide range of ages are needed. In a separate paper, Cabrera-Ziri et al. (in prep.) determine the SFH of 6 young (12--500 Myr) massive ($>10^6\mbox{ M}_\odot$) clusters from an ongoing spectroscopic survey. These result are consistent with the findings of \cite{Bastian:2013p2199}, who do not find any large age spreads in young massive LMC clusters, and they also disagree with GC formation scenarios that predict extended SFHs (e.g. \citealt{Conroy:2011p1997}). Finally, we conclude that \dinbas\ capabilities (i.e. SED fitting of just a few ages) are ideal for the study of the integrated spectra of young clusters, given that they reduce significantly the amounts of non-genuine components (i.e. ages) compared to traditional SED fitting algorithms, consequently simplifying the analysis of the results. | 14 | 4 | 1404.4056 |
1404 | 1404.3440_arXiv.txt | Obtaining accurate photometry of bright stars from the ground remains tricky because of the danger of overexposure of the target and/or lack of suitable nearby comparison star. The century-old method of the objective wire mesh used to produce multiple stellar images seems attractive for precision CCD photometry of such stars. Our tests on $\beta$ Cep and its comparison star differing by 5 magnitudes prove very encouraging. Using a CCD camera and a 20 cm telescope with objective covered with a plastic wire mesh, located in poor weather conditions we obtained differential photometry of precision 4.5 mmag per 2 min exposure. Our technique is flexible and may be tuned to cover as big magnitude range as 6 -- 8 magnitudes. We discuss the possibility of installing a wire mesh directly in the filter wheel. | \label{s0} Renewed interest in studies of bright stars in general stems from their suitability to long term spectroscopic monitoring with modest telescopes for asteroseismic purposes. As a byproduct of the extra-solar planet quest emerged the new generation of fiber-fed echelle spectrographs capable of measuring radial velocities of those stars accurate to meters per second. This opens a new window for studies of multiple/low amplitude coherent and stochastic (solar-type) oscillations of luminous stars. Both kinds of oscillations are of great use for asteroseismology, particularly to constrain the efficiency of convection and mixing in stellar interiors. However, precise mode identification demands knowledge of phase shifts between velocity and light curves, as well as color dependence of photometric amplitudes \citep{das02}. Accurate photometry of bright stars remains tricky because of danger of overexposure of the target and/or lack of suitable nearby comparison stars. Last half century produced relatively few long-term light curves for such stars. It may be argued that the best results can be obtained from Space and using wide angle cameras. This became the motivation for the constellation of BRITE nano-satellites in the process of launching (Orleański et al., 2010). In the present paper we investigate suitability of a venerable photographic technique to obtain good quality CCD photometry of very bright stars {\em from the ground}. For this purpose we applied a CCD camera fitted to a 20cm telescope with its objective covered with a dense wire mesh. Late XIX century attempts by astronomers to employ photographic plates for stellar photometry were hampered by the need to calibrate a non-linear response of the photographic emulsion to light. For {\em Carte du Ciel} Kapteyn in 1891 proposed to make alternate exposures with and without wire mesh cover of the objective to vary aperture (c.f. \citealt{wea46}). Later, \citet{her10} noted that a sufficiently dense wire mesh would produce multiple diffraction images for each star, thus alleviating the need for multiple exposures. He argued that the rate of illumination between different images of the same star would remain fixed. His idea applied either for the direct images of the sky or for images of the calibration source exposed on the edge of the plate. However, as far as we are aware, no images of sky taken through the wire mesh were reported in the electronic age of astronomy. | \label{s3} Several approaches have been utilized in the past for obtaining accurate photometry of bright stars using CCD detectors, including: \begin{enumerate} \item alternate long and short exposures - prone to residual bulk image on CCD chip and atmospheric condition changes. Additionally short exposure times are heavily affected by scintillation; \item snapshot observation technique \citep{man11} - requires precise and multiple telescope slews and is sensitive to atmospheric condition changes. Both (1) and (2) require photometric conditions; \item covering a fraction of the detector with a neutral density filter reduces useful telescope field of view by introducing a "penumbra" area and for a given filter yields limited dynamic range. \end{enumerate} The objective wire mesh technique described in Sect. \ref{s1} suffers from none of these drawbacks and produces useful CCD dynamic range between the target and the comparison star of up to at least 5 magnitudes, depending on selection of an appropriate fringe of the target star. Even a wider range of magnitude differences should be available by thinning of mesh wires, so that the low-order fringes become fainter. Thus our technique, combined with the appropriate exposure time, permits free choice of the comparison star to meet such criteria like scintillation time averaging or appropriate filling of the CCD pixel well. Results of Sect. \ref{s2} demonstrate that in this way excellent photometric precision may be reached in poor climate with inexpensive equipment. Immediate application of our technique would be for ground follow-up observations for BRITE constellation of satellites. Space photometry may reach several orders of magnitude better precision than possible from the ground. However, due to reliance on mechanical devices for accurate pointing its time span is limited and so is frequency resolution. Thus, for sufficiently large amplitude oscillations, ground observations still remain useful. The wire mesh does not have to be installed on the objective. It may be convenient to place it directly on the photometric filter. In that case the mesh cell size should be reduced proportionally to the $a/f$ ratio, where $f$ is the telescope effective focal length and $a$ is the distance between the mesh and the image. With an internal wire mesh each star fringe pattern is created by a different mesh section, but with a proper telescope tracking and non-rotating field of view this should always be the same section for a given star. Therefore, the relative intensity of diffraction fringes is preserved, making relative photometry still possible, but differs from star to star, which prevents from accurate photometric calibration. Our test of this variant of the wire mesh technique seems encouraging, but this concept requires further investigation. The wire mesh technique could be useful not just for BRITE follow-up observations, but could also provide parallel photometric observations for high resolution spectroscopic observations with larger telescopes, e.g. similar to our PST2 project. | 14 | 4 | 1404.3440 |
1404 | 1404.1445_arXiv.txt | In this paper, we focus on general features of quintessential inflation which is an effort to unify inflation and dark energy using a single scalar field. We describe a class of models of quintessential inflation which can give rise to the tensor to scalar ratio of perturbations consistent with recent BICEP2 measurements. The scale of inflation in the model is around the GUT scale and there is large parameter space consistent with the recent findings. | One of the most outstanding and clean predictions of inflationary paradigm is related to relic gravity waves \cite{Grishchuk:1974ny,Grishchuk:1977zz, Starobinsky:1979ty,Allen:1987bk,Sahni:1990tx,Souradeep:1992sm, Giovannini:1998bp,Giovannini:1999bh,Langlois:2000ns,Kobayashi:2003cn, Hiramatsu:2003iz,Easther:2003re,Brustein:1995ah,Gasperini:1992dp, Giovannini:1999qj,Giovannini:1997km,Gasperini:1992pa,Giovannini:2009kg, Giovannini:2008zg,Giovannini:2010yy,Tashiro:2003qp,Sahni:2001qp} which are generated quantum mechanically in the early Universe. The primordial tensor perturbations induce B mode polarization in microwave background spectrum such that the effect depends upon the tensor to scalar ratio of perturbations $r$. Since the effect was not observed , the tensor to scalar ratio was supposed to be negligibly small. However, the recent observations on CMB polarization has demonstrated that the effect is sizeable, namely, the scalar to tensor ratio of perturbations, $r=0.2_{-0.05}^{+0.07}$ \cite{Ade:2014xna} such that the scale of inflation is around the GUT scale. Quintessential inflation \cite{Peebles:1999fz,Sahni:2001qp,Sami:2004xk,Copeland:2000hn,Huey:2001ae, Majumdar:2001mm,Dimopoulos:2000md,Sami:2003my,Dimopoulos:2002hm,Rosenfeld:2005mt, Giovannini:2003jw,Dimopoulos:2002ug,Nunes:2002wz,Dimopoulos:2001qu, Dimopoulos:2001ix,Yahiro:2001uh,Kaganovich:2000fc,Peloso:1999dm, Baccigalupi:1998mn,Hossain:2014xha}, a unified description of inflation and dark energy using a single scalar field, is necessarily followed by kinetic regime responsible for blue spectrum of relic gravity waves \cite{Giovannini:1998bp, Giovannini:1999bh,Giovannini:1999qj,Sahni:2001qp}. These scenarios can be classified into Type I and Type II. In first type, we consider models for which the scalar field potential is exponentially steep for most of the history of universe and only at late times the potential turns shallow. In Type II, we place models with potentials, shallow at early times followed by steep behavior thereafter. Ideally, quintessential inflation requires a potential that could felicitate slow roll in the early phase followed by approximately steep exponential behavior such that the potential turns shallow only at late times. Steep nature of potential is necessitated for the radiative regime to commence and peculiar steep behavior is needed to realize the scaling regime. However, the generic potentials do not change their character so frequently, they rather broadly come into two said categories:\\ Type A: The inverse power law and $\cos$ hyperbolic potentials belong to this category. In this case one requires to assist slow roll by an extra damping at early times. In Randall-Sundrum scenario \cite{Randall:1999ee, Randall:1999vf}, the high energy corrections to Einstein equations \cite{Shiromizu:1999wj} give rise to brane-damping which assists slow roll along a steep potential \cite{Sahni:2001qp,Copeland:2000hn,Huey:2001ae, Majumdar:2001mm,Maartens:1999hf}. As the field rolls down its potential, high energy corrections cease leading to graceful exit from inflation. Unfortunately, the tensor to scalar ratio in this case is too large, $r\simeq 0.4$ to be consistent with observations and the steep brane-world inflation is therefore ruled out.\\ Type B: In this case, the field potential stays steep after inflation. In this case, the late time behavior can be achieved by invoking an extra feature in the potential. For instance, massive neutrino matter with non-minimally coupled to scalar field can give rise to minimum of the potential at late times when neutrinos become non-relativistic \cite{Wetterich:2013jsa,Hossain:2014xha}. In this paper, we shall describe a class of models of quintessential inflation of this category and look for parameter space which could comply with the recent measurement of B mode polarization spectrum\cite{Ade:2014xna}. | In this paper we have investigated a class of models that can successfully give rise to quintessential inflation. The Lagrangian of the single field system under consideration contains three free parameters $\tilde{\alpha}$, $\alpha$ and $\beta$ such that $\beta$ is related to scale of inflation and $\tilde{\alpha}$ defines the tensor to scalar ratio $r$ for a given number of efolds. As for $\alpha$, it is fixed by the post inflationary requirements, namely, nucleosynthesis constraint \cite{Ade:2013zuv}. For the observed values of $r$ from BICEP2 \cite{Ade:2014xna} and $\mathcal{N}=60$, the parameter $\tilde{\alpha}$ in the model ranges from $0.063$ to $0.183$ consistent with the BICEP2 measurements, (see Fig. \ref{fig:r_alpha}) such that the scale of inflation in this case is around the GUT scale. The distinguished feature of the model includes a blue spectrum of stochastic background of relic gravitational waves produced during inflation. The blue spectrum of relic gravity waves associated with the kinetic regime after inflation, is a generic feature of quintessential inflation irrespective of an underlying model \cite{Sahni:2001qp,Sami:2004xk, Giovannini:1998bp,Giovannini:1999bh,Giovannini:1999qj}. However, the amplitude of relic gravity waves naturally depends upon the tensor to scalar ratio of perturbations and we have quoted here $\Omega_{\rm GW}$ in accordance with the observed values of $r$. Fig. \ref{fig:RGW_r} shows the $r$ dependence of spectral energy density parameter ($\Om_{\rm GW}$). We should emphasize that we have neglected $n_\rmt$, the tilt of inflationary spectrum, in order to felicitate the analytical calculation. We reiterate that the blue spectrum here is nothing to do with blue tilt seen in BICEP2; the former is the consequence of kinetic regime which is a general feature in scenarios of quintessential inflation. We should also negligibly small values of running of the spectral index. note that the scenario under consideration predicts The BICEP2 findings, if confirmed, would rule out a large number of models including the currently favourite Starobinsky model. In purely theoretical perspective, the GUT scale of inflation as envisaged by the said measurements would throw a big challenge to model building in the framework of effective field theories. We hope that the forthcoming announcement from {\it Planck} collaboration and future observations would clarify the related issues and the same is eagerly awaited. | 14 | 4 | 1404.1445 |
1404 | 1404.3724_arXiv.txt | The relation between galaxies and dark matter halos is of vital importance for evaluating theoretical predictions of structure formation and galaxy formation physics. We show that the widely used method of abundance matching based on dark matter only simulations fails at the low mass end because two of its underlying assumptions are broken: only a small fraction of low mass ($<10^{9.5}\Ms$) halos host a visible galaxy, and halos grow at a lower rate due to the effect of baryons. In this regime, reliance on dark matter only simulations for abundance matching is neither accurate nor self-consistent. We find that the reported discrepancy between observational estimates of the halo masses of dwarf galaxies and the values predicted by abundance matching does not point to a failure of $\Lambda$CDM, but simply to a failure to account for baryonic effects. Our results also imply that the Local Group contains only a few hundred observable galaxies in contrast with the thousands of faint dwarfs that abundance matching would suggest. We show how relations derived from abundance matching can be corrected, so that they can be used self-consistently to calibrate models of galaxy formation. | \label{introduction} In the $\Lambda$CDM paradigm \citep[e.g.][]{Davis-1985, Frenk-1988}, stars form as gas cools in collapsed dark matter halos \citep[e.g.][]{White-1978, White-Frenk-1991}. The formation of galaxies involves both baryons and dark matter, but while only gas and stars are directly observable, numerical simulations of structure formation have largely been limited to the dark matter component, under the assumption that gravity is the only relevant force on large scales. A conceptually simple and yet very powerful method to connect galaxies and halos in this scenario, which does not require any detailed knowledge of the complex physics of galaxy formation, is abundance matching \citep[e.g.][]{Frenk-1988, Yang-2003, Kravtsov-2004, Vale-2006, Moster-2009, Guo-2010, Behroozi-2013, Moster-2013}. Assuming that a monotonic relation exists between some observable property of a galaxy (such as stellar mass) and some property of its dark matter halo (such as total mass), the relationship between both quantities can be computed from: \begin{equation}\label{eqn:abundance-matching} \int_{\rm{a_{h,min}}}^{\rm{a_{h,max}}} N_h(a_h) da_h = \int_{\rm{a_{\star,min}}}^{\rm{a_{\star,max}}} N_\star(a_\star) da_\star, \end{equation} where $N_h(a_h)$ and $N_\star(a_\star)$ are the numbers of halos and galaxies in the same volume, with properties $a_h$ and $a_\star$, whose maxima determine the upper limits $\rm{a_{h,max}}$ and $\rm{a_{\star,max}}$. For any lower limit, $\rm{a_{h,min}}$, a corresponding lower limit, $\rm{a_{\star,min}}$, is chosen such that Eq.~\ref{eqn:abundance-matching} is satisfied, and the average relation $a_\star(a_h)$ is then uniquely determined. Abundance matching is employed widely to infer quantities such as the stellar-to-total mass relation \citep[e.g.][]{Frenk-1988, Yang-2003, Guo-2010, Moster-2013}. It has also been used to constrain the mass of the Milky Way's halo from its stellar mass \citep{Guo-2010}, and to predict the total number of dwarf galaxies in the Local Group from an N-body simulation \citep{Garrison-Kimmel}. Generally regarded as assumption-free, abundance matching results are often interpreted as direct predictions of the underlying cosmological model. They have also been used as a benchmark for models of galaxy formation physics, such as semi-analytic models \citep[e.g.][]{Guo-2011}, and to calibrate hydrodynamic simulations \citep[e.g.][]{Scannapieco-2012, Munshi-2013}. However, models that reproduce the abundance matching relation for low-mass galaxies often require very strong feedback, which can result in an unrealistically high passive fraction \citep{Fontanot-2009, Weinmann-2012, Moster-2013}. Conversely, many hydrodynamical simulations that produce realistic dwarf galaxies appear to have halo masses significantly below those inferred by abundance matching \citep[e.g.][]{Sawala-Matter, Avila-Reese-2011}. In some cases, the stellar-to-total mass relation derived from abundance matching can also be compared directly to observations of individual galaxies. While they agree for galaxies in halos more massive than $\sim10^{12}\Ms$ \citep{Guo-2010}, discrepancies have been reported for lower mass halos \citep{Ferrero-2012}. In particular, dynamical mass estimates derived from stellar kinematics suggest stellar-to-total mass ratios for individual dwarf galaxies which are an order of magnitude higher than those inferred from abundance matching in $\Lambda$CDM. Recently, we have shown that simulations that model the evolution of dark matter and baryons as a single fluid subject only to gravity (henceforth referred to as ``Dark Matter Only'' or DMO simulations) do not produce the same abundance of halos as hydrodynamic simulations that include the full baryonic effects. In particular, the collapse and subsequent expulsion of baryons by feedback processes reduce the mass of individual halos \citep{Sawala-abundance}, a result that has since been reproduced and extended to higher masses \citep{Velliscig-2014, Cui-2014}. Here we use a new set of high-resolution hydrodynamical simulations of Local Group volumes to explore how the results of abundance matching are affected when the effects of baryons and the appearance of dark halos are included self-consistently. Unlike models that are calibrated to reproduce an abundance matching relation based on a DMO simulation {\it before} baryons are taken into account, the stellar-to-total mass relation produced in our simulations is consistent with the relation we derive from abundance matching {\it after} the effects of baryons are included. Furthermore, this relation agrees with the dwarf galaxy data, thus demonstrating that the reported high stellar-to-total mass ratios of these galaxies are entirely consistent with the $\Lambda$CDM model. This paper is organised as follows. In Section~\ref{sec:methods}, we describe the simulations on which our work is based. In Section~\ref{sec:results}, we describe our results: the fraction of halos that host galaxies is discussed in Section~\ref{sec:real-abundance}, the application to abundance matching in Section~\ref{sec:abundance-matching}, and the implications in Section~\ref{sec:implications}. We conclude with a summary in Section~\ref{sec:summary}. | \label{sec:summary} The commonly used method of abundance matching between N-body DMO simulations and the observed stellar mass function is not assumption free: it relies on the implicit assumption that structure formation can be represented by DMO simulations and that every halo hosts a galaxy. Using a set of high resolution simulations of Local Group volumes with and without baryons, we have shown that, at the low mass end, both of these assumptions are broken. Because of cosmic reionization, most halos below $3\times 10^9\Ms$ do not contain an observable galaxy. In this regime, the median stellar-to-total mass relation inferred directly from abundance matching only applies to (mostly unobservable) halos, but has no direct application to observable galaxies. Our simulations assume reionization in the optically thin limit without self-shielding, and radiative cooling which does not include molecular cooling at low temperatures. While these limitations may influence the impact of reionization on star formation rates in existing galaxies, the proto-galaxies that are prevented from star formation in our simulations reach neither the gas densities required for self-shielding nor the metallicities required for molecular cooling after reionization. Of course, our simulations cannot resolve the formation of Pop-III stars from primordial gas. Our results also assume that reionization is uniform and local variations might change the impact for individual galaxies, but probably not enough to affect our results substantially. By equating the cumulative abundances of objects from hydrodynamical and DMO simulations, we have shown how abundance matching results can be modified to account for baryonic effects. We stress that a reliance on the results from abundance matching obtained through dark matter only simulations can have serious consequences: it can lead to erroneous conclusions about cosmology, and, when used as a benchmark for hydrodynamic simulations or semi-analytical models, it can also lead to false conclusions about galaxy formation physics. Baryon effects have to be taken into account self-consistently for abundance matching to give a meaningful interpretation of the connection between galaxies and halos. | 14 | 4 | 1404.3724 |
1404 | 1404.5102_arXiv.txt | It has been around fifty years since R. K. Sachs and A. M. Wolfe predicted the existence of anisotropy in the Cosmic Microwave Background (CMB) and ten years since the integrated Sachs Wolfe effect (ISW) was first detected observationally. The ISW effect provides us with a unique probe of the accelerating expansion of the Universe. The cross-correlation between the large-scale structure and CMB has been the most promising way to extract the ISW effect from the data. In this article, we review the physics of the ISW effect and summarize recent observational results and interpretations. | \label{Sec:overview} After the discovery of the isotropic radiation of the Cosmic Microwave Background (CMB) of the Universe by A. A. Penzias and R. W. Wilson in 1965 \cite{PenziasWilson:65}, R. K. Sachs and A. M Wolfe predict the existence of anisotropy in the CMB associated with the gravitational redshift in 1967 \cite{SachsWolfe:67}. They fully integrate the geodesic equation in a perturbed Friedman-Robertson-Walker (FRW) metric in the fully general relativistic framework. The Sachs-Wolfe (SW) is the first paper that predicts the presence of the anisotropy in the CMB which now plays an important role for constraining cosmological models, the nature of dark energy, modified gravity, and non-Gaussianity of the primordial fluctuation. Let us begin by reviewing the history. In the first twenty years since the Sachs-Wolfe paper, most of the works were focused on the extension of the Sachs-Wolfe calculation to non-linear collapsed object \cite{ReesSciama:68, OstrikerVishniac:86}, or non-standard cosmological model such as topological defects \cite{KaiserStebbins:84}. Partridge and Wilkinson 1967 first gave a glimpse of the existence of inhomogeneity in the CMB temperature by using the Dicke Radiometer \cite{PartridgeWilkinson:67}. They found the temperature excess on the direction of the known quasar cluster position and considered it as the Rees-Sciama effect \cite{ReesSciama:68}. In the age of the COBE satellite, the Sachs-Wolfe paper attracted a huge attention. Most of the papers were focused on the theoretical prediction that was related to the observation; prediction of the amplitude of the quadrupole power for the SW effect. \cite{GoudaSugiyamaSasaki:91, GoudaSugiyama:91, GoudaSugiyama:92, BunnSugiyama:95}. Crittenden \& Turok 1996 pointed out that the gravitational potential may decay in the $\Lambda$ dominated Universe at $z < 1$ to produce the ISW signal \cite{CrittendenTurok:96}. They also proposed a novel method to detect the ISW effect by cross correlating the large-scale structure with the CMB. Kneissl et al. 1997 made an attempt to extract the ISW effect by cross correlating the CMB observed by the COBE with the ROSAT X-ray background \cite{Kneissl+:97}, Boughn \& Crittenden 2002 used the NVSS radio galaxies for the cross correlation \cite{BoughnCrittenden:02}, and Boughn et al. 1998 used the HEAO1 A2 X-ray background \cite{BoughnCrittendenTurok:98} but none of them could find the significant detection. The Sachs-Wolfe paper has attracted a renewed attention in the WMAP era. The first detection of the ISW was finally achieved by cross-correlating the WMAP first-year data with the number count of radio galaxies from the NVSS data, as well as with the HEAO1 A1 X-ray data \cite{BoughnCrittenden:04}. Subsequently a lot of detections with various mass tracers have been reported. In the early 2000's, much work was focused on obtaining cosmological constraint on dark energy models from the ISW effect, while in the late 2000's to present, more and more works studied various systematic effects which may enter in different ways for different measurement methods. In this paper, we review the ISW effect from theoretical derivation of the basic equations to the present cosmological interpretations. The paper is organized as follows. In section \ref{Sec:theory}, we revisit the derivation of the CMB anisotropy induced by the perturbation of the background geometry decomposed into scalar, vector and tensor modes. In section \ref{Sec:observation}, we discuss the statistical properties of the ISW effect and the method to measure it in the cross correlation with the large-scale structure. % We also discuss the possible systematic effects that affect our interpretations. In section \ref{Sec:models}, we provide cosmological applications of the ISW effect including constraints on dark energy and primordial non-Gaussianity. % In section \ref{Sec:summary}, we give a summary. | \label{Sec:summary} In this article we review the ISW effect induced by linear and non-linear structures of the Universe. It is well known that the Sachs-Wolfe effect and the ISW effect are simultaneously derived from the cosmological perturbation theory. The ISW effect can be generated by scalar, vector and tensor mode fluctuations of the metric but the significant contribution comes from the scalar mode, which is related to the time variation of the matter density fluctuation. The recent detection of the tensor-mode CMB polarization claimed by the BICEP2 collaboration \cite{BICEP2:14} can indeed be the polarization generated by the tensor-mode ISW effect from primordial gravitational waves. The time variation of the gravitational potential in the standard $\Lambda$CDM Universe comes from: 1) the coherent decay of the gravitational potential due to the accelerating expansion of the Universe, 2) the isotropic clustering inflow of mass lump toward the center of mass, and 3) the transverse motions of the clusters. The latter two effects are too tiny to detect with the current CMB experiments but it might be possible to detect with future experiments having finer angular resolutions such as ACTPol \cite{ACTPol:10}, SPTPol \cite{SPTPol:12} or COrE \cite{COrE:11}. We also review the observational studies of the ISW effect. The ISW effect was first detected by the cross correlation of the CMB observed by the WMAP with the number counts of the radio galaxies measured by the NVSS. The significance of detection was $2-3\sigma$ \cite{BoughnCrittenden:04}. Subsequently a number of detections have been reported with various tracers of dark matter using a variety of statistical methods. As the ISW effect reflects the large-scale fluctuations in the Universe, they can be used for constraining the cosmological models especially at low redshifts ($z<1$). The statistical error of the ISW effect is mostly dominated by sample variance but the detection significances vary among papers and are often not consistent with each other. It is partly due to the consequence of the different statistics they used, different treatments of the foreground, or different estimates or assumptions on the redshift distribution of galaxies. As we acquire more knowledge on those issues, we may be able to understand inconsistencies seen in table \ref{table:comparison}. All sky polarization data from the Planck will be useful for further studying the dust model of our Galaxy as well as the confirmation of the large amplitude of tensor mode fluctuation discovered by the BICEP2. In addition to this, complete BOSS spectroscopic galaxies and QSOs samples can be used for calibrating the redshift distribution of galaxies and QSOs. Therefore both the Planck and the BOSS data which are going to be delivered soon would allow us to extensively study the ISW effect with better understandings of the systematics. | 14 | 4 | 1404.5102 |
1404 | 1404.0511_arXiv.txt | {The oldest stars born before the onset of the main $s$-process are expected to reveal a pure $r$-process Ba/Eu abundance ratio.} {We revised barium and europium abundances of selected very metal-poor (VMP) and strongly $r$-process enhanced (r-II) stars to evaluate an empirical $r$-process Ba/Eu ratio.} {Our calculations were based on non-local thermodynamic equilibrium (NLTE) line formation for \ion{Ba}{ii} and \ion{Eu}{ii} in the classical 1D MARCS model atmospheres. Homogeneous stellar abundances were determined from the \ion{Ba}{ii} subordinate and resonance lines by applying a common Ba isotope mixture. We used high-quality VLT/UVES spectra and observational material from the literature.} {For most investigated stars, NLTE leads to a lower Ba, but a higher Eu abundance. The resulting elemental ratio of the NLTE abundances amounts, on average, log(Ba/Eu) = 0.78$\pm$0.06. This is a new constraint to pure $r$-process production of Ba and Eu. The obtained Ba/Eu abundance ratio of the r-II stars supports the corresponding Solar System $r$-process ratio as predicted by recent Galactic chemical evolution calculations of Bisterzo, Travaglio, Gallino, Wiescher, and K{\"a}ppeler. We present the NLTE abundance corrections for lines of \ion{Ba}{ii} and \ion{Eu}{ii} in the grid of VMP model atmospheres.} {} | \label{Sect:intro} Since the classical paper of \citet{B2FH} one traditionally believes that heavy elements beyond the iron group are produced in the neutron-capture nuclear reactions, which can proceed as the slow ($s$) or rapid ($r$) process. The distribution of the $s$-process abundances of the Solar System (SS) matter has been recognized as arising from a non-unique site. The isotopic distribution in the atomic mass range A = 90-208 is accounted for by the main $s$-process. It occurs in intermediate-mass stars of 1-4~$M_\odot$ during the asymptotic giant branch (AGB) phase, and it is rather well understood theoretically \citep[see, for example][and references therein]{1999ARA&A..37..239B}. The weak $s$-process runs in helium burning core phase of massive stars ($M > 10 M_\odot$) and it contributes to the nuclei up to A = 90 \citep{1989RPPh...52..945K}. The $r$-process takes place in an extremely n-rich environment. Astrophysical sites for the $r$-process are still debated, although they are likely associated with explosions of massive stars, with $M > 8 M_\odot$. We refer to the pioneering review of \citet{1978SSRv...21..639H} and also \citet{2004PhT....57j..47C} for further discussion on the $r$-process. In the Solar System matter, different isotopes were produced in differing proportions from the $s$- and $r$-process. For isotopes with A $> 90$, the $s$-abundances are evaluated from the Galactic chemical evolution (GCE) calculations. The difference between solar total and $s$-abundance is referred to as $r$-residual. The Solar system $r$-process (SSr) pattern is widely employed in stellar abundance comparisons \citep[see][and references therein]{Sneden2008}. However, different $s$-process calculations result in different SSr, in particular, for the chemical species with dominant contribution of the $s$-process to their solar abundances. This paper concerns with barium that together with Sr are the best observed neutron-capture elements in very metal-poor (VMP) and extremely metal-poor (EMP) stars. For example, the resonance lines of \ion{Ba}{ii} were measured in the [Fe/H]\footnote{In the classical notation, where [X/H] = $\log(N_{\rm X}/N_{\rm H})_{star} - \log(N_{\rm X}/N_{\rm H})_{Sun}$.} $\simeq -4.0$ stars in our Galaxy \citep{Francois2007} and also the classical dwarf spheroidal galaxies \citep{Tafelmeyer2010}. One of the most cited SSr is based on calculations of \citet[][stellar model, hereafter, A99]{Arlandini1999} who used stellar AGB models of 1.5 and 3~$M_\odot$ with half solar metallicity and predicted that 81~\%\ of the solar barium are of $s$-process origin (Table\,\ref{Tab:ba_isotope}). Very similar result was obtained by \citet[][hereafter, T99]{Travaglio1999} by the integration of $s$-abundances from different generations of AGB stars, i.e., considering the whole range of Galactic metallicities. \citet[][hereafter, B11]{Bisterzo2011} updated calculations of \citet{Arlandini1999} by accounting for the recent n-capture cross-sections, and they inferred a higher $s$-process contribution to the solar barium of 89~\%. Slightly lower solar $s$-abundance of Ba was obtained by \citet[][hereafter, B14]{2014arXiv1403.1764B} in their recent GCE calculations that considered the contributions from different generations of AGB stars of various mass. From one hand side, differences of 5 to 10~\%\ between different predictions are comparable with the quoted uncertainties in $s$-abundances (Table\,\ref{Tab:ba_isotope}). From other hand side, a small change in the $s$-abundance leads to significant change in the solar $r$-abundance of Ba and this has an important consequence for the Ba/Eu ratio of $r$-abundances, (Ba/Eu)$_r$. In contrast to Ba, solar europium is mostly composed of $r$-nuclei (Table\,\ref{Tab:ba_isotope}). The change in n-capture cross-sections shifts log(Ba/Eu)$_r$ from 0.93\footnote{In this study, the decomposition of the $s$- and $r$-process contributions is based on the meteoritic abundances of \citet{Lodders2009}.} (A99) down to 0.75 (B11). Updating the GCE calculations results in lower log(Ba/Eu)$_r$ = 0.87 (B14) compared with log(Ba/Eu)$_r$ = 0.96 of T99. It is worth noting, the SSr data cover the full range of predictions from the $r$-process models. In the classical waiting-point (WP) approximation, \citet{Kratz2007} inferred log(Ba/Eu)$_r \simeq$ 1, and log(Ba/Eu)$_r \simeq$ 0.8 was obtained in large-scale parameterized dynamical network calculations of \citet{Farouqi2010} in the context of an adiabatically expanding high-entropy wind (HEW), as expected to occur in core-collapse SNe. In this paper, we investigate whether the observed stellar abundances of Ba and Eu can constrain a pure $r$-process Ba/Eu ratio. Since the $s$- and $r$-process are associated with stars of different mass, their contribution to heavy element production varied with time. The Ba/Eu ratio is particularly sensitive to whether the $s$- or $r$-process dominated the nucleosynthesis. Old stars born before the onset of the main $s$-process should reveal more Eu relative to Ba compared with the SS matter. The existence of VMP stars enriched in the $r$-process element Eu was observationally established more than 30 years ago in a pioneering paper of \citet{Spite1978}, and the statistics was significantly improved in later studies \citep[for a review, see][]{Sneden2008}. Figure\,\ref{Fig:halo} shows the Ba/Eu ratios of a preselected sample of metal-poor (MP, [Fe/H] $< -1.5$) stars that reveal enhancement of Eu relative to Ba, with [Ba/Eu] $< -0.4$. The stars are separable into three groups, depending on the observed Eu abundance. \citet{HERESI} classified the stars, which exhibit large enhancements of Eu relative to Fe, with [Eu/Fe] $> 1$, as r-II stars. The Ba and Eu abundances of the r-II stars together with the sources of data are given in Table\,\ref{Tab:stars}. The stars of r-I type have lower Eu enhancement, with [Eu/Fe] = 0.3 to 1. The data for 32 r-I stars were taken from \citet{Cowan2002,Honda2004,HERESII,Ivans2006,Francois2007,Mashonkina2007,Lai2008}, and \citet{HE1219}. Here, we also deal with the 12 stars that reveal a deficiency of Eu relative to Fe, with $\mathrm{[Eu/Fe]} \le 0$ \citep{Honda2004,Honda2006,Honda2007,HERESII,Francois2007}. They are referred to as Eu-poor stars. The different groups have consistent within the error bars Ba/Eu ratios, independent of the observed Eu/Fe ratio, with the mean log(Ba/Eu) = 1.03$\pm$0.12, 1.08$\pm$0.13, and 1.14$\pm$0.08 for the r-II, r-I, and Eu-poor stars, respectively. Hereafter, the statistical error is the dispersion in the single star abundance ratios about the mean: $\sigma = \sqrt{\Sigma(\overline{x}-x_i)^2/(n-1)}$. The observed stellar abundance ratios are more than 1$\sigma$ higher compared with the Solar System $r$-process ratio log(Ba/Eu)$_r$ = 0.87 based on recent updated $s$-process calculations of B14, although they are consistent within the error bars with the SSr of T99. To make clear the situation with stellar abundances of Ba and Eu, we concentrate on the r-II stars that are best candidates for learning about the details of the $r$-process and its site. Their heavy element abundances are dominated by the influence of a single, or at most very few nucleosynthesis events. The first such object, {\sneden}, with [Fe/H] = $-3.1$ and [Eu/Fe] = 1.63, was discovered by \citet{Sneden1994}. At present, 12 r-II stars are known, and the fraction of r-II stars at [Fe/H] $< -2.5$ is estimated at the level of 5\,\%\ (Paper~II). Detailed abundance analysis of {\sneden} \citep{sneden1996} and another benchmark r-II star {\cayrel} \citep{hill2002} established the match of the stellar and solar $r$-process pattern in the Ba-Hf range suggesting that the $r$-process is universal. It produced its elements with the same proportions during the Galactic history. This conclusion was of fundamental importance for better understanding the nature of the $r$-process. Figure\,\ref{Fig:AbundancePattern} displays the heavy-element abundance patterns of four stars revealing the largest Eu enhancement, with [Eu/Fe] $\ge 1.5$. The data were taken from \citet[][{\sneden}]{Sneden2003}, \citet[][{\cayrel}]{2013A&A...550A.122S}, \citet[][HE\,1219-0312]{HE1219}, and \citet[][HE\,1523-091]{Sneden2008}. All these stars have very similar chemical abundance patterns in the Sr-Hf (probably Pt?) range suggesting a common origin of these elements in the classical $r$-process. Only single measurements or upper limits are available for the heavier elements up to Pb, not allowing any firm conclusion about their origin. Two of the four r-II stars have high abundances of Th, and they are referred to as actinide-boost stars. Figure\,\ref{Fig:AbundancePattern} displays also the SSr patterns from predictions of B14 and A99. Two sets of the solar $r$-abundances are consistent except the elements with significant contribution of the $s$-process to their solar abundances. This concerns, in particular, the light trans-Fe elements. For example, for Sr and Y, the $s$-process contribution exceeds 90\,\%. In such a case, the calculation of the $r$-residuals involves the subtraction of a large number from another large number, so that any small variation in one of them leads to a dramatic change in the difference. The uncertainty in the solar $r$-residuals does not allow to draw firm conclusions about any relation between the light trans-Fe elements in r-II stars and the solar $r$-process. For elements beyond Ba, the difference between SSr(B14) and SSr(A99) is notable for Ba, La, Ce, Ta, and Pb. The only measurement of stellar Ta \citep[][{\cayrel}]{2013A&A...550A.122S} favors the solar $r$-residual of B14. Due to the weakness of the lines of the lead in the optical wavelength region, Pb abundances have been measured only in very few metal-poor stars \citep[for a recent review, see][]{Roederer2009}. This paper focuses on stellar Ba and Ba/Eu abundance ratios. Our first concern is the line formation treatment. For cool stars, most abundance analyses are made under the assumption of local thermodynamic equilibrium (LTE), and Fig.\,\ref{Fig:halo} shows the LTE abundance ratios. In MP atmospheres, the departures from LTE can be significant due to a low number of electrons donated by metals, which results in low collision rates, and also due to low ultra-violet (UV) opacity, which results in high photoionization rates. Therefore, a non-local thermodynamic equilibrium (NLTE) line-formation modeling has to be undertaken. Each line of \ion{Ba}{ii} and \ion{Eu}{ii} consists of isotopic and hyper-fine splitting (HFS) components, and derived element abundances depend on which isotope mixture was applied in the calculations. Different $r$-process models predict different fractional abundances of the Ba isotopes, and different stellar Ba abundance analyses are based on different data on the Ba isotope mixture. We wish to evaluate systematic shifts in derived stellar Ba/Eu abundance ratios due to using different Ba isotope mixtures. \begin{table} % \caption{\label{Tab:ba_isotope} Solar System Ba and Eu isotope abundance fractions (\%) and $s$-process contributions (\%).} \centering \begin{tabular}{rrllcll} \hline\hline \noalign{\smallskip} & Total & \multicolumn{5}{c}{$s$-process} \\ \cline{3-7} \noalign{\smallskip} & & \multicolumn{2}{c}{1.5, 3~$M_\odot$ AGB models} & & \multicolumn{2}{c}{GCE calculations} \\ \noalign{\smallskip} \cline{3-4} \cline{6-7} \noalign{\smallskip} & L2009& A99 & B11 & & T99 & ~~B14 \\ \noalign{\smallskip} \hline \noalign{\smallskip} \iso{134}{Ba} & 2.4 & 100 & 100 & & ~~94 & ~~100 \\ \iso{135}{Ba} & 6.6 & ~~30 & ~~26 & & ~~22 & ~~~~28 \\ \iso{136}{Ba} & 7.9 & 100 & 100 & & ~~97 & ~~100 \\ \iso{137}{Ba} & 11.2 & ~~67 & ~~66 & & ~~58 & ~~~~63 \\ \iso{138}{Ba} & 71.7 & ~~94 & ~~86 & & ~~84 & ~~~~92 \\ \multicolumn{2}{l}{Total Ba} & 81$\pm6.7$ & 88.7$\pm5.3$ & & ~~80 & 85.2$\pm6.7$ \\ \noalign{\smallskip} \hline \noalign{\smallskip} \iso{151}{Eu} & 47.8 & ~~~6 & ~~~6 & & ~~~6 & ~~~~~6 \\ \iso{153}{Eu} & 52.2 & ~~~5 & ~~~6 & & ~~~5 & ~~~~~6 \\ \multicolumn{2}{l}{Total Eu} & ~~~6$\pm7$ & ~~~6$\pm0.3$ & & ~~~6 & ~~~6$\pm0.4$ \\ \noalign{\smallskip}\hline \noalign{\smallskip} \multicolumn{7}{l}{L2009 = \citet{Lodders2009}, } \\ \multicolumn{7}{l}{A99 = \citet{Arlandini1999}, B11 = \citet{Bisterzo2011},} \\ \multicolumn{7}{l}{T99 = \citet{Travaglio1999}, B14 = \citet{2014arXiv1403.1764B}.} \\ \end{tabular} \end{table} \begin{figure} % \resizebox{88mm}{!}{\includegraphics{baeu_halo_update2.ps}} \caption{\label{Fig:halo} The Ba/Eu abundance ratios of the r-II (filled circles), r-I (open rombs), and Eu-poor (asterisks) stars (for the sources of the data, see text). The error bars were computed as $\sigma_{\rm Ba/Eu} = \sqrt{\sigma_{\rm Ba}^2 + \sigma_{\rm Eu}^2}$, where the abundance errors available. The continuous and long-dashed lines indicate the SSr ratios, as predicted by GCE calculations of \citet{2014arXiv1403.1764B} and \citet{Travaglio1999}, respectively, while the short-dashed and dash-dotted lines correspond to the SSr of \citet{Arlandini1999} and \citet{Bisterzo2011}, respectively. The dotted line corresponds to the Solar System ratio \citep{Lodders2009}.} \end{figure} \begin{figure} % \centering \resizebox{88mm}{!}{\includegraphics{r2_rproc_update.ps}} \caption{\label{Fig:AbundancePattern} The heavy-element abundance patterns of {\sneden} (filled rombs) and the benchmark r-II stars CS\,31082-001, HE\,1523-0901, and HE\,1219-0312 (open triangles). The element abundances have been scaled to match Eu--Tm in {\sneden}. The dashed and continuous curves indicate the SSr abundance patterns calculated using the $s$-process predictions of A99 and B14, respectively. The dotted curve corresponds to the HEW $r$-process model of \citet{Farouqi2010}. } \end{figure} This paper is structured as follows. In Sect.\,\ref{Sect:nlte}, we describe NLTE calculations for \ion{Ba}{ii} and \ion{Eu}{ii} in VMP atmospheres and evaluate the HFS effects on Ba abundance determinations. Abundances of Ba and Eu of the selected r-II stars are revised in Sect.\,\ref{Sect:stars}. Section\,\ref{Sect:Conclusions} summarizes the obtained results. | \label{Sect:Conclusions} In this study, the Ba and Eu abundances of the eight r-II stars were revised by taking departures from LTE for lines of \ion{Ba}{ii} and \ion{Eu}{ii} into account, and accounting for HFS affecting the \ion{Ba}{ii} resonance lines with a common Ba isotope mixture. For most of the r-II stars, NLTE leads to a lower Ba, but a higher Eu abundance. Therefore, the Ba/Eu abundance ratios decrease on average by 0.21~dex when moving from LTE to NLTE We conclude that an adequate line-formation modelling for heavy elements is important for abundance comparisons between VMP stars, and in particular, giants. For stellar Ba abundance determinations, we recommend to use the subordinate lines of \ion{Ba}{ii}, which are nearly free of the HFS effects. At present, there is no consensus on the Ba isotope abundance fractions in the $r$-process, and abundances derived from the \ion{Ba}{ii} resonance lines vary by 0.15~dex, when applying different $r$-process models. We suspect that the total fractional abundance of the odd-A isotopes of Ba in old Galactic stars is related to the $r$-process abundances of the star. The observational evidence indicates that all the stars with high $f_\mathrm{odd}$-values, i.e., HD\,103095, HD\,84937, {\LyudmilasStar}, and the selected thick disk stars, are $r$-process enhanced, with [Eu/Fe] = 0.24 to 0.70 \citep{Mashonkina2006,Mashonkina2008,HE2327}. In contrast, both stars with low $f_\mathrm{odd}$, i.e., HD\,122563 and HD\,140283, are Eu-poor, with [Eu/Fe] = $-0.51$ \citep[][NLTE]{Mashonkina2008} and [Eu/Fe] $< -0.2$ \citep{Gallagher2010}, respectively. With the improved Ba and Eu abundances of the r-II stars, we constrain a pure $r$-process Ba/Eu abundance ratio to be log(Ba/Eu)$_r$ = 0.78$\pm$0.06. The obtained results support the solar $r$-residuals based on the chemical evolution calculations of \citet{Bisterzo2011} and \citet{2014arXiv1403.1764B}, and also the HEW $r$-process model by \citet{Farouqi2010}. For further constraining the $r$-process models, it would be important to determine Ba isotopic fractions of the r-II stars. This is a challenge for theory as well as observations. | 14 | 4 | 1404.0511 |
1404 | 1404.2452_arXiv.txt | New and complete multi-band light cur-ves of the oEA stars QY~Aql, BW~Del, TZ~Dra, BO~Her and RR~Lep were obtained and analysed with the Wilson-Devinney code. The light curves residuals were further analysed with the Fourier method in order to derive the pulsation characteristics of the oscillating components. All the reliable observed times of minimum light were used to examine orbital period irregularities. The orbital period analyses revealed secular changes for QY~Aql and BW~Del, while the Light-Time Effect seems to be the best explanation for the cyclic period changes in TZ~Dra and BO~Her. RR~Lep has a rather steady orbital period. Light curve solutions provided the means to calculate the absolute parameters of the components of the systems, which subsequently were used to make an estimate of their present evolutionary status. | \label{INTRO} Generally, eclipsing binary systems (hereafter EBs) offer unique information for the calculation of stellar absolute parameters and evolutionary status. Especially, the cases of binaries with $\delta$~Scuti components are extremely interesting, since they provide additional information (i.e. pulsation characteristics) for this part of the stellar lifetime. It has been shown that the $\delta$~Scuti stars in classical Algols (oEA stars) show difference in their pulsational characteristics from time to time due to mass gain \citep{MK04,MK07}. Therefore, the calculation of their absolute parameters and the identification of their oscillating characteristics help us to obtain useful conclusions for this `unstable' part of stellar lifetime. \citet{SO06a} and \citet{LI12} found a connection between orbital and pulsation periods in these systems and showed that binarity plays an essential role in the evolution of the components. Moreover, \citet{LI12} reported an empirical relation between evolutionary stage and dominant pulsation frequency of the $\delta$~Scuti stars in binaries, which differs significantly from that for single ones. Seventy four binaries with $\delta$~Scuti components have been discovered so far \citep{LI12}, but their number is increasing with a rapid rate. The present work is the continuation of the survey for candidate EBs with pulsating components \citep{LN09,LN12,LI12}. According to the Observed$-$Calculated times of minima variations (hereafter O$-$C) analysis, it is feasible to detect which physical mechanisms play a role in the period modulation (e.g. third body existence, mass transfer between the components) of a binary. On the other hand, from the light curve (hereafter LC) analysis it is possible to determine the Roche geometry of the EB (i.e. semi-detached, detached or contact configuration) or detect a third light. The solutions of these analyses are obviously qualitatively and in some cases also quantitatively (e.g. existence of a third body) connected, even though they are based on different methods. Five eclipsing systems candidate to include a $\delta$~Sct component, namely V345~Cyg, BW~Del, MX~Her, TW~Lac and AQ~Tau, were selected from the lists of \citet{SO06b} in order to check them for any possible pulsational behaviour. The results showed that only BW~Del exhibits pulsations, therefore systematic observations were performed in order to obtain its complete LCs. Moreover, short-periodic pulsations in the system LT~Her were suspected by Dr. Mkrtichian (private communication), based on his unpublished photometric observations, who kindly suggested us the system for further photometric observations. Our preliminary photometric analysis indeed confirmed the oscillating nature of the primary component and the results will be presented in a future work. Finally, the present work presents results for BW~Del and for other four confirmed cases of oEA stars, namely QY~Aql, TZ~Dra, BO~Her and RR~Lep, for which systematic observations were also made. For these five cases of oEA stars, initially, we performed LC analysis in order to determine their geometric and absolute parameters. Subsequently, frequency analysis on their LC residuals was made with the aim to reveal their main pulsational properties. In addition, since all systems, except for RR~Lep, present orbital period modulations, their O$-$C diagrams were also analysed. Finally, combining the derived information from the pre-mentioned analyses we obtained a more comprehensive view of these systems. The motivation for the present work was: (a) the lack of accurate and/or modern observations for these systems, especially in multiple bands, (b) their poor coverage of their LCs, (c) the lack of accurate pulsation characteristics and (d) the lack of interpretation of their orbital period changes. \textbf{QY~Aql}: The system has an orbital period of $\sim7.22956^{\rm d}$. The radial velocities of its primary component and the mass function of the system were calculated by \citet{ST46} and recalculated by \citet{LS71} who found $K_1$=36~km/s and $f(m)=0.035$~M$_{\odot}$, respectively. \citet{GM81}, based on the photographic LCs of \citet{WH45,WH48}, published revised photometric elements of the system but they disputed the past results of $K_1$ and $f(m)$, since both members turned out to be extremely massive. The spectral type of the system is F0 \citep[cf.][]{BU04,MA06}. Modern measurements of the system are given by the \textsl{ASAS} project \citep{PO05}, but they contain only a few points which do not cover the whole LCs. Finally, the pulsational behaviour of its primary was reported by \citet{LI12}. \textbf{BW~Del}: This EB ($P\sim2.42313^{\rm d}$) was generally neglected. The only available measurement concerns its F2 spectral type \citep[cf.][]{HA04,SK10}. \textbf{TZ~Dra}: The spectral type of this system is A7V \citep{HE60} and its period has a value $P\sim0.86603^{\rm d}$. \citet{RR90} reported for the first time that the period of the system is changing. \citet{RO03} noticed that small light variations, that can be connected either with spot activity or pulsations, occur in the system. A few years later, \citet{RO05}, \citet{MK05} and \citet{MK06} found $\delta$~Sct-type pulsations in the primary component with a pulsation period of $\sim28^{\rm min}$. \textbf{BO~Her}: The orbital period of this eclipsing pair is $\sim4.27283^{\rm d}$ and its spectral type is A7V \citep{HA84}. The primary's component oscillating nature was reported by \citet{SB07}, who found a dominant pulsation period of 1.7871$^{\rm hr}$. \textbf{RR~Lep}: The system has a period of $\sim0.91543^{\rm d}$. Photoelectric LCs were given by \citet{BO86}, \citet{AV89}, \citet{VA89} and \citet{SA89}. \citet{SA89} detected a light variation of $\sim45$~min, but they did not interpret it as a possible pulsation. A CCD LC of the system in $V$-filter was published by the \textsl{ASAS} project \citep{PO05}, but it is of low quality (i.e. small number of points, large photometric error and incomplete LC). \citet{DV09}, based on his CCD observations, found that pulsations occur in the system with a dominant frequency of 31.87~c/d ($\sim45$~min). The spectral type of the system has not been defined so far and ranges between A0-A7 in several catalogues and works \citep[cf.][]{PGC52,MA06,FA02,WR03,SS04}. The absolute parameters of all systems were calculated by \citet{BD80} (except for QY~Aql), based on the photometric parallax method, and \citet{SK90}, who used statistic relations (e.g. mass-radius, mass-luminosity). \begin{table} \centering \caption{Observations log of all observed systems.} \label{tab1} \scalebox{0.78}{ \begin{tabular}{l ccc c ccc} \tableline System & $m_{\rm min}^{\rm a}$& $S.T.$ & $F$ & $N$ & $hrs$ &$f_{\rm dom}$ & $Inst$ \\ & (mag) & & & & & (c/d) & \\ \tableline QY Aql & 11.4 & F0$^{\rm b}$& $BVI$& 36 & 211 & 10.656 & $K$\&$At$ \\ V345 Cyg & 11.3 & A1$^{\rm b}$& $B$ & 2 & 9 & -- & $K$ \\ BW Del & 11.4 & F2$^{\rm c}$& $BV$& 18 & 86 & 25.100 & $K$\&$At$ \\ TZ Dra & 9.6 & A7$^{\rm b}$& $BV$& 6 & 33 & 50.993 & $At$ \\ BO Her & 10.7 & A7$^{\rm d}$& $BVI$& 25 & 125 & 13.430 & $At$ \\ LT Her & 10.7 & A2$^{\rm b}$& $BV$ & 5 & 20 & 30.521 & $At$ \\ MX Her & 11.4 & F5$^{\rm b}$& $B$ & 2 & 12 & -- & $K$\&$At$ \\ TW Lac & 11.5 & A2$^{\rm c}$& $B$ & 2 & 10 & -- & $At$ \\ RR Lep & 10.2 & A7$^{\rm c}$& $BV$ & 9 & 30 & 33.280 & $At$ \\ AQ Tau & 12.0 & A5$^{\rm b}$& $B$ & 1 & 4.5 & -- & $At$ \\ \tableline \multicolumn{8}{l}{$^{\rm a}$\citet{WA06}, $^{\rm b}$\citet{MA06}, $^{\rm c}$\citet{SA11},}\\ \multicolumn{8}{l}{$^{\rm d}$\citet{HA84}} \end{tabular}} \end{table} \begin{table*}[t] \centering \caption{Detailed observations log of systems with a pulsating component.} \label{tab2} \scalebox{0.95}{ \begin{tabular}{l cccc cccc} \tableline System & Nights & Obs. dates & $T.S.$ &\multicolumn{3}{c}{Number of points/$sd$} & Comparison & $m_{\rm V}$ \\ \cline{5-7} & spent & & (d) & $B$ & $V$ & $I$ & stars & (mag) \\ \tableline QY Aql & 36 & 28/06-15/09 & 79 & 3255/3.8 & 3136/3.4 & 3159/3.2 &$C$: TYC 1618-1286-1 & 11.3$^{\rm a}$ \\ & & of 2011 & & & & &$K$: TYC 1618-0790-1 & 11.0$^{\rm b}$ \\ BW Del & 18 & 01/09-26/10 & 55 & 1791/3.8 & 1760/4.5 & -- &$C$: TYC 1635-1273-1 & 11.4$^{\rm a}$ \\ & & of 2011 & & & & &$K$: TYC 1635-1027-2 & 11.4$^{\rm a}$ \\ TZ Dra & 6 & 02/07-20/07 & 19 & 2108/4.3 & 2107/3.5 & -- &$C$: TYC 3529-0198-1 & 9.51$^{\rm a}$ \\ & & of 2008 & & & & &$K$: TYC 3529-0039-1 & 11.6$^{\rm a}$ \\ BO Her & 25 & 28/05-06/07 & 39 & 1992/3.8 & 1920/3.5 &1881/3.3 &$C$: TYC 2111-0124-1 & 11.3$^{\rm a}$ \\ & & of 2011 & & & & &$K$: TYC 2111-0128-1 & 12.3$^{\rm a}$ \\ RR Lep & 9 & 17/01-16/03 & 59 & 1035/2.8 & 991/2.7 & -- &$C$: TYC 5342-0022-1 & 9.6$^{\rm a}$ \\ & & of 2012 & & & & &$K$: TYC 5342-0128-1 & 10.4$^{\rm a}$ \\ \tableline \multicolumn{9}{l}{$^{\rm a}$\citet{HO00}, $^{\rm b}$\citet{HO98}} \end{tabular}} \end{table*} | \label{DIS} One newly discovered (BW~Del), and four already known (QY~Aql, TZ~Dra, BO~Her and RR~Lep) eclipsing systems with a $\delta$~Sct component were observed and analysed using modern analysis tools in order to obtain useful conclusions about their oscillating behaviour, geometrical shape, absolute parameters, evolutionary stage and orbital period modulations. These systems are confirmed as classical Algols with their primaries showing $\delta$~Sct type pulsations. Therefore, according to the definition given by \citet{MK04}, they can also be considered as oEA systems. LT~Her was also identified as an EB including a $\delta$~Sct type member, but the detailed results will be presented in the future. Four other EBs, candidates for including $\delta$~Sct components, namely V345~Cyg, MX~Her, TW~Lac and AQ~Tau, were also checked for pulsations but the results were negative. The primary component of QY~Aql is located beyond the TAMS and pulsates with a frequency of $\sim10.656$~c/d. The decreasing orbital period rate of the system is well explained with the mass transfer process from the secondary to the primary component, which is supported by its conventional semi-detached geometrical status, and the mass loss due to magnetic braking of its secondary, which was found to be at the giant stage of evolution. A mass loss rate of $3.2\times10^{-8}$~M$_{\odot}$/yr, typical for red giants \citep{HI01}, was estimated. Three pulsational frequencies were detected for the primary of BW~Del, with the most dominant one at 25.1~c/d. Based on the adopted mass and the derived radius, the star is located beyond the TAMS, but very close to it. The secondary component was found to be very evolved and it transfers material to the primary with a rate of $5\times10^{-8}$~M$_{\odot}$/yr. The primary component of TZ~Dra is a relatively fast pulsator with a frequency of $\sim50.99$~c/d and is located near the ZAMS. The frequency analysis results are in agreement with those of \citet{MK05}. Based on its frequency value and its evolutionary status we conclude that its oscillating lifetime must have started recently, according to the evolutionary stage-pulsation period empirical relation for this kind of stars \citep{LI12}. The secondary component of the system has filled its Roche lobe and is located beyond the TAMS. The cyclic changes of the system's period are caused probably due to a tertiary component with a period of $\sim62$~yr and a minimal mass of $\sim0.3~M_{\odot}$. On the other hand, the LC analysis did not reveal any third light. However, assuming that the third body is a MS star, and based to the mass-luminosity relation for dwarfs ($L\sim M^{3.5}$), we can calculate its luminosity and compare it with the absolute luminosity values of the binary's members (see Table~\ref{tab3}) by using the following formula: \begin{equation} L_{3,{\rm O-C}} (\%)=100 \frac{M_{3,{\rm min}}^{3.5}}{L_1+L_2+M_{3,{\rm min}}^{3.5}} \end{equation} We found that the expected luminosity contribution of such a third star should be $\sim$0.14\%, hence its light absence is plausible. However, according to the value of $\Delta Q$ ($\sim10^{50}~$g~cm$^2$), it is possible that the period changes can be caused due to magnetic influence of the secondary component. Applegate's mechanism predicts also brightness changes of the system, but this has not been verified so far. Therefore, future photometric observations covering several decades and/or astrometric observations are needed in order to conclude about the mechanism that forms the binary's orbital period. For the oscillating member of BO~Her we traced two pulsation frequencies with the dominant one at 13.43~c/d. This results agrees with that of \citet{SB07}. Due to the relatively high amplitude of this mode, its first two harmonics of its pulsation frequency were also detected in the frequency spectrum. The primary (pulsating) component of the system is located on the TAMS edge. On the other hand, the secondary is located far beyond the TAMS, being at the giant stage of evolution. A third body with a minimal mass of $\sim1.1~M_{\odot}$ and a period of $\sim31$~yr may exist around the EB, but we did not detect any additional luminosity in the LC analysis. Following the same method as for the case of TZ~Dra, we found an expected luminosity contribution $\sim$5\%, which is large enough to be detected photometrically. The most possible explanation for this disagreement could be either the non-MS nature of the third body (e.g. exotic object) or that the third body is in fact a binary with two low-mass and low-temperature components, providing lower luminosity in total, instead of a single star. Future spectroscopic and/or astrometric observations are desirable in order to solve this mystery. The primary component of RR~Lep pulsates in two modes with the dominant frequency at $\sim33.28$~c/d and it is located on the MS and very close to the TAMS. The present results regarding the frequency $f_1$ are in marginal agreement with those of \citet{DV09}, who found only one pulsation frequency of $\sim31.87$~c/d. However, our results are based on two-filter data which were obtained with better equipment in comparison with that used by \citet{DV09}. The secondary of the system is a rather more evolved star located above the TAMS. The O$-$C points distributions of TZ~Dra, BO~Her and RR~Lep do not show any secular period changes that can be connected with mass transfer. Very probably, these systems are at slow mass-accretion stage \citep{MK03} with a rate that cannot be detected with the current time coverage of minima timings. \citet{LI12}, based on the pulsational and absolute parameters of the $\delta$~Sct components of all known oEA stars, derived empirical relations between the dominant pulsation period $P_{\rm puls}$ and the gravity acceleration value $g$ and between the $P_{\rm puls}$ and the orbital period $P_{\rm orb}$ of the systems. Therefore, it is useful to check if the respective values of the pulsating components of the systems analysed herein follow these trends, and this is shown in Fig.~6. \begin{figure}[h] \centering \begin{tabular}{c} \includegraphics[width=7.8cm]{GP.eps}\\ \includegraphics[width=7.8cm]{PP.eps}\\ \end{tabular} \label{fig6} \caption{Positions of the pulsating members of the systems in the $g-P_{\rm puls}$ (upper) and $P_{\rm puls}-P_{\rm orb}$ (lower) diagrams for oEA stars with $\delta$~Sct members \citep{LI12}. The error bars of the periods' values are not shown due to scale reasons.} \end{figure} The pulsating stars of all systems seem to follow well the $P_{\rm puls}-P_{\rm orb}$ and $g-P_{\rm puls}$ trends, with the exception of the primary of QY~Aql in the $g-P_{\rm puls}$ diagram. This star is at the subgiant evolutionary stage. On the other hand, the sample, in which the empirical relation of $g-P_{\rm puls}$ is based, consists mostly of MS stars. Therefore, QY~Aql, with the longest orbital period and minimum $g$-value in the sample of \citet{LI12}, might have followed a different evolutionary track or another relation between $g-P_{\rm puls}$ for the evolved oEA stars has to be examined. Radial velocities measurements for all systems are needed in order to determine their absolute parameters and the $l$-degrees of their pulsation modes with higher certainty. Moreover, more precise photometric observations (e.g. space data) are expected to reveal additional pulsation frequencies that they could not be detected with the present instrumentation setup. Future surveys aiming to new discoveries of this kind of systems and long-term monitoring of the already known ones are highly encouraged in order to enrich our knowledge about the mass transfer implication in the pulsation mechanisms and, in general, about the stellar evolution of binaries with A-F components. | 14 | 4 | 1404.2452 |
1404 | 1404.1876_arXiv.txt | We describe an interferometric reflectometer method for passive detection of subsurface oceans and liquid water in Jovian icy moons using Jupiter's decametric radio emission (DAM). The DAM flux density exceeds 3,000 times the galactic background in the neighborhood of the Jovian icy moons, providing a signal that could be used for passive radio sounding. An instrument located between the icy moon and Jupiter could sample the DAM emission along with its echoes reflected in the ice layer of the target moon. Cross-correlating the direct emission with the echoes would provide a measurement of the ice shell thickness along with its dielectric properties. The interferometric reflectometer provides a simple solution to sub-Jovian radio sounding of ice shells that is complementary to ice penetrating radar measurements better suited to measurements in the anti-Jovian hemisphere that shadows Jupiter's strong decametric emission. The passive nature of this technique also serves as risk reduction in case of radar transmitter failure. The interferometric reflectometer could operate with electrically short antennas, thus extending ice depth measurements to lower frequencies, and potentially providing a deeper view into the ice shells of Jovian moons. | Subsurface oceans in Jupiter's icy moons could provide a present-day setting for extra-terrestrial life within our Solar System. Of the three Jovian icy moons, Europa is favored as having the greatest potential to sustain life, based on strong evidence for a persistent ocean directly in contact with rock. Galileo radio science measurements indicate Europa is differentiated, with a low density water-rich layer between 80 and 170 km thick (Anderson et al., 1998, Carr et al., 1998). Galileo magnetometry provides compelling evidence for a present-day ocean through the induced magnetic field (Kivelson et al., 2000, Zimmer et al., 2000). Estimates of the ice shell thickness of Europa are uncertain. Thermal models of the ice shell of Europa predict a thickness of $\leq$30~km (Ojakangas and Stevenson, 1989). Studies of Galileo spacecraft data have resulted in contradictory constraints for the ice shell thickness. Analysis of the Galileo magnetometer-derived oceanic conductivities, combined with radio Doppler data-derived interior models and laboratory conductivity vs concentration data, constrain the ice thickness to be $<$15~km with a best fit value of $\sim4$~km (Hand and Chyba, 2007). Galileo imaging of pits, domes, and dark spots provide an ice shell thickness constraint of 3-10~km (Pappalardo et al., 1998). Crater analyses, also obtained from Galileo images, constrain Europa's ice thickness to $>$3~km (Turtle et al., 2001) based on the need to isostatically support central features, and to at least 19-25~km thick from the thermal state inferred from depth-size relationships (Schenk, 2002). The most promising technique for direct detection of subsurface oceans in Jovian icy moons is ice-penetrating radar (IPR). A dual-frequency system, such as that described by Bruzzone et al., 2011, is capable of providing high-resolution images at shallow depths ($<$5~km) and characterize the depth of the ice up to 30~km with 100~m resolution. Unambiguous observation of a subsurface ocean demands that the detection technique have as high depth sensitivity as possible. To achieve this, the use of low frequencies ($<$30~MHz) has been proposed (Bruzzone et al, 2011). The main challenges involved with IPR are surface clutter and radio absorption of the ice, which can be reduced by use of low frequencies. However, the radio loud environment of Jupiter at frequencies $<40$~MHz requires a relatively strong transmitter. We explore a passive interferometric reflectometry technique that makes use of Jupiter's decametric (DAM) radio emission in the 1-40 MHz band. We argue that the DAM background that interferes low frequency IPR can be used as a source of ice depth sounding. This technique could be an attractive complement to a radar system because it can share the same dipole antenna and requires very low power passive components. Interferometric reflectometry could also extend the frequency band of observation to lower frequencies by operating as an electrically short dipole, further increasing the sensitivity to deep subsurface oceans. A passive measurement system could also serve as a backup to IPR in case of transmitter failure, thereby reducing the risks associated with the instrument. Interferometric reflectometry was first applied in the Dover Heights radio astronomical observatory in the 1940's (Bolton, 1982). In that setup, an antenna placed on a cliff observed both the direct emission of a radio source and its reflection on the sea surface. The signal was autocorrelated forming a virtual two-element interferometer. The baseline formed by the sea surface reflection provided one of the first demonstrations of radio emission from discrete sources (Bolton and Stanley, 1948) along with the first identification of cosmic radio sources including Centaurus A and the Crab Nebula (Bolton 1948). It is worth mentioning that this technique was born out of limited resources, not unlike the case for deep space probes. The interferometric reflectometry technique is currently applied in the measurement of snow depth using GPS signals (Larson et al., 2008, Gutmann et al., 2012). The interference between the GPS signal and its subsurface reflections modulates the signal to noise ratio with a sinusoidal wave whose frequency is directly proportional the snow depth (Larson et al., 2008). The technique has been successfully demonstrated and validated by comparison with other measurements (Gutmann et al., 2012). The geometry for the application of interferometric reflectometry to Jovian moon ice depth measurements is shown in Figure~\ref{fig:JANINE_concept}. Jupiter's radio emission arrives from a distance of $\gtrsim6\times10^{8}$~km to the vicinity of an icy moon. At the sub-Jovian point, where the spacecraft lies directly between Jupiter and the icy moon, an antenna receiver system records a sample of the decametric radio emission. The same emission strikes the surface of the icy moon and its echoes arrive at the spacecraft at a later time ($\sim$1~ms). The antenna beam pointed at the icy moon samples the echoed radio emission, which is cross-correlated with the direct emission to produce fringes. The cross-correlation peaks at delays corresponding to the moon surface and subsurface ice-water boundary reflection layers. The amplitudes of the cross-correlation peaks are related to the dielectric properties of the ice. Therefore, the cross-correlation has the potential to reveal the presence of sharp boundaries that would be associated with a subsurface ocean or liquid water deposits in the ice shell. \begin{figure*}[h!] \centering \includegraphics[width=0.9\linewidth]{./Figure_Concept_Dipole_BW.pdf} \caption{Passive detection of subsurface oceans in icy moons using Jupiter's radio emission and its echoes. The radio emission from Jupiter (shown as an arrow with dotted line) is sampled by the dipole antenna. The radio signal is then reflected from the surface of the icy moon (arrow with dashed-dotted line) as well as the subsurface ocean (arrow with dashed line). Both echoes are detected by the dipole antenna. The delays and amplitudes of the reflected signals are extracted by correlation with the direct emission.} \label{fig:JANINE_concept} % \end{figure*} In this paper, we will describe the physics of the passive interferometric reflectometer concept. In Section~2 we briefly review the properties of Jupiter's decametric radio emission. Section~3 gives a summary of the properties of Jovian moon ice shells. Section 4 describes the mathematical details of interferometric reflectometry and provides estimates for the sensitivity and resolution of the technique. Section 5 compares the expected sensitivity of interferometric reflectometry with ice penetrating radar. Section 6 summarizes our results and outlines the next steps in the development of this measurement technique. | We have provided the physical basis for a passive interferometric reflectometer that takes advantage of Jupiter's strong decametric emission to search for subsurface oceans in Jovian icy moons. We have shown that the absorptive properties of the ice could allow for enough of this signal to be reflected back to a spacecraft for passive observation of a subsurface ocean. The unambiguous detection of a subsurface ocean could be obtained with a relatively simple system consisting of a dipole antenna, a digitizer, and a correlator. The interferometric reflectometer concept could be used as a complementary system to an ice penetrating radar instrument by adding a passive device sharing the radar antenna. Radar provides its best measurements in the anti-Jovian side while the interferometric reflectometer works best in the sub-Jovian side. This is also a very low power system that could run in the background while other instruments are performing their measurements. There have been a number of simplifying assumptions in this work that will have to be studied in more depth. In particular, we have assumed that Jupiter's decametric radio emission behaves as white noise. Partial coherence in the bursts could improve the correlation compared to a white noise model. The autocorrelation behavior of the Jovian decametric emission can be constrained in a future study using data from low frequency arrays such as LOFAR and the LWA. Other features of Jovian decametric emission need to be included in the interferometric reflectometer measurement model. More detailed simulations using specific orbits and modeled behavior of the different components of the decametric emission and their spatio-temporal characteristics will be required. The Cyclotron Maser Instability model of decametric radiation (Treumann, 2006) claims that the decametric radiation originates in the poles of Jupiter and the emission propagates as a cone with wide opening angle but narrow width. The emission comes both from the north and south poles, which could potentially extend the region where interferometric reflectometry is applicable. Since some of the stronger emissions are sporadic, it also may be advisable to include a power meter that triggers the correlator when there is a sudden boost in the decametric flux density. Another possibility is that the interferometric reflectometer can work with an electrically short dipole. The dimensions of the antenna used for low frequency measurements can be a limiting factor in the design of an instrument. A resonant dipole antenna for frequencies around 3~MHz implies a dimension comparable to a half wavelength of $\lambda/2=50$~meters, which is impractical for a deep space probe. At a frequency of 30~MHz, a dipole antenna has dimension of $\lambda/2=5$~meters, which is more manageable. A 5~meter dipole operating as an electrically short antenna could extend observations of decametric radiation down to frequencies as low as 3~MHz. This would be particularly advantageous given the wide band over which decametric activity is observed and where the surface roughness and clutter effects are largely reduced. The extension to lower frequencies could provide significant improvements in sensitivity and should be studied in the future. The spatial, temporal, and spectral structure of the decametric emission needs to be studied in more detail for future instrument development. Many of the details left out in this first study could improve the estimates for this technique and open the way for a passive radio probe for Solar System exploration. | 14 | 4 | 1404.1876 |
1404 | 1404.4348_arXiv.txt | We present a comprehensive study of phase curves and secondary eclipses in the Kepler data set using all data from 16 quarters that were available in 2013-2014. Our sample consists of 20 confirmed planets with $R_p > 4 R_e$ , $P < 10d$, $V_{mag} < 15$. Here we derive their temperatures and albedos, with an eye towards constraining models for the formation and evolution of such planets. Where there was overlap our results confirm parameters derived by previous studies, whereas we present new results for Kepler 1b-8b, 12b-15b, 17b, 40b, 41b, 43b, 44b, 76b, 77b, and 412b derived in a consistent manner. We also present lightcurve analyses for Kepler 91b and Kepler 74b, which both show extra dimmings at times other than from the expected primary and secondary eclipses. Corrected for thermal emission we find most of the massive planets from our sample to be low in albedo ($<0.1$) with a few having higher albedo ($>0.1$). | \subsection{The \textsl{Kepler} mission} Studying extrasolar planets is one of the major frontiers of astronomy today. The field has transformed from simple identification to comprehensive categorization and characterization of exoplanets and exoplanetary systems. Analyses of data provided by NASA's \textsl{Kepler}\footnote{The data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST), which is managed by the Space Telescope Science Institute (STScI). STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support for MAST for non-HST data is provided by the NASA Science Mission Directorate via grant NNX13AC07G and by other grants and contracts. This paper includes data collected by the \textsl{Kepler} mission. Funding for the \textsl{Kepler} mission is also provided by the NASA Science Mission Directorate.} mission has revolutionized this field by compiling a statistically significant number of transiting planets and planetary candidates (e.g., \citealt{2010Sci...327..977B}; \citealt{2011ApJ...736...19B}; \citealt{2013ApJS..204...24B}; \citealt{2014ApJS..210...19B}; \citealt{2015ApJS..217...16R}; \citealt{2015ApJS..217...31M}). For example, \textsl{Kepler} data allowed researchers to discover Kepler 9b (\citealt{2010Sci...330...51H}), the first multi-planetary system outside our solar system; Kepler 10b (\citealt{2011ApJ...729...27B}), one of the first confirmed rocky planets outside the solar system; and Kepler 16b (\citealt{2011Sci...333.1602D}), the first circumbinary planet. More recently the \textsl{Kepler} team announced the discovery of potentially habitable worlds in the Kepler 62 (\citealt{2013Sci...340..587B}) and Kepler 69 (\citealt{2013ApJ...768..101B}) systems, and the near-Earth-sized planets Kepler-186f (\citealt{2014Sci...344..277Q}) and Kepler-452b (\citealt{2015AJ....150...56J}) in the habitable zones of their parent stars. Deeper analyses are possible using the exquisite \textsl{Kepler} data beyond merely detecting exoplanetary systems: researchers are now able to analyze large samples of planetary candidates to pin down occurrence rates such as $\eta_{earth}$ (e.g., \citealt{2012ApJS..201...15H}; \citealt{2013ApJ...767...95D}; \citealt{2013ApJ...766...81F}; \citealt{2013PNAS..11019273P}; \citealt{2015arXiv150604175B}), find non-transiting planets via transit timing variations (\citealt{2011ApJ...743..200B}), perform phase-curve analyses (\citealt{2013ApJ...771...26F}), and may eventually even be able to detect exomoons (\citealt{2013ApJ...770..101K}) and exotrojans (\citealt{2015arXiv150800427H}). For the close-in, and therefore hot, planets around bright, high-signal host stars in the \textsl{Kepler} data set, we are able to analyze secondary eclipses, i.e., the modulated flux from the star-planet system when the light (reflection and thermal emission) of the planet disappears during its passage behind the parent star. Differential measurements then help us to characterize physical parameters of the planet such as albedo and temperature. Measuring such quantities, together with complementary spectroscopic measurements, can provide contraints on models for the formation and atmospheric photochemistry of such close-in planets \citep[e.g.,][]{2010ApJ...717..496L, 2011ApJ...737...15M, 2011ApJ...738...72V} \begin{deluxetable}{lll} \tabletypesize{\scriptsize} \tablecaption{\textsl{Kepler} Quarters used in our analysis\label{tbl:used_quarters}} \tablewidth{0pt} \tablehead{ \colhead{KOI} & \colhead{LC Quarters} & \colhead{SC Quarters}} \startdata 1 &0,1,2,3,4,5,6,7,9,10,11,13,14,15 & 0,1,2,3,4,5,6,7,8,9,10,11,12,13,14,15 \\ 2 &0,1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 0,1,2,3,4,5,6,7,8,9,10,11,12,13,14,15 \\ 3 &0,1,2,3,4,5,6,8,9,10,12,13,14& 0,1,2,3,4,5,6,8,9,10,11,12,13,14 \\ 7 &0,1,2,3,4,5,6,7,9,10,11,13,15& 0,1,2,3,4,5,6,7,9,10,11,13,14,15 \\ 10 &0,1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 1,2,3,4,5,6,7,8,9,10,11,12,13 \\ 13 &0,1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 1,2,3,6,7,8,9,10,11,12,13,14,15 \\ 17 &0,1,2,3,4,5,6,8,9,10,12,13,14& 1,2,3,4,5,6,8,9,10,11,12 \\ 18 &0,1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 1,2,3,4,5,6,7,8,9,10,11,12 \\ 20 &0,1,2,3,4,5,6,7,9,10,11,13,14,15& 1,2,3,4,5,6,7,8,9,10,11 \\ 97 &0,1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 2,3,4,5,6,7,8 \\ 98 &0,1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 2,3,4,5,6,7,8,9,10,11,12 \\ 127 &1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 2,3,4,5,6,7 \\ 128 &1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 2,3,4,5,6,7 \\ 135 &1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 2,3,4,5,6,7,9,10,11,12,13,14,15 \\ 196 &1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 3,4,5,6,7 \\ 200 &1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 3,4,5,6,7 \\ 202 &1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 3,4,5,6,7 \\ 203 &1,2,3,4,5,6,8,9,10,12,13,14& 3,4,5,6,7,8,9,10,11,12,13,14 \\ 204 &1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& 3,4,5,6,7 \\ 428 &1,2,3,4,5,6,8,9,10,12,13,14& -- \\ 1658 &1,2,3,4,5,7,8,9,11,12,13,15& --\\ 2133 &0,1,2,3,4,5,6,7,8,9,10,11,12,13,14,15& -- \enddata \tablecomments{For our analysis we used all data available prior to data release Q16, August 2013} \end{deluxetable} \subsection{Transits and eclipses} Systems with transiting extrasolar planets can offer two important observational opportunities for deriving physical parameters of the planets. In \emph{primary transit} the planet crosses the star. From a broadband transit-lightcurve, in this case, one can measure the planetary radius $R_{p}$ in units of the stellar radius $R_{*}$. The depth of the transit is $\sim(R_{p}/R_{*})^{2}$, which for a Jupiter radius planet transiting a sun-like star, is of the order of $\sim 1\%$ (e.g., \citealt{2000ApJ...529L..41H}). If the geometry (inclination, eccentricity) is right the planet also disappears behind its host star in a so-called \emph{secondary eclipse}. For a 2000 K hot Jupiter-size planet, the typical flux deficit during secondary eclipse is $\sim 200$ ppm at $\sim 2 \mu m$ in the near-infrared and even larger at longer wavelengths, but considerably smaller at optical wavelengths (400-900 nm) at which \textsl{Kepler} observes. However, for host stars that are bright enough, \textsl{Kepler}'s outstanding sensitivity provides a direct measure of the planet's disk-averaged day side flux for some of its targets, particularly close-in gas giants -- so called 'Hot Jupiters' and 'Hot Neptunes'. Observing secondary eclipses combined with planetary phase curves can help us to characterize the planet and its atmosphere. For example, the depth of the secondary eclipse can constrain the albedo of the planet, while the timing and width of the secondary eclipse can help determine its orbital parameters. Comparing the amplitude of the reflected light in the phase curve with the depth of the secondary eclipse can constrain the day and night side temperatures, and therefore confirm the planetary nature of a candidate that is not self-luminous, and help to understand day to night side heat exchange. Measuring exoplanet eclipses, phase curves and albedo values yield information about the composition of their atmosphere, and day to night side temperature ratios yield the efficiency of energy transport and presence of possible temperature inversions (for details see, e.g., \citealt{2011ApJ...729...54C}). Several previous studies have focused on eclipses and phase curves in the \textsl{Kepler} database, either on select samples of objects (e.g., \citealt{2011ApJ...730...50K} 5 objects; \citealt{2012AJ....143...39C} 76 objects; \citealt{2013ApJ...772...51E} 8 objects) or for individual planets or candidates (e.g., \citealt{2012AandA...541A..56M}; \citealt{2013ApJ...764L..22M}). In the largest sample so far, \citealt{2012AJ....143...39C} modeled secondary eclipses and phase curves (only via a simple sinusoidal flux term applied to the lightcurves) of a uniform set of Hot Jupiter candidates from \textsl{Kepler} to derive albedos and thermal emission properties, and compared the results with stellar and planetary parameters. While our study is similar in scope to this study, we worked with much more data (15 quarters -- as available in August 2013 -- of data instead of 1), a slightly bigger radius range (4 Earth radii = 0.4 Jupiter radii, instead of 0.5 Jupiter radii as a lower limit), and a longer period range (10 days instead of 5 days). Furthermore we applied a fully physical model with all different phase curve components. Of the \citealt{2012AJ....143...39C} sample of 76 KOIs only 55 were successfully modeled with their methods. Of these 55 systems many still remain unconfirmed planetary candidates or turned out to be false postives (e.g., KOIs 102, 1419, 1459, 1541, 1543; http://exoplanetarchive.ipac.caltech.edu/, August 2015). For the confirmed planets there is significant overlap in our samples (KOIs 1, 2, 10, 13, 17, 18, 20, 97, 127, 128, 196, 202, 203, 204), but several planets in our study (KOIs 3, 7, 98, 135, 428 and 1658) were not covered by their analysis. Here we present initial results of a comprehensive and consistent study of secondary eclipses and phase curves using data from quarters 0 through 15 of \textsl{Kepler} lightcurves (see Table \ref{tbl:used_quarters}), that were available at the time of our analysis. In this paper we focus on the 20 confirmed (August 2013) planets in the sample of 489 Kepler Objects of Interest (KOI) with $R_p > 4 R_e$, $P < 10d$, and $V_{mag} < 15$. Consistent measurements of exoplanet phase curves not only allows us to break many current degeneracies in modeling the thermal and chemical structure of separate exoplanet atmospheres, but also enables us to compare results across the whole set of analyzed systems in a comprehensive way. Analyses like this will also help us prepare for future ground- and space-based facilities that increase the number of exoplanetary systems and the wavelength range of precision observations. | With our consistent analysis we were able to confirm and in most cases improve parameters derived by previous studies. We present new results for Kepler 1b-8b, 12b-15b, 17b, 40b, 41b, 43b, 44b, 76b, 77b, and 412b. \subsection{Comparison to other fitting routines - future improvements} For the cases of previously analyzed targets we were able to confirm (within less than 2 $\sigma$ for almost all cases) results derived from various previous or parallel analyses that used very different modeling approaches, from relatively simple boxcar fits (without a phase curve model) of only the secondary eclipse to very sophisticated Bayesian codes fitting all system parameters in an integrated way. The fact that we reproduce most of these results demonstrates the value of our compromise approach to combine a state-of-the-art phasecurve model including all important contributions with a robust least-squares fit to trade off between number of systems and computing time. Also our goal was to focus on eclipses and phasecurves while fixing all other parameters to previously derived values. This significantly reduces the number of potential degeneracies that usually call for these more elaborated methods. However, in order to increase the statistical relevance of our results we are currently working to extend the analysis to the whole sample of 489 Kepler Objects of Interest with $R_p > 4 R_e$ , $P < 10d$, $V_{mag} < 15$: we plan to apply EXONEST \citep{2014ApJ...795..112P}, a Bayesian model selection algorithm to the whole set of 489 candidates. With a sample of that size we hope to find statistically significant correlations of stellar and planetary parameters with the position of the planet, e.g., in albedo $vs$ incoming flux phase space (see Figure \ref{fig:corr_star}). \subsection{Correlations with system parameters} For the following analysis of a potential correlation of the albedo with stellar or planetary parameters (see e.g. Figure \ref{fig:corr_star}) we used the albedo values, that were corrected for thermal emission (assuming no redistribution; $f_{dist}=1/2$) that are shown in Table \ref{tab:corralb} (center). Our results confirm the general trend of relatively low albedos for most of the Hot Jupiters, but we also show outliers with higher albedos. We see no significant correlations in our data. Neither the stellar parameters ($[Fe/H]$ and $log(g)$, see Figure \ref{fig:corr_star}) nor the planetary characteristics (mass, radius, density and surface gravity) correlate with the derived parameters. When excluding the planets with large errors in the albedo, there are indications that massive planets, very dense and very bloated planets tend to be low in albedo -- i.e., density extremes produce low albedos. Even though we present a relatively large sample characterized in this comprehensive and consistent manner, our sample size is still too small to draw significant statistical conclusions. Future efforts will include analyzing all of the candidates as well as the complete Kepler dataset of 18 available quarters (Q0-Q17). With \textsl{TESS} \citep{2010AAS...21545006R} and \textsl{PLATO} \citep{2013arXiv1310.0696R,2015arXiv150303251H} on the horizon the future will bear an even bigger data set, marking great potential for characterizations in a similar way. Such future observations and analyses will include many more planets in a similar range of planetary and stellar parameters. A comparative analysis beyond their basic parameters of a large number of planets orbiting a variety of host stars will probe and eventually solve the fundamental underlying questions on planet formation, migration and composition. \begin{figure}[t!] \centering \includegraphics*[width=\textwidth]{f9_new.eps} \caption{Emission-corrected geometric albedo $A_{g,c}$ for $(f=1/2)$ versus the incident stellar flux for our sample of Kepler giant planets (compare to Fig. 1 in \citealt{2013ApJ...777..100H}). The colors represent the host star's metallicity [Fe/H] (see colorbar) and the size of the symbol corresponds to the host's surface gravity $log(g)$ (see legend). There is no obvious correlation with the distribution. The dashed line represents Jupiters's albedo of 0.52 for comparison; all of the Hot Jupiters in our study have lower albedos.} \label{fig:corr_star} \end{figure} | 14 | 4 | 1404.4348 |
1404 | 1404.4381_arXiv.txt | We apply a combination of N-body modeling techniques and automated data fitting with Monte Carlo Markov Chain uncertainty analysis of Keplerian orbital models to radial velocity data to determine long term stability of the planetary system GJ~581. We find that while there are stability concerns with the 4-planet model as published by \cite{forveille2011}, when uncertainties in the system are accounted for, particularly stellar jitter, the hypothesis that the 4-planet model is gravitationally unstable is not statistically significant. Additionally, the system including proposed planet g by \cite{vogt2012} also shows some stability concerns when eccentricities are allowed to float in the orbital fit, yet when uncertainties are included in the analysis the system including planet g also can not be proven to be unstable. We present revised reduced chi-squared values for Keplerian astrocentric orbital fits assuming 4-planet and 5-planet models for GJ~581 under the condition that best fits must be stable, and find no distinguishable difference by including planet g in the model. Additionally we present revised orbital element estimates for each assuming uncertainties due to stellar jitter under the constraint of the system being gravitationally stable. | In the last few years, results from the {\it Kepler} mission and radial velocity surveys have shown that small planets are significantly more common than giant planets. We are gradually discovering the properties and occurrence rates of the former population. Of particular interest are planets around M dwarfs because they contribute to our knowledge of planetary systems beyond those orbiting solar-type stars. The proximity of the habitable zone to the host star and the low stellar masses significantly facilitate the detection of lower mass, potentially habitable planets with the radial velocity method. GJ~581 is the first multiple system with all known planets having minimum masses smaller than that of Neptune, and its history as portrayed by the extensive publication record is rather tumultuous. This is in part due to the fact that although the system was discovered over eight years ago, the masses and orbital properties of some of the GJ~581 planets are still not well determined. The first detection of a planet orbiting GJ~581 with a period of 5.4 days (GJ~581~b) dates to 2005 \citep{bonfils2005}. Subsequent studies revealed planets with periods of 12.9 days (GJ~581~c) and 83.4 days (GJ~581~d, later revised to 66 days \citep{udry2007}) as well as a lower mass planet at 3.15 days (GJ~581~e; \citep{mayor2009}). Shortly after, \cite{vogt2010} reported the discovery of 2 additional planets at periods of 433 days (GJ~581~f) and 36.6 days (GJ~581~g). The existence of planets f and g has been the subject of some debate. \cite{forveille2011} (hereafter F2011) submitted a paper including new HARPS data that disputes this finding. \cite{vogt2012} (hereafter V2012) have analyzed F2011's claims, and have found that the reduced chi-squared values as published by F2011 cannot be reproduced under assumptions of either astrocentric or Jacobi coordinates without removal of at least 5 high-residual data points, and that the orbital properties as published by F2011 are dynamically unstable on a timescale of hundreds of years, due to the high reported eccentricity of planet e. V2012 claimed that due to the uncertainty in the determination of the eccentricities for this system, it was better modeled with circular orbits to address the issue of dynamical stability. Their justification is the principle that finding a lower overall reduced chi-squared value with fewer total fit parameters was favored by the principle of parsimony. Their approach also allows probing for a possible planet masked by an artificially high eccentricity detection of planet d. \cite{kress2013} found that relaxing the eccentricity of planet e to a lower yet still significant value could bring the system into dynamical stability. These findings were also reported by \cite{toth2013}, who studied key chaotic indicators and constrained the eccentricity of planet e to less than $0.2$. They also found that a 5 planet system including planet g could not be ruled out on the basis of dynamical stability. Additionally, \cite{gregory2011} applied a maximum likelihood method to the GJ~581 system and claims significant detection of a planet near $400$ days, but does not find significance for a planet near $33$ days. We carry out an analysis which combines modeling the published radial velocities for this system and N-body integrations, in order to take into account gravitational interactions between the planets. Our work examines the stability of the system with 4 and 5 planets as a function of eccentricity and orbital inclination (assuming all planets are co-planar) over a timescale of 10 Myr, whereas \citeauthor{toth2013} integrate their models over only 1 Myr. \citeauthor{mayor2009} find that, for some initial conditions, the system becomes unstable (with planet e escaping) even after a few Myr, further justifying our choice of a longer integration timescale. We describe the N-body simulations and our method for fitting the radial velocities in Section 2. Section 3 contains a discussion of computational issues. We report our results in Section 4. In Section 5 we address the impact on our results of including system inclination in our simulations. We discuss our results in the context of the habitable zone of the system in Section 6, and conclude in Section 7. | Using either the fit eccentricities relative to parameter space study determined critical values or N-body model lifetimes relative to some reasonable fraction of a solar lifetime as discrepancy variables, we can test the null hypothesis that the 4- and 5-planet models are stable. The hypothesis that the 4-planet Keplerian model for GJ~581 is stable fails rejection with a p-value of $0.59$ based on stability cutoffs on the results of an RV fit, and fails rejection with $p=0.50$ by sampling the posterior distribution for N-body model lifetime. The hypothesis that the 5-planet (including g) Keplerian model for GJ~581 is stable fails rejection with $p=0.21$ using eccentricity cutoffs and $p=0.10$ by sampling the posterior distribution. Neither the 4- nor 5-planet model can be categorically ruled out on stability grounds. When the constraint of stability is included in the model fits for 4 and 5 planet Keplerian models of GJ~581, the best fit for a constrained 4-planet model is 2.78 with 21 fit parameters and 240 data points. Similarly, 26 fit parameters for the 5-planet model leads to a best fit reduced chi squared value of 2.58. The F-value of the two distributions is 1.077, and the null hypothesis that they represent the same variance fails rejection with a p-value of .3. The semi-major amplitude of planet g is about $0.9 m/s$, compared to a mean instrumental error of $1.24 m/s$, a RMS residual of about $2.0 m/s$, and a stellar jitter of $1.4 m/s$. Previous studies have predicted a stellar jitter due to surface effects of $1.9 m/s$ \citep{vogt2010}. While planet g cannot be ruled out on stability grounds, it does pose greater stability concerns than a 4 planet model, and when stability is a factor in determining best fits to the radial velocity data does not result in enough of an improvement in the fit to pass an F-test. Given that the semi-major amplitude for planet g is lower than the instrumental error, the expected stellar jitter for a star of this type, and the residual noise in the fit, we cannot claim at this time that the data confirms the existence of planet g. The stable fraction of the posterior distribution of $10\%$ for the 5-planet model does indicate a strong potential for instability in a model for GJ~581 that includes planet g, and any further measurements used to validate the existence of g should include an in depth stability analysis. For additional multi-planet systems where stability is a concern, a key orbital parameter affecting stability is orbital eccentricity. One approach to including stability constraints in RV fits could be to start with a parameter space exploration focused primarily on limiting eccentricity, and then once upper bounds for reasonable eccentricities have been defined re-run the RV fit with limits on the eccentricities. This can be efficient, as it allows for a large sampling of the RV fit space, and greater detail in the shape of the posterior distribution; however, it can overestimate the stable portion of the posterior distribution, in particular missing cases where combined effects from multiple planets result in instability at lower eccentricities than from a single planet alone. Starting from the fit result using an MCMC algorithm and then sampling from that distribution as input to N-body simulation provides a better estimate of the likelihood a given planetary model is stable, but requires significantly greater computer time to build a detailed posterior distribution. With either approach, individual N-body simulations require long run-time with few objects and are not easily sped up with parallel computing, but both approaches require many such models to be run, all of which are independent of each other and parallelism can be used to speed up the calculation of the ensemble. Initial tests show that this approach could be applied to other multi-planet systems using modest computational hardware by taking advantage of GPGPU technology. | 14 | 4 | 1404.4381 |
1404 | 1404.1930_arXiv.txt | We present Wisconsin H-Alpha Mapper \sii $\lambda 6716$ and \Ha\ spectroscopic maps of the warm ionized medium (WIM) in the Scutum-Centaurus Arm at Galactic longitudes $310^\circ < l < 345^\circ$. Using extinction-corrected \Ha\ intensities (\IHac), we measure an exponential scale height of electron density-squared in the arm of $\HHa = 0.30 \kpc$ (assuming a distance of $3.5 \kpc$), intermediate between that observed in the inner Galaxy and in the Perseus Arm. The $\sii / \Ha$ line ratio is enhanced at large $|z|$ and in sightlines with faint \IHac. We find that the $\sii / \Ha$ line ratio has a power law relationship with \IHac\ from a value of $\approx 1.0$ at $\IHac < 0.2 \R$ (Rayleighs) to a value of $\approx 0.08$ at $\IHac \gtrsim 100 \R$. The line ratio is better correlated with \Ha\ intensity than with height above the plane, indicating that the physical conditions within the WIM vary systematically with electron density. We argue that the variation of the line ratio with height is a consequence of the decrease of electron density with height. Our results reinforce the well-established picture in which the diffuse \Ha\ emission is due primarily to emission from {\em in situ} photoionized gas, with scattered light only a minor contributor. | The warm ionized medium (WIM) traces energy transport in the diffuse interstellar medium (ISM) of star-forming galaxies through supernova-driven turbulence \citep{Armstrong:1995ho,Haverkorn:2008kt,Chepurnov:2010bt,Hill:2008gv,deAvillez:2012hp} and photoionization. Because of its large ($\sim 1 \kpc$; \citealt{Haffner:1999hi,Gaensler:2008bq,Savage:2009kj}) scale height, its weight is substantial, making it an important contributor to the hydrostatic pressure of the ISM \citep{Boulares:1990ba}. Evidence for the existence of the WIM comes from an absorption signature in the synchrotron spectrum \citep{Hoyle:1963ve}, pulsar dispersion \citep{Manchester:1981ge,Taylor:1993il}, Faraday rotation \citep{Mao:2010eg,Foster:2013gg}, and faint emission \citep{Reynolds:1973jo,Dettmar:1990tv} and absorption \citep{Savage:2009kj,Howk:2012gt} lines. Classical \ion{H}{2} regions, which are ionized by a local association of hot stars, are distinct from the WIM: the \ion{H}{2} regions have different spectral signatures \citep{Madsen:2006fw}, a much lower scale height ($\sim 50 \pc$ in the Milky Way), and a higher dust content \citep{Kreckel:2013fh,Rueff:2013ep}. However, like \ion{H}{2} regions, the WIM is associated with star formation. It is found in every star-forming galaxy, accounting for $59 \pm 19\%$ of the \Ha\ flux in normal star-forming galaxies \citep{Oey:2007cj} and $\gtrsim 90 \%$ of the \Hp\ mass \citep{Haffner:2009ev}. The WIM is primarily ionized by photons which escape from \ion{H}{2} regions and then travel through low-density pathways around neutral hydrogen established by stellar feedback and turbulence \citep{Reynolds:1990kw,Ciardi:2002bz,Wood:2010dg}. The \Ha\ emission line provides the bulk of the information about the distribution of the WIM, with \Ha\ from the WIM of the Milky Way detected in every direction \citep{Haffner:2003fe,Haffner:2009ev}. Up to $\approx 20 \%$ of the faint \Ha\ flux may be scattered light which originated in \ion{H}{2} regions, with the WIM contributing the remainder \citep{Reynolds:1973jo,Wood:1999bd,Witt:2010gd,Brandt:2012fw}. The relative intensities of collisionally-excited optical emission lines, primarily \nii $\lambda 6584$ and \sii $\lambda 6713$ (hereafter \nii\ and \sii), provide the best probes of the ionization state and temperature of the gas in the WIM. In general, the line intensity ratios $\sii/\Ha$ and $\nii/\Ha$ track each other and increase with decreasing \Ha\ intensity both in the Milky Way \citep{Haffner:1999hi,Madsen:2006fw} and in other galaxies; infrared and ultraviolet collisionally-excited lines trace similar physical trends \citep{Rand:2008dj,Rand:2011hd,Howk:2012gt}. Also, \oii\ is relatively bright while \oi, \oiii, and \Hei\ are faint in the WIM \citep{Domgorgen:1994kk,Reynolds:1995ki,Reynolds:1998ji,Mierkiewicz:2006ky}. In combination, these observations indicate that the WIM has a higher temperature ($8000 \K$ compared to $6000 \K$) but lower ionization state (O$^+$ and S$^+$, not O$^{++}$ and S$^{++}$, are the dominant ions) than classical \ion{H}{2} regions. The physical cause of the trends in temperature and ionization state inferred from the observed line ratios must relate to the heating and ionization of the gas. One would expect the radiation field to harden as it escapes \ion{H}{2} regions, but this would produce the opposite of the observed line ratio trends: a harder ionizing radiation field would produce higher ionization states in the WIM. Models in which the radiation field is dilute ($U \sim 10^{-3} - 10^{-4}$ photons per electron) --- allowing ions time to recombine before encountering an ionizing photon --- can explain the low ionization state \citep{Mathis:1986et}, but variations in $U$ cannot explain the observed constancy of $\sii/\nii$. A supplemental heating source which scales less steeply with density than photoionization heating (and therefore dominates at densities $n_e \lesssim 0.1 \cucm$) appears necessary to explain the inferred temperature trend \citep{Reynolds:1999jp}. In edge-on star-forming galaxies, a clear relationship between line ratios and height is observed \citep{Domgoergen:1997wh,Tullmann:2000td,Otte:2002dr,Hoopes:2003ec}. However, because the \Ha\ intensity also decreases with height, it is unclear whether the line ratios are fundamentally a function of height, density, or a combination of the two. In this paper, we use Wisconsin H-Alpha Mapper (WHAM) observations to investigate the physical properties of the WIM associated with the Scutum-Centaurus Arm, seen edge-on in the Galactic longitude range $310^\circ \lesssim l \lesssim 340^\circ$ in the local standard of rest velocity range $-80 \kms < \vlsr < -40 \kms$ \citep[e.~g.][]{Benjamin:2008uh}, analogous to the work done in the Perseus Arm in the outer Galaxy by \citet{Haffner:1999hi} and \citet{Madsen:2006fw}. We describe our observations in Section~\ref{sec:obs} and present them in Section~\ref{sec:maps}. We estimate an extinction correction from \ion{H}{1} data, tested against \Hb\ spectra in a few sightlines, in Section~\ref{sec:dust}. In Section~\ref{sec:height}, we measure the scale height of the WIM in the Scutum-Centaurus Arm. We present $\sii/\Ha$ line ratios and interpret them as a measurement of the temperature of the gas in Section~\ref{sec:lineratios}, addressing the degeneracy between height and density as underlying causes of the observed line ratio trends. We present conclusions in Section~\ref{sec:discussion}, arguing that scattered light is a minor contributor to the \Ha\ light attributed to the WIM \citep[in contrast to the model of][]{Seon:2012ci}. Finally, we summarize the paper in Section~\ref{sec:summary}. | \label{sec:discussion} The power-law relationship between $\sii/\Ha$ and \IHa, which is observed here as well as in local gas and the Perseus Arm, suggests that a relatively uniform physical process is at play across the observed range of \Ha\ intensities. The most likely explanation is a temperature effect, with higher $\sii/\Ha$ values corresponding to higher temperatures. A correlation between the $\nii/\Ha$ line ratio and the temperature of the emitting gas is now well established \citep{Haffner:1999hi,Collins:2001dz,Madsen:2006fw}. Ascribing the variation of $\sii/\Ha$ to temperature is less straightforward. The line ratio depends upon the temperature, sulfur abundance, and ionization state of the gas \citep{Otte:2002dr}: \begin{equation} \frac{\sii}{\Ha} = 7.49 \times 10^5 T_4^{0.4} \, e^{-2.14/T_4} \frac{\mathrm{H}}{\mathrm{H}^+} \frac{\mathrm{S}}{\mathrm{H}} \frac{\mathrm{S^+}}{\mathrm{S}}. \end{equation} The hydrogen ionization faction $\mathrm{H}^+ / \mathrm{H} \gtrsim 0.9$ in the WIM \citep{Reynolds:1998ji}, while we assume that the sulfur abundance $\mathrm{S}/\mathrm{H}$ is does not change appreciably in the diffuse ISM in the arm, leaving a dependence on the sulfur ionization fraction $\mathrm{S}^+/\mathrm{S}$ and the temperature. Sulfur can be in either the $\mathrm{S}^+$ or the $\mathrm{S}^{++}$ state in photoionized gas because its first ionization potential is less than that of hydrogen, while the second ionization potential is $0.9 \textrm{ eV}$ below the helium edge. Photoionization modelling of sulfur is difficult because of the unknown temperature dependence of its dielectronic recombination rate \citep{Ali:1991ba,Haffner:2009ev,Barnes:2014wc}. However, because the observed variations in $\sii/\Ha$ largely track variations in $\nii/\Ha$ \citep{Haffner:1999hi,Madsen:2006fw,Haffner:2009ev}, it is likely that they are primarily tracing temperature variations, while changes in $\sii/\Ha$ relative to $\nii/\Ha$ (or, equivalently, changes in $\sii/\nii$) trace changes in the sulfur ionization state. Our conclusion that $\sii/\Ha$ and, thus, the temperature of the gas depends primarily on \Ha\ intensity, not height, provides a clue about the non-photoionization heating of the WIM. The lower \Ha\ intensity presumably corresponds to a lower electron density in the emitting gas; if the lower intensity instead corresponded to a shorter path length, \sii\ would scale in the same way, leaving $\sii/\Ha$ unchanged. The temperature of the WIM is determined by the balance between heating and cooling, often parameterized as \citep{Reynolds:1999jp,Wiener:2013fl} \begin{equation} \Lambda n_e^2 = G_0 n_e^2 + G_1 n_e + G_2 + G_3 n_e^{-1/2}, \end{equation} where $\Lambda$ is the cooling function. Pure photoionization heating ($G_1 = G_2 = G_3 = 0$) cannot explain the observed $\nii/\Ha$ line ratio of the WIM at low \Ha\ intensities: an additional heating source parameterized by non-zero $G_1$ (such as dissipation of turbulence or photoelectric heating of dust grains), $G_2$ (magnetic reconnection), or $G_3$ (cosmic ray heating) is required \citep{Reynolds:1992ci,Reynolds:1999jp,Otte:2002dr,Wiener:2013fl,Barnes:2014wc}. Assuming such a heating source is important, our conclusion that the $\sii/\Ha$ depends more strongly on \IHa\ than on $|z|$ suggests that the heating mechanism at a given $n_e$ does not vary significantly with $|z|$. This argues against photoelectric heating because the FUV radiation field has a significantly smaller scale height ($\sim 300 \pc$) than $n_e$. Cosmic rays, on the other hand, most likely have a larger scale height than the gas, so $G_3$ varies more slowly with height than $n_e$. However, more detailed modelling and \nii\ data to measure the temperature independent of the sulfur ionization state are required to test this. The power-law relationship continues to $\IHac \approx 100 \R$ and then flattens. This change in slope suggests a physical change between the WIM and \ion{H}{2} regions. The $62$ sightlines with $\IHac > 100 \R$ regions have a mean $\langle \sii/\Ha \rangle = 0.082$ and a standard deviation $\sigma_{\sii/\Ha} = 0.037$, consistent with \ion{H}{2} region observations in other contexts. Alternatively, systematic uncertainty introduced by our extinction correction could produce this effect if the extinction to \ion{H}{1} column density ratio has increased scatter at large \ion{H}{1} columns. A spectroscopic survey of \Hb\ in the region we have mapped in \Ha\ would provide a more rigorous extinction correction. The scale height and midplane densities we observe in the Scutum-Centaurus Arm at Galactocentric radius $r_G \approx 6 \kpc$ are intermediate between those observed in the inner Galaxy at $r_G \approx 4-6 \kpc$ \citep{Madsen:2005ft}, locally, and in the Perseus Arm at $r_G \approx 10 \kpc$ \citep{Haffner:1999hi,Madsen:2006fw}. In the inner Galaxy, the scale height is $\HHa \sim 95-150 \pc$, while we find $\HHa \approx 300 \pc$ in Scutum-Centaurus, and $\HHa = 400 \pc$ in Perseus. The corresponding space-averaged rms midplane electron densities are $f_0^{1/2} n_{e,0} = 0.5-1 \cucm$, $0.2 \cucm$, and $0.11 \cucm$, using the same assumptions as \citet{Madsen:2005ft}. Our conclusion that the observed $\sii/\Ha$ line ratio depends primarily on the \Ha\ intensity of the sightline is relevant to recent claims that a significant fraction of the \Ha\ emission we attribute to the WIM in fact originated in higher-density regions and has been scattered into the beam by dust. \citet{Seon:2012ci} produced a model which they claim reproduces the high optical line ratios observed in the WIM with a large scattered light contribution. In their model, the diffuse \Ha\ emission primarily originates in late O and early B star \ion{H}{2} regions and is scattered by dust into high-latitude sightlines. O and B star \ion{H}{2} regions have typical line ratios $\sii/\Ha \le 0.15$ \citep{Reynolds:1988gp,Wood:2005cc,Madsen:2006fw}, inconsistent with the typical values of $\approx 0.4$ observed for the WIM either in local gas at high latitudes or in the Scutum-Centaurus or Perseus Arms. \citet{Seon:2012ci} argue that, in their picture, line ratios could increase with distance from the late O and early B \ion{H}{2} regions. However, our conclusion that the line ratio depends primarily on the electron density, not $|z|$ (Figure~\ref{fig:ratio}$c$ and $d$), qualitatively supports the interpretation that the variations in $\sii/\Ha$ trace temperature or density variations, not scattered light: we do not expect gas density variations at a fixed $|z|$ in the emitting regions of the WIM to correlate with distance from OB stars. | 14 | 4 | 1404.1930 |
1404 | 1404.2933_arXiv.txt | We examine the statistics of the low-redshift \lya\ forest from smoothed particle hydrodynamic simulations in light of recent improvements in the estimated evolution of the cosmic ultraviolet background (UVB) and recent observations from the Cosmic Origins Spectrograph (COS). We find that the value of the metagalactic photoionization rate (\gh) required by our simulations to match the observed properties of the low-redshift \lya\ forest is a factor of 5 larger than the value predicted by state-of-the art models for the evolution of this quantity. This mismatch in \gh\ results in the mean flux decrement of the \lya\ forest being underpredicted by at least a factor of 2 (a 10$\sigma$ discrepancy with observations) and a column density distribution of \lya\ forest absorbers systematically and significantly elevated compared to observations over nearly two decades in column density. We examine potential resolutions to this mismatch and find that either conventional sources of ionizing photons (galaxies and quasars) must be significantly elevated relative to current observational estimates or our theoretical understanding of the low-redshift universe is in need of substantial revision. | By virtue of its physical and chemical simplicity, the intergalactic medium (IGM) serves as an exquisite calorimeter, recording the instantaneous ionizing emissivity and heat produced by cosmic sources at each epoch. At $z<6$ the IGM is highly ionized (Gunn \& Peterson 1965), with a fluctuating residue of neutral hydrogen: the \lya\ forest \citep{lynds71}. After nearly two decades, the \lya\ forest remains the most well-understood and robust prediction of cosmological hydrodynamic simulations \citep[e.g][]{cen94, zhang95, miralda96, hernquist96, rauch97}. This robustness arises because the \lya\ forest is dominated by gas at moderate overdensity; gas that traces the dark matter (with only mild impact by pressure forces) and whose temperature is governed by the simple processes of photoionization heating and adiabatic cooling \citep[e.g.][]{weinberg98, peeples10a}. It is this simplicity that makes the calorimeter reliable: IGM models suggest the neutral fraction of gas at the cosmic mean density at $z\sim 3$ is $n_{HI}/n_{H} \sim 10^{-5.5}$ \citep[e.g.][]{kollmeier03}, and this low neutral fraction {\it must} be maintained by the background of photo-ionizing radiation produced by quasars and star-forming galaxies \citep{miralda90,hm96}, possibly augmented by other undiscovered sources. Determining the intensity of the ultra-violet background (UVB) --- specifically the hydrogen photoionization rate, \gh, is non-trivial. Because of its low surface brightness, the metagalactic UVB is not directly measured but is inferred through multiple independent channels, such as the quasar proximity effect \citep[e.g.][]{bechtold87} or the low surface brightness emission from the outskirts of galactic disks \citep[e.g.][]{adams11}. The measurements are intrinsically difficult and subject to significant uncertainties and challenges (e.g. anisotropic QSO radiation, uncertain local gas densities). Alternatively, one can {\it predict} the intensity and spectrum of the UVB by synthesizing measurements of all possible sources of ionizing flux and the absorption and re-emission of UV radiation by the IGM and high column-density absorbers \citep[][hereafter, HM01]{hm96,hm01}. While this procedure relies on a host of observational inputs, the most uncertain is the fraction $\fesc$ of ionizing photons that escape from star-forming galaxies. Direct (and difficult) ionizing continuum measurements suggest that $\fesc \sim10\%$ at $z \sim 3-4$ (e.g. Shapley et al. 2006, Vanzella et al. 2010) and is substantially lower at $z < 1$ \citep[e.g.][]{bridge10, barger13}. Exploiting our theoretical understanding of the IGM provides a third avenue for probing \gh. By forcing a match between the (more easily observed and well understood) opacity of the \lya\ forest and that from a theoretical IGM (typically taken directly from simulations) we have an independent determination that can be compared with both indirect measurements and the predicted UVB. There has historically been excellent agreement between the predicted UVB and the value inferred from the mean opacity. It was precisely the comparison between the predicted and observed forest opacity that provided strong arguments (now confirmed) for a ``high'' baryon density and low associated deuterium abundance \citep{rauch97,weinberg97}. In this paper, we demonstrate that this excellent agreement no longer holds at low redshift. Specifically, we will show a factor of $\sim 5$ discrepancy between the \gh\ predicted by the most sophisticated model of UVB evolution (Haardt and Madau 2012; hereafter HM12) and the value required to reproduce observed properties of the \lya\ forest. We show in Figure~\ref{fig:background} the predicted \gh\ from HM12 (black solid line) compared to observational determinations. The dashed line shows an independent model of the UVB from \cite{faucher-giguere09}, which overall is quite similar to that of HM12. The large open star, is the value reported here. \begin{figure}[H] \plotone{fig1.pdf} \caption{ The photoionization rate as a function of redshift for the HM12 UVB (solid) compared to observational constraints at $z=2-4$ (circles \citet{bolton07}, triangles \citet{becker07} and squares \citet{faucher-giguere08}) and the value we infer from our \lya\ forest modeling at $z=0.1$ (open star). The red triangle shows the low-redshift upper limit inferred by Adams et al. (2011) from non-detection of H$\alpha$ emission in the galaxy UGC 7321. The dashed line shows an alternative UVB model from \cite{faucher-giguere09}. The dotted line shows a model, discussed in \S 4.1, with a constant galaxy escape fraction $\fesc=15\%$.} \label{fig:background} \end{figure} We take advantage of new measurements of the column density distribution (CDD) of the low-redshift \lya\ forest from Cosmic Origins Spectrograph (COS) observations \citep{danforth14} to determine \gh\ by comparison with cosmological simulations of the IGM. We will further show that our conclusions would be very similar if we instead use the mean decrement as our observational measure. After describing the cosmological simulation that we use (\S 2) and inferring the value of \gh\ required to match the observed CDD (\S 3), we discuss (in \S 4) possible resolutions to the discrepancy between our inferred \gh\ and the value predicted by HM12, the ``photon underproduction crisis'' (PUC) of our title. None of these resolutions alone appears entirely satisfactory. The most exciting possibility is that this discrepancy is probing exotic sources of ionizing photons or novel heating mechanisms in the diffuse IGM operating far above the usual theoretical expectations. | The factor of 5 discrepancy between the value of \gh\ required to match cosmological models of the $z=0$ IGM to the observed mean decrement and CDD and that predicted by state-of-the art models for the evolution of the extragalactic UVB (HM12) highlights a significant gap in our current understanding of the sources of the UV background or the structure of the IGM, or both. We have discussed a number of possible resolutions, no one of which appears satisfactory. The {\it least} radical solution is to increase the mean $\fesc$ at low-redshift such that galaxies dominate the emissivity {\it and} simultaneously boost the quasar emissivity, though both of these changes oppose our current understanding of these sources. For the undaunted, extra photons from decaying dark matter or a drastic change to the physical structure of the IGM as predicted by LCDM may also be the resolution to the photon underproduction crisis. | 14 | 4 | 1404.2933 |
1404 | 1404.5514_arXiv.txt | Dynamic interactions between the two Magellanic Clouds have flung large quantities of gas into the halo of the Milky Way. The result is a spectacular arrangement of gaseous structures including the Magellanic Stream, the Magellanic Bridge, and the Leading Arm (collectively referred to as the Magellanic System). In this third paper of a series studying the Magellanic gas in absorption, we analyze the gas ionization level using a sample of \nsl\ \emph{Hubble Space Telescope}/Cosmic Origins Spectrograph sightlines that pass through or within 30\degr\ of the 21\,cm-emitting regions. We find that \msper\ (\ndet/\nsl) of the sightlines show UV absorption at Magellanic velocities, indicating that the total cross section of the Magellanic System is $\approx$11\,000 square degrees, or around a quarter of the entire sky. % Using observations of the \sit/\siw\ ratio together with \emph{Cloudy} photoionization modeling, we calculate the total gas mass (atomic plus ionized) of the Magellanic System to be $\approx$2.0$\times$10$^9$\msun\,($d$/55\,kpc)$^2$, % with the ionized gas contributing around three times as much mass % as the atomic gas. This is larger than the current-day interstellar \hi\ mass of both Magellanic Clouds combined, indicating that they have lost most of their initial gas mass. If the gas in the Magellanic System survives to reach the Galactic disk over its inflow time of $\sim$0.5--1.0\,Gyr, it will represent an average inflow rate of $\sim$3.7--6.7\smy, potentially raising % the Galactic star formation rate. However, multiple signs of an evaporative interaction with the hot Galactic corona indicate that the Magellanic gas may not survive its journey to the disk fully intact, and will instead add material to (and cool) the corona. | Like other massive spiral galaxies, the Milky Way has a gas-consumption time of a few billion years \citep[e.g.][]{La80}, far shorter than its age. Ongoing fuel replenishment (gaseous accretion) is therefore needed for it to sustain its star formation. Chemical and kinematic studies of stars in the solar neighborhood have shown that the Milky Way has been forming stars at a fairly constant rate for several Gyr \citep{Tw80, Bi00, Ch01}, so we know that fuel replenishment has to be occurring. However, several open questions surround {\it how} accretion onto the Milky Way and other $\approx\!L_\ast$ galaxies occurs. Is the accreting gas predominantly cold, warm, or hot? Does it arrive smoothly from the unenriched IGM or episodically following satellite-galaxy interactions? And do most accreting streams become disrupted by fluid instabilities before reaching the disk? % Hydrodynamical simulations offer one approach to answering these questions. They make clear predictions for the physical and chemical state of gas accreting onto galaxies \citep{Ke05, Ke09, Ba10, Fu11, St11, Fe12, VS12, Mu12, Jo12b, Vo12, Sh13}. However, clear observational evidence for gas accretion in external galaxies has been hard to find \citep[see][]{Rb11, Le13, Bu13}, largely due to a lack of kinematic information. Unambiguous \emph{spectroscopic} evidence for inflow onto external galaxies exists in only a small number of cases \citep[e.g.][]{Ru12}, and in these galaxies the inflow may be metal-enriched. Fortunately, the Milky Way offers a prime opportunity for studying the fueling of an $\approx\!L_\ast$ galaxy. We have detailed knowledge of the neutral-gas distribution around the Galaxy from \hi\ 21\,cm observations, and information on the ionized-gas distribution from a large body of absorption-line data taken from ultraviolet (UV) spectra of background QSOs. In both the 21\,cm and UV data, gaseous inflow onto the Galaxy can be seen directly among the high-velocity clouds (HVCs), interstellar clouds that are not co-rotating with the Galactic disk. First detected by \citet{Mu63}, HVCs are defined as having LSR velocities $|v_{\rm LSR}|\!>\!100$\kms\ \citep[or sometimes $>$90\kms; see reviews by][]{WW97, Ri06, Pu12}. While HVCs trace a variety of interstellar processes, \emph{infalling} HVCs can be identified by their negative velocities in the Galactic-Standard-of-Rest (GSR) frame and their low metallicities. Most \hi\ HVCs are now known to be within 5--20\,kpc of the Galactic disk, based on individual cloud distance measurements \citep{Wa01, Wa08, Th08, LH10, Sm11} and on the observation that the HVC sky-covering fraction $f_{\rm HVC}$ measured toward halo stars is similar to $f_{\rm HVC}$ measured toward AGN \citep{LH11, Le12}. However, the Magellanic Stream (MS), anchored by the Magellanic Clouds, stands as a notable exception. Although the distance to the MS is poorly constrained, it likely lies in the interval $\sim$55--200\,kpc. The SMC distance gives the lower limit, and we use the tidal models of \citet{Be12} to bound the upper limit, % although \citet{JL08} argue the MS distance near the south Galactic pole is unlikely to exceed 100\,kpc \citep[see also][]{BH13}. The Stream's origin in the Magellanic Clouds is supported by its spatial, kinematic, and chemical properties, but the mechanism by which it was removed from the Clouds has puzzled dynamicists since its discovery in 21\,cm observations over forty years ago \citep{Di71, WW72, Ma74}\footnote{Parts of the Stream were detected even earlier by \citet{Di65}, although the association with the Magellanic Clouds was not realized at that time.}. Recent models favor the tidal removal of much of the Stream during a close encounter between the two Magellanic Clouds $\approx$2\,Gyr ago \citep{GN96, Co06, Be10, Be12, DB11, DB12}, although ram-pressure stripping \citep{Me85, MD94, Ma05} and supernova-driven blowout of LMC material \citep[][hereafter N08]{SS03, LH07, Le09, Ni08} may have also contributed to its production. Tidal forces are strongly favored as the mechanism responsible for creating the Leading Arm (LA), the gaseous counterpart to the MS lying in front of the direction of motion of the Magellanic Clouds \citep{Wa72,Ma74,Pu98,Lu98}. To probe the Stream's physical and chemical conditions across its full length on the sky, we are conducting an absorption-line survey with both UV (\hst/Cosmic Origins Spectrograph) and optical (VLT/UVES) spectrographs. In Paper I \citep{Fo13a}, we presented new MS chemical abundance measurements along the sightlines to the AGN \object{RBS\,144} and \object{NGC\,7714}, and an upper limit on the metallicity toward the QSO \object{PHL\,2525}. Combined with earlier work on the \object{NGC\,7469} sightline \citep[][hereafter F10]{Fo10}, there are now three good measurements of $\approx$0.1 solar metallicity in the main body of the Stream, supporting the view that most of the MS was stripped from the SMC (not the LMC) 1.5--2.5\,Gyr ago. This is because the SMC had a metallicity of $\approx$0.1 solar at that time according to its age-metallicity relation \citep{PT98, HZ04}. However, in Paper II \citep[][see also Gibson et al. 2000]{Ri13}, we found a much higher metallicity of 0.5 solar along the inner-Stream sightline to QSO \object{Fairall\,9}, which passes through a MS filament that can be traced kinematically back to the LMC (N08). This shows that the bifurcation of the Stream, previously seen in its spatial extent \citep{Pu03a} and kinematics (N08), is also seen in its metal enrichment. This supports a dual origin for the Stream, with both the SMC and LMC contributing to its origin. We now turn our attention from the origin of the Stream to its fate, by observing its ionization level, which encodes information on the physical processes occurring as it interacts with the ambient plasma and radiation field. The MS and the LA contain both warm-ionized and highly-ionized material. The warm-ionized phase is seen in \ha\ emission \citep{WW96, Pu03b, BH13, Ba14} and absorption in low-ionization UV lines \citep[][F10, Paper I, Paper II]{Lu94, Lu98, Se03, Fo05a}. Photoionization and/or shock ionization followed by recombination may be responsible for exciting the \ha\ emission seen from the Stream, including the possibility of ionization by a Seyfert flare at the center of the Milky Way $\approx$1--3\,Myr ago \citep{BH13}. The highly ionized phase in the Stream is seen in \cf\ \citep[][F10, Paper II]{Lu94, Lu98} and \os\ \citep[][F10]{Se03, Fo05a} absorption. These high ions have column-density ratios consistent with an origin in the conductive or turbulent interfaces that surround the Stream, where the warm low-ionization gas adjoins the hot coronal plasma \citep{Fo05a}. In this paper, we present the first UV survey of the Stream's ionization properties, using \nsl\ COS sightlines passing through or within 30\degr\ of the MS, the LA, or the Magellanic Bridge (MB), the gaseous filament that connects the LMC and SMC. For brevity, we use the phrase ``the Magellanic System'' (MSys) to encompass the MS, MB, and LA. This is similar to the definition used by \citet{Br05} in their Parkes survey of 21\,cm emission from Magellanic gas, except that here we include the ionized outer regions of the MS that are not seen at 21\,cm. Furthermore, \citet{Br05} define the ``Interface Region'' of \hi-emitting gas lying in-between the Magellanic Clouds and the Stream proper. In this paper we include that region as part of the Stream, since the two principal MS filaments can be traced kinematically through it back to the Magellanic Clouds (N08), and hence there is no reason to treat it as a distinct object. Physically, the MSys refers to all gas that was stripped from the Magellanic Clouds at some point in the past (though not necessarily all at the same time). We define On-System and Off-System directions as those with and without a detection of \hi\ 21\,cm emission at Magellanic velocities. We also define a subset of ``LMC-Halo'' directions, that lie in the outer halo of the LMC; these are not part of the MS as traditionally defined, but are clearly Magellanic in origin, and hence are included in our analysis. We also include Wisconsin H-alpha Mapper (WHAM) observations of \ha\ emission from the Stream in six of the MS sightlines and three of the MB sightlines. This allows us to compare the optical, radio, and UV profiles of the Magellanic gas in the same directions. Throughout the paper we use solar (photospheric) elemental abundances from \citet{As09} and absorption-line data (rest wavelengths and oscillator strengths) from \citet{Mo03}. We present all velocities in the kinematical local standard-of-rest (LSR) frame, and all column densities are given in units of cm$^{-2}$. This paper is laid out as follows. \S2 describes the sample assembly, observations, data handling procedures, and absorption-line measurements. \S3 presents an overview of the UV absorption-line profiles and WHAM emission-line profiles. \S4 discusses the physical conditions of the Magellanic gas, presenting new constraints on the ionization level in both the low-ion and high-ion gas phases. In \S5 we derive and discuss the total (neutral plus ionized) gas mass in the MSys. In \S6 we discuss the accretion rate of Magellanic gas onto the Galaxy. In \S7 we briefly discuss the MSys in terms of intervening quasar absorption line systems. \S8 summarizes the main results. | We have presented an absorption-line survey of ionization in the MSys (MS, MB, LMC Halo, and LA) using a sample of medium-resolution (18\kms\ FWHM) \hst/COS UV spectra of \nsl\ background AGN lying within 30\degr\ of the 21\,cm emission from the MSys. These data are supplemented by Wisconsin H$\alpha$ Mapper (WHAM) \ha\ emisssion-line observations in \nwham\ directions, and 21\,cm emission spectra from the LAB survey, GASS survey, and additional Parkes telescope observations. The sightlines sample approximately four orders of magnitude of \hi\ column density, ranging from the dense cores of the MB to the diffuse outer layers of the MS and LA. This dataset has allowed us to characterize the physical properties of the MSys and its relationship to the extended Galactic halo. Our main results are the following. \begin{enumerate} \item{\bf Detection Rate}. UV absorption at Magellanic velocities is detected in \ndet\ out of \nsl\ directions (\msper\ detection rate). The detected lines include one or more of \sit\ $\lambda$1206, \siw\ $\lambda$1190, \cw\ $\lambda$1334, \sif\ $\lambda$1393, and \cf\ $\lambda$1548, and occasionally other lines including \oi\ $\lambda$1302. Since the 21\,cm-emitting regions of the MSys cover 2\,701 square degrees \citep{Ni10}, the total cross-section of the MSys is enormous, at $\approx$11\,000 square degrees, which is around a quarter of the entire sky. % The ionized (UV-absorbing) regions occupy $\approx$four times as much % area as the neutral (21\,cm-emitting) regions. \item{\bf Line ratios}. We measure the column-density ratios \sit/\siw, \sif/\siw, and \cf/\cw\ in the Magellanic gas. Among directions with 21\,cm detections (with log\,$N$(\hi)$\ga$18), all three ratios show significant anti-correlations with $N$(\hi) (Figure 4). In addition, \sif/\siw\ and \cf/\cw\ each show weak (but significant) anti-correlations with MS Longitude, such that the ionization level increases along the Stream from the Magellanic Clouds toward the tip. Given our current understanding of the Stream's orbit, in which the tip is the most distant region, this result is equivalent to an increase in ionization level with Galactocentric distance, consistent with the observed decline in $N$(\hi) along the MS. \item {\bf Ionization parameter and gas density.} Using \emph{Cloudy} photoionization models applied to \ncloudy\ Magellanic directions with measured \hi\ columns and \sit/\siw\ ratios, we derive ionization parameters log\,$U$ in the Magellanic gas ranging from --3.8 to --3.1 with a median value of --3.5. % This corresponds to an average Magellanic gas density log\,($n_{\rm H}$/cm$^{-3}$)$\approx$--1.8 given the % calculated density of the ionizing radiation field (including both Milky Way and Magellanic photons). \item {\bf Ionization level and warm \hw\ column.} The hydrogen ionization level $x_{\rm H~II}$ varies considerably between MSys regions, depending on the \hi\ column. In the MB, where log\,$N$(\hi)$\ga$20, $x_{\rm H~II}\!\approx\!20$\%, whereas the directions with log\,$N$(\hi)$\la$19.5 have $x_{\rm H~II}$ up to $\approx$98\%. % Since the Off-Stream directions have a larger cross-section than the On-Stream directions, the gas is predominantly ionized in most sightlines through the MSys, and even the phase referred to as ``low-ion'' phase is predominantly ionized. The warm \hw\ columns in the low-ion phase of the MSys take a fairly narrow range of values, between log\,$N$(\hw)$\approx$19.4--20.1, % even though the \hi\ columns in the same directions cover several orders of magnitude. \item {\bf Thermal pressure.} The MS thermal pressure $P/k$ calculated from the photoionization models varies from $\approx$30--250\,cm$^{-3}$\,K in the six Stream directions % analyzed. This places the Stream close to pressure equilibrium in a two-million-degree hydrostatic Galactic corona at 50--100\,kpc, where $P_{\rm corona}/k$=100--250\,cm$^{-3}$\,K in the isothermal models of \citet{Sg02}. \item {\bf Total mass}. We calculate the total (neutral plus ionized) gas mass of the MSys to be M(MSys)$\sim$2.0$\times$10$^9$\,($d$/55\,kpc)$^2$\msun, % by combining the \hi\ mass of 4.9$\times$10$^8$\msun\ \citep{Br05} with estimates of the \hw\ mass in 21\,cm-bright regions (9.5$\times$10$^8$\msun) and 21\,cm-faint regions (5.5$\times$10$^8$\msun). % The MS accounts for about half of the mass in the MSys. M(MSys) is over twice as large as the remaining interstellar \hi\ mass of the LMC and SMC combined, indicating \emph{these two dwarf galaxies have lost over two-thirds of their initial gas mass.} M(MSys) is also comparable to the mass of the hot Galactic corona out to 50\,kpc, M$_{\rm hot}\!\approx$1.7$\times$10$^9$\msun($r_{\rm hot}$/50\,kpc)$^3$($n_{\rm hot}$/10$^{-4}$\,cm$^{-3}$), although $r_{\rm hot}$ and $n_{\rm hot}$ are poorly constrained. The similarity between the mass of the MSys and the corona indicates that energetically, both components will be affected by their mutual interaction, causing the corona to cool as the Magellanic gas heats up. \item {\bf Mass flow rate onto the Galaxy.} The present-day infall rate of the MSys onto the Milky Way is $\approx$3.7--6.7\smy, % using an average galactocentric infall velocity of 100\kms, a MS distance of 55--100\,kpc, and the total gas mass derived above. % This is considerably larger than the inflow rate of all nearby HVCs combined and larger than the current SFR of the Galaxy ($\approx$1--2\smy). \emph{The MSys therefore has the potential to raise the global Galactic SFR}. However, this gaseous fuel faces a tortuous path to reach the disk: multiple signs of an evaporative encounter with the hot corona suggests that the MSys is disintegrating and replenishing the hot coronae with fresh material at large galactocentric radii. For the inflow rate across the entire halo to be in or close to equilibrium, cooler clouds must condense out of other, denser or more metal-enriched regions of the corona. The MSys can therefore be thought of as fueling the Galactic halo rather than directly fueling star formation in the disk. \end{enumerate} {\it Acknowledgments.} We thank David Nidever for providing the total area and velocity field of the \hi\ emission from the Magellanic System. Support for programs 11692, 12204, 12263, and 12604 was provided by NASA through grants from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS~5-26555. WHAM science and ongoing operations are supported by NSF award AST~1108911. The Parkes radio telescope is part of the Australia Telescope National Facility which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO. K.A.B. is supported through NSF Astronomy and Astrophysics Postdoctoral Fellowship award AST~1203059. The research was partially supported by the Japan Society for the Promotion of Science through Grant-in-Aid for Scientific Research 23740148. | 14 | 4 | 1404.5514 |
1404 | 1404.2428_arXiv.txt | Recently, neutron stars with the mass around $2M_{\odot}$ ($M_{\odot}$: the solar mass) have been reported. For example, the binary millisecond pulsar, J1614-2230, has the mass of $1.97 \pm 0.04 M_{\odot}$~\cite{Demorest:2010bx}, and the mass of pulsar, J0348+0432, is estimated to be $2.01 \pm 0.04 M_{\odot}$~\cite{Antoniadis:2013}. However, it is difficult to explain such heavy neutron stars by the equation of state (EoS) which have been calculated in mean-field theory so far, because the degrees of freedom of hyperons ($Y$) make the EoS very soft and the maximum mass of a neutron star is thus reduced. To solve this ``hyperon puzzle'', we examine in detail the extension of SU(6) spin-flavor symmetry based on the quark model to SU(3) flavor symmetry in determining the couplings of the isoscalar, vector mesons to the octet baryons. | We have calculated the EOS for neutron stars, using the popular RMF models as well as the QMC and CQMC models. As a result, We have found that the models except GM3, FSUGold and IU-FSU can explain the masses of J1614-2230 and/or J0348+0432 in SU(3) symmetry. The extension from SU(6) to SU(3) symmetry and the strange vector meson, $\phi$, are very significant in sustaining massive neutron stars. In addition, the variation of baryon structure in matter helps prevent the collapse of a neutron star. | 14 | 4 | 1404.2428 |
|
1404 | 1404.5664_arXiv.txt | Comet 17P/Holmes underwent a massive outburst in 2007 Oct., brightening by a factor of almost a million in under 48 hours. We used infrared images taken by the Wide-Field Survey Explorer mission to characterize the comet as it appeared at a heliocentric distance of 5.1~AU almost 3 years after the outburst. The comet appeared to be active with a coma and dust trail along the orbital plane. We constrained the diameter, albedo, and beaming parameter of the nucleus to 4.135 $\pm$ 0.610 km, 0.03 $\pm$ 0.01 and 1.03 $\pm$ 0.21, respectively. The properties of the nucleus are consistent with those of other Jupiter Family comets. The best-fit temperature of the coma was 134~$\pm$~11~K, slightly higher than the blackbody temperature at that heliocentric distance. Using Finson-Probstein modeling we found that the morphology of the trail was consistent with ejection during the 2007 outburst and was made up of dust grains between 250~$\mu$m and a few cm in radius. The trail mass was $\sim$ 1.2~-~5.3~$\times$~10$^{10}$~kg. | Comet 17P/Holmes (hereafter 17P) has undergone 3 massive outbursts since its discovery in 1892 \citep{1892Obs....15..441H}, most recently brightening by a factor of almost a million and becoming visible to the naked eye in 2007 Oct. The outbursts are likely thermally-driven since all three occurred 6-9 months after 17P passed through perihelion. Dynamically, 17P appears to be a typical Jupiter family comet (JFC) with a semi-major axis of 3.62 AU, eccentricity of 0.43, and inclination of 19$^{\circ}$.1. The comet is enigmatic given its propensity for unusually large outbursts but dynamical and physical properties similar to other JFCs. The material ejected during the 2007 outburst included gas species, dust grains, and macroscopic fragments (e.g. \citealt{2008ApJ...680..793D,2009P&SS...57.1162C,2009AJ....137.4538Y,2010Icar..208..276R,2010AJ....139.2230S}). Much of the smaller dust expanded in an almost spherical shell around the nucleus, while larger dust grains were observed to separate as a ``blob'' at a slower rate of $\sim$~120~-~135~m~s$^{-1}$ (e.g. \citealt{2008A&A...479L..45M,2009AJ....138..625L,2010MNRAS.407.1784H}). \cite{2010Icar..208..276R} and \cite{2012A&A...542A..73B} detected a slower moving core-component of the largest grains that separated from the nucleus at a relative velocity of $\sim$ 7-9 m s$^{-1}$. Large dust grains may persist in the vicinity or along the trail of a comet for years after ejection from the nucleus \citep{1990Icar...86..236S,1998ApJ...496..971L,2011ApJ...738..171B}. In this work we used infrared (IR) images obtained with the Wide-Field Infrared Survey Explorer (WISE) to examine the evolution of 17P several years after the 2007 outburst. | \subsection{Physical properties of the nucleus} The effective nucleus diameter of 4.135 $\pm$ 0.610 km calculated here is larger than previous estimates of 3.24 $\pm$ 0.02 km \citep{2006MNRAS.373.1590S} and 3.42 $\pm$ 0.14 km \citep{2009A&A...508.1045L}, derived when 17P appeared to be inactive or only weakly active prior to its 2007 outburst. Both of the previously mentioned results were determined from optical observations using an assumed albedo of 0.04. The diameter reported here is consistent with previously reported results when they are corrected using the NEOWISE-derived albedo. We derived an albedo of 0.03 $\pm$ 0.01. The albedo of the nucleus is consistent with those measured for other comets, which generally occupy a narrow range between 0.02 and 0.06 \citep{2004come.book..223L}. We note that we are unable to derive an albedo for material in the trail as we do not have simultaneous high signal-to-noise observations at optical wavelengths. \cite{2010ApJ...714.1324I} used optical, near-IR, and mid-IR observations to constrain the albedo of the ejecta within a few days of the outburst. They found that the albedo of the material (as observed at a phase angle of 16$^{\circ}$) decreased during their observations from 0.12 $\pm$ 0.04 to 0.032 $\pm$ 0.014 and suggested that sublimating volatiles would lower the albedo. \cite{2012ApJ...760L...2L} estimated the geometric albedo of the dust in the coma as 0.006 $\pm$ 0.002, also at a phase angle of 16$^{\circ}$. Such a low value is not unheard-of \citep{2002Sci...296.1087S,2004Icar..167...37N}, though it does not match well with results from \cite{2010ApJ...714.1324I}. The discrepancy may be due to different populations of grains dominating the thermal emission in the IR and the stellar extinction at optical wavelengths. Our result for the albedo of the nucleus is generally consistent with the albedo of the ejecta observed in 2007 when the albedo was determined from combined optical and mid-IR wavelengths. The beaming parameter of 1.03 $\pm$ 0.21 is consistent with the average value of 1.03~$\pm$~0.11 reported for 57 JFCs by the SePPCoN survey, which used thermal infrared measurements by the Spitzer Space Telescope \citep{2013Icar..226.1138F}. They found that beaming parameters for 57 JFC nuclei were approximately normally distributed and suggested that there appears to be little variation among bulk thermal properties of JFCs. As discussed in \cite{2013Icar..226.1138F}, a beaming parameter close to 1.0 implies low thermal inertia and little nightside emission. 17P appears typical in this regard. \subsection{Thermal emission in the coma} The best-fit coma temperature of 134~$\pm$~11~K is $\sim$~10\% warmer than the temperature expected for an ideal blackbody ($T_{BB}$) at a heliocentric distance of 5.13 AU (123~K assuming $T_{BB}$ $\propto$ 278~K~$r_{H}^{-0.5}$; \citealt{1992Icar..100..162G}). Previous results from observations taken close to the time of outburst ($r_{H}$ $\sim$ 2.4 AU) have also suggested that the dust temperature of the ejecta exceeded the local blackbody temperature of $\sim$ 180 K. \cite{2009AJ....137.4538Y} reported a dust temperature near the nucleus of 360 $\pm$ 40 K using near-IR observations obtained with the NASA Infrared Telescope Facility several days after the outburst, while mid-IR results obtained around the same time suggested cooler temperatures between 172 K and 200 K \citep{2010ApJ...714.1324I, 2009PASJ...61..679W}. Spitzer Space Telescope observations obtained on 2007 Nov.\ 10 resulted in an estimated temperature of 260 K for the near-nucleus dust \citep{2010Icar..208..276R}. Most of these results are higher than the estimated blackbody temperature to varying degrees, matching well with our results here. Numerous IR observations of comets have shown that it is common, perhaps even the norm, for comae and dust tail and trail temperatures to exceed the temperature expected for a co-located blackbody (e.g. \citealt{1988AJ.....96.1971T,1992Icar..100..162G,1998ApJ...496..971L,2000ApJ...538..428H}). Generally, excess emission at IR wavelengths is attributed to either small grains ($\lesssim$~1$\mu$m) that are unable to radiate efficiently at IR wavelengths, rough surfaced grains that are more emissive than the smooth spherical grains modeled by the blackbody temperature, or larger grains that maintain a thermal gradient across their surface \citep{1982come.coll..341C,1990Icar...86..236S,2001indu.book...95S}. In the case of 17P, all of these effects may be present. The nucleus was seen to remain active in the months and years following the outburst, likely releasing small dust grains from the surface (\citealt{2012AJ....144..138S}; Snodgrass, private communication). Based on results from the best-fit synchrone determined in section~\ref{sec:trail}, particles larger than a few cm in diameter would still be close enough to the nucleus to contribute to the excess thermal emission observed here. \subsection{An old trail} The morphology of the trail is consistent with being debris ejected during the 2007 outburst and observations by \cite{2013ApJ...778...19I} that showed large dust grains following the comet around aphelion in Oct.\ 2010. We constrained the range of particle diameters observed between $\sim$ 250 $\mu$m and a few cm. The larger grain diameters are consistent with the sizes of grains observed in a slow-moving ``core'' near the nucleus just a few days after the outburst, which were determined to be $\gtrsim$~200~$\mu$m \citep{2010Icar..208..276R,2012A&A...542A..73B}. We measured the flux along the trail using a box aperture that extends between 11$^{\prime\prime}$ and 880$^{\prime\prime}$ from the nucleus and has a width perpendicular to the length of the trail of 52$^{\prime\prime}$. The flux was calibrated and color-corrected as described in section~\ref{sec:wobs}. To correct for light potentially lost outside of the large aperture, we applied an aperture correction of -0.03~mag derived by \cite{2013AJ....145....6J}. To estimate the cross-section of material present we used the following relation from \cite{2005Icar..179..158M}: \begin{equation} \sigma_{\lambda} = \frac{F_{\lambda} ~ \Delta^{2}}{B_{\lambda}(T)} \label{eq:cross sec} \end{equation} where $\sigma_{\lambda}$ is the cross-section of material observed at wavelength $\lambda$ (in this case, 22~$\mu$m, or W4), F$_{\lambda}$ is the observed flux, $\Delta$ is the geocentric distance, and $B_{\lambda}(T)$ is the Planck function at temperature $T$. We were unable to constrain the temperature of the dust along the trail as the signal in W3 is too low to fit a Planck function to. We therefore assumed that the temperature is between the expected blackbody temperature of 123~K and the measured coma temperature of 134~K. This is consistent with the previously-discussed finding that many comet trails are at or exceed local blackbody temperatures. We also assumed that the temperature is constant along the trail and is not dependent on the size of the dust grains present. The cross-section of material was 1.5~$\times$~10$^{9}$~m$^{2}$ in the case of the local blackbody temperature or 10$^{9}$~m$^{2}$ in the higher temperature case. We used previously measured minimum and maximum particle sizes ($a-$, $a+$) of 250~$\mu$m and 3~cm, and assumed that the differential size distribution of particles follows a power law of the form $n(a)~da \propto a^{-q}~da$, with the value of $q$ set between 2.2 and 3.4, as measured by \cite{2010Icar..208..276R} and \cite{2012A&A...542A..73B}, respectively. The mean particle size within the trail was given by: \begin{equation} \bar{a} = \frac{\int_{a-}^{a+} \pi a^{3} n(a) da}{\int_{a-}^{a+} \pi a^{2} n(a) da} \label{eq:area} \end{equation} The mass within the observed trail was then given by: \begin{equation} M = \frac{4~\rho~\bar{a}~\sigma_{\lambda}}{3} \label{eq:mass} \end{equation} where $\rho$ is the bulk density of the grains and was assumed to be 1000~kg~m$^{-3}$ \citep{1991ASSL..167...19J}. The mass in the trail was $\sim$ 1.2 - 5.3 $\times$ 10$^{10}$ kg. This represented approximately 1~-~100\% of the total ejected mass \citep{2008ICQ....30....3S,2009AJ....138.1062S,2010Icar..208..276R,2010ApJ...714.1324I,2011ApJ...728...31L,2012A&A...542A..73B}. Assuming an average grain size of 200~$\mu$m, \citealt{2010Icar..208..276R} estimated the mass of the slow-moving core seen in 2007 to be $\sim$~4~$\times$~10$^{9}$ kg. \citealt{2012A&A...542A..73B} estimated the mass to be significantly higher at $\sim$~0.7-4~$\times$~10$^{11}$~kg by summing over an estimated particle size distribution with $a_{-}$ = 0.1~$\mu$m, 10 $< a_{+} <$ 1000 mm, and -3.3 $< q <$ -3.0. Thus, the dust trail observed by WISE represented 3~-~75\% of the core modeled by \citealt{2012A&A...542A..73B} in 2007. \subsection{Why did 17P outburst?} The overarching question remains to be answered: why does 17P undergo massive outbursts when most JFCs experience only mild mass loss? The diameter, beaming parameter, and albedo of the nucleus are similar to those of other JFC nuclei. The volatile species observed shortly after the outburst in 2007 similarly fail to provide any obvious clues about the cause of the outburst. Relative abundances of CN, C$_{2}$, C$_{3}$ and NH and the isotopic ratios of $^{12}$C/$^{13}$C and $^{14}$N/$^{15}$N in CN and HCN were similar to those observed for other comets \citep{2008ApJ...679L..49B,2009AJ....138.1062S}. Several species, including C$_{2}$H$_{6}$, HCN, CH$_{3}$OH, and C$_{2}$H$_{2}$, were enhanced with respect to H$_{2}$O although only by a factor of a few \citep{2008ApJ...680..793D}. The perihelion distance of 17P changed from 2.17 AU in 2000 to 2.05 AU in 2007 following a close encounter with Jupiter. The change resulted in a $\sim$~10\% increase in solar insolation at the surface. The small difference may have caused the thermal wave to propagate deeper than on previous perihelion passages, reaching previously unheated pockets of volatiles. A runaway exothermic phase transition of amorphous water ice to crystalline is probably insufficient to cause the outburst \citep{2010Icar..207..320K}. If supervolatiles such as CO and/or CO$_{2}$ are trapped within the amorphous ice and heated sufficiently, the resulting gas production may be able to drive such activity, if the gas can build up sufficient internal pressure \citep{2009AJ....138.1062S,2009ICQ....31...99S,2011Icar..212..847K,2012Icar..221..147H}. It is possible that the nucleus of 17P has unusually high tensile strength that allows gas pressure to build up in the interior before releasing the energy in a sudden outburst upon surface failure. \cite{2010Icar..208..276R} suggested that the nucleus must have a strength between 10~-~100~kPa in order to have survived the 2007 outburst. However, previous studies of other comets suggest much lower strengths for JFCs. 16P/Brooks 2 and D/1993 F2 (Shoemaker-Levy~9) both underwent tidal splitting during close encounters with Jupiter leading to estimates of 0.1 kPa and $\sim$ 0.38 kPa for the tensile strengths of the nuclei \citep{1985AJ.....90.2335S,1998P&SS...46...21S}. Based on observations of mini-outbursts, \cite{2008Icar..198..189B} estimated the strength of the sub-surface material of 9P/Tempel 1 to be not much more than 0.01 - 0.1 kPa, while \cite{2005Sci...310..258A} found that the strength of the surface must also be extremely low ($<$ 0.065 kPa). Only a few comets have estimated tensile strengths and are not necessarily representative of all JFCs. We note simply that the estimated tensile strength required of 17P is an order of magnitude higher than those estimated for other JFCs. | 14 | 4 | 1404.5664 |
1404 | 1404.2272_arXiv.txt | Previous studies have reported the existence of two counter-rotating stellar disks in the early-type spiral galaxy NGC\,7217. We have obtained high-resolution optical spectroscopic data ($R \approx 9000$) with the new fiber-based Integral Field Unit instrument VIRUS-W at the 2.7\,m telescope of the McDonald Observatory in Texas. Our analysis confirms the existence of two components. However, we find them to be co-rotating. The first component is the more luminous ($\approx$\,77\% of the total light), has the higher velocity dispersion ($\approx$\,170\,\kms) and rotates relatively slowly (projected $v_\mathrm{max} = 50$\,\kms). The lower luminosity second component, ($\approx$\,23\% of the total light), has a low velocity dispersion ($\approx$\,20\,\kms) and rotates quickly (projected $v_{max} = 150$\,\kms). The difference in the kinematics of the two stellar components allows us to perform a kinematic decomposition and to measure the strengths of their Mg and Fe Lick indices separately. The rotational velocities and dispersions of the less luminous and faster component are very similar to those of the interstellar gas as measured from the \oiii\ emission. Morphological evidence of active star formation in this component further suggests that NGC\,7217 may be in the process of (re)growing a disk inside a more massive and higher dispersion stellar halo. The kinematically cold and regular structure of the gas disk in combination with the central almost dust-free morphology allows us to compare the dynamical mass inside of the central 500\,pc with predictions from a stellar population analysis. We find agreement between the two if a Kroupa stellar initial mass function is assumed. \vspace{.25 cm} \noindent {}$^{\dagger}${This paper includes data taken at The McDonald Observatory of The University of Texas at Austin.}\\ {}$^{\diamond}${This paper contains data obtained at the Wendelstein Observatory of the Ludwig-Maximilians University Munich.} | Photometric studies have been decomposing galaxies into multiple stellar components for a long time now \citep[e.g.][]{de-Vaucouleurs1959}. The technique has become common practice in the attempt to disentangle the formation histories of galaxies. The problem of a kinematic decomposition, i.e.\ the detection of genuinely separate components in the line-of-sight velocity distributions (LOSVDs) --- especially in later type galaxies --- places higher demands on the data quality, both in terms of spectral resolution and signal-to-noise ratio. Nevertheless, disk-like structures have been detected in elliptical galaxies \citep{Bender1988a,Franx1988,Davies1988,Jedrzejewski1989,Scorza1990, Scorza1995} and spectroscopic surveys now provide statistics on the occurrence of kinematic subcomponents in early-type \citep{Krajnovic2011}, S0s \citep{Kuijken1996}, and spiral galaxies \citep{Pizzella2004}. \begin{figure*} \begin{center} \includegraphics[width=\textwidth, bb=0 30 750 610]{plots/finderchart2} \end{center} \caption{A \textit{gri} composite of NGC\,7217. The white box outlines 105\arcsec$\times$55\arcsec~field of view of the VIRUS-W IFU.\ We obtained the images with the Wendelstein Wide Field Imager at the new 2\,m Fraunhofer Telescope \citep{Gossl2012, Hopp2012} on the 25 and the 27 of October 2013. The inset in the lower right shows an HST F450W/F336W false color composite of the central 25\arcsec$\times$25\arcsec\ (Program 11128; PI David Fisher).} \label{fig:finderchart} \end{figure*} We have constructed a new, high-spectral resolution Integral Field Unit spectrograph called VIRUS-W \citep{Fabricius2008b,Fabricius2012b}, that is designed specifically to study the stellar LOSVDs of nearby disk galaxies. It offers a spectral resolution of $R \approx 9000$ ($\sigma_{inst}=15$\,\kms) in the optical which allows us to resolve the low velocity dispersions of a few tens of \kms\ typically found in disky systems, and to study the fine structure of the corresponding LOSVDs. As a test case, we observed the early-type spiral galaxy NGC\,7217. Its dynamical structure has been of particular interest because it hosts multiple rings. There are two star-forming rings with diameters of 63\arcsec\ and 156\arcsec, and a third inner dust ring with a diameter of 21\arcsec\ (\citealp{Buta1995}, hereafter B95). This inner nuclear ring marks a significant change in morphology: the outer flocculent spiral disappears completely and gives away to a smooth central light distribution \citep[][see \reffig{fig:finderchart}]{Fisher2008}. The ring locations seemingly correspond to resonances \citep{Verdes-Montenegro1995} but NGC\,7217 does not host any obvious, non-axisymmetric structure such as a stellar bar that could create corresponding resonances. Given its relative isolation \citep{Karachentseva1973} tidal effects caused by other galaxies can also be ruled out as a source of the resonances. However, B95 carried out an extensive photometric analysis of NGC\,7217 and through a Fourier analysis find a weak perturbation to the axisymmetry which may be a faded bar. Previous work has claimed that NGC\,7217 hosts a large-scale counter-rotating stellar disk \citep{Merrifield1994,Silchenko2000, Fabricius2012a}, a phenomenon that has been observed only in a handful of systems so far. The prototypical example of this class of systems is NGC\,4550, where its bimodal LOSVD reveals that $\approx$50\% of the stars are on retrograde orbits \citep{Rubin1992, Rix1992, Emsellem2004}. Only a few similar objects are known: \citet{Prada1996} claim that NGC\,7331 has a counter-rotating bulge, and~\cite{Zeilinger2001} describe stellar counter-rotation in NGC\,3521. NGC\,3593 hosts a counter-rotating component that dominates the light in the central regions \citep{Bertola1996}. Further examples include NGC\,4138 \citep{Jore1996, Haynes2000} and counter-rotation caused by interaction in NGC\,5719 \citep{Vergani2007, Coccato2011a}. It was hypothesized that the counter-rotation in NGC\,7217 may be the result of a minor merger event or the cold accretion of gas \citep{Merrifield1994, Pizzella2004}. The observed ring structure \citep{Lovelace1997} has been attributed to both the putative minor merger \citep{Silchenko2006} and instabilities created by counter-rotation. In an attempt to disentangle the two counter-rotating disks and to probe our ability to detect and to study kinematic substructure in stellar systems, we obtained observations of the central region of NGC\,7217, covering one of its two stellar rings (see \reffig{fig:finderchart}). We recover non-parametric LOSVDs with the proper treatment of nebular emission. We do not confirm the existence of counter-rotation in this galaxy. Rather we find a sub-dominant, kinematically cold and rapidly rotating stellar disk embedded in a higher dispersion, co-rotating essentially spherical stellar halo. In the next section we will briefly describe the characteristics of the new spectrograph that we use for this work. In \refsec{sec:observations}, we will then describe the observations. In \refsec{sec:data_reduction}, we discuss the basic data reduction, the algorithm that we use for the recovery of the non-parametric LOSVDs, the kinematic decomposition, and the method of direct spectral decomposition that we use to derive abundances of the two stellar components. In \refsec{sec:results}, we present kinematic maps for the two stellar components and the ionized gas. We also present maps for the line strength determinations of the two individual components. In \refsec{sec:tilted_ring} we analyse the gas velocity field through a tilted ring model and in \refsec{sec:mass_to_light} we derive a central mass-to-light ratio from the gas rotation and compare this value to the prediction from a stellar population analysis. We discuss the implications of our findings in \refsec{sec:discussion}, and summarize in \refsec{sec:conclusions}. | \label{sec:conclusions} Using the novel VIRUS-W spectrograph we have obtained moderately high resolution $R \approx 9000$ optical Integral Field Unit observations of NGC\,7217. Based on our high signal-to-noise data we are able to revisit the kinematic structure of this galaxy and to test previous claims of the existence of a counter-rotating stellar disk. Using a new algorithm, we derive non-parametrized line-of-sight velocity distributions and carry out double-Gaussian decompositions. We also use the methodology introduced by \citet{Coccato2011a,Coccato2013} to confirm our findings and to derive line strength indices for the two stellar components. \medskip \noindent Our main findings are: \begin{itemize} \item We confirm the existence of two dynamically distinct stellar components. In contrast to previous claims by \citet{Merrifield1994}, \citet{Silchenko2000}, and \citet{Fabricius2012a} we do not find them to be counter rotating. Rather we are able to decompose them into one hot, dominant, round, slowly-rotating component with a velocity dispersion of $\approx$170\,\kms, and a cold, co-rotating stellar disk with a velocity dispersion of $\approx$20\,\kms. \item The velocity and velocity dispersion fields of the cold stellar disk are very similar to those of the gas as derived from the \oiii\ emission lines. Together with the blue colours of the rings in this galaxy (B95) and the visible sites of active star formation, this supports a picture where the stellar disk is still in the process of regrowing. \item The kinematic position angles of the cold stellar disk component and the hot spheroidal component are identical within our measurement errors, rendering an external origin of the gas unlikely. \item We find an increase of velocity dispersion of the gas at the inner stellar ring. This may be a result of the resonant nature of the ring, but also a result of the ongoing star formation in that region. \item The two components are clearly separated in \MgbFe\ space in the sense that the disk component shows larger equivalent widths in $\langle \mathrm{Fe} \rangle$ and lower equivalent widths in \mgb\ that cannot be explained by an age difference. This points to a different star formation history, with a shorter-lived period of star formation in the spheroid and a later and/or longer-lived star formation in the disk. \item The Lick indices measured in the central regions of the hot component are more similar to those in the central regions of elliptical galaxies than to those in the bulges of spirals. \item We confirm the existence of a misalignment of the gas velocity field in the central arcseconds. We attempt a tilted ring model, but cannot confirm a 90\Deg\ angular separation between the central disk and the outer disk rotation axes as would be necessary for a polar ring that was described by~\cite{Zasov1997, Silchenko2000}. The maximum angular separation we find is 30\Deg. While this may be a consequence of the limited spatial resolution of our data, we stress that a natural explanation for the central deviation from the global rotation could lie in the weak break in axisymmetry that was reported by B95 \citep{Athanassoula1992}. \item The tilted ring analysis provides a rotation curve for the gas. This allows us to derive the total enclosed dynamical mass and a deprojected mass-to-light ratio of $\Upsilon^*_{dyn} = 4.5 M_{\astrosun}/L_{\astrosun}$. Using GALXEX FUV/NUV, SDSS, and IRAC I1,2,3 bands we also conduct a stellar population analysis inside the central 10\arcsec\ and find that the predicted mass-to-light ratio is in reasonable agreement with the dynamical value if a Kroupa IMF is assumed. \item The structural parameters of the two stellar components (scale lengths, ellipticities and position angles) are in good agreement with the values obtained from photometry (B95). This demonstrates that a kinematic decomposition is feasible for spheroid/disk systems. This method has the advantage of being completely model free; it does not rely on the extrapolation of model profiles. As such, it is fully complementary to photometric methods and allows us to test, for instance, assumptions where all the light that exceeds the inwards extrapolation of an outer exponential disk is attributed to the bulge of a system. It also enables us to probe the kinematic properties and to measure stellar population parameters beyond broadband colours. This extends the technique of the spectral decomposition \citep{Coccato2011a,Coccato2013} from the application to counter-rotating stellar disks to spheroid-disk systems, as long as the two components have sufficient separation in velocity dispersion. This makes it possible to probe the origins of individual components in a multicomponent system separately and unambiguously, rendering it a powerful tool to study the formation of such galaxies. \end{itemize} We suggest that the main bulk of stars in NGC\,7217, i.e.\ the spheroidal component, formed through a major merger. The merger remnant has photometric and spectroscopic properties more similar to those of an elliptical galaxy than to those of the bulges of spiral galaxies. The disk component formed after the merger, presumably from relatively primordial gas acquired through minor mergers or cold accretion from the intergalactic medium \citep{Mazzuca2006, Eliche-Moral2010}, or an external reservoir as suggested by the significant offset of the two stellar populations in the \MgbFe\ plane. | 14 | 4 | 1404.2272 |
1404 | 1404.0870_arXiv.txt | The Sunyaev-Zel'dovich effect provides a useful probe of cosmology and structure formation in the Universe. Recent years have seen rapid progress in both quality and quantity of its measurements. In this review, we overview cosmological and astrophysical implications of recent and near future observations of the effect. They include measuring the evolution of the cosmic microwave background radiation temperature, the distance-redshift relation out to high redshifts, number counts and power spectra of galaxy clusters, distributions and dynamics of intracluster plasma, and large-scale motions of the Universe. | The Sunyaev-Zel'dovich effect (SZE, \cite{zs69,sz70,sz72,sz80a}) is inverse Compton scattering of the Cosmic Microwave Background (CMB) photons off electrons in clusters of galaxies or any cosmic structures. It is amongst major sources of secondary anisotropies of the CMB on sub-degree angular scales. The most noticeable feature of the SZE is that its brightness is apparently independent of the source redshift $z$ because its intrinsic intensity increases with redshifts together with the energy density of seed (CMB) photons; otherwise observed brightness should decrease rapidly as $(1+z)^{-4}$. This makes the SZE a unique probe of the distant Universe. The SZE also has a characteristic spectral shape which helps separating it from other signals such as radio galaxies and primary CMB anisotropies. Recent developments of large area surveys by the South Pole Telescope (SPT) \cite{spt09,spt10,spt11,spt13a}, the Atacama Cosmology Telescope (ACT) \cite{act10,act11,act13a}, and the Planck satellite \cite{planck_earlysz,planck13a} have enlarged the sample of galaxy clusters observed through the SZE by more than an order of magnitude over the last decade as illustrated in Figure \ref{fig-nsz}. To date, the SZE by thermal electrons (thermal SZE) has been detected for about $1000$ galaxy clusters including more than 200 new clusters previously unknown by any other observational means. The imprint of yet unresolved smaller-scale cosmic structures has been explored by means of their angular power spectrum \cite{reichardt12,sievers13,planck_y} and the stacking analysis \cite{hand11,planck_group}. There have been reports of detections of the SZE by peculiar motions of galaxy clusters (kinematic SZE) either statistically \cite{hand12} or from a high-velocity merger \cite{sayers13b}. \begin{figure}[th] \centering\includegraphics[width=85mm]{figs/nsz_z.eps} \caption{ Redshift histograms of galaxy clusters with measured thermal SZE signals and redshifts from the literature. Solid line indicates 258 clusters detected in the ground-based surveys either by SPT \cite{spt09,spt10,spt11,spt13a} or by ACT \cite{act10,act11,act13a}, excluding overlaps, over a total of $\sim 3000$ deg$^2$. Dashed line shows 813 clusters detected by the Planck satellite over all sky \cite{planck_earlysz,planck13a}. For reference, hatched region marks 34 clusters with $>4\sigma$ SZE detections published as of 2002 (\cite{birkinshaw99,carlstrom02} and references therein) which consist mainly of X-ray luminous clusters.} \label{fig-nsz} \end{figure} The sensitivity of the SZE observations of individual clusters has also improved significantly, making it a useful tool for studying physics of intracluster plasma. In particular, the SZE data provide a direct measure of thermal pressure of electrons, which is highly complementary to X-ray observations. They allow us to study the distance-redshift relation (e.g., \cite{schmidt04,bonamente06}), three-dimensional structures \cite{defilippis05,sereno12}, and dynamics \cite{kitayama04,korngut11,planck_coma} of galaxy clusters. By means of the SZE, we are witnessing the high-mass end of structure formation in the Universe that in turn serves as a powerful probe of cosmology. Theoretical foundations and earlier observations of the SZE are reviewed extensively by \cite{sz80b,rephaeli95,birkinshaw99,carlstrom02}. In the present paper, we focus mainly on practical applications of the SZE that have become more feasible by recent observations, and discuss their cosmological and astrophysical implications. Unless explicitly stated otherwise and wherever necessary to assume specific values of cosmological parameters, we adopt a conventional $\Lambda$CDM model with the matter density parameter $\Omega_{\rm m}=0.3$, the dark energy density parameter $\Omega_{\Lambda}=0.7$, the baryon density parameter $\Omega_{\rm b}=0.045$, the Hubble constant $h_{70}=H_0/ (70 \mbox{\:km\:s$^{-1}$Mpc$^{-1}$})=1.0$, the dark energy equation of state parameter $w=-1.0$, the amplitude of density fluctuations $\sigma_8=0.8$, and the spectral index of primordial density fluctuations $n_{\rm s}=0.96$. | Extensive efforts well over four decades have now established the SZE as an indispensable tool in cosmology and astrophysics. Being one of the major foregrounds of the CMB, the SZE not only plays a key role in recovering correctly the primary anisotropies, but also offers unique cosmological tests on its own. They include measurements of the evolution of the CMB temperature, distances to high redshifts that are entirely free from the cosmic distance ladder, the absolute numbers and the power spectra of galaxy clusters, and large-scale motions of the Universe. It should be noted that their accuracy critically depends on our understanding of the physics of galaxy clusters and structure formation, which the SZE observations have also been improving, e.g., by finding high velocity cluster mergers, measuring pressure profiles, and detecting the gas in low-mass halos. Perhaps the most noticeable progress over the last decade or so is that the SZE measurements have started to achieve their own discoveries independently of any other means. This has made the SZE a truly complementary probe to X-ray observations in the studies of cosmic plasma. A number of outcomes from large area surveys and pointed observations by existing instruments are also underway. It is highly anticipated that future SZE measurements from both grounds and the space will continue to provide us new insights into our Universe. | 14 | 4 | 1404.0870 |
1404 | 1404.5949_arXiv.txt | Andrews-Lindsay 1 is a pertinent open cluster granted it may host the planetary nebula PHR 1315-6555, yet ambiguities linger concerning its fundamental parameters ($>50$\% scatter). New multiband $BVJHW_{1-4}$ photometry for cluster and field stars, in concert with observations of recently discovered classical Cepheids, were used to constrain the reddening and velocity-distance profiles along the sight-line. That analysis yielded the following parameters for the cluster: $E(J-H)=0.24\pm0.03$, $d=10.0\pm0.4$ kpc ($d_{JH}=9.9\pm0.6$ kpc, $d_{BV}=10.1\pm0.5$ kpc), and $\log{\tau}=8.90\pm0.15$. The steep velocity-distance gradient along $\ell \sim 305 \degr$ indicates that two remote objects sharing spatial and kinematic parameters (i.e., PHR 1315-6555 and Andrews-Lindsay 1) are associated, thus confirming claims that the PN is a cluster member (e.g., Parker et al.). The new distance for PHR 1315-6555 is among the most precise yet established for a Galactic PN ($\sigma/d=4$\%). | } The motivation behind the search for planetary nebulae (PNe) in star clusters is to secure precise parameters for the former. Those parameters include the PN's distance, physical dimensions, age, chemical composition, and potentially progenitor mass \citep{pa11,tu11,mb13}. A solid PN distance could likewise be used to calibrate relationships aimed at establishing distances to field PNe \citep{fr08}. Reliable distances are particularly desirable since precise parallaxes exist for a mere fraction of the nearest PNe \citep[][$\sigma/\pi\sim5$\%]{be09}, whereas distance estimates for the bulk of PNe display $20-30$\% uncertainties or larger \citep[][the latter's Fig.~6]{st08,gi11}. Admittedly, the quest to identify numerous cluster PNe has been hampered by several factors. Massive PNe associated with younger clusters feature short lifetimes ($10^3-10^4$ years). The matter is compounded by the paucity of old clusters that spawn lower-mass PNe exhibiting longer lifetimes. Star clusters rarely survive beyond $10^7$ years \citep{bb11}, and the majority which surpass that threshold host evolved constituents that terminate as SNe \citep[e.g.,][their Fig.~2]{ma07}. Furthermore, the bulk of cataloged Galactic PNe are members of the bulge \citep[][their Fig.~1]{ma07}, which underscores the pertinence of surveys aimed at discovering PNe throughout the Galactic disk where younger clusters reside \citep[][the Macquarie/AAO/Strasbourg H$\alpha$ PN Catalogue, see also \citealt{ja10,lu12}]{pa06}. \begin{table*}[!t] \begin{center} \small \caption{Parameters for Andrews-Lindsay 1 (AL1) \label{table1}} \begin{tabular}{lcllllc} \hline \hline Reference & Target & $d$ (kpc) & $E(B-V)$ & $\tau$ (Gyr) & RV (km/s) & [Fe/H] \\ \hline \citet{jp94} & AL1 & 7.57 & 0.72 & ... & ... & ... \\ \citet{ca95} & AL1 & 11.8 & 0.75 & 0.7 & ... & subsolar \\ \citet{cm04} & AL1 & $12\pm1$ & $0.7\pm0.2$ & 0.8 & ... & ... \\ \citet{fr04} & AL1 & ... & ... & ... & $40\pm10$ & $-0.51\pm0.30$ \\ \citet{fr04b} & AL1 & 9.35 & ... & 0.67 & ... & ... \\ \citet{ca05} & AL1 & $16.95^b$ & $0.34\pm0.05$ & $0.8\pm0.2$ & ... & ... \\ \citet{fr08} & PN & $9.7\pm3.1$ & $0.71^a$ & ... & $51.6\pm15.0$ & ... \\ & AL1 & ... & ... & ... & $50\pm10$ & ... \\ \citet{pa11} & AL1 & ... & ... &... & $57\pm5^a$ & ... \\ & PN & $10.4\pm3.4$ & $0.83\pm0.08$ &... & $58\pm2.5$ & ... \\ & & & & & $59\pm2$ & \\ \hline \end{tabular} \\ (a) See text. (b) Observations were undertaken during unsatisfactory weather conditions. \end{center} \end{table*} Consequently, few promising cases of PNe in Galactic clusters have been reported, and thus any \textit{bona fide} pairs\footnote{Perhaps an equally important endeavour is to eliminate unreliable calibrators and reputed associations \citep{ma07,ki08,mb13}.} are crucial for various research topics (e.g., constraining the impact of mass loss). The search for extragalactic cluster PNe has yielded similar results. \citet{ja13} observed 467 star clusters in M31 to detect PNe that share the former's velocity. That evaluation enables chance superpositions to be identified, in addition to contamination from unrelated emission sources along the sight-line (e.g., H II regions). \citet{ja13} concluded that 5 (of 270) globular clusters may host PNe, whereas those targets identified near open clusters likely constitute chance alignments. A subsample of the tentative Galactic PNe/open cluster pairs includes Abell 8 and Bica 6 \citep{bo08,tu11}, He 2-86 and NGC 4463 \citep{ma07,mb13}, and PHR 1315-6555 and Andrews-Lindsay 1 \citep{pa06,pa11,ma07,fr08}.\footnote{PHR 1315-6555 is likewise designated as PNG 305.3-03.1, and Andrews-Lindsay 1 is cataloged as ESO 96-SC04 and VdB 144.} Yet an important consideration arises when inferring the progenitor mass of those PNe from single stars near the cluster turnoff. The canonical hypothesis advocating that PNe stem from single stars does not readily explain their non-spherical morphologies or low formation rate \citep[][see also \citealt{dm13}]{ja13}. Indeed, the detection of PNe within globular clusters, which exhibit turnoff masses below the 1$M_{\sun}$ threshold predicted by PNe models, implies that the progenitor may have been augmented by mass transfer. Specifically, \citet{ja13} noted that 4 (of 130) Galactic globular clusters host PNe, with the cases split between PNe featuring non-spherical nebulae and high-mass central stars conducive to younger clusters. That evidence, in concert with the realization that three PNe are located in globular clusters hosting numerous X-ray sources, supports claims that multiplicity affects the formation of globular cluster PNe. In each of the aforementioned Galactic cases further independent research is required, with a validation standard on par with the magnitude of the discovery. For example, velocities measured for the PN NGC 2438 by \citet{pk96} indicated it was a cluster member (M46), whereas those by \citet{od63} suggested otherwise. Independent observations urged by \citet{ma07} and others were subsequently published by \citet{ki08}, and suggested that the pair constitute a chance alignment along the sight-line. Radial velocities for Abell 8 and Bica 6 likewise require confirmation, especially given their potential implications for Galactic dynamics \citep{tu13}. NGC 4463 may host a rare young PN (He 2-86) that exhibits sizable internal extinction \citep[][their comprehensive discussion in \S 5]{mb13}, and the cluster is pertinent for stellar evolution research. Specifically, NGC 4463 may host a F-supergiant (post onset of core helium burning), the PN, and a massive blue straggler \citep{al07}. Continuing the cluster PNe project spearheaded by \citet{mb13}, this study aims to bolster the link tying PHR 1315-6555 to Andrews-Lindsay 1, namely by resolving the faint cluster's presently ambiguous fundamental parameters (Table~\ref{table1}). New $BVJHW_{1-4}$ observations were analyzed to achieve that objective. $W_{1-4}$ are the four WISE mid-infrared passbands \citep{wr10}. | } New multiband ($BVJHW_{1-4}$) observations for Andrews-Lindsay 1 and the surrounding field were analyzed, with the objective of highlighting the correct set of cluster parameters (Table~\ref{table1}) and bolstering the link to PHR 1315-6555. Red clump stars identified in the near-infrared photometry were exploited to constrain the run of reddening with distance for the $\ell \sim 305 \degr$ sight-line (Fig.~\ref{fig-ce}). Those data were supplemented with information gleaned from Cepheid variables. The observations imply that the bulk of the reddening is foreground to Andrews-Lindsay 1, and is characterized by $E(J-H)=0.24\pm0.03$. The reddening along $\ell \sim 305 \degr$ increases in concert with distance until $\sim 4$ kpc, and thereafter remains nearly constant (Figs.~\ref{fig-ce}, \ref{fig-vd}). The corresponding optical color-excess matches that cited for the PN \citep{fr08,pa11}, to within the uncertainties (Table~\ref{table1}). The distance for Andrews-Lindsay 1 was inferred from a solar isochrone fit to cluster stars in color-magnitude ($BVJH$) diagrams (Figs.~\ref{fig-nircmd} and \ref{fig-bv}, $d_{JH}=9.9\pm0.6$ kpc, $d_{BV}=10.1\pm0.5$ kpc, and $\log{\tau}=8.90\pm0.15$). The mean distance derived lies between the \citet{jp94} and \citet{cm04} estimates. A velocity-distance correlation predicted for $\ell \sim 305 \degr$ from Galactic rotation generally agrees with the empirical trend delineated by Cepheids, PHR 1315-6555, and Andrews-Lindsay 1. Specifically, the former exhibit negative velocities conducive to nearer targets, whereas the latter two phenomena feature large positive velocities tied to distant objects (Fig.~\ref{fig-vd}). The steep correlation indicates that remote objects that are kinematically and spatially coincident along $\ell\sim305\degr$ are likely associated. The suite of evidence favors an association between the PN PHR 1315-6555 and open cluster Andrews-Lindsay 1 \citep{pa06,ma07,fr08,pa11}. \subsection*{{\rm \footnotesize ACKNOWLEDGEMENTS}} \scriptsize{DM is grateful to the following individuals and consortia whose efforts, advice, or encouragement enabled the research: G. Jacoby, D. Frew, A. Parker, P. Frinchaboy, 2MASS, P. Stetson (DAOPHOT), G. Cibis, F. van Leeuwen, F. Benedict, L. Kiss, W. Gieren, D. Balam, B. Skiff, L. Gallo, R. Thacker, Webda (E. Paunzen), W. Dias, CDS (F. Ochsenbein, T. Boch, P. Fernique), arXiv, and NASA ADS. This publication makes use of data products from the Wide-field Infrared Survey Explorer, which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by NASA; observations obtained at the Southern Astrophysical Research (SOAR) telescope (program ID: CN2013A-157), which is a joint project of the Minist\'{e}rio da Ci\^{e}ncia, Tecnologia, e Inova\c{c}\~{a}o (MCTI) da Rep\'{u}blica Federativa do Brasil, the U.S. National Optical Astronomy Observatory (NOAO), the University of North Carolina at Chapel Hill (UNC), and Michigan State University (MSU).} | 14 | 4 | 1404.5949 |
1404 | 1404.6434_arXiv.txt | We investigate the flux ratio between the 1335 \AA~ and 2326 \AA~ lines of singly ionized carbon in the extended narrow line regions of type 2 quasars at z$\sim$2.5. We find the observed CII $\lambda$1335 / CII] $\lambda$2326 flux ratio, which is not sensitive to the C/H abundance ratio, to be often several times higher than predicted by the canonical AGN photoionization models that use solar metallicity and a Maxwell-Boltzmann electron energy distribution. We study several potential solutions for this discrepancy: low gas metallicity, shock ionization, continuum fluorescence, and $\kappa$-distributed electron energies. Although we cannot definitively distinguish between several of the proposed solutions, we argue that a $\kappa$ distribution gives the more natural explanation. We also provide a grid of AGN photoionization models using $\kappa$-distributed electron energies. | One of the many advantages of looking to the high-redshift Universe is the red-shifting of the rest-frame ultraviolet (UV) emission into the optical observational regime, allowing us to access the unique physics that can be probed using the UV emission or absorption lines. In this paper we examine the ultraviolet CII $\lambda$1335 / CII] $\lambda$2326 flux ratio as a diagnostic of the excitation of extended narrow-line emitting gas associated with active galactic nuclei. Our interest in this flux ratio was initially motivated by the need for indicators of electron temperature (T$_e$) in the rest-frame UV spectral region, to be used with high-z galaxies for which the rest-frame optical temperature diagnostics (i.e., [OIII] $\lambda$4363 / [OIII] $\lambda$5007) have been red-shifted out of the optical observational regime. The large difference in excitation energies of the two CII lines (9.3 vs 5.3 eV) makes their flux ratio strongly sensitive to T$_e$, and hence the flux ratio potentially offers a means to determine T$_e$. Furthermore, the use of lines from a singly-ionized species ought to make this diagnostic readily accessible for both high- and low-excitation objects, AGN and star-forming objects, alike (cf. the [NeIV], [NeV] and OIII] high-ionization temperature diagnostics discussed by Humphrey et al. 2008: H08 hereinafter). In $\S$~\ref{cii_prob}~ we discuss the observational data used in this paper: the CII $\lambda$1335 / CII] $\lambda$2326 flux ratio in type 2, radio-loud quasars at z$\sim$2.5. These data are compared against different excitation models in $\S$~\ref{models}, and in $\S$~\ref{ngc1068}~ we discuss the case of the narrow line region of the nearby Seyfert 2 galaxy NGC 1068, which has been extensively observed at UV wavelengths. In $\S$~\ref{disc}~ we summarize and discuss our findings, and in appendix A1, we give a subset of the exploratory grid of models, using $\kappa$-distributed electron energies, that were computed during the preparation of this paper; the full grid has been made available online. \begin{figure} \special{psfile=plot1.ps hoffset=-20 voffset=-222 hscale=46 vscale=45} \vspace{3.5in} \caption{Results of AGN photoionization model calculations, showing the CII $\lambda$1335 / CII] $\lambda$2326 flux ratio vs. ionization parameter U. The loci of the photoionization models that use a Maxwell-Boltzmann distribution of electron energies are shown by solid black lines, and are plotted for gas metallicities of 1.0, 0.67, 0.1 and 0.01 times the solar value. Loci of photoionization models using solar gas metallicity together with $\kappa$-distributed electron energies are shown by dot-dashed blue lines, and are plotted for $\kappa$ = 40, 20, 10 and 5. In the interest of clarity, the sequences with $\kappa$ = 2.5 and 80 are not shown in this figure. The range of shock model predictions of Allen et al. (2008), plotted at an arbitrary value of U, are shown by the green vertical line. The red triangles show the CII $\lambda$1335 / CII] $\lambda$2326 flux ratios measured from the type 2 quasars, also plotted at arbitrary values of U.} \label{fig_models1} \end{figure} | \label{disc} While investigating the CII $\lambda$1335 / CII] $\lambda$2326 emission line ratio as a potential diagnostic of the temperature, metallicity and excitation in ionized nebulae at high redshift, we have identified a failure of canonical photoionization models of the narrow-line emitting ionized gas associated with several type 2, radio-loud quasars at z$\sim$2.5: the models substantially under-produce CII $\lambda$1335 relative to CII] $\lambda$2326. Given that the narrow line region of NGC 1068 shows a similar CII problem to the quasars (KC00), we suggest that this may be a general problem in fitting the narrow line emitting gases of active galaxies, rather than being an issue specific only to powerful quasars. We have considered several potential causes to explain the higher than predicted CII $\lambda$1335 / CII] $\lambda$2326 ratio in the type 2 quasars: substantially sub-solar gas metallicity; ionization by shocks; continuum fluorescence; or $\kappa$-distributed electron energies. We consider the hypothesis involving $\kappa$-distributed electron energies to be the most promising. It represents what we believe to be the simplest solution to the CII problem, insofar as it does not require the presence of a ionizing sources other than the radiation field of the AGN, nor does it require the summation of multiple models with different input parameters. Moreover, $\kappa$-distributed electron energies have been shown to be successful in solving other, long-standing problems in stellar-photoionized HII regions (see Nicholls et al. 2012, 2013; Binette et al. 2012), and are already widely used in studies of plasmas in the solar system (e.g. Livadiotis \& McComas 2011; Livadiotis et al. 2011). We find continuum fluorescence through CII $\lambda$1335, which enhances the flux of the line above that produced by simple AGN-photoionization, to be another interesting possibility. Given the presence of singly ionized carbon within the beams of AGN UV continuum radiation fields, it seems quite likely that CII $\lambda$1335 (and other resonant lines) could undergo some enhancement by this process. However, the strength of this effect depends on several difficult to constrain properties, such as the three-dimensional geometry of gas in the host galaxy, and the UV continuum luminosity of the hidden quasar nucleus. The z$\sim$2.5 type 2 quasars considered in this work are all radio-loud, and are near the top of the radio luminosity function for AGNs (e.g. Miley \& De Breuck 2008). There is a growing body of evidence to suggest that at least in some radio-loud active galaxies, shocks may contribute to the ionization of the extended narrow line emitting gas (e.g. Clark et al. 1998; Villar-Mart\'{i}n et al. 1999; Best, R\"ottgering \& Longair 2000; Bicknell et al. 2000; De Breuck et al. 2000; Inskip et al. 2002; H08; Humphrey et al. 2010). As such, it would seem entirely reasonable for shocks to contribute to the ionization of the extended narrow emission line regions of the distant quasars we have considered herein. However, this would need to be additional to a strong contribution from photoionization by the central AGN, in order to explain the flux ratios of the bright optical emission lines such as [OIII] $\lambda\lambda$4959,5007, [NeV] $\lambda$3426, H$\beta$, etc (see, e.g., H08). Low gas metallicity is the least appealing of the hypotheses we have considered. Although the AGN-photoionization models using low gas metallicity are able to match the observed CII $\lambda$1335 to CII] $\lambda$2326 flux ratios in the type 2 quasars, which is due to having higher electron temperature, we point out that low metallicities are at odds with the findings of several earlier studies which concluded that the UV-optical emission line spectra of type 2 quasars are best explained by photoionization models with solar or super-solar gas metallicities (e.g., Robinson et al. 1987; V01; H08; Villar-Mart\'{i}n et al. 2008). The presence of pockets of low metallicity gas, within higher metallicity gaseous halos (Tornatore et al. 2007; Cassata et al. 2013) may provide a means to explain the apparently discrepant implied gas metallicities. | 14 | 4 | 1404.6434 |
1404 | 1404.4577_arXiv.txt | {Several trends have been identified in the prompt gamma-ray burst (GRB) emission: e.g. hard-to-soft evolution, pulse width evolution with energy, time lags, hardness-intensity{$/$}-fluence correlations. Recently \textit{Fermi} has significantly extended the spectral coverage of GRB observations and improved the characterization of this spectral evolution.} {We want to study how internal shocks can reproduce these observations. In this model the emission comes from the synchrotron radiation of shock accelerated electrons, and the spectral evolution is governed by the evolution of the physical conditions in the shocked regions.} {We present a comprehensive set of simulations of a single pulse and investigate the impact of the model parameters, related to the shock microphysics and to the initial conditions in the ejecta.} {We find a general qualitative agreement between the model and the various observations used for the comparison. All these properties or relations are governed by the evolution of the peak energy and photon indices of the spectrum. In addition, we identify the conditions for a quantitative agreement. We find that the best agreement is obtained for (i) steep electron slopes ($p\ga 2.7$), (ii) microphysics parameters varying with shock conditions so that more electrons are accelerated in stronger shocks, (iii) steep variations of the initial Lorentz factor in the ejecta. When simulating short GRBs by contracting all timescales, all other parameters being unchanged, we show that the hardness-duration correlation is reproduced, as well as the evolution with duration of the pulse properties. Finally, we investigate the signature at high energy of these different scenarios and find distinct properties -- delayed onset, longer emission, and flat spectrum in some cases -- suggesting that internal shocks could have a significant contribution to the prompt LAT emission.} {Spectral evolution is an important property of GRBs that is not easily reproduced in most models for the prompt emission. We find that the main observed features can be accounted for in a quantitative way within the internal shock model. However the current uncertainties on shock acceleration in the mildly relativistic regime and relativistic ejection by compact sources prevent us from deciding if one or several of the proposed scenario are viable. It may be possible by combining observations over the whole spectral range of \textit{Fermi} to identify in the future specific signatures imprinted by this uncertain underlying physics.} | Since the launch of the \textit{Swift} (2004) \citep{gehrels:04} and \textit{Fermi} (2008) satellites, there is a significantly growing sample of gamma-ray bursts (GRBs) with a known redshift and a well characterized gamma-ray prompt emission \citep[see e.g. the recent review by][]{gehrels:13}. The high-energy domain ($>100$ MeV) is currently explored by \textit{Fermi}-LAT \citep{atwood:09}. The sample of detected bursts is still small but has allowed the identification of several important spectral and temporal properties \citep{omodei:09,zhang:11,LATcatalog:13}, that are summarized in \refsec{sec:LAT}. In the soft gamma-ray range, the GRB sample is much larger and not limited to the brightest bursts. Thanks to its large spectral range (8 keV-40 MeV), \textit{Fermi}-GBM \citep{meegan:09} has significantly improved the description of the GRB properties in the keV-MeV range. This effort follows the results already obtained by several past or current missions, especially BATSE (Burst And Transient Source Experiment) on board the \textit{Compton Gamma-Ray Observatory} \citep{kaneko:06}, \textit{Beppo-SAX} \citep{guidorzi:11}, and \textit{HETE-2} \citep{lamb:04,sakamoto:05}. Based on this large set of observations, our current know\-ledge of the spectral and temporal properties of the GRB prompt soft gamma-ray emission is summarized in \refsec{sec:obs}. The standard GRB model associates the prompt gamma-ray emission to internal dissipation within an ultra-relativistic outflow ($\Gamma \ga$ 100) ejected by a new-born compact source \citep[see e.g.][]{piran:99}. The nature of the dissipation mechanism and of the associated radiative process remains to be identified. In order to account for the observed short time scale variability ($\sim$ ms), the internal shock model \citep{rees:94}, where variations of the bulk Lorentz factor lead to the formation of shock waves within the ejecta, was proposed for the extraction of the jet kinetic energy. The dissipated energy is distributed between protons, electrons, and magnetic field; the prompt GRB emission model is due to the synchrotron radiation of shock accelerated electrons, with an additional component due to inverse Compton scatterings. Detailed calculations of the expected light curves and spectra are available \citep{kobayashi:97,daigne:98,bosnjak:09,asano:11} and show a good agreement with observations except for a notable exception, the low-energy photon index, which is usually observed to be larger than the standard fast cooling synchrotron slope $-3/2$. Several solutions have been proposed, as the role of inverse Compton scatterings in Klein Nishina regime \citep{derishev:01,bosnjak:09,nakar:09,wang:09,daigne:11} or the magnetic field decay in the shocked region \citep{derishev:07,zhao:14}. Other mechanisms could also play a role. Thermal emission is expected at the photosphere when the ejecta becomes transparent for its own radiation. Depending both on the efficiency of the acceleration process and of the non-thermal emission above the photosphere, this emission could be bright \citep{meszaros:00,daigne:02,hascoet:13}. It produces in principle a narrow quasi-Planckian component \citep{goodman:86,peer:08,beloborodov:11}; however different possible sub-photospheric dissipation processes may affect the spectrum, especially due to the comptonization, so that it appears as non-thermal \citep{thompson:94,rees:05,peer:06,giannios:07,beloborodov:10,vurm:11,toma:11,veres:12,veres:13}. The peak energy is governed by a detailed balance between the emission/absorption and scattering processes \citep{vurm:13} and can reproduce the observed values (\citealt{beloborodov:13}, see however \citealt{zhang:12}). The lateral structure of the jet may also affect the photospheric spectrum \citep{lundman:13,lazzati:13}. Magnetized ejecta offer a third possibility. A large initial magnetization may play a major role for the acceleration of the jet to relativistic speed \citep[see e.g.][]{begelman:94,daigne:02b,vlahakis:03,komissarov:09,tchekhovskoy:10,komissarov:10,granot:11} and is already invoked for this reason in some scenarios where the emission is due to the photosphere and/or internal shocks. However, if the ejecta is still magnetized at large distance, magnetic reconnection can provide a new dissipation process \citep{spruit:01,drenkhahnspruit:02,lyutikov:03,gianniosspruit:05,zhangyan:11,mckinney:12}. Compared to the previous possibilities, this model cannot provide yet detailed predictions for the GRB light curves and spectra \citep[see the preliminary calculation of the temporal properties by][]{zhang:13}. The photospheric emission should be present in all scena\-rios, even if very weak. On the other hand, magnetic reconnection requires a large magnetization at large distance which may prevent internal shock formation and propagation \citep{giannios:08,mimica:10,narayan:11}. Therefore, depending on the magnetization in the emission site, only one of the two mechanisms should be at work. Recent observations of two components in the soft gamma-ray spectrum of a few bright \textit{Fermi}/GBM bursts, one being quasi-Planckian and the other being non-thermal \citep{ryde:10,guiriec:11,axelsson:12,guiriec:13}, indicate that both the photosphere and either internal shocks or reconnection may be indeed at work in GRBs \citep{hascoet:13}. Aside from interpreting the light curves and spectra, a successful theoretical model should also reproduce the observed spectral evolution with time, which is is mainly governed by the evolution of the peak energy of the spectrum. It can be related either to the physics of the dissipative mechanism in the outflow, or to the curvature of the emitting surface. In the first case, the spectral evolution is due to an intrinsic evolution of the physical conditions in the flow, whereas it is a geometrical effect (delay, Doppler shift) in the second case. The spectral evolution in a pulse associated to the curvature effect has been studied by several authors and does not agree with observations \citep{fenimore:96,dermer:04,shen:05,shenoy:13}. Then, the spectral evolution has to be understood from the physics of the dissipative mechanism and may therefore represent an important test to discriminate between the different possible prompt emission models listed above. Regarding the photospheric emission, the spectral evolution has been computed only in the case of non-dissipative photospheres \citep{daigne:02,peer:08}. As mentioned above, this model cannot reproduce the observed spectrum. In the case of a dissipative photosphere, the peak energy of the spectrum is fixed by a complex physics \citep{beloborodov:13}, which makes difficult a prediction of the spectral evolution. It is usually assumed that modulations in the properties at the base of the flow will lead to the observed evolution (see for instance \citet{giannios:07} in the case where the dissipation is associated to magnetic reconnection). However, dissipative photospheric models require that the dissipation occurs just below the photosphere for the spectrum to be affected. It is not obvious that a change in the central engine leading to a displacement of the photosphere will affect the dissipation process in the same way so that it remains well located. Therefore, it remains to be demonstrated that these models can reproduce the observed spectral evolution. In the context of an emission produced above the photosphere, several authors have investigated the time development of the photon spectrum without specifying the dissipation mechanism and relate the observed spectral evolution to the evolution of the electron/photon injection rate and/or the decaying magnetic field \citep[e.g.][]{liang:97,stern:04,asano:09,asano:11}. It is encouraging that a reasonable agreement with observations is found in some cases. To reach a final conclusion, it is however necessary to carry such a study in the context of a physical model for the dissipation, which gives a prescription for the accelerated electrons and the magnetic field. It is still out of reach for reconnection models due to the lack of any spectral calculation. It has been done for the internal shock model using a very simple spectral calculation including only synchrotron radiation \citep{daigne:98,daigne:02}. Since these early calculations, the observational description of the spectral evolution has improved a lot, as well as the modeling of the emission from internal shocks. Therefore, we examine here the quantitative prediction of this model for the spectral evolution in GRBs. For the first time the detailed dynamical evolution is combined with the calculation of the radiative processes, and the outcome is confronted to the large set of observed properties summarized in \refsec{sec:obs} (e.g. hard-to-soft evolution, pulse width evolution with energy, time lags, hardness-intensity/fluence correlation). In \refsec{sec:internalshocks}, we present our approach, which is based on the model developed in \citet{bosnjak:09}. Following \citet{daigne:11} we define three \textit{reference cases}, which are representative of the different possible spectral shapes in the keV-MeV range, and present a detailed comparison of their temporal and spectral properties with observations. Then we investigate in \refsec{sec:EffectMicrophysics} and \refsec{sec:EffectDynamics} the effect on our results of different assumptions for the microphysics and dynamics of the relativistic ejecta. The specific signatures in the \textit{Fermi}-LAT range are presented in \refsec{sec:LAT}. We discuss our results in \refsec{sec:discussion} and conclude in \refsec{sec:conclusions}. | \label{sec:conclusions} Motivated by the results from the \textit{Fermi} satellite which significantly extends the spectral coverage of the GRB phenomenon and improves particularly the spectral analysis of the prompt emission, we investigated in this paper the origin of the observed spectral evolution in GRBs. We presented the results of a set of numerical simulations of the GRB prompt emission in the framework of the internal shock model. We made a detailed comparison of the model predictions with the observed temporal and spectral GRB properties in the soft gamma-ray range. We focussed on the simplest case of a single pulse burst associated to the synchrotron radiation from shock-accelerated electrons in the internal shocks formed after the collision between a 'fast' and a 'slow' region in an ultra-relativistic ejecta. We considered three reference case with a duration of $2-3$ s, an isotropic radiated energy of respectively $1.9\times 10^{52}$, $1.3\times 10^{52}$ and $1.3\times 10^{51}\, \mathrm{erg}$, a peak energy of 730, 640 and 160 keV, and a low-energy photon index of -1.5, -1.1 and -0.7. We show that many observed properties or common trends -- namely (i) the pulse asymmetry, (ii) the energy dependent pulse asymmetry (evolution of the pulse width with energy channel), (iii) the time lags between the light curves in different energy channels, (iv) the hard-to-soft evolution within pulses, (v) the hardness-intensity correlation, (vi) the hardness-fluence correlation -- can be accounted for and are governed by the details of the spectral evolution, i.e. the evolution of the peak-energy and the spectral slopes. We showed that there is a qualitative agreement between the model results for our three reference cases and the large set of observations listed above. With a comprehensive set of simulations, we demonstrated that a quantitative agreement can be achieved under some constraints on the model parameters. We distinguished between the effects of the microphysics (details of the energy distribution in shocked regions) and the dynamical parameters (initial conditions in the outflow). We found that the agreement with the observed spectral evolution can be significantly improved if (i) the distribution of shock-accelerated electrons is steeper than what is usually assumed, with a slope $p\ga 2.7$; (ii) the microphysics parameters vary with the shock conditions in a manner that reduces the dependency of the peak energy on the shock conditions. It is illustrated here by the case where the fraction of accelerated electrons increases for stronger shocks; (iii) the initial variations of the Lorentz factor in the outflow are steeper. An additional advantage of this assumption is the increase of the efficiency of internal shocks; (iv) the relativistic ejection proceeds with a constant mass flux rather than a constant kinetic energy flux. A drawback of this last possibility is a reduced efficiency of the shocks. As the microphysics parameters are not well constrained by the current stage of shock acceleration modelling in the mildly relativistic regime relevant for internal shocks, and as the initial conditions in the outflow are also poorly constrained due to many uncertainties regarding the mechanism responsible for the relativistic ejection by the central engine, we cannot conclude if one of these four possibilities may be expected or should be preferred. We also specifically investigated the impact of the duration of the relativistic ejection, as many of the properties listed above are known to evolve with pulse duration. The internal shock model naturally predicts a larger peak energy for short pulses, and possibly a harder photon index due to a deeper Klein-Nishina regime for inverse Compton scatterings. We showed that -- in agreement with observations -- this leads to a hardness-duration correlation and to the following consequences: pulses become more symmetric, with almost no evolution of the pulse width with energy, and with very short or zero lags. The prompt emission from short GRBs could then be due to the same mechanism as in long GRBs, but for different model parameters due to the fact that all timescales are contracted, probably because of a different central engine. Finally, we investigated the signature at high-energy (\textit{Fermi}-LAT range). In this domain, the observed flux is made of the high-energy tail of the synchrotron component and a new component produced by inverse Compton scattering. A direct comparison with \textit{Fermi}-LAT results is not possible as LAT bursts are among the brightest whereas we have simulated here average pulses. However, we note a qualitative agreement with data: due to the evolving efficiency of the scatterings -- they usually occur in the Klein-Nishina regime at early times and enter the Thomson regime during the pulse decay -- the resulting emission at high-energy can differ significantly from the keV-MeV range; specifically, the rise of the light curve is delayed and the emission lasts longer. This leads to a U-shape curve when plotting time lags with respect to the low-energy channel as a function of energy, in agreement with GBM+LAT observations. However, we do not have a quantitative agreement : the onset of the high-energy light curve is not delayed enough. Interestingly, some of the effects listed above -- a steeper electron slope, a varying electron acceleration fraction, and especially steeper variations of the initial Lorentz factor -- have also a positive impact on the properties of the high-energy emission. The time-integrated spectrum at high-energy depends strongly on the efficiency of the inverse Compton scatterings. In some cases, it is found to be very close of the extrapolation of the MeV component, possibly with a cutoff at high-energy; in other cases, it clearly shows an additional component, which can either be rising (photon index greater than $-2$) or flat (photon index close to $-2$). As there are significant differences between the various scenarios discussed in the paper, this motivates a specific comparison to \textit{Fermi}-LAT bursts which will hopefully provide diagnostics to distinguish among the various theoretical possibilities. This study illustrates the capacity of the internal shock model to reproduce most of the observed properties of the GRB prompt emission related to the spectral evolution, both for long and short bursts. Our conclusions are limited by many uncertainties in the ingredients of the model, namely the details of the microphysics in mildly relativistic shocks and the initial conditions in the GRB relativistic outflows. However, in a more optimistic view, we showed that this poorly understood physics may have a detectable imprint in GRB data, which should allow for some progress in the future. | 14 | 4 | 1404.4577 |
1404 | 1404.2351_arXiv.txt | Flux calibration of spectra in reverberation mapping (RM) is most often performed by assuming the flux constancy of some specified narrow emission lines, which stem from an extended region that is sometimes partially spatially resolved, in contrast to the point-like broad-line region and the central continuum source. The inhomogeneous aperture geometries used among different observation sets in a joint monitoring campaign introduce systematic deviations to the fluxes of broad lines and central continuum, and intercalibration over these data sets is required. As an improvement to the previous empirical correction performed by comparing the (nearly) contemporaneous observation points, we describe a feasible Bayesian method that obviates the need for (nearly) contemporaneous observations, naturally incorporates physical models of flux variations, and fully takes into account the measurement errors. In particular, it fits all the data sets simultaneously regardless of samplings and makes use of all of the information in the data sets. A Markov Chain Monte Carlo implementation is employed to recover the parameters and uncertainties for intercalibration. Application to the RM data sets of NGC 5548 with joint monitoring shows the high fidelity of our method. | Reverberation mapping (RM) is a well-established technique for the study of broad-line regions (BLRs) in active galactic nuclei (AGNs) with broad emission lines (\citealt{Blandford1982, Peterson1993}). With appropriate analysis, RM experiments divulge the geometry, kinematic, and ionization structure information of BLRs (e.g., \citealt{Brewer2011, Pancoast2012, Li2013}). Over the past two decades, the BLR size derived using the time delay between the continuum variation and the broad emission line response has been utilized with great success to measure the mass of the central supermassive black hole by combining it with the width of the broad emission line (e.g., \citealt{Peterson2004}). The tight relationship between BLR sizes and optical luminosities of AGNs plays a key role in the demography of supermassive black holes in large AGN surveys (e.g., \citealt{Bentz2013}, and references therein). At present there are $\sim$50 nearby Seyfert galaxies and quasars with RM measurements in the literature (e.g., \citealt{Bentz2013}), although a huge amount of effort has been invested in RM experiments. In practice, an RM campaign is quite observationally intensive and requires monitoring an object over a sufficient period with reasonable temporal resolution. Such high demand of time interval and sampling leads RM programs to be commonly undertaken by cooperative observations at multiple observatories, such as the well-known AGN Watch Project (\citealt{Peterson2002}) and MDM campaigns (\citealt{Denney2010, Grier2012}). Spectra calibration is most often based on the assumption that \Oiii~$\lambda$5007 line remains constant in flux over the timescale of interest and all spectra are scaled to an adopted absolute flux of \Oiii~$\lambda$5007, which can be measured on photometric nights (\citealt{Peterson1991, vanGroningen1992}). The problem that arises with using \Oiii~$\lambda$5007 for such a calibration is that its emission region (narrow-line region; NLR) is sometimes spatially resolved and the size is most likely comparable with or even larger than the aperture size, in contrast to the effectively point-like BLR and central continuum sources. Consequently, the inhomogeneous aperture geometries used among different observation sets in a joint monitoring campaign admit different amounts of light from the NLR, and therefore introduce systematic deviations to the fluxes of broad emission lines (\citealt{Peterson1995}). Similarly, this effect also influences the central continuum fluxes contaminated by the host galaxy starlight. \cite{Peterson1995} proposed an empirical correction to such an aperture effect by adopting one of the data set as standard, and applying a multiplicative scale factor and an additive flux adjustment to the other sets to bring the closely spaced measurements from the two sets into agreement. In reality, it is always impractical to base the correction on exactly contemporaneous observations. One has to relax such strict simultaneity and instead use pairs of observations that are closely separated (usually by more than one days), depending on the sampling of each data set. This apparently degrades the highest achievable temporal resolution. In this {\em Letter}, we describe a novel method for intercalibration of reverberation mapping data that obviates the need for (nearly) contemporaneous observations, naturally incorporates physical models of flux variations, and fully takes into account the measurement errors. The method is based on Bayesian statistics and is sufficiently elastic to automated program manipulation. | We propose a feasible Bayesian method for spectral intercalibration in a joint monitoring campaign based on the assumption of flux constancy of some specified narrow emission line (e.g., \Oiii~$\lambda5007$). Compared with the previous empirical method comparing the closely spaced measurement pairs, our new method obviates the requirement for (nearly) contemporaneity of observations and takes into account the measurement errors naturally. The Bayesian approach enables us to perform intercalibration on all the data sets simultaneously and self-consistently regardless of sampling rates, and therefore can cope with poorly sampled data. Application to the RM database of NGC 5548 from the AGN Watch Project shows the fidelity of our method and its capability to yield appropriate intercalibration where the previous method encountered difficulties. In conclusion we propose a road map for complete spectral calibration in RM campaigns in which one or more emission lines with constant flux are present: first employ the algorithm described by \cite{vanGroningen1992} to perform relative scaling based on the adopted emission line, which takes into account the zero-point wavelength-calibration errors between individual spectra and resolution differences; and then employ our method to perform intercalibration to correct for the effect of inhomogeneous apertures. | 14 | 4 | 1404.2351 |
1404 | 1404.2398_arXiv.txt | We have analyzed the distributions in the color-magnitude diagram (CMD) of a large sample of face-on galaxies to minimize the effect of dust extinctions on galaxy color. About 300 thousand galaxies with $log(a/b) < $ 0.2 and redshift $z < 0.2$ are selected from the SDSS DR7 catalog. Two methods are employed to investigate the distributions of galaxies in the CMD including 1-D Gaussian fitting to the distributions in individual magnitude bins and 2-D Gaussian mixture model (GMM) fitting to galaxies as a whole. We find that in the 1-D fitting only two Gaussians are not enough to fit galaxies with the excess present between the blue cloud and the red sequence. The fitting to this excess defines the centre of the green-valley in the local universe to be $(u-r)_{0.1} = -0.121M_{r,0.1}-0.061$. The fraction of blue cloud and red sequence galaxies turns over around $M_{r,0.1} \sim -20.1$ mag, corresponding to stellar mass of $3\times10^{10}M_\odot$. For the 2-D GMM fitting, a total of four Gaussians are required, one for the blue cloud, one for the red sequence and the additional two for the green valley. The fact that two Gaussians are needed to describe the distributions of galaxies in the green valley is consistent with some models that argue for two different evolutionary paths from the blue cloud to the red sequence. | The distribution of galaxies in the color-magnitude diagram (CMD) provides a powerful tool to investigate the evolution of galaxy populations. A remarkable feature of the CMD is a robust bimodality, which divides the galaxy population into a ``blue cloud'' (or blue sequence) and a ``red sequence''. The bimodality is seen in the optical colors \citep{Strateva_2001, Blanton_2003c}, UV$-$optical colors \citep{Wyder_2007}, the 4000 \r{A} break ($D_n4000$; \citealt{Kauffmann_2003a}), and spectral type \citep{Madgwick_2002}. Galaxies in ``red sequence" are quiescent, bulge-dominated galaxies \citep{Blanton_2009}, while the ``blue cloud" is characterized by star-forming, disk-dominated galaxies. Between these two sequences, there is a region called the ``green valley". \citet{Baldry_2004} explored the distribution of galaxies in the $(u-r)$ versus $M_r$ diagram for low-redshift SDSS samples. Their galaxies separate into ``blue cloud" and ``red sequence", and the distribution of $(u-r)$ color at each absolute magnitude bin is well fitted by the sum of two Gaussians. However, \citet{Wyder_2007} showed that the $(NUV-r)$ color distribution at each $M_{r}$ can not be fitted well by the sum of two Gaussians due to an excess of galaxies between the blue and red sequences. They utilized Balmer decrements and the Dust-SFH (star formation history)-Color relation \citep{Johnson_2006} to correct the extinction of each galaxy, there still remain galaxies in the green valley region between two sequences. Thus, galaxies in the green valley region may not be a simple mixture of blue and red galaxies. The understanding of the CMD color bimodality is also complicated by galaxy dust extinction. \citet{Salim_2009} found that many green valley galaxies are simply dust-obscured actively star-forming (SF) galaxies. However, there still exist 24 $\mu$m detected galaxies, some with LIRG-like luminosities, which have little current SF. They belong to green valley or even the red sequence because of their SF history, not just dust reddening. The CMD bimodality is already in place at $z\sim1$ \citep{Cooper_2006}, with color becoming bluer at higher redshift \citep{Blanton_2006, Willmer_2006}. Based on the DEEP2 and COMBO-17 surveys, Faber et al. (2007) argued that the number density of blue galaxies is more or less constant from $z\sim1$ to 0, while the number density of red galaxies has increased. This work supports that the red sequence has grown in mass by a factor of 3 since $z\sim1$. A plausible scenario is that the growth of red galaxies was triggered by quenching star formation in blue galaxies, which caused them to migrate into the red sequence \citep{Bell_2004}. In addition, galaxies may also be moving from the lower end of the red sequence to the blue cloud through accreting gas-rich dwarf galaxies \citep{Faber_2007}. Studies of galaxy morphologies show that red sequence is dominated by spheroidal galaxies with S{\'e}rsic index n=4, while the blue cloud is occupied by disk-dominated galaxies with S{\'e}rsic index n=1 \citep{Driver_2006}. \citet{Mendez_2011} investigated the morphologies of green valley galaxies from the AEGIS survey and found that most green valley galaxies are not classified as mergers and that the merger fraction in the green valley is lower than that in the blue cloud. \citet{Lackner_2012} presented a set of bulge-disc decompositions for a sample of 71825 SDSS main-sample galaxies and found that the majority of green valley galaxies are bulge+disc galaxies, and that the integrated galaxy color is driven by the color of galaxy disks. In general, blue galaxies with star formation being quenched will evolve from the blue cloud to the red sequence, passing through the green valley that thus represents an intermediate phase of this quench process. Different mechanisms have been proposed to cease star formation in blue galaxies, such as mergers \citep{Bell_2004, Hopkins_2010}, AGN feedback \citep{Croton_2006, Martin_2007, Schawinski_2010}, morphological quenching \citep{Martig_2009}, cold flows accretion and shock heating \citep{Dekel_2006, Cattaneo_2006} (see \citealt{Peng_2010} for a reccent review). \citet{Peng_2010,Peng_2012} investigated the quenched fraction of galaxies as a function of local density, stellar mass, and redshift. They parameterized galaxy quenching as fully separable ``environment quenching" and ``mass quenching", which are directly associated with the quenching processes of satellite and central galaxies in group. This model is successful in predicting the mass function of passive and star-forming galaxies. These effects may also be reflected in the color-magnitude distribution of galaxies, we will try to find some evidence of these affects in our analysis. In this paper, we focus on fitting CMD by the use of different methods. Our study aims at a better understanding of the ``green valley" and may answer the question whether the ``green valley" is dominated by one component. Since different quenching mechanisms will produce different distribution of galaxies in color-magnitude space, the fine structure of CMDs may also provide valuable clues about galaxy quenching mechanisms. Previous work always focus on the whole galaxy population without identification, which may conceal the fine structure of CMDs due to dust extinction, so we selected a nearly face-on galaxy sample to minimize the effect of dust. This paper is organized as follows. Section 2 describes the sample selection in this work. In Section 3, we investigate the color-magnitude distribution and show 1D and 2D Gaussian fitting results. Section 4 discuss the implication of the results. Section 5 is the conclusion. Throughout this paper, we assume a flat $\Lambda$CDM cosmology with a matter density $\Omega_{m}$=0.3, cosmological constant $\Lambda$=0.7 and Hubble's constant $H_0$=100$kms^{-1}Mpc^{-1}$ (i.e., $h=1$). | In this study we carried out Gaussian profile fittings to the CMD through two methods, i.e., the fitting to distributions in individual magnitude bins and to distributions in 2-D as a whole. The two Gaussian components in the CMD green valley are homologous to the excess in color distribution fitting, and explain why we cannot fit the color histogram well only two Gaussians. The green valley region is not dominated by a single Gaussian, but at least two Gaussian components. Galaxies in the green valley are composed of these two components plus the Gaussian tails of blue and red galaxies. By fitting the $(u-r)$ color distribution in each magnitude bin, our results are different from those of \citet{Baldry_2004}, but agree with the $(NUV-r)$ results of \citet{Wyder_2007}. As we select only nearly face-on galaxies which are not considered by \citet{Baldry_2004}, we attribute the difference to dust extinction, which seriously reddens the color of blue cloud galaxies and covers the excess in the green valley region. \citet{Kauffmann_2003a} showed that galaxies tend to divide into two distinct groups below and above a stellar mass of $3\times10^{10}M_\odot$. Galaxies below this mass limit tend to have younger stellar populations, while more massive galaxies tend to be older. We find that the fraction of blue cloud and red sequence are equal to each other around $M_{r,0.1}$$\sim$$-20.1\ mag$, corresponding to stellar mass about $M^*$ ($10^{10.5}M_{\sun}$). As shown in Figure 3, our result is consistent with their conclusion very well. If we assume that the density of galaxies represents their evolution time scale in color-magnitude space, the GMM results would intuitively show us the evolutionary paths of different galaxies: the components in Figure 6 would represent galaxies in different evolving phases, and the inclination of ellipses' major axes may suggest the galaxy evolving direction in color-magnitude space. The faint component in the green valley may be the early-quenching population, which quenched while galaxy are still small and grow mass along red sequence via ``dry" mergers \citep{Faber_2007}. In the faint end of the red sequence, these red low-luminosity galaxies tend to be in overdense regions \citep{Blanton_2006}, environment possibly plays a very important role in their evolution. If the two components in the green valley are dominated by two independent quenching processes, this scene would agree with \citet{Peng_2010} very well. According to the model of \citet{Peng_2010}, galaxies in the faint component are dominated by ``environment quenching", and galaxies in the bright one are dominated by ``mass quenching", which are associated with the quenching processes of satellite and central galaxies in group \citep{Peng_2012}. Since these two effects are fully independent of each other, they may produce the two independent Gaussian components we found in the CMD. As shown in Figure 7, the mean stellar mass of the bright component in the green valley region is about the characteristic mass $M^*$. The characteristic mass $M^*$ also corresponds to a dark halo mass $M_{shock}\sim10^{12}M_\odot$ based on the model of \citet{Dekel_2006}. In their model, galaxies with $M_{halo}\gtrsim10^{12}M_\odot$ will generate a steady shock in the gas accreting onto dark matter halo, the shock heats the gas and absolutely quenches star formation when AGNs begin to work. This process is strongly related to the mass of galaxy, which may dominate the evolution of bright galaxy component in green valley. \citet{Wong_2012} selected a local post-starburst galaxies (PSGs) sample from SDSS, those PSGs occupy the low-mass end of the "green valley" below the transition mass within the colour-stellar mass diagram (the same position as the faint component of green valley in Figure 7). They proposed those PSGs represent a population of galaxies which is rapidly transitioning between the star-forming and the passively evolving phases. \citet{Mendel_2013} select a sample of young passive galaxies from SDSS, which is identified based on the contribution of A-type stars to spectra and the relative lack of ongoing star formation. Most of these recently quenched galaxies have a stellar mass $> 10^{9.5}M_\odot$ and are predominantly early-type systems. \citet{McIntosh_2013} studied the recently quenched ellipticals (RQEs) with stellar mass $>10^{10}M_\odot$ and found a number of RQE properties are consistent with these galaxies being new remnants from a gaseous major merger. Their studies show that the low- and high- mass galaxies in green valley are very different, and suggest the green valley are dominated by different quenching processes, supporting our GMM fitting results. In Figure 8, it is notable that the blue peak of color bimodality is dominated by the bright component of green valley (the cyan line) for galaxies $M_r < -20.25$ $mag$, which imply that the blue galaxy population with $M_*>10^{10.5}M_\odot$ is distinct from galaxies which traditionally thought as blue cloud. This case is consistent with \citet{Schawinski_2014}, which proposed that the early- and late-type galaxies in green valley have two different evolutionary pathways. The evolving early-type galaxies generally have stellar mass $M_*>10^{10.5}M_\odot$, rapidly quenching star formation, moving out the blue cloud, into the green valley and to the red sequence as fast as stellar evolution allows \citep{Schawinski_2014}. It is still unclear about the quenching mechanisms of galaxies, differentiating such mechanisms requires more evidences which are beyond the scope of this work. | 14 | 4 | 1404.2398 |
1404 | 1404.2217_arXiv.txt | {} {We use optical integral field spectroscopy and 8\mi\ and 24\mi\ mid-IR observations of the giant \hii\ region NGC~588 in the disc of M33 as input and constraints for two-dimensional tailor-made photoionisation models under different geometrical approaches. We do this to explore the spatial distribution of gas and dust in the interstellar ionised medium surrounding multiple massive stars.} {Two different geometrical approaches are followed for the modelling structure: i) Each spatial element of the emitting gas is studied individually using models which assume that the ionisation structure is complete in each element to look for azimuthal variations across gas and dust. ii) A single model is considered, and the two-dimensional structure of the gas and the dust are assumed to be due to the projection of an emitting sphere onto the sky.} {The models in both assumptions reproduce the radial profiles of \hb\ surface brightness, the observed number of ionising photons, and the strong optical emission-line relative intensities. The first approach produces a constant-density matter-bounded thin shell of variable thickness and dust-to-gas ratio, while the second gives place to a radiation-bounded thick shell sphere of decreasing particle density. However, the radial profile of the 8\mic/24\mic\ IR ratio, depending on the gas and dust geometry, only fits well when the thick-shell model is used. The resulting dust-to-gas mass ratio, which was obtained empirically from the derived dust mass using data from {\em Spitzer}, also has a better fit using the thick-shell solution. In both approaches, models support the importance of the low surface-brightness positions on the integrated spectrum of the nebula, the chemical homogeneity, the ionisation-parameter radial decrease, and the robustness of strong-line methods to derive the equivalent effective temperature in extended regions. These results must be taken with care in view of the very low extinction values that are derived from the IR, as compared to that derived from the Balmer decrement. Besides, the IR can be possibly contaminated with the emission from a cloud of diffuse gas and dust above the plane of the galaxy detected at 250~\mic\ {\em Herschel} image.} {} | High surface brightness\hii\ regions in the optical spectral range are luminous tracers of both the physical properties and chemical abundances of the interstellar medium (ISM) of the galaxies where they are located. Indeed, \hii\ regions in star-forming galaxies are one of the most widely objects used to find out these properties throughout the Universe. In non-resolved \hii\ regions the optical spectra are collected by means of integrated fibres or long-slit techniques, while integral field spectroscopy (IFS) allows spectral measurements with 2D spatial resolution for objects in the Local Universe. The studies of the structure of \hii\ regions then become more complex, and therefore more elaborated depictions become now necessary. This opens the gate to a better understanding of the interplay between stars, dust, and gas appearing in different spatial positions and to assess if the assumptions made to study integrated observations in non-resolved objects lead to accurate determinations of their properties. The most challenging issues that come from the study of the two-dimensional structure of \hii\ regions appears as a consequence of the lack of spatial uniformity in many of their properties. For instance, it is shown by \cite{ercolano} and \cite{jamet08} that the distribution of the ionising stars in relation to the gas alters the ionisation structure and the electron temperature. Other authors (e.g., \citealt{gia04}) point out the relevance of gas density inhomogeneities, which also lead to inhomogeneities in the other physical properties derived from optical spectroscopy. \cite{marcelo} also show how the ionisation structure of the gas depends on the amount of available gas in different matter-bounded configurations, which gives place to large fractions of escaping ionising photons and then affects the observed emission-line ratios used to derive the physical properties and chemical abundances of the gas. Photoionisation models constitute powerful tools for the interpretation of the physics involved in \hii\ regions. Under appropriate geometrical considerations, models allow us to relate the collected observational information to both quantitative and qualitative characterisations of the studied objects and to derive their physical properties and chemical abundances. Models that are 3D are the most suitable appliances to spatially disentangle the effects of different ionising sources on the non-uniform surrounding gas. Nevertheless, the lack of observational information about the distribution of gas and stars along the line of vision prevents this kind of 3D model from fitting many optical IFS data on \hii\ regions most of the time. On the other hand, other techniques based on photoionisation models try to describe the observed spatial variations in ionised gaseous nebulae as a consequence of the projection of a 3D structure on a plane. This is the case of the codes, NEBU\_3D \citep{NEBU3D} or {\sc Cloudy 3D} \citep{C3D}. In a novel approach to reproduce IFS data, \cite{mod595} use 1D photoionisation models to fit the optical IFS \citep{relano10} and 8~\mic\ and 24~\mic\ mid-IR {\em Spitzer} bands properties of the Giant \hii\ Region (G\hii R) NGC~595 in the disc of M33. In that work, different annuli around the ionising source are defined in the area covered by a mosaic of several integral field unit (IFU) pointings, and their measured integrated properties are later fitted by the models. These models depict a uniform metallicity across a thin shell, whose optical and IR observed structure can be explained with azimuthal variations (i.e., in the plane of the galaxy) in some of the properties of the \hii\ region, as in the dust-to-gas ratio and the matter-bounded geometry. Since both azimuthal variations and projection effects are expected to co-exist as causes of the observed spatial variations throughout the distribution of gas and dust in \hii\ regions, it is necessary to explore both model strategies in well-known and characterised objects, as is the case of the G\hii R NGC~588, which is also in M33. Two-dimensional observations of NGC~588 in both optical and mid-IR are described in \cite{monreal11} (hereafter MI11); thus, a 2D observational characterisation of the properties of the ionised gas and the hot dust is possible. Besides, previous studies on this G\hii R, which are based on ground and spacecraft imaging in different bands from the UV up to the NIR, were used by \cite{jamet04} to study the location and nature of the ionising stellar population. The aim of this work is to study the possible causes of the observed spatial variations across NGC~588 for both the optical and mid-IR properties to test the solidness of the strong-line methods used in integrated observations. This is done to derive physical properties and chemical abundances by means of photoionisation models of the well-studied G\hii R NGC~588. To do so, we took two different assumptions: i) an improved version of the approach employed by \cite{mod595} for NGC~595, based on different models for the individual observed spatial elements which are used to explain the observations that are caused by azimuthal variations of its properties throughout the field of view and ii) a single model projected onto the sky which is used to explain the variations as a consequence of the perspective. In the next section, we describe the 2D structure of NGC~588 and the data sampling of the spatial distribution for both optical and mid-IR data studied in MI11 and the derivation of the integrated dust-to-gas ratio from dust temperature and total H{\sc i} mass. In Section 3, we present our models and in Section 4, we discuss our results. Finally, we summarise our results and conclusions in Section 5. \begin{figure*} \begin{minipage}{180mm} \centerline{ \psfig{figure=flujo_hb_elipse.ps,width=6cm,bb=55 30 525 565,clip=} \psfig{figure=flujo_8m_elipse.ps,width=6cm,bb=55 30 525 565,clip=} \psfig{figure=flujo_24m_elipse.ps,bb=55 30 525 565,width=6cm,clip=}} \label{annuli} \caption[]{Elliptical annular regions utilised to extract the radial profiles on top of the non-calibrated \hb\ (\emph{left}), 8~$\mu$m (\emph{middle}) and 24~$\mu$m (\emph{right}) emission distributions (see MI11 for more detailed plots in these bands). Units are arbitrary in logarithmic scale. The origin of coordinates is located at R.A. (J2000): 1h 32m 45.7s, DEC. (J2000): +30$^{\circ}$ 38'' 55.1' and is marked with a yellow "X". The three elliptical regions considered through the paper are labelled as "A" (\emph{red}), "B" (\emph{blue}), and "C" (\emph{green}) and sample the SE, W, and NE part of NGC~588, respectively. In all images, N points to up and E to left.} \end{minipage} \end{figure*} | The two-dimensional optical and mid-IR spatial structure of the G\hii R NGC~588 in the disc of M33 was studied by means of one-dimensional tailor-made photoionisation models. This object constitutes an appropriate target to study the spatial interplay between gas, dust, and stars and, thus, to explore the validity of different methods used for integrated observations because it is very well characterised in observations at several bands. The observational source for this study was the work made by MI11, who describe PMAS - CAHA 3.5 m optical IFU observations and compare them with {\em Spitzer} 8~\mi\ and 24~\mi\ mid-IR images. For this work, we also used {\em Spitzer} integrated data at 70~\mi\ and 160~\mi\ to derive the dust-to-gas ratio and {\em Herschel} 250~\mi\ to study the presence of diffuse dust in positions outside the G\hii R. To analyse the spatial variation of the observed main emission-line ratios and the IR emission, we followed the same procedure as \cite{mod595} for NGC~595 in which elliptical annuli were defined around the emission-peak. Although NGC~588 presents a ring-like morphology as seen in the \ha\ image, it does not have the same axial symmetry as NGC~595, so different regions were defined to reproduce the different observed patterns. Then, the emission in the three defined regions was co-added at different angular distances to the ionising clusters. Finally, the resulting radial profiles for \hb\ surface brightness and the emission-line ratios of [\oii]/\hb, [\oiii]/\hb, [\nii]/\ha, and [\sii]/\ha\ were taken to constrain the models that were made for the different annuli in the three regions. We adopted two different model strategies to find out the nature of the observed spatial variations across the field of view in this GH{\sc ii}R. On one hand, we considered each spatial element as having a complete ionisation structure that could be modelled independently. This model strategy is an improved version of the methodology described in \cite{mod595} for NGC~595 and accounts above all for azimuthal spatial variations of both the observed and the derived properties in two dimensions. The resulting geometry in these models consists of a hollow ellipsoid of revolution, whose MB thin-shell has most of the emission of the G\hii R. In this geometry, the optical emission is fitted by assuming a decreasing thickness and an increasing dust-to-gas ratio where most of the ionising photons are absorbed by dust and an additional fraction leaks from the nebula. On the other hand, we considered that the 2D observed properties of the nebula are due to the projection of a sphere onto the sky. This was modelled in a single photoionisation model using the {\sc pyCloudy} code \citep{C3D}. The resulting geometry in this model consists of a RB spherical thick-shell with decreasing particle density. In both approaches, both the \hb\ surface brightness and the main optical emission-line ratios were reproduced. The models also reproduce the radial variation of the {\em Spitzer} emission bands at 8~\mi\ and 24~\mi. However, the ratio between them is underestimated in the outer annuli in the thin-shell geometry due to the enhancement of the dust-to-gas ratio, which makes the emission at 24\mi\ to increase there. On the contrary, this ratio is well reproduced by the thick-shell model. We thus conclude that both projection and azimuthal spatial variations are possibly present in the properties of the gas and the dust in NGC~588. However, as the thick-shell approximation achieves a better fit to the mid-IR observations, the projection effects dominate over the azimuthal ones. This also shows the importance of fitting both optical and mid-IR spatial properties to better understand the spatial distribution of the region. A single-star SED, which is equivalent to the observed stellar clusters as demonstrated by \cite{jamet04}, was used in the models. This does not allow a detailed study of the inner ionisation structure in the G{\hii}R. However, the oxygen gas-phase abundance derived by the models [12+log(O/H) = 8.16, thin-shell, 8.3, thick-shell] for both assumptions is uniform across the nebula and agrees with the values measured by several authors (\citealt{vilchez88}, \citealt{jamet05}), who use the electron temperature based on optical collisionally excited emission lines. The conclusions are made by MI11. Under the assumption of a chemical homogeneity, we reproduced the very high values of \rdostres\ in the outer annuli of regions A and B only in the thin-shell models without considering high-velocity gas-shocks. Therefore, contrary to the rest of the nebula, azimuthal spatial variations of the geometry and dust-to-gas ratio are possibly behind the behaviour of these lines at these distances. According to both models, the estimation of $T_*$ from the $\eta'$ parameter is quite uniform across the nebula despite the high variation of several emission-line ratios involving both high- and low-excitation emission lines. This indicates that this method is very robust and independent of the assumed geometry of the models. However, this result needs to be confirmed with more observations of the spatial distribution of this parameter. The MB thin-shell models fit the optical observables by assuming an increasing dust-to-gas ratio, which go up to much higher values than that derived for the integrated nebula from IR emission at different bands. This estimation points to a value close to the standard Galactic value and was obtained empirically from the total dust mass, as found using the dust temperature derived from the ratio between 70~\mi\ and 160~\mi, and the total H{\sc i} mass. This value is consistent with the dust-to-gas ratio assumed in the thick-shell approach. However, the extinction value derived from this assumption is much lower than the estimates from the optical Balmer decrement from different authors. The 8 \mic\ image around the area of the G\hii R shows structures of the dust, which are not associated with the {\hii} region and which are confirmed with the {\em Herschel} image at 250 \mic. The presence of a cloud of diffuse gas and dust above the plane of the galaxy at the same position of NGC 588 could also explain the high values of the 8\mic/24\mic\ ratio. | 14 | 4 | 1404.2217 |
1404 | 1404.4027_arXiv.txt | Photospheric electric fields, estimated from sequences of vector magnetic field and Doppler measurements, can be used to estimate the flux of magnetic energy (the Poynting flux) into the corona and as time-dependent boundary conditions for dynamic models of the coronal magnetic field. We have modified and extended an existing method to estimate photospheric electric fields that combines a poloidal-toroidal (PTD) decomposition of the evolving magnetic field vector with Doppler and horizontal plasma velocities. Our current, more comprehensive method, which we dub the ``{\bf P}TD-{\bf D}oppler-{\bf F}LCT {\bf I}deal'' (PDFI) technique, can now incorporate Doppler velocities from non-normal viewing angles. It uses the \texttt{FISHPACK} software package to solve several two-dimensional Poisson equations, a faster and more robust approach than our previous implementations. Here, we describe systematic, quantitative tests of the accuracy and robustness of the PDFI technique using synthetic data from anelastic MHD (\texttt{ANMHD}) simulations, which have been used in similar tests in the past. We find that the PDFI method has less than $1\%$ error in the total Poynting flux and a $10\%$ error in the helicity flux rate at a normal viewing angle $(\theta=0$) and less than $25\%$ and $10\%$ errors respectively at large viewing angles ($\theta<60^\circ$). We compare our results with other inversion methods at zero viewing angle, and find that our method's estimates of the fluxes of magnetic energy and helicity are comparable to or more accurate than other methods. We also discuss the limitations of the PDFI method and its uncertainties. | In this paper, we use a specific set of anelastic pseudo-spectral \texttt{ANMHD} simulations \citep{Fan1999,Abbett2000, Abbett2004} of an emerging magnetic bipole in a convecting box \citep{Welsch2007} to test our improved electric field inversion technique. From \texttt{ANMHD} magnetic fields and plasma velocities, we know the actual electric fields, which we can compare to electric fields derived using the simulation's evolving magnetic field. In the past, this \texttt{ANMHD} simulation has been used for several studies of velocity field inversions \citep{Welsch2007, Schuck2008}, as well as in test-cases for the first electric field inversions \citep{Fisher2010, Fisher2012}. \cite{Fisher2012} showed that the PDFI-solution significantly improves the accuracy of the derived PI-solution, and the improvement from the knowledge of the Doppler velocity (PDI) is significantly more important than that of the horizontal velocity (PFI), at least in this example of magnetic flux emergence. In this paper, we perform a series of improvements to PDFI method beyond \cite{Fisher2012}: we expand the derivation of the non-inductive contribution to non-normal viewing angles, introduce spherical coordinates for the PTD solution (Appendix~\ref{sphere}) and significantly speed up the Poisson equation solutions by using \texttt{FISHPACK}. With the above upgrades, we find that the PDFI method yields a good estimate of the electric field, Poynting and helicity fluxes and is ready to be applied routinely to the observed vector magnetograms. We use a pair of \texttt{ANMHD} vector magnetograms, separated by $\Delta t=250$ s with a pixel size of $\Delta s=348.36$ km, and a LOS velocity map, observed at a specific viewing angle, to derive the electric field and the vertical Poynting flux for a given set of parameters (see Table~\ref{tparam}). To estimate the horizontal velocity field, $(V_x, V_y)$, we use the Fourier local-correlation tracking (\texttt{FLCT}) technique \citep{Welsch2004, Fisher2008} (\verb+http://solarmuri.ssl.berkeley.edu/~fisher/public/software/FLCT/C_VERSIONS/+). We denote the Gaussian window size scale (a parameter in FLCT) by $\sigma_{FLCT}$. To suppress noisy behaviour, we only calculate velocities where $|B_z|>370$ G, or $5\%$ of maximum $|B_z|$ \citep{Welsch2007}. In Table~\ref{tparam} we summarize all the input and output variables of the PDFI run and their typical ranges. Further, we vary the observed viewing angle $\theta$ within $[0,60]^{\circ}$-range to estimate the accuracy of the PDFI method at non-zero viewing angles. We also vary the other free parameters: the number of iterations in the perpendicularization technique ($N_{iter}$, see \S~\ref{vecEPI}), the width of the Gaussian window used in the \texttt{FLCT} estimate for the horizontal velocity ($\sigma_{FLCT}$, see below) and the width of polarity inversion line ($\sigma_{PIL}$, see \S~\ref{DEFzero}) to find the best values of the parameter set that yields the most accurate electric field solution (Right column). For the remainder of this paper, to assess the performance of our methods, i.e. of the reconstruction $u'$ of the \texttt{ANMHD} variable $u$, we use the following metrics: (1) a fraction of the integrated total $\displaystyle f(u,u')=\frac{\sum_i u'}{\sum_i u}$, (2) the slope or linear coefficient in the least-squares, polynomial fit, $a(u,u')$: $u' \approx a_0+a(u,u')*u$, (3) the linear Pearson correlation coefficient $\rho(u,u')=\frac{{\rm cov}(u,u')}{\sigma_u \sigma_{u'}}$, where $\sigma_u$ is a standard deviation of $u$, and (4) the normalized error of $u'$, $Err.=\sigma(u'-u)/\sigma_{u}$, $Err.=0.1$ means that the error of the reconstruction is $10\%$ relative to the characteristic range of $u$. Note that to exclude the weak-field background of $u$, where the observed magnetic field (e.g. in HMI/SDO) tends to be noisy, we estimated $f,a, \rho$ and $err.$ in locations where $|B|=({B_x}^2+{B_y}^2+{B_z}^2)^{1/2}>370$ G. In the HMI vector magnetogram case, a similar threshold will be determined from estimated errors in the magnetic field values. To summarize, the ideal reconstruction $u'$ of variable $u$ satisfies: $f(u,u')=1$, $a(u,u')=1$, $\rho(u,u')=1$ and $Err.=0.0$. \begin{table*}[h] \caption{Input and output variables of a set of PDFI runs for the \texttt{ANMHD} test case and their typical range. To test the PDFI method we varied the parameters within a shown range. We then used the best-value parameter shown in the parenthesis as a default.} \begin{center} \begin{tabular}{ccc} {\it Input} & Description & Observed Range \\ \hline ${B_{t,x},B_{t,y},B_{LOS}} (x,y)$ & Magnetic field & $[-6000,6000]$ Gauss \\ ${V_{LOS}}(x,y)$ &Doppler velocity & $[-0.5,0.2]$ km/s \\ $\theta$ & Viewing angle & $[0,60]^{\circ}$\\ \hline {\it Output} & & \\ ${V_{x},V_{y}}(x,y)$ & \texttt{FLCT} velocity field & $[-0.4,0.4]$ km/s \\ ${E_{x},E_{y},E_{z}}(x,y)$ & PDFI electric field & $[-1,1]$ V/cm \\ ${S_{z}(x,y)}$ & Poynting flux & $[-2,6]\times 10^{10}$ ergs/(s cm$^2$) \\ \hline {\it Parameters} & & Range (Best value) \\ $N_{iter}$ & No. of iterations in ${\bf E}\cdot{\bf B}$=0& $[0,50]$ iterations $(25)$ \\ $\sigma_{FLCT}$ & Gaussian window width &$[0,15]$ pixels $(15)$ \\ $\sigma_{PIL}$ & PIL width & $[0,2]$ pixels $(1)$\\ \end{tabular} \label{tparam} \end{center} \end{table*} \begin{figure*}[htb] \centering \resizebox{1.1\hsize}{!}{\includegraphics[angle=0]{edotb_sigmapil_iter.pdf}} \caption{Finding best set of parameters for PDFI: $N_{iter}$, $\sigma_{FLCT}$ and $\sigma_{PIL}$. {\it Left:} Dependence of the angle between E and B on the number of steps in the perpendicularization process, $\alpha_{EB}(N_{iter})$. {\it Middle:} Quality, $(f,a,\rho)$, of horizontal velocity components, $V_x$ (black) and $V_y$ (red), reconstructed using \texttt{FLCT} technique at different values of the Gaussian window width, $\sigma_{FLCT}$. {\it Right:} Black curves show the quality, $(f,a,\rho)$, of the $S_z$ reconstruction at different PIL widths, $\sigma_{PIL}$. Red, blue and yellow colors show the slopes $a(E_{x,PDFI},E_{x,ANMHD})$, $a(E_{y,PDFI},E_{y,ANMHD})$ and $a(E_{z,PDFI},E_{z,ANMHD})$ respectively, they quantify the faithfulness of the $E_x$, $E_y$ and $E_z$ reconstructions at different values of $\sigma_{PIL}$. Vertical dotted lines show the best values of parameters that we further use as default.} \label{fig_iter} \end{figure*} Figure~\ref{fig_iter} shows three panels that quantitavely justify selection of the best set of PDFI parameters (vertical dotted lines): number of iterations ($N_{iter}$, left panel), width of the Gaussian window ($\sigma_{FLCT}$, middle panel), and the PIL width ($\sigma_{PIL}$, right panel). The left panel, $\alpha_{EB}(N_{iter})$, shows dependence of the RMS angle between the electric and magnetic field vectors on the number of iterations using the perpendicularization technique (see \S~\ref{vecEPI}). When looking for the optimal number of iterations $N_{iter}$, we try to keep $N_{iter}$ as low as possible to achieve a high-speed performance, while still aiming for an angle close to $90^\circ$. For the \texttt{ANMHD} case, without perpendiculatization ($N_{iter}=0$), the angle between the magnetic and electric field vectors $\alpha_{EB}=75^{\circ}$. After only one iteration ($N_{iter}=1$), $\alpha_{EB}=87^{\circ}$. The angle slowly increases to $\alpha_{EB}=89.3^{\circ}$ by $N_{iter}=25$. The convergence rate slows down, and reaches $\alpha_{EB}=89.5^{\circ}$ by $N_{iter}$=50. We chose $N_{iter}=25$ as the optimal number of iterations, since above $25$ the convergence of $\alpha_{EB}$ toward $90^\circ$ is too slow to justify the additional computational effort. We also use $N_{iter}=25$ when applying the non-normal viewing angle technique (\S~\ref{nonzero}) for finding the needed scalar potential $\psi^D$. The Middle Panel of Figure~\ref{fig_iter} shows how the quality of the horizontal velocity reconstruction, using the \texttt{FLCT} technique in the \texttt{ANMHD} test case, depends on the Gaussian window width $\sigma_{FLCT}$. Since velocities parallel to the magnetic field do not affect the time evolution of the magnetic field and hence the fluxes of magnetic energy and helicity, in this plot we only compare components of the flow field that are perpendicular to ${\bf B}$. We find that $\sigma_{FLCT}<6$ pixels yield the worst agreement with the actual horizontal \texttt{ANMHD} plasma speed; $\sigma_{FLCT}=15$ pixels yields the best agreement, we therefore adopt it as the default value. For comparison, in \cite{Welsch2007} the optimal Gaussian window size was also chosen to be $\sigma_{FLCT}=15$ and in \cite{Fisher2012} $\sigma_{FLCT}=5$. Finally, the Right Panel of Figure~\ref{fig_iter}, ``Finding the best $\sigma_{PIL}$'', shows the quality of the PDFI electric field components (red, blue, orange) and Poynting flux (black) reconstruction for different $\sigma_{PIL}$'s, a free parameter that reflects the lack of confidence in the accuracy of the horizontal Doppler electric field away from PILs (see \S~\ref{pos}). The panel shows that $\sigma_{PIL}$ has a small effect on the quality of $E_y$, $E_z$ and $S_z$, above $\sigma_{PIL}=1$. We find that $\sigma_{PIL}\simeq1$ yields the best ratio between the total reconstructed and \texttt{ANMHD} Poynting fluxes: $\rho\simeq1,a\simeq1$. Using the best set of parameters ($N_{iter}=25$, $\sigma_{FLCT}=15$ and $\sigma_{PIL}=1$) found above, in \S~\ref{anmhd_esh}, we estimate the quality of the PDFI reconstructed electric fields, helicity and Poynting fluxes and also estimate the uncertainties in the results at a zero viewing angle. In \S~\ref{anmhd_ang} we describe how observing at the non-zero viewing angles affects the quality of the reconstruction. Finally in \S~\ref{anmhd_nind} we estimate the roles that different non-inductive Doppler- and FLCT contributions play in the reconstruction and compare current results to \cite{Fisher2012}. \subsection{Results: Electric Field, Poynting and Helicity Fluxes}\label{anmhd_esh} \begin{figure*}[htb] \centering \resizebox{0.9\hsize}{!}{\includegraphics[angle=0]{exeyez_display.pdf}} % \caption{Validation of the PDFI electric field at $\theta=0^\circ$: Actual ({\it top row}) and PDFI ({\it middle row}) electric field vector components, $[E_x, E_y,E_z]$, (left, middle, right) for the \texttt{ANMHD} test-case. {\it Bottom row:} Pixel-by-pixel scatter plots comparing the top two rows. The slopes of the linear fits and correlation coefficients are given in the top left corners.} \label{fig_el} \end{figure*} \begin{figure*}[htb] \centering \resizebox{0.98\hsize}{!}{\includegraphics[angle=0]{szdisplay_comp.pdf}} \caption{Validation of the PDFI Poynting flux at $\theta=0^\circ$. PDFI ({\it left}) and the actual ({\it middle left}) Poynting fluxes, $S_z$, for the \texttt{ANMHD} test-case, and also the pixel-by-pixel comparison between the two ({\it middle right}). The {\it far right} panel shows the same comparison, but instead of the current version of the PDFI we use the \cite{Fisher2012} method. } \label{fig_sz0} \end{figure*} Figures~\ref{fig_el} and~\ref{fig_sz0} show validation plots that compare electric field components $(E_x,E_y,E_z)$ and vertical Poynting fluxes derived from the PDFI-method with the actual \texttt{ANMHD} quantities. The PDFI-method reconstructs the \texttt{ANMHD} electric-field components quite well: the top and middle rows of Figures~\ref{fig_el} look almost identical. The slope of the linear fit to the reconstructed versus the actual electric-field component ranges from $a=0.94$ to $a=1.07$. The correlation coefficient, describing the quality of linear fit, in all cases is close to one. Using these electric-field components, we also find a good agreement for the vertical Poynting flux (Figure~\ref{fig_sz0}): the correlation coefficient $\rho=0.99$, the slope $a=0.98$ and the fraction, $f=1$. For comparison, \cite{Fisher2012} found slightly worse results: $\rho=0.97, a=0.94, f=0.9$ (Right Panel). It should be noted that the MEF method \citep{Longcope2004c} also accurately reconstructed the total Poynting flux in the tests by \cite{Welsch2007}, but the spatial correlation between the true and reconstructed fluxes was significantly worse, at $\rho=0.85$ (see their Figure 14). Decomposing the total Poynting flux into potential and free components (see \S~\ref{poynting}), we find that the PDFI reconstructs $100\%$ of both the potential and free components ($f=1$) and the slope between the reconstructed and PDFI is one ($a=1$). The free and potential components comprise $87\%$ and $13\%$ of the total unsigned Poynting flux respectively. We also test how sensitive the Poynting fluxes are to errors in the vertical Doppler velocity. We find that if there is a random Doppler velocity noise of $0.05$-km/s amplitude ($\bar{v}_z=0$), i.e. around $10$ to $20\%$ of the signal, then it does not significantly affect the Poynting flux: the error increases slightly from $0.14$ to $0.15$, the slope remains close to one ($a=0.98$) and the fraction $f=1.0$. However, if there is a bias Doppler velocity that increases all the velocities by $0.05$ km/s (towards the viewer) ($\bar{v}_z=0.05$), then the slope and the fraction increase to $a=1.2$ and the error is $Err.=0.2$. Similarly, if all velocities decrease by $0.05$ km/s ($\bar{v}_z=-0.05$), then the slope decreases to $a=0.8$ and the fraction to $f=0.8$. This test demonstrates how important it is to remove the Doppler velocity bias \citep{Welsch2013}, when inferring electric fields from the observations. In Figure~\ref{fig_hel0} we compare actual \texttt{ANMHD} and the PDFI helicity flux rates calculated from $\evec$. We find a good agreement between the two: the correlation coefficient $\rho=0.99$, the slope $a=1.08$, the fraction $f=1.1$ and the error $Err.=0.21$. For comparison, using \cite{Fisher2012} electric fields we find a very similar helicity flux rate, but with a slightly larger scatter ($Err.=0.23$): $\rho=0.97$, $a=0.95$, $f=0.9$. % The differences between our approach here and that of \cite{Fisher2012} are the adoption of the \texttt{\texttt{FISHPACK}} software to solve the two-dimensional Poisson equations, the ability to compute contributions to Doppler-shift electric fields from non-normal viewing angles, and a much more systematic and quantitative testing of the accuracy and robustness of the technique and its parameters. \begin{figure*}[htb] \centering \resizebox{1.0\hsize}{!}{\includegraphics[angle=0]{heldisplay_comp_fisher.pdf}} \caption{Validation of the PDFI helicity flux at $\theta=0^\circ$. See caption of Figure \ref{fig_sz0}.} \label{fig_hel0} \end{figure*} \subsection{Quality Of Electric Field and Poynting Flux Reconstructions At Non-Zero Viewing Angles}\label{anmhd_ang} \begin{figure*}[htb] \centering \resizebox{1.0\hsize}{!}{\includegraphics[angle=0]{anmhd_ex_angles_plot.pdf}} \caption{Validation of PDFI electric fields at viewing angles in the range of $\theta=[0^{\circ},60^{\circ}]$: quality of $E_x$ ({\it left}), $E_y$ ({\it middle}) and $E_z$ ({\it right}) reconstructions.} \label{fig_eang} \end{figure*} \begin{figure*}[htb] \centering \resizebox{1.0\hsize}{!}{\includegraphics[angle=0]{anmhd_sz_angles_plot.pdf}} \caption{Validation of PDFI Poynting fluxes at viewing angles in the range of $\theta=[0^{\circ},60^{\circ}]$: {\it Left and Middle:} Pixel-to-pixel scatter plots comparing actual and PDFI Poynting fluxes at viewing angles $\theta=15^{\circ}$ ({\it left}) and $\theta=50^{\circ}$ ({\it right}) . {\it Right}: Quality of $S_z$ reconstructions in the range of $\theta=[0^{\circ},60^{\circ}]$ .} \label{fig_szang} \end{figure*} To test performance of PDFI method at non-normal viewing angles $\theta$, we calculated the electric field (Figure~\ref{fig_eang}) and Poynting fluxes (Figure~\ref{fig_szang}) at values of $\theta$ ranging from $0$ to $60$ degrees. Figure~\ref{fig_eang} shows that at $\theta=0^\circ$ the error between reconstructed and the actual \texttt{ANMHD} electric fields is the smallest, and as the angle increases, the quality of the reconstruction decreases. For $\theta<30^{\circ},$ for all {\bf E}-components, PDFI correctly identifies the slope between reconstructed and the actual \texttt{ANMHD} electric fields, within a $10\%$ difference. At the largest angle, $\theta=60^\circ$, the slope is $a=0.83$ for $E_x$ and $a=0.67$ for $E_y$, and the error in these variables reaches up to $50\%$. In contrast to the horizontal electric field, the vertical component $E_z$ is relatively insensitive to the viewing angle: the slope $a$ varies within $[1.08,1.10]$. The latter is not surprising, since the Doppler contribution, which has most of the angular dependence, primarily constrains the horizontal field. For all ${\bf E}$-components the correlation coefficient at $\theta=[0,60]^\circ$ is quite high, $\rho>0.9$, implying that the estimates of the slope shown on the plot adequately describe the quality of the reconstruction. Figure~\ref{fig_szang} describes the quality of the vertical Poynting flux reconstruction at different viewing angles. Two left panels show a point-to-point comparison between the \texttt{ANMHD} and PDFI Poynting fluxes at $\theta=15^{\circ}$ (left panel) and $\theta=50^{\circ}$ (middle panel). At $\theta=15^{\circ}$ the agreement between the \texttt{ANMHD} and PDFI $S_z$ is very good: the slope $a=0.94, \rho=0.99, f=1, Err.=0.17$. Increasing the angle, at $\theta=50^{\circ}$ the scatter increases and the slope decreases to $a=0.64$, the error $Err.=0.44$ and PDFI method recovers $80\%$ of the total actual flux. The right panel summarizes the quality of the Poynting flux reconstruction for viewing angles within $[0,60^\circ]$-range. At close-to-normal angles, $\theta<20^{\circ}$, PDFI recovers more than $95\%$ of the total energy flux and the error between reconstructed and the actual Poynting fluxes is less than 15\% ($Err.=0.15$). At larger angles, $\theta<60^\circ$, $75\%$ of the total flux is recovered and the error $Err.<0.5$. For the helicity flux rate, at $\theta<60^\circ$, more than $90\%$ of the total flux is recovered correctly and the error $Err.=0.3$. \subsection{Comparison Of Electric Field Reconstruction Techniques}\label{anmhd_nind} In this section we analyze roles that the inductive and non-inductive components play in reconstructed vertical Poynting flux (Figure~\ref{fig_szcmp}), electric field and helicity fluxes (Figures~\ref{fig_helcmp} and \ref{fig_cmp}, Table~\ref{t1}). Figure~\ref{fig_szcmp} shows scatter plots comparing the actual vertical Poynting flux with the the Poynting fluxes derived with different reconstruction methods: (1) P, (2) PI, (3) PFI, (4) PDI, (5) PDFI, (6) PDFI at non-normal angle $\theta=30^{\circ}$, (7) FI, (8) DI, (9) DFI. The nomenclature that we use here is described in \S~\ref{esummary} and Table~\ref{table:esummary}. Using just the inductive part of the electric field (P), we reconstruct only $40\%$ of the total Poynting flux and the scatter from the linear dependence is large ($Err.=0.68$). Adding the ideal MHD assumption (PI) adds $30\%$ more Poynting flux, leading to much less scatter (larger $\rho$) and a smaller error ($Err.=0.48$). Inclusion of the non-ideal contribution due to horizontal plasma velocities, inferred from the \texttt{FLCT} (PFI), does not improve the solution. However, when we include the Doppler non-inductive component alone (PDI), we reconstruct $100\%$ of the flux, i.e. the role of the Doppler contribution is much higher than that from the \texttt{FLCT} horizontal velocities. Finally, when we add both the \texttt{FLCT} and the Doppler contributions (PDFI), the final Poynting flux is the closest to the actual $S_z$, with a slightly smaller error ($Err.=0.14$) than in the PDI case ($Err.=0.19$). On a separate note, if we use a non-PTD ideal inversion technique instead of the PTD, ${\bf E}=-{\bf V} \times \bvec$, $V_z=0$ (FI), we reconstruct only $20\%$ of the flux, i.e. the agreement between the reconstruction and the actual $S_z$ is poor ($Err.=0.95$). If we know the vertical Doppler velocity (DI), then $90\%$ of the actual flux is reconstructed and with a good agreement ($\rho=0.94, a=0.92, f=0.9, Err.=0.34$). With both vertical and horizontal velocities (DFI), we extract $100\%$ of the total Poynting flux, the slope $a=1.02$ and the error $Err.=0.28$. How is this different from $S_z$ from the PDFI? The PDFI electric field that includes both inductive and non-inductive contributions yields a factor of two smaller error in $S_z$ ($Err.=0.14$) than the DFI field, and shows less scatter, especially in the regions of strong Poynting fluxes, thus it better represents the vertical Poynting flux. For completeness, in Figure~\ref{fig_helcmp} we also compare helicity flux rates, $\frac{dH_R}{dt}$, from different reconstruction methods with the actual helicity flux rate. We remark that the crucial role the Doppler signal plays in reconstructing the Poynting and helicity fluxes in this case might be due to the process being modeled in the \texttt{ANMHD} simulation: an emerging magnetic bipole. It is possible that Doppler inputs to electric field estimates are less important when flux is not actively emerging. \begin{figure*}[htb] \centering \resizebox{1.0\hsize}{!}{\includegraphics[angle=0]{plot_sz_anmhd.pdf}} \caption{Comparison of the derived and the actual \texttt{ANMHD} vertical Poynting fluxes for different electric field inversion methods. Each method's results are shown in a separate panel with method's name indicated in the upper left corner and described in \S~\ref{esummary}. } \label{fig_szcmp} \end{figure*} \begin{figure*}[htb] \centering \resizebox{1.0\hsize}{!}{\includegraphics[angle=0]{plot_hel_all_anmhd.pdf}} \caption{Comparison of the derived and the actual \texttt{ANMHD} helicity flux rates for different electric field inversion methods. Each method's results are shown in a separate panel with method's name indicated in the upper left corner and described in \S~\ref{esummary}. } \label{fig_helcmp} \end{figure*} In Figure~\ref{fig_cmp} and Table~\ref{t1} we summarize the quality of different electric field inversion techniques in the same way we did for $S_z$ and $\frac{dH_R}{dt}$ for all the variables, $[E_x, E_y, E_z, S_z, \frac{dH_R}{dt}]$, where $[E_x, E_y, E_z]$ are the three components of the electric field, $S_z$ is the vertical Poynting flux, $ \frac{dH_R}{dt}$ is the helicity flux rate (see Eq.~(\ref{eq_hele}-\ref{eq_helv})). Besides the methods shown previously (see Figure~\ref{fig_szcmp}), we also add the method number 6, ``PDFI (Fisher)'', that compares variables reconstructed in \cite{Fisher2012} with the actual \texttt{ANMHD} values. The P-solution yields the worst reconstructions -- the slope between the actual and derived $E_x$ and $E_y$ is $0.26$ and $0.51$ respectively and the correlation coefficient in both cases is less than $0.8$. The quantity $E_z$, however, is reconstructed very well -- the slope is $1.10$ and the correlation coefficient is $0.93$. Adding the other ingredients only slightly improves the slope and the correlation coefficient of $E_z$. The Poynting and helicity fluxes, $S_z$ and $\frac{dH_R}{dt}$ respectively, are reconstructed poorly: the slopes are $0.4$ and $0.53$ respectively. Application of the ideal MHD constraint to the inductive solution (PI, case 2) significantly improves reconstruction of $E_x$ and $E_y$: the slope for $E_x$ increases from $0.26$ to $0.66$, and from $0.51$ to $0.63$ for $E_y$; this results in much better quality of the Poynting and helicity fluxes reconstructions: the slope for the Poynting flux changes from $0.4$ to $0.6$ from P- to PI-solution and for helicity flux - from $0.53$ to $0.93$; the error decreases by more than $20\%$. Addition of the non-inductive contribution from horizontal velocities (PFI, case 3), slightly improves the slope and correlation coefficient of $E_x$: the slope changes from $0.66$ to $0.78$ and the correlation coefficients - from $0.88$ to $0.91$. The error decreases as well. Inclusion of the non-inductive contribution from the vertical Doppler velocity, (PDI, case 4), provides much better improvement than the \texttt{FLCT} contribution (PFI) \citep{Fisher2012}. The slopes and correlation coefficients for both $E_x$ and $E_y$ increase significantly from $0.66$ and $0.63$ for the PI solution to $0.82$ and $0.96$, respectively. This results in much better reconstruction of $S_z$: the slope for $S_z$ changes from $0.6$ to $1.03$. The errors, especially for $E_y$, descrease, changing from $0.51$ to $0.24$. Finally, in the PDFI case (case 5) we add both non-inductive-, \texttt{FLCT} and Doppler contributions. We find that the PDFI method yields the best agreement with the \texttt{ANMHD} variables. The slopes for $E_x$ and $E_y$ are $0.94$ and $0.93$ with the correlation coefficients of $0.97$ and $0.98$ respectively. The Poynting and helicity fluxes have slopes of $0.98$ and $1.08$ with correlation coefficients of $0.99$ and $0.99$. The error varies from $0.14$ (for $S_z$) to $0.31$ (for $E_z$). To summarize, in terms of the fractions, PDFI method predicts roughly $100\%$ of the Poynting flux and $110\%$ of the helicity flux rate. It yields slopes and correlation coefficients that are similar to \cite{Fisher2012} (PDFI (Fisher), see case 6), but has smaller errors (compare $0.14$ vs. $0.23$ for $S_z$). At a non-normal viewing angle of $\theta=30^{\circ}$ (case 7), the agreement gets slightly worse, but still the slopes are close to one and the error, for example, in the Poynting flux is $28\%$. We reconstruct $90\%$ of the total Poynting flux and $110\%$ of the total helicity flux. Applying the same technique (\S~\ref{nonzero}) at normal viewing angle we find the same results as the ones derived using the normal viewing angle technique (\S~\ref{DEFzero}). We also calculate ${\bf E}, S_z$ and $dH_R/dt$ using the ideal non-PTD ${\bf E}=-{\bf V} \times \bvec$ formalism: the FI, DI and DFI (cases 8, 9 and 10). If only horizontal components of the velocity field are used and $V_z$ is set to zero, (FI), then we get the worst reconstruction - the slopes for $E_x$ and $E_y$ are less than $0.5$. The slope for the Poynting flux is $a=0.09$ and the fraction $f=0.5$. This is much worse than the PFI solution: $a=0.95$ and $f=1.0$. The difference between PFI- and FI-solutions might seem surprising at first, since both PFI and FI use the same information on the input. However while PFI solves the induction equation, the FI method does not. When we take the vertical velocity into account, (DI), the quantity $E_z$ is unknown, the quality of $E_x$ and $E_y$ improves slightly, raising from slopes of $a=0.49$ and $a=0.36$ to $a=0.35$ and $a=0.68$ respectively. As a result, the reconstruction of $S_z$ improves from $a=0.09$ to $a=0.92$ and the fraction is $f=0.9$. This is slightly worse than the $S_z$ from PDI: $a=1.03$ and $f=1.0$. For the non-inductive field estimates, we get the best agreement for $S_z$, when both horizontal and Doppler velocities are taken into account (DFI, case 10): $a=1.02, f=1.0$ \citep[See also][]{Ravindra2008}. Comparing the DFI with the PDFI reconstruction, we find that DFI is able to capture $E_y$, Poynting and helicity fluxes, although with larger errors than the PDFI, while $E_x$ and $E_z$ yield a poor reconstruction. In contrast, PDFI restores all three components of the electric field correctly, as well as Poynting and helicity fluxes, and with smaller errors. The reconstructions that include PTD (PI, PFI, PDI) also do a much better job than non-PTD solutions (FI, DI), when only one contribution to the velocity field (Doppler or FLCT) is available. We separately compare actual and calculated helicity flux rates and estimate the quality of the helicity flux rate reconstruction when, instead of the electric field (PI, PDI, PFI), only velocity estimates are used (FI or DI, see Eq.~(\ref{eq_helv})). When only the horizontal velocity field has been estimated (PFI or FI), then PFI yields much better reconstruction of helicity flux than FI: $a=0.95, f=1.0,Err.=0.28$ (PFI) vs. $a=0.61, f=0.5,Err.=0.55$ (FI). When only the Doppler velocity field is known (PDI or DI), then PDI also does a much better job than DI: $a=1.08, f=1.1, Err.=0.25$ (PDI) vs. $a=0.4,f=0.6, Err.=0.71$ (DI). Finally, if both horizontal and Doppler fields are known (PDFI or DFI), then both PDFI and DFI get similar values of helicity fluxes. However, the PDFI solution gets slightly better correlation coefficients and smaller errors and hence less scatter. Figure~\ref{fig_cmp} summarizes performance of all the methods in three panels (slope, correlation coefficient and error). For the PDFI method, shown with red squares, we get the most accurate electric fields, Poynting and helicity fluxes. The PDFI method outperforms ideal non-PTD methods (FI, DI, DFI). % \begin{figure*}[htb] \centering \resizebox{1.0\hsize}{!}{\includegraphics[angle=0]{comp_evec_auto_thr.pdf}} \caption{ Comparison of $[E_x, E_y, E_z, S_z, \frac{dH_R}{dt}]$, derived using different electric field inversion techniques (1-10), with the actual \texttt{ANMHD} quantities, where $[E_x, E_y, E_z]$ are the three components of the electric field, $S_z$ is the vertical Poynting flux, $\frac{dH_R}{dt}$ is helicity flux rate. The methods we use for calculating electric fields (1-10) are in detail described in \S~\ref{esummary}. {\it Left} panel shows slopes $a$ of the scatter-plots of reconstructed versus the actual \texttt{ANMHD} data. {\it Middle} panel shows correlation coefficients between the reconstructed and the actual data. {\it Right} panel shows errors $Err.=\sigma(u-u')/\sigma_{u}$, where $u$ and $u'$ are actual and derived values of one of the five analyzed quantities. Horizontal dotted lines correspond to the ideal reconstruction. The plotted values are given in Table~\ref{t1}. For description of $[a,\rho,Err.]$ see \S~\ref{anmhd_intro}. } \label{fig_cmp} \end{figure*} \renewcommand{\arraystretch}{1.6} \renewcommand{\arraystretch}{1.6} \begin{table*} [h] \caption{Comparison of $[E_x, E_y, E_z, S_z, \frac{dH_R}{dt}]$, derived using different electric field inversion techniques (1-10), with the actual \texttt{ANMHD} quantities. Each entry in the table has a form of $\left[\hspace{2.5pt}^{\rho}\hspace{1pt}a\,_{Err.}\right]$, where slope $a$, correlation coefficient $\rho$ and $Err.$ are described in \S~\ref{anmhd_intro}. The ideal reconstruction satisfies $\left[\hspace{2.5pt}^{\rho}\hspace{1pt}a\,_{Err.}\right]=(^{1.00} 1.00 _{0.00})$. For a plot, see Figure~\ref{fig_cmp}.} \small \begin{center} \begin{tabular}{l|@{}c@{}@{}@{}c@{}@{}@{}c@{}@{}@{}c@{}@{}@{}c@{}@{}} All &$E_x$ &$E_y$ &$E_z$ &$S_z$ &$\left(\frac{dH_{R}}{dt} \right)$\\ \hline (1) P &\hspace{2.5pt}$^{0.57}$\hspace{1pt}0.26\,$_{0.83}$\hspace{3pt} &\hspace{2.5pt}$^{0.78}$\hspace{1pt}0.51\,$_{0.64}$\hspace{3pt} &\hspace{2.5pt}$^{0.93}$\hspace{1pt}1.10\,$_{0.44}$\hspace{3pt} &\hspace{2.5pt}$^{0.78}$\hspace{1pt}0.40\,$_{0.68}$\hspace{3pt} &\hspace{2.5pt}$^{0.86}$\hspace{1pt}0.53\,$_{0.57}$\hspace{3pt} \\ (2) PI &\hspace{2.5pt}$^{0.88}$\hspace{1pt}0.66\,$_{0.59}$\hspace{3pt} &\hspace{2.5pt}$^{0.87}$\hspace{1pt}0.63\,$_{0.50}$\hspace{3pt} &\hspace{2.5pt}$^{0.96}$\hspace{1pt}1.06\,$_{0.43}$\hspace{3pt} &\hspace{2.5pt}$^{0.92}$\hspace{1pt}0.60\,$_{0.42}$\hspace{3pt} &\hspace{2.5pt}$^{0.95}$\hspace{1pt}0.93\,$_{0.34}$\hspace{3pt} \\ (3) PFI &\hspace{2.5pt}$^{0.91}$\hspace{1pt}0.78\,$_{0.42}$\hspace{3pt} &\hspace{2.5pt}$^{0.87}$\hspace{1pt}0.63\,$_{0.51}$\hspace{3pt} &\hspace{2.5pt}$^{0.96}$\hspace{1pt}1.06\,$_{0.32}$\hspace{3pt} &\hspace{2.5pt}$^{0.92}$\hspace{1pt}0.59\,$_{0.48}$\hspace{3pt} &\hspace{2.5pt}$^{0.96}$\hspace{1pt}0.94\,$_{0.28}$\hspace{3pt} \\ (4) PDI &\hspace{2.5pt}$^{0.93}$\hspace{1pt}0.83\,$_{0.37}$\hspace{3pt} &\hspace{2.5pt}$^{0.97}$\hspace{1pt}0.96\,$_{0.24}$\hspace{3pt} &\hspace{2.5pt}$^{0.96}$\hspace{1pt}1.07\,$_{0.32}$\hspace{3pt} &\hspace{2.5pt}$^{0.98}$\hspace{1pt}1.03\,$_{0.19}$\hspace{3pt} &\hspace{2.5pt}$^{0.98}$\hspace{1pt}1.09\,$_{0.25}$\hspace{3pt} \\ (5) PDFI &\hspace{2.5pt}$^{0.97}$\hspace{1pt}0.94\,$_{0.25}$\hspace{3pt} &\hspace{2.5pt}$^{0.98}$\hspace{1pt}0.93\,$_{0.19}$\hspace{3pt} &\hspace{2.5pt}$^{0.96}$\hspace{1pt}1.07\,$_{0.31}$\hspace{3pt} &\hspace{2.5pt}$^{0.99}$\hspace{1pt}0.98\,$_{0.14}$\hspace{3pt} &\hspace{2.5pt}$^{0.99}$\hspace{1pt}1.08\,$_{0.21}$\hspace{3pt} \\ (6) PDFI (Fisher) &\hspace{2.5pt}$^{0.91}$\hspace{1pt}0.83\,$_{0.41}$\hspace{3pt} &\hspace{2.5pt}$^{0.97}$\hspace{1pt}0.94\,$_{0.25}$\hspace{3pt} &\hspace{2.5pt}$^{0.96}$\hspace{1pt}1.05\,$_{0.31}$\hspace{3pt} &\hspace{2.5pt}$^{0.97}$\hspace{1pt}0.94\,$_{0.23}$\hspace{3pt} &\hspace{2.5pt}$^{0.97}$\hspace{1pt}0.95\,$_{0.23}$\hspace{3pt} \\ (7) PDFI ($\theta=30^\circ$) &\hspace{2.5pt}$^{0.95}$\hspace{1pt}0.93\,$_{0.32}$\hspace{3pt} &\hspace{2.5pt}$^{0.94}$\hspace{1pt}0.83\,$_{0.36}$\hspace{3pt} &\hspace{2.5pt}$^{0.92}$\hspace{1pt}1.11\,$_{0.47}$\hspace{3pt} &\hspace{2.5pt}$^{0.97}$\hspace{1pt}0.81\,$_{0.28}$\hspace{3pt} &\hspace{2.5pt}$^{0.96}$\hspace{1pt}1.08\,$_{0.33}$\hspace{3pt} \\ \hline (8) FI &\hspace{2.5pt}$^{0.68}$\hspace{1pt}0.49\,$_{0.73}$\hspace{3pt} &\hspace{2.5pt}$^{0.64}$\hspace{1pt}0.36\,$_{0.78}$\hspace{3pt} &\hspace{2.5pt}$^{0.84}$\hspace{1pt}0.68\,$_{0.54}$\hspace{3pt} &\hspace{2.5pt}$^{0.33}$\hspace{1pt}0.09\,$_{0.95}$\hspace{3pt} &\hspace{2.5pt}$^{0.85}$\hspace{1pt}0.61\,$_{0.55}$\hspace{3pt} \\ (9) DI ($\theta=0^\circ$) &\hspace{2.5pt}$^{0.52}$\hspace{1pt}0.35\,$_{0.86}$\hspace{3pt} &\hspace{2.5pt}$^{0.82}$\hspace{1pt}0.68\,$_{0.57}$\hspace{3pt} &\hspace{2.5pt}$^{--}$\hspace{1pt}--\,$_{--}$\hspace{3pt} &\hspace{2.5pt}$^{0.94}$\hspace{1pt}0.92\,$_{0.34}$\hspace{3pt} &\hspace{2.5pt}$^{0.73}$\hspace{1pt}0.40\,$_{0.71}$\hspace{3pt} \\ (10) DFI ($\theta=0^\circ$) &\hspace{2.5pt}$^{0.91}$\hspace{1pt}0.84\,$_{0.41}$\hspace{3pt} &\hspace{2.5pt}$^{0.97}$\hspace{1pt}1.04\,$_{0.26}$\hspace{3pt} &\hspace{2.5pt}$^{0.84}$\hspace{1pt}0.68\,$_{0.54}$\hspace{3pt} &\hspace{2.5pt}$^{0.96}$\hspace{1pt}1.02\,$_{0.28}$\hspace{3pt} &\hspace{2.5pt}$^{0.97}$\hspace{1pt}1.01\,$_{0.25}$\hspace{3pt} \\ \end{tabular}\normalsize \label{t1} \end{center} \end{table*} | \label{disc} We have modified the methods described in \cite{Fisher2012} for estimating electric fields from vector magnetic fields and Doppler velocities to incorporate non-normal viewing angles and the faster, more robust \texttt{FISHPACK} solver to produce a {\bf P}TD-{\bf D}oppler-{\bf F}LCT-{\bf I}deal-Ohm's solver (PDFI) that could be easily applied to observed data. We then used a pair of synthetic magnetograms extracted from MHD simulations, in which magnetic and Doppler velocity fields are known, to estimate the electric fields. Finally, we characterized the accuracy of the derived electric fields, as well as of the fluxes of magnetic energy and helicities. In this section we summarize the strengths and weaknesses of the PDFI technique, compare its results with other techniques and describe how it will be used to analyze HMI vector magnetograms and Doppler data. At zero viewing angle, the accuracy of the PDFI method is excellent. Using the PDFI method, the total Poynting flux, $S_z$, is estimated with an error of less than $1\%$ and the total helicity flux rate, $\frac{dH_R}{dt}$, with an error of $10\%$. The RMS of the integrands are $12\%$ and $26\%$ respectively. These estimates are more accurate, and much faster to derive than those from \cite{Fisher2012} (see Figures~\ref{fig_sz0} and \ref{fig_hel0} and Table~\ref{t1}).\footnote{For the Poynting flux the slope between the actual and the PDFI $S_z$ is $a=0.98$ in this paper versus $a=0.94$ in \cite{Fisher2012} and for the helicity flux rate, the slope $a=1.08$ in this paper versus $a=0.95$ in \cite{Fisher2012}; in both cases the scatter in the PDFIs is smaller than in \cite{Fisher2012}.} With increasing viewing angle $\theta$, the accuracy of the PDFI method slowly decreases. At $\theta<20^{\circ}$ PDFI recovers more than $95\%$ of the total Poynting flux with a slope of $a=0.98$. At $\theta<60^\circ$ PDFI recovers more than $75\%$ of $S_z$ with the slope of $a=0.6$ (see Figure~\ref{fig_szang}). As for the helicity flux rate, at $\theta<60^\circ$ more than $90\%$ of the total $\frac{dH_R}{dt}$ is recovered with the slope of $a=1.0$. In this paper we compare the quality of reconstructions of electric fields, Poynting and helicity fluxes from the PTD-based methods (PFI, PDI, PDFI), that explicitly enforce consistency of electric fields with evolution of $\vecB$, with non-PTD methods (FI, DI, DFI), like e.g. FLCT, that derive horizontal velocities by tracking the vertical magnetic field. We show that when both horizontal and Doppler velocity fields are known, the two approaches, PDFI and DFI, yield similar results for helicity and Poynting fluxes, with PTD-based PDFI method having slightly smaller errors. However, when either the Doppler or the FLCT velocity field is unknown, the PTD methods are much better in reconstructing both helicity and Poynting fluxes than non-PTD methods. For example, when only the horizontal velocity is known PFI reconstructs $70\%$ of total $S_z$ and $100\%$ of total $\frac{dH_R}{dt}$ while FI reconstruct only $20\%$ and $50\%$ respectively. Similarly, when only the Doppler velocity is known, then PDI reconstructs $100\%$ of total $S_z$ and $110\%$ of total $\frac{dH_R}{dt}$, while DI reconstructs $90\%$ and $60\%$ respectively. For the helicity flux rate, our results imply, that in order to correctly capture the helicity flux rate, one cannot only use the velocity field determined from the tracking of the vertical magnetic field (FI), but has to separately include the emergence term due to vertical velocity: the \cite{Demoulin2003} conjecture does not apply here, consistent with conclusions of \cite{Schuck2008,Liu2012,Ravindra2008}. We note, again, that the \texttt{ANMHD} data we analyze were drawn from a simulation of an emerging magnetic bipole, a configuration in which vertical flows (used as our Doppler velocity input) play particularly strong roles in the fluxes of magnetic energy and helicity. This might not be true in active regions when substantial amounts of new flux are not emerging. To put our results into context, in Table~\ref{table:comp_sz} we compare PDFI Poynting and helicity fluxes that we derive in this paper to those calculated with \texttt{DAVE4VM} and \texttt{DAVE+ANMHD} \citep{Schuck2008}. While \texttt{DAVE4VM} predicts roughly $75\%$ of the total Poynting flux and $95\%$ of the helicity flux rate, the PDFI method has a better performance, with less than $1\%$ error in the total Poynting flux and a $10\%$ error in the helicity flux rate. One should keep in mind, however, that \texttt{DAVE4VM}, unlike PDFI, does not take the Doppler velocity into account, hence it is fair to compare not the PDFI's, but the PFI's and \texttt{DAVE4VM}'s results. When we do that we find that the PFI performs similarly to \texttt{DAVE4VM}, and that both miss more than $25\%$ of the total Poynting flux. In contrast, PDFI, which includes both horizontal and Doppler velocities, yields an excellent agreement with the actual \texttt{ANMHD} quantities (see Table~\ref{table:comp_sz}). Note that the PDFI estimate for helicity flux is better than that of the \texttt{DAVE+ANMHD}, that takes the Doppler signal into account (see Fig. 14 in \cite{Schuck2008}). In addition, we remark that while the MEF method \citep{Longcope2004c} performed well in the tests by \cite{Welsch2007}, here we find the PDFI method to be superior, by several statistical measures, for the same test data (compare Figure 14 in \cite{Welsch2007} with Figure~\ref{fig_szcmp}). \begin{table*}[h] \caption{Comparison of accuracy of Poynting and rate of relative helicity fluxes estimates between the PDFI, PFI, \texttt{DAVE+ANMHD} and \texttt{DAVE4VM} \citep{Schuck2008} over $|B_z|>370G$. An ideal reconstruction satisfies $a=1, \rho=1, f=1$.} \small \begin{center} \begin{tabular}{p{2.4cm}|p{1.2cm}p{1.2cm}p{1.2cm}p{2.0cm}||p{1.2cm}p{1.2cm}p{1.2cm}p{2.0cm}} &PFI&\texttt{DAVE4VM} &PDFI & \small{\texttt{DAVE+ANMHD}} &PFI&\texttt{DAVE4VM} &PDFI & \texttt{DAVE+ANMHD} \\\hline & \multicolumn{4}{ c|| }{Poynting flux, $S_z$} & \multicolumn{4}{ c }{Helicity flux rate, $\frac{dH_R}{dt}$}\\ Slope, $a$ & $0.92$ & $0.71$ & $0.99$ & $0.99$ & $0.96$ & $0.9$ & $0.99$ & $1.16$ \\ Corr. coef., $\rho$ & $0.59$ & $0.83$ & $0.98$ & $0.96$ & $0.95$ & $0.94$ & $1.08$ & $0.96$ \\ Fraction, $f$ & $0.7$ & $0.76$ & $1.0$ & $0.99$ & $1.0$ & $0.94$ & $1.1$ & $1.46$ \\ \end{tabular} \label{table:comp_sz} \end{center} \end{table*} To estimate the speed of the PDFI method we did a series of inversion runs for the \texttt{ANMHD} dataset on a MacBook Pro laptop with 2GHz Intel Core i7 Processor and 8 GB 1333 MHz DDR3. For this test case, where $N_x=288$ and $N_y=288$ pixels, it takes $0.24$ seconds to calculate the plain P electric field, $1.2$ seconds for the PDF electric field and $7.3$ seconds for the ideal PDFI electric field ($N_{iter}=25$, see Figure~\ref{fig_iter}). One of the major weaknesses of the PDFI method, and also of any technique that uses the Doppler data, is a strong dependence on the Doppler bias velocity (see \S~\ref{anmhd_esh}). In this paper, to calculate PDFI electric fields, we assumed that we know the Doppler velocity field. In reality, however, the observed LOS velocity has a bias due to instrumental variations and the known correlation between the intensity and blue-shift, known as convective blueshift. Recently, \cite{Welsch2013} addressed this issue by presenting several methods to estimate the absolute calibration of LOS velocities in solar active regions near disk center that we apply to the HMI observations of the active region NOAA 11158 \citep{Kazachenko2014b}. Another weakness of the PDFI method is the absence of clear understanding of the extent to which additional information from other sources (besides Doppler and horizontal velocities) is necessary to fully specify the non-inductive part of the electric field. In the past \cite{Fisher2012} showed that the Doppler signal near PILs and horizontal velocities away from the PIL contain important non-inductive information for the electric field that cannot be derived from the Faraday's law. In this paper, we tested this approach at non-zero viewing angles and found a good agreement for viewing angles less than $50^\circ$ for the \texttt{ANMHD} test case. There may be additional degrees of freedom for the non-inductive electric field which are not fully captured by the Doppler plus transverse magnetic field flux-emergence contribution described in \cite{Fisher2012}. For example, rigid rotations of magnetic structures (e.g. sunspot rotation) that include a high degree of symmetry will have electric field components that are not fully captured by the PDFI formalism, because of the lack of change they produce in the vector magnetic field on the photosphere. Yet such motions transport significant energy and helicity into the corona \citep{Kazachenko2009}. For this reason, other tests of the emergence of a twisted flux tube from the interior into the solar atmosphere with different subsurface twisted flux tube configurations should be done. These tests will allow us formulate and then to incorporate any necessary additional electric fields corresponding to such horizontal vortical motions into the PDFI solution using observational information, such as observed sunspot penumbral motions as input. Furthermore, Doppler observations from a non-normal viewing angle can also be used to capture such horizontal motions, and help to recover the corresponding electric fields. | 14 | 4 | 1404.4027 |
1404 | 1404.1744_arXiv.txt | We report discovery of a dwarf galaxy in the Leo Triplet. Analysis of the neutral hydrogen distribution shows that it rotates independently of the tidal tail of NGC 3628, with a radial velocity gradient of 35--40\,km s$^{-1}$ over approximately 13\,kpc. The galaxy has a very high neutral gas content, explaining large part of its total dynamic mass -- suggesting a small dark matter content. As it is located at the tip of the gaseous tail, this strongly suggests its tidal origin. Should it be the case, it would be one of the most confident and closest (to the Milky Way) detections of a tidal dwarf galaxy and, at the same time, a most detached from its parent galaxy ($\approx$140 \,kpc) object of this type. | The idea of dwarf objects forming from the tidal debris left by galaxy mergers was first proposed by Zwicky~\cite{zwicky}, who suggested that interactions in systems of multiple galaxies can lead to an ejection of the tidal material and formation of an intergalactic structure, possibly even a dwarf galaxy. However, the "recycled" galaxies did not achieve much attention, apart from a symposium talk by Schweizer~\cite{schweizer}. The first object of this type was discovered by Mirabel et al.~\cite{mirabel}, who presented a photometric study of the Antennae galaxies, showing a tidal dwarf galaxy (TDG) formed from the collisional debris. Since then, many similar objects have been detected -- see eg. Brinks et al.~\cite{brinks}, or Duc et al.~\cite{vcc2062}. Recently, Kaviraj et al.~\cite{kaviraj} presented a study of a sample of 405 nearby TDG candidates, conducting a statistical analysis of their properties. Tidal dwarf galaxy candidates have been found in the Local Volume \citep[within 11 Mpc -- Hunter et al.][]{Hunter2000}, too. The M81 group hosts some of the closest examples of the TDGs. Small distance allowed to use the HST-based colour-magnitude diagrams \citep[Makarova et al.][]{Makarova2002} to analyse the star formation history of the TDG candidates and search for additional signs of the tidal origin. What makes the TDGs especially interesting is their mass composition. Whereas "normal" galaxies consist mostly of the dark matter (DM), TDGs do not; velocity of the DM particles in the galactic halo is much higher than the escape velocity of a TDG \citep[Bournaud][]{bournmass}, so they are not kinematically bound to it. Hence, such systems consist usually of the baryonic matter only. Additionally, as they are formed in the outer parts of the galactic disks, their metallicity is higher than in the non-tidal dwarfs. Only several TDGs were estimated to be heavy enough to contain significant non-baryonic fraction, but usual estimates suggest DM content at most similar to the baryonic mass -- far below the typical order of magnitude of difference in non-tidal dwarf systems \cite[Bournaud][]{bournmass}. Lack of the DM content and specific environment cause the evolution of the TDGs to be different from that of typical field galaxies, still to be studied and described. With the low dark matter content TDGs should also be more susceptible to the formation of galactic outflows driven by strong star formation. Alternately, different mass distribution may lead to a lower overall star formation and therefore to a low surface brightness nature of evolved TDG. TDGs are interesting not only because of their mass composition, but also because of their influence on the intergalactic environment. Tidal debris can interact with other group members, like in the case of the Leo Triplet galaxy NGC\,3627, known for its unusual magnetic field morphology \citep[Soida et al.][]{3627}. Recently, We\.zgowiec et al.~\cite{3627dwarf} suggested these peculiarities could be a result of a past collision with a dwarf galaxy. Thus, TDGs might play an important role in the further evolution of their progenitors. Galaxy systems with massive tidal tails and/or rings constitute favourable objects to search for the TDG candidates. One of the best examples of such objects is the Leo Triplet, a nearby group of galaxies, known for a large tidal plume extending eastwards from NGC\,3628. Originally described by Zwicky~\cite{zwicky}, the plume was later confirmed by photographic observations by Kormendy\&Bahcall~\cite{korm}. Neutral hydrogen studies by Rots~\cite{rots} and Haynes et al.~\cite{arecibo} revealed a thick {\rm H}{\sc i} structure, longer and wider than its optical counterpart. A detailed analysis of the {\rm H}{\sc i} distribution \citep[Stierwalt et al.][]{hinew} suggested numerous candidates for the non-tidal dwarf satellites. Recently, Nikiel--Wroczy\'nski et al.~\cite{leoeff} presented a study of the magnetic field in the Triplet. The authors suggested that the {\rm H}{\sc i} clump at the tip of the tidal tail could be a TDG. However, as pointed out in most of the TDG studies \citep[see eg. Kaviraj et al.][]{kaviraj} determination if a candidate is self-gravitating (galaxy), or rather a larger part of the tidal debris (that will never become a self-bound, independent object) is crucial. In this paper we use the archive neutral hydrogen and optical data to show that the velocity field of the TDG candidate detected in the Leo Triplet exhibits a velocity gradient and has a faint optical counterpart. These findings strongly support the idea of its independent rotation, thus confirming its identification as a galaxy. | \subsection{Stellar mass and age} Due to the very low surface brightness, the SDSS data do not allow to make a detailed fit to the spectral emission distribution (SED) that could be used to derive the star formation history and mass of the Leo-TDG. Still, it is possible to estimate some information from the photometry. Using the scaling relation from Bell et al.~\cite{Bell2003}, we can use {\it g'}\/ and {\it r'}\/ magnitudes and the resulting colour to get an estimate of the stellar mass. With {\it g'}\/ = 17.1$^{m}$, {\it r'}\/ = 16.65$^{m}$, {\it g'}--{\it r'}\/ = 0.45$^{m}$, and the values in Table~7 of Bell et al.~\cite{Bell2003} we calculate the M/L ratio = 3.12. With the measured L$_{r'}$ of $2.39\times 10^7$ L$_{\odot}$ this results in a stellar mass of $7.4\times 10^7$ M$_{\odot}$. A rough limit for the age of the dominant stellar population can be derived from comparison with the model integrated spectra. Assuming that the dwarf has at least some more or less recent star formation (given its large {\rm H}{\sc i} mass), we decided to use the \textsc{starburst99} code (Leitherer et al.~\cite{leith99, leith10}; Vazquez et al.~\cite{vazquez}) to model basic properties of the stellar population of Leo-TDG. Independently of the assumed metallicity (2 times solar to 1/20 solar) and star formation law (continous or instantanous), for the measured B$-$V = 0.62 we get a lower age limit of 1 Gyr (which is the limit of the published models). If we assume a moderate internal reddening of 0.3 mag, the age limits are from $3\times10^8$\,yrs for a star formation burst and solar metallicity to $\approx 10^9$\,yrs for 20\% of the solar metallicity. Obviously, while being relatively blue, the majority of the stars formed significantly more than $10^8$ years ago. For a more detailed analysis much better photometry is required. As estimated by Rots~\cite{rots}, the closest encounter between NGC\,3627 and NGC\,3628 may have happened $\approx 8 \times 10^{8}$\,yrs ago. Thus most of the stars in Leo-TDG (and probably the tidal dwarf itself) had to be formed shortly after the aforementioned collision of these galaxies. It is not likely that these stars formed in NGC\,3628 and have been later dragged away, as the distance from the parent object is very large. \subsection{Gas content} \label{himass} The gas mass of the Leo-TDG was estimated assuming M$_{\mathrm{H}_{\mathrm{I}}}$ [M$_{\odot}$] = 2.36$\times 10^{5} D^{2}_{\mathrm{Mpc}} \int S_{\nu} \mathrm{d}\nu$, where $\mathrm{S_{\nu} d\nu}$ is in Jy$\times$\,km\,s$^{-1}$ \citep[van Gorkom et al.][]{massest}. Using the distance of 12.15\,Mpc and total flux of 9.0$\pm$0.5\,Jy\,km\,s$^{-1}$ (see Sect.~\ref{momnt}), we obtained the total mass of the neutral hydrogen M$_{\mathrm{H}_{\mathrm{I}}} = 3.0$ -- $3.3 \times 10^{8} \mathrm{M_{\odot}}$. It is somewhat lower than the results from Stierwalt et al.~\cite{hinew}, but still of the same order of magnitude. The differences are most likely caused by the larger beamsize of the Arecibo telescope used by the authors of the former study, which causes confusion of the emission from the dwarf candidate by that from the tail. \subsection{Mass-to-light ratio and the total mass} The dynamical mass M$_{\mathrm{DYN}}$ of Leo-TDG can be derived from the rotational velocity at a given radius. For Leo-TDG, the radial velocity (not corrected for the inclination) gradient is about 35 -- 40\,km\,s$^{-1}$ over some 13\,kpc (with the tail contribution subtracted). The neutral hydrogen data do not allow to reliably estimate the turbulent component. Therefore, we decided to use a conservative assumption of 10 km\,s$^{-1}$ for a 1-dimensional turbulent contribution. If this is used to calculate the dynamical mass, one can obtain a total mass of some $7.9\times10^8 \mathrm{M}_{\odot}$. With the inclination unknown, this value can be treated as a lower limit of the dynamical mass. For a reasonable inclination of about $60\degr$ (based on the elongation of the the optical and {\rm H}{\sc i} shape of the dwarf), the total dynamical mass would rise to M$_{\mathrm{DYN}} = 1.41\times10^9 \mathrm{M}_{\odot}$. It should be strongly indicated here, that the dynamical mass estimate comes with a large uncertainty. As the dependence of the dynamical mass on the (unknown) inclination is given by is given by $M_{DYN} \propto 1\slash \sin(i)^{2}$, the dynamical mass would largely increase if only Leo-TDG is a more face-on-orientated system. In general, estimation of masses of the dwarf galaxies and their distributions is a complicated issue, as even if the inclination estimate is proper to some extent, the question of the finite disk thickness persists \citep[Rhee et al.][]{rhee}. The total baryonic content of Leo-TDG can be calculated as a sum of the stellar mass ($7.4\times 10^7 \mathrm{M}_{\odot}$) and gaseous component. Assuming a modest estimate of the molecular gas mass of $10-30\%$ of the {\rm H}{\sc i} mass \citep[as M$_{\mathrm{H}_{2}} \slash$M$_{\mathrm{H}_{\mathrm{I}}}$ for NGC\,3628 is equal to $\approx 20\%$, Obreschkow\&Rawlings][]{obr}, the total gas mass would be around 3.3 -- 4.3$\times 10^8$ M$_{\odot}$, so the total baryonic content M$_{\mathrm{BAR}}$ is $4.0$ -- $5.0 \times10^8 \mathrm{M}_{\odot}$. Estimate of M$_{\mathrm{DYN}}\slash$L$_{\mathrm{B}}$ can also be derived. L$_{\mathrm{B}}$ [L$_{\odot}$] is equal to $10^{-0.4\times(\mathrm{M} -\mathrm{M}_{\odot})}$. The B-band magnitude of the Sun is equal to 5.47 \citep[Cox][]{allen}. This yields the total B-band luminosity of 2.4$\times 10^7$\,L$_{\odot}$. {The M$_{\mathrm{DYN}}\slash$L$_{\mathrm{B}}$ is then 33 -- 59, and M$_{\mathrm{H}_{\mathrm{I}}}\slash$L$_{\mathrm{B}}$ is 12 -- 14. \subsection{Magnetic field} The resolutions used in our previous study \citep[Nikiel--Wroczy\'nski et al.][]{leoeff} -- 4$\arcmin$.3 in the radio continuum and 3$\arcmin$.5 in the {\rm H}{\sc i} data of Stierwalt et al.~\cite{hinew} gave no grounds to reject the coincidence of the {\rm H}{\sc i} and radio continuum-emitting regions. There were also no reliable optical images available. All this was suggestive for the existence of the magnetic field in Leo-TDG. With almost 4 times smaller beam of the {\rm H}{\sc i} data analysed in this work and using our optical image we could state that the radio peak is shifted by approximately 1' westwards from the neutral gas peak and seems to be located outside the optical emission. A large fraction of the radio continuum emission may be thus due to a background source. In the light of the new data we need to revise the estimate of the magnetic field strength. Setting the upper limit to the radio emission of 3.0 mJy/beam at the position of gaseous and{ optical feature implies the total magnetic field in the TDG to be B$_{\mathrm{TOT}} \le$ 2.8\,$\mu$G. The magnetic and cosmic-ray energy density amounts therefore to P$_{\mathrm{B+CR}}\le 6.8 \times 10^{-13}$ erg$\cdot$cm$^{-3}$. \subsection{TDG or a non-tidal, LSB galaxy?} Leo-TDG shares many of its characteristics with the TDGs \citep[Kaviraj et al.][]{kaviraj}. It is rather bluer than its supposed progenitor \citep[0.65 compared to 0.8 for NGC\,3628, Paturel et al.][]{hyperleda}, it is located exactly at the tip of the tidal tail, and has mass of some $10^{8}\mathrm{M_{\odot}}$, typical for such objects. On the other hand, if identified as a TDG, the discussed object would be the tidal dwarf most distant from its parent object, with a calculated separation of some 140 -- 150\,kpc, while 95\% of the TDG candidates do not lie more than 20\,kpc from their progenitors \citep[Kaviraj et al.][]{kaviraj}. Compared to the statistical sample, Leo-TDG is dim, as it contains less stars than typical TDG candidates. Among the most distinct features of this galaxy are its low surface brightness ($\mu_{\mathrm{B}}=25.55$) and very high abundance of the neutral gas. Because of that, we compare its properties not only with the TDGs, but also Low Surface Brightness (LSB) galaxies. We decided to take a galaxy (F563--1) from the samples collected by de~Blok et al. \citep{deblok, deblok2} and the "dark" LSB NGC\,3741 ~\citep[Begum et al.][]{n3741, begum}. As a comparison TDG, we have chosen the "old TDG" VCC\,2062 \citep[Duc et al.][]{vcc2062}. The data for the selected objects (TDGs and LSBs) are shown in Table \ref{compar}.\\ \begin{table*}[ht!] \caption{\label{compar} Parameters of Leo-TDG compared to TDG's and LSB's} \begin{center} \begin{tabular}{ccccc} \hline \hline Name & Leo-TDG & F563--1 & VCC\,2062 & NGC\,3741\\ Type & TDG & LSB galaxy & TDG & LSB (very dark)\\ Opt. size [kpc] & 7.5 & 3.4$^{1}$ & 0.7$^{1}$ & 1.7\\ {\rm H}{\sc i} size [kpc] & 13 & 16$^{1}$ & 4.2 & 14.6 \\ $\mu_{\mathrm{B}}$ [mag arcsec$^{-1}$] & 25.55 & 23.79 & 24.85 & 24.91\\ B--V & 0.615 & 0.58 & 0.35 & 0.36$^{2}$\\ Total mass [$\mathrm{M_{\odot}}$]& 7.9 -- 14.1$\times 10^{8}$ &$3.9\times 10^{10}$ & 3 -- 4$\times 10^{8}$& $4\times 10^{9}$ \\ Gas content [$\mathrm{M_{\odot}}$] & 3.3 -- 4.3$\times 10^{8}$ & $1.5\times 10^9$ & $0.8\times 10^{8}$ &$1.6\times 10^{8}$ \\ Stellar content [$\mathrm{M_{\odot}}$] & $7.4\times 10^7$ & $2.3\times 10^8$ $^3$ & 0.2 -- 0.7$\times 10^{8}$ & $1.4\times 10^{7}$ \\ M$_{\mathrm{H}_{\mathrm{I}}}\slash$L$_{B}$ & 12 -- 14 & 2.06 & 3 & 6.26 \\ M$_{\mathrm{B}}$ & $-12.97$ & $-16.7$ & $-13$ & $-13.13$\\ M$_{\mathrm{DYN}}\slash$M$_{\mathrm{BAR}}$ & 1.6 -- 3.5 & $\approx$17$^3$ & 2 -- 4 & 24\\ M$_{\mathrm{DYN}}\slash$L$_{\mathrm{B}}$ & 33 -- 59 & 50.1 & $\approx$10 & 149\\ \hline \hline \end{tabular} \end{center} $^{1}$Derived from its angular size\\ $^{2}$From Taylor and Webster \cite{tw}\\ $^{3}$Derived from the M/L ratio calculated basing on Bell \cite{Bell2003} \end{table*} The table clearly shows that the detected galaxy shares parameters of both TDGs and LSBs. In fact, it is not the only one dwarf system that is considered to be either a TDG or an LSB - likewise is the VCC\,2062 in the Virgo Cluster \citep[Duc et al.][]{vcc2062}. Both galaxies share similar characteristics: they are dim, low--mass systems with low surface brightness. They have a significant neutral hydrogen halo, showing signs of rotation, independent from the tidal arc movement. The velocity gradients are -- to the limits of inclination -- similar. However, the sizes of the {\rm H}{\sc i} haloes are different, as the one of VCC\,2062 is just 4.2\,kpc -- approximately three times smaller than that of Leo-TDG. The main difference between the Leo-TDG and non-tidal LSBs is the dominance of the gas content in the former one. In most of the LSBs, gas is not a dominant component: M$_{\mathrm{H}_{\mathrm{I}}}\slash$L$_{\mathrm{B}}$ is close to 1, and M$_{\mathrm{H}_{\mathrm{I}}}\slash$M$_{\mathrm{DYN}}$ is less than 10\% \citep[de Blok][]{deblok}. In case of Leo-TDG, {\rm H}{\sc i} dominates over other fractions, manifesting as a very high M$_{\mathrm{H}_{\mathrm{I}}}\slash$L$_{\mathrm{B}}$ (12 -- 14) and M$_{\mathrm{H}_{\mathrm{I}}}\slash$M$_{\mathrm{DYN}}$ of 25 -- 55\%. Such high neutral gas content causes M$_{\mathrm{DYN}}\slash$L$_{\mathrm{B}}$ to be higher than in VCC\,2062, while M$_{\mathrm{DYN}}\slash$M$_{\mathrm{BAR}}$ is very similar. As shown by de Blok~\citep{deblok2}, non--tidal LSB's have rather high M$_{\mathrm{DYN}}\slash$M$_{\mathrm{H}_{\mathrm{I}}}$ ratios, which is different from the case of the gas-dominated Leo-TDG. M$_{\mathrm{DYN}}\slash$M$_{\mathrm{BAR}}$ of the non-tidal galaxies are also much higher than 1.6 -- 3.5 estimated for Leo-TDG \citep[Bournaud et al.][]{bournmass}. All these features strongly favour scenario of the tidal origin of Leo-TDG. | 14 | 4 | 1404.1744 |
1404 | 1404.0574_arXiv.txt | We report the results of our analysis of new high resolution spectra of 37 late-F to early-G dwarf stars for the purpose of deriving their Li abundances. Most of the stars were selected from the large Valenti and Fischer compilation and had unknown Li abundances prior to the present study. When the new data are combined with data from our previous studies on this topic and analyzed in a similar way, we find, again, that stars with planets near the solar temperature are deficient in Li relative to a comparison set of stars. A similar result is obtained when we combine our data with a large database of stellar Li abundances from the literature. | In this study we again revisit the question of a possible correlation between the presence of Doppler-detected planets and stellar Li abundance. Several studies \citep{is04,tk05,gg08,is09,gg10,del14} indicate that stars with planets (SWPs) have lower Li abundances compared to stars without detected planets over a limited range in effective temperature (T$_{\rm eff}$) near the solar value. However, other studies \citep{lh06,bau10,gh10,ram12} have failed to confirm this pattern. Therefore, despite having received attention for about a decade from several inpedendent groups, this question is still unsettled. This question is important, because the Li abundance in a star's atmosphere is sensitive to a number of processes. These include gradual destruction of Li by canonical convective mixing in a star's envelope \citep{pin97}, enhanced destruction of Li from rotationally-induced mixing \citep{pin90}, and increase \citep{lg03,ash05} or decrease \citep{tv12} of surface Li abundance from accretion of planetary material. Rotational mixing (and the associated destruction of Li) can be enhanced by external torques, such as from the presence of stellar companions \citep{rd95}, a migrating planet \citep{cas09}, or a protoplanetary disk \citep{bou08}. It is this last mechanism that is most relevant to the present study. The formation of Doppler-detectable planets is more likely if a star is accompanied by a long-lived protoplanetary disk, and the disk is likely to slow the rotation of the star \citep{matt12}. \citet{gg11} showed that SWPs tend to have smaller vsini values than non-SWPs, confirming the link between planet formation and stellar rotation. The question of possible differences in Li abundance between SWPs and non-SWPs is therefore an important part of this puzzle and is an important test of the protoplanetary disk-stellar spin-down model. However, disentangling the various factors that influence Li abundance is a difficult task that requires large sample sizes. In \citet{gg10}, we compared the Li abundances of 50 SWPs and 49 comparison stars; the number of stars we actually employed in the analysis was less than this, as stars having only upper limits on Li abundance were not used in the comparison. In the present study, we seek to improve on that study by increasing the number of Sun-like comparison stars with Li abundance determinations. Our primary list for selecting targets is \citep{vf05}. The stars in that study have been searched for planets with the Doppler method, and the stars also have accurately known atmospheric parameters. However, most of them are still lacking Li abundance determinations; the spectra used by Fischer and Valenti did not include the Li feature at 6707 \AA. The purpose of the present study is to test again the claim that the Li abundances of Sun-like SWPs are different than those of similar stars without known planets. In Section 2 we describe our new spectroscopic observations and Li abudance analyses. In Section 3 we compare SWPs and stars without detected planets. We present our conclusions in Section 4. | We present the results of our analysis of high quality spectra of 37 late-F to early-G dwarfs, observed for the purpose of measuring Li in them. When combined with a large homogeneous sample of similar stars from the literature, we are able to confirm our previous findings from \citet{gg10} that the Li abundances of SWPs with T$_{\rm eff} \sim 5700$ K are smaller than those of stars without detected planets. In particular, SWPs with $5600 <$ T$_{\rm eff} < 5800$ K are deficient in Li by about 0.5 dex relative to comparison stars with similar properties. There is weaker evidence that SWPs with T$_{\rm eff} > 6100$ K are also deficient in Li. Our results generally confirm other recent independent studies of Li abundances in SWPs \citep{tk10,del14}. Additional observations of SWPs are needed for T$_{\rm eff} < 5600$ K and T$_{\rm eff} > 6100$ K to test whether Li abundance deviates significantly compared to non-SWPs in these regions. Observations of additional stars known not to have planets are required over the full temperature range studied here. Work is also required on theoretical predictions of the amount of Li depletion expected from the interactions between a protoplanetary disk and its central star (e.g., \citet{bou08}). Any successful model will need to be able to explain the observed pattern of Li abundance differences between SWPs and non-SWPs with T$_{\rm eff}$ as well as the scatter. | 14 | 4 | 1404.0574 |
1404 | 1404.5956_arXiv.txt | We present a catalogue of photometric and structural properties of 228 nuclear star clusters (NSCs) in nearby late-type disk galaxies. These new measurements are derived from a homogeneous analysis of all suitable WFPC2 images in the {\it HST} archive. The luminosity and size of each NSC is derived from an iterative PSF-fitting technique, which adapts the fitting area to the effective radius ($r_{\rm eff}$) of the NSC, and uses a WFPC2-specific PSF mo\-del tailored to the position of each NSC on the detector. The luminosities of NSCs are $\leq\!10^8L_{\rm V,\odot}$, and their integrated optical colours suggest a wide spread in age. We confirm that most NSCs have sizes similar to Globular Clusters (GCs), but find that the largest and brightest NSCs occupy the regime between Ultra Compact Dwarf (UCD) and the nuclei of early-type galaxies in the size-luminosity plane. The overlap in size, mass, and colour between the different incarnations of compact stellar systems provides a support for the notion that at least some UCDs and the most massive Galactic GCs, may be remnant nuclei of disrupted disk galaxies. We find tentative evidence for the NSCs' $r_{\rm eff}$ to be smaller when measured in bluer filters, and discuss possible implications of this result. We also highlight a few examples of complex nuclear morphologies, including double nuc\-lei, extended stellar structures, and nuclear $F606W$\,excess from either recent (circum-)nuclear star formation and/or a weak AGN. Such examples may serve as case studies for ongoing NSC evolution via the two main suggested mechanisms, namely cluster merging and {\it in situ} star formation. | Driven mostly by advances in the spatial resolution of modern telescopes over the last decades, it has now become firmly established that nuclear star clusters (NSCs) are an important morphological component of all types of galaxies \cite[e.g.][]{Phillips96, Carollo98, Boeker02,Boeker04,Cote06,Georgiev09b,Turner12}. The connection between the formation and evolution of NSCs and their host galaxies is a much-discussed topic of modern astrophysics. In particular, it is an open question whether NSCs are an essential ingredient for (or an intermediate step towards) the formation of a supermassive black hole (SMBH) in the galaxy nucleus \citep{Neumayer&Walcher12}. This question has been brought into focus by the realization that in many galaxies, both NSC and SMBH co-exist \citep{Seth08,Graham&Spitler09}, and that the few known SMBHs in bulge-less disks all reside in NSCs \citep{Filippenko&Sargent89, Shields08, Satyapal08, Satyapal09, Barth09, Secrest12}. The debate on the interplay between NSCs and SMBHs has been fuelled further by the finding that both types of "central massive object" (CMO) appear to grow in a way that is correlated with the growth of their host galaxies. This correlation has been induced from a number of so-called scaling relations, which demonstrate the dependence of CMO mass on various properties of the host galaxy. More specifically, the mass of both SMBHs \cite[e.g.][]{Ferrarese&Merritt00, Gebhardt00,Haering&Rix04} and NSCs \cite[e.g.][]{Wehner&Harris06, Rossa06, Ferrarese06} appears to correlate with the mass of the host galaxy bulge \cite[see][for a recent summary of this topic]{Scott&Graham13}. The most promising way to investigate the driving mechanism(s) behind these scaling relations is perhaps the study of late-type disk galaxies which are believed to be the most "pristine" galaxies which have not (yet) experienced any significant build-up of either bulge or CMO, and should therefore be well-suited to investigate the early stages of their (co)evolution. Understanding the origin and evolution of NSCs may also shed light on the nature of other massive compact stellar systems such as Globular Clusters (GCs) and Ultra Compact Dwarf galaxies (UCDs). There are numerous suggested scenarios for the origin of UCDs, including them being the extreme end of the GC luminosity function \cite[e.g.][]{Drinkwater00,Mieske02,Mieske12}, the end product of star cluster merging \cite[e.g][]{Kroupa98,Fellhauer&Kroupa02a,Kissler-Patig06,Bruens11}, the former nuclei of now dissolved galaxies \cite[e.g.][]{Bekki01,Bekki&Freeman03,Ideta&Makino04, Pfeffer&Baumgardt13}, or a combination of these mechanisms \cite[e.g.][]{Mieske06,Hilker09,DaRocha11,Brodie11,Norris&Kannappan11}. In particular, expanding the sample of NSCs with well-characterized sizes and stellar populations is needed to provide empirical constraints on the ``stripped dwarf galaxy'' scenario. The latter has recently received observational support from an overlap in the properties of UCDs and dwarf galaxy nuclei, which appear to show similar trends in their internal velocity dispersions \cite[e.g.][]{Drinkwater03,Chilingarian11,Frank11}, size-luminosity and color-magnitude relations \cite[e.g.][]{Cote06,Rejkuba07,Evstigneeva08,Taylor10}, their luminosity-weighted integrated ages and metallicities \cite[e.g.][]{Paudel10,Paudel11,Francis12,Madrid13}, and dynamical mass-to-light ratios \citep{Hasegan05,Hilker07,Mieske08,Taylor10} which seem to suggest unusual stellar mass functions \cite[e.g.][]{Mieske&Kroupa08,Dabringhausen09,Dabringhausen12,Marks12,Bekki13} or the presence of dark matter, most likely in the form of a SMBH \citep{Mieske13}. All these observations hint at a close connection between UCDs (and high-mass GCs) and NSCs, which should be further tested by comparison to NSC covering as wide a range in size and mass as possible. Further connections can be provided from utilizing the high spatial resolution of HST, to enable the investigation of internal spatial variations of the stellar populations of such systems \cite[e.g.][]{Kundu&Whitmore98,Larsen01,Strader12,Sippel12,Wang&Ma13,Puzia14}. Given that the typical sizes of NSCs fall into the range between a few pc and a few tens of pc \citep{Boeker04,Cote06,Turner12}, measuring their effective radii (and accurately separating their light from the surrounding, often complex, galaxy structure) requires HST resolution in all but the closest galaxies. We have therefore explored the HST/WFPC2 Legacy archive to analyse all available exposures of spiral galaxies within $\leq40$\,Mpc, and to derive the structural and photometric properties of the identified NSCs. Taking advantage of the accurate instrument knowledge gained following nearly 20 years of WFPC2 observations, our work expands on previous studies by i) significantly increasing the number of NSCs with accurate size and flux measurements, ii) improving the accuracy of previous photometric measurements by using updated PSF-fitting techniques, and iii) increasing (by a factor of three) the number of NSCs within an expanded morphological range of late-type spiral galaxies. This will allow to study evolutionary trends with the Hubble type of the host galaxy. Our work is organized as follows. In Section\,\ref{Sect:Data-Reduction-Analysis} we describe the galaxy sample and NSC identification (Sect.\,\ref{Sect:Gal-samp}). Image processing and combination is discussed in Section\,\ref{Sect:Image processing and combination}. Section\,\ref{Sect: NSCs} details the PSF-fitting techniques to derive the NSCs' structural parameters (Sect.\,\ref{Sect: Measuring NSC sizes}) and photometry (Sect.\,\ref{Sect: NSC photometry}). The limitations and uncertainties of the measured sizes, ellipticities and photometry are discussed in Section\,\ref{Sect: Uncertainties} and comparison with earlier work is performed in Section\,\ref{Sect:Comparison to Previous Studies}. Analysis of the general properties of the NSCs' size and luminosity distributions are presented in Sections\,\ref{Sect: Size Distribution} and \ref{Sect: NSC stellar pops}. We discuss the implications for the formation and evolution of massive compact stellar systems (\S\,\ref{Sect: NSCs_GC_UCDs}), the growth of NSCs (\S\,\ref{Sect: multiple NSCs}), and their coexistence with weak AGNs (\S\,\ref{Sect: AGNs}). Finally, we summarize our results in \S\,\ref{Sect:summary}. | \label{Sect:summary} Understanding the formation and evolution of nuclear star clusters (NSCs) promises fundamental insight into their relation to other massive compact stellar systems. This is because systems such as ultra compact dwarf (UCD) galaxies or massive globular clusters (GCs) harboring multiple stellar populations possibly originate as the former nuclei of now defunct satellite galaxies. On the other hand, NSCs often coexist with central black holes at the low-mass end of the SMBH mass range (around $\simeq 10^6 M_{\odot}$). It is actively debated what the role of NSCs is in the growth of such black holes and the fuelling of energetically weak ``mini-AGNs''. To address these questions, it is important to provide reliable measurements of the stellar populations properties of NSCs (age, metallicity, mass), as well as of their structural parameters for as many NSCs as possible, to provide observational constraints to the growing body of theoretical work addressing the above topics. The NSC catalogue presented in this work is the first step in this direction. It provides the largest and most homogeneously measured set of structural and photometric properties of nuclear star clusters in late-type spiral galaxies, derived from HST/WFPC2 archival imaging. We have searched the HST legacy archive for all late-type spirals within 40 Mpc (see Sect.\,\ref{Sect:Gal-samp}, Table\,\ref{Table:Galaxy sample}) that were observed with WFPC2. We have identified 323 such galaxies with suitable images of their nuclear region. More than two thirds (228/323) of these show an unambiguous NSC. We have used a state-of-the-art customized PSF-fitting technique to derive robust measurements of their effective radii and luminosities. For the PSF-fitting of each NSC, detailed in Section\,\ref{Sect: Measuring NSC sizes}, we employ {\sc TinyTim} PSF models tailored to the pixel location on the respective WFPC2 detector. We use the {\sc ishape} software \citep{Larsen99} to perform a $\chi^2$ minimisation fitting between the observed NSC profile and the PSF model convolved with an analytical model cluster profile. During this step, the fitting radius is iteratively adapted to the size of the NSC, minimizing the impact of circum-nuclear structures. We use the best-fit {\sc ishape} model to derive both the structural parameters ($r_{\rm eff}$, $PA$, and $\epsilon$) and the photometry of each NSC. For the latter, we use the latest \citep{Dolphin09} transformations, together with the measured NSC color, if available (see \ref{Sect: NSC photometry}). The complete catalogue of all measured structural and photometric properties of the 228 NSCs analysed in the various WFPC2 filters is provided in the online version of this paper. Table\,\ref{Table:Galaxy sample} contains the basic properties of the sample galaxies. Effective radius measurements for different filters and best fit analytical profiles are provided in Table\,\ref{Table:reff}, while the ellipticities and position angles of all NSCs are provided in Table\,\ref{Table:ellpa}. Calibrated and foreground reddening corrected NSC model magnitudes in the WFPC2 magnitude system are listed in Table\,\ref{Table: F300-814W_mag}, and their magnitudes in the Johnson/Cousins system in Table\,\ref{Table: UBVI_mag}. We caution that the latter magnitudes should be used with caution, because their accuracies depend on the available colour information for the respective NSC. Our main results can be summarized as follows: \begin{itemize} \item{ {\bf Sizes and structure:} we find that the measured sizes of NSCs in late-type spiral galaxies cover a wide range. Most NSCs have $r_{\rm eff}$ of a few pc, typical for Milky Way GCs. However, the $r_{\rm eff}$ distribution includes NSCs as large as a few tens of pc, i.e. comparable to some UCDs. There is tentative evidence for a smaller mean NSC size at bluer wavelengths, possibly caused by the presence of a weak AGN and/or a young stellar population that is more concentrated than the bulk of the NSC stars. On the other hand, some NSCs appear to be about 30\% {\it larger} when observed in the $F606W$ filter, compared to measurements in other filters. We discuss that this could be due to H$_\alpha$ line emission from (circum-)nuclear star formation. The NSCs in our late-type galaxy sample fall into the size-luminosity parameter space between early-type nuclei, UCDs, and massive GCs, which we interpret as support for the formation of these dense stellar systems from remnant nuclei of disrupted satellite galaxies (see \S\,\ref{Sect: NSCs_GC_UCDs}). The majority of the NSCs in our sample are best fit with King profiles with a high-concentration index ($C\equiv r_t /r_c = 100$), which makes them structurally similar to UCDs (see \S\,\ref{Sect: Size Distribution}). } \item {\bf Stellar populations:} the colour-colour and colour-magnitude diagrams (Fig.\,\ref{fig:WFPC2_CCD} and \ref{fig:NSCs CMD}) show that NSCs span a wide range in age and/or metallicity of their stellar populations. This agrees with previous studies which have suggested that NSCs are likely to experience recurring star formation events and/or accretion of other stellar clusters. A comparison to SSP models shows that NSCs also span a wide range in mass, from a few times $10^4$ to a few times $10^8\cal{M}_\odot$. Unfortunately, stellar mass estimates from optical photometry alone are rather uncertain due to the strong variation in $M/L$ as a function of age. As discussed in \S\,\ref{Sect: NSC stellar pops}, a small contribution from a young (e.g. 0.5 - 1\,Gyr) stellar population can be as luminous as a (much more massive) older (e.g. 10\,Gyr) stellar population, thus strongly biasing the derived cluster age. Combination of this catalogue with UV- and/or near-infrared data would provide much more robust mass, metallicity, and age estimates for NSCs. \item {\bf Double nuclei and nuclear disks:} We find a number of galaxies hosting double nuclei, i.e. two star clusters with comparable luminosity which are separated by only a few tens of parsecs. We regard these cases as plausible examples for the ongoing process of clusters merging in the galaxy nucleus (\S\,\ref{Sect: multiple NSCs}). We also find examples for small-scale circum-nuclear disk (aligned with host galaxy disk), which we interpret as evidence for NSC growth via gas accretion \citep{Seth06}. A systematic search for such morphological features in HST images can provide important constraints on each of the proposed NSC formation channels, and a statistical analysis of their frequency will be the topic of a follow-up paper. \item {\bf Active nuclei:} We also analysed a small comparison sample of weak AGNs. In some of these cases, we find faint unresolved nuclear emission in the residuals of the best-fit cluster model which are most likely caused by the AGN. The only AGN with complete colour information (NGC\,1042) deviates from SSP model predictions, suggesting that the AGN emission significantly affects the NSC colour. We therefore argue that such PSF-fitting techniques can be used to search for so far undetected nuclear activity, or at least to define promising target samples for spectroscopic searches for similar weak AGNs. \end{itemize} | 14 | 4 | 1404.5956 |
1404 | 1404.0432_arXiv.txt | The stability of dark matter is normally achieved by imposing extra symmetries beyond those of the Standard Model of Particle Physics. In this paper we present a framework where the dark matter stability emerges as a consequence of the Standard Model symmetries. The dark matter particle is an antisymmetric tensor field (analogous to the one used for spin-1 mesons in QCD), singlet under the Standard Model gauge group. The Lagrangian possesses an accidental $Z_2$ symmetry which makes the dark matter stable on cosmological time scales. Interactions with the Standard Model fields proceed through the Higgs portal, which allows the observed dark matter abundance to be generated via thermal freeze-out. We also discuss the prospects for observing this dark matter particle in direct detection experiments. | 14 | 4 | 1404.0432 |
||
1404 | 1404.5942_arXiv.txt | We test a new ``hybrid'' scheme for simulating dynamical fluid flows in which cylindrical components of the momentum are advected across a rotating Cartesian coordinate mesh. This hybrid scheme allows us to conserve angular momentum to machine precision while capitalizing on the advantages offered by a Cartesian mesh, such as a straightforward implementation of mesh refinement. Our test focuses on measuring the real and imaginary parts of the eigenfrequency of unstable axisymmetric modes that naturally arise in massless polytropic tori having a range of different aspect ratios, and quantifying the uncertainty in these measurements. Our measured eigenfrequencies show good agreement with the results obtained from the linear stability analysis of \citet{kojima1986} and from nonlinear hydrodynamic simulations performed on a cylindrical coordinate mesh by \citet{WTH1994}. When compared against results conducted with a traditional Cartesian advection scheme, the hybrid scheme achieves qualitative convergence at the same or, in some cases, much lower grid resolutions and conserves angular momentum to a much higher degree of precision. As a result, this hybrid scheme is much better suited for simulating astrophysical fluid flows, such as accretion disks and mass-transferring binary systems. \end {abstract} | \subsection{Context} Binary star systems, especially those containing compact components, are of great current interest in astrophysics. Binaries containing white dwarf components, in particular, are quite common because white dwarfs represent the most frequent endpoint for stellar evolution. Even double white dwarf (DWD) binaries, which are formed through common envelope evolution, are estimated to number $\sim 2.5 \times 10^8$ in our Galaxy \citep{Nelemans2001}. When gravitational radiation drives these binaries to a semi-detached state, these mass transferring systems are thought to be progenitors to both type Ia supernovae \citep{Geier2007,Rosswog2009,Fryer2010,Schaefer2012} and to hydrogen-poor R Coronae Borealis (RCB) stars \citep{Webbink1984,Staff2012}. \citet{Yungelson2004} review many possible evolutionary paths of binary systems. \citet{Tylenda2011} point to the evolution of V1309 Sco as an example of a system in which an actual merger has been witnessed observationally. Our desire is to employ computational fluid techniques to model, in a self-consistent manner and to a high degree of accuracy, mass-transfer in a wide variety of interacting binary star systems over hundreds, if not thousands, of orbits. Examples of our efforts, to date, include \citet{DSouza2006}, \citet{Motl2007}, \citet{Even2009}, and \citet{Marcello2012}. Related work by other groups includes \citet{Benz1990}, \citet{Fryer2006,Fryer2010,Fryer2012}, \citet{Yoon2007}, and \citet{Raskin2012}. Such a capability would allow us not only to better understand the behavior of binaries that are dynamically unstable toward merger or tidal disruption of the donor, but also to examine how spin-orbit coupling -- for example, the exchange of angular momentum between the donor star and a disk surrounding the accretor -- facilitates dynamical stability and leads to long phases of quasi-steady mass transfer. With such a tool we could simulate how slow accretion can bring an initially sub-Chandrasekhar-mass accretor to the brink of critical collapse; how a transition from sub- to super-Eddington accretion rates affects common-envelope development and evolution; and the steady-state structure of mass-transferring AM CVn type binaries. In simulating these systems numerically, a faithful representation of the flow will be achieved only if the grid resolution is sufficiently high across a range of dynamically interesting flow regions. These regions can vary in structure as well as in identity over time, so the grid needs to adapt accordingly. For example, even when only considering double degenerate binaries, the smallest length scales (surface layers of both stars, scale height of the disk, and fluid flow through the L1 Lagrange point) can be tiny compared to the binary separation. At the other extreme, an envelope consisting of an optically thick atmosphere engulfing both stars can develop shortly after accretion begins, quickly filling the original computational domain. This ``common envelope'' structure may expand to a size many times larger than the binary separation. Adaptive mesh refinement (AMR) techniques can be called upon to provide an appropriately high degree of spatial resolution in various, as well as in time-varying, regions of the flow. Astrophysical simulation codes that employ AMR include FLASH \citep{Fryxell2000}, ZEUS \citep{Hayes2006}, PLUTO \citep{Mignone2007}, and Scorpio, recently developed by Marcello et al. (2014, in preparation). AMR techniques can be straightforwardly implemented on a Cartesian mesh but are more difficult to employ across curvilinear grids. Despite this difficulty, AMR has been implemented on both spherical and cylindrical coordinate grids \citep{Fryxell2000}. On the other hand, advection schemes implemented on Cartesian meshes are most naturally designed to conserve linear momentum, rather than angular momentum. Because binary evolutions can be faithfully followed through hundreds of orbits only if the simulation conserves angular momentum to a high degree of accuracy, in the past we have chosen to use a cylindrical computational grid, which more naturally facilitates conservation of orbital angular momentum. However, even a cylindrical grid does not match the symmetry of each individual binary component. And in the absence of a mesh refinement capability we have not had the full freedom to distribute resolution where it is needed. Alternative curvilinear meshes designed to address some of the above outlined shortcomings have been used in previous disk and torus simulations. For example, \citet{Zink2008} discuss a multi-patch technique meant for spherically symmetric or axisymmetric simulations, and \citet{Fragile2009} introduce a ``patched-sphere'' mesh similar to the multi-patch technique. But, as with a cylindrical coordinate grid, such schemes do not easily accommodate mesh refinement techniques. Ultimately, \citet{Zink2008} say that Cartesian mesh-refinement is likely better suited for problems like binary mergers. Historically, therefore, it has been difficult to achieve both high -- and adaptive -- spatial resolution while at the same time achieving a high degree of angular momentum conservation. The hybrid advection scheme designed by \citet{Call2010} and implemented here, allows us to have our cake and eat it, too. It facilitates conservation of angular momentum to machine accuracy on a refined Cartesian mesh. \citet{Mignone2012} state that their method, implemented in the PLUTO code, allows for conserving angular momentum on a Cartesian mesh to machine precision. However, it appears as though this only applies to local ``shearing-box'' models. Their scheme breaks down the azimuthal fluid velocity into two pieces, an average plus a residual term. The average fluid velocity is handled in a linear step by moving the fluid in the direction of the orbital motion. The residual velocity is then handled in the standard way. Only the residual portion is subject to the Courant condition, so this leads to larger time steps, and the apparent motion of the fluid through the grid is much smaller, leading to less numerical dissipation. However, this method requires that the average angular motion be parallel to one of the coordinate bases describing the grid and that the relevant grid direction use periodic boundary conditions. This is not the case for a global simulation performed on a Cartesian mesh, as we are doing in this work. The method implemented in PLUTO would be a very suitable choice for simulating axisymmetric tori as we have done in this work. It would not, however, provide a significant advantage for binary mass-transfer simulations, even if performed on a cylindrical grid, as the fluid is largely rotating with a uniform angular velocity. While our new method will allow us to conserve angular momentum at a level that is necessary to faithfully follow interacting binary simulations through thousands of orbits, it does not ease the computational burden of simulating the large number of time steps needed for such a simulation. In particular, the Courant limit on the size of individual time steps remains an impediment. But by facilitating the straightforward implementation of AMR, our new scheme simplifies the task of load balancing and thereby enhances a code's ability to efficiently and more fully use the capabilities of massively parallel computers, allowing simulations to be carried through thousands of orbits. As our discussion in \S \ref{octopus} indicates, the port of our new hybrid scheme to Octopus is yet another step toward achieving this goal. \subsection{Overview of this Work} In this paper we demonstrate the utility of the hybrid scheme by focusing on a quantitative analysis of nonaxisymmetric, dynamical instabilities that arise spontaneously in Papaloizou-Pringle tori \citep{PP1984}, hereafter referred to as PP tori. Each PP torus is a non-self-gravitating, differentially rotating, geometrically thick, axisymmetric disk in orbit about a central point mass. Its internal structure is defined by a balance between gas pressure gradients and gradients in the effective potential. The vertical thickness of the disk/torus relative to its radial extent is determined by the choice of the polytropic index for the gas and an initial angular momentum distribution. (See \S2 for details.) These configurations are suitable for demonstrating the capabilities of our hybrid scheme because: They each have a simple analytically definable initial state; while each initial state is axisymmetric, the system is unstable to the development of nonaxisymmetric structure, hence, its evolution has a fully three-dimensional character; the eigenvector of the most unstable mode for each chosen initial configuration, while not known analytically, should be well defined and its measured properties -- for example, its complex eigenfrequency -- should be reproducible and independent of the specific numerical scheme that is used to perform the dynamical simulation. At the same time, the PP torus provides a good test for hydro codes such as ours because of the challenges it provides. A Cartesian mesh is not ideally suited for the initially axisymmetric torus problem, and thus gives the hybrid scheme an opportunity to prove its worth, for example, by partially overcoming the spurious $m=4$ modes that are excited by the structure of an underlying Cartesian grid. Each of our simulations is carried out on a rotating and refined Cartesian grid. The hydrodynamic code we are using (see \S\S 2.2-2.3) has AMR capabilities, but for simplicity we have chosen not to activate the AMR feature. Instead, for each simulation the volume of the grid that is occupied by the initial torus is resolved using a time-invariant, fixed level of refinement (LOR). The effect of grid resolution is assessed by repeating individual simulations several times, using a different (fixed) number of refinement levels. (Typically, we employ 4, 5, or 6 LOR.) In addition, for each initial state and for each specified LOR, the dynamical evolution is carried out using two separate advection schemes: (1) A traditional ``Cartesian'' scheme in which the $x$, $y$, and $z$ components of the linear momentum are advected across the refined Cartesian grid; and (2) our new ``hybrid'' scheme in which radial momentum and angular momentum -- instead of the $x$ and $y$ components of the momentum -- are advected across the refined Cartesian grid. In total, results from 23 simulations are reported here. For each simulation, the real and imaginary components of the complex eigenfrequency of the fastest growing unstable mode are measured and, as appropriate, compared with previously published results from the nonlinear hydrodynamic simulations performed on a uniformly zoned cylindrical mesh by \citet*{WTH1994} and from the linear stability analysis presented by \citet{kojima1986}. In an effort to eliminate any subjective bias that might be introduced into the measurement of these eigenfrequencies and, at the same time, to facilitate future efforts to reproduce our results, we introduce a mathematically prescriptive method for quantifying both the value of and uncertainty in each eigenfrequency measurement. In doing this we are able to meaningfully assess the performance of our new hybrid advection scheme, relative to the performance of the traditional Cartesian advection scheme. We show that qualitative convergence is achieved with the hybrid scheme at the same, or sometimes at significantly lower, grid resolutions. At the same time, we show that the hybrid scheme allows conservation of the system's total angular momentum to machine accuracy. As explained above, this is a highly desirable feature that is not possible to achieve using a familiar Cartesian advection scheme. This is perhaps the most significant attribute of our hybrid scheme. Historically, the expectation has been that angular momentum conservation can be achieved when modeling an astrophysical disk or binary system only if one adopts a coordinate grid -- for example, cylindrical or spherical coordinates -- whose underlying basis vectors accommodate the curvilinear features of the flow. Our hybrid scheme facilitates conservation of angular momentum on a Cartesian grid. The hydrocode that has been used to carry out the primary set of simulations reported in this paper employs OpenMP to enable multiple execution threads within a single, multi-core compute node. All of the models in our primary set of simulations -- totaling 22 in number and using up to 6 LOR -- fit within a single node of our high-performance computing system. In \S3 of this paper we present results from one simulation that was conducted on a rotating Cartesian grid with 7 LOR. This single simulation was run using Octopus, a closely aligned hydrocode built on top of High Performance ParalleX (HPX), a newly emerging parallel runtime system. | Our ultimate goal is to model in as realistic a manner as possible the dynamical evolution of mass-transferring binary systems. This can only be accomplished if the hydrodynamic code that is used to perform each simulation conserves angular momentum extremely well. We also need to have the flexibility of AMR to adequately resolve spatial features across many orders of magnitude in length scales simultaneously. The hybrid scheme described here allows us to conserve angular momentum to high accuracy on a refined Cartesian mesh, facilitating the use of AMR. Our hybrid scheme is an implementation of the theoretical formulation developed by \cite{Call2010}, which shows that we have the freedom to choose different coordinate bases for the transport velocity relative to the grid and the advected momentum quantities. In the past, these chosen basis sets typically have been the same -- resulting in the advection of Cartesian momentum components on a Cartesian mesh, or cylindrical momentum components on a cylindrical mesh. In the hybrid scheme implemented here, we have chosen to advect cylindrical momentum components across a rotating Cartesian mesh. This allows us to conserve angular momentum to machine precision while capitalizing on the advantages of a Cartesian mesh, such as mesh refinement. In order to test this method, we followed the development of nonaxisymmetric instabilities in massless PP tori having $n=3$ and $q=2$ (uniform specific angular momentum). This is a well-defined, fully three-dimensional problem with a reproducible solution. We evolved seven different initial tori with aspect ratios ranging from $R_-/R_+=0.1$ to $0.7$. We chose to evolve two particular models, $R_-/R_+ = 0.3$ and $0.7$, using several different grid resolutions. We compared our results to the linear stability analysis of \cite{kojima1986} and to the nonlinear hydrodynamics results of \cite{WTH1994}. Our code achieved good agreement with results from these previous studies. We also introduced a prescriptive method for measuring the real and imaginary parts of the eigenfrequency of unstable modes, attaching an uncertainty to those measurements. This was done in an effort to increase transparency, reduce the influence of human judgment, and facilitate the reproducibility of these simulations. Through this work we have illustrated the utility of the PP tori as a new test problem, to be added to the standard suite of hydrodynamic test problems, that provides a means for measuring the ability of a particular code to correctly transport and conserve angular momentum. A comparison of the resolution dependence of the hybrid scheme compared to the Cartesian momentum advection scheme shows that the hybrid scheme achieves qualitative convergence at grid resolutions that are equal to or lower than the Cartesian scheme. Specifically, we observe that in the $R_-/R_+ = 0.7$ torus, the hybrid scheme achieves qualitative convergence at only 5 LOR, whereas the Cartesian advection scheme required 7 LOR to achieve the same convergence -- requiring a factor of $\sim 30$ more computation zones. The hybrid scheme also reduces the level of unphysical $m=4$ distortions that characteristically appear in simulations involving angular motion across a Cartesian grid. Here we have demonstrated the utility of the hybrid scheme, which is only one very specific implementation of the formalism presented by \citet{Call2010}, which can be applied in a fully relativistic generalized coordinate system. We acknowledge valuable interactions that we have had with J. Frank, G. Clayton, P. M. Motl, and H. Kaiser over the course of this project. This work has been supported, in part, by grants ACI-1246443 and AST-1240655 from the U.S. National Science Foundation, in part, by NASA ATP grant NNX10AC72G, and, in part, by U.S. Department of Energy grant DE-SC0008714. This research also has been made possible by grants of high-performance computing time at HPC@LSU through the allocation hpc\_dwd\_amr, and across LONI (Louisiana Optical Network Initiative), especially award loni\_lsuastro11. \begin{figure} \centering \includegraphics[height=\textheight,width=\textwidth,keepaspectratio=true]{f3.pdf} \caption{Data from the Cartesian advection simulation of Model 3 ($R_-/R_+ = 0.3$) is used to demonstrate four diagnostic diagrams, providing a direct comparison to Figure 2 in \cite{WTH1994}. (a) A ``$D_m - t$'' diagram showing the Fourier amplitude of modes $m=1,2,3,4$ at the radius of pressure maximum in the equatorial plane. (b) A ``$\phi_m-t$'' diagram showing the phase angle of the $m=1$ mode, again at the radius of pressure maximum in the equatorial plane. (c) A ``$D_m - r$'' diagram showing the amplitude of the $m=1$ mode as a function of radius in the equatorial plane at time $t=2.5\ t_\mathrm{orb}$. (d) A ``$\phi_m-r$'' diagram showing the azimuthal location of the density maximum ($m=1$) as a function of radius in the equatorial plane at time $t=2.5\ t_\mathrm{orb}$.} \label{4_panel} \end{figure} \begin{figure} \centering \includegraphics[scale=1.8]{f4.pdf} \caption{Data from a simulation of Model 3 ($R_-/R_+ = 0.3$) using the hybrid scheme and 5 LOR. (a) Shows a $D_m - t$ plot, with vertical lines marking the starting and ending points of the region used to determine $\bar{S}$; the portion of the curve used to measure this slope is shown in bold. A dashed line with a slope equal to the $\bar{S}$ is also shown. (b) Shows the windowed slope measurement, $S(t)$, with vertical lines marking the starting and ending points of the region used to measure $\bar{S}$. The horizontal dashed line identifies the measured slope. } \label{postprocessing1} \end{figure} \begin{figure} \centering \includegraphics[scale=2]{f5.pdf} \caption{Same as Figure \ref{postprocessing1}, but from a simulation of Model 7 ($R_-/R_+ = 0.7$) using the Cartesian momentum advection scheme and 5 LOR.} \label{postprocessing2} \end{figure} \begin{figure} \centering \includegraphics[angle=90,scale=1.3]{f6.pdf} \caption{A series of $D_m - t$ plots from every initial model listed in Table \ref{table:modeltable} labeled with either ``H'' or ``C'' for simulations using the hybrid scheme or the Cartesian momentum advection scheme, respectively, followed by the model number. All simulations were performed using 5 LOR. In each case the portion of the plot used to measure $\bar{S}$ is highlighted in bold, and a dashed line with the measured slope is also plotted.} \label{postprocessing} \end{figure} \begin{figure} \centering \includegraphics[scale=2]{f7.pdf} \caption{Comparison of the imaginary component ($y_2$) of the eigenfrequency of various unstable modes between this work and \citet{WTH1994} and \citet{kojima1986}. The red, green, and blue dashed curves connect discreet points from Kojima's linear analysis for $m=1,2,$ and 3 respectively. The points marked with red boxes, green circles, and blue triangles show values measured in this work for $m=1,2,$ and 3 respectively; open symbols represent the Cartesian momentum advection scheme and filled symbols represent the hybrid scheme. Red and green diamonds represent $m=1$ and 2 results published in WTH, respectively. Our measured growth rates show good agreement with both previous studies. As described in the text, error bars on data points from this work represent the relative quality of measurements.} \label{y_2comparison} \end{figure} \begin{figure} \centering \includegraphics[scale=2]{f8.pdf} \caption{As in Figure \ref{y_2comparison}, but showing the real component ($y_1$) of the eigenfrequency of various unstable modes. Our measured frequencies show good agreement with both previous studies.} \label{y_1comparison} \end{figure} \begin{figure} \centering \includegraphics[scale=1.8]{f9.pdf} \caption{ The accumulated change in Model 7's total angular momentum $(L-L_0)$, measured relative to its initial value, $L_0$, is shown as a function of time, $t/t_\mathrm{orb}$, for 6 different simulations -- 3 different levels of refinement and using Cartesian momentum advection (dashed curves) or the hybrid scheme (solid curves). In the case of both schemes, the red, green, and blue curves show data from simulations conducted with 5, 4, and 3 LOR respectively. When using the Cartesian momentum advection scheme, the level of angular momentum conservation shows clear resolution dependence. The hybrid scheme conserves angular momentum at a level set by machine truncation error.} \label{angmom_comparison} \end{figure} \begin{figure} \centering \includegraphics[scale=1.6]{f10.pdf} \caption{Data from Model 3 ($R_-/R_+ = 0.3$) simulations. Each of the six panels shows the $D_m - t$ plot for $m=1$ (solid curve) and $m=4$ (dashed curve) obtained from, as labeled, either the hybrid scheme or the Cartesian advection scheme for 4, 5, or 6 LOR. While the time-dependent growth of the unstable $m=1$ mode is very similar in all cases, the amplitude of $m=4$ appears to be strongly resolution dependent, reflecting the 4-fold symmetry of the Cartesian grid structure.} \label{model_3_compare_LOR} \end{figure} \begin{figure} \centering \includegraphics[scale=2]{f11.pdf} \caption{The same information as shown in Figure \ref{model_3_compare_LOR}, but grouped differently. Each frame shows all three levels of refinement for either $m=1$ (top row) or $m=4$ (bottom row), using either the hybrid scheme (left column) or Cartesian momentum advection scheme (right column). The two upper panels ($m=1$) show that both the hybrid and Cartesian seem to converge, as the lower resolution simulations display noisier $D_m - t$ plots. The lower two panels show how the amplitude of $m=4$ fluctuations decreases dramatically with increasing grid resolution.} \label{model_3_compare_mode} \end{figure} \begin{figure} \centering \includegraphics[scale=1.4]{f12.pdf} \caption{The same information as shown in Figure \ref{model_3_compare_LOR}, but grouped differently. Each frame compares hybrid (solid curve) to Cartesian (dashed curve) momentum advection schemes, for either $m=1$ (left column) or $m=4$ (right column), and at 4, 5, and 6 levels of refinement (bottom, middle, and top rows, respectively). The left column ($m=1$) illustrates how the difference between the Cartesian and hybrid scheme plots differs less with increasing resolution, suggesting that they are both converging to the same answer. The right column ($m=4$) shows that, at each specified resolution, the hybrid scheme generally produces slightly lower amplitude $m=4$ distortions than the Cartesian advection scheme.} \label{model_3_compare3} \end{figure} \begin{figure} \centering \includegraphics[scale=1.8]{f13.pdf} \caption{Data from Model 7 ($R_-/R_+ = 0.7$) simulations. Each of the six panels shows the $D_m - t$ plot for $m=3$ (solid curve) and $m=4$ (dashed curve) obtained from, as labeled, either the hybrid scheme or the Cartesian advection scheme for 4, 5, or 6 LOR.} \label{model_7_compare_LOR} \end{figure} \begin{figure} \centering \includegraphics[scale=2]{f14.pdf} \caption{The same information as shown in Figure \ref{model_7_compare_LOR}, but grouped differently. Each frame shows all three levels of refinement for either $m=3$ (top row) or $m=4$ (bottom row), using either the hybrid scheme (left column) or Cartesian momentum advection scheme (right column).} \label{model_7_compare_mode} \end{figure} \begin{figure} \centering \includegraphics[scale=1.6]{f15.pdf} \caption{The same information as shown in Figure \ref{model_7_compare_LOR}, but grouped differently. Each frame compares hybrid (solid curve) to Cartesian (dashed curve) momentum advection schemes, for either $m=3$ (left column) or $m=4$ (right column), and at 4, 5, and 6 levels of refinement (bottom, middle, and top rows, respectively).} \label{model_7_compare3} \end{figure} \begin{figure} \centering \includegraphics[scale=1.7]{f16.pdf} \caption{ $D_m - t$ plots showing the development of the $m=3$ mode from seven different simulations of Model 7 ($R_-/R_+ = 0.7$). The hybrid scheme simulations (solid curves) show qualitative convergence at 5 LOR, while the Cartesian momentum advection scheme (dashed curves) does not converge until 7 LOR, requiring a factor of $\sim 30$ more computational zones.} \label{7LOR_compare_3} \end{figure} \begin{figure} \centering \includegraphics[scale=2]{f17.pdf} \caption{ $D_m - t$ plots showing the development of the $m=4$ distortion in Model 7 ($R_-/R_+ = 0.7$). At each specified LOR, the hybrid scheme (solid curves) shows much lower levels of development of this unphysical distortion than the Cartesian advection scheme (dashed curves). } \label{7LOR_compare_4} \end{figure} \begin{figure} \centering \includegraphics[scale=0.3]{f18.pdf} \caption{Mass density plots from the Cartesian momentum advection simulations of the Model 7 ($R_-/R_+ = 0.7$) torus, showing the progression from $m=4$ dominated evolutions at lower resolutions to $m=3$ dominated evolutions at the highest resolution. Each slice shows the mass density in the $z=0$ plane, and is labeled with the number of LOR used in the simulation. Data for 4, 5, 6, and 7 LOR are taken from $t=1.8t_\mathrm{orb}$, $2.1t_\mathrm{orb}$, $3.0t_\mathrm{orb}$, and $2.5t_\mathrm{orb}$ respectively. Note that the simulations are the same as shown in Figure 2, although from different times in the evolutions. } \label{state_data} \end{figure} \appendix | 14 | 4 | 1404.5942 |
1404 | 1404.7495_arXiv.txt | Motivated by the order-of-magnitude difference in the frequency of giant planets orbiting M dwarfs inferred by microlensing and radial velocity (RV) surveys, we present a method for comparing the statistical constraints on exoplanet demographics inferred from these methods. We first derive the mapping from the observable parameters of a microlensing-detected planet to those of an analogous planet orbiting an RV-monitored star. Using this mapping, we predict the distribution of RV observables for the planet population inferred from microlensing surveys, taking care to adopt reasonable priors for, and properly marginalize over, the unknown physical parameters of microlensing-detected systems. Finally, we use simple estimates of the detection limits for a fiducial RV survey to predict the number and properties of analogs of the microlensing planet population such an RV survey should detect. We find that RV and microlensing surveys have some overlap, specifically for super-Jupiter mass planets ($m_p \gtrsim 1~M_{\rm Jup}$) with periods between $\sim 3-10~$years. However, the steeply falling planetary mass function inferred from microlensing implies that, in this region of overlap, RV surveys should infer a much smaller frequency than the overall giant planet frequency ($m_p\gtrsim 0.1~M_{\rm Jup}$) inferred by microlensing. Our analysis demonstrates that it is possible to statistically compare and synthesize data sets from multiple exoplanet detection techniques in order to infer exoplanet demographics over wider regions of parameter space than are accessible to individual methods. In a companion paper, we apply our methodology to several representative microlensing and RV surveys to derive the frequency of planets around M dwarfs with orbits of $\lesssim 30~$years. | \label{sec:introduction} The field of exoplanets has reached the state where statistically significant samples of planets have been detected via multiple different methods: radial velocities, transits, microlensing, and direct imaging. Each of these techniques are sensitive to different, yet complementary regions of parameter space, both in terms of the orbits and physical properties of planets, as well as their host stars. This means that it should be possible to derive generalized constraints on the demographics of planets over a broader region of parameter space than is available to each method individually by synthesizing their various data sets. This is not a trivial task, as each method of exoplanet discovery introduces unique observational biases and selection effects that have to be reconciled before comparing directly with detections from other methods. Various RV and microlensing surveys have provided independent constraints on the demographics of exoplanets \citep{2005ApJ...622.1102F,2008PASP..120..531C,2008A&A...487..373S,2009A&A...493..639M,2010Sci...330..653H,2010PASP..122..905J,2010ApJ...720.1073G,2010ApJ...710.1641S,2011arXiv1109.2497M,2013A&A...549A.109B,2011Natur.473..349S,2012Natur.481..167C}. These two methods are generally complementary. RV surveys are sensitive to low-mass planets at short periods and larger planets at longer periods (ultimately limited to planets with periods less than the duration of the survey), while microlensing is sensitive to planets with projected separations near the Einstein radius of their host stars, which generally corresponds to larger separations than are covered by RV surveys ($\sim 4~$AU for a Solar mass star, and scaling as the square-root of the mass of the host star). With regards to giant planets, this complementarity of RV and microlensing surveys is important. The two techniques are generally probing regions of parameter space where we expect planet formation to operate differently. RVs are most sensitive to orbits interior to the snow line, whereas microlensing is more sensitive to orbits exterior to the snow line. The snow line is the point in protoplanetary disks at (and beyond) which water ices can deposit out of the nebular gas \citep{1981PThPS..70...35H}. Past the snow line, deposition of water ices increases the surface density of the disk by a factor of $\sim 2-3$, increasing the isolation mass by a factor of $\sim 4-5$, allowing for gas giant formation by core accretion within the observed, short lifetimes ($\sim 1-10~$Myr) of gaseous disks \citep{2008ApJ...673..502K,1995Natur.373..494Z,2006ApJ...651.1177P}, at least for Solar-mass stars. Giant planets detected by RV are thus generally expected to have migrated while those detected by microlensing presumably formed \emph{in situ}. Therefore, by combining constraints on exoplanet demographics inferred by the two methods, we can get a complete picture of giant planet formation and migration, and disentangle these two competing physical effects. RV and microlensing surveys also differ in terms of host stars: RV samples are targeted, but this is not the case for microlensing. The rarity and unpredictability of microlensing events render a targeted survey ineffective, so the general observing strategy is to instead monitor large numbers of stars in the most dense fields towards the Galactic bulge. Such serendipitous microlensing events are dominated by M dwarf lenses, as these are more common than higher mass lenses. In terms of giant planets, which we define as having masses $m_p \gtrsim 0.1~M_{\rm Jup}$ (see the companion paper, \citet{clanton_gaudi14b}, for a discussion of what exactly constitutes a ``giant planet''), comparing detection results from RV and microlensing surveys is particularly interesting, because giant planets formed by core accretion are predicted to be rare around M dwarfs \citep{2004ApJ...612L..73L}. This prediction has been claimed to be confirmed by RV surveys of M dwarfs, who find a paucity of giant planets \citep{2010PASP..122..905J,2013A&A...549A.109B} relative to Solar-mass stars. However, because these surveys are mostly sensitive to separations less than the snow line, it is not clear if this is a consequence of a lack of formation or a lack of migration. Indeed, microlensing surveys \citep{2010ApJ...720.1073G,2012Natur.481..167C} find that giant planets beyond the snow line are more common by an order of magnitude than short period giant planets probed RV surveys, which on its face is in contradiction to the generic prediction of core accretion. Motivated by this order of magnitude difference in giant planet frequency, here we develop the formalism and methodology for comparing and synthesizing the exoplanet detection results from RV and microlensing surveys. In a companion paper, we employ this methodology to determine if these two independent measurements of giant planet frequencies are consistent with the same population of planets. We provide a high-level overview of the RV and microlensing techniques in \S~\ref{sec:mlens_rv_techniques}, highlighting the key differences necessary to understand before comparing exoplanet demographics derived independently from these two methods. In \S~\ref{sec:OoM} we provide an order of magnitude estimate that demonstrates that the ``typical'' microlensing planet, defined such that its parameters place it in the peak region of sensitivity for microlensing surveys, should be marginally detectable by RVs. We derive distributions of the period and velocity semi-amplitude expected from this ``typical'' microlensing planet, allowing for circular and eccentric orbits in \S~\ref{sec:typical_dists}. In \S~\ref{sec:marg} we develop a procedure to marginalize over all microlensing parameters, including the projected separation, the host star/planet mass ratio, host star mass, and lens distances to produce a posterior distribution of radial velocity observables for a population of planets analogous to that inferred from microlensing. We show this marginalized distribution and compute the number of planets per star a RV survey of M dwarfs (with some fiducial sensitivity and number of epochs per star) should detect, as well as the number of long-term RV trends per star resulting from planets, in \S~\ref{sec:results} and provide a discussion of our conclusions in \S~\ref{sec:discussion}. | \label{sec:discussion} Our order of magnitude calculation demonstrates that an RV survey with an average number of epochs per star $N=30$, average precision $\sigma=4~{\rm m~s^{-1}}$, and time baseline $T=10~$yr, should be able to detect the typical microlensing planet ($M_l\sim0.5~M_{\odot}$, $M_p\approx 0.26~M_{\rm Jup}$, $r_{\perp}\sim 2.5~{\rm AU}$) on a circular orbit with an inclination equal to the median value of a uniform prior in $\cos{i}$ at a SNR of about 5. However, if we allow the orbital parameters of this typical microlensing-detected planet to vary, we find that it is not always detectable via RVs. This motivates our study to map analogs of the population of planets inferred from microlensing into the RV observables $K$ and $P$ to determine the number of planets per star RV surveys should detect and see as trends. In doing so, we are able to show that although the regions of planet parameter space where microlensing and RV surveys are sensitive are largely disjoint, there is some overlap for giant planets with masses $m_p\gtrsim 10^2~M_{\oplus}$ and with orbital periods between $\sim 3-10~$years. An RV survey with a duration of $\sim 10~$years has significant overlap with microlensing surveys in terms of period. We find that the median period of analogs to the planets inferred by microlensing is $P_{\rm med}\approx 9.4~$years, with a 68\% interval of $3.35\leq P/{\rm yr}\leq 23.7$. However, due primarily to a steeply declining planetary mass function, we find that the median velocity semi-amplitude of such planets is $K_{\rm med}\approx 0.24~{\rm m~s^{-1}}$, with a 68\% interval of $0.0944\leq K/{\rm m~s^{-1}}\leq 1.33$. The steeply declining planetary mass function measured by microlensing is primarily responsible for the shape of the posterior distribution of $K$ for these planets, and puts a majority of the planet population inferred from microlensing out of reach for RV surveys. The detection sensitivity of RV surveys to such planets is thus ultimately limited by their velocity precision. The best precisions of state-of-the-art RV surveys is $\sim$few ${\rm m~s^{-1}}$ (including both instrumental errors as well as stellar jitter), so the bulk of the the planet population inferred from microlensing is inaccessible by RVs. In order for RV surveys to access the bulk of these planets, precisions of $\sim 0.1~{\rm m~s^{-1}}$ need to be achieved. However, in order to detect the large population of giant planets ($m_p\gtrsim 0.1~M_{\rm Jup}$) inferred from microlensing, precisions of only $\sim 1~{\rm m~s^{-1}}$ are needed. We find that orbital eccentricity and inclination also significantly impact the mapping of microlensing parameters into the RV observables $K$ and $P$, although to a lesser extent than the planetary mass function. Allowing for eccentric orbits greatly increases the area of allowed parameter space (see figures \ref{fig:kp_contours_typical} and \ref{fig:kp_contours_ecc_test}) and variations in the inclination increase the spread of $K$ values directly, since $K \propto \sin{i}$ for a given period, host star mass and mass ratio. Nonetheless, we find that variation of the orbital parameters tend not to shift the median $K$ and $P$, but mostly serve to increase the spread in these observables. Despite the fact that RV and microlensing surveys are sensitive to different parameters, have very different targeting/observing strategies, and thus suffer from different selection effects, we have demonstrated that it is possible to statistically compare detection results from these two completely independent techniques. We find that a RV survey monitoring a sample of 100 M dwarfs with a mass distribution covering the interval $0.07\leq M_{\star}/M_{\odot}\leq 1.0$ uniformly in $\log{M_{\star}}$, and with an average number of epochs per star $N=30$, average sensitivity $\sigma=4~{\rm m~s^{-1}}$, and time baseline $T=10~$yr, should on average detect $4.9^{+4.6}_{-2.6}$ giant planets and identify $2.4^{+2.4}_{-1.4}$ as long-term trends at a SNR of 5 or higher if analogs of the planets detected by microlensing orbit local M dwarfs. In practice, real RV surveys could have different sensitivities than those implied by our fiducial survey. A careful comparison of microlensing and RV detection results will require knowledge of the specific sensitivities of all stars monitored by RVs, but the methodology of comparison will be very similar to that laid out in this study. In a companion paper, we do just this, comparing microlensing results with the M dwarf RV studies of \citet{2010PASP..122..905J} and \citet{2013A&A...549A.109B}. | 14 | 4 | 1404.7495 |
1404 | 1404.2279_arXiv.txt | % We show that the PeV neutrinos detected by IceCube put unique constraints on ``secret'' interactions of neutrinos with the cosmic neutrino background (C$\nu$B). The coupling must be $g <0.03$ for the mediating boson mass $m_{X} \lesssim 2$ MeV, $g/m_{X} < 5$ GeV$^{-1}$ for $m_{X} \gtrsim 20$ MeV, and $g/m_{X} < 0.07$ GeV$^{-1}$ in between. We also investigate the possibility that neutrino cascades degrade high-energy neutrinos to PeV energies by upgrading C$\nu$B where the energy flux of PeV neutrinos can coincide with the Waxman--Bahcall bound or the cosmogenic neutrino flux for protons, thanks to energy conservation. However, a large coupling is required, which is disfavored by laboratory decay constraints. The suppression of PeV--EeV neutrinos is a testable prediction for the Askaryan Radio Array. | Recently IceCube reported the detection of two PeV neutrinos and 26 additional events, more than expected from atmospheric backgrounds \cite{IC13:1stPeV,IC13:Sci}. The arrival direction is consistent with the isotropic distribution, suggesting that at least some of the events are of cosmological origin. The Hillas condition to accelerate primary cosmic-rays up to $\sim 100$ PeV allows a dozen possibilities \cite{lah+13,Anchordoqui+13}, such as gamma-ray bursts \cite{Waxman_Bahcall97,Murase+06,gz07,Murase_Ioka13}, active galactic nuclei \cite{ste13,Murase+14}, galaxy clusters and groups \cite{Murase+08,Murase+13:hadronuclear}, star-forming galaxies \cite{Loeb_Waxman06,Murase+13:hadronuclear,kat+13}, and heavy dark matter \cite{Feldstein+13,Esmaili_Serpico13,Ema+13}. We are witnessing the birth of high-energy neutrino astrophysics. The IceCube events remind us of Supernova (SN) 1987A, which placed unique limits on the properties of neutrinos, especially ``secret'' interactions of neutrinos with the cosmic neutrino background (C$\nu$B) \cite{Kolb_Turner87,Manohar87}. The neutrino--neutrino interactions \cite{Bialynicka-Birula64,Bardin+70}, even stronger than the weak interactions of the standard model, remain largely unconstrained below the electroweak energy scale \cite{Bilenky_Santamaria99,Lessa_Peres07,Laha+13} because of their weakness and the difficulties in focusing the neutrino beam \cite{T2K14}. High-energy neutrinos are attenuated by C$\nu$B if the cross section is large enough \cite{Weiler82}. The much longer distance and higher energy of the IceCube events than those of SN 1987A can tighten restrictions on the secret interactions \cite{Keranen98}, as well as neutrino decays \cite{bae+12,Pakvasa+13}, leptoquark couplings \cite{Barger_Keung13}, and so on. The IceCube neutrinos may even result from secret interactions of neutrinos. By energizing C$\nu$B, high-energy neutrinos may develop cascades in intergalactic space, like gamma-ray cascades \citep[e.g.,][]{Coppi_Aha97,Murase+12,Inoue_Ioka12}. Since such self-interactions conserve the total energy, the neutrinos keep the energy flux while reducing the typical energy. Thus, this scenario can naturally account for a possible ``{\it coincidence problem}'': why the observed neutrino flux is comparable to the Waxman--Bahcall bound \cite{Waxman_Bahcall99} or equivalently the cosmogenic neutrino flux at EeV energies produced by ultrahigh-energy cosmic-ray (UHECR) protons \citep[e.g.,][]{bz69,Yoshida_Teshima93,Takami+09}. The lack of $>2$ PeV events indicates either a soft spectrum or a break at several PeV \cite{IC13:Sci,lah+13}, implying different processes at PeV and EeV energies. It is a coincidence that two different processes separated by three orders of magnitude in energy give almost the same flux. We use $(H_0,\Omega_m,\Omega_{\Lambda}) =(72\ {\rm km}\ {\rm s}^{-1}\ {\rm Mpc}^{-1}, 0.27, 0.73)$ and $c=\hbar=k=1$. | Requiring that the observed neutrinos are not affected, we obtained astrophysical constraints on the secret interactions. Although our work is greatly simplified, detailed studies are possible, including various interaction types, scattering angles of particles, and the neutrino--antineutrino difference. Although a large coupling is already disfavored by the laboratory decay constraints, an appealing point of the cascade scenario is that the PeV flux can coincide with the Waxman--Bahcall bound. We must tune the mass to bring a peak to PeV energies, but the neutrino flux is the same for different parameters. The cascade scenario predicts the suppression of $>$PeV neutrinos, whatever the source is. Future neutrino detectors such as the Askaryan Radio Array \cite{ARA12} can test this possibility. \ack We thank J.~Beacom, K.~Blum, H.~Kodama, K.~Kohri, T.~Moroi, K.~Ng, Y.~Okada, K.~Omukai, and H.~Takami, and also A.~Ishihara, K.~Mase, and S.~Yoshida for holding the workshop ``Cosmic Neutrino PeVatron (NuPeV 2014).'' This work is supported by KAKENHI 24000004, 24103006, 26287051.\\ {\it Note added}: As this paper was being completed, we learned of an independent study by Ng and Beacom \cite{Ng+14}, which will be submitted to arXiv simultaneously. | 14 | 4 | 1404.2279 |
1404 | 1404.1673.txt | We present high S/N UV spectra for eight quasars at $z\sim3$\ obtained with VLT/FORS. The spectra enable us to analyze in detail the strong and weak emission features in the rest-frame range 1300-2000 \AA\ of each source (\ciii, \siiii, \aliii, \siii, \civ\ and blended \siiv+\oiv). Flux ratios \aliii/\siiii, \civ/\aliii, \siiv+\oiv/\siiii\ and \siiv+\oiv/\civ\ strongly constrain ionizing photon flux and metallicity through the use of diagnostic maps built from {\sc cloudy} simulations. The radius of the broad line region is then derived from the ionizing photon flux applying the definition of the ionization parameter. The \rb\ estimate and the width of a virial component isolated in prominent UV lines yields an estimate of black hole mass. We compare our results with previous estimates obtained from the \rb\ -- luminosity correlation customarily employed to estimate black hole masses of high redshift quasars. | \defcitealias{negreteetal12}{Paper I} \defcitealias{negreteetal13}{Paper II} Measuring relevant physical parameters from the observed broad-line spectra of quasars is still an open challenge. Identification and intensity measurements of the strongest emission lines has made possible a rough inference about typical conditions in the emitting gas from the earliest days of quasar spectroscopy. The first intermediate redshift quasars discovered in the 1960s showed a fairly high ionization spectrum with prominent lines of \civ, and \heiiuv\ in addition to strong Balmer lines seen in the lower redshift sources. Photoionization by a central continuum source was considered the principal heating mechanism of the emitting gas. Significant \ciii\ emission suggested electron densities (\ne) in the range $10^9 - 10^{10}$ \cm3. The observed intensity ratio \ciii/\civ\ indicated ionization parameter ($U$; defined by Eq. \ref{eq:u} later in this paper) values of the order of $10^{-1}$. This photoionization scenario was successful in explaining at least some quasar optical and UV spectra (see the review by \citealt{davidsonnetzer79} for a synopsis). More recent work recognized the existence of several problems with the original scenario. Low ionization lines (LILs), and especially \feii\ are too strong to be explained by a photoionized region of moderate density and column density (see for example \citealt{dumontmathez81, joly87, collinsouffrinetal88, dumontcollinsouffrin90b}). These authors stressed that the LILs required a denser, low-temperature environment. Even more recently, \citet{baldwinetal96,laoretal97b,baldwinetal04} point towards high density at least for the LIL emitting zone. This low ionization broad line region (LIL BLR) has very similar properties to the O{\sc i} and Ca{\sc ii} emitting region identified by \citet{matsuokaetal08}. The region where these LILs are produced cannot emit much \ciii\ if the electron density exceeds 10$^{11}$ \cm3. BLR conditions are certainly complex and the assumption of a single emitting region cannot explain both LILs and high ionization lines (HILs) in all quasars \citep[][and references therein]{marzianietal10,wangetal11}. In \citet{negreteetal12}, hereafter \citetalias{negreteetal12}, we report an analysis based on several diagnostic ratios used to constrain density, ionization parameter and metallicity in the BLR of two sources that are representative of extreme Population A (narrow Line Seyfert 1 - NLSy1) sources. These sources showed weak \ciii\ emission (relative to \siiii) which simplified our interpretation of the emission line spectra. Diagnostic ratios indicate a very dense (\ne$\sim 10^{12}$ \cm3), low ionization ($U\sim10^{-2.7}$) region that, in a photoionization scenario, is expected to also emit \feii\ and \caii\ lines \citep[e.g.,][]{bruhweilerverner08,matsuokaetal08}. In \citet{negreteetal13}, hereafter \citetalias{negreteetal13}, we show that \ciii\ is not associated with the high density region. However \ciii\ is strong in most sources, dominating the emission of the 1900\AA\ blend. The presence of strong \ciii\ indicates that the BLR cannot be characterized anymore by a narrow range of density and ionization parameter: a gradient of ionization, of density, or both may be present. Since the BLR is not spatially resolved, the meaning of emission line ratios becomes much more ambiguous, and we cannot obtain meaningful single-value measures of \nh\ and ionization. Nonetheless, the ionizing photon flux (i.e., the product \nhu) can be retrieved with an accuracy comparable to estimates obtained from reverberation mapping, and then used to estimate \R\ \citepalias{negreteetal13}. In this paper we apply the method discussed in \citetalias{negreteetal13} to a pilot sample of high redshift ($z \sim$ 3) sources. The technique we present should allow one to easily compute \R\ for large samples of high $z$\ sources. Moderate resolution and high S/N dedicated observations allowed us to detect and measure faint and blended lines in order to analyze all physical information that can be retrieved from rest frame UV spectra of high-$z$ quasars. The pilot sample helps us to explore challenges that exist for studying high-$z$\ sources using only rest-frame UV spectra. Estimation of redshift cannot rely on measures of low-ionization narrow lines that are the most credible diagnostics at low z. \citep{eracleoushalpern03,huetal08}. Narrow lines are weak and often undetectable in very high luminosity sources \citep{marzianietal09}, a phenomenon sometimes called the \oiiiopt\ Baldwin effect; Zhang et al. 2011). On the other hand, the criteria for population and spectral type identification were set from properties of the \hb\ spectral spectral range \citep{sulenticetal00a} that is customarily not available for high $z$\ quasars. We present the spectra of 8 pilot sources obtained with the VLT/FORS1 in Section \ref{sec:observations}; Section \ref{sec:reduction} summarizes data reduction including details of redshift estimation (Section \ref{z}). Before discussing the analysis of the data, we present a synopsis of quasar systematics along the so-called ``eigenvector 1'' (E1), with special attention on the interpretation of the line profiles. In Section \ref{sec:data_analysis} we apply the interpretation of the line profiles derived from low-$z$ sources to fit the line profiles in this high z sample. In Section \ref{sec:component_analysis} we discuss the population assignment and the properties of each source, while Section \ref{sec:physical_conditions} explores BLR physical conditions, with special reference to metallicity issues (Section \ref{metals}). Section \ref{sec:results} presents results from application of the photoionization method to our pilot sample. Section \ref{rblr} presents derivations of the BLR radius (\rb; its distance from the ionizing source) and and black hole mass (\mbh). Section \ref{sec:discussion} compares our results with previous work and Section \ref{sec:conclusions} considers the the prospects for application of our technique for single epoch \mbh\ estimates in high redshift samples. All the computations were made considering $H_0$ =70 km s$^{-1}$ Mpc$^{-1}$ and relative energy densities $\Omega_\Lambda=0.7$ and $\Omega_\mathrm{M}=0.3$. | \label{sec:conclusions} In this paper we presented new observations of eight high redshift quasars. The spectra were meant to provide high S/N, moderate resolution data on which the \siiv, \civ, \siiii\ and \aliii\ emission line profiles could be accurately analyzed. Line profile fits permitted us to isolate a specific component whose intensity ratios were used to derive values of the ionizing photon flux. These results allowed us to compute the product \nhu\ and hence the size of the BLR and the central black hole mass. The method described in this paper rests on the assumption of photoionization as the mechanism of gas heating, on the assumption of isotropic luminosity, and on line ratios predicted by {\sc cloudy} simulations. We found that the \mbh\ derived from the computed \rb\ and from the virial assumption are in good agreement with the ones derived from the luminosity-size relation. The photoionization method explored in this paper offers an estimate of \rb\ for each quasar, with some advantages on the \rb\ valued derived from the luminosity-size correlation. The luminosity correlation suffers from large scatter and is simply extrapolated to very high luminosity without any support since there are, unfortunately, no conclusive results on reverberation of high luminosity quasars even if heroic efforts are underway \citep[e.g.][]{treveseetal07,kaspietal07,bottietal10}. We repeat that our \mbh\ and \rb\ results are based on the product \nhu\ and not on values of \nh\ and of $U$\ taken separately. To apply the photoionization method in the most effective way, determining \nhu\ with the lowest uncertainty, spectral data should be of moderate resolution ($\lambda/\Delta \lambda \sim 1000$) as well as of high S/N. If the \siii\ line can be measured in an accurate way, it would be possible to derive reliable estimates of $Z/Z_\odot$. Instead, we used \siiv+\oiv/\civ\ to constrain metallicity. Only for extreme Pop. A sources, when \siiii\ $\gtsim$ \ciii, it is possible to estimate the \nh, $U$\ and metallicity with very high S/N spectra. The present exploratory analysis emphasized several sources of uncertainty. However, the parameter needed for \rb\ and \mbh\ computation, the product \nhu, seems to be fairly stable and well-defined. Even with an error of a 0.3 in logarithm, the square root will be subject to a 0.15 uncertainty in logarithm, much lower than the uncertainty associated with the \rb\ -- luminosity correlation. The large intrinsic spread of the correlation at low luminosity, its uncertain extrapolation at very high luminosity make preferable a one-by-one determination based on physical properties of the emitting regions. | 14 | 4 | 1404.1673 |
1404 | 1404.6988_arXiv.txt | {\normalsize Recent BICEP2 detection of low-multipole B-mode polarization anisotropy in the cosmic microwave background radiation supports the inflationary universe scenario and suggests a large inflaton field range. The latter feature can be achieved with axion fields in the framework of string theory. We present such a helical model which naturally becomes a model with a single cosine potential, and which in turn reduces to the (quadratic) chaotic inflation model in the super-Planckian limit. The slightly smaller tensor/scalar ratio $r$ of models of this type provides a signature of the periodic nature of an axion potential. % We present a simple way to quantify this distinctive feature. As axions are intimately related to strings/vortices and strings are ubiquitous in string theory, we explore the possibility that cosmic strings may be contributing to the B-mode polarization anisotropy observed. } \vspace{1cm} \begin{flushleft} \today \end{flushleft} \end{center} \end{titlepage} \setcounter{page}{1} \setcounter{footnote}{0} \tableofcontents \parskip=5pt | Observations of the Cosmic Microwave Background (CMB) radiation provides strong support for the inflationary universe scenario that explains the origin of the hot big bang beginning of the universe \cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi}. CMB temperature anisotropies measured by COBE, WMAP and Planck are in excellent agreement with predictions of the simplest inflationary models \cite{Bennett:1996ce,Bennett:2012zja,Ade:2013uln}, which also predict a scale-invariant spectrum of gravitational waves \cite{Starobinsky:1979ty} that came from the quantum fluctuation of space-time during inflation \cite{Crittenden:1993wm,Frewin:1993dq,Harari:1993nb}. The gravity waves generate primordial polarization, including the so-called B-mode pattern \cite{Kamionkowski:1996zd,Seljak:1996gy,Zaldarriaga:1996xe,Kamionkowski:1996ks}. In particular, the recent data from BICEP2 \cite{Ade:2014xna} is consistent with predictions of the chaotic inflation (the quadratic version) model \cite{Linde:1983gd}. On the other hand, the B-mode signal seen by BICEP2 can also contain small contributions from other potential sources such as cosmic strings. The primordial microwaves are linearly polarized, which may be decomposed into E-mode and B-mode. Quantum (scalar) fluctuation of the inflaton (the scalar field responsible for inflation) leads to E-mode polarization while the quantum fluctuation of space-time metric during the inflationary epoch leads to both E- and B-mode polarizations. If the recent BICEP2 observation of the B mode polarization in CMB is confirmed, it would imply the presence of the tensor mode quantum fluctuation and so gravity is quantized. The only known consistent perturbatively quantized gravity theory is the (super)string theory. If string theory is the theory of nature, it should be able to explain the inflationary universe. Although there are a number of explicit realizations of the inflationary universe scenarios in string theory, a typical range of the inflaton field $\phi$, i.e., the field range that $\phi$ evolved during the inflationary epoch, is $\Delta \phi < M_{pl}$, where $M_{pl}$ is the reduced Planck mass. This property simply follows from the compactification of the extra dimensions present in string theory \cite{Baumann:2014nda}. Following from the Lyth bound \cite{Lyth:1996im}, \begin{equation} \frac{\Delta \phi}{M_{pl}} \ge N_{e} \sqrt{r/8} \end{equation} where $N_e$ is the number of e-folds of inflation and $r$ is the tensor/scalar ratio. Taking $60 \ge N_e \ge 40$, we find that typical values of $r$ satisfies $r < 0.005$ for $\Delta \phi < M_{pl}$. This is much smaller than $r \simeq 0.2$ observed by BICEP2 \cite{Ade:2014xna}. One way to obtain a relatively large $r$ value in string theory is to employ the axion fields for inflation, which has been explored under the name of natural inflation \cite{Freese:1990rb,Adams:1992bn,Randall:1995dj} or axion monodromy \cite{Silverstein:2008sg,McAllister:2008hb,Marchesano:2014mla}. A very good feature of axion-generated inflation is the presence of the (approximate) shift symmetry: $\phi \rightarrow \phi \, +$ constant, a property noted some years ago \cite{Adams:1992bn}. As the inflaton, this property can protect the axions from the so called $\eta$ problem, namely having too steep a potential. Being angular (or phase) fields, axions can perform helical-like or similar motions to extend the field range, thus allowing a larger effective field range for the inflaton \cite{Kim:2004rp,Kaloper:2008fb}. Since axions are ubiquitous in string theory, such scenarios should actually appear rather naturally. So it is not difficult to come up with models that can fit the existing data. Here we present a simple two-axion helical model. The model reduces to a single cosine potential, which in turn reduces to a model closely resembles the quadratic version of chaotic inflation. Following Ref\cite{Kallosh:2014vja}, we see that this feature is quite natural in a large class of axionic models in the supergravity framework. Because of the periodic nature of an axionic potential, this cosine model can have a slightly smaller value of $r$ than that from chaotic inflation. This deviation is quite distinctive of the periodic nature of the inflaton potential. Here we present a simple way to search for this periodic feature by quantifying its deviation from the $\phi^2$ chaotic inflation \cite{Linde:1983gd}, a model that fits the existing data reasonably well. All physical parameters such as the runnings of the power spectra indices can be expressed in terms of $r$ and the scalar power spectrum index $n_s$, including the parameter $\hat \Delta$, which measures the deviation of the cosine model from $\phi^2$ chaotic inflation, $$ {\hat \Delta} = 16 \Delta = r+ 4(n_s-1) $$ where $\phi^2$ chaotic inflation has $\hat \Delta=0$. We show that other quantities such as runnings of spectral indices have very simple dependences on $\hat \Delta$. We see that a downward shift of $r$ from its value in the $\phi^2$ model by as large as ${\hat \Delta}= - 0.03$ (or about $20\%$ of $r$) is possible. As the data improves, a negative value of $\hat \Delta$ can provide a distinctive signature for a periodic axionic potential for inflation. However, when there are many axions, or axion-like fields, with a variety of plausible potentials, the possibilities may be quite numerous and so predictions may be somewhat imprecise. Here, we like to point out that the presence of axions would easily lead to cosmic strings (i.e., vortices, fundamental strings and D1-strings), which may provide a relatively clean signature of string theory scenarios for the inflationary universe. This is especially relevant if cosmic strings come in a variety of types and tensions and maybe even junctions. It so happens that cosmic strings will generate some B-mode polarization as well. This provides the motivation to further explore cosmic strings along this direction. Although the helical inflation model generates a $r$ value consistent with the BICEP2 result, it typically yields a slightly smaller value. This may leave room for B-mode contributions coming from cosmic strings. Recent analyses of the BICEP2 and POLARBEAR data \cite{Ade:2014afa} suggests that some cosmic string contribution to the B-mode polarization is possible in the fitting of the B-mode power spectrum \cite{Lizarraga:2014eaa,Moss:2014cra}. In any case, since the B-mode spectrum from cosmic strings \cite{Pogosian:2003mz} is different from the B-mode spectrum from inflation, better B-mode polarization data will either provide evidence of cosmic strings or put a tight bound on its contribution. This paper is organized as follows. In Sec. 2, we present the helical model. Since the scale of the inflaton potential is essentially at the the grand unified theory (GUT) scale, it is possible that a phase transition may have taken place during inflation. We make some preliminary observations on this issue. Since the helical model reduces to a single cosine potential, we suggest a simple way to pick out the key feature of a periodic potential by comparing it to $\phi^2$ chaotic inflation. This analysis is presented in Sec. 3. In Sec. 4, we discuss some properties of cosmic strings in relation to the B-mode polarization. We give our conclusions in Sec. 5. | We find that the helical model reproduces the (quadratic version) chaotic inflation model while the $r$ value may be a little smaller due to the periodic nature of the axion potential. We quantify this deviation, a distinctive feature of axionic inflation, and argue that it can be tested and measured with better data. Phase transition during inflation will probably have little impact on $r$, or if anything, tends to lower its value. The data may allow some room for cosmic string contribution to the B-mode polarization. Since the B-mode power spectrum from cosmic strings differs substantially from that from inflation, better data offers the hope of detecting cosmic string signals or provides a tight bound on the possible cosmic string contribution. We thank David Chernoff, John Ellis, Mark Hindmarsh, Renata Kallosh, Andrei Linde, Levon Pogosian, Yoske Sumitomo, Alex Vilenkin and Ira Wasserman for discussions. {\it While this paper was in preparation, Ref\cite{Choi:2014rja} by Choi, Kim and Yun appeared with the same helical model described here. After our paper, we note that Ref\cite{Kappl:2014lra,Ben-Dayan:2014zsa,Long:2014dta} also discuss the helical or a similar model. See also Ref\cite{Kaloper:2014zba,Harigaya:2014eta,Hebecker:2014eua,Higaki:2014pja,Bachlechner:2014hsa} for related proposals. The revised version of Ref\cite{Kallosh:2014vja} gives a simple nice realization of the helical model within supergravity. } | 14 | 4 | 1404.6988 |
1404 | 1404.1420_arXiv.txt | The particular giant X-ray bump of GRB 121027A triggered by \emph{Swift} is quite different from the typical X-ray flares in gamma-ray bursts. There exhibit four parts of the observed structural variabilities in the rise and decay phase of the bump. Considering the quality of four parts of the data, we can only analyze the data from about 5300 s to about 6100 s in the bump using the stepwise filter correlation method (Gao et al. 2012), and find that the $86^{+5.9}_{-9.4}~\rm s$ periodic oscillation may exist, which is confirmed by the Lomb-Scargle method (Scargle 1982). Furthermore, a jet precession model (Liu et al. 2010) is proposed to account for such a variability. | Gamma-ray bursts (GRBs) are the most dramatic astronomical phenomena in the universe (Zhang \& M{\'e}sz{\'a}ros 2004; M{\'e}sz{\'a}ros 2006). The prompt emissions of GRBs last from a few milliseconds to thousands of seconds. The statistics of their durations are shown as a bimodal distribution (Kouveliotou et al. 1993), and therefore GRBs can be classified as short- and long-duration GRBs. The progenitors of them are believed to be the mergers of two compact objects (see e.g., Eichler et al. 1989; Paczy{\'n}ski 1991; Narayan et al. 1992) or the collapsars of massive stars (see e.g., Woosley 1993), respectively. Despite of the different progenitors, a rotating black hole (BH) surrounded by a neutrino-dominated accretion flow (NDAF, see e.g., Popham et al. 1999; Narayan et al. 2001; Gu et al. 2006; Liu et al. 2007, 2008, 2010a, 2010b, 2012a, 2012b, 2013; Sun et al. 2012; Li \& Liu 2013; Kawanaka et al. 2013; Xue et al. 2013) will be formed and therefore power GRBs via the neutrino annihilation or the BZ mechanism (Blandford \& Znajek 1977). A tilted accretion disc surrounding a supermassive BH leads to the precession of the BH and results in an S- or Z-shaped jet as observed in galaxies (e.g., Lu \& Zhou 2005). The jet precession caused by the system of the BH and disc can also explain some periodic variabilities of X-ray binaries, such as SS 433 (Sarazin et al. 1980). Quasi-periodic feature observed by BATSE in the gamma-ray lightcurves motivates the idea that the GRB jet may be precessed (see, e.g., Blackman et al. 1996; Portegies Zwart et al. 1999). In the central engine of GRBs, either the misalignment of angular momenta of two compact objects or the anisotropic fall-back mass in collapsar may induce the precession between the BH and the accretion flow. According to the Bardeen-Petterson effect (Bardeen \& Petterson 1975), we suggest that in the jet precession model the BH can capture the inner part of the NDAF to conform with the direction of the angular momentum, and the outer part of the NDAF drives the BH and inner part to precess (Liu et al. 2010a; Liu \& Xue 2012). The model can be used to explain the temporal structure and spectral evolution of GRBs (Liu et al. 2010a), to simulate almost all types of the gamma-ray lightcurve of GRBs (Portegies Zwart et al. 1999; Lei et al. 2007), and to predict the intensities of the gravitational waves from GRBs (Romero et al. 2010; Sun et al. 2012). However, compared with the observations of the gamma-ray emission, there is little X-ray observational evidence of precession in GRBs. In the paper, we will analyse the variability of the X-ray bump in GRB 121027A and discuss its possible origin by using our jet precession model. In Section 2, we describe the \emph{Swift}/XRT observations of GRB 121027A and use the stepwise filter correlation method, coupled with the Lomb-Scargle method, to present the analysis of the quasi-periodic X-ray lightcurve from about 5300 s to about 6100 s since trigger. In Section 3, the jet precession model is introduced. For the reasonable properties of the BH in the centre of collapsar, the model can explain the X-ray lightcurve of GRB 121027A. The conclusions and discussion are presented in Section 4. | We roughly analyze the timescales of oscillations in the four parts of the bump, which indicate that there may exist the time evolution in the bump of GRB 121027A, and find there may exist $86^{+5.9}_{-9.4}~\rm s$ periodic signals from about 5300 s to about 6100 s by using the SFC method. The result is confirmed by the Lomb-Scargle method. If the bump originates from the fall-back accretion process, we argue that the quasi-periodic oscillations may be caused by the jet precession in the BH-NDAF system. Our model (Liu et al. 2010a) can explain the behaviours, and the final properties of the BH are well consistent with the collapsar models. Thus it is a possible method to test or estimate the mass of the BH in the centre of a GRB due to the quasi-period oscillations of the lightcurve in the future observations. For the prompt emission, Gao et al. (2012) performed SFC method to 266 GRBs in BATSE sample. Although no quasi-periodic oscillations was claimed in their work, they indeed found that the majority of the bursts had clear evidence of containing a ``slow" variability component superposed on a rapidly varying time sequence. Furthermore, we searched all the \emph{Swift}/XRT samples. The light curves of most flares and afterglows are smooth, thus the quasi-periodic oscillations do not exist. For the ultra-long GRBs, we found a possible sample in the data of GRB 101225A. We used the SFC method to analyze its data from about 4950 s to about 7300 s, and did not find the periodic signals. Furthermore, Fan et al. (2005) suggested that if jet powering the late X-ray flares is launched via magnetic processes, such as GRB 050724, the radiation of the flares is expected to be linearly polarized. As well as the bump of GRB 121027A, given the requirement for accretion rate ($\sim 1.8 \times 10^{-4}~M_\odot~\rm s^{-1}$) in the jet precession model, the jet is possibly dominated by the magnetic field. The bump including quasi-periodic signals may be one of the astronomical candidate sources of linearly polarized. Future GRBs observations by the POLAR detector (Bao et al. 2012) may test this possibility. | 14 | 4 | 1404.1420 |
1404 | 1404.0801_arXiv.txt | \label{sec:introduction} The Near-Earth asteroid (NEA) population has a relatively large amount of data compared to other small body populations, including detailed information on asteroid figures and binary structure, often made possible through the combination of lightcurve and radar techniques. Observers have discovered a wide and complex set of asteroid systems that before this study have not been tied together into a coherent theory. The emergence of radiative forces as a major evolutionary mechanism for small bodies, in particular for NEA systems due to their small size and proximity to the Sun, makes the development of such a theory possible. [Figure 1] A simple model of NEA evolution constructed from the Yarkovsky-O'Keefe-Radzievskii-Paddack (YORP) and binary YORP (BYORP) effects, ``rubble pile'' asteroid geophysics, and gravitational interactions can incorporate all of the diverse observed asteroid classes as shown in Figure 1: synchronous binaries, doubly synchronous binaries, contact binaries, asteroid pairs, re-shaped asteroids, and stable ternaries. \subsection{Observed NEA Classes} \label{sec:observedclasses} [Table 1] Binary asteroid systems comprise a significant fraction ($15 \pm 4\%$) of the NEA population~\citep{Margot2002, Pravec2006} and include all compositional classes and size scales~\citep{Pravec2007}. Most of these systems are synchronous binaries--the orbital and secondary spin periods are equal, but the primary has a faster spin rate. Observed synchronous systems have mass ratios~$\lesssim0.2$, a system semi-major axis of $1.5$ to $3$ primary diameters, and a possibly elongated secondary and a nearly spherical primary with a distinctive shape characterized by an equatorial bulge~\citep{Pravec2006, Pravec2007}. The system has a positive free energy, but the tidally locked secondary inhibits disruption. Migration to the inner solar system from the main belt as binaries, binary creation via collision amongst NEAs, and binary creation via tidal disruption from close planetary flybys are not efficient enough mechanisms to create this population nor match the observed synchronous binary properties~\citep{Margot2002, Walsh2008a}. Several theories attempt to explain this binary population by a YORP-induced rotational fission process, but do not capture all properties of synchronous binaries and do not predict the other NEA systems that are seen~\citep{Scheeres2007b, Scheeres2009a, Scheeres2009b, Walsh2008b}. The~\citet{Walsh2008b} theory requires rotational fission induced ``landslides'' that re-shape the primary, then enter into orbit. Secondaries are built from collections of ``landslide'' material in orbit after many ``landslide'' events, and consequently YORP cycles--the length of the process to rotationally accelerate an asteroid to spin fission from its current state under the YORP effect. However, we will show that material entering orbit via rotational fission will almost always escape on timescales much shorter than a YORP cycle. Furthermore,~\citet{Holsapple2010} using continuum approximations of granular theory finds that mass loss would not occur at the equator of small, critically spinning asteroids, but that their shapes would deform, elongating the body until interior failure. These deformations in the shape of the body allow YORP to continue to increase the angular momentum without significant changes to the spin rate, even slightly decreasing the spin rate in some cases.~\citet{Scheeres2011} reports a similar analytic finding that when cohesive theory is considered, failure will most likely occur along interior planes. The analytic theory in~\citet{Scheeres2009b} describes the first stage of the model proposed herein, where a chaotic binary system is immediately formed from the rotational fissioning of a ``rubble pile.'' A rotational fission model related to the one proposed in this work has been implicated in the formation of asteroid pairs \citep{Pravec2010}--two asteroids with heliocentric orbits that in the recent past ($\lesssim 10^6$ yrs) intersect deep within the other's Hill radius and with small relative speeds \citep{Vokrouhlicky2008}. Asteroid pairs are observed in the Main Asteroid Belt with similar sizes to NEAs, but they have not been observed in the NEA population. The theory outlined in this paper predicts them; asteroid pairs are difficult to detect in the NEA population because their orbits are rapidly perturbed and smaller initial asteroids fission into even smaller secondaries for the same mass ratios as the Main Belt asteroid population. Other observed distinct dynamical and morphological classes include doubly synchronous binaries, high-$e$ binaries, ternaries, contact binaries, and re-shaped asteroids. We describe each in turn. Doubly synchronous binaries: all spin rotation periods are equivalent to the orbital revolution period. They also have high mass ratios $\gtrsim0.2$ and a system semi-major axis of $2$ to $8$ primary diameters \citep{Pravec2007}. These systems are difficult to detect because of an observational bias in light curve data; doubly synchronous systems and elongated single objects appear similarly. Contact binaries: bimodally-shaped asteroids observed as two similar-sized components resting on each other, which implies a formation mechanism that brings the two components together very gently. Contact binaries comprise a significant fraction ($>9\%$) of the NEA population \citep{Benner2006}. High-$e$ binaries: low mass ratio binary systems distinct from the synchronous binaries, because they are asynchronous and have high eccentricities \citep{Taylor2008}. Ternary systems: large primary orbited by two smaller satellites. The primaries are spinning faster than the orbital rates and the mass ratio is low ($<0.1$) \citep{Brozovic2009}. Re-shaped asteroids: single bodies similar to the primaries of the synchronous binary class--an oblate shaped figure with an equatorial bulge. For reference, examples of each NEA class are given in Table 1. \subsection{Motivation} \label{sec:motivation} A collisionally evolved asteroid can be modeled as a ``rubble pile"--a collection of gravitationally bound boulders with a distribution of size scales and very little tensile strength between them~\citep{Michel2001, Richardson2005, Tanga2009}. ``Rubble pile'' morphology has been closely examined by the Hayabusa mission to Itokawa, as shown in Fig. 2, which has no obvious impact craters and appears as collection of shattered fragments of different size scales~\citep{Fujiwara2006}. Mass and volume measurements from the NEAR Shoemaker flyby of Mathilde~\citep{Yeomans1997} and radar observations of 1999 KW$_4$~\citep{Ostro2006} determine mean densities that are lower than their constitutive elements, which is evidence of voids and cracks in the structures of these bodies. Asteroids with diameters larger than $\sim200$ m rarely spin with periods less than $\sim2.2$ hours, which corresponds with the critical disruption spin rate of self-gravitating, ``rubble pile'' bodies~\citep{Pravec2007B}. Both theoretical modeling and direct observation indicate that asteroids within a size range of $\sim100$ m to $\sim10$ km have ``rubble pile'' geophysics. The details of the rotational fission process determine the initial conditions for the binary system. The torque from the YORP effect will increase the centrifugal accelerations acting on each ``rubble pile'' component. There is a specific spin rate at which each component of the body will go into orbit about the rest determined by the largest separation distance of the mass centers of the fissioned component and the remainder of body~\citep{Scheeres2009a}. The smaller component is now the secondary, and the remainder is the primary, both in orbit about each other. The motivation for this study was to determine what happens dynamically after a rotational fission event. This paper will utilize some important concepts throughout that will be briefly introduced here and further defined later. The mass ratio is defined as the secondary mass divided by the primary mass. The primary of a binary system is always larger than the secondary, so the mass ratio is a number between $0$ and $1$. Secondary fission is rotational fission of the secondary induced via spin-orbit coupling and occurring during the chaotic binary stage of low mass ratio evolution creating chaotic ternaries. Ternary systems have three members. The components in decreasing mass are labeled primary, secondary, and tertiary, however the two smaller members are referred collectively as secondaries. Secondaries may escape the system if the system has a positive free energy. The free energy of an asteroid system is the sum of the kinetic and mutual potential energies of the system (both rotational and translational) neglecting the self-potentials of each body. | \label{conclusion} The evolution of NEA systems is driven by four important processes: initial rotational fission, secondary fission, impacts, and solar gravitational perturbations. The lower the mass ratio, the faster the spin rate required for initial rotational fission, and thus the more energy in the eventual binary system. The free energy transitions from positive to negative at a mass ratio of $0.2$ for the spherical end state, this divides the evolution of rotationally fissioned systems into two paths as shown in Fig. 1. Secondary fission can occur before low mass ratio systems are ejected. Enough energy is transferred to the secondary via spin-orbit coupling so that it undergoes rotational fission and creates a chaotic ternary as shown in Fig. 2. Secondary fission grows increasingly likely as mass ratio decreases, since the initial energy in the system increases and rotational energy transferred to the secondary is more effective on a less massive secondary. Chaotic ternaries are formed from secondary fission and evolve quickly back into a chaotic binary state, however impacts dissipate energy and produce more stable binaries. Escape of ternary members can also stabilize the system. Solar gravitational perturbations are important in changing the eccentricity and are responsible for both stabilizing and destabilizing binary systems. NEAs are actively evolving systems driven by these four processes and the observed asteroid classes are stages in this evolution. Radiative processes dominate the evolution of the NEA population from the Yarkovsky effect which drives small Main Belt asteroids into resonances with Jupiter pushing them into the NEA population, to the YORP effect which dominates their spin evolution and forces them to disrupt forming asteroid systems, to the BYORP effect which drives these systems back together or apart. The lives of NEAs are exciting--each asteroid may go through many iterations of the cycle shown in Fig. 1 taking different paths each time. \appendix \label{sec:appendix} | 14 | 4 | 1404.0801 |
|
1404 | 1404.2729_arXiv.txt | We discuss a simple and fast method for estimating masses of early-type galaxies from optical data and compare the results with X-ray derived masses. The optical method relies only on the most basic observables such as the surface brightness $I(R)$ and the line-of-sight velocity dispersion $\sigma_p(R)$ profiles and provides an anisotropy-independent estimate of the galaxy circular speed $V_c$. The mass-anisotropy degeneracy is effectively overcome by evaluating $V_c$ at a characteristic radius $R_{\rm sweet}$ defined from {\it local} properties of observed profiles. The sweet radius $R_{\rm sweet}$ is expected to lie close to $R_2$, where $I(R) \propto R^{-2}$, and not far from the effective radius $R_{\rm eff}$. We apply the method to a sample of five X-ray bright elliptical galaxies observed with the 6-m telescope BTA-6 in Russia. We then compare the optical $V_c$-estimate with the X-ray derived value, and discuss possible constraints on the non-thermal pressure in the hot gas and configuration of stellar orbits. We find that the average ratio of the optical $V_c$-estimate to the X-ray one is equal to $\approx 0.98$ with $11 \%$ scatter, i.e. there is no evidence for the large non-thermal pressure contribution in the gas at $\sim R_{\rm sweet}$. From analysis of the Lick indices H$\beta$, Mgb, Fe5270 and Fe5335, we calculate the mass of the stellar component within the sweet radius. We conclude that a typical dark matter fraction inside $R_{\rm sweet}$ in the sample galaxies is $\sim 60\%$ for the Salpeter IMF and $\sim 75 \%$ for the Kroupa IMF. | Being the most massive galaxies in the local Universe, giant elliptical galaxies provide a natural laboratory to study galaxy formation, assembly and evolution processes. The current paradigm of galaxy formation is the hierarchical scenario which suggests that early-type galaxies have complex merging histories of assembling most of the mass through accretion of small galaxies with rare major merger events \citep[e.g.][]{de.Lucia.Blaizot.2007, Naab.et.al.2007}. Accurate mass determinations and disentangling a luminous and dark matter components at different redshifts are the key steps towards a consistent theory for elliptical galaxies formation. Determining the mass profile of early-type galaxies is a notoriously difficult problem as there are no dynamical tracers with the known intrinsic shape and structure of orbits, so that circular velocity curves of elliptical galaxies cannot be measured directly. A number of methods are in use for constraining the mass of early-type galaxies and the shape of dark matter halos, each having its own set of assumptions and limitations. Comparison of the mass profiles obtained from different independent techniques is necessary to get reliable estimates. It also helps to control the systematic uncertainties, inherent in all methods, as well as leads to interesting constraints on properties of elliptical galaxies. One of the mass estimation techniques comes from X-ray observations of extended hot X-ray-emitting coronae of massive elliptical galaxies. It is a powerful tool to probe the mass distribution over several decades in radius: from $\sim0.1 R_{\rm eff}$ out to $\sim10R_{\rm eff}$. In this approach spherical symmetry of a galaxy and hydrostatic equilibrium of the gas are commonly assumed. While the spherical symmetry approximation introduces only a small bias, if any \citep[e.g.][]{Piffaretti.et.al.2003, Boute.Humphrey.2012c}, validity of the hydrostatic equilibrium assumption is the subject of debate. When one is able to quantify deviations from hydrostatic equilibrium, it allows to estimate (although indirectly) pressure of the non-thermal gas motions. Most simulations suggest that in relaxed systems hydrostatic approximation works well, with non-thermal support at the level of 5\% to 35\% of the total gas pressure \citep[e.g.][]{Nagai.Vikhlinin.Kravtsov.2007}. When X-ray observations are combined with, for instance, optical data on the stellar kinematics, then comparison between the X-ray gravitating mass profile and the optical mass allows one to estimate the magnitude of the non-thermal motions of the hot gas, to constrain the mass-to-light ratio, to disentagle stellar and dark matter contributions to the total gravitating mass profile and to characterize the distribution of stellar orbits. Although elliptical galaxies suffer from a lack of `ideal' traces like disc rotation curves in spiral galaxies and there is an inherent degeneracy between anisotropy and mass, studies of stellar kinematics and dynamics provide the tools for measuring the gravitating mass profile with sufficient accuracy (up to $\sim 15 \%$, \citealt{Thomas.et.al.2005}). Methods based on the Schwarschild modeling of stellar orbits in axisymmetric (or even triaxial) potentials are considered to be the state-of-the-art techniques in this field. The most sophisticated approaches operate with full information on the line-of-sight velocity distribution including higher-order moments. The orbit-based methods allow to infer not only the total mass profile, but also to measure the dark matter content, derive mass-to-light ratios and get the distribution function of stellar orbits. Among the drawbacks of these methods are the high computational cost and the necessity to have high-quality observational data. So only nearby elliptical galaxies can be studied by means of Schwarschild modelling, and for a large sample of objects, especially with noisy photometric and/or kinematical data, such an approach is not justified. In this paper we discuss a simple approach for estimating the mass from the stellar kinematics \citep{2010MNRAS.404.1165C, Lyskova.et.al.2012} that relies only on the most basic observables such as the surface brightness and line-of-sight velocity dispersion profiles. By design the method is simple and fast and has a modest scatter ($\disp \Delta V_{c} / V_{c} \sim 5-10 \% $, \citealt{Lyskova.et.al.2012}). This makes it suitable for large samples of elliptical galaxies even with limited and/or noisy observational data. Of course, the method is not intended to replace a thorough investigation of each indvidual galaxy. We apply the method to a small and rather arbitrarily selected sample of massive elliptical galaxies located at the centers of groups and clusters, and bright in X-rays. The surface brightness and projected velocity dispersion profiles up to several effective radii have been measured with optical long-slit spectroscopic facilities on the 6-m telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences (SAO RAS). Using publicly available Chandra data we also derive the X-ray mass profile and compare it with simple optical estimates. The paper is organized as follows. In Section~\ref{sec:method}, we provide a brief description of the method. We apply it to real elliptical galaxies in Section~\ref{sec:analysis}, starting with the illustration of the method on the example of the extensively studied giant elliptical galaxy M87 in Section~\ref{subsec:M87}. Details on the observations of the sample of early-type galaxies are presented in Section \ref{subsec:observations}. We derive circular speed estimates from optical and X-ray analyses and estimate stellar contrubutions in Section \ref{subsec:rotcurve}. Results are summarized in Section \ref{sec:discussion}, and Section \ref{sec:conclusion} contains conclusions. | \label{sec:conclusion} We discuss a simple mass estimator that allows one to estimate the circular speed $V_c$ from {\it local} properties of the surface brightness and the line-of-sight kinematics at a characteristic radius where the $V_c$-estimate is largely insensitive to the unknown anisotropy of stellar orbits. Although the method is designed for non-rotating spherical galaxies, we extend it also to mildly rotating axisymmetric and slowly-rotating triaxial ones, substituting $\sigma_p(R)$ in equation \ref{eq:main} or \ref{eq:agd_simple} with $\disp V_{\rm rms}(R)=\sqrt{\sigma_p^2(R)+V_{\rm rot}^2(R)}$, where $\disp V_{\rm rot}(R)$ is the rotational velocity. Tests on the sample of massive simulated galaxies show that the recovered from $I(R)$ and $\sigma_p(R)$ measured along apparent major and minor axes of the galaxy circular speed is almost unbiased with the RMS-scatter of $\sim 5 \%$. We apply the method to M87 and compare our simple mass estimate with circular speed profiles derived from X-rays and the state-of-the-art Schwarzschild modeling, thus revisiting the results of \cite{2008MNRAS.388.1062C,2010MNRAS.404.1165C}. At the sweet radius $R_{\rm sweet}=141''$ we derive $V_c^{\rm opt}=524$ $\kms$, that agrees well with the circular speed obtained in \cite{Murphy.et.al.2011}. After comparing the optical $V_c$-estimate with the X-ray derived one, we conclude that at the sweet radius the non-thermal contribution to the total gas pressure is $\sim 25 \%$. The true value of the non-thermal contribution in M87 could be even lower, since X-ray data near the sweet radius are affected by the shock \citep{Forman.et.al.2007}. We observe a sample of five X-ray bright elliptical galaxies with the 6-m telescope of the SAO RAS and measure the surface brightness, line-of-sight velocity and velocity dispersion distribution of stars up to two effective radii along one or two slits. We apply our simple method to estimate the circular speed and compare it with the circular speed measurements based on the X-ray analysis of Chandra data. We conclude that optical and X-ray $V_c$-estimates agree with each other remarkably well implying the sample averaged non-thermal pressure support of $\disp \sim 4\% \pm 20 \% $ , i.e. to be consistent with zero. From deep long-slit spectral data obtained with SCORPIO/BTA we derive high-precision Lick indices profiles out to $\sim 2$ effective radii, which in turn used to estimate the radial variations of the stellar population mass-to-light ratios and the dark matter fraction within $R_{\rm sweet}$, typical value of the latter is $\sim 60 \%$ for the Salpeter IMF and $\sim 75 \%$ for the Kroupa IMF. | 14 | 4 | 1404.2729 |
1404 | 1404.2942_arXiv.txt | We study the evolution of disc galaxies in group environments under the effect of both the global tidal field and close-encounters between galaxies, using controlled $N$-body simulations of isolated mergers. We find that close-range encounters between galaxies are less frequent and less damaging to disc galaxies than originally expected, since they mostly occur when group members have lost a significant fraction of their initial mass to tidal stripping. We also find that group members mostly affect disc galaxies \emph{indirectly} by modifying their common global tidal field. Different initial orbital parameters of group members introduce a significant ``scatter'' in the evolution of general properties of disc galaxies around a ``median'' evolution that is similar to when only the effect of the global tidal field is included. Close-encounters introduce a high variability in the properties of disc galaxies, even \emph{slowing} their evolution in some cases, and could wash out correlations between galaxy properties and the group total mass. The combined effect of the global tidal field and close-encounters appears to be inefficient at forming/enhancing central stellar bulges. This implies that bulges of S0 galaxies should be mostly composed by young stars, which is consistent with recent observations. | Galaxies experience different environments during their lifetimes, inhabiting regions of the Universe ranging in density from ``average'' to medium and high in environments such as groups and clusters of galaxies \citep[e.g.,][]{lacey1994,zhao2003,berrier2009,mcgee2009}. In these regions galaxies evolve under the influence of a combination of environmental and internal processes. Galaxies infalling onto a halo experience the effects of the global tidal field of their host \citep[e.g.][]{gnedin2003} which can vary gradually in intensity along their orbits or more abruptly if the halo grows rapidly or the galaxy's stellar mass grows significantly after accretion \citep[e.g][]{neistein2011}. Galaxies can also experience gravitational interactions with other halo galaxies or substructures in the form of direct mergers or close-encounters with low/high relative velocity \citep{moore1996,mastropietro2005}. For galaxies in dense environments, both the cold and hot reserves of gas can be removed hydrodynamically by the hotter intra-group or intra-cluster medium via processes such as ram-pressure stripping \citep{gunn1972} and starvation \citep{larson1980,balogh2009}, which can limit severely the capacity of galaxies to continue forming stars after accretion. Internal galactic processes, such as gas cooling, star formation, supernovae feedback, active nuclei feedback, formation of central bars also affect the evolution of galaxies and their contributions can be triggered/affected by environmental factors. Environmental effects are believed to play a crucial role in the evolution of galaxies in dense regions of the Universe. However, the precise mechanisms through which they affect galaxies remain unclear \citep[see the review by][]{weinmann2011-2}. They are thought to be the culprits behind the lower star formation, lower fraction of disc morphologies and redder colours observed in galaxies belonging to groups and clusters in comparison to regions of lower density \citep[e.g.,][]{dressler1997,lewis2002,gomez2003,balogh2004,weinmann2006,mcgee2008}. Recent studies, based on cosmological simulations combined with semi-analytic models, claim that a significant fraction of galaxies that currently reside in massive clusters have been accreted as part of galaxy groups \citep[][but see also \citeauthor{berrier2009}~\citeyear{berrier2009}]{mcgee2009,delucia2012}. These results emphasise the concept of galactic ``pre-processing'' \citep{zabludoff1998,lisker2013}. In this scenario, a considerable fraction of galaxies observed at present time in clusters, have had most of their properties shaped within group environments. This highlights the importance of understanding in detail the role of the environment in the evolution of galaxies within groups. However, theoretical studies on environmental effects within galaxy groups remain surprisingly few, while most of the theoretical effort has been largely focused on galaxies residing in either massive clusters \citep[e.g.][]{moore1999,gnedin2003,mastropietro2005} or in Milky Way-like haloes \citep[e.g.][]{mayer2001,mayer2001-2,klimentowski2009,kazantzidis2011}. Most of the recent studies on the evolution of galaxies in group-like environments by means of observations or numerical simulations \citep[][]{feldmann2011,tonnesen2012,bahe2013,bekki2013,taranu2013,vija2013,wetzel2013,ziparo2014} have concentrated on the formation of elliptical galaxies, the effect of ram-pressure stripping on galaxies by the intergalactic medium, and on the efficiency of galaxies to form stars while they inhabit a group. Interestingly, little attention has been paid to understand better the relative contribution of different environmental (and environmentally-triggered) processes (e.g. driven by the global tidal field, close-encounters between galaxies, rapid halo growth) to the overall evolution of galaxies in groups. This is a fundamental step in order to comprehend how galaxies currently residing in massive clusters have been ``pre-processed''. Close-encounters between galaxies are considered one of the main contributors to their evolution. Within groups, these encounters take place at lower relative velocity with respect to interactions within more massive clusters. Given the long duration of the gravitational perturbation, galaxy close-encounters in groups can induce significant changes in the structure of galaxies. \citet[][]{bekki2011} investigate the combined effect of the global tidal field of groups and repetitive close-encounters on the evolution of spiral galaxies, using chemodynamical simulations. They find that the combined tidal interaction can have a significant impact on the morphology of galaxies, inducing series of bursts of star formation that increase the stellar mass in the central bulges. These results suggest that groups are suitable environments where S0 galaxies can be formed, in line with implications from observational evidence \citep{wilman2009,just2010}. As mentioned above, however, the relative contribution of the global tidal field and close-encounters to the general evolution of galaxies in groups remains unclear. Group members (except the galaxy under study) are usually modelled as point mass objects. In addition, the use of standard Schechter functions for the stellar mass distribution of group members does not allow the evolution that group members experience in terms of their orbits, structure and mass content to be followed consistently, starting from stellar masses consistent with recent observations. Moreover, by adopting a significantly higher number of group members in comparison to observations \citep[as done for example in][]{bekki2011}, the contribution of close-encounters to galaxy evolution could be largely overestimated. All these factors are potentially relevant for the way galaxies interact with each other, and perhaps more importantly, for how galaxies might also affect the environment they inhabit. In groups, the co-evolution between a single galaxy and its environment should be more significant than in clusters, because of its larger contribution to the total mass of the environment \citep[see][]{aceves2013}. In the previous paper of this series \citep{villalobos2012}, we explore the evolution of disc galaxies within the global tidal field of a galaxy group by means of $N$-body simulations of isolated mergers, covering an ample parameter space of different orbits, disc inclinations, galaxy-to-group mass ratios, and presence of a central bulge. We find that the galaxy-to-group mass ratio and the initial inclination of disc galaxies at the time of accretion play a fundamental role in the evolution of galaxies within a group \citep[see also][]{bekki2013}. Specifically, we find that disc galaxies start suffering significant evolution due to the global tidal field only after the mean density of the group (within the orbit of the galaxy) exceeds 0.3--1 times the disc galaxy central density. Different inclinations of disc galaxies at the time of accretion cause them to experience significantly different structural evolution. In particular, retrograde infalls (inclination of 180$\degr$ with respect of their orbital plane) allow accreted galaxies to retain their initial disc structure for a significantly longer time in comparison to prograde infalls (inclination of 0$\degr$). Additionally, we find that the global tidal field of a group environment is not efficient at either inducing the formation of central bulges in stellar discs or enhancing existing ones at the time of accretion. This result suggests that the global tidal field alone cannot explain the formation of S0 galaxies in groups, and that additional processes (e.g., close-encounters, internal processes) are required to explain the formation of additional bulge stars. Finally, we find that more massive galaxies suffer more tidal stripping due to the fact that the stronger dynamical friction acting on them drags them rapidly to the densest region of the group, where they are exposed to stronger tidal forces. In this paper, we present the results of controlled $N$-body simulations of infalling disc galaxies as they are accreted onto a group-size halo that contains a central galaxy and a population of satellite galaxies. We explore different orbital parameters (consistent with cosmological simulations) for the infalling disc galaxy and various spatial, velocity and mass distributions for the population of satellite galaxies. Our main goal is to determine the contribution of low-velocity close-encounters to the evolution of galaxies within a group environment. Our approach is to compare simulations of the combined influence of global tidal field and close-encounters to simulations where only the influence of the global tidal field is included. As in the first paper of this series, our group environments are built from idealised initial conditions and do not account for hierarchical growth. In particular, our simulations do not consider cases where galaxies are accreted and evolve as part of ``sub-haloes'' within a group halo\footnote{\citet{delucia2012} estimate that about half of low ($9 < \log[M_*/M_{\sun}] < 10$) and intermediate ($10 < \log[M_*/M_{\sun}] < 11$) mass group members are accreted onto their final group when they are satellite galaxies. The remaining half being accreted while being central galaxies or isolated ones.}. Note also that our simulations do not include gaseous components and star formation, since we focus on the effect of close-encounters (and the group tidal field) on the stellar content of galaxies that is \emph{already present} at the time they are accreted onto a group environment. We plan to explore the effect of internal processes and the evolution of star formation in galaxies within groups in the next paper of this series. The layout of this paper is as follows: Section~\ref{sec-setup} describes the set-up of our experiments; Section~\ref{sec-descrip} describes our findings; in Section~\ref{sec-discussion} we discuss our results and in Section~\ref{sec-conclusions} we give our conclusions. | \label{sec-conclusions} In this work, we study the evolution of disc galaxies inhabiting a group environment, using controlled collisionless $N$-body simulations of isolated mergers. Our goal is to estimate the relative contributions of both the global tidal field and close-encounters between group members to the evolution of disc galaxies. We probe a parameter space that covers a number of relevant aspects of the galaxy-group interaction. Regarding the disc galaxies, we explore different initial inclinations (with respect to the disc rotation), different orbital eccentricities (consistent with cosmological simulations), and the presence of a central stellar bulge in the disc. Regarding the population of satellite galaxies in the group, we explore different number of members (consistent with observations of groups), variations in their stellar mass content (also consistent with observations), and variations in the total mass of the satellite population. Our fiducial disc galaxy resembles a bulge-less, slightly less massive Milky Way. Our main results are the following: \begin{itemize} \item Close-encounters in a group environment with the most likely number of members are found to be rare and gravitationally weak. Most of the encounters between galaxies occur at separations $\sim$50~kpc and relatively late in the simulations, when group members have already lost a sizeable fraction of their initial mass, reducing significantly their capacity to affect gravitationally the disc galaxy. Thus, in our simulations, the \emph{direct} effect of close-encounters between galaxies in groups is much less relevant than the influence of the global tidal field of the group. \item Within the group environment, the influence of other members on a given galaxy is found to be mostly \emph{indirect} by altering the global tidal field, as group members transfer to it mass (via tidal stripping) and energy/angular momentum (via dynamical friction). The mass added to the central region of the group shortens the merger timescale of an infalling galaxy, causing it to experience more structural transformations and mass stripping. On the other hand, the addition of energy/angular momentum provokes a displacement of the densest region of the group, causing a delay in the evolution of an infalling galaxy, as the time between pericentric passages about the group centre increases. \item We find that the evolution of general properties of disc galaxies (such as merger timescale, stellar mass loss, scale-length, mean thickness) caused by other group members is highly variable, depending mostly on the initial orbital parameters (positions and velocities) of group members. However, the ``median'' evolution of disc galaxies is found to be reasonably well approximated by the effect only due to the global tidal field (i.e. in absence of close-encounters). Interactions with group members introduce a ``scatter'' around the ``median'' evolution of disc galaxies, whose amplitude depends mostly on the combined total mass of the group members, their number and mass distribution. \item Counter-intuitively, disc galaxies can sometimes be \emph{less} affected by environmental effects after interacting with group members, i.e. galaxies can retain their initial structure and mass content for a longer time. The high variability of the properties of disc galaxies due to the influence of group members could also wash out underlying correlations with environmental properties, such as the group mass. \item The effect of the global tidal field combined with close-encounters between galaxies is inefficient at inducing or enhancing the formation of central bulges in stellar discs present at accretion time. This confirms our previous conclusion that, if S0 are formed from spiral galaxies preferentially in group environments, then their often large central bulges should contain mostly relatively young stars in comparison to their stellar discs. This result is found to be consistent with recent observations of S0 galaxies in the Fornax cluster. \item Finally, we find that prescriptions based on simulations that do not account for close encounters (e.g. those done to obtain merger timescales) remain valid and can indeed be used in the framework of galaxy formation models. More sophisticated implementations of those prescriptions could include a scatter dependent e.g. on the total stellar mass, mass distribution and number of nearby galaxies. \end{itemize} In the first two papers of this series we have studied the influence of the global tidal field and that of close-encounters on the evolution of disc galaxies within group environments, estimating their relative contributions. Thus far, we have focused on their effect on the stellar content of galaxies that is already present when disc galaxies are accreted onto a group. In the third paper of this series, we will examine the contribution of the gaseous components of disc galaxies in groups. | 14 | 4 | 1404.2942 |
Subsets and Splits
No community queries yet
The top public SQL queries from the community will appear here once available.