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1206
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1206.2614_arXiv.txt
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We are carrying out a bibliographic compilation of near-infrared (NIR) ($0.7-5.0$ $\mu$m) spectroscopic studies available for stars in the Galactic O Star Catalog (GOSC, Ma\'iz Apell\'aniz et al. 2004). This compilation allows us to quantify the precise degree of knowledge about NIR spectral information for GOSC sources, such as band coverage, spectral resolution, equivalent-width measurements, etc. This bibliographic compilation has a clear next step toward the development of a new catalog of O-type stars observed only in the NIR, which will be annexed to the GOSC. In this poster paper we present preliminary results derived from a set of different attributes extracted from the retrieved papers.
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\noindent There are many open questions about the spiral structure and stellar distribution in the Milky Way (MW). While the structure outlined by star forming regions (SFR) and molecular clouds is relatively well known, it is difficult to be drawn using only massive stars and young open clusters. In spite of the progress made, the spiral structure beyond 2~kpc from the Sun, as well as that on the far side of the Galaxy, is poorly determined. L\'epine et al. (2011) show how dramatic is this situation. Their Figures 7 and 9 plot the distribution of CS and maser sources (tracers of SFRs) in the Galactic Plane, in contrast with stellar optical/NIR tracers like Cepheids and young open clusters. Comparing these two distributions, we can infer that about 90\% of the spiral structure as traced by the youngest stellar populations is still completely unknown. On the other hand, the knowledge of the number and distribution of massive stars is intrinsically important because these stars play a crucial role in the dynamic and chemical evolution of the MW. Modern deep NIR sky surveys (e.g. VVV, Minniti et al. 2010; UKIDSS, Hewett et al. 2006) are opening a new window of galactic explorations. We now have the opportunity to discover an abundant population of hidden massive stars from the stellar candidates selected from those surveys. Thus, NIR spectroscopy, through spectral classification, is an indispensable tool allowing us to determine the nature of these candidates. At present, the primary source of knowledge about massive stars (with 99\% completeness at $B < 8$), is the {\it Galactic O Star Catalog} (GOSC, Ma\'iz Apell\'aniz et al. 2004). The GOSC collects information for the optically brightest galactic O-type stars (370 objects), providing coordinates, spectral types, optical and near-infrared photometry, and other useful information. Sota et al. (2008) presented the second version of GOSC, which seeks to extend the catalog to $B < 14$. Moreover, we are part of the {\it Galactic O-Stars Spectroscopic Survey} (GOSSS, Ma\'iz Apell\'aniz et al. 2011), this all-sky intermediate-resolution spectroscopic optical survey of all O-type stars, is designed to revolutionize the spectral classification system through high-quality and homogeneous spectroscopic observations of more than one thousand O-type stars. Based on GOSSS data, Sota et al. (2011) re-discussed the spectral classification system and presented a new atlas of O Stars. In the NIR regime, the general spectroscopic classification system for O-type stars is based on a few atlases (e.g. Hanson et al. 1996, 2005). These atlases constitute the tool to perform a Morgan-Keenan (MK) process of spectral classification in the NIR regime in a similar way as in the optical. Unfortunately, the quality and incompleteness of the published NIR spectroscopic atlases is not enough to reproduce the MK system in this regime, and many ambiguities are detected. We are planning a new NIR spectroscopic survey in order to establish the basis of a spectroscopic classification system for O-type stars in NIR comparable in quality and number of standards to that performed in the optical by Sota et al. (2011). As a first step, we are now retrieving bibliographic spectroscopic resources for the GOSC sources in the NIR domain, which allow us to quantify the precise degree of knowledge about NIR spectral information for GOSC entries, like band coverage, spectral resolution, equivalent-width measurements, etc. This bibliographic compilation has a clear next step toward the development of a new catalog of O-type stars observed only in the NIR (GOSC-IR), which will be annexed to the GOSC. In this poster paper we present preliminary results derived from a set of different attributes extracted from the retrieved papers. \begin{figure}[!ht] \centering \includegraphics[width=0.45\textwidth]{Torres-Simon-fig1.eps} \includegraphics[width=0.45\textwidth]{Torres-Simon-fig2.eps} \caption{ {\it Left:} Distribution of the number of published papers with NIR spectroscopy for O-type stars by year interval. The first published paper showing an $I$-band photographic spectrum of an O-type star corresponds to {\it Spectre infrarouges de quelques \'etoiles des premiers types entre 6500 et 8800 A} by Andrillat \& Houziaux (1967). {\it Right:} Distribution of the number of O-type stars by number of papers. Note that almost 50\% of GOSC sources have no published NIR spectroscopic information.} \label{fig1} \end{figure}
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1206.2614
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1206
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1206.5805_arXiv.txt
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This paper examines star formation (SF) in relatively massive, primarily early-type galaxies (ETGs) at $z\sim 0.1$. A sample is drawn from bulge-dominated \galex/SDSS galaxies on the optical red sequence with strong UV excess and yet quiescent SDSS spectra. High-resolution far-UV imaging of 27 such ETGs using \hst\ ACS/SBC reveals structured UV morphology in 93\% of the sample, consistent with low-level ongoing SF ($\sim 0.5 M_{\odot} {\rm yr}^{-1}$). In 3/4 of the sample the SF is extended on galaxy scales (25--75 kpc), while the rest contains smaller (5--15 kpc) SF patches in the vicinity of an ETG--presumably gas-rich satellites being disrupted. Optical imaging reveals that {\it all} ETGs with galaxy-scale SF in our sample have old stellar disks (mostly S0 type). None is classified as a true elliptical. In our sample, galaxy-scale SF takes the form of UV rings of varying sizes and morphologies. For the majority of such objects we conclude that the gas needed to fuel current SF has been accreted from the IGM, probably in a prolonged, quasi-static manner, leading in some cases to additional disk buildup. The remaining ETGs with galaxy-scale SF have UV and optical morphologies consistent with minor merger-driven SF or with the final stages of SF in fading spirals. Our analysis excludes that {\it all} recent SF on the red sequence resulted from gas-rich mergers. We find further evidence that galaxy-scale SF is almost exclusively an S0 phenomenon ($\sim 20\%$ S0s have SF) by examining the {\it overall} optically red SDSS ETGs. Conclusion is that significant number of field S0s maintain or resume low-level SF because the preventive feedback is not in place or is intermittent. True ellipticals, on the other hand, stay entirely quiescent even in the field.
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\label{sec:intro} The processes that control global star formation (SF) in galaxies are at the core of many current studies that aim to understand galaxy formation and evolution. Efforts are being made both to recognize a range of relevant processes and, perhaps more challenging, to establish which processes dominate in different types of galaxies, in different environments, and at different cosmic epochs. Since at the present time the galaxies are still growing and transforming, many of the mechanisms of SF regulation can be studied at lower redshifts where the data are of much higher quality. However, despite major advances in this area, there are still many open questions. One of the most enduring puzzles is presented by the evolution of massive early-type galaxies (ETGs), which comprise elliptical and lenticular (S0) galaxies. Their overall old stellar populations (e.g., \citealt{trager,delucia06}) and consequently the lack of current SF are to first order at odds with the otherwise successful model of hierarchical assembly of galaxies \citep{kauffmann96} in which the massive galaxies formed most recently. While their morphological and kinematical transformation from disk galaxies into spheroid-dominated ones is very well explained by dissipative major mergers (especially for massive ellipticals, \citealt{barnes}), the question of their subsequent quiescence is a separate question, and has started to be addressed only more recently with the introduction of various non-stellar feedback processes, especially the feedback from active galactic nuclei (AGN) (\citealt{springel}). It now appears that in addition to shutting down of star formation more or less concurrently with the morphological transformation, another feedback mechanism is required to keep an ETG free from subsequent SF \citep{croton1}. This requirement for a {\it maintenance} (or {\it preventive}) feedback is especially strong for field ETGs, which could be expected to continue accreting cold gas from the intergalactic environment (e.g., \citealt{gabor}). Such accretion is probably the primary source of gas for actively star forming galaxies as well \citep{keres}, without which they would not be able to sustain their observed star formation rates (SFR) for longer than a few Gyr \citep{larson80,kennicutt94,bauermeister}. While ETGs are {\it mostly} quiescent almost by definition, the question is are their {\it entirely} quiescent? What fraction is? If they exhibit SF, is it a new episode due to a fresh supply of gas, i.e., is galaxy being ``rejuvenated'' or are we seeing remnants of the original disk SF? Is, in the former case, the SF present because of the failure or absence of preventive feedback mechanism? Does SF tend to be wide-spread, as in spiral galaxies, or confined to circumnuclear regions? The dominance of old populations in ETGs makes the detection of relatively small amounts of SF intrinsically difficult. However, if SF could be systematically detected in ETGs, or, more generally, on the ``red sequence'', it would represent a potentially powerful way of identifying processes that lead to or prevent star formation. In contrast, these regulative processes are more difficult to study in actively star-forming (``blue cloud'') galaxies because of the high ``background'' of normal SF. The presence of SF, especially in the {\it central} regions, has been firmly established in some nearby ETGs; e.g., in the SAURON survey of 48 S0s and ellipticals \citep{combes,temi,shapiro,crocker}. Star formation on {\it galaxy scales}, which we refer to to as the {\it extended SF}, has received less attention, to some extent because of the difficulties in detecting it using optical methods. Detection of ionized emission can be challenging in ETGs where the strong continuum from old stars lowers the equivalent widths. Also, ionized emission can arise from a number of sources not associated with SF \citep{sarzi10}. Mid-IR emission (e.g., 24 $\mu$m from {\it Spitzer}/MIPS) in ETGs can have an order of magnitude stronger contribution from intermediate-age and older stellar populations than from the young stars \citep{salim09,kelson}. The most promising method is the UV emission from young massive stars, which is an order of magnitude more sensitive to recent SF than the blue optical flux \citep{kauff07}, and probes timescales that are closer to the current SF than those probed by the blue light from less massive, longer lived stars ($\sim 100$ Myr vs.\ $\sim 1$ Gyr)\footnote{We note that far-UV (FUV, $\lambda \sim 1500\AA$) is easier to interpret as SF than the longer wavelength near-UV (NUV, $\lambda \sim 2300\AA$, which in ETGs primarily comes from main sequence turn off stars and is strongly dependent on the stellar metallicity \citep{donas,dorman03,smith12}.}. Large-scale detection and characterization of ETGs in the UV was facilitated with the UV surveys of \galex \citep{martin05}. Initial \galex\ studies (e.g., \citealt{yi,rich}) based their approach on selecting large statistical samples of galaxies on the {\it red optical sequence}, which is where most ETGs are found. While most {\it optical} red galaxies remained red in the UV-optical colors, a significant fraction exhibited a UV excess, which \citet{yi} interpreted as low levels of ongoing SF\footnote{This underlines the importance of specifying the type of color (optical vs.\ UV-optical) when referring to the red sequence, as the optical red sequence will contain both truly quiescent galaxies and those with small relative amounts of SF.}. Since these studies culled their samples from SDSS spectroscopic survey in which galaxies are typically found at $z\sim 0.1$, little could be said about the morphology of the UV light (\galex\ has a resolution of 5$"$) and consequently whether this UV excess actually arose from young stars. Subsequent efforts focused on the UV morphology to confirm the presence of the {\it extended} SF in the optical red sequence and/or among individual ETGs \citep{donovan,thilker10,cortesehughes,sr2010,marino1,lemonias}. However, many questions remained: what is the origin of the star-forming gas in ETGs?; has this SF started recently or is it prolonged?; is SF related to the processes of disk building?; how the SF relates to the two types of ETGs: lenticulars (S0s) and ellipticals (Es). This paper and the accompanying work (Fang et al.\ 2012, Paper II) address these questions using a sample of 29 ETGs selected from SDSS and \galex\ surveys for which detailed far-UV images were obtained with the \hst. The initial analysis of this sample was presented in \citet{sr2010} (hereafter SR2010), where it was shown that these ETGs, selected on the basis of the presence of a strong UV excess, exhibit clear signatures of extended SF. The current work expands on the morphological and size-related aspects of the analysis of the \hst\ sample and appends it with the analyses of the general population of ETGs from SDSS and \galex. Paper II tackles star formation histories of this sample using surface brightness photometry, and also discusses the selection of more complete samples of ETGs with extended SF. The paper is structured as follows. Selection of the \hst\ sample is explained in \S\ \ref{sec:sample}, and the resulting UV observations and optical imaging data from SDSS and WIYN are described in \S\ \ref{sec:data}. Results are given in the subsequent five sections. In \S\ \ref{sec:uv}, \ref{sec:ha}, \ref{sec:opt} we provide morphological analysis based on the UV, \ha, and optical imaging, respectively, which shows that the majority of the sample has SF on large scales (tens of kpc), and that the SF is found exclusively in S0s and not in ellipticals. This is followed by two sections that place our sample in context. In \S\ \ref{sec:uv_opt} we analyze the relation between the UV and optical morphologies, showing that the star forming gas is preferentially acquired subsequent to a galaxy getting onto a red sequence, while in \S\ \ref{sec:s0_vs_e} we further explore differences between S0 and elliptical galaxies in terms of SF and determine the incidence rate of SF in the two types in the overall population. Discussion of the results in view of evolutionary scenarios is given in \S\ \ref{sec:disc}, and the findings are summarized in \S\ \ref{sec:conclusions}. Cosmology parameters $\Omega_m=0.3$, $\Omega_\Lambda=0.7$, $H_0= 70\, {\rm km\, s^{-1}\, Mpc^{-1}}$ are assumed throughout.
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\label{sec:conclusions} This study explores the morphology of $z\sim0.1$ optical red sequence galaxies (primarily early-type) with UV-detected star formation. ETGs were selected to have strong UV excess yet weak central ionized emission. Such selection encompasses the large majority of ETGs with extended, galaxy-scale SF. Here are the main findings: \begin{enumerate} \item There are two modes of SF in our sample. The dominant one is extended on scales similar to or larger than the optical extent of a galaxy (19 of 25, or 76\%). The secondary mode is one in which SF is concentrated in regions smaller than the optical extent. Such small-scale SF may represent cases in which the ETG assimilates a gas-rich dwarf. The remaining conclusions refer only to 19 ETGs with extended SF (ESF-ETGs). \item None of the ESF-ETGs is optically classified as a true elliptical galaxy. All show the presence of stellar disks, usually of S0 type. Based on the analysis of the general SDSS/\galex\ sample (for which Hubble types are not available) the incidence of extended SF is highest among the ETGs with $10.6<\log M_*<11.1$ and declines at both higher and lower masses; it also declines in ETGs with higher optical concentrations. These trends independently suggest that the extended SF is primarily a phenomenon of central (non-satellite) S0 galaxies. Tentatively, extended SF is estimated to be present in no more than 2\% of massive ellipticals (with fraction closer to zero not being excluded), as opposed to $\sim 20$\% for S0s. The latter is similar to the incidence rate of XUV disks in spirals, hinting at similar fueling mechanism(s). \item In all but one case the SF in ESF-ETGs takes place in UV rings with diameters of tens of kpc. The star-forming rings are not conspicuous in the optical (except one) and less than a half of hosts have an optical stellar bar. The latter indicates that the UV rings are not necessarily maintained by bar resonances. The dominance of UV rings is partially due to our selection criteria. However, even if we allow all ETGs with extended SF not covered by our selection to have non-ring morphology (e.g., in-filled disks), the UV ring incidence would still dominate ($\sim 2/3$). The morphology of UV rings varies and may indicate different sources of star forming gas or modes of SF. \item The morphological analysis suggests that the recent or ongoing IGM accretion is the likeliest dominant mechanism for the source of gas in ESF-ETGs having wide UV rings, while the fading of the original SF (the latest stage of quenching) may be responsible for SF in disks with central UV holes. The latter also show signatures of optically fading spiral structure. Galaxy interactions may fuel SF in several other cases, including the two extreme disks (diameters $\sim 70$ kpc), but in general the disturbances are not visible in the optical but only in the UV, suggesting at most very minor mergers. The ESF-ETGs with extreme disks resemble the giant LSB galaxies Malin 2 and UGC6614 and together with them may represent a distinct class of giant LSBs, which we call S0+LSBs. \item Analysis of disk sizes shows that in roughly half of the ESF-ETGs the optical disks are $\sim 50\%$ larger than in the underlying population of quiescent ETGs of the same mass, suggesting that significant disk build-up occurred in ESF-ETGs after the galaxy concluded the original epoch of SF. Such disk enhancement is consistent with long-lasting and relatively smooth accretion from the IGM and not the more recent merger-supplied gas. ESF-ETGs, whether their disks have been enhanced in the optical or not, on average have a larger ratio of UV to optical size than the comparable green valley galaxies, suggesting that the source of gas is more likely to be external (assuming that most green valley galaxies are fading onto the red sequence due quenching that cuts off the infalling gas). \item Altogether the IGM accretion may be responsible for SF in 55\% of the sample, followed by fading of the original SF in 25\% and minor mergers in the remaining 20\%. External origin of gas (from accretion and minor mergers) thus together accounts for 3/4 of the sample. The IGM accretion in these galaxies may have been present ever since the major SF ended, i.e., these galaxies may never have been truly passive but are instead suspended in the green valley (quasi-static IGM accretion). This is consistent with the fact that the plane of new SF coincides with the old disk. \end{enumerate} The results at which we arrive here point to a range of mechanisms driving SF in ETGs. Keeping in mind that our sample does not include filled-in disks, we find that the extended SF in most (but not all) cases looks like it is the result of a gradual, non-merger process. This is in general agreement with the results presented in Paper II, where it is shown that UV and optical colors do not favor bursty SF for most of our sample. Our conclusions primarily hold for extended SF. Small-scale central or circumnuclear SF, which our selection disfavors but could be as frequent as the extended SF, might instead preferentially be the result of minor mergers. The widespread occurrence of galaxy-scale SF in ETGs reported in this paper may appear to contrast the results from the SAURON sample, in which the extended SF is reported for only two out of 48 galaxies (NGC2974 and NGC2685). The primary reason behind this is in SAURON sample selection--which is 1/2 S0s and 1/2 ellipticals by design, with a disproportionate fraction being cluster galaxies. When taken into account that the extended SF is common only in {\it field} S0s, the small numbers of extended SF in the SAURON sample are not surprising. The incidence of galaxy-wide SF will certainly be higher in the ATLAS$^{\rm 3D}$ sample, which includes all ETGs within 42 Mpc, but even then it will be the UV, rather than the emission line maps, that will have the requisite surface brightness sensitivity to detect low levels of SF. The absence of true ellipticals from our sample of ESF-ETGs is intriguing and underlines the importance of distinguishing between ellipticals and S0s, especially when considering SF. Results from other studies seem to support this: ellipticals of intermediate and high mass, whether rotating fast or slow, do not show strong evidence of either extended or central current SF. The traditional view of all non-dwarf ellipticals as being ``red and dead'' is therefore basically upheld even in field populations. In contrast, many (perhaps even the majority, if the process is episodic) field S0s maintain or reacquire some level of SF activity.
| 12
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1206.5805
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1206
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1206.2939_arXiv.txt
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We present an update to the search for a non-trivial topology of the universe by searching for matching circle pairs in the cosmic microwave background \cite{Cornish:2003db} using the WMAP 7 year data release. We extend the exisiting bounds to encompass a wider range of possible topologies by searching for matching circle pairs with opening angles $10^\circ\le\alpha\le 90^\circ$ and separation angles $11^\circ\le\theta\le180^\circ$. The extended search reveal two small anomalous regions in the CMB sky. Numerous pairs of well-matched circles are found where both circles pass through one or the other of those regions. As this is not the signature of any known manifold, but is a likely consequence of contamination in those sky regions, we repeat the search excluding circle pairs where both pass through either of the two regions. We then find no statistically significant pairs of matched circles, and so no hints of a non-trivial topology. The absence of matched circles increases the lower limit on the length of the shortest closed null geodesic that self-intersects at our location in the universe (equivalently the injectivity radius at our location) to $98.5\%$ of the diameter of the last scattering surface or approximately $26$Gpc. It extends the limit to any manifolds in which the intersecting arcs of said geodesic form an angle greater than $10^\circ$.
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The search for a non-trivial topology of the universe has enjoyed a long and fascinating history. Using different methods -- from searching for specific topologies to the more general circles in the sky approach \cite{Cornish:1997ab} -- the cosmic microwave background (CMB) has been analyzed extensively, looking for any signs that light from the same object reaches us by more than one path (see \cite{Levin:2001fg} for a review of the various suggested methods, including \cite{Cornish:1997hz,Cornish:1997ab,Cornish:1997rp,Cornish:2003db,ShapiroKey:2006hm, Bond:1996qs, Bond:1997ym, Souradeep:1998yk, Bond:1999te, Bond:1999tf, LachiezeRey:1995kj, Lehoucq:1996qe, Luminet:1999qh, Uzan:1998hk, Uzan:1999de, Lehoucq:2000hf, Riazuelo:2003ud,Weeks:2003xq,Riazuelo:2006tb,Niarchou:2006be,Niarchou:2007nn} ). So far, all specialized efforts to detect specific topologies as well as the search for matching opposing circles in the sky have failed to detect any sign of a non-trivial topology of our universe. The circles-in-the-sky method, which we adopt in this paper, is based on the following intuitive picture. For illustrative purposes, assume that the true topology of the universe is a 3-torus, with unit cell size smaller than the Hubble horizon (see Figure~\ref{fig:schematic_setup}). This can be thought of as a tiling of flat space by identical cubes. An observer, such as ourselves, performing a series of cosmic microwave background (CMB) observations somewhere in one of the cubes, has clones identically located in each of the other cubes performing the identical series of observations. Centered around the observer is the 2-sphere of the surface of last scattering at $z\approx 1100$, with CMB fluctuations imprinted on it. Around the clone of the observer on the right, there is another 2-sphere of the surface of last scattering. The intersection of both 2-spheres is given by a circle. Both observers will look at the same ring of temperature fluctuations -- albeit from different ``sides''. Both observers are in fact identical, so an observer will see a matching pair of circles: one to the left and to the right. Hence, comparing temperature fluctuations along circles potentially yields information about the topology of our universe. Going away from toroidal geometries, it becomes immediately clear that the separation angle $\theta$ (the angle between the centers of the pair of matching circles) need not be $180^\circ$. Also, depending on the orientability or non-orientability of the manifold, circle pairs might have matching temperature fluctuations either both going clockwise around the circles (non-orientable) or one going clockwise and the other anti-clockwise (orientable). So far, the search for matching anti-podal circles, i.e. circles with separation angles $180^\circ$, or nearly anti-podal circles \cite{Cornish:2003db}, like other topology searches, has only yielded lower limits on the size of the Universe, and then only for ``nearly flat'' topologies. \begin{figure} \includegraphics[width=\linewidth]{horizon_torus} \caption{Schematic geometry of the circles-in-the-sky method. If the topology of the universe is a torus, the observer, located at the center of the 2-sphere of the surface of last scattering, sees matching circles of temperature fluctuations on opposite sides of the CMB sky.} \label{fig:schematic_setup} \end{figure} In this work, we apply the circles-in-the-sky statistics to searches for circles pairs of all opening angles $10^\circ\le\alpha\le90^\circ$, and integer separation angles $11^\circ\le\theta\le 180^\circ$ with both orientations. This extends the previous searches~\cite{Cornish:2003db,ShapiroKey:2006hm} to cover almost all possible topologies. We find what seems to be a systematic effect at two special positions in the sky that produces spurious signals for osculating circles. Otherwise, we see no evidence of non-trivial topology.
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We employed the circles-in-the-sky statistics first devised in \cite{Cornish:2003db}, looking for pairs of matching circles of opening angles $10^\circ<\alpha<90^\circ$ and separation angles $11^\circ\le\theta\le180^\circ$. We positioned the circle centers on a grid with $N_\text{side}=128$, but computed the statistics on the full $N_\text{side}=512$ CMB map. While the WMAP 7 year data brought quite some improvements in the noise of the $\Smax$ statistics (c.f. Figure 2 in \cite{Cornish:2003db}), we find no hints of a non-trivial topology of the universe (see Figures~\ref{fig:w_orientable_exclude_galactic_anticenter2_5_exclude_special_position2_5} and \ref{fig:w_non_orientable_exclude_galactic_anticenter2_5_exclude_special_position2_5b}). The new search covered a much wider range of possible topologies, and by extending the search to circles with opening angles as small as $10^\circ$, we have extended the previous bound on the size of the Universe to $98.5\%$ of the diameter of the last scattering surface, or approximately $26$Gpc. There are systematic effects coming from both members of a circle pair touching either the galactic anti-center or the position $l=109.44^\circ$, $b=27.8^\circ$ (see Figures~\ref{fig:w_orientable_skymap} and \ref{fig:w_non_orientable_skymap}). As these positions appear both when looking for orientable and non-orientable manifolds, they cannot be of topological origin, but point towards a contamination of the map at these positions. The galactic anti-center region contains significant amounts of galactic emission. While the ILC maps used in this analysis attempt to remove most of this emission, the correlated residuals are a likely source of contamination in the circle searches. We are looking forward to the data release of the Planck mission, which will offer an exciting new, sharper view of the surface of last scattering, allowing for a better search of signs of non-trivial topology by removing noise particularly at smaller separation angles $\alpha<30^\circ$. Further advances in computing power will enable a search on a full $N_\text{side}\ge 512$ grid of circle positions.
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1206.2939
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1206
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1206.5034_arXiv.txt
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We compare the cosmological first-order post-Newtonian (1PN) approximation with the relativistic cosmological linear perturbation theory in a zero-pressure medium with the cosmological constant. We compare equations and solutions in several different gauge conditions available in both methods. In the PN method we have perturbation equations for density, velocity and gravitational potential independently of the gauge condition to 1PN order. However, correspondences with these 1PN equations are available only in certain gauge conditions in the perturbation theory. Equations of perturbed velocity and the perturbed gravitational potential in the zero-shear gauge exactly coincide with the Newtonian equations which remain valid even to 1PN order (the same is true for perturbed velocity in the comoving gauge), and equations of perturbed density in the zero-shear gauge and the uniform-expansion gauge coincide to 1PN order. We identify other correspondences available in different gauge conditions of the perturbation theory.
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Einstein's gravity is generally accepted as the gravity to handle astronomical phenomena. The theory holds a remarkable track record in the solar-system test based on vacuum Schwarzschild solution and the parameterized post-Newtonian approximation where the gravitational fields are supposed to be weak. It is true that Einstein's theory has not failed in any experimental test based on modern scientific and technological development up till today, but it is also true that there has been no experimental test of the theory in the strong gravitational field and in large scale even including the galactic scale. Einstein's gravity is generally accepted in cosmology mainly based on its successes in other astronomical and Earth bound tests and the theory's own prestige associated with Einstein's fame and historical legacy. A self-consistent treatment of cosmological world model is possible in Einstein's gravity. Without a lead by Einstein's gravity, however, the spatially homogeneous and isotropic cosmological world model based on Newton's gravity is known to be incomplete and indeterminate (Layzer 1954; Lemons 1988). Despite such troubles in the background world model, evolution of perturbations in Newton's theory is known to be quite successful in reproducing the corresponding results in Einstein's gravity (Lifshitz 1946; Bonnor 1957; Noh \& Hwang 2004; Jeong et al. 2011). Considering the action-at-a-distance nature of Newton's gravity, such a coincidence is a non-trivial result. Differences between the two theories, however, appear as the scale approaches the horizon. We will address this issue in this work. If we accept Einstein's theory in analyzing the large-scale cosmic structure in current era, we have two methods available. One well known method is the perturbation theory where all dimensionless deviations in the metric and the energy-momentum tensor from the background world model are assumed to be small. If we accept only linear order deviations we have the linear perturbation theory. The perturbation theory assumes small deviation but is fully relativistic, generally applicable in all scales including the super-horizon scale and to particles with relativistic velocities (Lifshitz 1946; Harrison 1967; Nariai 1969; Bardeen 1980; Peebles 1980; Kodama \& Sasaki 1984; Bardeen 1988; Mukhanov et al. 1992; Ma \& Bertschinger 1995). The other less known method is the cosmological post-Newtonian (PN) approximation where all dimensionless deviations are assumed to be weakly relativistic with ${GM \over Rc^2} \ll 1$ ($M$ and $R$ are characteristic mass and length scales) and for a virialized system ${v^2 \over c^2} \ll 1$ ($v$ is the characteristic velocity involved). The first-order PN (1PN) approximation makes expansion up to ${GM \over Rc^2} \sim {v^2 \over c^2}$ order. The PN approximation assumes small relativistic effects and is applicable only in sub-horizon scale. But the equations derived are applicable to fully nonlinear situations (Chandrasekhar 1965; Futamase 1988; Tomita 1988; Futamase 1989; % Tomita 1991; Futamase 1993a, 1993b; Shibata \& Asada 1995; Asada \& Futamase 1997; Tanaka \& Futamase 1999; Hwang et al. 2008, PN2008 hereafter). Therefore, the two methods are complementary to each other. If we encounter cosmological situations where both nonlinearity as well as relativistic effects are important we may need full-blown numerical relativity implemented in cosmology. Currently such a general relativistic numerical simulation in cosmology is not available. The nonlinear perturbation analysis, being based on perturbative approach, is not sufficient to handle the genuine nonlinear aspects of structure formation accompanied with self-organization and spontaneous formation of structures. In order to handle the relativistic nonlinear process in cosmology we believe the post-Newtonian approach is currently practically relevant to implement in numerical simulation. We can find cosmological situations where the cosmological post-Newtonian approach, being weakly relativistic but fully nonlinear, might have important applications. Especially, the current cosmological paradigm favors a model where the large-scale structures (requiring the relativistic treatment) are in the linear stage, whereas small-scale structures are apparently in fully nonlinear stage. The often adopted strategy is to assume the small-scale nonlinear structures are fully under control by the Newtonian gravity. In the galactic and cluster scales we have the general relativistic measure ${GM \over Rc^2} \sim {v^2 \over c^2} \sim 10^{-6} - 10^{-4}$, thus small but nonvanishing, and indeed the 1PN (weakly relativistic) assumption is quite sufficiently valid. Thus, we believe the 1PN approach would be quite relevant to estimate the general relativistic effects in the nonlinear clustering processes of the galaxy cluster-scale and the large-scale structures. In this work we will compare the two relativistic methods in the matter dominated era: the 1PN method vs. linear perturbation theory. The 1PN approximation is based on previous studies in Chandrasekhar (1965) and PN2008, whereas the linear perturbation theory is based on previous studies in Bardeen (1988), Hwang (1994), and Hwang \& Noh (1999). We will compare the equations and solutions derived in the two methods. In both methods we have gauge degrees of freedom which need to be fixed by the gauge conditions. The Newtonian perturbation theory appears as the zeroth-order PN (0PN) approximation. Thus, we naturally also have Newtonian theory for the comparison. To 1PN order we will show that the equations for density, velocity and gravitational potential do not depend on the gauge conditions with each variables gauge-invariant to 1PN order. In the perturbation theory, however, the perturbation variables for density, velocity, potential, curvature, and other kinematic variables (like expansion and shear) do depend on the gauge conditions adopted. Our emphasis in this work is on the correspondences between the PN variables and perturbation theory variables based on different gauge conditions. For the velocity and gravitational potential the Newtonian (0PN) equations are valid even to 1PN order, whereas for the density variable we have 1PN correction terms: see Equations (\ref{PN-delta-eq1})-(\ref{PN-U-eq}). The zero-shear gauge will be shown to be distinguished by showing the perturbed velocity and gravitational potential having exact correspondences with the 0PN (thus 1PN as well) equations, and the perturbed density showing correct correspondence with the 1PN equation. The zero-shear gauge can be contrasted with the comoving gauge where the perturbed density and velocity have exact Newtonian correspondences even to the second order perturbations in all scales whereas the gravitational potential vanishes to the linear order. Other PN correspondences of the perturbation theory are summarized in Section \ref{sec:correspondences}.
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1206.5034
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1206.2834.txt
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This paper presents a first observational investigation of the faint Of?p star NGC 1624-2, yielding important new constraints on its spectral and physical characteristics, rotation, magnetic field strength, X-ray emission and magnetospheric properties. Modeling the spectrum and spectral energy distribution, we conclude that NGC 1624-2 is a main sequence star of mass $M\simeq 30~M_\odot$, and infer an effective temperature of $35\pm 2$~kK and $\log g=4.0\pm 0.2$. Based on an extensive time series of optical spectral observations we report significant variability of a large number of spectral lines, and infer a unique period of $157.99\pm 0.94$~d which we interpret as the rotational period of the star. We report the detection of a very strong - $5.35\pm 0.5$~kG - longitudinal magnetic field $\bz$, coupled with probable Zeeman splitting of Stokes $I$ profiles of metal lines confirming a surface field modulus $\langle B\rangle$ of $14\pm 1$~kG, consistent with a surface dipole of polar strength $\gtrsim 20$~kG. This is the largest magnetic field ever detected in an O-type star, and the first report of Zeeman splitting of Stokes $I$ profiles in such an object. We also report the detection of reversed Stokes $V$ profiles associated with weak, high-excitation emission lines of O~{\sc iii}, {which we propose may form in the close magnetosphere of the star.} We analyze archival Chandra ACIS-I X-ray data, inferring a very hard spectrum with an X-ray efficiency $\log L_{\rm x}/L_{\rm bol}=-6.4$, a factor of 4 larger than the canonical value for O-type stars and comparable to that of the young magnetic O-type star $\theta^1$~Ori C and other Of?p stars. Finally, we examine the probable magnetospheric properties of the star, reporting in particular very strong magnetic confinement of the stellar wind, with $\eta_*\simeq 1.5\times 10^4$, and a very large Alfven radius, $R_{\rm Alf}=11.4~R_*$.
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The detection \citep[e.g.][]{2009MNRAS.400L..94G}, empirical characterization \citep[e.g.][]{2011MNRAS.416.3160W} and theoretical modelling \citep[e.g.][]{2012MNRAS.tmpL.433S} of a growing sample of magnetic O-type stars is leading to a new, refined picture of the scope and impact of magnetic fields in high-mass stars. O-type stars are unique laboratories for investigating the physics of stellar magnetism. Magnetic fields have clear influence on their rotation rates \citep[rotation periods of most detected magnetic O stars are significantly longer than those of non-magnetic O stars of similar spectral types; e.g.][]{2009MNRAS.392.1022U, 2007MNRAS.381..433H, 2010MNRAS.407.1423M}. Evolutionary models \citep[][]{2003A&A...411..543M, 2004A&A...422..225M} and, recently, observations (Briquet et al., MNRAS, submitted) of massive stars suggest that the internal rotation profile is strongly modified by the presence of a magnetic field, enforcing essentially solid-body rotation throughout the bulk of the outer radiative zone. Magnetic fields have clear and fundamental effects on the structure, dynamics and heating of the powerful radiative winds of O stars \citep[e.g.][]{2002ApJ...576..413U, 2012MNRAS.tmpL.433S}. The lives of magnetic O-type stars are therefore expected to differ significantly from those of their non-magnetic brethren. The subject of the present paper, NGC 1624-2\footnote{According to the numbering system of \citet{1979A&AS...38..197M}.}, is an Of?p star \citep{Walbetal10a} and the main ionizing source of the open cluster NGC 1624 (the H~{\sc ii} region S212). It is also one of only eight O-type stars in which magnetic fields have been detected with confidence. The classification Of?p was introduced by \citet{1972AJ.....77..312W} to describe spectra of early O-type stars exhibiting the presence of C~{\sc iii} $\lambda 4650$ emission with a strength comparable to the neighbouring N~{\sc iii} lines. Well-studied Of?p stars are now known to exhibit periodic spectral variations (in Balmer, He~{\sc i}, C~{\sc iii} and Si~{\sc iii} lines), narrow P Cygni or emission components in the Balmer lines and He~{\sc i} lines, and UV wind lines weaker than those of typical Of supergiants (see \citet{2010A&A...520A..59N} and references therein). With our report of a detection of a magnetic field in NGC 1624-2, magnetic fields have now been firmly detected in all of the 5 known Galactic members of this class - HD 191612: \citet{2006MNRAS.365L...6D}, \citet{2011MNRAS.416.3160W}; HD 108: \citet{2010MNRAS.407.1423M}; HD 148937: \citet{2008A&A...490..793H}, \citet{2012MNRAS.419.2459W}; CPD -28\, 2561: \citet{2012IBVS.6019....1H}, Wade et al. (2012b); NGC 1624-2: This paper - prompting the inference that there is a direct physical relationship between the magnetic field and the Of?p characteristics. According to the recent analysis by \citet{2011MNRAS.411.2530J}, NGC 1624 is a young open cluster located significantly above the Galactic plane. Their analysis yields a heliocentric distance of $6.0\pm 0.8$~kpc \citep[in agreement with that found in the pioneering study of][]{1979A&AS...38..197M} and an age of no greater then $\sim 4$~Myr. In addition to NGC 1624-2, which is by far the brightest cluster member, 3 other apparently bright optical sources are located within 2 arcmin of the cluster centre. These sources are reported to be an early B-type main sequence star, and two probable F giants (which are probably not physically associated with the cluster). The current paper provides a first observational characterization of NGC 1624-2, the faintest of the known Galactic Of?p stars ($V=11.8$). In particular we report the physical parameters of the star, its rotational period, the detection of its magnetic field, and a preliminary characterization of its magnetic and magnetospheric characteristics. In Sect. 2 we describe the spectroscopic and spectropolarimetric observations upon which we base our results. In Sect. 3 we provide a short overview of the spectral properties of the star. In Sect. 4 we derive the physical properties of the star and its wind. In Sect. 5 we determine the spectral variation period of the star. In Sect. 6 we describe the magnetic field diagnosis and our constraints on the surface magnetic field. In Sect. 7 we characterize its X-ray properties based on archival data. In Sect. 8 we describe the probable magnetospheric properties of the star. In Sect. 9 we discuss our results and summarize our conclusions. % ======================================================================== % ======================================================================== %\begin{table} %\caption{\label{spectroscopy}Log of spectroscopic observations used to determine the period. $B$ refers to the spectral range that includes He\,{\sc i}~$\lambda$4471, He\,{\sc ii}~$\lambda$4542, He\,{\sc ii}~$\lambda$4686, and H$\beta$; $V$ to the range that includes %He\,{\sc i}~$\lambda$5876; and $R$ to the range that includes H$\alpha$. J stands for the low-resolution observations from \citet{2011MNRAS.411.2530J}, G for the intermediate-resolution GOSSS observations \citep{2011hsa6.conf..467M}, and N for the high-resolution NoMaDS observations \citep{2011arXiv1109.1492M}. {\bf [Unable to get bibtex to mark these as 2011a/b. Will need to modify .bbl file just prior to submission.]}} %\begin{center} %\begin{tabular}{cccc} %Date & $B$ & $V$ & $R$ \\ %\hline %2006-09-06 & J & J & J \\ %2006-09-08 & J & J & J \\ %2007-01-26 & J & J & J \\ %2008-10-14 & G & & G \\ %2009-09-28 & G & & \\ %2009-09-29 & G & & \\ %2009-09-30 & G & & \\ %2009-11-01 & G & & G \\ %2009-11-03 & G & & G \\ %2009-11-17 & J & J & J \\ %2009-11-24 & G & & G \\ %2011-08-22 & N & N & N \\ %2011-09-09 & & G & G \\ %2011-10-02 & N & N & N \\ %2011-10-03 & N & N & N \\ %2011-11-03 & N & & \\ %2011-11-10 & & N & N \\ %2011-11-13 & & N & N \\ %2011-11-17 & & N & N \\ %2011-11-18 & N & & \\ %2011-11-19 & N & & \\ %2012-01-15 & & N & N \\ %2012-01-20 & & N & N \\ %2012-01-22 & & N & N \\ %2012-01-24 & & N & N \\ %2012-01-27 & & N & N \\ %2012-02-11 & G & G & G \\ %2012-02-15 & G & G & G \\ %2012-03-04 & & N & N \\ %2012-03-12 & N & & \\ %2012-03-22 & & N & N \\ %\hline %Total N & 7 & 13 & 13 \\ %Total G & 9 & 3 & 7 \\ %Total J & 4 & 4 & 4 \\ %\hline %\end{tabular} %\end{center} %\end{table} \begin{table} \caption{\label{spectroscopy}Log of spectroscopic observations used to determine the period. $B$ refers to the spectral range that includes He\,{\sc i}~$\lambda$4471, He\,{\sc ii}~$\lambda$4542, He\,{\sc ii}~$\lambda$4686, and H$\beta$; $V$ to the range that includes He\,{\sc i}~$\lambda$5876; and $R$ to the range that includes H$\alpha$. J stands for the low-resolution observations from \citet{2011MNRAS.411.2530J}, G for the intermediate-resolution GOSSS observations \citep{2011hsa6.conf..467M}, and N for the high-resolution NoMaDS observations \citep{2011arXiv1109.1492M}.} % {\bf [Unable to get bibtex to mark these as 2011a/b. Will need to modify .bbl file just prior to submission.]} \begin{center} \begin{tabular}{cccc | cccc} Date & $B$ & $V$ & $R$ & Date & $B$ & $V$ & $R$ \\ \hline 2006-09-06 & J & J & J & 2011-11-13 & & N & N \\ 2006-09-08 & J & J & J & 2011-11-17 & & N & N \\ 2007-01-26 & J & J & J & 2011-11-18 & N & & \\ 2008-10-14 & G & & G & 2011-11-19 & N & & \\ 2009-09-28 & G & & & 2012-01-15 & & N & N \\ 2009-09-29 & G & & & 2012-01-20 & & N & N \\ 2009-09-30 & G & & & 2012-01-22 & & N & N \\ 2009-11-01 & G & & G & 2012-01-24 & & N & N \\ 2009-11-03 & G & & G & 2012-01-27 & & N & N \\ 2009-11-17 & J & J & J & 2012-02-11 & G & G & G \\ 2009-11-24 & G & & G & 2012-02-15 & G & G & G \\ 2011-08-22 & N & N & N & 2012-03-04 & & N & N \\ 2011-09-09 & & G & G & 2012-03-12 & N & & \\ 2011-10-02 & N & N & N & 2012-03-22 & & N & N \\ 2011-10-03 & N & N & N & {\bf Total N} & 7 & 13 & 13 \\ 2011-11-03 & N & & & {\bf Total G} & 9 & 3 & 7 \\ 2011-11-10 & & N & N & {\bf Total J} & 4 & 4 & 4 \\ \hline \end{tabular} \end{center} \end{table} \begin{figure*} \begin{centering} \includegraphics[width=16cm]{twoepochsplotsm.ps} %\vspace{3in} \caption{\label{spectral lines}Two NoMaDS spectra of NGC 1624-2 (top, obtained on 22 Aug 2011, high state, phase 0.92; middle, obtained on 11 Nov 2011, low state, phase 0.43). The bottom spectrum represents the difference (high minus low). The spectra have been convolved to a resolving power $R=10\, 000$ for display purposes.} \end{centering} \end{figure*} \begin{figure} \begin{centering} \includegraphics[width=8cm]{CiiiNiii.eps} %\vspace{3in} \caption{\label{CNIII}Of?p-diagnostic emission lines of C~{\sc iii} and N~{\sc iii} near maximum emission at the full ESPaDOnS resolving power ($R=65\,000$). Note the remarkable composite emission profiles with broad and narrow components, not seen in previously investigated members of the Of?p category.} \end{centering} \end{figure}
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\label{conclusion} \begin{figure} \includegraphics[width=85mm]{Halpha.ps} \caption{\label{halpha}Peak H$\alpha$ emission of NGC 1624-2 (solid black line) compared to peak emission of HD 191612 (dashed blue line) and $\theta^1$~Ori C (dot-dashed red line). } \end{figure} We have discovered an extraordinarily strong magnetic field (maximum mean longitudinal magnetic field $\bz=5.35\pm 0.5$~kG, corresponding to a dipole of surface polar strength $\sim 20$~kG) in the Of?cp star NGC 1624-2 that distinguishes it qualitatively from other known magnetic O-type stars. Of particular interest is the presence of clear Zeeman signatures in individual spectral lines, and the apparent detection of resolved Zeeman splitting (corresponding to a maximum mean magnetic field modulus of $\langle B\rangle =14\pm 1$~kG). The detectability of Zeeman splitting should in principle allow a much stronger constraint on the geometry and topology of the star's magnetic field as compared to any other magnetic O-type star, once additional observations are acquired. We also suggest that NGC 1624-2 may be a good target for future transverse Zeeman effect (Stokes $Q$ and $U$) observations. With its sharp, magnetically-split lines, the $QU$ signatures could potentially have amplitudes comparable to Stokes $V$. Using an extensive spectroscopic data set, we performed a first determination of the physical properties of the star. We confirm that it is a main sequence object with a mass comparable to those of the other Of?p stars. While the models used to infer the stellar properties included non-LTE effects, they did not directly include effects of the magnetic field. Such effects may include desaturation of profiles of spectral lines (modifying line blanketing), and introduction of Lorentz forces, both of which may lead to modification of the hydrostatic structure of the atmosphere \citep[e.g.][]{2009IAUS..259..407S, 2009IAUS..259..405S}. While the intrinsic uncertainties associated with the derived quantities are probably sufficiently large to dominate the observable consequences of these phenomena, considering the remarkable strength of the magnetic field of NGC 1624-2 their potential importance should be explored in more detail. Indeed, the spectrum of NGC 1624-2 exhibits a number of peculiarities that distinguish this star from other Of?p stars. First, we have observed unprecedented composite profiles of the C~{\sc iii} $\lambda 4650$ complex, with narrow and broad components. We hypothesize that the broad components are photospheric as in normal Of or Ofc stars \citep{Walbetal10a}, while the narrow components typical of Of?p spectra are magnetospheric. Consistent with this interpretation, we observe that only the narrow components appear to vary with phase. Moreover, the narrow C~{\sc iii} emission does not disappear at minimum, so unlike most other members of this class, NGC~1624-2 remains Of?p at both extreme phases. We have also found that the absorption lines of NGC 1624-2 are very narrow, and that the widths of C~{\sc iv} $\lambda\lambda 5801, 5811$ can be reproduced essentially by magnetic broadening. This is in stark contrast to other Of?p stars for which the profiles of these lines are clearly dominated by turbulent broadening at the level of several 10s of \kms. The origin of this difference is currently unknown. We have also used the strong variations of various spectral absorption and emission lines to infer a unique and unambiguous spectral variation period of $157.99\pm 0.94$~d. Based on the observed behaviour of all other known magnetic O-type stars, it is reasonable to assume this to be the stellar rotation period (an assumption that {will} be tested by acquisition of additional magnetic field measurements). This implies that NGC 1624-2 rotates very slowly, spinning once in approximately one-half year. Such a conclusion is in good agreement with the negligible $v\sin i$ inferred from modelling the magnetically-split line profiles. It has been established that all known magnetic O-type stars have diverse but relatively long periods of rotation, from about a week \citep[HD 148937; ][]{2008AJ....135.1946N} to perhaps more than 50 years \citep[HD 108; ][]{2001A&A...372..195N}. The most common mechanism invoked to produce such slow rotation is magnetic braking, i.e. the shedding of rotational angular momentum via the stellar wind and enhanced lever arm provided by the magnetic field. If we employ the braking model described by \citet{2009MNRAS.392.1022U} (Eq. 25 of that paper), we can {roughly} compute the braking timescale $\tau_{\rm spin}$ as a function of the magnetic field strength, stellar mass and radius, mass loss rate and terminal velocity. Using the physical parameters reported in Table~\ref{param_summary} and moment of inertia coefficient $k\sim0.1$ \citep{2004A&A...424..919C}, we obtain a spin-down time for NGC 1624-2 of 0.24\,Myr. As discussed in Sect. 1, the estimated maximum age of the cluster NGC 1624 is no greater than 4~Myr, i.e. no more than $\sim 17\tau_{\rm spin}$. Such an age is, however, easily sufficient for the star to have braked from an initial short rotation period to its current very long period. (For example, if we assume the star was initially rotating at critical ($P_{\rm crit}=1.1$~d), the time required to slow the rotation to 158~d would be just $\ln (158/1.1)=4.96\tau_{\rm spin}=1.2$~Myr.) {On the other hand, if the cluster age is significantly younger than 4 Myr this would place constraints on the initial rotational speed of the star (requiring a slower initial rotation) or a reconsideration of the origin of the current slow rotation. These conclusions are subject to the assumptions and limitations of the braking model (e.g. aligned magnetic and rotation axes computed in 2D) and the uncertainties of the input parameters ($B_{\rm d}$, $M_*$, $R_*$, $\dot M$, $v_\infty$ and $k$).} {In our analysis of the optical spectrum of NGC 1624-2, we derived a upper limit on the photospheric N abundance: {([N/H]$\ltsim$0.3)}. We consider this upper limit to be uncertain due to our present inability to directly include the influence of the strong magnetic field on the line formation in our NLTE spectrum synthesis model. Accurate knowledge of the surface N abundance represents an important constraint on the interior rotation profile of a magnetic early-type star. As reported by \citet{2011A&A...525L..11M}, when magnetic braking occurs in a massive star characterized by internal differential rotation, a strong and rapid mixing occurs in layers near the surface. This results in enhanced mixing, resulting in enriched surface abundances of nitrogen relative to similar models with no magnetic braking. However, when solid-body rotation is imposed in the interior, the star is slowed so rapidly that surface enrichments are in fact smaller than in similar models with no magnetic braking. Therefore magnetic braking can enhance the surface N abundance or, in contrast, have little effect on it, depending on the internal rotation profile of the star. \citet{2012A&A...538A..29M} investigated the N surface enrichment of 6 known magnetic O-type stars, including 3 Of?p stars. Depending on the assumed initial rotation velocity of the Of?p stars, they found that that they display surface nitrogen abundances consistent with those of non-magnetic O stars (if their rotation was initially rapid, $v_{\rm rot}\sim 300$~\kms), or enhanced relative no non-magnetic O stars (if their rotation was initially modest, $v_{\rm rot}\sim$ a few times $10$~\kms). Assuming NGC 1624-2 has a main sequence age no greater than 4 Myr, comparison of our preliminary N abundance with the figures presented by \citet{2012A&A...538A..29M} indicates that (1) NGC 1624-2 has a N abundance slightly lower than non-magnetic O stars with similar positions on the HR diagram; (2) that the abundance is low relative to that expected for a star of this age and mass if it was rotating initially at high velocity; (3) but consistent with the expected abundance if the star was rotating initially at lower velocity. However, as discussed above these results are quite tentative, and require more detailed modelling and a more robust determination of the age of the star before firm conclusions can be drawn.} As a consequence of its intense magnetic field, NGC 1624-2 is expected to host a magnetospheric volume substantially larger than any other magnetic O-type star. The inferred magnetic wind confinement parameter, $\eta_\star=1.5\times 10^4$, is 300 times larger than that of the Of?p star with the next-strongest field, HD 191612. This leads to a predicted Alfven radius of 11.4~$R_\star$ (versus 2.2~$R_*$ in the case of HD 191612). This much larger volume of confined plasma should result in much stronger magnetospheric emission \citep[e.g. according to the mechanisms discussed by][]{2012MNRAS.tmpL.433S}. Indeed, the H$\alpha$ emission of NGC 1624-2 is found to be substantially stronger than that observed in any other magnetic O-type star. For comparison, and as illustrated in Fig.~\ref{halpha}, the maximum EW of H$\alpha$ in the spectrum of HD 191612 is about 4~\AA\ \citep[e.g.][]{2011MNRAS.416.3160W}, while that of $\theta^1$~Ori C is about 2~\AA\ \citep[][]{2008A&A...487..323S}. The peak EW of the H$\alpha$ line of NGC 1624-2 is {26~\AA}, {\em 6.5 times greater than HD 191612 and 13 times greater than that of $\theta^1$~Ori C.} Modeling of this remarkable H$\alpha$ emission is an urgent priority, and first attempts are underway by Sundqvist, ud Doula et al. (priv. comm.). It is expected that such strong magnetic wind confinement in the presence of such a powerful wind should lead to intense X-ray emitting shocks. Analysis of archival {\em Chandra} ACIS-I X-ray observations indicates a hard and luminous X-ray spectrum ($\log(L_\mathrm{X}/L_\mathrm{bol})\sim -6.4$), qualitatively consistent with theoretical expectations as well as the behaviour of $\theta^1$~Ori C. However, the current observations are more or less equally consistent with a relatively hard, weakly extinguished source, or a relatively soft, highly extinguished source. New higher-quality X-ray observations will be required to draw useful quantitative conclusions about the X-ray properties of NGC 1624-2. Due to the paucity of magnetic data, the magnetic topology and geometry of NGC 1624-2 are at present only weakly constrained. Although the geometry (i.e. the inclination angle $i$ and the magnetic obliquity $\beta$) is normally inferred from the variation of the longitudinal magnetic field, observations of other magnetic O stars (e.g. $\theta^1$~Ori C, HD 191612, HD 57682) show that a clear relationship exists between the EW variations of optical emission lines diagnostic of the wind (e.g. H$\alpha$, He~{\sc i} $\lambda 5876$, He~{\sc ii} $\lambda 4686$) and the longitudinal field variation. In particular, the emission extrema of these stars all correspond to extrema of the (sinusoidally-varying) longitudinal field, with the emission maximum corresponding to the maximum unsigned longitudinal field. If we assume that the magnetic topology of NGC 1624-2 is roughly dipolar, we can expect with reasonable confidence that the ESPaDOnS observation, acquired at phase 0.96 near emission maximum, corresponds to the approximate maximum of the longitudinal field. The longitudinal field measured from the Narval observation, acquired at phase 0.28, is not very different from that of the ESPaDOnS measurement, suggesting that the variation of the longitudinal field is not very large, and moreover that it does not change sign as the star rotates. This implies that we view essentially only one magnetic hemisphere during the rotation of the star, and consequently $i+\beta<90\degr$. This conclusion is supported by the single-wave nature of the EW variations. When both magnetic hemispheres are visible during the stellar rotation \citep[i.e. $i+\beta\gg 90\degr$, as in the case of HD 57682; ][ and MNRAS, submitted]{2009MNRAS.400L..94G}, the EW variations of the emission lines exhibit a double-wave variation. In contrast, smaller values of the sum $i+\beta$ produce single-wave variations \citep[e.g. $\theta^1$~Ori C, HD 191612; ][]{2008A&A...487..323S,2007MNRAS.381..433H}. Clearly, NGC 1624-2 holds a special place amongst the known magnetic O-type stars. Its extreme magnetic field and puzzling spectral peculiarities have much to teach us, and demand immediate attention. Plans for observational follow-up - including optical spectroscopy, spectropolarimetry, and photometry; UV spectroscopy; and X-ray spectroscopy - are already underway, as are theoretical investigations of its H$\alpha$ and X-ray properties to better understand the magnetospheric geometry, structure and energetics. %\begin{figure} %\includegraphics[width=80mm]{fig_xrays.ps} %\caption{\label{fig:xray} ACIS-I spectra of NGC\,1624-2, binned to obtain at least one count per bin (black dots, error bars and shaded area). The three models described in Tab.~\ref{tab:xray} are shown in red (solid), green (dashed) and blue (dotted). The C-statistic for the fits are 13.7, 20.5 and 21.6, respectively. $N_H$ is $10^{22}$~cm$^{-2}$.} %\end{figure}
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1206.2108_arXiv.txt
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We present new far-infrared ($70-500\,\mu$m) Herschel PACS and SPIRE imaging observations as well as new mid-IR Gemini/T-ReCS imaging (8.7 and $18.3\,\mu$m) and spectroscopy of the inner Lindblad resonance (ILR) region ($R<2.5\,$kpc) of the spiral galaxy NGC~1365. We complemented these observations with archival Spitzer imaging and spectral mapping observations. The ILR region of NGC~1365 contains a Seyfert 1.5 nucleus and a ring of star formation with an approximate diameter of 2\,kpc. The strong star formation activity in the ring is resolved by the Herschel/PACS imaging data, as well as by the Spitzer $24\,\mu$m continuum emission, [Ne\,{\sc ii}]$12.81\,\mu$m line emission, and 6.2 and $11.3\,\mu$m PAH emission. The AGN is the brightest source in the central regions up to $\lambda \sim 24\,\mu$m, but it becomes increasingly fainter in the far-infrared when compared to the emission originating in the infrared clusters (or groups of them) located in the ring. We modelled the AGN unresolved infrared emission with the {\sc clumpy} torus models and estimated that the AGN contributes only to a small fraction ($\sim 5\%$) of the infrared emission produced in the inner $\sim 5\,$kpc. We fitted the non-AGN $24-500\,\mu$m spectral energy distribution of the ILR region and found that the dust temperatures and mass are similar to those of other nuclear and circumnuclear starburst regions. Finally we showed that within the ILR region of NGC~1365 most of the on-going star formation activity is taking place in dusty regions as probed by the $24\,\mu$m emission.
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\begin{table*} \centering \begin{minipage}{100mm} \caption{Summary of Herschel photometry observations}\label{tab:her_obs} \begin{tabular}{lcccc} \hline Obsid & Instrument & Bands & Start time & Duration\\ & & (\micron) & (UTC) & (s)\\ \hline 1342201436 & SPIRE-P & 250, 350, 500 & 2010-07-14 20:20:45.0 & 999\\ 1342222495 & PACS-P & 70, 160 & 2011-06-11 12:59:25.0 & 1217\\ 1342222496 & PACS-P & 70, 160 & 2011-06-11 13:20:45.0 & 1217\\ 1342222497 & PACS-P & 100, 160 & 2011-06-11 13:42:05.0 & 1217\\ 1342222498 & PACS-P & 100, 160 & 2011-06-11 14:03:25.0 & 1217\\ \hline \end{tabular} \end{minipage} \end{table*} The fueling of the nuclear and circumnuclear activity of galaxies has been a topic of extensive discussion. Such activity not only includes the accretion of matter onto a supermassive black hole with the accompanying active galactic nucleus (AGN), but also the presence of intense nuclear and circumnuclear starbursts. Both types of activity require gas to be transported from the host galaxy on physical scales of a few kiloparsecs down to less than one kiloparsec for the nuclear starburst activity and even further in for the nuclear activity. Interactions, mergers, and large-scale bars, among others, have been proposed as possible mechanisms to transport gas from kiloparsec scales to the nuclear and circumnuclear regions \cite[see the review by][and references therein]{Jogee2006}. In isolated galaxies with a large-scale bar the gas is believed to flow inwards between corotation and the inner Lindblad resonance (ILR). Indirect evidence of this is the presence of star formation rings and hot spots near the ILR of barred galaxies \citep[see the review of][]{Buta1996}. The direction of the flows inside the ILR is generally outwards. This implies the existence of a gravity torque barrier at this resonance. However, if there is a ``spatially-extended'' ILR region, this translates into the existence of an inner ILR (IILR) and outer ILR (OILR). Numerical simulations predict that in this case a gas ring can be formed in between these two resonances. Furthermore, in the scenario of a double ILR the dynamical decoupling of an embedded nuclear bar can drive the gas further in \citep{Shlosman1989, Hunt2008, GarciaBurillo2009}. The combined action of gravity torques due to embedded structures (bars-within-bars) and viscous torques could be a viable mechanism to drive the gas to the inner few parsecs and feed the AGN \citep[][]{GarciaBurillo2005, Hopkins2011}. NGC~1365 is a giant isolated barred galaxy at a distance of 18.6\,Mpc \citep[][therefore $1\arcsec=90$\,pc]{Lindblad1999}. This galaxy hosts a Seyfert 1.5 nucleus \citep{Schulz1999}. \citet{Jungwiert1997} showed that in NGC~1365 there is also a nuclear bar ($R<10\arcsec$) embedded in the the large-scale bar ($R\sim 100\arcsec$). \citet[][]{Lindblad1996} used hydrodynamical simulations to reproduce the kinematics and the offset dust lanes along the large-scale bar of this galaxy with an outer ILR at a radius of $R_{\rm OILR}= 27\arcsec = 2.4\,$kpc. There is also an inner ILR at a radius of $R_{\rm IILR} \sim 0.3\,$kpc. Henceforth we refer to the ILR region as the region interior to the OILR of NGC~1365. As predicted by simulations, there is a ring of star formation inside the ILR region of this galaxy. The star formation activity in the central regions of NGC~1365 has been revealed by the presence of hot spots \citep{Sersic1965}, intense H$\alpha$ emission \citep{Alloin1981, Kristen1997}, non-thermal radio continuum sources associated with H\,{\sc ii} regions and supernova remnants \citep{Sandqvist1995, Forbes1998}, large amounts of molecular gas \citep{Sakamoto2007}, and point-like and diffuse extended X-ray emission not associated with the AGN \citep{Wang2009}. Moreover, there is evidence that a significant fraction of the circumnuclear star formation activity might be obscured based on the prominent dust lane crossing the nuclear region (see Fig.~\ref{fig:PACS100_large}, right panel) and the bright and extended mid-infrared (mid-IR) emission in this region \citep{Telesco1993, Galliano2005}. \begin{figure*} \subfloat{\includegraphics[width=0.46\textwidth]{figure1a.ps}} \subfloat{\includegraphics[width=0.51\textwidth]{figure1b.eps}} \caption[PACS100_large]{Herschel/PACS $100\,\mu$m image (left panel) of NGC~1365. The image is shown in a square root intensity scale. The displayed FoV is $360\arcsec \times 360\arcsec$ and matches that of the optical image shown in the right panel. The latter reproduces the BVR image shown in figure~1 of \cite{Elmegreen2009} (reproduced by permission of the AAS), for an easy comparison of the large-scale optical and far-IR morphologies. The square shows a FoV$=60\arcsec\times60\arcsec$ as in Fig.~\ref{fig:IRcentralimages}. The small ellipse on the optical image represents the approximate size of the ILR region studied in this work, whereas the large ellipse shows the corotation radius. The horizontal bar represents 10\,kpc. For both images, orientation is north up, east to the left. } \label{fig:PACS100_large} \end{figure*} In this paper we present new far-infrared (far-IR) imaging observations of NGC~1365 performed with the ESA Herschel Space Observatory \citep{Pilbratt2010} and new mid-IR imaging and spectroscopy obtained with the camera/spectrograph Thermal-Region Camera Spectrograph \citep[T-ReCS;][]{Telesco1998} instrument on the Gemini-South telescope. The Herschel images were obtained using the Photodetector Array Camera and Spectrometer \citep[PACS;][]{Poglitsch2010} and the Spectral and Photometric Imaging REceiver \citep[SPIRE; ][]{Griffin2010} instruments. We also use archival Spitzer data taken with the Infrared Array Camera \citep[IRAC;][]{Fazio2004}, the Multi-Band Imaging Photometer for Spitzer \citep[MIPS; ][]{Rieke2004}, and the InfraRed Spectrograph \citep[IRS; ][]{Houck2004} instruments. Using IR observations to study the nuclear and circumnuclear activity in the ILR region of NGC~1365 is crucial because the central region is crossed by a prominent dust lane that obscures from our view in the optical a significant fraction of emission produced there. This paper is organized as follows. We describe the observations in Section~2. We analyze the AGN IR emission and the spatially resolved IR emission in Sections~3 and 4, respectively, and we summarize our conclusions in Section~5. \begin{figure*} \subfloat{\includegraphics[width=0.3\textwidth]{figure2a.ps}} \hspace{0.3cm} \subfloat{\includegraphics[width=0.3\textwidth]{figure2b.ps}} \hspace{0.3cm} \subfloat{\includegraphics[width=0.3\textwidth]{figure2c.ps}} \subfloat{\includegraphics[width=0.3\textwidth]{figure2d.ps}} \hspace{0.3cm} \subfloat{\includegraphics[width=0.3\textwidth]{figure2e.ps}} \hspace{0.3cm} \subfloat{\includegraphics[width=0.3\textwidth]{figure2f.ps}} \subfloat{\includegraphics[width=0.3\textwidth]{figure2g.ps}} \hspace{0.3cm} \subfloat{\includegraphics[width=0.3\textwidth]{figure2h.ps}} \hspace{0.3cm} \subfloat{\includegraphics[width=0.3\textwidth]{figure2i.ps}} \caption[IRcentralimages]{IR view of the approximate extent of the ILR region of NGC~1365. We show the Spitzer/IRAC images at 3.6, 4.5, 5.8, and $8\,\mu$m, the Spitzer/MIPS image at $24\,\mu$m, the Herschel/PACS 70, 100, and $160\,\mu$m images and the Herschel/SPIRE image at $250\,\mu$m. We mark as filled stars the positions of the brightest mid-IR clusters (M2, M3, M4, M5, M6, see text for more details and also Fig.~\ref{fig:closeups}) detected from the ground by \cite{Galliano2005}, the position of the AGN as a filled dot, and the position of the L4 H$\alpha$ hot spot from \cite{Alloin1981} as a filled square (see Table~\ref{tab:sources}). Orientation is north up, east to the left. The images are shown in a square root intensity scale.} \label{fig:IRcentralimages} \end{figure*} \section[]{Observations} \subsection{Herschel/PACS and SPIRE imaging} We obtained Herschel far-IR imaging observations of NGC~1365 using PACS at 70, 100, and $160\,\mu$m and SPIRE at 250, 350, and $500\,\mu$m. The data are part of the guaranteed time program entitled ``Herschel imaging photometry of nearby Seyfert galaxies: testing the coexistence of AGN and starburst activity and the nature of the dusty torus'' (PI: M. S\'anchez-Portal). The PACS observations were carried out using the standard scan map mode that takes two concatenated scan line maps at 45$^{\circ}$ and 135$^{\circ}$ (in array coordinates), at a speed of 20\,arcsec/sec, each one with 26 lines of $9\arcmin$ length and cross-scan step of $20\arcsec$. This mode produces a rather homogeneous exposure map within a square region of about $7\arcmin \times 7\arcmin$. The set of maps were duplicated to observe through both the 70 \micron~(``blue'') and 100 \micron~(``green'') filters. Therefore the galaxy was observed twice through the 160 \micron~(``red'') filter. With the SPIRE photometer we observed the three available bands simultaneously using the ``large map'' mode, with two nearly orthogonal scan maps (2 scan lines each), at a scan speed of 30\,arcsec/sec, and three repetitions. The homogeneous exposure area for scientific use is approximately $8\arcmin \times 8\arcmin$. Table~\ref{tab:her_obs} gives the summary of the observations. \begin{table*} \centering \begin{minipage}{100mm} \caption{Mid and far-IR aperture photometry of the central region of NGC~1365}\label{tab:photometry} \begin{tabular}{lcccccc} \hline Instrument & $\lambda_{\rm c}$ & Pixel size & FWHM & $f_\nu$ (r=15'') & $f_\nu$ (r=30'')\\ & $(\mu{\rm m})$ & (arcsec) & (arcsec) & (Jy) & (Jy)\\ \hline MIPS & 24 & 2.45 & 5.9 & 7.2 & 8.7\\ PACS & 70 & 1.4 & 5.6 & 85.5 & 102.7\\ PACS & 100 & 1.7 & 6.8 & 110.1 & 141.7\\ PACS & 160 & 2.8 & 11.3 & 87.4 & 123.7\\ SPIRE & 250 & 6.0 & 18.1 & 33.4 & 53.3 \\ SPIRE & 350 & 10.0 & 25.2 & 10.4 & 21.0 \\ SPIRE & 500 & 14.0 & 36.9 & -- & 5.6\\ \hline \end{tabular} Note.--- The reported values of the FWHM are the nominal values. The errors in the aperture photometry are dominated by the photometric calibration uncertainty of the instruments that is typically 10\%. \end{minipage} \end{table*} We reduced the data with the Herschel Interactive Processing Environment (HIPE) v8.0.1 and Scanamorphos \citep{Roussel2012} v15. For the PACS instrument, we used HIPE and the Calibration Database V32 to build Level 1 products. These included detecting and flagging bad pixels, converting the ADU readings to flux units (Jy/pixel), and adding the pointing information. We did not attempt to perform deglitching at this stage to prevent the bright AGN nucleus to be affected by the MMT deglitching process. The final maps were built from the Level 1 products using Scanamorphos, which performs a baseline subtraction, correction of the striping effect due to the scan process, removal of global and individual pixel drifts, and finally the map assembly using all the nominal and cross-direction scans. For SPIRE we used the standard (small) map script and Calibration Database v8.1. The processing included detection of thermistor jumps in the time line, frame deglitching, low pass filter correction, conversion of readings to flux units (Jy/beam), temperature drift and bolometer time response corrections, and addition of pointing information. We built the final maps using the ``na\"ive'' scan mapper task. Colour corrections (for PACS, see \citealt{Poglitsch2010}; please refer to the \textit{Observer's Manual} for the SPIRE ones) are small for blackbodies at the expected temperatures (e.g., \citealt{PerezGarcia2001}) and have been neglected. More details on the processing of Herschel data are given in S\'anchez-Portal et al. (in preparation). Figure~\ref{fig:PACS100_large} shows the fully reduced PACS $100\,\mu$m image of NGC~1365 together with the optical BVR image from \cite{Elmegreen2009}. Fig.~\ref{fig:IRcentralimages} shows the PACS images together with the SPIRE $250\,\mu$m with a field of view (FoV) covering the approximate extent of the ILR region of NGC~1365 (see also Section~\ref{sec:alignment}). We performed aperture photometry on all the Herschel images using two different radii, $r=15\arcsec$ and $r=30\arcsec$. The latter encompasses the ILR region, whereas the former includes mostly the ring of star formation. Table~\ref{tab:photometry} lists the measured flux densities. \subsection{Gemini/T-ReCS imaging and spectroscopy}\label{sec:trecs} We obtained mid-IR imaging of NGC~1365 using T-ReCS on the Gemini-South telescope on September 8 and 9, 2011 as part of proposal GS-2011B-Q-20 (PI: N. Levenson). We used the Si-2 ($\lambda_{\rm c}=8.74\,\mu$m) and the Qa ($\lambda_{\rm c}=18.3\,\mu$m) filters and mid-IR standard observation techniques. The plate scale of the T-ReCS imaging observations is 0.089\arcsec/pixel with a FoV of $28.5\arcsec \times 21.4\arcsec$. The total integration times (on-source) were 145\,s and 521\,s in the Si-2 and Qa filters, respectively. The Qa filter observations were split between the two nights, whereas those in the Si-2 filter were done on the second night. We observed standard stars immediately before or after the science observations in the same filters, to both flux-calibrate the galaxy observations and to estimate the unresolved nuclear emission (see below). The observations were diffraction limited, with a full width half maximum (FWHM) of $0.34\arcsec$ in the Si-2 filter and $0.55-0.58\arcsec$ in the Qa filter, as measured from the standard star observations. We reduced the imaging data using the CanariCam data reduction package RedCan (Gonz\'alez-Mart\'{\i}n et al., in preparation). We refer the interested reader to this work and \cite{RamosAlmeida2011AGN} for details. Fig.~\ref{fig:trecs} shows the fully reduced T-ReCS Qa image resulting from the combination of the data taken during the two observing nights. The T-ReCS $8.7\,\mu$m image (not shown here) shows a similar morphology. We also retrieved archival T-ReCS spectroscopic observations in the $N$-band ($\sim 8-13\,\mu$m) using a $0.35\arcsec$ slit width as part of proposal GS-2009B-Q-19 (PI: M. Pastoriza). The total on-source integration time was 600\,s. We reduced the galaxy and corresponding standard star observations using the RedCan package following the steps described in \cite{AAH11AGN}. Finally we extracted the nuclear spectrum as a point source. To estimate the nuclear unresolved emission from the mid-IR imaging data we followed the point spread function (PSF) scaling technique implemented by \cite{RamosAlmeida2009,RamosAlmeida2011AGN}. This unresolved emission which is assumed to represent the torus emission (see Section~\ref{sec:torusfit}). To do so, we scaled the observation of the corresponding standard star to the peak of the nuclear emission of the galaxy in each of the two filters. This represents the maximum contribution of the unresolved source (100\%), whereas the residual of the total emission minus the scaled PSF accounts for the extended emission. In both filters we found that a 90\% PSF scaling provided a realistic estimate of the extended emission. The estimated errors in the T-ReCS unresolved flux densities reported in Table~\ref{tab:AGNfluxes} are 15\% and 25\% in the Si-2 and Qa filters, respectively, and account for both the flux calibration and PSF subtraction uncertainties \citep[see][for more details]{RamosAlmeida2009}. The unresolved $8.7\,\mu$m emission computed this way is in good agreement with the flux density at the same wavelength obtained from the T-ReCS nuclear spectrum. \subsection{Archival Spitzer/IRAC and MIPS imaging} We retrieved from the Spitzer archive imaging data of NGC~1365 observed with all four IRAC bands (3.6, 4.5, 5.8, and $8\,\mu$m) and with MIPS at $24\,\mu$m (Program ID: 3672, PI: J. Mazzarella). These observations were part of The Great Observatories All-Sky LIRG Survey \citep[GOALS, see][]{Armus2009}. We retrieved the basic calibrated data (BCD) from the Spitzer archive. The BCD processing includes corrections for the instrumental response (pixel response linearization, etc.), flagging of cosmic rays and saturated pixels, dark and flat fielding corrections, and flux calibration based on standard stars. We combined the BCD images into mosaics using the MOsaicker and Point source EXtractor (MOPEX) software provided by the Spitzer Science Center using the standard parameters. The final mosaics were repixeled to half of the original pixel size of the images, that is, the IRAC mosaics have $0.6\arcsec$/pixel, whereas the MIPS $24\,\mu$m mosaic has $1.225\arcsec$/pixel. In Fig.~\ref{fig:IRcentralimages} we show the Spitzer images with a FoV covering the approximate extent of the ILR region, as done with the PACS images and the SPIRE $250\,\mu$m image. The angular resolutions of the IRAC images are between 1.7 and $2\arcsec$ (FWHM) and that of the MIPS $24\,\mu$m is $5.9\arcsec$. \subsection{Archival optical ground-based imaging} We retrieved from the ESO archive optical images obtained with the WFI instrument on the MPG/ESO 2.2m telescope using the narrow-band Halpha/7 filter ($\lambda_{\rm c}=6588.3$\AA, FWHM=70\AA) obtained as part of proposal 065.N-0076 (PI: F. Bresolin). We combined a total of 6 images, each of them with a 350\,s exposure. The plate scale of the images is $0.238\arcsec$/pixel. The filter contains the H$\alpha$ and [N\,{\sc ii}] emission lines plus adjacent continuum. Since we use this image for morphological purposes only, we did not attempt to either subtract the continuum or calibrate it photometrically. The positions of the H$\alpha$ hot spots identified by \cite{Alloin1981} in the central region of NGC~1365 are displayed in Fig.~\ref{fig:closeups} (upper panel). \begin{figure} \hspace{0.5cm} \resizebox{0.8\hsize}{!}{\rotatebox[]{0}{\includegraphics{figure3.ps}}} \caption[]{Gemini/T-ReCS image of the central region of NGC~1365 in the Qa ($\lambda_{\rm c}=18.3\,\mu$m) filter. The image has been rotated to the usual orientation of north up, east to the left, and smoothed with a Gaussian function. We mark the positions of the AGN (filled dot) as well as the mid-IR clusters (star symbols) identified by \citet{Galliano2005}.} \label{fig:trecs} \end{figure} \begin{figure} \hspace{0.5cm} \resizebox{0.8\hsize}{!}{\rotatebox[]{0}{\includegraphics{figure4a.ps}}} \hspace{0.5cm} \resizebox{0.8\hsize}{!}{\rotatebox[]{0}{\includegraphics{figure4b.ps}}} \hspace{0.5cm} \resizebox{0.8\hsize}{!}{\rotatebox[]{0}{\includegraphics{figure4c.ps}}} \caption[]{Close-ups of the inner $36\arcsec \times 36\arcsec$ region showing the location of the AGN (filled dot), the mid-IR clusters M2...M8 (filled star symbols) of \cite{Galliano2005}, the radio sources (open circles) of \cite{Sandqvist1995}, and the H$\alpha$ hot spots (filled squares) of \cite{Alloin1981} and \cite{Kristen1997}. The upper panel is the optical ESO Halpha/7 narrow-band image, the middle panel is the IRAC $8\,\mu$m image, and the lower panel the PACS $70\,\mu$m image. The beam of the images is shown on the lower left corner of each image, approximated as the FWHM of a Gaussian function.} \label{fig:closeups} \end{figure} \subsection{Alignment of the images}\label{sec:alignment} The alignment of the Spitzer/IRAC and the Gemini/T-ReCS images is straightforward because the AGN and the mid-IR clusters detected by \cite{Galliano2005} are clearly identified in all these images. The AGN is also bright in the optical image and therefore we used it as our reference. In Fig.~\ref{fig:closeups} we present a close-up of the ILR region in the IRAC $8\,\mu$m band and the optical narrow-band H$\alpha$ image. We marked the positions of the bright (designated as M4, M5, and M6) and faint mid-IR (designated as M2, M3, M7, and M8) clusters and the nucleus using the relative positions given by \cite{Galliano2005}. We also indicated the positions of radio sources and H$\alpha$ sources (see Section~\ref{sec:morphology_detailed}). Although the Spitzer/MIPS $24\,\mu$m image has a poorer angular resolution when compared to that of IRAC, the AGN is still sufficiently bright at $24\,\mu$m that can be distinguished from the mid-IR clusters. The alignment of the Herschel/PACS images is not as simple because the AGN does not appear to be a bright source in the far-IR. We used the astrometry information in the Herschel image headers and the optical position of the nucleus of NGC~1365 given by \cite{Sandqvist1995}: RA(J2000)=$03^{\rm h}33^{\rm m}36.37^{\rm s}$ and Dec(J2000)=$-36^{\circ}08\arcmin25.5\arcsec$, for the initial alignment. These coordinates placed the AGN to the southwest of the bright source detected in the PACS images and in the MIPS $24\,\mu$m image that appears to coincide with the region containing clusters M4, M5, and M6. Additionally we compared the morphologies of the PACS $70\,\mu$m and the MIPS $24\,\mu$m images as they have similar angular resolutions (see Table~\ref{tab:photometry}). We used the mid-IR source located $\sim 17\arcsec$ N, $16\arcsec$ E of the AGN identified in the IRAC and MIPS images and also seen in the PACS $70\,\mu$m and $100\,\mu$m images for a finer alignment. This source appears to be coincident with the L4 H\,{\sc ii} region or H$\alpha$ hot spot (see Fig.~\ref{fig:closeups} and Table~\ref{tab:sources}) identified by \cite{Alloin1981} and also seen in our archival H$\alpha$ image. The PACS $160\,\mu$m image was aligned relative to the other PACS images with the astrometry information in the headers. Finally since the bright mid-IR clusters cannot be resolved in the SPIRE $250\,\mu$m image, we placed the center of image at the position of the AGN. We note that the alignment of the PACS images in Figs.~\ref{fig:IRcentralimages} and \ref{fig:closeups} is only good to within 1 pixel in each band. Since the main goal of this work is to study the processes giving rise to the IR emission within the ILR region of NGC~1365, in Fig.~\ref{fig:IRcentralimages} we show the aligned IR images with a FoV covering the approximate extent of this region. We do not show the SPIRE 350 and $500\,\mu$m images because of the small number of pixels covering the ILR region of NGC~1365. \begin{table} \centering \begin{minipage}{70mm} \caption{AGN emission}\label{tab:AGNfluxes} \begin{tabular}{lcc} \hline Wavelength & $f_\nu$ & Method \\ ($\mu$m) & (mJy)\\ \hline 8.7 & $203\pm30$ & Imaging (unresolved) \\ 13.0 & $400\pm60$ & Spectroscopy \\ 18.3 & $818\pm205$ & Imaging (unresolved) \\ 24 & $1255^{+783}_{-500}$ & BC fit \\ 70 & $734^{+1482}_{-422}$ & BC fit\\ 100 & $271^{+632}_{-163}$ & BC fit\\ 160 & $<78$ & BC fit\\ \hline \end{tabular} References.--- The quoted uncertainties for the BC fit fluxes are the $\pm 1\sigma$ uncertainties, as discussed in Sections~\ref{sec:torusfit} and \ref{sec:farIRAGNemission}. \end{minipage} \end{table} \subsection{Spitzer/IRS spectral mapping} We retrieved from the Spitzer archive low spectral resolution ($R \sim 60-126$) observations (Program ID: 3269, PI: J. Gallimore) of NGC~1365 obtained with the spectral mapping capability of IRS. These observations were part of the Spitzer study of the spectral energy distributions (SED) of the $12\,\mu$m sample of active galaxies \citep{Gallimore2010}. The observations were obtained with the Short-Low (SL1; $7.5-14.3\,\mu$m and SL2; $5.1-7.6\,\mu$m) and Long-Low (LL1; $19.9-39.9\,\mu$m and LL2; $13.9-21.3\,\mu$m) modules. The plate scales of the SL and LL modules are $1.8\arcsec$/pixel and $5.1\arcsec$/pixel, respectively. \begin{figure} \resizebox{0.96\hsize}{!}{\rotatebox[]{-90}{\includegraphics{figure5.ps}}} \vspace{-1.5cm} \caption[]{Spitzer/IRS SL1+SL2 $5-15\,\mu$m spectra of selected regions (see Table~\ref{tab:spectroscopy}) normalized at $13\,\mu$m. The spectrum of the AGN has been shifted up for clarity. The most prominent spectral features are marked.} \label{fig:SLspectra} \end{figure} \begin{table} \centering \caption{Summary of sources in the circumnuclear region}\label{tab:sources} \begin{tabular}{lccccccc} \hline Name & Spectral &Rel. Position & Ref. \\ & Range & arcsec, arcsec & \\ \hline M2, L2 & mid-IR, H$\alpha$ & -4.7, -5.1 & 1, 2\\ M3, L3 & mid-IR, H$\alpha$ & -5.4, -2.6 &1, 2, 3\\ M4, D & mid-IR, radio & 0.4, 7.1 & 1, 4\\ M5, E, L12 & mid-IR, radio, H$\alpha$ & 2.8, 10.0 & 1, 4, 3\\ M6, G & mid-IR, radio & 4.8, 7.0 & 1, 4 \\ M7, H& mid-IR, radio & $^*$ & 1, 4\\ M8, L1 & mid-IR, H$\alpha$ & $^*$ & 1, 2\\ L11 & H$\alpha$ &10, 15 & 3\\ L4 & H$\alpha$ & 17, 16 & 3\\ A & radio &-4.1, -2.4 & 4\\ \hline \end{tabular} The positions are relative to that of the AGN and correspond to those of the first listed source. $^*$The positions shown in Figs.~3 and 4 are estimated from the mid-IR images of Galliano et al. (2005). References.--- 1. Galliano et al. (2005). 2. Kristen et al. (1997). 3. Alloin et al. (1981). 4. Sandqvist et al. (1995). \end{table} The data were processed using the Spitzer IRS pipeline. The IRS data cubes were assembled using {\sc cubism} \citep[the CUbe Builder for IRS Spectra Maps,][]{Smith2007CUBISM} v1.7 from the individual BCD spectral images obtained from the Spitzer archive. {\sc cubism} also provides error data cubes built by standard error propagation, using, for the input uncertainty, the BCD-level uncertainty estimates produced by the IRS pipeline from deviations of the fitted ramp slope fits for each pixel. We used these uncertainties to provide error estimates for the extracted spectra, and the line and continuum maps \citep[see][for full details]{Smith2007CUBISM} discussed in the next two sections. \subsubsection{Extraction of the 1D spectra}\label{sec:1Dspectra} We used {\sc cubism} to extract spectra of regions of interest using small apertures taking advantage of the angular resolution of the spectral mapping observations obtained with the SL1+SL2 modules \citep[$\sim 4\arcsec$ FWHM, see][]{Pereira2010IRSmapping}. We used square or rectangular apertures in the original orientation of the SL data cubes with sizes of two or three pixels (see Table~\ref{tab:spectroscopy} for the extraction apertures used for each region). The selected regions include the AGN, and the regions containing the M2+M3, M4, and M5+M6 mid-IR clusters (see Fig.~\ref{fig:closeups} for the positions) identified by \cite{Galliano2005}. We note that we did not attempt to apply a point source correction to the SL $3.7\arcsec \times 3.7\arcsec$ spectrum of the AGN because we are mostly interested in the extended features, that is, the polycyclic aromatic hydrocarbon (PAH) features and the [Ne\,{\sc ii}]$12.81\,\mu$m fine structure line. Fig.~\ref{fig:SLspectra} shows the SL1+SL2 spectra of the selected regions normalized at $13\,\mu$m. Finally we extracted the integrated spectrum of the region covered by the observations, i.e. the central $27.8\arcsec \times 24.0\arcsec$. We measured the fluxes and the equivalent widths (EW) of the [Ne\,{\sc ii}]$12.81\,\mu$m emission line and the 6.2 and $11.3\,\mu$m PAH features fitting Gaussian profiles to the lines and lines to the local continuum. Our measurements of the [Ne\,{\sc ii}] flux for clusters M4, M5, M6 are in good agreement with those reported by \cite{Galliano2008} from ground-based high angular resolution observations. Since the PAH features are broad, it has been noted in the literature that Gaussian profiles might not be appropriate to measure their flux because a large fraction of the energy in these bands in radiated in the wings. A Lorentzian profile might be a better approximation \citep[see][]{Galliano2008PAH} to measure their flux, and therefore we repeated the line fits with this method. Table~\ref{tab:spectroscopy} lists the measurements for the extracted spectra for the Gaussian profiles. To illustrate the effect of using different profiles, in this table we give the measured 6.2 to $11.3\,\mu$m PAH ratio for the two profiles and the selected regions. \begin{table*} \centering \begin{minipage}{150mm} \caption{Measurements from the Spitzer/IRS SL1+SL2 spectra}\label{tab:spectroscopy} \begin{tabular}{lccccccccc} \hline Region & Extraction & \multicolumn{2}{c}{$6.2\,\mu$m PAH feature} & \multicolumn{2}{c}{$11.3\,\mu$m PAH feature} & [Ne\,{\sc ii}]$12.81\,\mu$m & $S_{\rm Si}$ & \multicolumn{2}{c}{$\frac{{\rm PAH}6.2}{{\rm PAH}11.3}$}\\ & Aperture & flux & EW & flux & EW & flux & & G & L\\ \hline AGN & $3.7\arcsec \times 3.7\arcsec$ & 8.6 & 0.09 & 5.3 & 0.13 & 1.7 & $-0.10$ & 1.6 & 1.2\\ M2+M3 & $5.6\arcsec \times 3.7\arcsec$ & 21.0 & 0.50 & 12.7 & 0.61 & 4.6 & $-0.24$ & 1.7 & 1.1\\ M4 & $3.7\arcsec \times 3.7\arcsec$ &10.7 & 0.45 & 5.8 & 0.63 & 2.3 & $-0.85$ & 1.8 & 1.7\\ M5+M6 & $5.6\arcsec \times 3.7\arcsec$ & 20.0 & 0.45 & 10.2 & 0.40 & 6.5 & $-0.81$ & 2.0 & 1.5\\ Integrated & $27.8\arcsec \times 24.0\arcsec$ & 265.0 & 0.42 & 161.0 & 0.49 & 60.5 & $-0.45$ & 1.6 & 1.3\\ \hline \end{tabular} Notes.--- Fluxes (observed, not corrected for extinction) are in units of $\times 10^{-13}\,{\rm erg \, cm}^{-2}\,{\rm s}^{-1}$ and EW in units of $\mu$m for measurements done with Gaussian profiles. The typical errors are 10\% for the fluxes and $0.05\,\mu$m for the EW. $S_{\rm Si}$ is the strength of the silicate feature (see Section~\ref{sec:spectralmaps} for the definition). The ratio of the $6.2\,\mu$m PAH flux to the $11.3\,\mu$m PAH feature flux is given for fits to the features done with Gaussians (G) and Lorentzian (L) profiles (see Section~\ref{sec:1Dspectra} for more details). \end{minipage} \end{table*} \begin{figure} \resizebox{1.\hsize}{!}{\rotatebox[]{-90}{\includegraphics{figure6.ps}}} \caption[]{Spitzer/IRS observed (not corrected for extinction) spectral maps built with {\sc cubism} from the SL data cubes. The maps of the $6.2\,\mu$m PAH feature (lower left), $11.3\,\mu$m PAH feature (upper right), and [Ne\,{\sc ii}]$12.81\,\mu$m (lower right) are in units of $ 10^{-7}\,{\rm W}\,{\rm m}^{-2}\,{\rm sr}^{-1}$, and the $5.5\,\mu$m continuum map (upper left) in units of ${\rm MJy}\,{\rm sr}^{-1}$. Only pixels with relative errors of less than 20\% are displayed. The position of the AGN (filled dot) was made to coincide with the peak of the $5.5\,\mu$m continuum emission. The positions of the M2...M6 mid-IR clusters (filled star symbols) of \cite{Galliano2005} are also plotted. Orientation is north up, east to the left. The apparently faint $11.3\,\mu$m PAH emission in the region of clusters M4, M5, and M6 is due to foreground extinction (see also the left panel of Fig.~7, and Sections~\ref{sec:morphology_detailed} and \ref{sec:silicatefeature}).} \label{fig:spectralmaps} \end{figure} \subsubsection{Spectral Maps}\label{sec:spectralmaps} We used {\sc cubism} to construct spectral maps of the most prominent features in the SL data cubes, namely, the $6.2\,\mu$m and $11.3\,\mu$m PAH features, and the [Ne\,{\sc ii}]$12.81\,\mu$m fine structure line. The technique used here was very similar to that of \cite{AAH09} and involves integrating the line flux over the user-defined emission line regions. Note that, unlike the line measurements in the previous section, the features are not actually fitted with Gaussian or Lorentzian profiles, and therefore these maps are only used for morphological purposes. Since {\sc cubism} does not fit or deblend emission lines, the SL [Ne\,{\sc ii}]$12.81\,\mu$m map includes some contribution from the $12.7\,\mu$m PAH feature. We also built a continuum map at $5.5\,\mu$m. The Spitzer/IRS SL spectral maps shown in Fig.~\ref{fig:spectralmaps} were trimmed and rotated to the usual orientation of north up, east to the left. We used the associated uncertainty maps produced by {\sc cubism} to compute the relative errors of the spectral maps. Finally we constructed the map of the silicate feature, which is shown in Fig.~\ref{fig:silicatesPAHratio}, following the technique of \cite{Pereira2010IRSmapping}. Briefly, it involves fitting the silicate feature from 1D spectra extracted in $2{\rm pixel} \times 2{\rm pixel}$ boxes from the IRS SL data cubes. The map is then constructed by moving by 1 pixel in the x and y directions until the FoV of the IRS SL data cubes is completely covered. We measured the apparent strength of the silicate feature in the IRS spectra following \cite{Spoon2007}: $S_{\rm Si} = \ln f_\nu({\rm obs})/f_\nu({\rm cont})$, where $f_\nu({\rm obs})$ is the flux density at the feature and $f_\nu({\rm cont})$ is the flux density of the underlying continuum. The latter was fitted as a power law between 5.5 and $14\,\mu$m. We evaluated the strength of the silicate feature at $10\,\mu$m. When the silicate strength is negative, the silicate feature is in absorption, whereas a positive silicate strength indicates that the feature is seen in emission. The uncertainties of the measured strengths in the spectral map and the extracted 1D spectra of previous section are $\pm 0.05$. \begin{figure} \resizebox{1.\hsize}{!}{\rotatebox[]{-90}{\includegraphics{figure7.ps}}} \vspace{-2.5cm} \caption[]{Spitzer/IRS spectral maps of the strength of the silicate feature $S_{\rm Si}$ (left panel) and the observed (not corrected for extinction) ratio of the $6.2\,\mu$m to the $11.3\,\mu$m PAH feature emission (right panel). Negative values of $S_{\rm Si}$ indicate that the feature is observed in absorption. Orientation is north up, east to the left. Symbols are as in Fig.~\ref{fig:spectralmaps}.} \label{fig:silicatesPAHratio} \end{figure}
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In this paper we have studied the IR emission associated with the star formation activity in the ILR region of NGC~1365, as well as the IR emission of the AGN. To this end we have analyzed new far-IR ($70-500\,\mu$m) Herschel/PACS and SPIRE imaging, and high angular resolution ($\sim 0.4\arcsec$) Gemini/T-ReCS mid-IR imaging and spectroscopy of this galaxy. We have also made use of archival Spitzer/IRAC and MIPS imaging and IRS spectral mapping data. Our main findings for the inner $D \sim 5\,$kpc region of NGC~1365 are: \begin{itemize} \item The new Herschel PACS imaging data at 70, 100, and $160\,\mu$m reveal that the ring of star formation in the ILR region is bright in the far-IR. The AGN is the brightest mid-IR source in the inner 2\,kpc up to $\lambda\simeq 24\,\mu$m, but it becomes increasingly fainter in the far-IR when compared with the mid-IR clusters or groups of them in the ring. \item The 24 and $70\,\mu$m emissions as well as the [Ne\,{\sc ii}]$12.81\,\mu$m line and PAH features trace the star-forming ring in the ILR region and have morphologies similar to the CO ``twin-peaks''. This all indicates that there is intense on-going star formation taking place in the inner few kpc of NGC~1365. \item The unresolved near and mid-IR nuclear emission and mid-IR spectrum (i.e., AGN-dominated emission) of NGC~1365 are well reproduced with a relatively compact torus (outer radius of $R_{\rm o}=5^{+0.5}_{-1}\,$pc) with an opening angle of $\sigma_{\rm torus}=36^{+14}_{-6}$deg, and an AGN bolometric luminosity $L_{\rm bol}({\rm AGN})=2.6\pm0.5\times 10^{43}\,{\rm erg \, s}^{-1}$ using the {\sc clumpy} torus models. These parameters are in good agreement with independent estimates in the literature. \item Using the fitted torus model we quantified the AGN emission in the far-IR. The AGN only contributes at most 1\% of the $70\,\mu$m emission within the inner 5.4\,kpc ($r=30\arcsec$), and less than 1\% at longer wavelengths. At $24\,\mu$m the AGN accounts for $\sim 15\%$ of the emission in the same region. We estimated that the AGN bolometric contribution to the $3-1000\,\mu$m luminosity in the inner 5.4\,kpc is approximately 5\%. \item The non-AGN 24 to $500\,\mu$m SED of the ILR region (inner 5.4\,kpc) of NGC~1365 is well fitted with a combination of two modified blackbodies with warm and cold temperatures of 54\,K and 24\,K, respectively. However, the cold dust component accounts for most of total dust mass inferred in this region ($M_{\rm dust}({\rm ILR})= 7\times 10^7\,{\rm M}_\odot$) and has a temperature similar to that of other nuclear and circumnuclear starbursts of similar sizes and IR luminosities. \item From the comparison between the SFR from H$\alpha$ (unobscured) and the SFR from $24\,\mu$m (obscured) we infer that up to $\sim 85\%$ of the on-going SFR inside the ILR region of NGC~1365 is taking place in dust-obscured regions in the ring of star formation. \end{itemize}
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1206.2108
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1206
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1206.1989_arXiv.txt
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Reliable measurements of the solar magnetic field are still restricted to the photosphere, and our present knowledge of the three-dimensional coronal magnetic field is largely based on extrapolation from photospheric magnetogram using physical models, e.g., the nonlinear force-free field (NLFFF) model as usually adopted. Most of the currently available NLFFF codes have been developed with computational volume like Cartesian box or spherical wedge while a global full-sphere extrapolation is still under developing. A high-performance global extrapolation code is in particular urgently needed considering that Solar Dynamics Observatory (SDO) can provide full-disk magnetogram with resolution up to $4096\times 4096$. In this work, we present a new parallelized code for global NLFFF extrapolation with the photosphere magnetogram as input. The method is based on magnetohydrodynamics relaxation approach, the CESE-MHD numerical scheme and a Yin-Yang spherical grid that is used to overcome the polar problems of the standard spherical grid. The code is validated by two full-sphere force-free solutions from Low \& Lou's semi-analytic force-free field model. The code shows high accuracy and fast convergence, and can be ready for future practical application if combined with an adaptive mesh refinement technique.
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1206.1989
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1206
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1206.4048_arXiv.txt
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{ Breakthrough direct detections of planetary companions orbiting A-type stars confirm the existence of massive planets at relatively large separations, but dedicated surveys are required to estimate the frequency of similar planetary systems. To measure the first estimation of the giant exoplanetary systems frequency at large orbital separation around A-stars, we have conducted a deep-imaging survey of young (8--400~Myr), nearby (19--84~pc) A- and F-stars to search for substellar companions in the $\sim$10--300~AU range. The sample of 42 stars combines all A-stars observed in previous AO planet search surveys reported in the literature with new AO observations from VLT/NaCo and Gemini/NIRI. It represents an initial subset of the International Deep Planet Survey (IDPS) sample of stars covering M- to B-stars. The data were obtained with diffraction-limited observations in $H$- and $K_{\mathrm{s}}$-band combined with angular differential imaging to suppress the speckle noise of the central stars, resulting in typical 5$\sigma$ detection limits in magnitude difference of 12~mag at 1\as, 14~mag at 2\as and 16~mag at 5\as which is sufficient to detect massive planets. A detailed statistical analysis of the survey results is performed using Monte Carlo simulations. Considering the planet detections, we estimate the fraction of A-stars having at least one massive planet (3--14~\MJup) in the range 5--320~AU to be inside 5.9--18.8\% at 68\% confidence, assuming a flat distribution for the mass of the planets. By comparison, the brown dwarf (15--75~\MJup) frequency for the sample is 2.0--8.9\% at 68\% confidence in the range 5--320~AU. Assuming power law distributions for the mass and semimajor axis of the planet population, the AO data are consistent with a declining number of massive planets with increasing orbital radius which is distinct from the rising slope inferred from radial velocity (RV) surveys around evolved A-stars and suggests that the peak of the massive planet population around A-stars may occur at separations between the ranges probed by existing RV and AO observations. Finally, we report the discovery of three new close M-star companions to HIP~104365 and HIP~42334. }
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\label{sec:introduction} An extensive population of exoplanets has been discovered down to sub-Jovian masses and at orbital separations below 5~astronomical units (AU) based on large-scale radial velocity (RV) surveys \citep{mayor2008,marcy2008} and on transit surveys, which are now identifying hundreds of new candidates \citep{cabrera2009,borucki2011}. These indirectly detected planets provide invaluable information on the distribution of close orbit planets \citep{cumming2008} and on their frequency around nearby stars covering a large range of masses \citep{marcy2008,johnson2010b}. Planets at orbital radii larger than 5~AU, comparable to the locations of the giant planets in the Solar System, remain outside the range of detection of these methods, and, consequently, little is known about the extrasolar population of wide orbit planetary systems. Direct detection with high-contrast imaging provides a method to explore the wide orbit planet population. Several adaptive optics (AO) surveys, concentrating on Solar-type stars, have been conducted to search for low-mass substellar companions around nearby young stars. Because direct imaging needs to overcome the large contrast ratio between the star and a potential substellar companion, the existing surveys were focused on FGKM stars, around which young objects down to a few masses of Jupiter (\MJup) would be detectable at separation larger than a few tens of AUs. The majority of these surveys reported no \citep{masciadri2005,biller2007,lafreniere2007b,kasper2007,leconte2010,janson2011,delorme2012} or few \citep{chauvin2010} substellar companions to nearby stars and have placed upper limits on the population of massive planets at large orbital separation. Some objects were nonetheless discovered \citep[e.g.][]{lafreniere2008,thalmann2009,biller2010}, demonstrating that these objects exist, but are indeed rare \citep{nielsen2010}. Recent breakthrough detections around young A-stars -- HR~8799 \citep{marois2008a,marois2010}, $\beta$~Pictoris \citep{lagrange2009a,lagrange2010} and Fomalhaut \citep{kalas2008,janson2012} -- have provided new perspectives on the search for companions orbiting more massive stars. The results of RV surveys of the evolved counterparts of A-stars \citep{johnson2007} have also provided intriguing results \citep{johnson2010a,bowler2010a,johnson2011}, with significant differences in the planet population compared to RV-detected exoplanets around lower mass stars. Statistical analysis of the results clearly shows a higher frequency of planets around more massive stars and larger masses for the detected planets \citep{johnson2010b}. Moreover, theoretical work on planet formation by core-accretion similarly shows the same kind of correlation between the star and planet mass \citep{kennedy2008,alibert2011}. Although more technically challenging in terms of observations, these recent results suggest that more massive stars may present more favorable targets in terms of exoplanet detections. And estimations of planet yield show that A-stars will be even more favorable for future large direct imaging surveys \citep{crepp2011} with dedicated upcoming instruments, such as Gemini/GPI \citep{macintosh2008} and VLT/SPHERE \citep{beuzit2008}. To pursue the possibility of a higher frequency of massive planets around early-type stars, we have constructed a sample of 42 young A--F stars observed at high contrast to be sensitive to massive planets. We report new measurements for 39 stars, and we include 3 A-stars from the literature that have been observed in previous surveys. This survey intends to start defining the population statistics of massive planets and brown dwarfs (BDs) at orbital separations in the tens of AUs from their parent A-stars. This sample is an initial subset of the International Deep Planet Survey (IDPS) spanning M- to B-stars, the results of which will be presented in a forthcoming publication (Galicher et al. 2012, in preparation). The selection of the target sample is explained in Sect.~\ref{sec:sample_selection}. In Sect.~\ref{sec:observations_data_reduction} the observing strategy, observations and data reduction steps are detailed. The detection limits of the survey are derived in Sect.~\ref{sec:results}, and we describe the identification of the candidate companions detected in the data. In Sect.~\ref{sec:statistical_analysis} we use the detection limits to perform a statistical analysis of the survey, from which we derive a first estimation of the planetary systems frequency at large orbital radii around A-stars. Finally, we discuss our results and conclusions in Sect.~\ref{sec:discussion_conclusions}.
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\label{sec:discussion_conclusions} \begin{figure} \centering \includegraphics[width=0.5\textwidth]{fig10} \caption{Comparison of the mean probability of detection for the present survey (plain orange contours) with a simulated survey with GPI of the same 42 targets (dashed blue contours). The mean probability of detection is obtained using Monte-Carlo simulations as described in Sect.~\ref{sec:mean_probability_detection}. In addition to $\beta$~Pic~b (orange square) and the HR~8799 planets (red circles), the giant planets detected around old A stars by \citet{johnson2010a,johnson2011} have been overplotted (purple triangles). The dotted purple line shows the median detection threshold of the \citet{bowler2010a} RV survey around old A-type stars. The HR~8799 planets are represented at their \emph{projected} physical separation because the true physical separation is not yet known precisely. Note that Fomalhaut b has not been included due to its uncertain nature \citep[see e.g.][]{janson2012}.} \label{fig:future_surveys} \end{figure} \begin{table*} \caption{Samples and methods comparison.} \label{tab:samples_methods_comparison} \centering \begin{tabular}{lcccccc} \hline\hline Sample & Host mass & Tech. & Frequency & Sep. range & Planet mass limit & Reference \\ & (\MSun) & & & (AU) & (\MJup) & \\ \hline F5--A0 & $\sim$1.5--3.0 & AO & 5.9--18.8\% & 5--320 & 3--14\tablefootmark{a} & this work \\ Evolved A & 1.3--1.9 & RV & $11 \pm 2$\% & 0.1--3 & $>$0.2--1.3\tablefootmark{b} & \citet{johnson2010b} \\ \hline K7--F2 & 0.7--1.5 & AO & $<20$\% & 25--856 & $>$4 & \citet{nielsen2010} \\ FGK & 0.7--1.3 & RV & $6.5 \pm 0.7$\% & 0.01--3 & $>$0.5--0.9\tablefootmark{c} & \citet{johnson2010b} \\ \hline M5--M0 & 0.2--0.6 & AO & $<20$\% & 9--207 & $>$4 & \citet{nielsen2010} \\ M & 0.1--0.7 & RV & $2.5 \pm 0.9$\% & 0.01--3 & $>$0.1--0.5\tablefootmark{c} & \citet{johnson2010b} \\ \hline \end{tabular} \tablefoot{ \tablefoottext{a}{The detection limit is a function of separation, as detailed in Fig.~\ref{fig:mean_proba_detection}. The frequency estimation is based on a flat distribution for the planets mass.} \tablefoottext{b}{The detection limit is a function of separation, as reported in \citet{bowler2010a}.} \tablefoottext{c}{Scaled from the limits in \citet{bowler2010a} using the fixed RV amplitude cutoff of $K > 20$~m~s$^{-1}$ reported in \citet{johnson2010b} and the ratio of the average mass of the bin to the average mass of the evolved A-stars, taken to the power of 2/3.}} \end{table*} The present survey results on A-star wide orbit planetary systems frequency and population distribution can be compared and combined with previous results to investigate the effects of host star mass, the predictions of theoretical models, the impact of debris disks, and the current state of the planet population across a range of orbital separations. Table~\ref{tab:samples_methods_comparison} summarizes the frequencies and survey sensitivities reported in both AO (current work and \citealt{nielsen2010}, which combines the results of \citealt{masciadri2005}, \citealt{biller2007}, and \citealt{lafreniere2007b}) and RV surveys \citep{johnson2010a}. These studies were selected for comparison, since each covers a large sample and considers the same 68\% confidence interval. For the sample with the most similar target mass range, the evolved A-stars, the AO measurement of the wide orbit planetary systems frequency range of 5.9--18.8\% encompasses the close orbit planet frequency of $11 \pm 2$\%. Direct imaging searches covering similar orbital separations to this study, but targeting lower mass stars \citep{nielsen2010} find an upper limit of $<$20\%, while RV searches for close orbit planets measure a declining frequency for the successively lower mass samples \citep{johnson2010b}. From the current statistics on the wide orbit imaged planet population, it is not possible to determine how the planetary systems frequency scales with host star mass. Theoretical models have predicted a rising \citep{kennedy2008}, peaked \citep{ida2005}, or declining \citep{kornet2006} planet frequency on mass, with models incorporating different treatments of factors such as the location and evolution of the snow line, the initial disk size and its dependence on host star mass, and the orbital migration of planets. Since the survey was designed to detect planets, the full sample is extremely sensitive to brighter brown dwarf companions in wide orbits (probability of detection $>90$\% in $\sim$75--300 AU, see Fig.~\ref{fig:mean_proba_detection}), enabling an initial assessment of the brown dwarf systems frequency around stars more massive than the Sun. Among the 42 targets, one brown dwarf was previously identified around HR~7329 \citep{lowrance2000,neuhauser2011}, yielding a brown dwarf systems frequency of $2.8_{-0.9}^{+6.0}$\% (see Sect.~\ref{sec:estimation_planet_frequency}). This low level of brown dwarf companions is consistent with the $3.2^{3.1}_{-2.7}$\% frequency measured for a larger scale survey of young (3--3000~Myr), solar-type (F5--K5) stars covering a wider range of separations (28--1590~AU) with an AO survey \citep{metchev2009}. Theoretical models of the formation of giant planets by gravitational instability have proposed that the directly imaged planets may be the lowest mass products of the same process that forms binary stars and brown dwarfs \citep{kratter2010}, and these simulations predict a higher proportion of brown dwarfs than planets. While the statistics of the current study are limited, the similarity of the brown dwarf and planetary systems frequencies suggests that the imaged planets may have formed through a different process such as core accretion \citep[e.g.][]{pollack1996}. An investigation into the conditions for gravitational instability for the specific case of the HR~8799 planets \citep{nero2009} found that the likelihood of this formation mechanism increased with radius towards the outermost planets. Other possible explanations for the wide orbits of some of the directly imaged planets include scattering to larger orbital radii from interactions in multiple planet systems \citep[e.g.][]{veras2004}, outward migration in a resonance with multiple planets \citep[e.g.][]{crida2009}, and outward migration through a planet-disk interaction \citep[e.g.][]{veras2004}. The full sample of 42 young, A- and F-stars includes 17 systems with dusty debris disks sustained by the ongoing collisional grinding of planetesimals into smaller particles \citep{backman1991} or an event such as a catastrophic collision of planets \citep{cameron1997,melis2010}. The 3 targets in this sample with imaged planets or brown dwarfs all reside in systems encircled by dust disks. Additionally, Fomalhaut (not included in this study) would provide another example of a planet-disk system \citep{kalas2008} if the presence of a planet was confirmed \citep[see e.g.][]{janson2012}. While not all A-stars with debris disks harbor massive giant planets in wide orbits, the frequency of planetary systems appears higher among the targets with excess emission from dust. The combination of resolved disk maps and planet imaging is an especially powerful tool to understand the planet-disk interactions that may sculpt young planetary systems \citep[e.g.][]{liou1999,kuchner2003,wyatt2006,quillen2006}. The structure of the Fomalhaut disk inner edge has been compared with dynamical models of planets with different masses \citep{chiang2009}, and the size and shape of the HR~8799 disk have been compared with models of the dynamically cleared zones and orbital migration history \citep{patience2011}. Advances in disk imaging with ALMA and planet imaging with extreme AO will enable more detailed studies to investigate planet-disk interactions which may generate structures such as asymmetries and clumps. By combining the results and sensitivities of existing direct imaging A-star planet searches with those of RV programs targeting retired A-stars, a comprehensive summary of the currently known A-star planet population is given in Fig.~\ref{fig:future_surveys}. Analysis of the current data already shows distinct differences. The close orbit A-star planets are best fit by a distribution that is flat or rising with increasing orbital separation \citep{bowler2010a} and rising with increasing planet mass \citep{bowler2010a,johnson2010b}. In contrast, the wide orbit planet data, incorporating the outer cutoff implied by the outermost HR~8799 planet, are consistent with a distribution that is declining with increasing orbital separation. This result is independent of whether the distribution of planets is rising, flat, or declining with planet mass, as shown in the three panels of Fig.~\ref{fig:powerlaws} (the allowed region always has a negative power law index). Thus, the existing surveys have identified the boundaries of planets around A-stars -- from $\sim$0.6~AU \citep{bowler2010a} to $\sim$70~AU -- and have indicated that the data cannot be fit by a single distribution of properties. Figure~\ref{fig:future_surveys} also presents the sensitivity of a simulated survey with the upcoming high-contrast imager Gemini/GPI \citep{macintosh2008}. Using simulated contrast curves for that instrument, we performed Monte-Carlo simulations similar to the ones described in Sect.~\ref{sec:mean_probability_detection} on the 42 targets of our sample. They show that extreme AO imaging will provide the crucial coverage to connect the close orbit and wide orbit populations and to reveal the full distribution of planets as a function of separation and planet mass. Extreme AO data will determine if the transition from a flat or rising population of close orbit planets to a declining population of wide orbit planets is smooth or discontinuous and may discover a peak in the separation distribution. The total number of planets detected in extreme AO surveys, combined with trends measured for close orbit A-star planets, can be used as a test of formation models \citep[e.g.][]{crepp2011}. Since the upcoming extreme AO instrument include integral field units, it will be possible to investigate the nature of the exoplanet atmospheres in addition to the population statistics. Spectra of the currently imaged planetary mass companions have revealed differences with brown dwarfs of similar temperatures \citep[e.g.][]{janson2010,patience2010,barman2011a,barman2011b,skemer2011}, and upcoming AO observations will further explore the architectures and atmospheres of exoplanets around A-stars.
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1206.4048
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1206
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1206.6887.txt
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We present new, full-orbit observations of the infrared phase variations of the canonical hot Jupiter HD 189733b obtained in the 3.6 and 4.5~\micron~bands using the \emph{Spitzer Space Telescope}. When combined with previous phase curve observations at 8.0 and 24~\micron, these data allow us to characterize the exoplanet's emission spectrum as a function of planetary longitude and to search for local variations in its vertical thermal profile and atmospheric composition. We utilize an improved method for removing the effects of intrapixel sensitivity variations and robustly extracting phase curve signals from these data, and we calculate our best-fit parameters and uncertainties using a wavelet-based Markov Chain Monte Carlo analysis that accounts for the presence of time-correlated noise in our data. We measure a phase curve amplitude of $0.1242\%\pm0.0061\%$ in the 3.6~\micron~band and $0.0982\%\pm0.0089\%$ in the 4.5~\micron~band, corresponding to brightness temperature contrasts of $503\pm21$~K and $264\pm24$~K, respectively. We find that the times of minimum and maximum flux occur several hours earlier than predicted for an atmosphere in radiative equilibrium, consistent with the eastward advection of gas by an equatorial super-rotating jet. The locations of the flux minima in our new data differ from our previous observations at 8~\micron, and we present new evidence indicating that the flux minimum observed in the 8~\micron~is likely caused by an over-shooting effect in the 8~\micron~array. We obtain improved estimates for HD~189733b's dayside planet-star flux ratio of $0.1466\%\pm0.0040\%$ in the 3.6~\micron~band and $0.1787\%\pm0.0038\%$ in the 4.5~\micron~band, corresponding to brightness temperatures of $1328\pm11$~K and $1192\pm9$~K, respectively; these are the most accurate secondary eclipse depths obtained to date for an extrasolar planet. We compare our new dayside and nightside spectra for HD~189733b to the predictions of 1D radiative transfer models from \citet{burrows08}, and conclude that fits to this planet's dayside spectrum provide a reasonably accurate estimate of the amount of energy transported to the night side. Our 3.6 and 4.5~\micron~phase curves are generally in good agreement with the predictions of general circulation models for this planet from \citet{showman09}, although we require either excess drag or slower rotation rates in order to match the locations of the measured maxima and minima in the 4.5, 8.0, and 24~\micron~bands. We find that HD~189733b's 4.5~\micron~nightside flux is $3.3\sigma$ smaller than predicted by these models, which assume that the chemistry is in local thermal equilibrium. We conclude that this discrepancy is best-explained by vertical mixing, which should lead to an excess of CO and correspondingly enhanced 4.5~\micron~absorption in this region. This result is consistent with our constraints on the planet's transmission spectrum, which also suggest excess absorption in the 4.5~\micron~band at the day-night terminator.
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Observations of eclipsing extrasolar planetary systems, in which the planet periodically passes in front of and then behind its host star, have proven to be a powerful diagnostic tool for studies of exoplanetary atmospheres. Because the probability of a transiting geometry scales as $R_{\star}/a$ where $R_{\star}$ is the radius of the host star and $a$ is the planet's semi-major axis, the majority of currently known transiting planet systems have orbital periods of just a few days. At these distances, the time scale for the planet to achieve synchronous rotation is short compared to the typical ages of the systems, leading to the prediction that a majority of these transiting planets should be tidally locked \citep{bodenheimer01}. In this paper we focus on the class of gas giant planets known as ``hot Jupiters", which typically have orbital periods on the order of $1-3$ days and atmospheric temperatures ranging between $1000-3000$~K. %These objects represent ideal test cases for exoplanet atmosphere studies, as their relatively large radii and high temperatures allow for precise characterization using existing ground- and space-based facilities. One fundamental question for these planets is what fraction of the incident flux absorbed on the planet's day side is subsequently transported around to the night side. Atmospheric circulation models predict that these planets should develop a broad superrotating (eastward) equatorial jet that circulates energy between the day and night sides \citep[e.g.,][]{showman02,showman09,showman11,langton08,dobbs10,heng11,rauscher11}. Depending on the relative strengths of these winds, these planets could exhibit large gradients in both temperature and composition between the two hemispheres. We can constrain the efficiency of the day-night circulation by measuring changes in the infrared brightness of the planet as a function of orbital phase; the day-night brightness contrast can then be translated into a day-night temperature contrast. There are currently well-characterized phase curves published for seven planets, including $\upsilon$~And b \citep{harrington06,crossfield10}, HD~189733b \citep{knutson07,knutson09a}, HD 149026b \citep{knutson09c}, HD 80606b \citet{laughlin09}, HAT-P-7b \citep{borucki09,welsh10}, CoRoT-1b \citep{snellen09}, and WASP-12b \citep{cowan12}, with more sparsely sampled phase curves for three additional planets (51 Peg b, HD 209458b, and HD 179949b) from \citet{cowan07}. These data indicate that hot Jupiters display a diversity of circulation patterns, ranging from relatively small day-night temperature gradients (e.g., HD 189733b) to large temperature gradients (e.g., WASP-12b). Of these four systems, HD~189733b stands out both as having the best-characterized phase variation, and also as the only system with phase curve observations at more than one wavelength. We know more about this planet's atmosphere than that of any other extrasolar planet; key results include the detection of a high-altitude haze \citep{pont08,sing11} and sodium absorption \citep{redfield08,huitson12} in its visible-light transmission spectrum, as well as more controversial detections of methane, carbon monoxide, and water absorption in its infrared transmission spectrum \citep{swain08,sing09,desert09,gibson11a,gibson11} and carbon dioxide absorption in its dayside emission spectrum \citep{swain09}. Several recent ground-based studies \citep{swain10,waldmann12} have also reported detections of methane emission from the planet's day side, although these results have been the subject of some debate \citep{mandell11}. HD 189733b's dayside emission spectrum has also been characterized in the near- and mid-infrared using both \emph{Spitzer} IRAC photometry \citep{charbonneau08} and IRS spectroscopy \citep{grillmair08}, with additional constraints on its variability in the 8~\micron~IRAC band from \citet{agol10}. Despite the extent of the data available for this planet, there are still a number of open questions regarding the properties of its atmosphere. The single largest outstanding question centers on the issue of whether or not the chemistry is in equilibrium \citep[e.g.,][]{moses11,visscher11}; it is likely that the atmospheric circulation plays an important role in shaping this chemistry \citep[e.g.,][]{cooper06}. Although we have observational constraints on relative abundances for the planet's day side and the day-night terminator, we know very little about the properties of HD~189733b's night side. In this paper we present new full-orbit phase curve observations for HD~189733b in the 3.6 and 4.5~\micron~bands obtained with the \emph{Spitzer Space Telescope} during its extended warm mission. We combine these data with previous observations at 8.0 and 24~\micron~to provide the first detailed characterization of its emission spectrum as a function of orbital phase. Our data include two secondary eclipses and one transit in each band, which we use to derive improved estimates for the planet's orbital ephemeris and dayside emission spectrum.
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In this paper we present new full-orbit and near-continuous phase curve observations of the hot Jupiter HD~189733b in the 3.6 and 4.5~\micron~\emph{Spitzer} bands, which allow us to characterize the atmospheric circulation patterns and corresponding chemistry of HD~189733b's atmosphere. These data include one transit and two secondary eclipses in each bandpass, which we use to derive an improved estimate for the planet's orbital ephemeris and wavelength-dependent transmission and emission spectrum in these bands. We confirm that the planet's 4.5~\micron~transit depth is $3\sigma$ smaller than at 3.6~\micron, consistent with the presence of excess CO at the day-night terminator, although our precision is comparable to that reported by \citet{desert09}. We obtain improved estimates for HD~189733b's dayside planet-star flux ratio of $0.1466\%\pm0.0040\%$ in the 3.6~\micron~band and $0.1787\%\pm0.0038\%$ in the 4.5~\micron~band, corresponding to brightness temperatures of $1328\pm11$~K and $1192\pm9$~K, respectively; these are the most accurate secondary eclipse depths obtained to date for an extrasolar planet. Our new 3.6~\micron~secondary eclipse depth is $7.5\sigma$ smaller than the value reported in \citet{charbonneau08}, but we find that the uncertainties in this previous measurement, which assumed Gaussian noise, are likely underestimated. We conclude that there is no convincing evidence for time variability in the measured secondary eclipse depths or times, consistent with the upper limits derived by \citet{agol10} using a more extensive 8~\micron~data set. We combine our new 3.6 and 4.5~\micron~phase curves with previously published observations at 8.0 and 24~\micron~in order to characterize the planet's emission spectrum as a function of orbital phase. We find that the times of minimum and maximum flux occur several hours earlier than predicted for an atmosphere in radiative equilibrium, consistent with the eastward advection of gas by an equatorial super-rotating jet. The locations of the flux minima in our new data differ from our previous observations at 8~\micron, and we present new evidence indicating that the flux minimum observed in the 8~\micron~is likely caused by an over-shooting effect in the 8~\micron~array. We fit the planet's dayside spectrum with a 1D radiative transfer model from \citet{burrows08} where the amount of energy transported to the night side is left as a free parameter and find that the corresponding nightside spectrum is in good agreement with our measured values. This serves to validate studies \citep[e.g.,][]{cowan11} that seek to constrain circulation patterns on hot Jupiters based on dayside spectra alone. We then compare our phase curves to the predictions of 3D general circulation models from \citet{showman09} and find that we require models with either a slower-than-synchronous rotation rate or increased drag at the bottom of the atmosphere in order to match the small measured offsets in the locations of the phase curve maxima and minima at 4.5, 8.0, and 24~\micron. We also find that HD~189733b's 4.5~\micron~nightside flux is $3.2\sigma$ smaller than predicted by these models, which assume that the chemistry is in local thermal equilibrium. We conclude that this discrepancy is best-explained by vertical mixing, which should lead to an excess of CO and correspondingly enhanced 4.5~\micron~absorption in this region. This result is consistent with our constraints on the planet's transmission spectrum, which also suggest excess absorption in the 4.5~\micron~band at the day-night terminator. Looking ahead, it is clear that the questions regarding atmospheric chemistry and circulation patterns on tidally locked planets will continue to recur as the field shifts towards studies of smaller and more earth-like worlds. Current studies of the atmospheres of low-mass planets such as GJ~1214b \citep{bean10,bean11,berta11,desert11b} and GJ~436b \citep[e.g.,][]{stevenson10,beaulieu11,knutson11} focus almost exclusively on systems with M star primaries, as the lower stellar effective temperature and smaller stellar radius result in proportionally larger transit and secondary eclipse depths. For late M stars, the location of the habitable zone is within the region in which we would expect tidal locking to occur \citep{kasting93}; it is therefore likely that the first atmosphere studies of potentially habitable worlds with the \emph{James Webb Space Telescope} will focus on these tidally locked systems.
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1206.6798_arXiv.txt
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Variable NaI absorption lines have been reported in a number of type Ia supernovae (SNeIa). The presence of this circumstellar material suggests that cataclysmic variables (CVs) with a giant donor star may be the progenitors of these SNeIa (Patat et al. 2007). We present echelle spectra of the CV QU Carinae which strengthen the connection between CVs of the V Sge class, the Accretion Wind Evolution scenario, variable wind features, variable NaI absorption, and SNIa. This thread not only provides insight into the spectral peculiarities of QU Car, but also links SNeIa as a class with their parent systems.
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Although the progenitors of Supernovae Ia (SNeIa) are still controversial, it is accepted that, at least some of the events, originate from detached or semi-detached binary star systems in which at least one of the two components is a massive C-O white dwarf (WD). In semi-detached systems, either via Roche lobe overflow of the companion or via a wind, the WD accumulates H or He-rich material which is then burned to C and O. Under the right conditions (which are primarily controlled by the stability of the mass-accretion rate onto the WD) the net WD mass reaches the Chandrasekhar mass limit initiating a series of thermonuclear reactions eventually leading to a SNeIa. Although outlining this scenario seems to be rather straightforward, its specifics are far from being understood or defined, especially the nature of the mass-losing star. Among the most promising SNeIa progenitors are the V Sagittae (V Sge)-type cataclysmic variables (CVs) \citep{1998PASP..110..276S}. These are semi-detached binaries consisting of a WD and a giant or main sequence donor star, with high mass transfer rates ($\sim$10$^{-7}$-10$^{-5}$M$_{\sun}$/yr). This rate allows for stable nuclear burning on the surface of the WD, which can take place when a relatively massive (0.7-1.2 M$_\sun$; van den Heuvel et al. 1992) WD accretes near Eddington rates ($\sim$10$^{-6} M_{\sun}$/ yr). The process of a WD accumulating mass in a V~Sge-type system, gradually reaching the critical Chadrasekhar mass, is described in detail by the accretion wind evolution scenario (AWE), coined by Hachisu \& Kato (2003a and references therein). AWE successfully reproduces the long-term light curve of the prototype of the category, V~Sge and its Large Magellanic Cloud twin RX~J0513.9-6951 (Hachisu \& Kato 2003b) accounting for the bright and faint states of the systems and the transitions between them. When the accumulated envelope on the WD reaches a critical mass, the WD atmosphere expands, generating a massive wind with $\dot{M}_{wind}\sim$10$^{-7}$M$_{\sun}$yr$^{-1}$. This wind drives a Kelvin-Helmholtz instability at its interface with the disk, peeling off the surface layer of the disk, which completely obscures soft x-rays from escaping from the binary. With time (and accumulated material) $\dot{M}_{wind}$ increases reaching $\sim$10$^{-5}$M$_{\sun}$yr$^{-1}$, in which case the mass of the WD envelope is further reduced by wind mass losses. The chromosphere of the companion is then eroded by the wind to the point where the donor star shrinks inside its Roche lobe and accretion (through L1) temporarily stops. The wind mass loss is then decreased, eventually reaching zero, and the system transitions to a faint state. During this faint state, soft x-rays emerge. The system gradually recovers to the bright state when the donor star regains contact with its Roche lobe and accretion commences; the cycle starts again. Other than the prototype, there are 4 suspected members in this category \citep{1998PASP..110..276S}, but other than V Sge itself, their properties are not well-known. Recently, the Galactic CV QU Carinae (hereafter QU Car) was suggested to also be a V Sge star (Kafka et al. 2008--hereafter Paper~I). The system displays strong and pronounced outflows in the UV and in optical Carbon lines, reaching velocities of at least 5700km s$^{-1}$ (Drew et al. 2003 and Paper~I). At the same time, strong [OIII]~5007 emisson confirms the presence of a circumbinary nebula (Paper~I). QU Car, being the brightest member in the V~Sge category, may be a galactic twin of supersoft X-ray sources. Here, we present new high resolution spectra of QU Car revealing the highly variable nature of the system on timescales of years, and we discuss its SNeIa progenitor status.
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In V Sge, the optical photometric bright state lasts for about 170 days and the faint state for about 130 days (Robertson et al. 1997; \citep{1999A&AS..139...75S}. Figure~1 presents part of the long-term light curve of QU Car: visual data are from the database of the American Association of Variable Star Observers (AAVSO)\footnote{American Association of Variable Star Observers (AAVSO), Henden, A.A., 2009, Observations from the AAVSO International Database, private communication. http://www.aavso.org} and the V magnitude data are from the All Sky Automated Survey (ASAS)\footnote{All Sky Automated Survey (ASAS) http://www.astrouw.edu.pl/asas, Pojmanski (1997).}. The full (more than 20 years) AAVSO visual light curve is also presented and discussed in Paper~I. The ASAS V-mag light curve has a detection limit of V$\sim$14 mag. Sporadic coverage does not allow for a detailed study of the short-term behavior of the system, however the light curves do indicate the occasional presence of ``faint'' states, where the brightness drops below visual 12 mag (also see Paper~I). Albeit short, those faint states can last for $\sim$100~days before the system recovers to its usual V ``bright'' magnitude of $\sim$11.5. The inset in figure~\ref{lc} is a closeup of the light curves at the times of our 2006/2007 observations, also presented in Paper~I, where the system was in its bright V$\sim$11.5 magnitude, exhibiting low-amplitude erratic variations which are normal for nova-like CVs. No light curves exist for the times of the echelle data of 2010/2011. Therefore, we will use the spectral appearance of QU Car from (Paper~I) as a guide on the behavior of the system in its bright state and as a comparison for the new spectroscopic data. Overall, the characteristics of the bright states of QU Car have been explored spectroscopically in \citep{1982ApJ...261..617G}, in \citep{2003MNRAS.338..401D}, and in Paper~I. As discussed in Paper~I, in the bright state the system exhibits variable and irregular Balmer emission lines, HeI and HeII emission and one of the strongest Bowen blend (CIII/NIII/OII at 4640-4650$\AA$) among all CVs. The broad CIV 5807 $\AA$ emission is accompanied by pronounced P-Cygni profiles reaching velocities of $\sim$5700 km/sec. Excess of CNO-processed material in the optical spectra likely originates from a wind expelled from the binary's C-O WD, in agreement with the AWE scenario and with relevant UV observations \citep{2003MNRAS.338..401D}. It is critical to stress that such a wind, carrying away the outer layers of the disk and of the donor star, is necessary in the AWE scenario in order for the WD to eventually reach the Chandrasekhar mass; otherwise the WD envelope expands and forms a circumbinary envelope, eventually leading to merging of the two stellar components instead of a SNeIa \citep{2008ApJ...679.1390H}. The presence of [OIII] 5007$\AA$ and [NII] 6584$\AA$ at all four epochs of the bright state in Paper~I, is the imprint of a nebula in the vicinity of the binary and a relic of material expelled from it. In following sections we will discuss the time variability of the QU Car spectrum in our high resolution data. Gilliland \& Phillips (1982) provided an orbital period of 10.9 hr for QU Car, however this period was not confirmed in Paper I and there is not yet a reliable ephemeris for the binary. The spectral variability discussed in this paper may be due to the varying viewing angle of the different line-forming regions as the system revolves. On the other hand the changes may be due to variations in the strength and motion of the wind. We argue for the latter interpretation, but we admit that the lack of an ephemeris precludes our testing for orbital dependences. Since we don't have flux-calibrated echelle spectra in hand, we use IRAF's {\it continuum} function to normalize the spectra in each echelle order. This allows comparison of the spectra at different epochs, however it smeared out features that can be as wide as the echelle orders themselves. In following subsections we will discuss various emission and absorption features seperately. Finally, to address the origin of various components, we need to assess the accuracy with which we can measure the radial velocities of the lines and the possible systematic errors affecting their respective values. It is always possible that instrumental effects (echelle spectrograph flexure) introduce pixel shifts and pseudo-variations in the radial velocities of the lines. Taking comparison lamp spectra before and after each object exposure is a common procedure to minimize this effect; nevertheless, an independent confirmation is necessary. In the spectrum of QU~Car, there are two diffuse interstellar bands (DIBs) at 5870$\AA$ and 5797$\AA$ at adjacent blue orders, the oxygen A and B bands and the telluric bands at red orders. We have used the narrow DIB features to ascertain the accuracy of our wavelength calibrations inasmuch as the radial velocity shifts in the Na I D absorption occurs at different temporal epochs and could conceivably be due to slight changes in the wavelength calibration due to the instrument configuration. Measurement of the DIBs radial velocities confirms the accuracy of our radial velocity determinations within an RMS error of $\sim$$\pm$0.10 km/sec. \subsection{Emission lines} Table 2 lists the equivalent widths (EWs) of the most prominent features in our echelle spectra. By comparing to the emission line strengths in 2006/2007 (Table 3 of Paper~I) we see that the new EWs are in general consistent with those measured in 2006/2007 using lower resolution data. Both the 2006/2007 EWs and the new data are characterized by occasional changes of over 2$\times$ between nights, and sometimes within a night, mixed with intervals of relative stability. Note in particular that in 2006/2007 the Bowen blend was nearly constant over 6 months at EW$\sim$3.2\AA, which is a little weaker than in 1979/1980 (Gilliland \& Phillips 1982). In our 2010 spectra this feature is much weaker, and has practically disappeared by E3. The Balmer lines are also weak at this time, but HeII is stronger than the 2006-2011 average. Apparently these various emission lines originate in quite different regions of the system, which can vary independently. \begin{figure} \includegraphics[width=10cm, angle=0.0]{kafka.figure1.ps} \caption{Visual AAVSO and V Mag. ASAS light curves of QU Car. This plot demonstrates that bright and faint states are present in the long-term optical behavior of QU Car, although the transitions are not well defined due to lack of continuous monitoring. The four epochs of the 2006-2007 low-resolution spectroscopic observations of paper~1 are highlighted in the inset, to demonstrate that QU Car was at its ``bright'' magnitude at the time of the relevant observations. Photometric data were not available during our 2010-2011 observations. \label{lc}} \end{figure} \begin{figure} \includegraphics[width=12cm, angle=0.0]{kafka.figure2.ps} \caption{Balmer lines for epochs E1-E6 (E4 is missing, since the relevant data are not used in this work). The Balmer line profiles seem to be smoother in 2010, with multiple absorption components in 2011. A stationary redshifted emission component at $\sim$-135km/sec is representing stationary circumbinary gas. Finally, the blueshifted and redshifted absorption troughs in the 2010 H$\beta$ line likely represents material expelled from the binary. \label{balmer}} \end{figure} \begin{figure} \includegraphics[width=12cm, angle=0.0]{kafka.figure3.ps} \caption{Same as in figure~\ref{balmer}, but for the Bowen blend and HeII4686 emission. The dashed horizontal lines guide the eye to the normalized continuum. Notice the attenuated Bowen blend lines in 2010 with respect to their appearance in 2011.\label{bowen}} \end{figure} Figures 2 and 3 show the line profiles of some of the emission lines, where the higher spectral resolution of the echelle data reveal considerable details compared to earlier studies. At times the H$\alpha$ profile resembles the line profiles in V Sge (Gies, Shafter \& Wiggs 1998; Robertson, Honeycutt \& Pier 1997), with multiple components whose relative strengths vary. Note that the H$\alpha$ profile sometimes shows weak blue-shifted absorption at a velocity of $\sim$-300 km s$^{-1}$ This wind-induced P-Cygni feature was also seen in the H$\alpha$ profiles of Paper I. For all four of the 2011 spectra (E5 and E6) the H$\alpha$ emission line is accompanied by 7 narrow weak absorption features. These features are identical in these 4 spectra (separated by 3 nights), and do not appear in any of the 2010 spectra (E1-E3). They do not agree with the wavelengths of telluric water features (e.g. Vince \& Vince 2010) nor do they appear in our standard stars (or any of the stars observed the same night in neighboring parts of the sky, bracketing the QU Car observations). In the absence of an accurate ephemeris, we have few clues to the physical locations of those lines. The fact that the EWs of HeII and HeI do not follow changes in the Balmer and Bowen blend suggests that those lines have different origins in the system (as expected from their different excitation potentials). Using radial velocity arguments Gilliland and Phillips (1982) argue that CIII~4650$\AA$ should be the dominant species in the Bowen blend. Furthermore, UV spectra \citep{2003MNRAS.338..401D} indicate that carbon is overabundant in QU Car. A Carbon overabundance in QU Car is also supported by the presence of the strong CIV~5801,5812$\AA$ emission, in the Gilliland \& Phillips (1982), Drew et al. (2003) and Paper~I optical spectra. Drew et al. (2003) argue that this carbon enhancement should originate from the envelope of the donor star, suggesting that the star is an early-R type. However, the observed variations in the Bowen blend and in CIV in our spectra argue against such an interpretation. Emission from the secondary star could not have been diminished in one epoch of our observations but still be present at a different one. In 2010 both the Bowen blend and CIV are significantly reduced in strength; this is an indication that the lines are variable in nature. A plaussible interpretation is that carbon originates from the atmosphere of the the white dwarf of the binary, which is attenuated in the faint state. In 2011, when the CIII/NIII/OII line strength of the binary returned to its ``bright'' state value, weak CIV~5801,5812$\AA$ emission is also present. In our 2006/2007 spectra, we also detected variable [OIII]~5007$\AA$ and [NII]~6584$\AA$ emission, indicative of the presence of a nebula around the system. The [OIII]~5007$\AA$ lines appeared to have two components, ranging in velocities from -500km/sec to 370km/sec - probably representing the front and back of an expanding shell or an outflow. In our new echelle data, only one (weak) component of the [OIII]~5007$\AA$ line is present; its strength is reduced with time (perhaps with orbital phase) between E1 and E2. In 2011, the [OIII]~5007$\AA$ emission is present only during E5, with no accompanying [NII]~6584$\AA$ emission. This erratic variation in the oxygen emission line strength does not seem to be associated with variations in the strength of the Bowen blend. However, we are in need of a better ephemeris for the system in order to correlate the observed oxygen emission to a specific location and process on the binary. \subsection{H$\beta$ components: evidence for circumstellar material?} We now turn our attention to the mysterious absorption components bracketing H$\beta$ in 2010. Having excluded contamination in our spectra (flat fielding artifacts, data reduction faux pas or background reduction residuals), the features are convincingly real. Since there are no strong telluric features at this part of the spectrum (nor is there any chance for this absorption to come from the interstellar medium, especially considering its variable nature and appearance), the features should originate from the binary. The red component is stationary at $\sim$1110km/sec. The blue component is present only during two subsequent nights of our observations (second spectrum in E1 and first spectrum in E2, at -370km/sec and -1050km/sec respectivelly). Considering that our exposure times are 2000 sec, we are surprised to not detect this red component in all spectra in E1 and in E2. Among CVs, those features are unique to QU Car. However, similar absorption lines sometimes appear in other accreting sources: complex blueshifted and redshifted absorption components emerge in the Balmer lines of Herbig Ae/Be stars (e.g. Guimaraes et al. 2006). In this case, the redshifted absorption is usually interpreted as being due to material accreted onto the star from its inner disk (sometimes via magnetic fields) and the blueshifted one as being due to mass loss (outflows). Both processes can happen simultaneously (e.g. Natta, et al. 2000). Taking the Ae/Be work as a starting point, we could attribute the blue components to ejecta from QU Car, while the redshifted components originate on an inflow from the cirumstellar medium. It is difficult to imagine a mechanism that allows for both processes simultaneously, unless the flows are magnetically controlled (as is the case in many T Tauri stars). However, there is no evidence that QU Car contains a highly magnetic white dwarf (i.e., no Zeeman or cyclotron features). Furthermore, it seems unlikely that this material represents a jet. Jet-like features in CVs sometimes appear in dwarf novae in outburst or in some novae, both times representing material ejected from the binary during explosions (e.g. Cowley et al. 1998). In semi-detached binaries, optical jets also appear in symbiotic stars and in low mass x-ray binaries as {\it emission} components moving away from Balmer lines with velocities reaching 4000km/sec ( e.g. Cowley et al. 1998). Perhaps the gas motions in QU Car are so turbulent and complex that portions of the flow can be both red and blue shifted in front of the same continuum source. It is unfortunate that we lack optical photometry (or any light curves) at the time of our spectroscopic observations (stressing again the importance of having simultaneous multiwavelength observations for a comprehensive understanding of QU Car). \subsection{ NaI D absorption: connection with the supernovae Ia?} The unusual appearance of the NaID lines is revealed in figure~4: in 2010, each component consists of weak blueshifted emission accompanied by redshifted absorption! Those lines (and the adjacent HeI 5876 emission) are at the center of the order during all our observing runs, ruling out radial velocity variations introduced by pixel shifts due to poor echelle order distortion correction at the edges of each order. The EW and RV values are given in Table~3, where missing entries indicate that the line was not detected. The emission components are very narrow and appear slightly blueshifted with respect to the line's rest wavelength. Furthermore, this emission is quite variable and present only during 2010. No correlation of the line strength and radial velocity was found. In CVs, the NaID lines are usually in absorption, arising from the photosphere of the donor star. Sodium is usually ionized in accretion disks; its presence in emission in QU Car is indicative of lower than usual temperatures. This suggests a different optical state for the system, perhaps similar to V Sge's "faint" state, where the mass transfer rate is significant lower, in agreement with the reduced strength of the Balmer and the Bowen blend line profiles. We are intrigued by the "stationary" (constant velocity) nature of this NaID emission component over the four nights of our observations in 2010. {\em If} this line has an origin on the accretion disk or on the hotspot, it may indicate that the QU Car's disk outer edge reaches the center of mass of the binary, in agreement with QU Car being an unusually bright object. The RVs of both NaD absorption components are constant at ~$\sim$-7 km s$^{-1}$ during E1-E3. Eight months later these RVs are again constant but at ~$\sim$-13 km s$^{-1}$ during E5-E6. Both components of the NaID are well displaced from their rest wavelength as they accelerate between the 8 months of observations. In the bright state, variable NaID is also present (e.g. Paper~I), however, with the lower resolution of Paper~I it was not possible to resolve individual absorption components. Furthermore, in our echelle spectra, we cannot kinematically associate any other emission line components with the absorption features. Since neither component of the NaID doublet is at the local standard of rest velocity, we explore the possibility that they are interstellar in origin, following the Galactic rotation at the location of QU Car. \citep{2003MNRAS.338..401D} estimate the galactic rotation velocity at $\sim$2kpc along the line-of-sight to QU Car to be -18.3km/sec and attribute FeII~1608$\AA$ and SII~1250$\AA$ bluehifted line components to interstellar absorption. As a side-note, the minimum distance to QU Car is suggested to be 500pc (Linnell et al. 2008 and references therein), which would imply less interstellar absorption. However, for our purposes, the distance to QU Car and amount of interstellar abosprtion is irrelevant since we are discussing radial velocity {\it variations} in the absorption lines, not their absolute values.\footnote{ It is also possible that an interstellar component is present and unresolved in the observed lines, however it will simply introduce an error in the estimate of the line EW.} In this case, the variable velocities of NaID components between 2010 and 2011 argue against an interstellar origin. Moreover, NaID absorption could originate on the donor star of the binary. Because of the absence of other photospheric identification features from this star (e.g. FeI, CaI, MgII, Ti or even Balmer absorption) we rule out this possibility. In CVs NaI absorption lines can be generated in nova explosions (e.g. Shore et al. 2011); however QU Car has no indications of recent erruptions. A possible scenario comes from the supernova community: \citep{2007Sci...317..924P} describes variable NaI D absorption in the vicinity of SN 2006X, originating from circumbinary clouds, heated by the SN and interacting with the explosion ejecta. Patat et al. (2007) argue that the NaI was present in the SNeIa progenitor as circumbinary material. UV radiation from the SNeIa explosion ionized this material, which slowly recombined post-SNeIa to produce the observed variable NaID absorption. The absence of CaII H\&K absorption lines was attributed to the lack a radiation field hot enought to significantly ionize Ca. Variable NaID absorption has now been detected in a group of SNeIa (Sternberg et al. 2011; Simon et al. 2009; Simon et al.2007), favoring the single degenerate scenario (C-O WD + giant donor star) for a SNeIa progenitor. Circumstellar clouds continuously replenished by a wind from the giant donor star provide a plausible mechanism for the formation of NaID lines in QU Car. Therefore, our detection of variable NaID lines around QU~Car provides, for the first time, a bridge between this class of SNe and their elusive progenitors. \begin{figure} \includegraphics[width=10cm, angle=0.0]{kafka.figure4.ps} \caption{Same as figure~3, for the NaD line region.} \end{figure}
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1206.6798
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1206.1642.txt
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We consider a class of metric $f(R)$ modified gravity theories, analyze them in the context of a Friedmann--Robertson--Walker cosmology and confront the results with some of the known constraints imposed by observations. In particular, we focus in correctly reproducing the matter and effective cosmological constant eras, the age of the Universe, and supernovae data. Our analysis differs in many respects from previous studies. First, we avoid any transformation to a scalar-tensor theory in order to be exempted of any potential pathologies (e.g. multivalued scalar potentials) and also to evade any unnecessary discussion regarding frames (i.e. Einstein {\it .vs.} Jordan). Second, based on a robust approach, we recast the cosmology equations as an initial value problem subject to a modified Hamiltonian constraint. Third, we solve the equations numerically where the Ricci scalar itself is one of the variables, and use the constraint equation to monitor the accuracy of the solutions. We compute the ``equation of state'' (EOS) associated with the modifications of gravity using several inequivalent definitions that have been proposed in the past and analyze it in detail. We argue that one of these definitions has the best features. In particular, we present the EOS around the so called ``phantom divide'' boundary and compare it with previous findings. \bigskip {\bf Keywords:} modified gravity, equation of state, cosmological parameters
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\label{sec:introduction} Astronomical observations based on type Ia supernovae (SNIa) together with the assumption that the Universe is homogeneous and isotropic at large scales led to the conclusion that the Universe is currently expanding in an accelerated way~\cite{Perlmutter1999,Riess1998,Amanullah2010}. This phenomenon can be most easily explained by appealing to the existence of a cosmological constant $\Lambda$ (sometimes termed {\it dark energy}). This constant along with the introduction of dark matter (DM) apparently needed in many regions of the Universe (galaxies and clusters) have originated what is called today the $\Lambda CDM$ paradigm. This paradigm has also successfully explained most of the current details of the Cosmic Background Radiation (CBR or CMB) in the framework of general relativity~\cite{WMAP}, as well as other important features of the Universe at large scales~\cite{LSS} (for a thorough review see Ref.~\cite{Weinberg2012}). However, despite of the simplicity and success of this paradigm, several theoretical as well as epistemological arguments have been put forward as objections against such a simple model of the Universe. For instance, as concerns the DM hypothesis, one of the the main criticisms is that its nature (i.e. its quantum and classical properties) is not well understood (if at all) yet. That is, apart from the gravitational evidence, there is no further strong reason supporting its existence. Since several experiments have failed so far to detect the proposed DM particles, skepticism keeps growing in this direction. On the other hand, the cosmological constant has been historically regarded as ``suspicious'' by several detractors (including Einstein himself; see Refs.~\cite{Lambda} for a review), although some of its apparent drawbacks are based more on prejudices than on strong and well grounded physical arguments~\cite{Bianchi2010}. In any case, the discomfort that $\Lambda$ has produced in the spirit of some people has led to consider more complicated alternatives, of which, is fair to say, none is regarded today as a more serious candidate for dark energy (DE) than $\Lambda$ (the BigBOSS experiment~\cite{BB} has been designed to shed light in this direction). Among these alternatives are the so called modified theories of gravity (MTG) as opposed to general relativity (GR). Some of these theories have been also proposed to substitute DM and even as models for inflation. Perhaps the most popular MTG over the past ten years and the one we focus in this article are $f(R)$ metric theories, where an {\it a priori} arbitrary function of the Ricci scalar $R$ replaces $R$ itself in the gravitational Lagrangian. This kind of MTG were conceived originally in order to create a late accelerated effect without a cosmological constant or as inflationary model without an extra scalar field (see Refs.~\cite{f(R),Capozziello2008a,Sotiriou2010,deFelice2010} for a detailed review). Notwithstanding, despite of the promising $f(R)$ models first proposed to replace $\Lambda$ \cite{accexp}, a cumulative evidence, both theoretical and observational, has been found against most of them. However, new models have been proposed to overcome the initial difficulties, some better motivated that others but none introducing a new fundamental principle that can be used as a guiding line; indeed they have rather been constructed by trial and error. The simplest (non-trivial) choice $f(R)=R$ was historically favored by Einstein since mathematically led to second order partial differential equations (PDE's) which could easily reduce to the Newtonian theory in the week field limit. General mathematical and physical conditions are usually demanded in order to avoid pathologies in the models. For instance, the conditions $f_{RR}>0$ and $f_R>0$ (where the subindex indicate derivative with respect to $R$) seem to be required for stability considerations and to ensure a positive definite effective gravitational constant, respectively. It is however not clear if those conditions are really necessary, and in many studies they are not imposed. Therefore, in most cases, ``handcraft'' has been used to design a particular $f(R)$ model based on heuristic arguments that might account for the actual phenomenology when the full fledge model is submitted to a detailed scrutiny. The general trend so far is that no single model is able to explain most of the current observations, but only some aspects of them. That is, most of the $f(R)$ models fail miserably when they are preempted as models for all the dark substance (both DE and DM) and when taking into account the Solar System tests as well. Even when considered only as DE models, most of them fail, with the exception of some notable cases. Of course it could well happen (but perhaps not very desirable) that dark matter, dark energy and modifications of the laws of gravity in some combination are required by nature in order to fully understand our Universe. In fact, this is the approach we pursue here at the cosmological level except that we do not include explicitly a cosmological constant, but rather, the models themselves give rise to an effective $\Lambda$ which vary slightly around the required value at late (i.e. present) epochs of the Universe. We then consider $f(R)$ theories simply as a model for explaining the accelerated expansion of the Universe and introduce a dark matter component in the same way as in the $\Lambda CDM$ paradigm. Nevertheless, things turn out to be not so simple, given that the proposed models can disturb the successes of GR. Any proposed specific $f(R)$ model has not only to satisfy the cosmological observations, but all the gravitational observations at all scales. These include the Solar System experiments, the existence of physically acceptable compact objects, the binary pulsar, etc. As today, there is no single MTG model that replaces successfully GR and explains as should be, all the observations for which it was designed originally. After the first $f(R)$ models were proposed to explain the accelerated expansion of the Universe, among them the ``historic'' $f(R)= R - \mu^4/R$, a sequence of papers appeared where the constraints imposed by the Solar System were taken into account \cite{solarsystem1}. Without reaching a clear consensus on the issue, it seemed that such models were not viable. One of the arguments put forward to establish that conclusion was based on the fact that such theories can be shown to be dynamically equivalent to a Brans-Dicke (BD) theory with $\omega=0$. Since such a value for $\omega$ in these theories gives rise to a post-Newtonian parameter $\gamma=1/2$, which conflict with $\gamma\sim 1$ favored by the Solar System tests, then at first sight the analysis suggested that all $f(R)$ theories were excluded blatantly as viable theories. Much later, it was recognized that such an argument should be used with care in view that $f(R)$ theories are not equivalent to the standard BD theory with $\omega=0$, but to a BD theory with a potential. Therefore, depending on the mass of the effective scalar, the theory at hand could pass or fail the Solar System tests~\cite{solarsystem2}. Although it is now recognized that many of the $f(R)$ models give rise to $\gamma\sim 1/2$, and are therefore ruled out, some others, due to the effective mass of the scalar, might produce a successful phenomenology. This success depends on whether or not the scalar field which is associated with the model at hand can act as a {\it chameleon}~ \cite{chameleon,Hu2007}, a mechanism that appears in some scalar-tensor theories of gravity which allows them to satisfy the local tests and the possibility of producing the required cosmological effects~\cite{Khoury}. Now, as concerns the cosmological restrictions on these theories, Amendola {\it et al.}~\cite{Amendola2007a,Amendola2007b} have devised criteria of quite general applicability that allows to discard many of the proposed $f(R)$ models. In short, their analysis shows for a large class of models that they either produce an accelerated expansion at recent times but fail to generate a correct matter-dominated era (the scale factor behaves as radiation in GR) or the opposite. In fact, when viewed from the past to the present, many of such models cross from the radiation-dominated era (deceleration epoch) to an effective dark-energy era (acceleration epoch) with a very short or not even existent matter-domination era. Such models are therefore incompatible with the CBR observations, the age of the Universe and the structure formation at large scales~\cite{Amendola2007a,Amendola2007b}. This issue was, however, not free of debate either~\cite{Capozziello2007,Amendola2007c}. To make things even more confusing in this matter, additional skepticism was raised about the viability of such kind of theories when a further test was performed on several cosmologically successful $f(R)$ models. This time, the test consisted in analyzing the possibility that such models allowed the construction of solutions representing realistic (or at least idealized) neutron stars. In a first attempt to do so, Kobayashi \& Maeda~\cite{Kobayashi2008} showed that idealized neutron stars (specifically, incompressible compact objects) were not allowed by the Starobinsky model~\cite{Starobinsky2007} since a singularity in the Ricci scalar developed within the object. Later, Babichev \& Langlois~\cite{Babichev2009} criticized such conclusion and argued that the singularity was only due to the use of an incompressible fluid and when a similar situation was analyzed with a compressible gas (e.g. a polytrope) no singularity was found. Finally, Upadhye \& Hu~\cite{Upadhye2009} argued that a chameleon effect was the responsible of avoiding the formation of singularities in compact objects and not the use of more realistic equation of state. In a more recent work by us~\cite{Jaime2011}, while we arrived to the same primary conclusion of Refs.~\cite{Babichev2009,Upadhye2009} on that no singularities were necessarily formed, we criticized some very basic aspects of the analysis common to all of the three above investigations concerning the existence of neutron stars~\cite{Kobayashi2008,Babichev2009,Upadhye2009}. To be specific here let us mention that their analysis relied fundamentally on the fact that a potential associated with a scalar-tensor counterpart of the Starobinsky model could drive the scalar field to a point where the corresponding Ricci scalar diverged. The main point of our criticism, noticed earlier in~\cite{Goheer2009,Miranda2009}, is that such a potential has unpleasant features as it is multivalued and therefore, the transformation to the STT is not well defined. Since the conclusions reached on those papers depend crucially on the use of such a pathological potential we consider them not very trustworthy, even if the dynamics is supposed to take place in a region where the potential is single valued. Moreover, due to the fact that many of the controversies regarding this subject have been the result of using the STT in the Einstein or Jordan frames (see also Ref.~\cite{Multamaki2006} for a similar criticism), we strongly suggested to abandon such an approach and treat $f(R)$ models in all applications without performing the transformation to any frame of STT. In the particular case of static and spherically symmetric spacetimes, we devised a very transparent and simple approach that allowed us to deal with compact objects. Using a particular case of the Starobinsky model we concluded that no singularities were found. We also studied the Miranda {\it et al.} model~\cite{Miranda2009} (which was previously shown to be free of singularities following the STT approach but with a single-valued potential), and arrived to the same conclusion. We emphasized the advantages of using that robust approach over the STT technique, and we argued that even in the cases where the STT transformation is well defined it is not particularly useful and that the original variables are required anyway in order to interpret the solutions correctly. Furthermore, and as we mentioned above, dealing with the STT approach opens the way to the long-standing controversy about which frame (Einstein or Jordan) is the physical one~\cite{frames}. This discussion is in some cases semantic and in many others completely ill founded and corrupted. In our treatment we don't even need to deal with it at all since we consider the theory directly as it emerges from the original action without performing any formal or ``rigorous'' identification or transformation with any other theory, frame or variables. In the present article we review, in light of our robust approach, the cosmological analysis of some of the apparently least problematic $f(R)$ models, in the sense that they seem to pass the cosmological and Solar System tests (although the model of Ref.~\cite{Miranda2009} is still on debate). As we will show, our treatment allows to handle the equations most as in the case of GR, which in turn, makes possible the use of simple techniques of numerical relativity to monitor the accuracy of the solutions. The paper is organized as follows, Section II introduces the $f(R)$ theory, the general equations and the approach under which we will treat them. In Section III we present four inequivalent definitions of the ``energy-momentum'' tensor of {\it geometric dark energy} that we have identified in the literature and which give rise to three inequivalent equations of state (EOS) used in cosmology. Section IV displays in detail the three specific $f(R)$ models that we submit to a cosmological analysis. In Sections V and VI we focus on the Friedmann-Robertson-Walker (FRW) cosmology and analyze the EOS that arises in the particular $f(R)$ models we treat, as well as the relative abundances of the different matter-energy components in order to explicitly show the matter and modified-gravity (geometric dark energy) domination epochs. In order to have some insight about the viability of the solutions, we compare the results with the successful GR-$\Lambda CDM$ scenario. We also compute the luminosity distance and confront the results with the historic SNIa data~\cite{Riess1998} and the UNION 2 compilation~\cite{Amanullah2010}. The age of the Universe that arises from these models is also estimated. Finally, Section VII concludes with a summary and a discussion. An Appendix displays the dimensionless form of the cosmological equations used for numerical integration and the numerical test used to check the accuracy of our solutions. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%% f(R) GRAVITY THEORIES %%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
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\label{sec:discussion} Modified theories of gravity, like metric $f(R)$ theories, have been analyzed thoroughly in recent years mainly to explain the accelerated expansion of the Universe which was inferred from the measurements of luminous distance of SNIa, while some other theories have been put forward to avoid the need of dark matter. In this article we have focused mainly on the mechanism to produce an accelerated expansion. Such behavior can be most easily accounted within the general theory of relativity and by (re)introducing the cosmological constant which was proposed by Einstein almost one hundred years ago. The motivation behind the proposal for modifying GR was to avoid the introduction of such constant and then to circumvent the apparent problems associated with it. Furthermore, such an alternative exempts us from adding new fields (scalar or otherwise) like quintessence or k-essence, in order to explain that acceleration. Nevertheless this alternative adds, to our opinion much more troubles than solutions. While many specific $f(R)$ models are able to produce an accelerated expansion similar to the $\Lambda CDM$ model, they have, at the same time, spoiled many of the successes of GR or are inconsistent with other features of cosmology, like an adequate matter dominated era. Only a few models have succeeded in explaining, at least partially, the actual cosmological evolution of the Universe without disturbing, for instance, the predictions at Solar System scales. Since, $f(R)$ theories do not introduce a new fundamental principle of nature, there is then, not a deep criteria that favors one among the apparently viable $f(R)$ models. Basically some kind of ``handcraft'' have been used so far to mold specific $f(R)$ or in other cases even reconstruction methods~\cite{Dunsby2010,reconstruction}. At any rate, simplicity would be in favor of GR with $\Lambda$. In this article we argued that even if this kind of modified theories can be internally consistent, some special care has to be taken into account when they are analyzed. For instance, it was ``discovered'' that $f(R)$ theories can be equivalent to scalar-tensor theories of gravity. This identification is free of inconsistencies provided that the mapping between both representations is well defined. For that it is required that the function $f_R$ be a monotonic function of the Ricci scalar $R$. Several of the apparently successful models fail to fulfill this condition in general, and yet, different authors have used the STT representation. One of the consequences is that the scalar field potential turns to be multivalued. This pathology has contributed to create confusion in the subject, in addition to the already existing confusion between frames in STT. We have emphasized that in order to avoid such potential drawbacks $f(R)$ theories should be treated using the original variables. This is not only possible, but in our view, it turns to be much more transparent, even if the equations seem more involved at first sight. Although some other people share this view and have used the original variables in several applications, we give a step forward and propose here a general system of equations simpler than the one usually used, and when applied particularly to cosmology, it reduces further and can be treated numerically as an initial value problem, where the initial values are restricted by a modified Hamiltonian constraint. Previously we presented a system of equations that can be used to construct compact objects in static and spherically symmetric spacetimes and showed the way to treat them numerically~\cite{Jaime2011}. In the current article we analyzed three specific $f(R)$ models that are viable, at least in the background, since they provide an adequate matter dominated behavior followed by a correct accelerated expansion. It has been argued that the Miranda {\it et al.} model~\cite{Miranda2009} is ultimately incompatible when cosmology is analyzed at the level of perturbations or in the Solar System~\cite{delaCruz2009}, but very likely a deeper analysis is required in order to rule out this model completely~\cite{Miranda2009b}. In the future we plan to analyze some other models, like the exponential ones~\cite{Zhang2006,exponential,gravwaves1} which can be viable as well. For each of the three classes of $f(R)$ models that we considered here, we analyzed three possible inequivalent equations of state associated with the modified gravity and which have been studied in the past by several authors. Such EOS arise from general EMT that represent the modified gravity or geometric dark energy, although such interpretation is to be handled with care as several authors have warned before. We have identified at least four inequivalent definitions of the EMT which lead to the three EOS just alluded. Two of such EMT's have the unpleasant feature that are not conserved. One of the other four EOS can diverge at some red-shift. Thus the more appealing definition comes from Recipe I whose EMT is conserved and the corresponding EOS behaves appropriately. We emphasized that it is important to come to an agreement on which definition of the EOS will be ultimately compared with the cosmological observations. This is crucial in view that the forthcoming data might differentiate between them as precision is increased~\cite{BB,obsEOS}. It is natural to ask about the new predictions that $f(R)$ theories can make in addition to explaining the accelerated expansion while reproducing, in the viable cases, the previous successes of GR and in particular the $\Lambda CDM$ model. There are indeed several predictions that, in principle, would allow us to distinguish between GR and $f(R)$ theories. One of them is the additional polarization modes of gravitational waves like the ``breathing'' mode~\cite{gravwaves1,gravwaves2}. At present time the current gravitational-wave detectors are not sensitive enough to detect gravitational radiation and therefore, this feature cannot be used as a tool to falsify alternative theories of gravity yet, but in a near future this will be certainly possible~\cite{Will}. The binary pulsar on the other hand can be an excellent tool to do so at present time. We know that the binary pulsar can constrain STT even if they succeed in passing the Solar System tests~\cite{Damour96}, thus, the same might happen for $f(R)$ theories. At the cosmological level, $f(R)$ theories predict a large-scale integrated Sachs--Wolf effect which is different from the one predicted in GR~ \cite{Zhang2006,Zhao2010,ISW,Bertschinger2008}. This is because the total EMT that one can associate with such theories is not necessarily ``isotropic'' at the time of recombination and thus the metric potentials of scalar perturbations in the Newtonian gauge are not equal (in absolute value). Nevertheless, due to the so called cosmic variance, it turns difficult to use it as a further constraint for some values of the model parameters. On the other hand, weak and strong gravitational lensing as well as the growth of matter perturbations can also be affected when those potentials are different ~\cite{Zhang2007,Zhao2010,Bertschinger2008,HS&Staromodel}, and so they can further constrain the specific models. In this article we have not attempted to best-fitting the parameters using current observations, but rather using the $\Lambda CDM$ models as a standard, taking a zero spatial curvature as a prior. The former analysis will be pursued in a future work. %Finally, it can be interesting to search for an $f(R)$ theory that can explain in a ``universal'' fashion the rotation curves of spiral galaxies. %This amounts to look for an $f(R)$ theory that reduces in some limit to the Milgrom's Modified Newtonian Dynamics~\cite{MOND} %without changing the parameters of the theory for each galaxy. This unpleasant adjusting is a common drawback for many of the %relativistic theories or dark matter models proposed so far to explain such curves and which is very often overlooked or dismissed. %%%%%%%%%%%%%%%%%%%% %%% APPENDIX %%% %%%%%%%%%%%%%%%%%%%% \newpage
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1206.1806_arXiv.txt
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{ We report on follow-up observations of 20 short-duration gamma-ray bursts (GRBs; $T_{90}<2$~s) performed in $g'r'i'z'JHK_s$ with the seven-channel imager GROND between mid-2007 and the end of 2010. This is one of the most comprehensive data sets on GRB afterglow observations of short bursts published so far. In three cases GROND was on target within less than 10 min after the trigger, leading to the discovery of the afterglow of GRB 081226A and its faint underlying host galaxy. In addition, GROND was able to image the optical afterglow and follow the light-curve evolution in further five cases, GRBs 090305, 090426, 090510, 090927, and 100117A. In all other cases optical/NIR upper limits can be provided on the afterglow magnitudes. After shifting all light curves to a common redshift we find that the optical luminosities of the six events with light curves group into two subsamples. GRBs 090426 and 090927 are situated in the regime occupied by long-duration events (collapsars), while the other four bursts occupy the parameter space typical for merger events, confirming that the short-burst population is contaminated by collapsar events. Three of the aforementioned six bursts with optical light curves show a break. In addition to GRBs 090426 and 090510 (paper I, II), also for GRB 090305 a break is discovered in the optical bands at 6.5 ks after the trigger. For GRB 090927 no break is seen in the optical/X-ray light curve until about 150~ks/600~ks after the burst. The GROND multi-color data support the view that this burst is related to a collapsar event. For GRB 100117A a decay slope of its optical afterglow could be measured. For all six GRBs at least a lower limit on the corresponding jet opening angle can be set. Using these data, supplemented by a about 10 events taken from the literature, we compare the jet half-opening angles of long and short bursts. We find tentative evidence that short bursts have wider opening angles than long bursts. However, the statistics is still very poor.}
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Gamma-Ray Bursts (GRBs) show a bimodality in their duration distribution, separated in the CGRO/BATSE data at $T_{90}=2$~s, with the peak of the short-burst population at $T_{90}\sim$0.5~s and the long-burst population at $\sim$30~s (\citealt{Kouveliotou1993,Sakamoto2011ApJS}). Historically, bursts are still devided into long and short based on the BATSE scheme, even though the shape of the bimodal distribution is energy-dependent, in particular peaking for \swift/BAT at $T_{90}\sim0.5$~s and $\sim70$~s, respectively (\citealt{Sakamoto2011ApJS}). According to the current picture, long bursts originate from the collapse of massive stars into black holes (\citealt{MacFadyen1999}) or into rapidly spinning, strongly magnetized neutron stars (e.g., \citealt{Usov1992,Mazzali2006}). Short bursts are instead commonly attributed to the merger of compact stellar objects (e.g., \citealt{Paczynski1986,Nakar2007}). The physical association of long bursts with the collapse of massive stars has been well established (e.g., \citealt{ZehKH2004,Hjorth2003Natur,Pian2006Natur,Ferrero2006,WB2006, Fruchter2006}). However, the observational situation with short bursts is less clear. Until 2005 no afterglow of a short burst had ever been detected, while for the long burst sample at that time many important discoveries had already been made (redshifts, supernova light, collimated explosions, circumburst wind profiles). The first well-localized short burst (GRB 050509B; \citealt{Gehrels2005Natur}) was seen close in projection to a massive early-type galaxy (\citealt{Hjorth2005,Bloom2006ApJ638}), supporting the model that compact stellar mergers are the progenitors of short-duration gamma-ray bursts. However, since then the observational progress has been rather modest when compared to the long-burst population (for a review \citealt{Gehrels2009,Berger2011NewAR}). There are mainly two reasons for this situation. Firstly, compared to long bursts there is a substantially smaller detection rate of short bursts. Secondly, short-burst afterglows are rarely brighter than $R=20$ even minutes after a trigger (e.g., \citealt{Kann2010,Kann2011}). This general faintness makes their discovery and detailed follow-up very challenging. However, only the precise detection of the afterglow, with sub-arcsec accuracy, enables a secure determination of a putative GRB host galaxy and its redshift, while the X-ray plus optical light curves provide information about the processes that take place after the explosion, clues about the physics of the central engine, and the properties of the environment of the progenitor. Rapid follow-up observations of these events are therefore very important to gain as much observational data as possible. Since there is a substantial overlap between the long and the short-burst duration distribution, the simple devision between long and short is only a first guess about the true origin of a burst under consideration. Several other phenomenological properties of the bursts and their afterglows have to be considered in order to reveal the nature of their progenitors (\citealt{Zhang2007,Zhang2009,Kann2011}). Thereby, of special interest are the circumburst density profiles, the afterglow luminosities, and the outflow characteristics that might be shaped by or related to the physical properties of the GRB progenitors. Theoretical studies suggest that long GRBs are followed by more luminous afterglows than short bursts, mainly due to the expected difference in the circumburst density around the GRB progenitors (\citealt{Panaitescu2001}). Also the circumburst density profile is an indicator on the nature of the explosion (e.g., \citealt{Schulze2011}). In addition, the distribution function of the jet-opening angles of long and short bursts should be different from each other since an extended massive envelope collimates the escaping relativistic outflow \citep{Zhang2004ApJ608}, while the lack of such a medium in the case of merger events might allow for wider jet-opening angles \citep{Aloy2005,Rezzolla2011}. Any short-burst afterglow that adds information here is naturally of great interest. Here we report on the results of the first 3.5 years of follow-up observations of short-duration GRBs using the optical/NIR seven-channel imager GROND (\citealt{Greiner2007Msngr,Greiner2008}) mounted at the 2.2-m ESO/MPG telescope on La Silla (Chile). GROND is in continuous operation since mid-2007. Since then it observes every burst with a declination $\lesssim +35^\circ$, providing a complete sample of events observed with the same instrument at the same telescope. The capability of GROND to observe in seven bands simultaneously, from $g'$ to $K_s$, does not only provide the opportunity to follow the color evolution of an afterglow but also allows for a stacking of all bands; in particular a white-light image in $g'r'i'z'$ reaches a fainter detection threshold. In addition, GROND's routine operation in Rapid Response Mode in principle allows us to start observations within minutes after a trigger, catching also afterglows even if they are fading rapidly. In this work, we summarize the detections and upper limits for 20 short burst afterglows in $g'r'i'z'JHK_s$. First results have already been published in \citet[][ in the following paper I]{NicuesaGuelbenzu2011a} and \citet[][ in the following paper II]{NicuesaGuelbenzu2012a}. Here we add detailed information on all individual bursts. In particular, we compare the afterglow luminosities with those of their long-burst relatives. We also include X-ray data in order to extend this discussion to the high-energy band. If possible, based on our optical data, we derive the spectral energy distribution (SED) of the afterglows and give an estimate of the corresponding jet half-opening angles. Throughout the paper, we adopt a concordance $\Lambda$CDM cosmology ($\Omega_M=0.27$, $\Omega_{\Lambda}=0.73$, $H_0=71$~km/s/Mpc; \citealt{Spergel2003}), and the convention that the flux density is described as $F_\nu (t)\propto t^{-\alpha}\,\nu^{-\beta}$. In cases where no redshift is known for a burst, we adopt a redshift of $z$=0.5, as it is justified based on the redshift distribution of short bursts detected by \swift \ by the end of 2010 (\citealt{Leibler2010}, their table 1).
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1206.3803_arXiv.txt
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We derive an analytic phenomenological expression that predicts the final mass of the black-hole remnant resulting from the merger of a generic binary system of black holes on quasi-circular orbits. Besides recovering the correct test-particle limit for extreme mass-ratio binaries, our formula reproduces well the results of all the numerical-relativity simulations published so far, both when applied at separations of a few gravitational radii, and when applied at separations of tens of thousands of gravitational radii. These validations make our formula a useful tool in a variety of contexts ranging from gravitational-wave physics to cosmology. As representative examples, we first illustrate how it can be used to decrease the phase error of the effective-one-body waveforms during the ringdown phase. Second, we show that, when combined with the recently computed self-force correction to the binding energy of nonspinning black-hole binaries, it provides an estimate of the energy emitted during the merger and ringdown. Finally, we use it to calculate the energy radiated in gravitational waves by massive black-hole binaries as a function of redshift, using different models for the seeds of the black-hole population.
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Black-hole (BH) mergers play a central role in gravitational-wave (GW) astrophysics, because they are expected to be among the main sources for existing and future detectors. More specifically, the LIGO/Virgo detectors~\citep{ligo,virgo} are expected to detect mergers of stellar-mass BHs happening within several hundred Mpc, when operating in their advanced configurations. Similarly, future space-based detectors such as LISA~\citep{lisa} or DECIGO~\citep{decigo} will detect mergers of massive BHs (MBHs) up to redshifts as high as $z\sim 10$ or beyond. Even intermediate-mass BHs (IMBHs), provided they exist, will be within reach of GW detectors, e.g. IMBH-MBH binaries will be detectable by LISA or DECIGO, while IMBH-IMBH binaries will be detectable with DECIGO or with the planned ground-based Einstein Telescope~\citep{ET1,ET2}. Given their relevance for GW astrophysics, it is not surprising that BH binaries have received widespread attention over the past few years. Because a detailed understanding of the dynamics of these systems is crucial in order to predict accurately the gravitational waveforms, which, in turn, is necessary to detect the signal and extract information on the physical parameters of the binaries, numerical simulations have been performed by a number of groups for a variety of mass ratios, BH spin magnitudes and orientations [see~\citet{pfeiffer2012} for a recent review]. However, even today, numerical-relativity simulations are computationally very expensive and not able to cover the full seven-dimensional space of parameters of quasi-circular BH binaries. Fortunately, phenomenological models have been very successful at reproducing many aspects of the dynamics of BH binaries as revealed by the numerical simulations. For instance, hybrid ``phenomenological waveforms''~\citep{ajith2008,phenom}, i.e., templates that represent phenomenological combinations of Post-Newtonian (PN) and numerical-relativity (NR) waveforms, can reproduce with high precision the NR waveforms for a wide range of binary parameters. Similar results are achieved by the effective-one-body (EOB) model, which attempts to reproduce not only the gravitational waveforms, but also the full dynamics of BH binaries during the inspiral, merger and ringdown phases, by resumming the PN dynamics~\citep{BD99,DamourResummedWfms}, and more recently the self-force dynamics~\citep{BBL12}. Other aspects of the dynamics of BH binaries have been phenomenologically understood by using combinations of PN theory, symmetry arguments, as well as hints from the test-particle limit and fits to numerical simulations. For instance, the final spin magnitude of the BH remnant can be predicted by a number of phenomenological formulas~\citep{rezzolla08a, rezzolla08b, tic08, BKL, kesden, rezzolla08c, spin_formula}, starting from the configuration of the binary either at small separations $r\lesssim 10 M$, or at large separations\footnote{For MBHs, the latter are roughly the separations at which the dynamics starts being dominated by GW emission, and represent therefore the separations at which these phenomenological formulas should work in order to be useful in cosmological contexts.} $r\sim 10^4 M$. These formulas also predict the orientation of the final spin with good accuracy when applied to small-separation binaries, while the formula of \citet{spin_formula} is also accurate when the binary has a large separation, e.g. $r\sim 10^4 M$, in a large portion of the parameter space~\citep{spin_formula,emanuele}. Similar phenomenological formulas have also been proposed for the recoil imparted to the final BH remnant from the anisotropic emission of GWs~\citep{her07b, koppitz2007, rezzolla08b, superkicks, gon07, kick_RIT1, lou09, kick_RIT3, kick_goddard1, kick_goddard2, kick_goddard3}. Because most of the anisotropic GW emission takes place as a result of the strongly nonlinear merger dynamics, these recoil formulas are not predictive, as they depend on quantities that can only be derived with full NR simulations, but they are still useful in the statistical studies usually performed in a cosmological context~\citep{mymodel,statistical_kick,lou12}. The dependence of the final mass of the BH remnant on the binary's initial parameters has also been investigated systematically in the literature~\citep{tic08,boy08,reisswig09,kesden,lou10}\footnote{An initial expression for the radiated energy was also suggested by~\citet{faithful_templates}, but was restricted to nonspinning binaries and based on early NR calculations.}, but the knowledge of this dependence is far less detailed. For instance, the formula of~\citet{tic08} [who built upon previous work by \citet{boy08}] is calibrated to reproduce NR results for comparable-mass binaries, but does not have the correct test-particle limit and is therefore inaccurate for binaries with small mass ratios. The formula of~\citet{kesden}, on the contrary, has the correct test-particle limit, but does not reproduce accurately the NR results for comparable-mass binaries. Finally, the formula of \citet{lou10} depends, for generic binary configurations, on quantities that can only be calculated using full NR simulations, and is therefore only useful in statistical studies. We here introduce a new phenomenological formula for the final mass of the BH remnant (Section \ref{sec:derivation}), which, by construction, reproduces both the test-particle limit and the regime of binaries with comparable masses and aligned or antialigned spins, which has been extensively investigated by NR calculations. In Section \ref{sec:comparison} we show that this novel formula reproduces accurately all of the available NR data (even for generic spin orientations and mass ratios), both when applied to small- and large-separation binary configurations. Furthermore, in Section \ref{sec:applications}, we consider three different areas where our formula can be useful: \textit{(i)} we show that it can help reduce the phase error of the EOB waveforms during the ringdown; \textit{(ii)} we combine it with the results of \citet{LBB12} for the self-force correction to the binding energy of nonspinning BH binaries and derive an estimate for the energy emitted during the merger and ringdown by nonspinning binaries; \textit{(iii)} using a semi-analytical galaxy-formation model to follow the coevolution of MBHs and their host galaxies, we use our formula to predict the energy emitted in GWs by MBH binaries as a function of redshift, and show that these predictions are strongly dependent on the model for the seeds of the MBH population at high redshifts. Our final conclusions are drawn in Section \ref{sec:conclusions}. Throughout this paper, geometrized units $G=c=1$ are used.
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\label{sec:conclusions} We have presented a novel algebraic formula to measure the energy radiated by coalescing binary systems of BHs with generic spin magnitudes and orientations and arbitrary mass ratios. Our expression uses information on the binary configuration at an arbitrary separation and reproduces correctly the two regimes in which the radiated energy is known best, namely, the test-particle limit (which is known analytically) and the comparable-mass case (which has been extensively investigated with NR simulations over the last few years). Because it smoothly interpolates these two regimes, we expect our formula to work reasonably well also for intermediate mass ratios. Indeed, we have verified that it reproduces the results of all the NR simulations published so far in the literature, including those with unequal masses, to within an error which is comparable to the typical errors of the simulations. In addition, we have checked that our formula works equally well when applied to binaries starting at small separations (i.e., $r\lesssim 10M$) and at large separations (i.e., $r\sim 10^4 M$), thus opening up the possibility of using our expression also in cosmological contexts. The algebraic nature of our expression makes it a useful tool in a variety of contexts that range from GW physics to cosmology. As representative examples we have discussed three different applications, namely: \textit{(i)} we have shown that, when combined with the results of \citet{LBB12} for the self-force correction to the binding energy of nonspinning BH binaries, the new formula provides an estimate for the energy emitted during the merger and ringdown, and that this estimate confirms the conjecture that the results of perturbative calculations may be successfully extrapolated to comparable-mass binaries when expressed in terms of the symmetric mass ratio $\nu$; \textit{(ii)} we have shown that the new formula can help reduce the phase error of the EOB waveforms during the ringdown; \textit{(iii)} using a semi-analytical galaxy-formation model to follow the coevolution of MBHs and their host galaxies, we have used our formula to predict the energy emitted in GWs by MBH binaries as a function of redshift, and found that these predictions strongly depend on the scenario adopted for the MBH seeds at high redshifts, thus making GW emission a powerful cosmological diagnostic. Additional uses of the new formula can be easily considered and a particularly relevant one is the impact of the mass loss on the accretion disk surrounding the MBH binary. The dynamics of the disk, in fact, can change considerably as a result of the very rapid change in the gravitational mass of the system, with the formation of large shocks, which are potentially detectable via their electromagnetic emission~\citep{Oneill2009,Rossi2010,Zanotti2010}. As a final remark we note that because our approach exploits knowledge derived from NR simulations, the accuracy of the final-mass formula can be improved as additional and more precise NR simulations, especially with highly-spinning BHs, become available.
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1206.3803
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1206
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1206.4060.txt
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We present the {\em Spitzer Extragalactic Representative Volume Survey} (SERVS), an $18\,\rm{deg}^2$ medium-deep survey at 3.6 and $4.5\, \mu$m with the post-cryogenic {\em Spitzer Space Telescope} to $\approx$\,2\,$\mu$Jy ($AB=23.1$) depth of five highly observed astronomical fields (ELAIS-N1, ELAIS-S1, Lockman Hole, Chandra Deep Field South and XMM-LSS). SERVS is designed to enable the study of galaxy evolution as a function of environment from $z\sim 5$ to the present day, and is the first extragalactic survey that is both large enough and deep enough to put rare objects such as luminous quasars and galaxy clusters at $z\stackrel{>}{_{\sim}}1$ into their cosmological context. SERVS is designed to overlap with several key surveys at optical, near- through far-infrared, submillimeter and radio wavelengths to provide an unprecedented view of the formation and evolution of massive galaxies. In this article, we discuss the SERVS survey design, the data processing flow from image reduction and mosaicking to catalogs, and coverage of ancillary data from other surveys in the SERVS fields. We also highlight a variety of early science results from the survey.
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\label{sec:intro} \setcounter{footnote}{0} Progress in extragalactic astronomy has been greatly enhanced by deep surveys such as the Great Observatories Origins Deep Survey (GOODS, \citealt{Dickinson+03}), the Cosmic Evolution Survey (COSMOS, \citealt{Sanders+07}), the Galaxy Mass Assembly ultradeep Spectroscopic Survey (GMASS, \citealt{Cimatti+08}), the {\em HST} Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey (CANDELS, \citealt{Grogin+11,Koekemoer+11}), that have allowed us to study the evolution of galaxies from the earliest cosmic epochs. However, a limitation of such surveys is the relatively small volumes probed, even at high redshifts: for example, \cite{Ilbert+06} find noticeable field-to-field variations in redshift distributions in the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS\footnote{www.cfht.hawaii.edu/Science/CFHTLS}) in fields of $0.7-0.9\,$deg$^2$. Until lately, the combination of depth and area required to map a large volume ($\sim 1\,$Gpc$^3$) of the high-redshift Universe at near-infrared wavelengths, where the redshifted emission from stars in distant galaxies peaks, has been prohibitively expensive in telescope time. Two recent developments have now made this regime accessible. On the ground, the availability of wide-field near-infrared cameras has greatly improved the effectiveness of ground-based near-infrared surveys in the $1 - 2.5\, \mu$m wavelength range. In space, the exhaustion of the cryogenic coolant of the {\em Spitzer Space Telescope} opened up an opportunity to pursue large near-IR surveys using the two shortest wavelength channels (IRAC1 [3.6] and IRAC2 [4.5]) of the Infrared Array Camera (IRAC, \citealt{Fazio+04}) in the post-cryogenic or ``warm'' mission that were much larger than was feasible during the cryogenic mission. The {\em Spitzer Extragalactic Representative Volume Survey} (SERVS), a {\em Spitzer} ``Exploration Science'' program, stems from these two developments. SERVS is designed to open up a medium-depth, medium-area part of parameter space in the near-infrared (see Figure~\ref{fig:areadepth}), covering 18\,deg$^2$ to $\approx 2\, \mu$Jy in the {\em Spitzer} [3.6] and [4.5] bands. These observations required 1400hr of telescope time and covered five highly observed astronomical fields: ELAIS-N1 (hereafter EN1), ELAIS-S1 (ES1), Lockman Hole (Lockman), Chandra Deep Field South (CDFS) and XMM-large-scale structure (XMM-LSS). The five SERVS fields are centered on or close to those of corresponding fields surveyed by the shallower {\em Spitzer} Wide-area Infrared Extragalactic Survey (SWIRE; \citealt{Lonsdale+03}), and overlap with several other major surveys covering wavelengths from the X-ray to the radio. Of particular importance is near-infrared data, as these allow accurate photometric redshifts to be obtained for high redshifts \citep{vanDokkum+06,BvDC08,Ilbert+09,Cardamone+10}: SERVS overlaps exactly with the 12\,deg$^2$ of the VISTA Deep Extragalactic Observations (VISTA VIDEO, Jarvis et al.\ 2012, in preparation) survey ($Z, Y, J, H$ and $K_s$ bands) in the south, and is covered by the UKIRT Infrared Deep Sky Survey (UKIDSS DXS, \citealt{Lawrence+07}) survey ($J$, $K$) in the north. SERVS also has good overlap with the {\em Herschel} Multitiered Extragalactic Survey (HerMES, \citealt{Oliver+12}) in the far-infrared, which covers the SWIRE and other {\em Spitzer} survey fields, with deeper subfields within many of the SERVS fields. Sampling a volume of $\sim 0.8$\,Gpc$^3$ from redshifts 1 to 5, the survey is large enough to contain significant numbers of rare objects, such as luminous quasars, ultraluminous infrared galaxies (ULIRGs), radio galaxies and galaxy clusters, while still being deep enough to find $L^*$ galaxies out to $z\approx 5$ (see for example \citealt{Falder+11}, and \citealt{Capak+11} who find two galaxies in the $z=5.3$ cluster bright enough to be detected by SERVS at 4.5$\mu$m.) For comparison, the largest structures seen in the Millennium simulation at $z\sim 1$ are of the order of $100$ Mpc \citep{Springel+05}, which subtends $3^{\circ}$ at that redshift, so each SERVS field samples a wide range of environments. By combining the five different fields of SERVS, the survey effectively averages over large-scale structure, and presents a representative picture of the average properties of galaxies in the high redshift Universe. {\em Spitzer} observations of the five SERVS fields are presented in detail in Section~\ref{sec:observations}. Image processing is detailed in Section~\ref{sec:processing}, focusing on the mosaicking process and uniformity of coverage. Section~\ref{sec:catalogs} presents the extracted SERVS catalogs, as well as an assessment of overall data quality, detection limits and expected number counts. Section~\ref{sec:ancillary} gives an overview of the ancillary data available at different wavelengths in the five fields. Preliminary science results and science goals are described in Section~\ref{sec:science_goals}. A summary of the SERVS data at hand is provided in Section~\ref{sec:summary}. \begin{figure*} \begin{minipage}[b]{1\linewidth} \centering \vspace{1.5cm} \begin{tabular}{cc} \hspace{-1cm}\includegraphics[trim=0cm 0.7cm 0cm 0cm,clip=true,scale=0.75]{f1.jpg} & \hspace{0.5cm}\includegraphics[trim=0cm 0.37cm 0cm 0cm,clip=true,scale=1.115]{f2.jpg}\\ \end{tabular} \end{minipage} \caption{Area versus depth for SERVS compared to other surveys at wavelengths of $\approx 3.6 \mu$m (\emph{left panel}) and $\approx 4.5 \mu$m (\emph{right panel}). For consistency, the depth shown is the $5\,\sigma$ limiting flux for point sources, excluding confusion noise ($\sigma_{\rm pp}$ as described in Section~\ref{subsec:data_analysis}), calculated from the {\em Spitzer} performance estimation tool (http://ssc.spitzer.caltech.edu/warmmission/propkit/pet/senspet /index.html) in each case. The surveys are ({\it from left to right}): GOODS, the {\em Spitzer} follow-up to the CANDELS {\em HST} survey (Cosmic Assembly Near-IR Deep Extragalactic Legacy Survey, \citealt{Grogin+11,Koekemoer+11}), the {\em Spitzer} Extragalactic Deep Survey (SEDS, Program identifier - hereafter PID - 60022, 61040, 61041, 61042, 61043, P.I.\ G.\ Fazio), the {\em Spitzer} IRAC/MUSYC Public Legacy in E-CDFS (SIMPLE) survey ({\em Spitzer}, PID 20708), the {\em Spitzer} Ultra Deep Survey (SpUDS, PID 40021, P.I.\ J.S.\ Dunlop), S-COSMOS, the {\em Spitzer} Deep Wide-Field Survey (SDWFS, \citealt{Ashby+09}), the {\em Spitzer}-HETDEX Exploratory Large Area (SHELA, PID 80100, P.I.\ C.\ Papovich) Survey, SWIRE, the SPT-{\em Spitzer} Deep Field (SSDF, PID 80096, P.I.\ S.\ Stanford) and the Wide-Field Infrared Explorer (WISE, \citealt{Wright+10}).} \label{fig:areadepth} %\vspace{0.2cm} \end{figure*} \vspace{0.5cm}
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\label{sec:summary} The \emph{Spitzer} Extragalactic Representative Survey (SERVS) is designed to open up a medium-depth, medium-area part of parameter space in the near-infrared, covering $18\,\rm{deg}^2$ to $\approx$\,2$\,\mu$Jy in the {\em Spitzer} [3.6] and [4.5] bands in five highly observed astronomical fields (EN1, ES1, Lockman Hole, CDFS and XMM-LSS). The five SERVS fields are centered on or close to those of corresponding fields surveyed by the shallower SWIRE fields, and they overlap with several other major surveys covering wavelengths from the X-ray to the radio. Of particular importance are near-infrared data, as these allow accurate photometric redshifts to be obtained. SERVS overlaps exactly with the 12 deg$^2$ of the VISTA VIDEO survey in the south, and is covered by the UKIDSS DXS survey in the north. SERVS also has good overlap with HerMES in the far-infrared, which covers SWIRE and other {\em Spitzer} survey fields, with deeper subfields within many of the SERVS fields. Sampling $\sim 0.8$\,Gpc$^3$ and redshifts from 1 to 5, the survey is large enough to contain significant numbers of rare objects, such as luminous quasars, ultra-luminous infrared galaxies, radio galaxies and galaxy clusters, while still being deep enough to find $L^*$ galaxies out to $z\approx 4$. In this article, we have described the \emph{Spitzer} observations, the data processing and the wealth of ancillary data available in the fields covered by SERVS. Mosaics and catalogs will be made available to the community in the summer of 2012 through the {\em Infrared Science Archive} (IRSA). \vspace{0.2cm}
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1206.4060
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1206
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1206.5422_arXiv.txt
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We present \Spitzer\ IRAC (2.1 sq. deg.) and MIPS (6.5 sq. deg.) observations of star formation in the Ophiuchus North molecular clouds. This fragmentary cloud complex lies on the edge of the Sco-Cen OB association, several degrees to the north of the well-known $\rho$~Oph star-forming region, at an approximate distance of 130~pc. The Ophiuchus~North clouds were mapped as part of the \Spitzer\ Gould Belt project under the working name `Scorpius'. In the regions mapped, selected to encompass all the cloud with visual extinction $A_V>3$, eleven Young Stellar Object (YSO) candidates are identified, eight from IRAC/MIPS colour-based selection and three from 2MASS~$K_S$/MIPS colours. Adding to one source previously identified in L43 \citep{chen09}, this increases the number of YSOcs identified in Oph~N to twelve. During the selection process, four colour-based YSO candidates were rejected as probable AGB stars and one as a known galaxy. The sources span the full range of YSOc classifications from Class~0/1 to Class~III, and starless cores are also present. Twelve high-extinction ($A_V>10$) cores are identified with a total mass of $\sim 100$~\Msun. These results confirm that there is little ongoing star formation in this region (instantaneous star formation efficiency $<0.34$\%) and that the bottleneck lies in the formation of dense cores. The influence of the nearby Upper~Sco~OB association, including the 09V~star $\zeta$~Oph, is seen in dynamical interactions and raised dust temperatures but has not enhanced levels of star formation in Ophiuchus~North.
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\label{sect:introduction} The Ophiuchus~North (Oph~N) molecular clouds lie $20^\circ$ above the Galactic plane in the direction of the Galactic Centre. They are part of the same filamentary cloud complex as the well-studied Ophiuchus~L1688 and L1689 clouds, but lie several degrees to the north, on the boundary of the constellation of Ophiuchus with Scorpius. Like the Ophiuchus~~L1688 and L1689 clouds, they are illuminated from the northwest by the Upper~Scorpius subgroup of the Sco-Cen OB association. The region has been studied little. Early low-resolution CO mapping \citep{degeus90,degeus92} showed the filamentary structure of the clouds in the Ophiuchus % region (mirrored in the extinction maps published by \citealt{dobashi05} and \citealt{rowlesfroebrich09}), and suggested a shock origin due to expanding shells surrounding the Upper~Sco subgroup. A detailed study of the molecular clouds was made in $^{13}$CO by \citet{nozawa91} which gives an excellent overview of the cloud complex, and its relationship to the Ophiuchus cores, IRAS sources, and the Sco-Cen OB association. They find that the region contains some 23 $^{13}$CO clouds containing 51 $^{13}$CO cores, with a total mass of 4000\Msun\ and typical core densities of $N_{{\mathrm H}_2} \sim 3\times 10^3\hbox{ cm}^{-3}$. The dense cores ($N_{\mathrm{H}_2} \sim 10^4\hbox{ cm}^{-3}$) and velocity structure were subsequently followed up in C$^{18}$O by \citet{tachihara00b,tachihara00a,tachihara02}. \citet{nozawa91} identify only thirteen YSOs associated with the cores, pointing to a low star formation efficiency of 0.3\%. Distance estimates for the Oph~N molecular cloud complex come from its relationship with the stars in Upper Sco (US) as, from extinction, the molecular clouds lie in front of and distributed through the OB population \citep{degeus89}. Hipparcos parallaxes place Upper Sco at a mean distance of $145\pm2$~pc \citep{dezeeuw99}, with a line-of-sight extent of $\pm 17$~pc assuming the $14^\circ$ spatial extent is reproduced in the third dimension. This places an effective upper limit on the clouds' distances of 162~pc. Extinction-based distance modulus estimates suggest the clouds are distributed between 80~pc (near side) and 170~pc (far side), centre 125~pc \citep{degeus89} or, slightly further away, 120~pc (near side) to 200~pc (far side) \citep{straizys84}. These distances are consistent with the Hipparcos data. Looking at individual stars, \citet{degeus90} suggest that the western clouds (our OphN~4,5,6 their Complex~2) lie in front of $\chi$~Oph (150~pc) and the northeastern (OphN~1/2, Complex~4) in front of $\xi$~Oph (130~pc). Recent estimates for the better-studied L1688 Ophiuchus cloud range from 120--145~pc \citep{wilking08}. There is certainly no reason to believe that the Ophiuchus clouds all lie at the same distance, but for convenience we assume a working distance of 130~pc, which is consistent with the above estimates. In this paper, we present mid-infrared, \Spitzer\ Space Telescope observations of the high column density regions of Ophiuchus~North. The observations and data reduction are described in Sect.~\ref{sect:observations}. Results, including source statistics, YSO candidates, extinction maps and large-scale emission, are given in Sect.~\ref{sect:results} with comments on individual regions in Appendix~\ref{sect:regions}. The results are discussed in Sect.~\ref{sect:discussion} and summarised in Sect.~\ref{sect:summary}. \section[]{Description of Observations} \label{sect:observations} We observed Ophiuchus~North in the mid-infrared as part of the \Spitzer\ legacy program ``Gould's Belt: star formation in the solar neighbourhood'' (\SpitzerGB). The clouds were mapped under the working name `Scorpius' (Sco) and appear under this name in the \SpitzerGB\ catalogues. We present them here as the Ophiuchus~North (OphN) clouds in line with previous nomenclature \citep{nozawa91,tachihara00a, tachihara00b,tachihara02} which reflects their location predominantly within the constellation of Ophiuchus. Only our regions OphN~5,6 and LDN~43 lie beyond the constellation boundary in Scorpius. The \SpitzerGB\ program aimed to complete the mapping of local star formation started by the \Spitzer\ ``From Molecular Cores to Planet-forming Disks'' (c2d) project \citep{c2d,evans09} by targetting the regions IC5146, CrA, Scorpius (renamed Ophiuchus~North), Lupus II/V/VI, Auriga, Cepheus Flare, Aquila (including Serpens South), Musca, and Chameleon to the same sensitivity and using the same reduction pipeline \citep{gutermuth08,harvey08,kirk09,peterson11,spezzi11}. Images were made at 3.6/4.5/5.8/8.0\micron\ with the Infrared Array Camera (IRAC; \citealt{fazio04}) and 24,70 and 160\micron\ with the Multiband Imaging Photometer for \Spitzer\ (MIPS; \citealt{rieke04}). With an 85cm mirror, IRAC observes with an angular resolution of $2''$ whereas MIPS is diffraction limited with $6''$, $18''$ and $40''$ resolution at 24, 70 and 160\micron\ respectively. For our observations, we targetted small regions encompassing the $A_V>3$ contours from the optical extinction map of \citet{dobashi05}, as shown in Fig.~\ref{fig:coverage}. The area in Oph~N which lies above $A_V>3$ is fragmentary and scattered over an area of 20 sq. deg. Some of these regions (L158,L204,L146/CB68,L234E,L260,L43) had already been mapped as part of the \Spitzer\ ``Cores to Disks'' project \citep{c2d}. These were avoided by the \Spitzer\ Gould Belt team to avoid unnecessary duplication of observations, as the two projects work to the same target sensitivities. Most of these c2d data are incorporated in this study of Oph~N. The exception is L43 which is presented separately by \cite{chen09}. Ultimately, seven new areas were mapped by \SpitzerGB\ with IRAC and MIPS. The \citet{dobashi05} and \citet{rowlesfroebrich09} extinction maps, which are not limited to the IRAC observations but extend across the entire area covered by MIPS, confirm that all $A_V > 3$ \citep[measuring from ][]{dobashi05} or $A_V > 4$ \citep{rowlesfroebrich09} regions in these filaments were observed by IRAC with the exception of two small clouds $0.3^\circ$ to the south of OphN~6 (included in the MIPS map) and $0.4^\circ$ to the north of CB68 \citep[][ core q2]{tachihara00a}. Details of the datasets included for each region are listed in Table~\ref{tbl:aors}, including the observation dates, Astronomical Observation Request (AOR) identification numbers, program identification (30574 for \Spitzer\ Gould Belt, 139 for ``Cores to Disks''), and duration. Table~\ref{tbl:assc} gives the associated Lynds Dark Clouds \citep{lynds62} and molecular cores \citep[C$^{18}$O, ][]{tachihara00b} for each region. The overlap area covered in all four IRAC bands is slightly smaller than the area covered in any single band because the array design leads to an offset between the 3.6/5.8\micron\ and 4.5/8.0\micron\ maps. The final area covered by MIPS is much larger than with IRAC because of the long MIPS scans, with almost all the IRAC areas covered by MIPS as shown in Fig.~\ref{fig:coverage}. In total, 2.1 square degrees were covered by IRAC and 6.5 with MIPS. This is roughly half the area of 14.4~sq.deg. covered by MIPS in the main Ophiuchus clouds \citep{padgett08}. The observations of each area were split between two epochs to allow removal of transient objects, with the second epoch maps offset from the first epoch maps so that bad pixels do not always lie in the same sky positions. The IRAC observations were taken with a total integration time of 48 seconds per point, split equally between the two epochs, and an offset of $10''$ in array coordinates between epochs. Short integrations in High Dynamic Range (HDR) mode were also taken for all regions in which bright YSOs were expected (the exception was L234E). The \SpitzerGB\ MIPS observations were taken in fast-scanning mode, stepping by $240''$ cross-scan to fill gaps in the coverage, and ($125''$,$80''$) between the two epochs in order to provide full 70\micron\ coverage with only half the array working. MIPS total integration times were 32.4s at 24 and 70\micron\ and 6.2 seconds at 160\micron. % The ``Cores to Disks'' observations of the small cores L234E and CB68 were taken in MIPS photometry mode. Each core was observed in 2 epochs. At 24\micron\, 1 cycle of 3 seconds was taken at each epoch for a total integration time of 84s. At 70\micron\, 3 cycles of 3 seconds were taken at each epoch for a total integration time of $\sim 94$s. \subsection{Data reduction} Data reduction was carried out using the c2d pipeline as described in the delivery documentation (\citealt{c2ddel}; see also \citealt{harvey06} and \citealt{rebull07} for details of IRAC and MIPS processing). The Basic Calibrated Data (BCD) from the \Spitzer\ Science Center were processed by the \SpitzerGB\ team to remove artifacts ( eg. bad pixels, jailbar effects, muxbleed column pulldown due to bright sources) and apply the location-dependent photometric corrections. \subsubsection{Mosaics} Mosaicking was carried out using the SSC ``Mopex'' code. For IRAC mosaics, the high dynamic range (HDR) data were included in the final map, which improves the dynamic range and allows for the inclusion of otherwise saturated sources at the expense of slightly increased noise levels. For both IRAC and MIPS, both epochs were combined. The \SpitzerGB\ Oph~N data were reduced to form six separate maps, numbered OphN~1--OphN~6 in order of decreasing Galactic longitude. These were originally mapped as Sco~1--6 but relabelled `OphN' in line with the overall cloud nomenclature change. OphN~1,4,5 and 6 each incorporated a single \SpitzerGB\ AOR area. For OphN~2, the c2d cloud L204C was mosaicked with the \SpitzerGB\ data. For OphN~3, the MIPS observations for two of the \SpitzerGB\ regions and the overlapping c2d region L158 were merged into a single mosaic. The IRAC observations for OphN~3 remain split into OphN~3a (including L158) and OphN~3b. The two separate c2d cores L234E and CB68 were mosaicked individually. These pairings are indicated in Table~\ref{tbl:aors}. A finding chart showing the six \SpitzerGB\ regions and the two complementary c2d areas overlaid on the \citet{dobashi05} extinction map is shown in Fig.~\ref{fig:coverage}. \subsubsection{Catalogues} \label{sect:catalogues} Source extraction was carried out in each of the 4 IRAC bands on the combined mosaics (2 epoch plus HDR), plus MIPS 24\micron\ on the combined epoch mosaics, using c2dphot, a photometry tool based on DoPHOT \citep{harvey06, schechter93}. Sources believed to be real at one or more wavelengths were band-filled at the missing wavelengths, first in IRAC bands 1--4 then MIPS~24\micron\ then MIPS~70\micron\ by fitting a point spread function (PSF) to images at the missing wavelengths, and these fluxes are included in the catalogues with bandfilling noted in the data quality flags. Fluxes were extracted from the 2MASS point source catalogue \citep{skrutskie06} using a $2''$ position matching criterion. The 70\micron\ source list was separately matched to the shorter wavelength catalogue with an $8''$ position matching criterion. With a larger pixel size of $10\arcsec$ at 70\micron\ (compared to $1.2\arcsec$ for IRAC and $2.55\arcsec$ at 24\micron) sources often matched more than one shorter-wavelength source. The 70\micron\ fluxes were assigned to the best candidate by hand matching the spectral energy distribution (SED), with the remaining possible 70\micron\ emitters also noted in the catalogue. Source matching was not attempted at 160\micron\ as the spatial resolution is low at this wavelength ($\sim 40''$) and most bright regions are saturated. For further details on source extraction see the c2d delivery documentation \citep{c2ddel}. \subsubsection{160 micron data} The 160\micron\ data obtained in fast scan mode mapping does not have enough redundancy to fill in completely all the gaps due to a dead readout and the intermediate gap between the array detector rows \citep{mipshandbook}. Furthermore with only 2 epochs, effectively 2 pointings per pixel, the effects of hard radiation hits and saturation translates into small regions without data. To deal with these gaps and to preserve as much as possible of the diffuse emission, the images are first resampled from 15\arcsec\ to 8\arcsec\ pixels and then a 5~pixel by 5~pixel median filter is applied to the image. The net effect is a slight redistribution of surface brightness ($\sim 15\%$) and smearing of the original beam from 40\arcsec\ to about 1\arcmin\ in size. 160\micron\ data are not available for the two c2d cores CB68 and L234E. Fluxes at 160\micron\ are determined using aperture photometry with a 32\arcsec\ radius aperture, an annulus from 64-128\arcsec\, and an aperture correction of 1.97. Flux density uncertainties are about 20\% below 5~Jy and 30\% for higher flux densities \citep{rebull10}. Visual inspection of each 160\micron\ source in the original unsmoothed image is used to determine whether the object is cleanly detected or contaminated by data dropouts.
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\label{sect:summary} We have surveyed the high extinction regions ($A_V > 3$) of the Oph~N molecular cloud complex in the mid-infrared with \Spitzer\ IRAC and MIPS. There is little active star formation in Oph~N. Twelve YSOcs are identified in total. Eight of these are identified from their IRAC-MIPS colours. An additional Class~I protostar exists in the L43 cloud in the centre of this region \citep{chen09}. Three further YSO candidates are found based on red $[K-24]$ colours. Four sources initially selected as YSO candidates were rejected as AGB stars and one as a galaxy. Of the twelve remaining YSOcs, three are Class~I, one flat-spectrum, seven Class~II and one Class~III. All except one Class~III and two Class~II sources were previously known. The region as a whole has a low star formation efficiency of $<0.34\%$. The high extinction regions in Oph~N are fragmented into twelve small ($\sim 0.2$~pc), scattered and low-mass ($\sim 10\Msun$ or less) cores, of the kind which might form one or two low-mass stars. Three of these (OphN~1, CB68 and L43) currently contain Class~I sources. The remainder are starless. The interstellar medium in this region is influenced by the Upper~Sco subgroup of the Sco-Cen OB association. There is evidence for dynamical interaction in the OphN~4 and OphN~5 bubbles, elevated temperatures in OphN~5 and 6, and irradiated cloud edges throughout the region. The bulk of the gas mass is at low column density with only a few low-mass cores surviving to form a few YSOs. This is very different from the situation in nearby Ophiuchus~L1688, which contains hundreds of YSOs \citep{padgett08}, but similar to the Lupus clouds \citep{tachihara96,merin08}. As the impact of radiation and winds from Upper~Sco must have been similar in both regions, the initial conditions must have been very different. Whereas the main Ophiuchus complex contained hundreds of solar masses of dense gas, Oph~N only contained a few low-mass high-density cores. It seems likely that the small star-forming molecular cores in Oph~N are the surviving remnants of the GMC which produced Upper~Sco and are now slowly being ablated by the massive stars. The low star formation rate is due to the lack of dense cores.
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1206.5422
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1206
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1206.5939_arXiv.txt
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{} {The \hii\ regions LMC N191 and SMC N77 are among the outermost massive star-forming regions in the Magellanic Clouds. So far, few works have dealt with these objects despite their interesting characteristics. We aim at studying various physical properties of these objects regarding their morphology (in the optical and Spitzer IRAC wavelengths), ionized gas emission, nebular chemical abundances, exciting sources, stellar content, age, presence or absence of young stellar objects, etc. } {This study is based mainly on optical ESO NTT observations, both imaging and spectroscopy, coupled with other archive data, notably Spitzer images (IRAC 3.6, 4.5, 5.8, and 8.0\,$\mu$m) and 2MASS observations. } { We show the presence of two compact \hii\ regions, a low-excitation blob (LEB) named LMC N191A and a high-excitation blob (HEB) named SMC N77A, and study their properties and those of their exciting massive stars as far as spectral type and mass are concerned. We also analyze the environmental stellar populations and determine their evolutionary stages. Based on Spitzer IRAC data, we characterize the YSO candidates detected in the direction of these regions. Massive star formation is going on in these young regions with protostars of mass \ab\,10 and 20 \sm\ in the process of formation. } {}
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The compact \hii\ regions residing in the Magellanic Clouds are interesting in the context of massive star formation in these neighboring galaxies. Typical Magellanic Cloud \hii\ regions are giant complexes of ionized gas with sizes of several arc minutes, corresponding to physical scales of more than 50\,pc and are powered by a large number of exciting stars. In contrast, Magellanic Cloud compact \hii\ regions are small regions mostly $\sim$\,5\frac\, to 10\frac\, in diameter, corresponding to $\sim$\,1.5 to 3.0\,pc and excited by a much smaller number of massive stars. There are two types of compact \hii\ regions, high-excitation blobs \citep[HEBs, for a review see][]{MHM10b} and low-excitation blobs \citep[LEBs,][]{Meynadier07}. The members of the first group are often observed lying adjacent or projected onto giant \hii\ regions and are younger than the associated giant \hii\ regions. Do HEBs indeed belong to the same region of the Magellanic Clouds at which the giant \hii\ regions have formed or is the association between these two types of \hii\ regions a line-of-sight effect? If they are associated, are HEBs powered by triggered, second-generation massive stars? Why has star formation not proceeded in a single burst although massive stars are believed to form in the dense core of giant molecular clouds? These are some interesting questions, the answers to which will be helpful for better understanding massive star formation in the Magellanic Clouds. A problem is that HEBs are not numerous, and moreover, few of them have been studied individually in detail. \\ This paper is devoted to a first detailed study of two compact \hii\ regions, one in the Large Magellanic Cloud (LMC) \hii\ region N191 and the other in the Small Magellanic Cloud (SMC) N77 \citep[][]{Henize56}. Among the LMC \hii\ regions listed by \citet[][]{Henize56}, N191 is one of the outermost, lying below the bar, at a distance of \ab\ 200\min\ (\ab\ 3 kpc in projection) from the famous 30 Doradus. N191 appears as an elongated structure, with two components N191A and N191B in the Henize catalog. Here we essentially study the brightest component N191A, also known as DEM L 64b \citep[][]{Davies76}. N77 is one of the most northern \hii\ regions of the SMC; it is situated at a distance of \ab\ 25\min\ (\ab\ 440 pc in projection) from the pre-eminent SMC \hii\ region N66 \citep[][and references therein]{MHM10a}. SMC N77 is identified in the optical survey of \citet[][]{Davies76} as DEM S 117. \\ Few works have been devoted to these two \hii\ regions despite their interesting characteristics. LMC N191 belongs to the OB association LH 23 \citep[][]{Lucke70}. It was also detected as IRAS source 05051-7058 \citep[][]{Helou88}. The compact \hii\ region SMC N77 seems to coincide with the stellar association B-OB 24 \citep[]{Battinelli91}. It was identified in the infrared as IRAS source 01011-7209 \citep[][]{Helou88} and as source \#48 in the ISO 12\,$\mu$m catalog \citep[][]{Wilke03}. Furthermore, LMC N191 and SMC N77 were part of a Spitzer study of compact \hii\ regions by \citet[][]{Charmandaris08} and have been included in several radio continuum surveys of the Magellanic Clouds \citep[][]{Filipovic95,Filipovic02}. Both compact \hii\ regions are associated with molecular clouds. The giant molecular cloud LMC N J0504-7056 is centered at 130\frac\ south of N191 \citep[][]{Fukui08}. Moreover, the OB association LH 23 and the \hii\ region are related to this molecular cloud \citep[][]{Fukui08,Kawamura09}. A small molecular cloud has been detected near the position of SMC N77 \citep[][]{Mizuno01}. \\ This paper is arranged as follows. Section 2 presents the observations, data reduction, and the archive data (Spitzer data, 2MASS data). Section 3 describes our results (overall view, extinction, nebular emission, stellar content and chemical abundances). Section 4 presents our discussion, and finally our conclusions are summarized in Section 5.
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This paper presented the first detailed study of LMC N191A and SMC N77A using imaging and spectroscopy in the optical obtained at the ESO NTT as well as Spitzer and 2MASS data archives. The two objects are among the outermost star-forming regions of the Magellanic Clouds. We derived several physical characteristics of these regions and their powering sources. The compact \hii\ region N191A, \ab\,10\frac\ (2.4 pc) in diameter, belongs to a small class of ``low-excitation blobs'' in the Magellanic Clouds. In contrast, SMC N77A, \ab\,20\frac\ (5.8 pc) in size, belongs to the ``high-excitation blob'' family. The class of compact \hii\ regions in the Magellanic Clouds is not very populated. Therefore new members provide additional data for improving our knowledge of their characteristics and their formation processes. Higher resolution observations are necessary to deepen the study of these objects.
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1206.5939
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1206
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1206.0031_arXiv.txt
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{ % To understand the feedback of black holes on their environment or the acceleration of ultra-high energy cosmic rays in the present cosmic epoch, a comprehensive inventory of radio galaxies in the local universe is needed. This requires an all-sky catalog of radio-emitting galaxies, that hitherto has not been available.} { We present such an all-sky sample. Our catalog allows one to build volume-limited subsamples containing all low-power radio galaxies, similar to the prototypical low-power radio galaxies Cen~A or M87, within some hundred Mpc.} { We match radio emission from the NVSS and SUMSS surveys to galaxies of the 2MASS Redshift Survey (2MRS) using an image-level algorithm that properly treats the extended structure of radio sources. } { The bright master sample we present contains 575 radio-emitting galaxies with a flux greater than 213~mJy at 1.4~GHz. Over 30\% of the galaxies in our catalog are not contained in existing large-area extra-galactic radio samples. We compute the optical and radio luminosity functions and the fraction of radio galaxies as a function of galaxy luminosity. 94\% of the radio galaxies within $z=0.03$ are of Hubble type E/S0. The local galaxy density in a sphere of 2~Mpc centered on the radio galaxies is 1.7 times higher than around non-radio galaxies of the same luminosity and morphology, which is a statistically significant enhancement ($>3\sigma$).} { Our sample presents the deepest all-sky catalog of low-power radio galaxies. The observed enhancement of the galaxy density around radio galaxies suggests a causal relation between external galaxy properties, such as environment or merger history, and the formation of powerful jets in the present universe. Since the enhancement is observed with respect to galaxies of the same luminosity and Hubble type, it is not primarily driven by black hole mass. Our automated matching procedure is found to select radio-emitting galaxies with high efficiency (99\%) and purity (91\%), which is key for future processing of deeper, larger samples.}
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1206.0031
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1206
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1206.5570_arXiv.txt
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We examine the properties and evolution of a simulated polar disc galaxy. This galaxy is comprised of two orthogonal discs, one of which contains old stars (old stellar disc), and the other, containing both younger stars and the cold gas (polar disc) of the galaxy. By exploring the shape of the inner region of the dark matter halo, we are able to confirm that the halo shape is a oblate ellipsoid flattened in the direction of the polar disc. We also note that there is a twist in the shape profile, where the innermost 3~kpc of the halo flattens in the direction perpendicular to the old disc, and then aligns with the polar disc out until the virial radius. This result is then compared to the halo shape inferred from the circular velocities of the two discs. We also use the temporal information of the simulation to track the system's evolution, and identify the processes which give rise to this unusual galaxy type. We confirm the proposal that the polar disc galaxy is the result of the last major merger, where the angular moment of the interaction is orthogonal to the angle of the infalling gas. This merger is followed by the resumption of coherent gas infall. We emphasis that the disc is rapidly restored after the major merger and that after this event the galaxy begins to tilt. A significant proportion of the infalling gas comes from filaments. This infalling gas from the filament gives the gas its angular momentum, and, in the case of the polar disc galaxy, the direction of the gas filament does not change before or after the last major merger.
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Classical disc galaxies consist of flattened distributions of stars, gas, and dust, co-rotating with the angular momentum vectors more or less aligned;\footnote{In the presence of warps though, there may be some misalignment driven by external torques \citep{Roskar2010}.} polar disc galaxies\footnote{Originally called `multi-spin galaxies' by \citet{Rubin1994}.} are significantly more rare, possessing two orthogonal discs. An original survey of galaxies consisting of orthogonal structures of stars was carried out by \citet{Whitmore1990} who identified a sample of polar {\it ring} galaxies. The `primary' discs in such systems are typically classified as lenticulars (S0), or Ellipticals, with a second, younger, structure aligned with angular momentum axis of the primary disc. Fewer than 1\% of S0 galaxies contain a polar structure. Subsequent studies of NGC4650A revealed a subclass of structures where the orthogonal distribution of stars was not a {\it ring} but a second fully-formed continuous {\it disc} structure, \citep{Iodice2002, Gallagher2002,Swaters2003, Maccio2006, Brook2008}. Systematic studies suggest that their stellar populations, light distribution, and gas-phase characteristics, are not dissimilar from those of classical discs - e.g., exponential light profiles \citep{Schweizer1983} and \textsc{HI} mass-to-optical luminosities typical of mass to B-band luminosity typical of late-type spirals \citep{Huchtmeier1997, Arnaboldi1997, Sparke2000, Cox2006}. Thus there exist {\it polar ring} galaxies as well as {\it polar disc} galaxies. We will discuss the latter type in the following paper. On of the important uses of polar disc galaxies is a direct result of having orthogonal discs at the bottom of a halo potential well. One stellar disc historically allowed astronomers to calculate the existence of dark matter halos, (see \citet{Begeman1991} for a detailed discussion). \citet{Iodice2003} note that the existence of two orthogonal discs means that the shape of the inner region of the halo can be recovered from observations. There are several theories proposed to explain the origin of polar ring and polar disc galaxies: \begin{enumerate} \item {\bf Mergers} : As envisioned by \citet{Bekki1997, Bekki1998} and \citet{Bournaud2003}, polar rings can eventuate for very specific collision configurations where the collision is `head-on' and the initial angular momentum low. \item {\bf Tidal Accretion} : Proposed by \citet{Schweizer1983} and simulated by \citet{Reshetnikov1997}, this scenario involves the capture of a gas-rich secondary by the (future) polar disc host halo. The interaction-induced polar ring is formed more readily than in the aforementioned merger-induced scenario \citep{Bournaud2003, Combes2006}. \item {\bf Infall from Filaments} : If gas falls into a galaxy along cosmic filaments that are inclined to the stellar disc, a polar disc may form \citep{Maccio2006} \& \citep{Brook2008}. \citet{Combes2006} and \citet{Bournaud2003} suggest such cosmological infall is a viable alternative to tidal accretion for the NGC~4650A system e.g. \citep{Iodice2006, Spavone2010, Spavone2011}. \item {\bf Resonance} : This approach assumes that the polar disc is formed via resonant coupling to a triaxial halo potential \citep{Tremaine2000}. In this picture, a disc is taken to lie in the plane of symmetry of the triaxial halo and the halo tumbles with respect to the disc. As the tumbling slows, the stars in the disc can get trapped in resonance with the halo. A stellar orbit inclined to the disc precesses at a slow retrograde rate and the star can be pushed into a polar orbit. This method has been invoked to best explain equal mass stellar disc systems. \end{enumerate} A seminal review of the field is provided by \citet{Combes2006}, to which all interested readers should refer for a rich survey of the field. The filamentary infall scenario (iii, above) of \citet{Maccio2006} has gathered momentum over the past few years as being the favoured origin of polar disc systems. This can be traced, in part, to the fact that the alternatives do not account fully for the range of observations -- in particular, those which show that the polar disc is often of high mass, possesses extended and continuous structure, and lacks obvious evidence of starbursts that might be associated with polar structure formation. Scenario (i) can form a polar disc/ring galaxies when the relative initial velocities are low (e.g. \citet{Bournaud2003} ) and by (ii) in loose groups of galaxies such as UGC9796 \citep{Spavone2011}. Polar disc/ring galaxies are found in low density environments as the polar structure is destroyed by mergers \citep{Maccio2006}. This suggests that each type of formation may exist in the universe. Closely related to \citet{Maccio2006} scenario (iii, above), \citet{Brook2008} suggest that the polar disc forms as a result of the last major merger changing the angular momentum of the stars and gas of the stellar disc, and then subsequent gas continues to fall in along the old trajectory and builds up the polar disc.\footnote{Alternatively, the gas falling along filaments could change direction.} The basic observable characteristics of their simulated polar disc system were presented, including star formation rates, circular velocity profiles, and structural parameters of the discs, along with the aforementioned putative origin scenario. The exact mechanism by which the dark halo aligns to the orthogonal discs and how the new infalling gas becomes misaligned to the orthogonal discs was {\it not} \rm studied by \citet{Brook2008}. In what follows, we rectify this by examining the underlying physics governing the misalignment of the discs in the \citet{Brook2008} simulation, tracing its temporal evolution and association with large-scale structure, in order to put their origin scenario to the test, quantitatively. The simulation and its basic properties are reviewed in \S~2 and \S~3, while the time evolution of the host halo's shape is detailed in \S~4. The metallicity and age gradients of the discs are then confronted with recent observational work (\S~5), while the time evolution of the dark halo, gas, and stellar angular momenta is derived in \S~6. Our conclusions and suggestions for future directions are made in \S~7.
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We have extended the preliminary analysis of the polar disc structure first introduced by \citet{Brook2008}. We have examined the properties of the dark matter halo shape, the formation process of this object and the metallicity and stellar age gradients of the discs. \citet{Iodice2003} point out that polar disc galaxies provide an unparalleled chance to study the shape of the inner region dark matter haloes using observations in a way normally only possible in simulations. The two orthogonal discs have a circular velocity defined by the potential, which is dominated by dark matter. By comparing the circular velocity of orthogonal discs we estimate the flattening of the dark matter in the plane where both discs are seen edge on and find that the halo is flattened in the direction of the polar disc, as found by Iodice et al. Specifically the axial ratio is found to be c/a=0.9. Directly measuring the dark matter halo in the simulation we find the same result and compare find the axial ratios are extremely similar using both methods, i.e. c/a=0.93. The value quoted is for the inner region of the halo, at the radius of the polar disc itself. In the outer region of the disc the sphericity declines to 0.67. We concentrated on the inner region in this work because that is the region probed by the baryons. This value is closely matched by simulations of dark matter haloes, such as the galaxy size halos with cooling simulated by \citet{Kazantzidis2004} and shown in their figure 2, which shows an almost spheroidal inner region, with a sphericity of 0.9 at 10\% if the virial radius, and an outer region of c/a=0.6, confirming the work of \citet{Springel2004} and \citet{Allgood2006} amongst others. \citet{Allgood2006} has categorised halo shapes from cosmological simulations using statistical samples. Using their results we expect our halo to have a c/a ratio of approximately 0.65 at z=0.17 which is consistent with the polar disc galaxy halo; the value of c/a for the `classical' disc shown in Fig. \ref{Fig:darkshapedisc} at the viral radius is 0.7. The shape profile, both in angle and sphericity remain stable from outside the inner region to the virial radius. This differs from the work of \citet{Bailin2005} and \citet{Allgood2006} who find a change of angle and shape with radius. This feature is unusual but polar discs are unusual objects and this may be characteristic. \citet{Maccio2006} did not see the same level of coherence in their polar disc galaxy, however. Compared to observational estimates, e.g. \citet{Sackett2000}, however, these halos appear very spherical. Studies of polar ring/disc galaxies (see \citet{Sackett1994}, \citet{CombesA1996}, \citet{Iodice2003}) suggest a highly flattened halo of between 0.3 and 0.4. \citet{Sackett2000} does, however, provide some suggestions as to why this difference occurs. Estimates of c/a for halos based on different observations vary wildly depending on the chosen metric, such that $c/a = 0.5 \pm 0.2$. In the outer region our halo falls within this range, 0.65, while in the inner region ($R<15$ kpc) where the observational estimates probe we see $c/a$=0.9 to 0.75. The simulation outcome is also close to the result of \citet{Hoekstra2004} based on weak lensing who estimate a halo sphericity of $0.67^{+0.09}_{-0.07}$, the upper range of which is close to the sphericity of our halo at a radius of 14 kpc. As for the extremely elliptical shape found for polar ring/disc galaxies, the methods used may simply not give good estimates, indeed, \citet{Sackett2000} notes the difficulty in measuring halo sphericity observationally. We have concluded that the polar disc of this system is not caused directly by a merger or stripping. There is no single merger or stripped satellite which contains enough gas to form the massive polar disc. By process of elimination, the polar disc must form from inflowing cold gas, as in \citet{Maccio2006}. We attempted to identify the mechanism by which this cold gas flows into the system and gives rise to the polar structure, asking, why the filaments should change direction relative to the old stellar disc. The structure of the polar disc seems to be a result of the direction of gas infall from the filament. We have shown that there is not a single moment when the polar disc `forms' but, instead, there is a constant evolution of the angular momentum of the infalling gas and stars. The angular momentum of the last major merger is almost exactly orthogonal to the filament and we feel that this is the most likely origin of the polar disc structure. In this scenario, put forwards by \citet{Brook2008}, the major merger reorientates the angular momentum of the old stars and the cold gas already present in the disc. Gas continues to flow in from the large scale structure unaffected by this merger and so builds up the polar disc. We see that the major merger affects the stars most strongly without any significant influence on the dark halo, and the cold gas recovers its old angular momentum within a Gyr. However, this is not the whole story as subsequent to the major merger there is further evolution in the orientation of the entire galaxy system, potentially due to tumbling or torques from the halo. The growing alignment of the polar disc with the intermediate axis of the dark halo does not seem affected by the major merger. We also provide a basic comparison of the polar disc galaxy to a classical disc galaxy to better isolate the origin of the extreme behaviour of the polar disc. We also confirm that the polar disc exhibits the characteristics of a disc built up from the inside out, both from the metallicity gradient and the stellar age profile. Clearly though, any conclusions for their formation, when based upon a single simulated system, must be taken as tentative, at best. We need to (and will) repeat this analysis on a statistically significant sample of simulated polar disc systems.
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1206.5570
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1206
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1206.7096_arXiv.txt
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Many astrophysical systems of interest, including protoplanetary accretion disks, are made of turbulent magnetized gas with near solar metallicity. Thermal ionization of alkali metals in such gas exceeds non-thermal ionization when temperatures climb above roughly 1000~K. As a result, the conductivity, proportional to the ionization fraction, gains a strong, positive dependence on temperature. In this paper, we demonstrate that this relation between the temperature and the conductivity triggers an exponential instability that acts similarly to an electrical short, where the increased conductivity concentrates the current and locally increases the Ohmic heating. This contrasts with the resistivity increase expected in an ideal magnetic reconnection region. The instability acts to focus narrow current sheets into even narrower sheets with far higher currents and temparatures. We lay out the basic principles of this behavior in this paper using protoplanetary disks as our example host system, motivated by observations of chondritic meteorites and their ancestors, dust grains in protoplanetary disks, that reveal the existence of strong, frequent heating events that this instability could explain.
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In this paper, we describe an exponential instability that acts to narrow and strengthen current sheets in partially ionized gas with a resistivity inversely dependent on temperature. While the physics that we describe pertains to any magnetized system with both current sheets and a conductivity that increases strongly enough with temperature, this problem arises specifically in the context of protoplanetary disks, where the formation of high-temperature minerals such as chondrules and crystalline silicates suggests strong, intermittent heating events. As the local temperature in such a disk climbs above $\sim 1000$~K, the dominant source of free electrons becomes the thermal ionization of alkaline metals. These temperatures are still well below the ionization energy, so the argument of the exponential in the Boltzmann term is quite small; thus the ionization fraction $x_e=n_e/n_n$ depends steeply on temperature. We have found that this results in startling new behavior with positive temperature fluctuations increasing conductivity, concentrating current sheets, and positively feeding back on the temperature through enhanced local Ohmic heating. This is quite different from classical reconnection, with electrical short circuits effectively forming in these regions. This effect is the opposite of the more commonly assumed anomalous resistivity that increases in the reconnection region \citep{KrallLiewer71,SatoHayashi79,Yamada10}. (Classical reconnection was applied to disks using general energetic arguments by \citealt{King10}.) While our mechanism narrows current sheets, similarly to ambipolar diffusion \citep{BZ94}, the mechanism is Ohmic resistivity rather than a drift of the charge carriers relative to the neutral gas. Our mechanism is also different from previous work on partially ionized reconnection because that focused on the transition to the collisionless regime \citep{Malyshkin11, Zweibel11}. Observations of protostellar disks have revealed their integrated properties, including masses \citep{bs93} and accretion rates \citep{h98}. Compositional gradients in the dust \citep{vb04} can be detected, as well as the difference between predominantly amorphous and crystalline mineral structures at different radii \citep{Waelkens96,Malfait98}. The observed presence of crystalline minerals at large radii \citep{Sargent09} suggests the need for a heating mechanism active in the disk far from the parent star. Meanwhile, we have direct evidence of conditions in the protosolar disk. Laboratory measurements of textural, mineralogical, chemical, and isotopic properties of meteoritic chondrules and calcium-aluminum rich inclusions, as well as related high-temperature materials in comet samples \citep{Brownlee06,Zolensky06,Nakamura08,Simon08}, give strong constraints on their local formation history and environment. They represent melts condensed and cooled from temperatures of 1500--1800~K at rates of around 100--1000~K/hour \citep{RadomskyHewins90,LofgrenLanier90,Connolly98,sk05,Ebel06}: far faster than disk dynamical timescales, but far slower than the free-space cooling time of millimeter sized objects. The source of these processed materials in unclear. While there is a vast reservoir of gravitational potential energy in the disk, tapping it through an accretion flow to create hot regions of finite size, as appears to be needed, is non-trivial. The primary source of angular momentum transport in protoplanetary disks appears to be the conversion of orbital kinetic energy into magnetic energy through the magnetorotational instability \citep[MRI;][]{Velikhov59,Chandrasekhar61,bh91,bh98}, with gravitational instability playing a larger role at earlier times \citep[e.g][]{Lodato04}. The ionization structure of the disk midplane may introduce a magnetically dead zone at some radii and periods of disk history \citep{Gammie96}, with reduced though non-zero turbulent viscosity \citep{Fleming03,Oishi09}. Its exact structure and history remains controversial \citep{Glassgold97,Sano00,Ilgner06,Umebayashi09,Turner09}. MRI draws on the huge amount of energy contained in the differential rotation of the disk to drive magnetohydrodynamical (MHD) turbulence. The turbulence will dissipate that energy into heat. However, the heating will occur intermittently, not uniformly. MHD turbulence forms current sheets \citep{Parker72,Parker94,1997PhR...283..227C} that dissipate energy at far greater than the average rate, and can provide locations for magnetic reconnection to occur. \citet{Romanova11a} noted in passing the volume filling nature of this mechanism \citep[also see][]{Romanova11}. \citet{Hirose11} used a moderate resolution simulation to demonstrate that such current sheets forming in the atmosphere above a dead zone can locally heat gas well above the radiative equilibrium temperature. Those calculations, along with preliminary work of our own \citet{McNally12a,McNally12b} suggests that such current sheets may well heat the gas up to temperatures sufficiently high for the instability described here to set in. This is hardly the first suggestion that electromagnetic fields can drastically alter the temperature profile in a protoplanetary disk. \citet{Levy89} examined reconnection heating in disks as a chondrule formation mechanism. However, they worked before the nature of the turbulence driving angular momentum transport was understood. Therefore, they reasoned by analogy to the Sun that a stratified, convective, magnetized flow would drive reconnection in the low density region above it. Thus, they only considered coronal heating many scale heights above the surface of the disk. However, we now understand that the turbulence in the disk is probably not convective but rather driven by the differential rotation acting through the field, so that intermittent dissipation will occur throughout the disk. The short-circuit instability is in many ways similar to lightning \citep{Whipple66,Horanyi95,Pilipp98,DC2000,Muranushi10,Muranushi12}: a rapid local increase in the ionization fraction leads to high currents and a dramatic release of energy. However, there are significant differences in that the increase in ionization is thermal rather than due to electric fields strong enough to directly induce ionization breakdown. Further, the dynamo magnetic fields act as a current source rather than a voltage source. Electrical-short like behavior is only possible due to the residual current that flows in the low temperature regions around the instability, maintaining a non-trivial Ohmic electric field (see \Sec{short:statement}). This allows us to bypass the need to generate. electric fields strong enough to directly ionize the gas, which is a non-trivial challenge to lightning models. Another related proposed mechanism is melting of charged dust by acceleration through standard reconnection regions \citep{Lazerson10}. In this paper, we lay out the basic principles of this novel behavior, and explore the implications for heating in dusty protoplanetary disks in more detail in a companion paper \citet[][hereafter Paper II]{McNally12b}. Closely related behavior in planetary atmospheres was found by \cite{Menou12}, who called it the thermo-resistive instability; though that term was also used by \cite{Price12} for a system where the resistivity increases with temperature. In Sect.~\ref{sec:temp} we describe the physical principles at play. In Sect.~\ref{sec:eqns} we lay out a numerical approach to modeling this behavior, whose results are described in Sect.~\ref{sec:results}, and discussed in Sect.~\ref{sec:discuss}.
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\label{sec:discuss} We have shown that in slab-symmetric reconnection, a strong inverse temperature dependence of the resistivity can lead to an instability that concentrates the current in a narrow, high temperature, low resistivity sheet. This scenario is the polar opposite of the more common situation where the resistivity appears to increase inside current sheets through some anomalous resistivity \citep{KrallLiewer71,SatoHayashi79}. Unlike many reconnection scenarios, the inward transport of the magnetic field in our case is resistive rather than advective in nature, sidestepping issues of fluid pile up that occur with advective field transport, as demonstrated by \Fig{rhoplot}, where the central density declines with time. However, rather than speeding the dissipation of magnetic energy into heat, in one dimension the total dissipation rate actually falls, thanks to the formation of small volumes with very low resistivity, even though heating increases sharply within those regions. This raises interesting questions about the structure of magnetic turbulence in three-dimensional systems with similar temperature dependent resistivities. The instability does not grow extremely fast as can be seen in \Fig{timenorm}, which shows that the instability growth rate estimate $t_c$ is significantly slower than the background resistive broadening time of the current sheet $t_{\eta}$. Nevertheless, as seen clearly in the third column of \Fig{BJT1k}, the instability can grow on timescales of tens of resistive times. For this to occur, the strength of the magnetic field must allow rapid heating, and external large-scale dynamics must not tear the current sheet apart. Both of these conditions appear reasonable for disks subject to the MRI \citep{Sano2007}. Further, it is clear from \Fig{BJT1k} that a fully self-consistent analysis is needed for any MRI active region in a protoplanetary disk whose magnetic field and temperature flirt with $\beta \sim 1-4$ and $T \sim 1000$~K. We have only considered growth of the instability in approximations to current sheets that the MRI generates in the absence of our instability. Such a self-consistent approach will be difficult considering the large range in dissipation parameters that must be resolved and the large spatial scale separation between the turbulence (larger than $\ell_0$) and the narrowed current sheets (much smaller than $\ell_0$, Fig \ref{satplot}). We expect such self-consistent systems to show the concentration of current into localized regions with high current and temperature (either two-dimensional sheets or one-dimensional tubes), with the rest of space taken up by almost force-free magnetic fields. As the concentrated current regions have low resistivity, they could potentially have long lifetimes, perhaps much longer than that associated with the high wavenumber tail of a subsonic turbulent cascade. Where the magnetic field energy does not exceed equipartition with the fluid kinetic energy, the bending of the magnetic field will create new current structures that fence in the magnetic field configuration, with the potential for extremely large and localized Lorentz forces because of the highly concentrated current densities. While we have shown that, in sufficiently restricted circumstances, this instability occurs for any inverse relationship between the resistivity and the temperature, we have also shown that it can take a prohibitive time to set in. In practice it appears that the growth rate estimate of \Eq{instabgrowth}, while sometimes an overestimate, is accurate within factors of a few. Unfortunately, evaluating it requires that the instability be growing, and any initial transients can result in strong overestimates of the growth rate (see \Fig{timenorm}, early times). We explore the peak temperatures achieved by this instability further in Paper II, in particular by including radiative transfer and a fuller treatment of thermal ionization. However the significance for protoplanetary temperature structures in the inner MRI active disk is already clear from the work presented here. Although we have considered this effect from the perspective of protoplanetary disks, it should occur in any system with an adequately strongly increasing conductivity dependence on temperature when compared to available cooling. Candidates include cool stellar surfaces with poorly ionized hydrogen, and even planetary atmospheres, as recently suggested by \citet{Menou12}.
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1206.7096
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1206
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1206.2350_arXiv.txt
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The disruption of stars by supermassive black holes has been linked to more than a dozen flares in the cores of galaxies out to redshift $z \sim 0.4$. Modeling these flares properly requires a prediction of the rate of mass return to the black hole after a disruption. Through hydrodynamical simulation, we show that aside from the full disruption of a solar mass star at the exact limit where the star is destroyed, the common assumptions used to estimate $\md$, the rate of mass return to the black hole, are largely invalid. While the analytical approximation to tidal disruption predicts that the least-centrally concentrated stars and the deepest encounters should have more quickly-peaked flares, we find that the most-centrally concentrated stars have the quickest-peaking flares, and the trend between the time of peak and the impact parameter for deeply-penetrating encounters reverses beyond the critical distance at which the star is completely destroyed. We also show that the most-centrally concentrated stars produced a characteristic drop in $\md$ shortly after peak when a star is only partially disrupted, with the power law index $n$ being as extreme as -4 in the months immediately following the peak of a flare. Additionally, we find that $n$ asymptotes to $\simeq -2.2$ for both low- and high-mass stars for approximately half of all stellar disruptions. Both of these results are significantly steeper than the typically assumed $n = -5/3$. As these precipitous decay rates are only seen for events in which a stellar core survives the disruption, they can be used to determine if an observed tidal disruption flare produced a surviving remnant. We provide fitting formulae for four fundamental quantities of tidal disruption as functions of the star's distance to the black hole at pericenter and its stellar structure: The total mass lost, the time of peak, the accretion rate at peak, and the power-law index shortly after peak. These results should be taken into consideration when flares arising from tidal disruptions are modeled.
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Supermassive black holes (SMBHs) have been found to reside at the centers of most galaxies. These black holes are orbited by a cluster of stars that interact with one another gravitationally through stochastic encounters. Occasionally, an encounter will shift a star onto an orbit that takes it within its tidal radius, defined as the distance at which the black hole's tidal forces would overcome the star's self-gravity at its surface \citep{Frank:1976tg}. A fraction of the star's mass then becomes bound to the black hole, and proceeds to fall back towards the star's original pericenter, eventually forming an accretion disk that results in a luminous flare with a luminosity comparable to the Eddington luminosity. The standard model of tidal disruption presumes that the star is completely destroyed, resulting in approximately half of the star's original mass falling back onto the black hole, with the debris possessing a variety of orbital periods resulting from a spread of orbital energy that is ``frozen in'' at pericenter. First described in \cite{Rees:1988ei}, the rate of fallback has been estimated both through increasingly sophisticated numerical simulations and analytical models. While previous results have provided reasonable models for the fallback resulting from the complete disruptions of stars at the tidal radius $\rt = \rs (\mh/\ms)^{1/3}$, where $\ms$ and $\rs$ are the mass and radius of the star and $\mh$ is the mass of the black hole, they completely neglect partial stellar disruptions, in which a stellar core survives the encounter and only a fraction of the star's mass becomes immediately bound to the black hole. These events are likely to be much more common than their complete disruption counterparts, both for the reason that the rate of encounters interior to the pericenter distance $r_{\rm p}$ scales as $r_{\rm p}$ \citep{Hills:1988br}, and also that the disrupted star may return on subsequent orbits and be subject to disruption and/or further tidal dissipation. Additionally, many previous studies have focused on stars of a single structural profile, usually selected to match the familiar profile of our own Sun. However, standard stellar mass functions predict that low-mass main sequence (MS) stars are more common \citep[e.g.][]{Kroupa:1993tm}, and thus may contribute significantly to the overall disruption rate. These stars are significantly less centrally concentrated than their solar mass brethren. The dynamics of stellar tidal disruption have been modeled by many authors using simple analytical arguments \citep{Rees:1988ei,Phinney:1989uw,Lodato:2009iba,Kasen:2010ci}, increasingly complex dynamical models \citep{Luminet:1985wz,Carter:1985ti,Luminet:1986ch,Diener:1995ui,Ivanov:2001fva}, and hydrodymical simulations utilizing either an Eulerian \citep{Evans:1989jua,Khokhlov:1993cu,Khokhlov:1993bj,Diener:1997kw,Guillochon:2009di} or Lagrangian \citep{Nolthenius:1982dn,Bicknell:1983dn,Laguna:1993cf,Kobayashi:2004kq,Rosswog:2008gc,Lodato:2009iba,RamirezRuiz:2009gw,Rosswog:2009gg,Antonini:2011ia} approaches. Very few of these studies have presented the effect varying $r_{\rm p}$ on the amount of mass lost by the star, $\Delta M$, or the effect on $\md$, the rate at which the mass liberated from the star returns to pericenter. Given that the viscous time is expected to be significantly shorter than the period of the returning debris, this $\md$ is expected to track the luminosity $L(t)$ closely. As the number of observed disruptions increases, and as the cadence and quality of data improves, it becomes increasingly more important to improve models of $\md$ for disruptions of all kinds. In this paper, we present the results of 43 hydrodynamical simulations at high-resolution representing the disruption of both low-mass and high-mass MS stars. This provides, for the first time, a complete picture of the feeding of SMBHs by the disruption of MS stars. While the expected trend of smaller mass accretion rates for progressively more grazing encounters is reproduced, our study reveals several surprises on how disruptions work, particularly on the effect of stellar structure and how the fallback rate scales for both grazing and deep encounters. Contrary to what is expected from the freezing model, in which only the distribution of mass at pericenter is considered, the non-linear response of the star to the tidal field is found to play a crucial role in determining $\md$. Our simulations show that the simple models previously employed to predict the rate of fallback do not capture the full dynamics of the problem, and are only appropriate for anything other than the full disruption at exactly the tidal radius. We find that the decay rate of $\md$ does not settle to a constant value until a few months after the disruption for all disruptions by black holes with $\mh > 10^{6} \msun$, implying that the range of characteristic decay rates used to identify tidal disruption flares should be widened to include events that may not follow the fiducial $t^{-5/3}$ decay rate. For partial disruptions, the decay rate at a few years after the disruption depends crucially on the hydrodynamical evolution of the debris stream. This means that simulations must cover more than a few stellar dynamical timescales after the disruption, with the final functional form of $\md$ not being established until the star is many hundreds of tidal radii away from the black hole. And while we do find that there are differences between the fallback functions calculated for the disruptions of profiles characteristic of low- and high-mass stars, the mass-radius relationship of MS stars results in a family of fallback curves that are difficult to distinguish from one another for stars of $0.3 \msun \gtrsim \ms \gtrsim 1.0 \msun$ without considering secondary features related to the shape of the fallback curves themselves, such as the decay rate of $\md$, characterized by a time-dependent power-law index $n(t)$. This paper is organized as follows. We describe our numerical method, initial models, and method for calculating $\Delta M$ and $\md$ in Section \ref{sec:method}. The results of these simulations and how they improve our understanding of stellar tidal disruptions is described in Section \ref{sec:results}. A discussion of the general trends and their effect on the observable features of tidal disruptions is presented in Section \ref{sec:discussion}. Finally, we provide fitting formula to four characteristic variables describing disruptions of stellar profiles characteristic of low- and high-mass stars in Appendix \ref{sec:appendix}.
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\label{sec:discussion} The results of our extensive parameter study produce a number of unexpected trends as compared to the predictions presented by previous work. In the previous section we attempted to explain the observed scalings, and how these features arise as a result of the interaction between the black hole and a potentially surviving stellar core. In what follows, we explain how these newly discovered features can be used to constrain the type of disrupted star and how it was disrupted. \subsection{Can $\gamma$ and $\beta$ be determined a posteriori?}\label{sec:aposteriori} Given our predicted $\md$, is it possible to determine either the stellar structure or the impact parameter from the light curve produced by a tidal disruption event? The conversion efficiency between mass accreted by the black hole and the light emitted is somewhat uncertain, and depends on factors such as the black hole's spin and how the accretion rate compares to $\dot{M}_{\rm Edd}$ \citep{Ulmer:1999jja,Beloborodov:1999wb,Strubbe:2009ek,Lodato:2010ic}. However, as the efficiency cannot be larger than unity, and as flares are typically observed in the decay phase, we can only place lower limits on the amount of mass accreted by a black hole to produce a given flare \citep{Gezari:2008iv}. Thus, at the very least, our predicted $\Delta M$ (Figure \ref{fig:mlost}) can be used to exclude events for $\beta$ less than some critical value, given the mass of the star. Figures \ref{fig:n} and \ref{fig:peak} present four additional quantities that enable us to classify tidal disruptions based on the properties of observed tidal disruption flares. Two of these quantities, $\mdpeak$ and $\tpeak$, are only available to us for flares in which the peak of the accretion rate is clearly observed \citep{Gezari:2012fk}, but both $n(t)$ and $\ninf$ are measurable for flares that are observed long after peak \citep{Komossa:1999uy,Komossa:2004dr,Gezari:2006gd,Gezari:2008iv,Cappelluti:2009jl,vanVelzen:2011gz,Cenko:2012fg}. If the mass of the black hole is known with some certainty, one may be able to infer both $M_{\ast}$ and $\beta$ by simply measuring $\mdpeak$ and $\tpeak$ and comparing to our resultant $\md$, which at first glance appear to form distinct sequences for $\gamma = 4/3$ and $\gamma = 5/3$ stars (Figure \ref{fig:structure}, left panel). However, this is only true assuming that centrally concentrated stars have the same mass and radius as stars of near constant density. The transition from stars that are well-modeled by a $\gamma = 4/3$ polytrope to a $\gamma = 5/3$ polytrope is also accompanied by a decrease in radius such that all stars with mass $0.25 M_\odot < M_\ast < M_\odot$ have the same central density \citep{Kippenhahn:1990tm}. Adjusting the radii and mass of our $\gamma = 5/3$ models to the mass and radius of a $0.25 M_\odot$ star \citep{Tout:1996waa}, we find that the sequence of $\md$ functions for $1.0 M_\odot$ and $0.25 M_\odot$ stars lie on top of each other (Figure \ref{fig:structure}, right panel), making the determination of the disrupted mass of a star somewhat degenerate with its structure. This motivates us to look for other features of $\md$ that may uniquely identify either $\gamma$ or $\beta$. If we consider the power-law of the rate of decline $n$ after peak, we find that there is a distinguishing feature between $\gamma = 5/3$ and $\gamma = 4/3$ models at $\sim 0.5$ dex after $\tpeak$. Whereas $\gamma = 5/3$ stars quickly converge to $n \simeq -5/3$, $\gamma = 4/3$ models show a characteristic drop, with $n$ being as large as -4 for some encounters (Figure \ref{fig:n}, left panel). This feature is most prominent for intermediate $\beta$ in which $\sim 50\%$ of the star's mass is removed during the encounter, and represents the strong influence of the dense stellar core, which acts to drag material deeper within the black hole's potential before tidal forces are capable of removing it. In addition to being more centrally-concentrated to begin with, an additional component that likely contributes to this observed drop is the adiabatic response of the surviving core. For $\gamma = 5/3$ stars, the removal of mass results in the inflation of the star, whereas $\gamma = 4/3$ exhibit the opposite behavior, shrinking dramatically in response to the loss of mass \citep{Hjellming:1987ci}. This enhances the core's influence during the encounter in the phase where the core's mass is changing, slowing the reduction in the core's effective gravity, and thus pulling even more matter to higher binding energies. The recently observed flare PS1-10jh presented in \cite{Gezari:2012fk} shows a clear drop in the accretion rate with respect to the canonical $t^{-5/3}$ decline rate expected from the freezing model. In the freezing model, it is impossible to produce a decline feature steeper than $t^{-5/3}$ within any part of $\md$, as we explained in Section \ref{sec:characteristic}. As many tidal disruption flares may show this characteristic drop in $\md$, a clearly-resolved peak can be used to compare to the subsequent decay phase for a precise determination of $n(t)$. For events in which the peak is not clearly observed, and for which the signal-to-noise is too small to permit an accurate determination of $n(t)$, the asymptotic slope $\ninf$ of $\md$ can still provide additional information about the star that was disrupted. As shown in the right panel of Figure \ref{fig:n}, $\ninf$ can be used to distinguish between partial and full disruptions. The fact that $\ninf$ assumes values that are significantly steeper than -5/3 may indicate that additional tidal disruption flares have been found observationally, but subsequently discarded and/or ignored due to the mismatch between the measured $n$ and $-5/3$ \citep{vanVelzen:2011gz}. This implies that some supernovae that have been observed at the centers of galaxies may in fact be misidentified partial tidal disruptions. \subsection{Future work} As found in previous work \citep{Faber:2005be,Guillochon:2011be}, there is a change in surviving star's orbital energy after the encounter, with the change in energy being comparable to the star's initial self-binding energy. This change in energy, combined with the star's initial orbital energy, leads to a shift in the entire $\dmde$ distribution, which can affect the fallback of material for $E \sim \Delta E_{\rm orb}$, or for $t \gtrsim \mh R_{\ast}^{3/2} G^{-1/2} M_{\ast}^{-3/2} \sim 100$ years, given that $\Delta E_{\rm orb} \sim G M_{\ast}/R_{\ast}$. As the star's initial orbital energy may not be zero and can be comparable to $\Delta E$ itself, and thus the final binding energy of the star depends on its initial orbital energy, we have presented our $\dmde$ and $\md$ curves with this change in energy removed. As a result, our plots show the fallback rate that would be expected if the final star were to remain on a parabolic trajectory, as our initial conditions assume. While these kicks that are typically of the order of star's own escape velocity may be important in determining the further fate of the star and whether it will suffer additional disruptions, they are not expected to affect the first century of a flare's evolution, of which only the first few years are accessible to currently available transient surveys. Even if $\md$ is directly related to the properties of the star being disrupted, the luminosity of the accretion disk $L$ may not directly follow $\md$. The primary factors that affect the link between $\md$ and the bolometric $L$ are the viscous evolution of the disk and the size of the disk \citep{RamirezRuiz:2009gw}, although other processes may strongly affect the amount of light observed in a single band, especially in the optical/UV where dust extinction can play a vital role. Disk viscosity can only affect $L$ for $t \lesssim \tau_{\rm visc}$, in which its primary affect is to delay emission at early times. However, once $t > \tau_{\rm visc}$, $L$ is expected to track $\md$ closely. As the material is delivered to the disk at $r \simeq r_{\rm p}$, the ratio of $\tau_{\rm visc}$ to $\tpeak$ is \begin{align} \frac{\tau_{\rm visc}}{\tpeak} = 3.2 &\times 10^{-2} \beta^{-3} \left(\frac{B_\gamma\left(\beta\right)}{0.1}\right)^{-1} \times\nonumber\\ &\left(\frac{\mh}{10^{6} \msun}\right)^{-1/2} \left(\frac{\alpha}{0.1}\right)^{-1} \left(\frac{M_\ast}{M_\odot}\right)^{1/2}, \end{align} where $B_\gamma$ is a fitted parameter derived from our simulations, ($B_\gamma \sim 0.1$ for most $\beta$, see Appendix \ref{sec:appendix}) and $\alpha$ is the parameterized $\alpha$-disk scaling coefficient, where we have taken the scale-height ratio $h/r = 1$. If $\tau_{\rm visc}/t_{\rm peak} \gtrsim 1$, the accretion is spread over longer timescales, resulting in a power-law decay index $n = -1.1$ \citep{Cannizzo:1990hw}. This may affect the light curve shape in the earliest phases of the fallback (prior to peak) where $t \ll \tau_{\rm visc}$, and thus the early evolution of $L(t)$ may not follow the functional forms of $\md$ presented here for $t \ll t_{\rm peak}$. But as observations of tidal disruption flares in the decay phase seem to be consistent with the canonical $n = -5/3$ decay law, it is clear that $\md$ and $L(t)$ must be closely coupled on year-long timescales. An ingredient that the set of simulations presented in this paper do not include is the inclusion of general relativistic effects, which become important for very deeply penetrating encounters. Qualitatively, for both spinning and non-spinning black holes, general relativity is expected to result in more mass loss and a spreading of mass in $\dmde$ as compared to Newtonian encounters, as its primary effect is to bend the star's path such that it spends a larger fraction of time near the black hole where tidal forces are important \citep{Luminet:1985wz,Kobayashi:2004kq}. The only numerical provenance for how the metric may affect the feeding rate comes from low-resolution simulations performed by \cite{Laguna:1993cf}, which find a slight increase in $\mdpeak$ for increasing $\beta$, but much less than the predicted $\beta^{3}$ scaling. If the black hole has non-zero spin, the resulting $\dmde$ depends on the orientation of the star's angular momentum vector as compared to the black hole's spin vector \citep{Haas:2012ci}. A spinning black hole permits deeper encounters that don't result in the star being immediately swallowed \citep{Kesden:2012cn}, provided that the two angular momentum vectors are aligned, and also should affect the final binding energy distribution, with co- and counter-rotational encounters resulting in smaller and larger $\Delta M$, respectively \citep{Diener:1997kw,Ivanov:2006fe,Kesden:2012kv}. However, as the fraction of disruptions in which non-Newtonian metrics can affect the dynamics is $\sim r_{\rm s}/r_{\rm t}$, which is $\sim 5\%$ for a $10^{6} \msun$ black hole and $\sim 20\%$ for a $10^{7} \msun$ black hole, the majority of tidal disruption events are well-represented by a Newtonian approximation to the black hole's gravity. Lastly, the absence of hydrogen in spectra taken of the tidal disruption event PS1-10jh \citep{Gezari:2012fk} strongly suggests that the disruption of stars that are not on the MS may contribute significantly to the overall rate of tidal disruption. As we show, the structure of the star that is disrupted is clearly imprinted upon $\md$, providing valuable additional information that can be used to distinguish between candidate disruption victims. We explore the disruption of post-MS stars in a companion paper using a method similar to what is presented here \citep{MacLeod:2012cd}. The discovery of flaring black hole candidates in nearby galaxies will continue to elucidate the demography of the AGN population \citep{DeColle:2012bq}. Whereas AGN are supplied by a steady stream of fuel for hundreds or even thousands of years, tidal disruptions offer a unique opportunity to study a single black hole under a set of conditions that change over a range of timescales. There are, of course, rapidly varying stellar-mass black hole candidates in X-ray binaries within our own Galaxy. But for SMBHs, tidal disruption events offer the firmest hope of studying the evolution of their accretion disks for a wide range of mass accretion rates and feeding timescales. The simulations and resultant $\md$ curves presented here are crucial for determining the properties of the black hole itself, as an incomplete model of a stellar disruption can result in much uncertainty in how the black hole converts matter into light. For a disruption with a well-resolved light curve, our models permit a significant reduction of the number of potential combinations of star and black hole properties, enabling a better characterization of SMBHs and the dense stellar clusters that surround them.
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1206.2350
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1206
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1206.2559.txt
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%context {CAL\,83 is a prototype of the class of Super Soft X-ray Sources (SXS). It is a binary consisting of a low mass secondary that is transferring mass onto a white dwarf primary and is the only known SXS surrounded by an ionisation nebula, made up of the interstellar medium (ISM) ionised by the source itself. We study this nebula using integral field spectroscopy. } %aims {The study of ionised material can inform us about the source that is responsible for the ionisation, in a way that is complementary to studying the source directly. Since CAL\,83 is the only SXS known with an ionisation nebula, we have an opportunity to see if such studies are as useful for SXSs as they have been for other X-ray ionised nebulae. We can use these data to compare to models of how CAL\,83 should ionise its surroundings, based on what we know about the source emission spectrum and the physical conditions of the surrounding ISM.} %Methods {With the VIMOS integral field spectrograph we obtained spectra over a $25 \times 25$\arcsec field of view, encompassing one quarter of the nebula. {\bf Emission line maps -- H\,{\sc i}, \Hetwo, \OthreeC, \NtwoC, and \StwoC\ -- }are produced in order to study the morphology of the ionised gas. We include CAL\,83 on diagrams of various diagnostic ion ratios to compare it to other X-ray ionised sources. Finally we computed some simple models of the ionised gas around CAL\,83 and compare the predicted to the observed spectra. } %Results {CAL\,83 appears to have a fairly standard ionisation nebula as far as the morphology goes: the edges where H is recombining are strong in the low stage ionisation lines and the central, clumpy regions are stronger in the higher stage ionisation lines. But the He\,{\sc ii} emission is unusual in being confined to one side of CAL\,83 rather than being homogeneously distributed as with the other ions. We model the CAL\,83 nebula with {\sc cloudy}\,\normalfont {\bf using model parameters for SXSs found in the literature}. The He\,{\sc ii} emission does not fit in with model predictions; in fact none of the models is able to fit the observed spectrum very well.} %Conclusion {The spectral line images of the region surrounding CAL\,83 are revealing and instructive. However, more work on modelling the spectrum of the ionised gas is necessary, and especially for the high-ionisation level emission from CAL\,83. In particular, we wish to know if the He\,{\sc ii} emission {\bf and the other nebular lines are powered by the same ionising source.}} % Max 6 keywords!
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\label{sect:intro} Between 1979 and 1981, the Columbia Astrophysics Laboratory (CAL) carried out a systematic soft X-ray survey of the large Magellanic cloud (LMC) using the \textit{Einstein} \textrm{Observatory}. This resulted in the detection of 97 X-ray sources \citep*{Long1981}. Some of the detected sources are characterised by an unusually soft spectrum in which little or no radiation at energies above $\sim0.5$ keV is detected. Source no. 83, being one of a small number of such sources, became known as CAL\,83 (also known as LHG83, $RX J0543.7$-6822, and $1E0543.8$-6823) and it is now regarded as the prototype of the class of supersoft X-ray sources (SXSs). In 1990, the first X-ray all-sky survey was performed with \textrm{ROSAT} and many more similar sources were found in the Galaxy, LMC, and M31 (\citealt*{Kahabka1997}, \citealt{Parmar1998}). These sources are characterised by a luminosity of $\sim$\,$10^{37-38}$ \ergs\ and effective temperatures in the range of $\sim$\,$2-6\times10^5$ K ($kT\simeq17$--50 eV) \citep*{Rappaport1994}. Van den Heuvel et al.\ (1992) were the first to propose a model for SXSs: an accreting white dwarf of mass 0.7--1.2$M_{\odot}$ accompanied by a normal star of mass 1.5--2$M_{\odot}$. According to their model the supersoft X-ray emission and high luminosities are a result of steady nuclear burning of hydrogen accreted onto the white dwarf. The mass transfer from the main-sequence star to the white dwarf occurs on a thermal timescale via Roche-lobe overflow, with rates of $\sim1-4\times10^{-7}M_{\odot}$yr$^{-1}$. The dominant optical light source is the white dwarf, which, with its accretion disc, completely outshines the donor star. The nuclear burning on the surface of the white dwarf allows it to retain the accreted mass, making it a possible progenitor of accretion-induced collapse \citep{Heuvel1992} or Type Ia supernovae \citep{Rappaport1994, Nelson1996, Hachisu1996}. CAL\,83 has a time-averaged luminosity $L_{X}$\,(0.15--4.5\,keV) of $3.2\times10^{37}$ \ergs\ and is identified with a variable B$\sim$16.8 star with a blue continuum and strong narrow ($\sim$$305$ \kms), variable \Hetwo\ emission \citep{Cowley1998}. It is, in fact, a mass-transferring binary, with a secondary mass of $\sim0.5M_{\odot}$ \citep{Cowley1998} and a white dwarf primary mass of $1.3\pm0.3 M_{\odot}$ \citep{Alcock1997,Lanz2005}. The orbital period of the system is 1.04 days \citep{Smale1988}. While CAL\,83 is one of a dozen luminous, extremely soft SXSs now known in our Galaxy and the Magellanic Clouds, so far it is the only one known to have an ionisation nebula. Its inhomogeneous nebula has bright emission extended over 37\arcsec, with fainter emission extending out beyond twice that, and a total mass of $\sim150M_{\odot}$ (\citealt*{Remillard1995} = RRM95); obviously the nebula is ionised interstellar medium (ISM), rather than being intimately related to CAL\,83 itself. SXSs emit copious quantities of photons in the range 20--200 eV, hence the radiation can ionise any gas surrounding them, creating an ionisation nebula. We expect these ionisation nebulae to be distinct from classic H\,\scriptsize{II}\normalsize\ regions where the ionisation is the result of photo-ionisation by massive O-type stars by the absorption of higher energy photons while the lower energy photons escape. In comparison, in the case of the SXSs it is the other way around: the lower energy photons do the ionising while the higher energy photons escape (RRM95). Models for SXS nebulae by \citet{Chiang1994} predict that these should be distinct from other astrophysical nebulae, in particular, \Othree\ and \Hetwo\ should be far brighter than in classic \Htwo\ regions, and the radial gradients of these and other lines much more gradual. These models have not yet been tested on real data of SXSs, and CAL\,83 presents a good opportunity to do so. Using the VLT VIMOS \citep{Lefevre2003} integral field unit (IFU) we have observed one field around CAL\,83, obtaining spatial and spectral information about its surrounding ISM. In this paper we report on our findings, presenting emission line flux maps, spectra, line ratio maps, and a comparison of our results to models of ionised ISM created with the {\sc cloudy} code \citep{Ferland1998}. %####################################################################################### %****************************************************************** % DATA REDUCTION %******************************************************************
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In this paper, we have presented flux maps made from the fitting the emission lines of \OthreeB, \OoneA, \Ha\ and \StwoA\ showing the morphology of these ions in our FoV. The morphology in \OthreeB\ and \Ha\ match those found by RRM95 who imaged the full nebula in these ions. We find the edge of the RRM95 nebula is where low ionisation ions peak, while inside this the ionisation state is higher. We have also presented, for the first time, evidence of an \Hetwo\ region around CAL\,83. The \HetwoX\ emission peaks at the position of CAL\,83, but it has a distinctly asymmetrical distribution around the central star. We did not detect any \Heone\ emission. We estimated an average value for the electron density $n_e$ in our FoV of $\sim$10 cm$^{-3}$ which is consistent with the value found by RRM95. In addition to the flux maps, we also show flux {\em ratio} maps which characterise zones with different ionisation degrees in our FoV. The flux ratio values of four interesting positions in our FoV were used to place CAL\,83 in the diagnostic diagrams of \citet{Veilleux1987} which is used to distinguish between AGNe and \Htwo\ region-like objects (i.e. star forming galaxies). This has not been done before for a SXS, as CAL\,83 is the only know SXS surrounded by an ionisation nebula. We show that CAL\,83 has characteristics of both AGNe and \Htwo\ region-like objects and does not seem to be distinguishable from ULXs. We have also placed CAL\,83, both observed and modelled, on a plot of \OthreeB/\Hb\ vs \Hetwo/\Hb, developed especially for SXSs by \citet{Chiang1994}. CAL\,83, observed and modelled, differ slightly in the \Hetwo/\Hb\ ratio and more in the \OthreeB/\Hb\ ratio, and is situated with the ULXs and starburst galaxies rather than with the \Htwo\ regions. Finally, we have presented a comparison between our observations and model calculations for nebulae surrounding SXSs using the {\tt CLOUDY} ionisation code. Improvement on the models of Rappaport et al. (1994a) was achieved by utilising LMC abundances rather than Galactic. We found that none of the models presented matches our observations completely. Keeping in mind that the observed nebula is not even close to being homogeneous, as is assumed in the modelling, one could argue that the models are perhaps too simplistic. The modelling inconsistencies do not, however, affect the observationally oriented results presented in this paper. Perhaps the discrepancies are telling us something interesting and potentially important about the physical processes that we do not yet understand. To get to the bottom of this, new observations covering the entire CAL\,83 nebula should be compared to more detailed models which perhaps take into account a certain degree of inhomogeneity. %Finally, we have presented a comparison between our observations and model calculations for nebulae surrounding SXSs using the {\sc cloudy} ionisation code. We improve on the models of \citet{Chiang1994} because we use the LMC abundances rather than the Galactic. We found that none of the presented models match our observations completely. Keeping in mind the observed nebula is not at all homogeneous like assumed in the modeling, one could argue the models are perhaps too simplistic. However, there are some similarities between our results and the predictions by the models, the clearest being the observation that \Othree\ is by far the strongest emitting ion in the nebula followed by \Ha\ which is also predicted by the models. The modelling inconsistency does however not affect our purely observational results of this paper. Perhaps the discrepancies are telling us something interesting and potentially important about the physical processes that we don't yet understand. To get to the bottom of this, new observations covering the whole nebula should be compared to more detailed models which maybe take into account a certain degree of inhomogeneity. %#######################################################################################
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1206.2559
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1206
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1206.0025_arXiv.txt
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Neutron matter presents a unique system for chiral effective field theory (EFT), because all many-body forces among neutrons are predicted to next-to-next-to-next-to-leading order (N$^3$LO). We present the first complete N$^3$LO calculation of the neutron matter energy. This includes the subleading three-nucleon (3N) forces for the first time and all leading four-nucleon (4N) forces. We find relatively large contributions from N$^3$LO 3N forces. Our results provide constraints for neutron-rich matter in astrophysics with controlled theoretical uncertainties.
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1206.0025
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1206
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1206.0739_arXiv.txt
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Observations of black hole binaries (BHBs) have established a rich phenomenology of X-ray states. The soft states range from the low variability, accretion disc dominated thermal (TD) state to the higher variability, non-thermal steep power law state (SPL). The disc component in all states is typically modeled with standard thin disc accretion theory. However, this theory is inconsistent with optical/UV spectral, variability, and gravitational microlensing observations of active galactic nuclei (AGNs), the supermassive analogs of BHBs. An inhomogeneous disc (ID) model with large ($\simeq 0.4 \rm$ dex) temperature fluctuations in each radial annulus can qualitatively explain all of these AGN observations. The inhomogeneity may be a consequence of instabilities in radiation-dominated discs, and therefore may be present in BHBs as well. We show that ID models can explain many features of the TD and SPL states of BHBs. The observed relationships between spectral hardness, disc fraction, and rms variability amplitude in BHBs are reproduced with temperature fluctuations similar to those inferred in AGNs, suggesting a unified picture of luminous accretion discs across orders of magnitude in black hole mass. This picture can be tested with spectral fitting of ID models, X-ray polarization observations, and MHD simulations of radiation-dominated discs. If BHB accretion discs are indeed inhomogeneous, only the most disc-dominated states (disc fraction $\gtrsim 0.95$) can be used to robustly infer black hole spin using current continuum fitting methods.
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Spectral and timing observations of black hole X-ray binaries (BHBs) in the energy range $\sim 1-100$ keV have established characteristic outburst ``states'' common to many sources \citep[e.g.,][hereafter RM06]{remillardmcclintock2006}. The ``thermal'' (TD) state has little variability, and a large fraction of the integrated flux is well described by a standard thin black hole accretion disc model \citep[NT,][]{shaksun1973,novthorne}. The ``low/hard'' state is characterized by a hard power law spectrum with a weak, soft disc component. It is often modeled as a thin disc truncated well outside of the black hole marginally stable orbit. Inside of this location, the accretion flow is assumed to be geometrically thick and radiatively inefficient \citep{esinetal1997}. The ``steep power law'' state (SPL) is characterized by a soft (photon index $\Gamma > 2.4$) power law extending unbroken to $\lesssim 1$ MeV \citep{groveetal1998}. Its physical origin remains uncertain. This common interpretation of the X-ray states relies on the validity of the NT model. However, this model is theoretically inconsistent in the TD and SPL states when the disc extends down to the marginally stable orbit of the black hole. In these luminous inner regions, radiation pressure provides the vertical support against gravity, and the NT model is both thermally \citep{shakurasunyaev1976} and viscously \citep{lightmaneardley1974} unstable. Global MHD simulations of thin discs \citep[][and references therein]{pennaetal2010,sorathiaetal2012}, which include the physics of angular momentum transport via the magnetorotational instability \citep{mri}, do not include radiation pressure and therefore cannot test the validity of the NT model. Local MHD simulations find that small patches of radiation-dominated discs are thermally stable \citep{hirosestable2009}, but suggest that the viscous (Lightman-Eardley) instability may operate \citep{hiroseetal2009}. This instability has also been proposed as a model for spectral state transitions in GRS 1915+105 \citep{bellonietal1997,neilsenetal2012}. The NT model is disfavored by optical/UV observations of active galactic nuclei (AGNs) with similar $L/L_{\rm edd}$ to the soft, luminous states of BHBs. The model underproduces the observed UV emission \citep[][]{zhengetal1997} and requires a relativistic mechanism to explain the simultaneous variability observed at well separated optical wavelengths \citep[][]{kroliketal1991}. Recent microlensing observations find that quasar accretion discs are a factor of $\sim 4$ larger than predicted by the NT model \citep[][and references therein]{jimenezvicenteetal2012}. \citet[][hereafter DA11]{dexteragol2011} showed that a disc with large, local temperature fluctuations ($\simeq 0.4$ dex) can explain all of these observations. The fluctuations could be driven by instabilities in radiation pressure dominated discs, in which case they would operate in BHBs as well. In this Letter, we propose a novel interpretation of the soft, luminous TD to SPL states of BHBs in terms of this inhomogeneous accretion disc model (ID, summarized in \S \ref{sec:model}). We show how interpreting ID spectra with the standard NT model can explain the X-ray spectral and variability properties of soft BHBs (\S \ref{sec:states}). In \S\ref{sec:power-law-tail} we assess the requirements for the ID model to fit observed BHB X-ray spectra. Inhomogeneity also may have important implications (\S \ref{sec:impl-cont-fitt}) for the continuum fitting method for measuring black hole spin \citep[CF, ][and references therein]{mcclintocketal2011}. We discuss the prospects for testing BHB accretion disc inhomogeneity observationally and theoretically in \S\ref{sec:discussion}.
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\label{sec:discussion} X-ray spectral and timing observations of BHBs have established a set of states common to many objects. The physical origin of the SPL state and the cause of state transitions remain poorly understood. We have shown that the spectral and rms variability properties of soft, luminous BHBs (TD to SPL states) can be explained by varying one parameter: the amplitude of local temperature fluctuations in an inhomogeneous accretion disc. Comparable temperature fluctuations in quasar accretion discs can resolve a number of important observational puzzles in AGNs (DA11). The ID model then provides a unified picture of luminous black hole accretion discs across many orders of magnitude in black hole mass. The most likely physical processes for driving large temperature fluctuations are radiation pressure dominated disc instabilities. Local MHD simulations are consistent with the thin disc prediction for the relationship between surface density and stress (or temperature) when time-averaged \citep{hiroseetal2009}. The simplest interpretation is that the disc is viscously unstable. If so, it would likely produce temperature fluctuations qualitatively similar to those in the ID model. This can be tested with radially-extended radiation MHD simulations. Alternatively, large fluctuations could also occur in magnetically-dominated discs \citep[e.g.,][]{begelmanpringle2007}, where the disc radiation energy content is de-coupled from its stability. The ID model is insufficient to fit the X-ray spectra observed in both AGNs and BHBs. The broad ID spectrum produces some, but not all, of the emission typically attributed to the high energy tail, and some prescription for the remaining emission is still required. This could be from Comptonization in a compact corona or from hot gas inside the disc marginally stable orbit \citep{zhudavisetal2012}. As with other models, the high energy tail normalization must increase with decreasing disc fraction in order to explain SPL spectra. Quantitative fitting of the ID spectral model to observations will be carried out in future work, with various prescriptions for the high energy tail. Spectral models could use the colour-corrected blackbody used here, or could incorporate the results of sophisticated accretion disc atmosphere calculations \citep[BHSPEC,][]{davisetal2005,davisetal2006}. If BHBs are indeed inhomogeneous, spin measurements from current CF methods using the NT model are potentially subject to large systematic uncertainties (spin differences $\Delta a \sim 0.3-0.9$ for ID spin $a=0$, Figure \ref{deltaa}). These can be much larger than current statistical uncertainties and systematic errors from ignoring emission from the plunging region \citep[$0.2-0.3$,][]{kulkarnietal2011,nobleetal2011}, except in the most disc-dominated states (disc fraction $\gtrsim 0.95$ or $\sigma_T \lesssim 0.15$). Therefore, if BHB discs are significantly inhomogeneous, spin measurements using CF methods should be restricted to the most disc-dominated states. The ID model also generically predicts that the inferred spin using NT models should increase with decreasing disc fraction. \citet{steinersilver2009,steineretal2010,steineretal2011} have found that inferred spins in a few sources do not change significantly over a wide range of $L/L_{\rm edd}$ or $f_{\rm SC}$. Since $f_{\rm SC}$ is a proxy for disc fraction, this latter conclusion is in apparent conflict with our Figure~\ref{deltaa} \citep[particularly Figure 1 of][]{steinersilver2009}. However, the magnitude of the inferred spin differences when fitting NT models to ID spectra depends sensitively on the high energy tail model assumed. Further, the inferred spin from fitting ID models is nearly constant outside of a critical range in disc fraction. Observations of IDs entirely above (below) this range in disc fraction could find a constant, correct (incorrect) spin value, although there is no evidence for sharp changes in spin from CF at disc fractions $\simeq 0.95$. Fitting X-ray spectra with the ID model will allow estimates of $a$ and $\sigma_T$ in a variety of states, and test the model. Requiring the spin to be independent of spectral state may constrain $\sigma_T$ and the degree of inhomogeneity in BHB accretion discs. X-ray polarization measurements may provide another test for inhomogeneous accretion discs in BHBs. Assuming a semi-infinite scattering-dominated atmosphere \citep{chandrasekhar1950}, we have calculated polarized images and spectra via ray tracing using the codes \textsc{geokerr} \citep{dexteragol2009} and \textsc{grtrans} \citep{dexter2011}. We account for all relativistic effects, including the rotation of the polarization vector \citep{connorspiranstark1980,agolphd}. With these assumptions, the time-averaged ID polarization is unchanged from the NT model: $1-5\%$ peak polarization, increasing strongly from face-on to edge-on viewing, and also increasing at lower black hole spin \citep{schnittmankrolik2009}. X-ray polarization may then provide an additional constraint on the BH spin, independent of the location of the spectral peak. In addition, the polarization angle and degree vary by factors of a few on the fluctuation timescale. Both the degree of polarization and its time-variability increase with increasing $\sigma_T$. These estimates ignore returning radiation, as well as coronal \citep{schnittmankrolik2010} and plunging region \citep{agolkrolik2000} emission. These effects could cause additional differences between the polarization signatures of the ID and NT models. The ID model used here assumes an optically thick disc. The NT effective optical depth is \citep[e.g.,][]{abramowiczfragile2011}: \begin{equation} \tau_{\rm NT} \sim 1 \left(\frac{L}{L_{\rm edd}}\right)^{-2} \left(\frac{R}{10 M}\right)^{93/32}, \end{equation} \noindent with $G=c=1$ and ignoring the effects of general relativity and of the inner disc edge on the optical depth. If the disc is inhomogeneous, portions will become optically thin. Assuming independent temperature and density fluctuations, we estimate the fraction of the disc with $\tau_{\rm ID} < 1$ by integrating the log-normal distribution, $f(w)$, for $\tau_{\rm ID} (w) = w^{2.3} \tau_{\rm NT} < 1$. For $\sigma_T < 0.1$, $\gtrsim 99\%$ of the disc is optically thick where the bulk of the luminosity is produced for $L/L_{\rm edd} \le 1$. At larger $\sigma_T$, the inner disc will become optically thin ($\gtrsim 50\%$ by area at the peak emission radius for $\sigma_T > 0.35$ and $L/L_{\rm edd}=1$). If this emission is similar to that from the plunging region \citep{zhudavisetal2012}, it could lead to the observed increasing high energy tail emission with decreasing disc fraction. Comptonization from hot electrons in optically thin regions may also effectively truncate the disc spectrum as inferred in the SPL state \citep[][]{donekubota2006}.
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1206.0739
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1206
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1206.1215_arXiv.txt
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We present $H$-band polarimetric imagery of UX Tau A taken with HiCIAO/AO188 on the Subaru Telescope. UX Tau A has been classified as a pre-transitional disk object, with a gap structure separating its inner and outer disks. Our imagery taken with the 0.15$\arcsec$ (21 AU) radius coronagraphic mask has revealed a strongly polarized circumstellar disk surrounding UX Tau A which extends to 120 AU, at a spatial resolution of 0.1$\arcsec$ (14 AU). It is inclined by 46$\degree$ $\pm$ 2$\degree$ as the west side is nearest. Although SED modeling and sub-millimeter imagery suggested the presence of a gap in the disk, with the inner edge of the outer disk estimated to be located at 25 -- 30 AU, we detect no evidence of a gap at the limit of our inner working angle (23AU) at the near-infrared wavelength. We attribute the observed strong polarization (up to 66 \%) to light scattering by dust grains in the disk. However, neither polarization models of the circumstellar disk based on Rayleigh scattering nor Mie scattering approximations were consistent with the observed azimuthal profile of the polarization degrees of the disk. Instead, a geometric optics model of the disk with nonspherical grains with the radii of 30 $\mu$m is consistent with the observed profile. We suggest that the dust grains have experienced frequent collisional coagulations and have grown in the circumstellar disk of UX Tau A.
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A protoplanetary disk around a young stellar object is the site of planetary formation. The core accretion model (e.g., \cite{Nakagawa83}) and the gravitational instability model (e.g., \cite{Boss98}) are two possible methods by which the process of planet formation might occur. In the core accretion model, a key process is dust grain growth by collisional coagulations. Several fundamental processes of the dust growth have been proposed under simple assumptions. Dust grains with sizes less than a few tens of micrometers orbit a central star with the same velocity as gas (\cite{Adachi76}). They frequently collide and coagulate with each other. Simultaneously they begin to settle into an equatorial plane of the disk (\cite{Nakagawa81}). Dust grains with different settling velocities also experience collisional coagulations. As a result, planetesimals (radii $\sim$ 10 -- 100 km; \cite{Kokubo98}) are formed in the mid-plane of the disk. Bodies between a few tens of micrometers and a few kilometers are dominated by gas drag. Such bodies are thought to migrate towards the central star in a short timescale (e.g., \cite{Nakagawa86}). This rapid inward migration poses a challenge to planet formation theory. By contrast, for planetesimals larger than a few kilometers, gravitational interactions are significant and relative velocities between the bodies increase due to mutual perturbations. This situation allows faster growth of the larger bodies and leads to formation of planetary embryos (radii $\sim$ 1000 km; \cite{Kokubo00}). The core accretion model is one of the plausible model of planetary formation. However it is still debated how sub-micron sized grains are transformed into kilometer sized bodies. An infrared excess in the spectral energy distribution of a young stellar object provides indirect evidence for the presence of a circumstellar disk. The spectral energy distributions of T Tauri stars with continuous, optically thick disks have strong infrared excesses in the near-infrared to far-infrared wavelengths (\cite{Williams11}). Recently, objects with large continuum excesses in the mid- to far-infrared wavelengths but no excess in the near-infrared wavelengths have been discovered, called 'transitional disk objects' (TDOs; \cite{Calvet02}, \cite{Calvet05}). Such an object is considered to lie in the transition state between classical T Tauri stars (CTTSs) and weak-line T Tauri stars (WTTSs). It is expected that the inside of the disk has been cleared out by dust accumulations and/or formations of protoplanets. Photo-evaporation process is also proposed as a mechanism to create a transitional disk. On the other hand, 'pre-transitional disk objects' (PTDOs) are considered to be in the evolutionary phase before reaching the transitional disk phase. They have mid- to far-infrared excesses similar to TDOs, however, they also show small excess in the near-infrared wavelengths (\cite{Esp07}). The slight near-infrared excess implies that optically thick material remains in the innermost part of the disk. \citet{Esp10} proposed that the planetary formation process is expected to progress in the gap structures between the inner and outer disks. High spatial resolution coronagraphic imaging polarimetry is one way to directly diagnose young circumstellar disks. The degree of polarization depends on scattering angle, grain size, and composition. By constructing the polarization profile of the disk, we are able to investigate size and composition of dust grains. \citet{Silber00} conducted infrared polarization imaging observations of a circumbinary disk around GG Tau with NICMOS mounted on the Hubble Space Telescope. The circumbinary disk shows strong polarization degrees of $\sim$ 50\% at 1 $\mu$m wavelength. The polarization azimuthal profile indicated Rayleigh-like scattering from dust grains with sub-micron size. \citet{Hashimoto11} conducted $H$-band polarization imaging observations of AB Aur with HiCIAO/AO188 on the Subaru Telescope and revealed spiral structure in the outer part and the double ring structure at the inner part of the circumstellar disk. There was a number of studies in the last years, using the same observational technique for similar sources, e.g., \authorcite{Apai04} (\yearcite{Apai04}; TW Hya), \authorcite{Oppenheimer08} (\yearcite{Oppenheimer08}; AB Aur), \authorcite{Perrin09} (\yearcite{Perrin09}; AB Aur), \authorcite{Quanz11} (\yearcite{Quanz11}; HD100546), \authorcite{Quanz12} (\yearcite{Quanz12}; HD97048). UX Tau is a T Tauri multiple system (\cite{Jones79}) in the Taurus molecular cloud (distance $\sim$ 140 pc; \cite{Elias78}). It consists of a primary star (UX Tau A), with UX Tau B separated by 5.86$\arcsec$ and UX Tau C separated by 2.63$\arcsec$ from the primary. UX Tau B is itself a binary system with the separation of 0.14$\arcsec$. The spectral type of UX Tau A is K2 (\cite{Kraus09}). Its spectral energy distribution shows a slight excess in the near-infrared wavelengths and significant excesses in the mid- and far-infrared wavelengths. These characteristics indicate that UX Tau A has an optically thick inner disk separated from an optically thick outer disk by a gap, i.e. it is a pre-transitional disk object (\cite{Esp10}). Model fits of the SED suggested that the outer edge of the inner disk is located at $<$ 0.21 AU and the inner wall of the outer disk is located at 30 AU from the central star (\authorcite{Esp10} \yearcite{Esp10}, \yearcite{Esp11}). The disk around UX Tau A was spatially resolved by Sub-millimeter Array at 880 $\mu$m wavelength, with the spatial resolution of 0.3$\arcsec$ (\cite{And11}). They found a dust-depleted disk cavity around the central star, and estimated the inner edge of the outer disk to be located 25 AU from the central star. UX Tau A shows no 10 $\mu$m silicate emission (\cite{Esp10}), implying a lack of small dust grains in the disk. These features suggest dust grain growth in the circumstellar disk of UX Tau A. We conducted polarization imaging observations of UX Tau A in the $H$-band (1.6 $\mu$m) and investigate the collisional coalescence process of the dust grains in its disk.
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With HiCIAO/AO188 mounted on the Subaru Telescope, we carried out $H$-band polarimetric imaging observations of UX Tau A, which has been classified as a pre-transitional disk object. The observation revealed a circumstellar disk around UX Tau A at the spatial resolution of 0.1$\arcsec$ beyond 23 AU from the central star. The disk extends to 120 AU in radius. The disk is inclined by 46$\degree$ $\pm$ 2$\degree$ as the west side is near. We did not detect evidences of the inner disk and the gap-like structure. The notable feature of the circumstellar disk is the huge variation of the polarization degrees. It varies from 1.6 to 66 \%. We constructed several polarization models of the circumstellar disk based on the Rayleigh scattering and Mie scattering approximations. However neither models were consistent with the observational azimuthal profile of the polarization degrees. Focusing on geometric optics, we built the polarization model of the geometrically thin disk with nonspherical grains with the radii of 30 $\mu$m. The model reproduced well the observational azimuthal profile of the polarization degree. We suggest that UX Tau A has a geometrically thin disk containing the nonspherical dust grains with the radii of 30 $\mu$m. Such a disk with nonspherical large dust grains is consistent with the core accretion model of planetary formation process. At 40 AU from UX Tau A, the dust grains can grow up to 100 $\mu$m in radius by collisional coagulation and settle toward the mid-plane with the timescale of 10$^{5}$ years at the earliest. Observational evidence of large dust grains as well as the gap structure in the circumstellar disk provides robust signatures of planetary formation process in the UX Tau A system. \bigskip We thank the telescope staff members and operators at the Subaru Telescope. This work is partly supported by the JSPS-DST collaboration. E.L.T. gratefully acknowledges support from a Princeton University Global Collaborative Research Fund grant and the World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan. J. Carson gratefully acknowledges support from NSF grant AST-1009203.
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1206.1215
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1206
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1206.6603_arXiv.txt
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{We present the XMM-Newton temperature profiles of 12 bright ($L_X>4\times10^{44}$ erg s$^{-1}$) clusters of galaxies at $0.4<z<0.9$, having an average temperature in the range $5\la kT\la 11$~keV.} {The main goal of this paper is to study for the first time the temperature profiles of a sample of high-redshift clusters, to investigate their properties, and to define a universal law to describe the temperature radial profiles in galaxy clusters as a function of both cosmic time and their state of relaxation.} {We performed a spatially resolved spectral analysis, using Cash statistics, to measure the temperature in the intracluster medium at different radii.} {We extracted temperature profiles for the clusters in our sample, finding that all profiles are declining toward larger radii. The normalized temperature profiles (normalized by the mean temperature $T_{500}$) are found to be generally self-similar. The sample was subdivided into five cool-core (CC) and seven non cool-core (NCC) clusters by introducing a pseudo-entropy ratio $\sigma=(T_{IN}/T_{OUT})\times(EM_{IN}/EM_{OUT})^{-1/3}$ and defining the objects with $\sigma<0.6$ as CC clusters and those with $\sigma\ge0.6$ as NCC clusters. The profiles of CC and NCC clusters differ mainly in the central regions, with the latter exhibiting a slightly flatter central profile. A significant dependence of the temperature profiles on the pseudo-entropy ratio $\sigma$ is detected by fitting a function of $r$ and $\sigma$, showing an indication that the outer part of the profiles becomes steeper for higher values of $\sigma$ (i.e. transitioning toward the NCC clusters). No significant evidence of redshift evolution could be found within the redshift range sampled by our clusters ($0.4<z<0.9$). A comparison of our high-z sample with intermediate clusters at $0.1<z<0.3$ showed how the CC and NCC cluster temperature profiles have experienced some sort of evolution. This can happen because higher $z$ clusters are at a less advanced stage of their formation and did not have enough time to create a relaxed structure, which is characterized by a central temperature dip in CC clusters and by flatter profiles in NCC clusters.} {This is the first time that a systematic study of the temperature profiles of galaxy clusters at $z>0.4$ has been attempted. We were able to define the closest possible relation to a universal law for the temperature profiles of galaxy clusters at $0.1<z<0.9$, showing a dependence on both the relaxation state of the clusters and the redshift. }
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Clusters of galaxies are the largest virialized structures in the Universe, arising from the gravitational collapse of high peaks of primordial density perturbations. They represent unique signposts where the physical properties of the cosmic diffuse baryons can be studied in great detail and can be used to trace the past history of cosmic structure formation \citep[e.g.][]{peebles93, coles95, peacock99, rosati02, voit05}. As a result of adiabatic compression and shocks generated by supersonic motion during shell crossing and virialization, a hot thin gas permeating the cluster's gravitational potential well is formed. This gas reaches temperatures of several $10^7$ K and therefore emits mainly via thermal bremsstrahlung and is easily detectable in the X-rays. The gas key observable quantities are its density, temperature, and metallicity. Assuming hydrostatic equilibrium, the gas temperature and density profiles allow one to derive the total cluster mass and thus to use galaxy clusters as cosmological probes \citep[e.g.][]{voit05}. Temperature and density profiles can also be combined to determine the intracluster medium (ICM) entropy distribution, which provides valuable information on the cluster thermodynamic history and has proven to be a powerful tool in investigating non-gravitational processes \cite[e.g.][]{pratt06}. Radial temperature profiles of the hot ICM in galaxy clusters are therefore very important for studying the gravitational processes responsible for large-scale structure formation and non-gravitational energy input into the ICM. Early measurements of cluster temperature profiles were obtained by ASCA and Beppo-SAX \citep[e.g.][]{ikebe97,markevitch98,ettori00,finoguenov01,nevalainen01,degrandi02}. In particular, using ASCA data, \citet{markevitch98} obtained temperature profiles for a sample of 32 nearby clusters, which showed significant declines with radius between $r=0.1r_{vir}$ and $r=0.6r_{vir}$. These authors also found that in clusters without obvious mergers, the radial temperature profiles outside the cool cores were similar when normalized to the virial radius. Consistent results were obtained using the spatially resolved spectroscopic capabilities of BeppoSAX. In their analysis of 21 clusters, \citet{degrandi02} found declining temperature profiles, in good agreement with \citet{markevitch98} outside the cores [$r > 0.15-0.2r_{vir}$], although they were less peaked at the center. De Grandi and Molendi also found that for $r>0.2r_{180}$, where the gas can be treated as a polytrope, the polytropic index derived for cool-core (CC) clusters is significantly flatter than the index derived for non cool-core (NCC) clusters, corresponding to $1.20\pm0.06$ for CC clusters and to $1.46\pm0.06$ for NCC clusters. More recently, Chandra and {\em XMM-Newton} allowed measuring the temperature profiles of galaxy clusters with better accuracy and improved spatial resolution, especially in their central regions \citep[e.g.][]{vikhlinin05,vikhlinin06b,baldi07,pratt07,leccardi08a}, without the complications introduced by the point spread function (PSF) of ASCA and BeppoSAX. In particular, using Chandra data, \citet{baldi07} were able to study the temperature profiles of 12 clusters in the redshift range $0.1<z<0.3$, dividing the clusters into CC and NCC clusters in a systematic fashion. These authors found that the profiles in the inner $0.1r_{180}$ were showing a positive gradient $kT(r)\propto r^\mu$, with $\mu\sim0.25$ in CC clusters and $\mu\sim0$ in NCC clusters. Moreover Baldi et al. found that the outer region profiles were significantly steeper in the NCC systems than in the CC systems. This behavior agrees with the recent findings of \citet{arnaud10}, who derived a universal pressure profile of galaxy clusters (including CC and NCC clusters) in the REXCESS local sample ($z<0.2$) with low dispersion, especially in the external regions. Steeper temperature gradients in the profiles of NCC systems than in those of CC systems are indeed expected if this pressure profile, analogous for both cluster categories, is combined with the density profiles \citep[steeper in CC and flatter in NCC clusters, see e.g.][]{cavaliere09}. Probing a similar redshift range, \citet{leccardi08a} analyzed a sample of $\approx$50 galaxy clusters observed by {\em XMM-Newton} to measure their radial temperature profiles. In agreement with previous results, these authors found a decline of the temperature in the $0.2-0.6r_{180}$ range. In contrast with \citet{degrandi02} and \citet{baldi07}, Leccardi and Molendi did not find any evidence of a dependence in the slope of the outer regions from the presence or absence of a cool core, nor did they find evidence of profile evolution with the redshift out to $z\approx0.3$. Cosmological simulations have also tackled the measure of temperature profiles in galaxy clusters. However, the existing simulations \citep[e.g.][]{loken02,roncarelli06} are not able to account for the physics at small radii (i.e. in the cluster core regions) and are not considering the evolution of the profiles through cosmic time. In this work we aim at studying for the first time the temperature profiles in a sample of galaxy clusters at $z>0.4$, i.e. in a redshift range where no comprehensive study is present in the literature. The plan of the paper is as follows. In \S~2 we describe how we selected our {\em XMM-Newton} sample and the data reduction procedure we followed. We describe our spectral analysis strategy in \S~3, with particular focus on the background treatment procedure (\S~3.1) and the correction we applied to take into account the PSF of {\em XMM-Newton} (\S~3.2). In \S~4 we present the results obtained in our analysis, presenting the temperature profiles of individual clusters in the sample and the classification method we used to divide them into CC and NCC (\S~4.1), their profiles by the overdensity radius $r_{500}$ and by the average temperature $T_{500}$, and the properties of self-similarity they exhibit (\S~4.2). We also present several fits of the normalized profiles of the total sample (\S~4.2.1), of the CC and NCC clusters samples individually (\S~4.2.2), and of the external regions of the profiles (\S~4.2.3). In \S~5, we discuss our results and compare our sample with the \citet{leccardi08a} intermediate redshift cluster sample (\S~5.1) to attempt to define a universal law for temperature profiles for galaxy clusters at $0.1<z<0.9$ (\S~5.2). Our conclusions are summarized in \S~6. We adopt a cosmological model with $H_0=70$ km/s/Mpc, $\Omega_m=0.3$, and $\Omega_\Lambda=0.7$ throughout the paper. Confidence intervals are quoted at $1\sigma$ unless otherwise stated. \begin{figure} \includegraphics[width=9.2cm, angle=0]{histo_z.ps} \caption{Redshift distribution of the {\em XMM-Newton} galaxy cluster bright sample presented in this paper.} \label{histoz} \end{figure} \begin{table*} \caption{{\em XMM-Newton} archival observations of the 12 bright galaxy clusters analyzed in this paper.\label{exposures}} \centering \begin{tabular}{lcrrccccc} \hline Cluster & $z$ & Obs. Date & Obs. ID & $t_{clean,MOS1}$ & $t_{clean,MOS2}$ & $t_{clean,pn}$ & $N_H$ & No. of radial bins\\ & & & & (ksec) & (ksec) & (ksec) & (10$^{20}$ cm$^{-2}$) & considered \\ \hline\hline \object{A851} & 0.407 & 2000 Nov 06 & 0106460101 & 41.7 & 41.1 & 30.3 & 1.0 & 7 \\ \hline \object{RXCJ0856.1+3756} & 0.411 & 2005 Oct 10 & 0302581801 & 24.2 & 23.9 & 14.7 & 3.2 & 4 \\ \hline \object{RXJ2228.6+2037} & 0.412 & 2003 Nov 18 & 0147890101 & 24.6 & 24.4 & 19.3 & 4.3 & 9 \\ \hline \object{RXCJ1206.2-0848} & 0.440 & 2007 Dec 09 & 0502430401 & 29.1 & 28.8 & 20.6 & 3.7 & 8 \\ \hline \object{IRAS09104+4109} & 0.442 & 2003 Apr 27 & 0147671001 & 12.1 & 12.3 & 8.3 & 1.4 & 4 \\ \hline \object{RXJ1347.5-1145} & 0.451 & 2002 Jul 31 & 0112960101 & 32.6 & 32.4 & 27.5 & 4.9 & 9 \\ \hline \object{CLJ0030+2618} & 0.500 & 2005 Jul 06 & 0302581101 & 15.0 & 14.0 & - & 3.7 & 3 \\ & & 2006 Jul 27 & 0402750201 & 27.0 & 27.4 & 19.6 & & \\ & & 2006 Dec 19 & 0402750601 & 28.5 & 28.6 & 20.6 & & \\ \hline \object{MS0015.9+1609} & 0.541 & 2000 Dec 29 & 0111000101 & 30.6 & 30.1 & 20.4 & 4.0 & 7 \\ & & 2000 Dec 30 & 0111000201 & 5.5 & 5.4 & - & & \\ \hline \object{MS0451.6-0305} & 0.550 & 2004 Sep 16 & 0205670101 & 25.5 & 26.1 & 19.4 & 3.9 & 5 \\ \hline \object{MACSJ0647.7+7015} & 0.591 & 2008 Oct 09 & 0551850401 & 52.9 & 53.7 & 32.3 & 5.4 & 6 \\ & & 2009 Mar 04 & 0551851301 & 33.2 & 34.6 & 18.4 & & \\ \hline \object{MACSJ0744.9+3927} & 0.698 & 2008 Oct 17 & 0551850101 & 40.4 & 41.4 & 22.4 & 5.7 & 6 \\ & & 2009 Mar 21 & 0551851201 & 63.8 & 67.0 & 35.8 & & \\ \hline \object{CLJ1226.9+3332} & 0.890 & 2001 Jun 18 & 0070340501 & 10.3 & 10.7 & 5.3 & 1.8 & 4 \\ & & 2004 Jun 02 & 0200340101 & 67.2 & 67.5 & 53.4 & & \\ \hline \end{tabular} \end{table*}
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We analyzed an {\em XMM-Newton} sample of 12 bright ($L_X>4\times10^{44}$ erg s$^{-1}$) galaxy clusters in the redshift range $0.4<z<0.9$ with an average temperature $kT>4.5$~keV. This sample was extracted from the {\em XMM-Newton} sample analyzed in BAL12 by selecting all clusters with at least 3000 {\tt MOS} net counts to obtain radial temperature profiles with a sizeable number of radial bins and reasonable errors on the temperature ($\sigma_{kT}/kT<15\%$). Taking advantage of EPIC {\em XMM-Newton}'s high throughput and effective area, which makes it an ideal instrument for performing a spatially resolved spectral analysis, this paper aimed at a systematic study of the temperature profiles in galaxy clusters at $z>0.4$, which is not currently provided in the literature. The results we obtained can be summarized as follows: \begin{itemize} \item We extracted temperature profiles for the 12 clusters in our sample. The cluster extension ranged from $\sim500$~kpc to $\sim1.3$~Mpc from the center. All profiles were found to be declining toward larger radii.\\ \item The temperature profiles of the galaxy clusters in our sample, normalized by the mean temperature $T_{500}$, were found to be generally self-similar and could be well described by a function obtained by adapting the relation of \citet{vikhlinin06b} derived for lower redshift clusters.\\ \item We divided the sample into five CC and seven NCC clusters by introducing a pseudo-entropy ratio $\sigma$ and defining a threshold $\sigma=0.6$ between CC and NCC clusters with the latter having $\sigma\ge0.6$. The profiles of the two subsamples were found to be different mainly in the inner regions, with the inner slope parameter assuming a positive value in the CC clusters ($\alpha=1.14\pm0.59$), to fit the temperature drop in the center, and a value consistent with zero in the NCC clusters ($\alpha=0.18\pm0.66$), to fit their flat central profile. The large errors on the measurement of the external slope $\beta_0$ gave inconclusive results on the differences between the samples. Fitting the external regions ($r>0.15r_{500}$) with a simpler function yielded no significant difference between the slopes of CC and NCC clusters ($\beta_0=0.39\pm0.12$ and $\beta_0=0.37\pm0.06$, respectively). The lack of any significant difference between the two samples could be attributed to the small sample size, and therefore to the few data points available for the fit, especially for the CC clusters.\\ \item We introduced a function of both $r$ and $\sigma$ to fit the data points of CC and NCC clusters together. In this case, the improved statistics allowed us to detect a significant dependence of the temperature profiles on the pseudo-entropy ratio $\sigma$, showing an indication that the outer part of the profiles becomes steeper for higher values of $\sigma$ (i.e. transitioning toward the NCC clusters). This behavior would agree with the results obtained in lower redshift galaxy cluster samples \citep[e.g.][]{baldi07} and with the universal pressure profile derived by \citet{arnaud10} in galaxy clusters in the REXCESS local sample ($z<0.2$), which show a low dispersion, especially in the external regions.\\ \item In all attempted fits, no evidence of redshift evolution could be found within the redshift range sampled by our clusters ($0.4<z<0.9$).\\ \item We compared our sample with the intermediate cluster sample of \citet{leccardi08a} at $0.1<z<0.3$, finding significant differences in both the CC and the NCC cluster samples. In particular, we found that CC clusters at $z>0.4$ shows a temperature dip in the center less deep with respect to CC clusters at $z<0.3$, while NCC clusters at $z>0.4$ showed steeper temperature profiles with respect to NCC clusters at $z<0.3$. This can be caused by the circumstance that higher $z$ clusters are at a less advanced stage of their formation and did not have enough time to create a large massive hot gas halo comparable in size and mass with that of lower redshift clusters.\\ \item We defined for the first time the closest possible relation to a universal law for the temperature profiles of galaxy clusters at $0.1<z<0.9$. This relation shows a dependence on the state of relaxation of the clusters and the redshift. \end{itemize} Although some of the results obtained in this paper can in principle be biased by possible selection effects introduced in extracting the sample from the XMM archive (\S~\ref{sampsel}), we stress that this is the first time that a systematic study of the temperature profiles in galaxy clusters at $z>0.4$ has been attempted. Additional deep XMM-Newton and Chandra observations, and most likely a new generation of X-ray observatories, would be needed to improve the current knowledge of the temperature distribution in the hot gas of galaxy clusters at high redshift.
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1206.6603
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1206
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1206.4789_arXiv.txt
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{We present the analysis of supernova remnants (SNRs) and candidates in M\,31 identified in the \xmm\ large programme survey of M\,31. SNRs are among the bright X-ray sources in a galaxy. They are good indicators of recent star formation activities of a galaxy and of the interstellar environment in which they evolve.} {By combining the X-ray data of sources in M\,31 with optical data as well as with optical and radio catalogues, we aim to compile a complete, revised list of SNRs emitting X-rays in M\,31 detected with \xmm, study their luminosity and spatial distribution, and understand the X-ray spectrum of the brightest SNRs. } {We analysed the X-ray spectra of the twelve brightest SNRs and candidates using \xmm\ data. The four brightest sources allowed us to perform a more detailed spectral analysis and the comparison of different models to describe their spectrum. For all M\,31 large programme sources we searched for optical counterparts on the \ha, [\sii], and [\oiii] images of the Local Group Galaxy Survey.} {We confirm 21 X-ray sources as counterparts of known SNRs. In addition, we identify five new X-ray sources as X-ray and optically emitting SNRs. Seventeen sources are no longer considered as SNR candidates. We have thus created a list of 26 X-ray SNRs and 20 candidates in M\,31 based on their X-ray, optical, and radio emission, which is the most recent complete list of X-ray SNRs in M\,31. The brightest SNRs have X-ray luminosities of up to $8 \times 10^{36}$~erg~s$^{-1}$ in the 0.35 -- 2.0~keV band. } {}
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The Andromeda galaxy (M\,31) is the largest galaxy in the Local Group and the nearest spiral galaxy to the Milky Way. Its size and mass are comparable to those of our Galaxy. Therefore, this archetypical spiral galaxy provides us a unique opportunity to study and understand the nature and the evolution of a galaxy like our own. Various authors have studied the star formation history in different regions of M\,31 using observations with both the \hubble\ Space Telescope (\hst) and large ground-based telescopes (e.g., the Local Group Galaxy Survey [LGGS] performed at the Kitt Peak National Observatory [KPNO] and the Cerro Tololo Inter-American Observatory, Williams \citeyear{2003AJ....126.1312W}; Massey et al.\ \citeyear{2006AJ....131.2478M}). Deep {\sl HST} photometry has shown that the mean age of the disk of M\,31 is $\sim6 - 8$~Gyr \citep{2006ApJ...652..323B} and that the average metallicity is [Fe/H] $\approx$ --0.6~dex \citep{2003A&A...405..867B}. \citet{2003AJ....126.1312W} measured a mean star formation rate of about 1~$M_{\sun}$~yr$^{-1}$ in the full disk of M\,31. \citet{2010A&A...517A..77T} have recently studied the dust distribution and computed the de-reddened \ha\ distribution in the disk of M\,31. They derived a star formation rate of 0.27~$M_{\sun}$~yr$^{-1}$ for the radial range of $6 < R < 17$~kpc with an increase to about twice the mean value at around $R = 10$~kpc. Although the current star formation rate in M\,31 is lower than that of the Milky Way, M\,31 seems to have undergone more active star formation periods. In addition to the well-known dust ring at a radius of $\sim$10~kpc \citep{1984A&AS...55..179B,1993ApJ...418..730D} with enhanced star formation, \citet{2006Natur.443..832B} found an inner dust ring with a radius of 1 -- 1.5~kpc, which has apparently been created in an encounter with a companion galaxy, most likely with M\,32. First observations of individual sources in M\,31 in X-rays were performed with the {\sl Einstein} Observatory \citep{1979ApJ...230..540G} in the energy band of 0.2 -- 4.5~keV and yielded the first catalogues of X-ray sources in the field of M\,31 \citep{1979ApJ...234L..45V,1991ApJ...382...82T}. In the 1990s, {\sl ROSAT} \citep{1982AdSpR...2..241T} observed M\,31 in the 0.1 -- 2.4~keV band and revealed a total of 560 sources in the field of M\,31 (Supper et al.\ \citeyear{1997A&A...317..328S}, hereafter SHP97; Supper et al.\ \citeyear{2001A&A...373...63S}, hereafter SHL01). The next generation X-ray satellites \chandra\ X-ray Observatory \citep{2002PASP..114....1W} and X-ray Multi-Mirror Mission \citep[\xmm,][]{2001A&A...365L...1J}, which were both launched in 1999, have significantly improved spatial and spectral resolutions with respect to the prior X-ray telescopes. They have also performed several observations of M\,31 and allowed both to obtain a comprehensive list of X-ray sources and to study individual sources \citep[e.g.,][]{2001A&A...378..800O,2002ApJ...577..738K, 2002ApJ...578..114K,2005A&A...434..483P,2008A&A...480..599S, 2008ApJ...689.1215B}. The entire galaxy M\,31 was observed by \xmm\ in a large programme (LP) between June 2006 and February 2008 with the European Photon Imaging Cameras \citep[EPICs,][]{2001A&A...365L..18S,2001A&A...365L..27T} as prime instruments. The \xmm\ source catalogue with 1897 sources has been published by Stiele et al.\ (\citeyear{2011A&A...534A..55S}, hereafter SPH11). Supernova remnants are the aftermath of stellar explosions releasing a large amount of energy in galaxies. The spherically expanding blast wave shock produces a cavity in the interstellar medium with a very low-density, high-temperature interior, which predominantly emits soft X-ray radiation. In addition, relativistic electrons and heavier particles in SNRs emit synchrotron emission which can be detected in radio and in some cases also in X-rays. After a few thousand years, the SNR becomes radiative, i.e., the radiative losses in SNR shocks expanding into the ambient ISM become non-negligible and the shell emits energy as UV and optical line emission. If a neutron star is created in the supernova explosion a pulsar and/or a pulsar wind nebula (PWN) can be found inside the SNR, in which particles are accelerated in the strong magnetic field of the neutron star and thus non-thermal emission is produced. SNRs in M\,31 were mainly detected in the optical \citep[e.g.,][]{1980A&AS...40...67D,1981AJ.....86..989D,1981ApJ...247..879B} and in combined optical and radio studies \citep{1993A&AS...98..327B}. The X-ray survey performed with {\sl ROSAT} led to the detection of 16 X-ray SNRs \citep{2001A&A...373...63S}, while 21 were detected and identified with \xmm\ \citep{2005A&A...434..483P}. \citet{2002ApJ...580L.125K} presented the first resolved X-ray image of an SNR in M\,31 taken with \chandra, while \citet{2003ApJ...590L..21K} and \citet{2004ApJ...615..720W} reported on the discovery of new SNRs in M\,31 based on \chandra\ data. In this paper, we present the study of all X-ray sources that have been suggested to be SNRs or candidates of SNRs in M\,31 based on a complete survey of M\,31 performed within the framework of the LP of \xmm\ (SPH11). Through a detailed study of each SNR and candidate detected by \xmm, we have obtained an improved sample of SNRs in M\,31 consisting of X-ray or optically confirmed SNRs as well as bona-fide candidates. We have performed a spectral analysis of the bright SNRs and candidates in the catalogue of SPH11. Using optical data of the LGGS, we have searched for optical counterparts of the X-ray SNRs and candidates showing optical \ha, [\sii], and [\oiii] emission from the radiative shock. We newly calibrated the LGGS data to obtain optical fluxes and computed the [\sii]/\ha\ flux ratio, which is an indicator of SNR emission in the optical. This study has also allowed us to identify sources that are most probably no SNRs. In addition, we have performed statistical studies using the revised list of SNRs and candidates in M\,31 detected with \xmm.
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\subsection{Cumulative X-ray luminosity distribution}\label{cumxlf} \begin{figure} \centering \includegraphics[width=0.5\textwidth,bb=50 0 440 355,clip=]{xlf_0.35_2.0.eps} \caption{ Cumulative luminosity distribution of X-ray SNRs and candidates in M\,31 listed in Tables \ref{snrlist} and \ref{candlist} for the band of 0.35 -- 2.0~keV (solid line). For comparison, the distribution of SNRs in M\,33 is also shown \citep[dash-dotted,][]{2010ApJS..187..495L}. The dotted lines show the power-law fits. The power-law index $\alpha$ is given. } \label{xlf} \end{figure} We estimated the luminosity of the SNRs and candidates listed in Tables \ref{snrlist} and \ref{candlist} and thus obtained a cumulative luminosity distribution of the SNRs in M\,31. For the M\,31 SNRs we converted the \xmm\ count rates into flux by assuming a thermal spectrum with $kT = 0.2$~keV, absorbed by a foreground column density of \nh(MW) $= 0.7 \times 10^{21}$~cm$^{-2}$ and an additional \nh(M31) $= 1.0 \times 10^{21}$~cm$^{-2}$ similar to values of the fits of the four brightest sources [SPH11] 969, 1050, 1066, and 1234. We simulated the flux assuming different models by varying the temperature in the range of $kT = 0.1 - 0.3$~keV and the foreground absorption \nh(M31) $= 0 - 2 \times 10^{21}$~cm$^{-2}$ for the observed count rates and obtained an uncertainty of $\sim$20\%. Assuming CIE or NEI model results in a difference in flux of $\sim$10\% (see Table \ref{10661234}). As most of the sources are too faint to distinguish between CIE and NEI, we assume the same {\tt APEC} model for all sources. For the cumulative luminosity distribution plot shown in Fig.\,\ref{xlf}, the luminosities in the energy band of 0.35 -- 2.0~keV were calculated with the foreground absorption set to zero. We also included the source CXOM31\,J004247.82+411525.7, which was identified as an SNR by \citet{2003ApJ...590L..21K} but not resolved with \xmm\ (see Sect.\,\ref{compchan}). For comparison, the luminosities for M\,33 SNRs taken from \citet{2010ApJS..187..495L} are plotted as well. These luminosities were converted from \chandra\ count rates by \citet{2010ApJS..187..495L} assuming a thermal plasma model with $kT = 0.6$~keV and an absorbing \nh\ $= 5.0 \times 10^{20}$~cm$^{-2}$, which correspond to the best fit values for the brightest SNRs in M\,33 observed with \chandra. As one can see in Fig.\,\ref{xlf} the slope of the cumulative luminosity distribution of X-ray SNRs in M\,31 and M\,33 are comparable. Both distributions can be fitted with a power law with an index of $\alpha \approx -1$. The distribution in M\,33 seems to deviate from this power-law distribution for luminosities $> 5 \times 10^{35}$~erg~s$^{-1}$ showing an excess. The slight difference in the shape of the cumulative luminosity distribution might indicate that the fraction of more luminous SNRs in M\,33 is higher than in M\,31. There are 24 sources brighter than $5 \times 10^{35}$~erg~s$^{-1}$ in M\,31 and nine in M\,33. Above $10^{36}$~erg~s$^{-1}$, there are 13 sources in M\,31 and seven in M\,33 \citep[see also][]{2005AJ....130..539G,2010ApJS..187..495L}. The number of SNRs in M\,33 is higher than what is expected if we simply scale with the total mass of the galaxy, as M\,33 is about 10 times less massive than M\,31 \citep{2003MNRAS.342..199C,2010MNRAS.406..264W}. M\,33 is a typical flocculent spiral galaxy with discontinuous spiral arms, in which star formation regions are found. In contrast to grand design spirals such as the Milky Way or M\,31, in which density waves are believed to produce the spiral arms, gravitational instabilities together with turbulence in the interstellar medium seem to be the origin of the spiral structure and thus the on-going star formation in flocculent galaxies \citep[and references therein]{1984ApJ...282...61S,2003ApJ...593..333E}. The star formation rate in the disk of M\,31 is 0.27~$M_{\sun}$~yr$^{-1}$ for $6 < R < 17$~kpc \citep{2010A&A...517A..77T} corresponding to a star formation rate per unit area of $\Sigma_{\rm SFR} = 0.4 M_{\sun}$~Gyr$^{-1}$~pc$^{-2}$. This value is about six times lower than in M\,33, for which a star formation rate per unit area of $\Sigma_{\rm SFR} = 2 - 3 M_{\sun}$~Gyr$^{-1}$~pc$^{-2}$ has been measured \citep{2009A&A...493..453V}. The higher star formation rate implies a higher rate for the occurrence of core-collapse SNRs in M\,33. In addition, one should note that our list of SNRs in M\,31 is based on \xmm\ observations whereas the M\,33 SNRs have been detected in a survey performed with \chandra. Not only was the \chandra\ survey of M\,33 deeper, but the superior angular resolution of \chandra\ made it possible to detect smaller and thus most likely younger SNRs in M\,33, which would not have been resolved and classified as an SNR in an \xmm\ observation. However, the \chandra\ M\,33 survey only observed the inner part of the galaxy inside the $D_{25}$ ellipse, whereas the \xmm\ M\,31 survey fully covered the $D_{25}$ ellipse. \subsection{Radial distribution} \begin{table*} \caption{ \label{info} Geometric parameters of M\,31 and M\,33 used for the calculation of the radial SNR number density distribution. } \centering \begin{tabular}{cccccc} \hline\hline Galaxy & Distance\tablefootmark{a} [kpc] & Position Angle\tablefootmark{b} & Inclination Angle\tablefootmark{b} & Corrected $D_{25}$\tablefootmark{b} & $R_{25}$ [kpc] \\ \hline M\,31 & 744 & 35\degr\ & 71\degr\ & 155\farcm5 & 16.8 \\ M\,33 & 805 & 23\degr\ & 54\degr\ & 52\farcm6 & 6.2 \\ \hline \end{tabular} \tablefoot{ \tablefoottext{a}{\citet{2010A&A...509A..70V} for M\,31 and \citet{2009MNRAS.396.1287S} for M\,33.} \tablefoottext{b}{\citet{1991trcb.book.....D}.} } \end{table*} In the shock waves of SNRs, particles can gain energies up to 10$^{15}$~eV or higher due to diffusive shock acceleration \citep[][] {1978MNRAS.182..147B,1978MNRAS.182..443B,1978ApJ...221L..29B,2005JPhG...31R..95H}. Therefore, SNRs together with pulsars are thought to be the primary sources of Galactic cosmic rays (CRs). The distribution of SNRs and pulsars in our Galaxy is a crucial basis for the understanding of the CR distribution. \citet{1977ApJ...217..843S} studied the radial dependence of the Galactic SNR surface density using observational SNR data of \citet{1974PASJ...26..255K} and pulsar data of \citet{1977NASCP.002..265S} and showed that it can be described as $\propto x^{\alpha} \exp{(-x/R)}$ with $x$ being the radial distance to the Galactic centre. Using only Galactic shell-type SNRs, \citet{1989PASP..101..607L} found a peak in the surface density distribution at 4 -- 6~kpc distance from the Galactic centre. \citet{1998ApJ...504..761C} re-analysed the Galactic SNR data and suggested to use a dependence of the type $\propto sin(\pi x + \theta) \exp{(-x/R)}$ for the radial surface density distribution. They obtained a scale length of $\sim7$~kpc and a maximum of the distribution at about 5~kpc. Based on the obtained distribution of SNRs in our Galaxy, the spectral distribution of Galactic CRs can be modelled to explain the CR spectrum up to 10$^{15}$ -- 10$^{16}$~eV \citep[see, e.g.,][]{2005JPhG...31R..95H}. However, our vantage point is not ideal to study the source distribution in our Galaxy. To understand the distribution of SNRs in a spiral galaxy, it is thus necessary to study the most nearby spiral galaxies M\,31 and M\,33. \begin{figure} \centering \includegraphics[width=0.4\textwidth,bb=63 0 435 354,clip=]{radial_m31_X_snrs.eps} \\[2mm] \includegraphics[width=0.4\textwidth,bb=63 0 435 354,clip=]{radial_m33_X_snrs.eps} \caption{ Surface density of SNRs and candidates in M\,31 (this work) and M\,33 \citep{2010ApJS..187..495L} plotted over the radius normalised to $D_{25}/2$ \citep[$D_{25}$ = 155.5\arcmin, 52.6\arcmin\ for M\,31 and M\,33, respectively,][]{1991trcb.book.....D}. The radial distance in kpc is given along the upper x-axis. Dotted lines show the fitted model function $f(x) = C x^{\alpha} \exp{(-\beta x)}$, dashed lines show the model function for the Milky Way normalised to the maximum of M\,31 or M\,33. } \label{radial} \end{figure} In order to obtain the radial distribution of SNRs in M\,31, the positions of the SNRs and candidates in Tables \ref{snrlist} and \ref{candlist} as well as the source CXOM31\,J004247.82+411525.7 were first corrected for projection and their galactocentric distances were computed. \citet{2011A&A...534A..55S} have also presented the radial distribution of SNRs and candidates in M\,31 detected in the \xmm\ LP. For comparing M\,31 and M\,33, the radial distances were normalised to $R_{25} = D_{25}/2$. The positions of sources were binned into equidistant radial bins and the surface density was calculated for each annulus. The parameters used, i.e., distance, inclination angle, position angle, and $D_{25}$ are listed in Table \ref{info}. The SNR surface densities are plotted against the normalised radial distance in Fig.\,\ref{radial}. We fitted the obtained radial surface density distribution of the SNRs in each galaxy with the model distribution introduced for the radio selected SNRs in the Milky Way \begin{equation}\label{xexp} f(x) = C\,x^{\alpha} \exp{(-x/R)} \end{equation} with a maximum at a radial distance of several kpc from the Galactic centre as originally suggested by \citet{1977ApJ...217..843S}. The parameters of the distribution obtained by \citet{1998ApJ...504..761C} for our Galaxy are $C = 136.5, \alpha = 2.00$, and $R = 0.14$. The best-fit model curves according to Eq.\,\ref{xexp} are plotted in Fig.\,\ref{radial} with dotted lines. This radial SNR distribution in M\,31 and M\,33 is different from the mass distribution of these galaxies derived from the rotation curves \citep[e.g.,][and references therein]{2003MNRAS.342..199C,2010A&A...511A..89C}. The SNR distribution rather seems to follow the distribution of stars and gas, similar to what had been suggested for the Milky Way. For comparison, we also plot the distribution for the Milky Way normalised to the fitted maxima of the distributions in M\,31 and M\,33. The fitted curve indicates a maximum at about 5.5~kpc and 2~kpc for M\,31 and M\,33, respectively, corresponding to $\sim$0.3~$R_{25}$ for both galaxies. The distribution in M\,31 seems to be almost flat for $<$ 17~kpc $\approx R_{25}$ and falls exponentially outside $\sim1.0~R_{25}$, while the distribution in M\,33 falls off exponentially for $>$ 4~kpc $\approx 0.65~R_{25}$. This behaviour of the SNR distribution in M\,31 seems to be correlated with the distribution of gas in M\,31, which is known to have ring-like structures consisting of many spiral arms between a radius of $\sim5$~kpc to 20~kpc, with the most prominent ring found at a radius of $\sim10$~kpc \citep{1981PASJ...33..449S,1991ApJ...372...54B}.
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1206.4789
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1206
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1206.1353_arXiv.txt
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{\it Hubble Space Telescope} spectroscopy of the Seyfert 1.5 galaxy, NGC 3227, confirms previous reports that the broad H${\alpha}$ emission line flux is time variable, decreasing by a modest ${\sim}$ 11\% between 1999 and 2000 in response to a corresponding ${\sim}$ 37\% decrease in the underlying continuum. Modeling the gas distribution responsible for the broad H${\alpha}$, H${\beta}$ and H${\gamma}$ emission lines favors a spherically symmetric inflow as opposed to a thin disk. Adopting a central black hole mass of 7.6 ${\times}$ 10$^{6}$ M${_{\sun}}$, determined from prior reverberation mapping, leads to the following dimensions for the size of the region emitting the broad H${\alpha}$ line; an outer radius ${\sim}$ 90 l.d and an inner radius ${\sim}$ 3 l.d. Thus, the previously determined reverberation size for the broad line region (BLR) consistently coincides with the inner radius of a much {\it larger} volume of ionized gas. However, the {\it perceived} size of the BLR is an illusion, a consequence of the fact that the emitting region is ionization bounded at the outer radius and diminished by Doppler broadening at the inner radius. The actual dimensions of the inflow remain to be determined. Nevertheless, the steady state mass inflow rate is estimated to be ${\sim}10^{-2} ~{\rm M_\odot~yr^{-1}}$ which is sufficient to explain the X-ray luminosity of the AGN in terms of radiatively inefficient accretion. Collectively, the results challenge many preconceived notions concerning the nature of BLRs in active galactic nuclei.
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NGC 3227 is an early type SAB(s)a spiral galaxy located at a distance of 20.8 Mpc \citep{Tul88}. The galaxy harbors a Seyfert 1.5 nucleus \citep{ Ho97} whose broad Balmer emission lines have long been known to be time variable. Since the broad line region (BLR) is unresolved, the spatial distribution of emission cannot be directly measured. However, reverberation mapping, which refers to measuring the time delay in the response of broad emission lines to time variable illumination from the central continuum source \citep{Pet93}, has yielded a variety of estimates for the size of the broad line region \citep{Pet85, Win95, Pet04, Kas05, Den09, Den10}. Even though the broad line emission emanates from a finite volume, reverberation mapping yields only a single size. Despite numerous attempts to model the reverberation phenomenon \citep[e.g.,][and references therein]{Ede88, Rob90, Wel91, Hor04, Kor04, Pan11}, the question remains {\it what exactly does the reverberation size refer to; the inner radius of the BLR, the outer radius or a luminosity weighted radius?} Emission line profile fitting is a complementary technique used to constrain the physical dimensions of the BLR in low luminosity AGNs (LLAGNs) with known central masses. Gravity dominates the kinematics of the gas because LLAGNs radiate well below the Eddington limit with insufficient radiation pressure to drive an outflow. Consequently, knowing the relationship between velocity and radius allows the broad emission line profiles to be modeled revealing the shape and size of the BLR. This technique has led to large BLR size discrepancies for the LLAGNs M81, NGC 3998 and NGC 4203 \citep{Dev07, Dev11a, Dev11b}. Specifically, the dimensions of the BLR deduced from emission line profile fitting are much larger than expected based on the reverberation size--UV luminosity correlation of \cite{Kas05}. However, \cite{Kas05} note that the correlation appears to break down for AGNs with low UV luminosities, comparable to those measured for M81, NGC 3998 and NGC 4203, so the comparison may not be meaningful. On the other hand, the reverberation size has actually been measured for the LLAGN in NGC 3227, most recently by \cite{Den10}. The principle aim of this paper, therefore, is to make the first direct comparison of the reverberation size with the dimensions of the BLR in NGC 3227 deduced from emission line profile fitting. The larger context for this investigation is to understand why it is that BH masses estimated using the reverberation radius require a considerable factor of ${\sim}$ 5.5 correction in order to place them on the M$_{\bullet}$--${\sigma}_* $ relation defined by BH masses measured directly using gas and stellar kinematics \citep{Onk04}. The mass of the BH in NGC 3227 has been measured using three different methods; stellar kinematics \citep{Dav06}, gas kinematics \citep{Hic08} and most recently via reverberation mapping \citep{Den10}, collectively yielding mass estimates in the range (0.7 -- 2) ${\times}$ 10${^7}$ M$_{\odot}$. When combined with the X-ray luminosity the range of BH masses indicate that NGC 3227 is radiating at ${\leq}$0.4\% of the Eddington luminosity limit \citep{Win11, Xu11, Vas09}. The layout of the paper is as follows. In Section 3, the broad emission lines seen in NGC 3227 are evaluated in the context of inflow and accretion disk models. Some physical properties of the BLR are presented in Section 4 including the size, structure and source of ionization. Conclusions follow in Section 5. We begin, however, with Section 2 and a review of the emission lines observed in the nucleus of NGC 3227.
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Spectroscopic observations of NGC 3227 with the {\it Hubble Space Telescope} have revealed a time variable single-peak broad H${\alpha}$ emission line profile which has been successfully modeled as a steady state inflow. The {\it perceived} dimensions of the BLR correspond to an outer radius of ${\sim}$90 l.d. and an inner radius of ${\sim}$3 l.d. that coincides with the gas reverberation radius. However, the {\it perceived} small size for the BLR is an illusion, a consequence of the fact that the emitting region is ionization bounded at the outer radius and diminished by Doppler broadening at the inner radius. The actual dimensions of the inflow are likely much larger and remain to be determined. If the electron density is high, ${n_e}$ ${\geq}$ 10$^6$ cm$^{-3}$, as suggested by the absence of similarly broad [O I] and [O III] emission lines, then the luminosity in the broad H${\alpha}$ emission line observed in the year 2000 leads to a steady state inflow rate ${\sim} 10^{-2} ~{\rm M_\odot~yr^{-1}}$ which is sufficient to explain the X-ray luminosity of the AGN, if the AGN is powered by radiatively inefficient accretion.
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1206.1353
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1206
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1206.5957_arXiv.txt
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The extremely efficient process of resonant Compton upscattering by relativistic electrons in high magnetic fields is believed to be a leading emission mechanism of high field pulsars and magnetars in the production of intense X-ray radiation. New analytic developments for the Compton scattering cross section using Sokolov \& Ternov (S\&T) states with spin-dependent resonant widths are presented. These new results display significant numerical departures from both the traditional cross section using spin-averaged widths, and also from the spin-dependent cross section that employs the Johnson \& Lippmann (J\&L) basis states, thereby motivating the astrophysical deployment of this updated resonant Compton formulation. Useful approximate analytic forms for the cross section in the cyclotron resonance are developed for S\&T basis states. These calculations are applied to an inner magnetospheric model of the hard X-ray spectral tails in magnetars, recently detected by RXTE and INTEGRAL. Relativistic electrons cool rapidly near the stellar surface in the presence of intense baths of thermal X-ray photons. We present resonant Compton cooling rates for electrons, and the resulting photon spectra at various magnetospheric locales, for magnetic fields above the quantum critical value. These demonstrate how this scattering mechanism has the potential to produce the characteristically flat spectral tails observed in magnetars.
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1206.5957
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1206
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1206.0269_arXiv.txt
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We report the discovery of two new halo velocity groups (Cancer groups A and B) traced by 8 distant RR Lyrae stars and observed by the Palomar Transient Factory (PTF) survey at ${\rm R.A}\sim129\arcdeg$, ${\rm Dec}\sim20\arcdeg$ (${\rm l}\sim205\arcdeg$, ${\rm b}\sim32\arcdeg$). Located at 92 kpc from the Galactic center (86 kpc from the Sun), these are some of the most distant substructures in the Galactic halo known to date. Follow-up spectroscopic observations with the Palomar Observatory 5.1-m Hale telescope and W. M. Keck Observatory 10-m Keck I telescope indicate that the two groups are moving away from the Galaxy at $\bar{v}^A_{gsr} = 78.0\pm5.6$ {\kms} (Cancer group A) and $\bar{v}^B_{gsr} = 16.3\pm7.1$ {\kms} (Cancer group B). The groups have velocity dispersions of $\sigma^A_{v_{gsr}}=12.4\pm5.0$ {\kms} and $\sigma^B_{v_{gsr}}=14.9\pm6.2$ {\kms}, and are spatially extended (about several kpc) making it very unlikely that they are bound systems, and are more likely to be debris of tidally disrupted dwarf galaxies or globular clusters. Both groups are metal-poor (median metallicities of ${\rm [Fe/H]^A = -1.6}$ dex and ${\rm [Fe/H]^B=-2.1}$ dex), and have a somewhat uncertain (due to small sample size) metallicity dispersion of $\sim0.4$ dex, suggesting dwarf galaxies as progenitors. Two additional RR Lyrae stars with velocities consistent with those of the Cancer groups have been observed $\sim25\arcdeg$ east, suggesting possible extension of the groups in that direction.
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} In the last decade, deep and wide-area astronomical surveys, such as the Sloan Digital Sky Survey (SDSS; \citealt{yor00}) and Two-Micron All Sky Survey, (2MASS; \citealt{skr06}) have enabled detection of substructures in the halo both as stellar overdensities in space and as moving groups. The tidal streams of the disrupting Sagittarius dwarf galaxy \citep{igi94} are the best example of such substructures, with streams wrapping around most of the sky \citep{maj03}. Other known substructures include the Pisces Overdensity (``Clump J'' in \citealt{ses07}; \citealt{wat09}; \citealt{ses10a}), the GD$-$1 stream \citep{gd06}, the Orphan stream \citep{gri06, bel07}, and several other overdensities and streams\footnote{A fairly extensive list of halo streams and overdensities can be found at \url{http://homepages.rpi.edu/~newbeh/mwstructure/MilkyWaySpheroidSubstructure.html.}}. While many halo substructures have been discovered so far, the motivation for finding and characterizing more of them remains strong. For example, at larger galactocentric distances ($R_{\rm GC}\gtrsim15$ kpc), the Galactic gravitational potential is dominated by dark matter and observations of more distant tidal streams can be used to constrain the shape, orientation, and mass of the Milky Way's dark matter halo (e.g., \citealt{lm10}, \citealt{wat10}). Other motivations for completing the census of halo substructures are to better constrain the severity of the so-called ``missing satellite problem'' \citep{kly99, moo99} and the contribution of the accretion of dwarf satellite galaxies to the formation of the Galactic halo. Finding distant halo substructures, however, is not an easy task. One approach to this problem is to look for spatial groups of (non-variable) sources that are consistent in color-magnitude space with, for example, an old population of stars at a fixed distance (e.g., \citealt{wil05, gri06, bel07}). While this technique has been employed quite successfully in the past, its strongest limitation when looking for distant halo substructures using ground-based imaging data is the separation of stars from galaxies at faint magnitudes. At magnitudes fainter than $r\sim21$ (corresponding to a heliocentric distance of $\sim30$ kpc for metal-poor main-sequence turn-off stars), the number of field Milky Way stars per unit magnitude decreases with increasing magnitude (i.e., the number density profile of Galactic halo stars steepens with distance from the Galactic center from a power-law with index $n=-2.6$ to that with $n\lesssim-3.8$; \citealt{sji11, dbe11}), while the number of galaxies per unit magnitude increases. In addition, morphological separation of stars and galaxies (e.g., SDSS resolved-unresolved source classification; \citealt{lup02}) becomes increasingly unreliable towards faint magnitudes as the signal-to-noise ratio decreases. The increasing contribution of galaxies in point-source catalogs at faint magnitudes causes a sharp increase in the noise from which the signal of a true low surface brightness stellar system must be detected. Instead of looking for clustering in samples of non-variable tracers, which may get contaminated by galaxies at faint magnitudes, we should look for clustering in samples of variable tracers with distinctive light curves (e.g., RR Lyrae stars) that are difficult to confuse with galaxies. RR Lyrae stars have several advantages when used to map the Galactic halo because they i) are old, evolved stars and therefore trace old stellar populations \citep{smi95}; ii) have distinct light curves which make them easy to identify given multi-epoch observations (peak-to-peak amplitudes of $V\sim1$ mag and periods of $\sim 0.6$ days); and iii) are bright, standard candles ($M_V=0.6$ mag at ${\rm [Fe/H]}=-1.5$ dex, with $\sim7\%$ uncertainty in distance) that can be detected at large distances (5-120 kpc for $14 < V < 21$). The steepening of the stellar density profile beyond 30 kpc is actually beneficial for searches that employ RR Lyrae stars because it reduces the pool of stars that can form false spatial groups, and therefore increases the contrast between stars associated with halo substructures and stars associated with the smooth halo (e.g., the Pisces Overdensity shown in Figure 11 of \citealt{ses10a}). Due to the steepening of the stellar density profile beyond 30 kpc, groups of distant RR Lyrae stars are more likely to be real halo substructures than chance associations of stars. While the probability of chance association at large distances and small angular scales is very small, it is not zero and spectroscopic followup is still needed to test whether the stars in a spatial group also form a moving group. In this paper we use the line of reasoning presented above, and successfully used by \citet{kol09} and \citet{ses10b}, as a motivation to study a group of 8 RR Lyrae stars found during a preliminary search for halo substructures in regions of the sky observed by the Palomar Transient Factory survey\footnote{\url{http://www.astro.caltech.edu/ptf}}. The data set and the algorithm used to select RR Lyrae stars are described in Sections~\ref{database} and~\ref{selection}. The spectroscopic observations, their reduction, and measurement of line-of-sight velocities and metallicities are described in Section~\ref{spectro_obs}. The results are discussed in Section~\ref{discussion} and our conclusions are presented in Section~\ref{conclusions}.
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} We confirm the existence of two kinematic groups in the direction of the Cancer constellation (${\rm R.A}\sim129\arcdeg$ and ${\rm Dec}\sim20\arcdeg$, or ${\rm l}\sim205\arcdeg$ and ${\rm b}\sim32\arcdeg$), located at 92 kpc from the Galactic center (86 kpc from the Sun). These groups, tentatively named Cancer groups A and B, are moving at $\bar{v}^A_{gsr} = 78.0\pm5.7$ {\kms} (Cancer group A) and $\bar{v}^B_{gsr} = 16.3\pm3.8$ {\kms} (Cancer group B). The groups have velocity dispersions smaller than 15 {\kms}, are spatially extended (about several kpc), metal-poor (median metallicities of ${\rm [Fe/H]^A = -1.6}$ dex and ${\rm [Fe/H]^B=-2.1}$ dex), and have a metallicity spread of $\sim0.4$ dex. These results suggest that the observed groups are debris of tidally disrupted dwarf galaxies, possibly near the apocenters of their orbits. Whether these groups are related to known substructures in the Galactic halo is unclear at this point, and answering this question may require extensive orbit modeling and comparisons with simulations (e.g., as done by \citealt{joh12} and \citealt{car12}). Observations of two RR Lyrae stars (RR9 and RR10 in Tables~\ref{table-positions} and~\ref{table-data}) obtained after this paper was submitted for peer-review, may help with the modeling of Cancer groups' orbits. The two RR Lyrae stars have velocities of 88 {\kms} and 38 {\kms} that are consistent (to within the uncertainties) with the mean velocities of Cancer groups A and B, and are at similar distances ($\sim76$ kpc from the Sun). They are offset $\sim25\arcdeg$ east of the Cancer groups and may indicate an eastward extension of these groups. However, due to the sparse coverage in PTF of the sky between Cancer groups and RR Lyrae stars RR9 and RR10, we are unable to verify at this point whether RR9 and RR10 are truly related to the Cancer groups or not. Initially, this work was motivated by a hypothesis that distant halo substructures may be found by simply looking for distant spatial groups of RR Lyrae stars. Based on this work and previous work by \citet{kol09} and \citet{ses10b}, we conclude that this is indeed an efficient approach to finding and following-up halo substructures. So far, all of the distant (galactocentric distances greater than 80 kpc) RR Lyrae stars located in spatial groups have proven to be members of a halo substructure (e.g., Cancer groups in this work and RR Lyrae stars in the Pisces Overdensity; \citealt{kol09, ses10b}). Since the density profile of the relatively smooth, inner halo steepens beyond 30 kpc \citep{sji11, dbe11}, it may be that the majority, if not all, of RR Lyrae stars beyond $\sim30$-$40$ kpc are part of some halo substructure. Spectroscopic followup of distant RR Lyrae stars not studied in this work may provide more data to support or refute this hypothesis, and we plan to follow this strategy with other RR Lyrae stars observed by the Palomar Transient Factory.
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1206.0269
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1206
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1206.0928_arXiv.txt
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We have determined the metallicity (O/H) and nitrogen abundance (N/O) of a sample of 122751 Star Forming Galaxies (SFGs) from the Data Release~7 of the Sloan Digital Sky Survey (SDSS). For all these galaxies we have also determined their morphology and obtained a comprehensive picture of their Star Formation History (SFH) using the spectral synthesis code STARLIGHT. The comparison of the chemical abundance with the SFH allows us to describe the chemical evolution of the SFGs in the nearby universe (z $\leq$ 0.25) in a manner which is consistent with the formation of their stellar populations and morphologies. A high fraction (45\%) of the SFGs in our sample show an excess of abundance in nitrogen relative to their metallicity. We also find this excess to be accompanied by a deficiency of oxygen, which suggests that this could be the result of effective starburst winds. However, we find no difference in the mode of star formation of the nitrogen rich and nitrogen poor SFGs. Our analysis suggests they all form their stars through a succession of bursts of star formation extended over a few Gyr period. What produces the chemical differences between these galaxies seems therefore to be the intensity of the bursts: the galaxies with an excess of nitrogen are those that are presently experiencing more intense bursts, or have experienced more intense bursts in their past. We also find evidence relating the chemical evolution process to the formation of the galaxies: the galaxies with an excess of nitrogen are more massive, have more massive bulges and earlier morphologies than those showing no excess. Contrary to expectation, we find no evidence that the starburst wind efficiency decreases with the mass of the galaxies. As a possible explanation we propose that the lost of metals consistent with starburst winds took place during the formation of the galaxies, when their potential wells were still building up, and consequently were weaker than today, making starburst winds more efficient and independent of the final mass of the galaxies. In good agreement with this interpretation, we also find evidence consistent with downsizing, according to which the more massive SFGs formed before the less massive ones.
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\label{sec:1} One of the most important achievements of modern astronomy is the discovery of the process of nucleo-synthesis of chemical elements in stars. Once properly understood, this process, coupled with the concept of stellar evolution regulated by the different masses of stars, gives us a unique insight about the chemical evolution of galaxies \citep{EP78,brodie91,Zaritsky94,Coz98,Coz99,henry00,pilyugin03,pilyugin11,TP11}. For instance, it is now well accepted that oxygen and sulfur are two elements produced by massive stars ($M \ge 8 M_\odot$), while nitrogen is mostly a product of lower mass stars \citep{renzini81,McCall85,EvansDopita85,garnett90}. Consequently, due to the longer time passed on the main sequence by stars with decreasing masses, we would expect some time delay between the enrichment of oxygen and that of nitrogen in galaxies having different ages \citep{matteucci85,garnett90,molla06,richer08}. Assuming the initial mass function (IMF) does not vary between spiral galaxies, such time delay, when properly documented, may thus reveal something fundamental about how these systems formed their stars; for example, allowing to distinguish between constant star formation over the formation by a sequence of stellar bursts \citep{lehnert96,Coz99,tremonti04}. To determine the ``normal'' or standard chemical evolution of Star Forming Galaxies (SFGs) in the nearby universe (z $\leq$ 0.25), the observation and study of a large, homogeneous and statistically representative sample is required. This is where the Sloan Digital Sky Survey (SDSS) project becomes so valuable \citep{york00,hogg01,pier03}. By applying different automatic algorithms to the enormous data bank produced by the SDSS it is now possible to retrieve one of the largest and homogenous sets of spectral line ratios necessary to estimate the basic chemical abundances, and to describe the chemical evolution process of SFGs in a manner consistent with their star formation histories and morphologies. In recent articles \citep{tremonti04,nagao06,izotov06,yin07} data from the SDSS were already used to verify the consistency of the different methods devised in the past to determine the abundance of elements. One of the difficulties encountered in these studies is related to the rarity of the [OIII]$\lambda4363$ line. Theoretically, this line was recognized as crucial in order to determine an accurate temperature for the gas in HII regions. Unfortunately, [OIII]$\lambda4363$ can only be observed in very low metallicity SFGs, which form a minute fraction of the SDSS galaxies. This fact emphasizes the importance of developing different empirical methods like $R_{23}$ or $R_{3}$ to obtain chemical information for a significantly larger sample of galaxies \citep{Pagel79,McCall84,VC92,Thurston96,thuan10}. Recently, some authors \citep[e.g.,][]{nagao06,yin07} advocated that we need to modify some of the empirical methods applied in previous abundance studies using only galaxies where [OIII]$\lambda4363$ was observed. These authors based their claims on the fact that they found apparent significant differences in the abundances determined when they use their new calibrations. However, such point of view is somewhat problematic, as it assumes all galaxies follow the same chemical evolutionary pattern, independently of their mass or morphology, and assumes no evolution with redshift, while these are two assumptions that need to be verified separately. Moreover, other researchers in the field \citep[e.g.,][]{tremonti04,izotov06} that also tested thoroughly the empirical methods using SDSS data, have demonstrated that in general new empirical relations show results that are in good agreement with what was found before. For these reasons, but also for comparison sake with what was done in \citet{Coz99}, we choose for our study to apply the same empirical relations that were used before, but limiting our chemical study to the two most important abundance ratios, O/H and N/O, which were shown by \citet{izotov06} to be less dependent on the method adopted to determined the gas temperature. Another important difficulty encountered in chemical evolution studies of SFGs is related with the contamination of emission lines by absorption features produced by the underlying older stellar populations. In our research we have solved this problem by subtracting a stellar population template from each spectrum, as determined by the spectral synthesis code STARLIGHT \citep{cid05}. This method also has the advantage that through the fitted templates the star formation histories (SFHs) of the SFGs can be deduced, and other metallicity-independent parameters like the stellar velocity dispersion, that combined with the effective radius can yield an estimate of the mass of the bulge. In our study we use these new information in parallel with the morphologies which were determined independently to complete our view about the chemical evolution and formation process of SFGs in the nearby universe. This study is organized in the following manner. In Section~\ref{sec:2} we describe how our sample of SFGs was constructed and how the data for our analysis were obtained. In Section~\ref{sec:3} we present our results for the chemical abundances, and show how they varied with the mass and morphology of the galaxies. In the same section, we also explore the relation between the chemical abundances and the SFHs, and compare our results using STARLIGHT with some relevant models from Starburst 99. In Section~\ref{sec:4} we discuss our observations and propose a new interpretation. Our main conclusion can be find in Section~\ref{sec:5}. Many of our results were verified using statistical tests, which were regrouped in Appendix A.
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\label{sec:5} In \citet{Coz97}, it was reported that UV bright starburst galaxies are deficient in oxygen compared to normal late-type spirals. In \citet{Coz99} it was also reported that the same sample of starburst galaxies show a possible excess in nitrogen abundance. In the present study, using a much larger (122751) and generally selected sample of SFGs (that is, the only restriction was that they do not show the presence of an AGN), we find exactly the same phenomena, suggesting that this is a common trait of star forming galaxies, and not a peculiar characteristics of starburst galaxies. Our analysis of the SFH using STARLIGHT, in good agreement with Starburst 99 models, suggests that the depletion in oxygen and the relative excess in nitrogen are most probably due to the effect of intense starburst winds \citep{Heckman90,lehnert96,Heckman03,tremonti04,Strickland09} happening during a prolonged sequence of bursts of star formation. According to our analysis, all the SFGs form their stars through a sequence of bursts and what produces the chemical differences is a variation in intensity of the bursts. The SFGs experiencing the more intense bursts (with a median intensity above a particular threshold), or those that have experienced more intense bursts in the past, all show the effect of starburst winds: they are deficient in oxygen and relatively rich in nitrogen (N/O). We illustrate our model in Figure~\ref{fig:21}. Contrary to expectation, however, we find no evidence that the efficiency of the starburst wind decreases with the mass of the galaxy. Instead, our data suggests the intensity of the bursts grows with the mass, and with it the depletion of oxygen and the excess of nitrogen. To explain this observation, we propose that the abundance characteristics of these galaxies took shape during their formation process, when their gravitational wells were still forming, and consequently less massive and susceptible to decrease the efficiency of starburst winds. Consistent with this interpretation we have shown that the galaxies with an excess of nitrogen are the more massive, have bigger bulges and have an earlier morphological type than those without an excess. We have also found evidence consistent with downsizing, suggesting that the most massive galaxies formed first at higher redshift. Considering the generality of our analysis we conclude that the formation process of the SFGs is an open process, the galaxies loosing mass and energies to their environment.
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1206.0928
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1206
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1206.5020_arXiv.txt
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We use Galaxy Zoo 2 visual classifications to study the morphological signatures of interaction between similar-mass galaxy pairs in the Sloan Digital Sky Survey. We find that many observable features correlate with projected pair separation; not only obvious indicators of merging, disturbance and tidal tails, but also more regular features, such as spiral arms and bars. These trends are robustly quantified, using a control sample to account for observational biases, producing measurements of the strength and separation scale of various morphological responses to pair interaction. For example, we find that the presence of spiral features is enhanced at scales $\la 70$ \hkpc, probably due to both increased star formation and the formation of tidal tails. On the other hand, the likelihood of identifying a bar decreases significantly in pairs with separations $\la 30$ \hkpc, suggesting that bars are suppressed by close interactions between galaxies of similar mass. We go on to show how morphological indicators of physical interactions provide a way of significantly refining standard estimates for the frequency of close pair interactions, based on velocity offset and projected separation. The presence of loosely wound spiral arms is found to be a particularly reliable signal of an interaction, for projected pair separations up to $\sim 100$ \hkpc. We use this indicator to demonstrate our method, constraining the fraction of low-redshift galaxies in truly interacting pairs, with $M_* > 10^{9.5} M_{\sun}$ and mass ratio $< 4$, to be between 0.4 -- 2.7 per cent.
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When galaxies approach one another, their mutual gravitational attraction can result in substantial disruptions to their morphologies, such as tidal arms, counter arms, bridges and tails. Many studies have shown, both analytically and using numerical simulations, that galaxies of similar mass can provoke dramatic disturbances in the stellar distributions of one another, with the details depending on the orbital parameters of the interaction (e.g., \citealt{toom1972,barn1992,howa1993,gerb1994,barn2011}). Gravitational perturbations can also redistribute the gas content of galaxies, potentially leading to changes in their star-formation properties (e.g., \citealt{nogu1988,barn1996}). These effects are strong functions of pair separation, and hence should be most obvious after the galaxies' first pass, and particularly around times of closest approach. Many interactions are ultimately likely to result in the complete morphological transformation of the galaxies involved (e.g., \citealt{toom1977,hopk2008}). However, more subtle effects are expected both earlier, while the galaxies are on their initial approach, and at times of wide separation between passes (e.g., \citealt{pere2006,lotz2008,stru2011}). The expected strength and prevalence of pair interactions mean they are potentially important for determining the properties and evolution of the galaxy population. It is therefore critical that we test our theoretical expectations of the effects of such interactions by studying representative samples of interacting systems. Furthermore, we may utilise the observed frequency of interacting pairs to discriminate between the details of cosmological galaxy formation models. A relatively straightforward, and physically motivated, definition of `interacting' galaxies is a pair for which the tidal force experienced across one of the galaxies, averaged over its internal dynamical time, $F_\rmn{t}$, is at least some specified fraction of the gravitational force binding the outer regions of that galaxy, $F_\rmn{g}$. By this definition, all interacting pairs should produce significant internal dynamical effects, which would have otherwise been absent, in at least one of the member galaxies. In practice, however, it is difficult to measure the forces involved. We may estimate them by studying the effects of an interaction, but the relative orientations and types of the galaxies in each pair, as well as observational limitations, lead to large variations in the apparent effects for interactions of a given strength, $F_\rmn{t}/F_\rmn{g}$. A more convenient definition, which is roughly equivalent, although only statistically applicable, is that a pair are `interacting' if their gravitational influence upon one another could have observable effects in a favourable orientation and mix of galaxy types. For example, a pair of elliptical galaxies might not display signatures of an interaction in a given observation, but would still count as `interacting' by this definition if the tidal forces they are experiencing would have been sufficient to produce an observable signature in a pair of spirals. The observational details of a particular dataset therefore fix the minimum $F_\rmn{t}/F_\rmn{g}$ probed. This definition removes much of the incompleteness associated with only considering pairs with observational signs of interaction, but of course only a fraction of such 'interacting' galaxies will possess observational signatures. As we shall see later in this paper, it is nevertheless possible to constrain the fraction of galaxies interacting according to this definition. Studies of galaxy pairs typically discuss close pairs, bound pairs, merging pairs, or pairs with observational disturbances, but often mix the usage and definitions of these classes. Interacting pairs are closely related to bound pairs, for which the sum of the gravitational potential energy and kinetic energy of the pair is negative. However, not all interacting pairs are bound, particularly where both are part of a larger system, such as a galaxy cluster. Likewise, not all bound pairs will be experiencing significant tidal interactions. Interacting pairs are also closely related to mergers. Galaxies in bound pairs will typically merge on relatively short time-scales if they are experiencing significant tidal interactions, as the kinetic energy of the pair orbit is transferred to deforming and heating the internal mass distribution of each galaxy (e.g. \citealt{stru1999}). We can identify close pairs of galaxies which are likely to be sufficiently near to one another such that they are interacting (and potentially bound and will eventually merge) using projected distance and line-of-sight velocity \citep{char1991}. However, this approach suffers from significant contamination and incompleteness (with respect to the above definitions of physically meaningful interacting, bound or merging selections), due to a lack of full spatial information and the inverse relationship between relative velocity and separation for loose pairs (i.e. very close pairs can have very large relative peculiar velocities, so appear significantly separated in redshift space). Observational signatures of interactions, for example visual classifications, quantitative morphological measurements or induced star formation, may be used to improve the selection of truly interacting galaxies. However, as often one wishes to study the physical effect of interactions, one must be careful to avoid a circular argument. The early atlases of \citet{voro1959,voro1977} and \citet{arp1966} clearly demonstrated that interactions between galaxies can have profound effects on their morphologies, providing examples of bridges, tails, distorted spiral patterns, and other features. These morphological changes where observed between pairs up to over 100 \hkpc, as is the case for the bridging filament of Arp 295. The restricted three-body simulations of \citet{toom1972} clearly demonstrated that these features are the result of strong tidal forces between the interacting galaxies. That galaxy interactions can also induce star-formation was first suggested by \citet{lars1978}, who found that the scatter in the UBV colours of interacting galaxies from the Arp atlas was significantly larger than normal galaxies from the \emph{Hubble Atlas} \citep{sand1961}. Similar evidence for interaction induced star-formation has also been found over a wide range of the energy spectrum, from near-UV to radio (e.g. \citealt{kenn1984,keel1985,bush1986,kenn1987,bush1987,humm1990}). Studies which use quantitative measures of morphology find signs of interaction at fairly small projected separations. Using the CAS method of \citet{cons2003} to measure galaxy asymmetry ($A$) and concentration ($C$), \citet{hern2005} found that both these quantities increase, relative to isolated galaxies, for galaxy pairs with separations less than the photometric diameter of the primary. \citet{depr2007} reliably identified interacting pairs with projected separations up to $\rp \la 50$ \hkpc using $A>0.35$ and visual confirmation, for a sample of pairs with line-of-sight velocity differences $\Delta V < 500$ \kms.\footnote{Throughout this introduction, separations quoted from other studies have been converted to units of \hkpc as necessary.} Similarly, \citet{elli2010} show that asymmetry increases for $\rp \la 50$ \hkpc for a sample of pairs with $\Delta V < 200$ \kms. Meanwhile, studies which probe the effects of tidal interaction through star formation modulations find changes up to larger projected separations. \citet{niko2004} demonstrate an increase in star formation at at $\rp \la 70$ \hkpc for early and mixed type pairs, and at $\rp \la 430$ \hkpc (their maximum separation probed) for late type pairs. They also find that pairs with $\rp \la 110$ \hkpc show a strong increase in central concentration, suggestive of nuclear starbursts. \citet{li2008} find a star formation increase for close pairs, with star formation rate (SFR) enhanced by a factor of $1.5$ at $\rp \la 140$ \hkpc to a factor of $4$ at $\rp \la 30$ \hkpc. This strong dependence on $\rp$, is contrasted with a weak dependence on luminosity ratio, with the star formation enhancement being strongest in lower luminosity galaxies. \citet{elli2008}, \citet{roba2009} and \citet{patt2011} all find a strong increase in SFR for $\rp<40$ \hkpc, while \citet{patt2011} also see a smaller increase up to at least $\rp < 80$ \hkpc (their maximum separation probed) for $\Delta V < 200$ \kms pairs. There is also evidence that equivalent levels of tidally induced star formation require smaller $\rp$ in denser environments (e.g., \citealt{lamb2003,alon2004}). Some of the variation in the separation scale at which different studies begin to see the effects of tidal interactions is likely due to differences in the mass and luminosity ranges of the samples, as well as the methods used. However, it appears clear that the effects of tidal interactions are found at larger projected separations when identified by induced star formation (up to $80$ to $100$ \hkpc), compared to quantitative measurements of asymmetry (up to $\sim 30$ \hkpc). This is consistent with the results of \citet{lotz2008}, which used simulations to show that quantitative morphological methods for finding merging galaxies, such as $A$, Gini and $M_{20}$, are most sensitive for galaxies undergoing close passages and during the post merger phase. Induced star formation, on the other hand, will be evident between passes, when the galaxies achieve a wide separation before falling back towards one another, or in galaxies which have experienced a close encounter but will not merge. Note that in dry mergers there may be no star-formation signature of the interaction, and morphological features will typically only be observable for short times \citep{bell2006}. As mentioned previously, interacting galaxies can produce distinctive morphological features such as tidal arms, counter arms, bridges and tails, which are best classified visually. These features are extremely reliable for discriminating truly interacting galaxies from interlopers in close pair catalogues. Features such as two loosely-wound tidal arms may not be detectable by quantitative morphology methods because these galaxies may not appear to be sufficiently asymmetric or disturbed, especially between the first and second pass when the galaxies may appear to be widely separated. One of the advantages of using visual morphological classifications over automated methods is the ability to identify very faint and subtle features. Tidal features are known to become rapidly undetectable as a function of time and survey imaging depth (e.g., \citealt{bell2006,scha2010}), however we find that the Galaxy Zoo classifications used in this paper are extremely sensitive to faint features. Furthermore, as we will show, by studying the occurrence of such features in a statistical sense, and making weak assumptions concerning the observability of physical interactions, we are able to make decisive statements concerning the prevalence of interactions. The visual classification of peculiar, disturbed and interacting galaxies has a long history, beginning with \citet{hubb1926}. The catalogues of \citet{voro1959,voro1977}, \citet{sand1961}, and \citet{arp1966} complied together a significant number of galaxies with obvious tidal features. Such work continues to be valuable today, for example \citet{brid2010} uses visual classifications of galactic bridges and tails in the CFHT Legacy Survey to study the evolution of the galaxy interaction fraction (GIF) with redshift. These galaxies were either isolated merger remnants or fairly close interacting pairs, due to their requirement that galaxy pairs be connected by a bridge. \citet{naka2003} and \citet{fuku2007} visually classified a subsample of $\sim 2500$ bright galaxies from the Sloan Digital Sky Survey (SDSS; \citealt{york2000}) imaging of SDSS galaxy objects, finding that $\sim 1.5$ per cent of galaxies in their nearby magnitude-limited sample show morphological indications of interaction. \citet{nair2010} provide an impressive catalogue of detailed visual classifications for over $14\,000$ bright SDSS galaxies, which includes information regarding tidal tails and other indicators of interaction. The Galaxy Zoo project has enabled visual classification to be performed for extremely large samples, allowing the continued use of this valuable technique with modern surveys. \citet{skib2009} obtained the marked correlation function for the Galaxy Zoo 1 merger classification likelihood and found that it increases sharply in their closest $\rp$ bin (of 170 \hkpc width), and found evidence that most of this increase was for pairs with $\rp \la 30$ \hkpc. Taking an alternative approach, \citet{darg2010a} and \citet{darg2010b} imposed thresholds to select Galaxy Zoo 1 galaxies with high merger classification likelihoods and study their frequency and properties. Most of these galaxies are either highly disturbed systems or very close pairs. While these studies have been successful at identifying a subset of interacting pairs, they primarily select galaxies which have relatively small projected separations and so do not typically identify interacting pairs which are at large projected separation between their first and second close passes. In this paper we use visual classifications from Galaxy Zoo 2 to study what morphological changes are taking place in interacting galaxy pairs as a function of physical projected separation ($\rp$) and line of sight velocity difference ($\Delta V$). We then use these results to estimate the frequency of pair galaxy interactions in the local universe. In Section \ref{data} we describe the data set and sample selection, in Section \ref{method} we outline the methods employed, in Section \ref{results} we present our results, and in Section \ref{disc} we summarise our results and discuss their implications. A $\Lambda CDM$ cosmology is assumed throughout, with $\Omega_{\Lambda}$ = 0.7, $\Omega_{m}$ = 0.3, and $h_{70} = H_0 / (70$ \kms Mpc$^{-1})$.
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\label{disc} In this paper we have examined a variety of morphological signatures of interaction between galaxy pairs, and demonstrated how the trends in these observable features, as a function of projected separation, can provide a refined estimate of the frequency of pair interactions in the galaxy population. We consider an `interacting' galaxy to be one which has experienced a significant tidal force, compared to its gravitational binding force, averaged over the previous dynamical time (see Section \ref{intro}). The tidal force deemed `significant' depends upon the properties of a given observational dataset. We began by presenting our sample and the methods we employ, and particularly, in Section \ref{morphprob}, discussing the information provided by Galaxy Zoo 2 and its interpretation in terms of the probability that a given galaxy is observed to possess a particular set of morphological features. We also presented a method to correct for `projection bias', an effect whereby the signal of certain morphological features may depend on the apparent separation of galaxy pairs, even in the absence of any possible physical associations between pair members. In Section \ref{oddclass} we considered questions from GZ2 designed to identify \emph{Odd} features, including answers, such as \emph{Merger}, which were intended to identify signs of interaction. We found that these classifications suffer from a strong projection bias. For example, galaxies with small projected separation, but very large velocity offsets, tend to have a spuriously large \emph{Merger} signal. Furthermore, the way in which the question for these \emph{Odd} features was arranged, allowing only one of the available options to be selected, results in cross-talk between the \emph{Odd} categories, in terms of both true signal and projection bias, which complicates their interpretation. We find that this projection bias is also present for the GZ1 \emph{Merger} classification, and shows a behaviour very similar to its GZ2 equivalent. Previous studies which used the GZ1 \emph{Merger} class to identify merger candidates \citep{darg2010a,darg2010b} will have suffered to some degree from this issue, but the effect is probably relatively small due to their use of vote fraction thresholds and the fact that the low-$\Delta V$ galaxy pairs have a larger mean \emph{Merger} vote fraction at most projected separations relative to the control sample pairs. In future iterations of Galaxy Zoo, and other visual classification efforts, it would be preferable to keep questions regarding different types of features distinct, or allow multiple answers to be selected for a single question when the relevant features are not mutually exclusive. Nevertheless, the \emph{Odd} classifications do provide useful information on the reality of galaxy interactions, particularly once a correction for the projection bias is applied by reference to a control sample of pairs with large velocity offsets. As discussed by \citet{darg2010a} with relation to GZ1, the GZ2 \emph{Merger} class (and also \emph{Odd=Yes}) primarily selects interacting galaxies at small projected separations. It therefore mainly probes close passes and the later stages of mergers. We further find that the \emph{Irregular} and \emph{Disturbed} classes can identify interacting galaxies with very small projected separations, which may be either at an advanced stage of merging or aligned along the line-of-sight. We have searched all the GZ2 classifications for trends with projected separation, and find significant signals with respect to the presence and form of spiral arms. The observability of spiral arms (\emph{Spiral=Yes}), and particularly the dominant \emph{2 Arms} class, is enhanced for close pairs on a scale of $\rp \la 70$ \hkpc. \citet{darg2010a} find that the spiral-to-elliptical ratio for galaxies classified as mergers in GZ1 is approximately twice the global ratio, and in \citet{darg2010b} conclude that this is due to the longer time-scale over which spiral mergers are detectable compared to elliptical mergers. Our results show that this ratio can also be at least partly explained by the enhancement and formation of spiral arms in interacting galaxies. More unusual spiral arm features also present a trend with $\rp$. The occurrence of \emph{One Arm} spirals dramatically increases for small separations (on a scale of $\rp \la 20$ \hkpc), while \emph{Loose Winding Arms} show the strongest increase (operating on an intermediate scale of $\rp \la 30$ \hkpc). There are two principal ways in which spiral-like features can be created through tidal interactions. Tidal perturbations can instigate or amplify instabilities in gas disks, leading to the formation or enhancement of star formation in spiral arms similar to those seen in isolated galaxies (e.g., as seen by \citealt{xu2010}). Tidal forces can also strip stars and gas out of the galaxies, forming tidal tails, counter tails and bridges, which may or may not harbour star-formation (e.g., \citealt{mull2011}). We appear to detect both signatures: an enhancement of `normal' spiral arm features, occurring at large projected separations, with the signatures of stronger tidal interactions becoming increasingly prevalent at smaller separations. The tidal nature of the \emph{Loose Winding Arms} features is confirmed by examination of typical images, such as those in Table \ref{lwaimages}. It is clear that many of the galaxies which have significant probability of \emph{Loose Winding Arms}, especially those with higher stellar mass, are red, early-type galaxies. This indicates that the loose spiral features that are observed in these galaxies, and probably also those same features in star-forming, late-type galaxies, are the result of tidal stripping. In galaxies with sufficient cold gas, there will almost certainly be star formation in these tidal spiral features, and indeed in Table \ref{lwaimages} we see that several of the galaxies possess very blue loose spiral arms. Simulations (e.g., \citealt{toom1972,howa1993,barn2011}) indicate that tidal features, such as those apparently identified by the GZ2 \emph{Loose Winding Arms} class, are indicative of a stage between close passes, primarily between 1st and 2nd pass, when pairs can still attain relatively large separations. The \emph{Loose Winding Arms} features are thus probing the early stage of mergers and pair interactions. The onset and appearance of tidal arms are known to depend on the geometry of the encounter. Numerical studies demonstrate that in-plane, prograde encounters produce the most symmetrical two-sided disturbances, while polar encounters give the most one-sided disturbances, and retrograde encounters are the last to make tidal tails (e.g., \citealt{thom1989,howa1993,barn2009,barn2011}). Retrograde encounters also produce the greatest increase in star formation efficiency \citep{cox2008}. Considering this, the galaxies which are selected as having \emph{2 arms} and \emph{Loose Winding Arms} in GZ2 are likely the result of prograde, in-plane encounters, while galaxies identified as displaying \emph{1 arm} are likely the result of polar or retrograde encounters. Our observed separation scales for these different features are consistent with this interpretation (see Table~\ref{fitparam}). When comparing our results to other studies which look at the onset of tidally induced changes in interacting galaxies, we find that answers to Galaxy Zoo questions regarding spiral arms detect changes at separations similar to studies of tidally induced star formation, as discussed in Section \ref{intro}. The \emph{Loose Winding Arm} class begins to detect interacting galaxies around $\rp \la 120$ \hkpc, which is similar to the separation scale associated with induced star formation \citep{niko2004,li2008,patt2011}. The star formation detected at these large separations is relatively weak, while a strong increase is observed for $\rp \la 40$ \hkpc \citep{li2008,elli2008,roba2009,patt2011}. This corresponds to scale for which we observe an enhancement of the \emph{Merger} class. Quantitative morphological measurements also typically present signals on these scales \citep{hern2005,depr2007,elli2010}. Logically, kinematic disturbance must precede morphological disturbance and so it might seem that star formation, if it is triggered by kinematical perturbations, should be an earlier indication of interaction than morphology (e.g., \citealt{byrd1992}). However, this paper shows that some morphological signatures are as sensitive as enhanced star-formation, and more unambiguously related to interaction. In Section \ref{bars} we found that the likelihood of a bar being observed decreases sharply for pairs with projected separations $\rp \la 20$ \hkpc. Bars are thought to be created through periodic orbital resonance \citep{bour2002} and are known to initiate radial gas inflows, which in the end act to destroy or weaken the bar structure \citep{pfen1990} (for a recent review see \citet{sell2010}, or more comprehensively \citet{sell1993}). Gas inflows, perhaps together with the enhancement of bar features, caused by tidal perturbations in the early stages of major interactions may similarly act to rapidly destroy any pre-existing or transient bars (e.g., \citealt{bere2003,dima2007}). Our results suggest that this is indeed the case, with the appearance of bars being strongly suppressed in close pairs, in agreement with other recent studies by \citet{mend2011} and \citet{lee2012}. Eventually, the violent reorganization of stellar orbits in the later stages of many major interactions (i.e. mergers) must act to erase any orbital resonances which created the bar. In Section \ref{probpairs} we focus on using the presence of \emph{Loose Winding Arms} to identify probable interacting galaxies. These criteria are then used in Section \ref{interlopers} to constrain the frequency of galaxy pair interactions, without requiring an arbitrary cut-off in projected separation or any further corrections for contamination of our close pair sample. We find that the fraction of galaxies with $M_* > 10^{9.5} M_{\sun}$ and $0.01 < z < 0.05$ that are in truly interacting pairs with $\Delta V < 500$ \kms is in the range $0.5 \pm 0.1$--$2.1 \pm 0.6$ per cent. The limits correspond to assuming the maximum and minimum permitted probability, $P_\rmn{int obs}$, that interacting galaxies produce observable \emph{Loose Winding Arms} features, respectively. We expect simple extensions of our technique to lead to significantly tighter confidence intervals in future work. It is difficult to precisely compare our estimate of the interacting galaxy fraction to other studies, due to the range of different methods employed. Although the close pair fraction is mostly constant with luminosity \citep{patt2008}, the limiting mass-ratio and projected separation, varying definitions for selecting pairs, and many other subtleties, make comparisons complicated. Given that the estimate in this paper is derived from a relatively simple demonstration of combining close pair and morphological information, we defer such involved comparisons to future work. Nevertheless, we note that the major interaction fractions quoted by other recent studies: e.g., $1.1 \pm 0.5$, $2.1 \pm 0.1$, $1.6 \pm 0.1$, $1.3 \pm 0.1$ by \citet{bell2006,patt2008,domi2009,xu2012}, respectively, are neatly bracketed by our estimate of 0.4 -- 2.7 per cent. This lends support for our use of GZ2 \emph{Loose Winding Arms} as indicators of pair interactions, and encourages confidence in the method described in Section \ref{interlopers}, and the various assumptions we have made.
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1206.5020
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1206
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1206.6501.txt
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We make a revision of the stability criteria for equatorial circular orbits, obtained from the epicyclic approximation, which is widely used in Newtonian models for axisymmetric galaxies. We find that, for the case of thin disk models, the criterion for vertical stability must be reformulated, due to the discontinuity in the gravitational field. We show that, for a model characterized by a surface mass density $\Sigma$, the necessary and sufficient condition to have vertically stable circular orbits is that $\Sigma>0$. On the other hand, the criterion for radial stability is the same as in thick diks, i.e. that the (radial) epicyclic frequency squared is positive. As an application, we present finite thin disk models for nine galaxies, as superpositions of members of the Morgan \& Morgan family (in Newtonian version), which can be considered as stable configurations in a first approximation. Also, as an additional product of this study, we show that any galactic model with a thin disk component admits a wide variety of integrable disk-crossing orbits, which are determined approximately by a third integral of motion of the form $Z\Sigma^{1/3}$, where $Z$ is the z-amplitude of the motion.
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\label{sec:intro} It is usually accepted that many galaxies in the universe are nearly axisymmetric, with a mass distribution formed by several components: a thin disk, a central bulge and a surrounding halo. In consequence, there is a number of mass models incorporating one or all of these features, depending on the particular case (\cite{freeman, kent1986, kent1987, sofue2008}). For example, there exists a number of galaxies in the Ursa Major cluster which can be modeled, at the scale of the optical radius, with only the thin disk component (\cite{ngc}), suggesting that they obey the so-called \emph{maximum disk hypothesis} (\cite{GD}). If we decide to continue using Newtonian theory beyond the optical radius, presumably we have to introduce a dark halo, but the disk component still provides a significant contribution, taking into account that the main part of the stellar population is located there. For this reason, thin disk models have been an issue of interest in galactic dynamics (see for example \cite{hunter2, morgan, gonzalez-reina, pedraza} as well as \cite{GD} and references therein). Thin disks have also been used to model self-gravitating rings (\cite{letelier1994, letelier2007}), with applications in other branches of astrophysics. These models can be used to describe accretion disks, stellar structures with central black holes (see \cite{sesu,LOP}) or planetary rings. For example, the last issue was partially encompassed in \cite{pn}, where a study of the linear stability of the monopole-ring system was performed by using superpositions of Morgan \& Morgan disks. Similar studies were conducted by the authors in the Newtonian realm of galactic dynamics (\cite{javier, pedraza}). A fundamental step in the formulation of galactic models is the stability analysis. In fact, the stability analysis could suggest sometimes the introduction of new features in a given model in order to, presumably, obtain a more realistic representation. Consider for example the study conducted by \cite{ostriker} where it is shown that a flattened system of self-gravitating particles, initially supported against gravity by rotation, does not maintain its discoidal form in the course of time. They suggested that the introduction of a spherical halo with mass of the order of the disk mass (or greater) improves substantially the stability of the disk, as it was corroborated by simulations. But, on the other hand, consider also the counter argument made by \cite{kalnajs2}, in which it is suggested (i) that the stability problems can be overcome by improving the features of the the inner parts of galactic models (for example, by considering hot centers or small bulges) and (ii) that a halo with a scale length larger than the disk and more massive than it, does not contribute significantly to the stability. This discussion has its roots in the fact that internal kinematics of self-gravitating disks determines the stability of the system: cool disks, which are mainly formed by stars in circular motion (the ones considered by \cite{ostriker}), require a prominent halo in order to avoid instabilities, contrary to the case of hot disks (the ones considered by \cite{kalnajs2}) which can be supported by the random motions of the stars (for a recent review of disk instabilities see \cite{sellwood2011}). A detailed knowledge of the orbital features associated with a given galactic model provides the basis for the stability analysis. Usually this knowledge can be achieved once we have at hand the distribution function (DF) of the model. In general, the obtention of the DF associated to a particular model is not an easy task (except for those models which are defined by an analytical DF) and, in consequence, the corresponding stability analysis is far from trivial. However it is possible to perform a first test of stability without knowing the DF, by using the so-called \textit{epicyclic approximation} (\cite{GD}), i.e. by performing a linear stability analysis of fixed points of the effective potential. Once it is verified that the model is linearly stable, it deserves to perform more conclusive stability tests based on statistical mechanics. The linear stability analysis can be thought as a first test to evaluate how realistic any particular model is. In disk galaxies many stars are on nearly circular orbits, so we have to demand, as a basic requirement, that any galactic model must be characterized by allowing the existence of stable circular motion, especially in regions where the stellar population is maximum, i.e. the equatorial plane. The epicyclic approximation provides a formalism to study motion in the equatorial plane and leads to the establishment of simple stability criteria. However, when one deals with models incorporating a razor-thin disk, this method needs to be reformulated in view of the mathematical features introduced by the thin disk component. Next, we will briefly illustrate this problem. Consider an axisymmetric potential-density pair (APDP) with mass distribution $\rho(R,z)$ and gravitational potential $\Phi(R,z)$, where $R$ and $z$ are the usual cylindrical coordinates. The motion of test particles can be described by the Hamiltonian (\cite{GD}) \begin{equation} H = \frac{P_R^2 + P_z^2}{2} + \Phi_{eff}(R, z), \end{equation} where $P_R=\dot{R}$, $P_z=\dot{z}$ and $\Phi_{eff}$ is the so-called effective potential, defined as \begin{equation} \Phi_{eff}(R, z) = \Phi(R, z) + \frac{l^2}{2R^2}. \end{equation} Here, $l$ represents the $z$-component of the angular momentum, often called azimuthal angular momentum, which is a first integral of motion. The resulting equations of motion can be written as \begin{equation}\label{Req} \ddot{R} = -\frac{\partial \Phi_{eff}}{\partial R}, \end{equation} \begin{equation}\label{zeq} \ddot{z} = - \frac{\partial \Phi_{eff}}{\partial z}. \end{equation} In particular, the equatorial circular orbits (i.e., belonging to the plane $z = 0$) correspond to the minimum of $\Phi_{eff}$, obtained by setting $\partial \Phi_{eff}/\partial R=0$ and $z=0$, which lead to the relation \begin{equation} \left.\frac{\partial \Phi}{\partial R}\right|_{z=0} - \frac{l^2}{R^3}=0. \end{equation} The value of the $R$-coordinate that is solution of the above equation is the radius of the circular orbit with angular momentum $\ell$ (in this case $R$ is called the \emph{guiding-center radius}, which we will denote as $R_{o}$). If this orbit (or any other in the equatorial plane) suffers a small perturbation, one would expect that the resulting motion is not very different from the original one (say, a nearly circular orbit), in order to guarantee the stability of the entire configuration. Strictly speaking, a nearly circular orbit is defined as a trajectory with coordinates $(R,z)$ very close to $(R_{o},0)$, so we can expand $\Phi_{eff}$ near this minimum and neglect cubic and higher terms in the expansion, that is, \begin{equation}\label{expansion} \Phi_{eff}\approx\Phi_{eff}(R_{o},0)+\frac{\kappa^{2}}{2}(R-R_{o})^{2}+\frac{\nu^{2}}{2}z^{2}, \end{equation} where $\kappa$ and $\nu$ are defined as \begin{equation}\label{frecuencias} \kappa^{2}\equiv\left.\frac{\partial^{2}\Phi_{eff}}{\partial R^{2}}\right|_{(R_{o},0)}, \qquad \nu^{2}\equiv\left.\frac{\partial^{2}\Phi_{eff}}{\partial z^{2}}\right|_{(R_{o},0)}. \end{equation} By introducing (\ref{expansion}) in the equations of motion (\ref{Req})-(\ref{zeq}), it can be seen that $R-R_{o}$ and $z$ evolve like the displacements of two harmonic oscillators with frequencies $\kappa$ (epicyclic frequency) and $\nu$ (vertical frequency), respectively, once it is guaranteed that $\kappa^2> 0$ and $\nu^2> 0$. In this case the corresponding circular orbit is said to be \textit{stable} (\cite{GD}). So we say that, in a first approximation, the conditions for stability of the self-gravitating structure located at the equatorial plane (which it is assumed to be composed principally by particles describing nearly circular motions) is that the epicyclic and vertical frequencies squared are both positive. Note that, whenever one deals with razor-thin disks, it is not possible to define a Taylor expansion of the effective potential around the point $(R_o, 0)$. This procedure, which leads to the analysis of vertical and epicyclic frequencies around the circular orbit (\cite{GD}), is only valid for potentials that are smooth or that have at least continuous second derivatives. This means that the analysis of vertical frequencies of thin disk potentials, as did in \cite{pedraza, javier, ngc, pn}, is not a reliable indicator of the vertical stability of the corresponding circular orbits. The analysis must take into account the discontinuity in the partial derivative of the potential with respect to $z$ due to the surface distribution of matter in the equatorial plane. In the following sections we will show that, when considering a razor-thin disk, it can be constructed a vertical stability criterion in terms of the first derivative of the potential, whereas the criterion of the epicyclic frequency remains unchanged (secs. \ref{sec:thin} and \ref{sec:vert-stab}). Also we will present some consequences of this approach in the description of disk-crossing orbits. In particular we will show that many of them can be described by an approximate third integral of motion depending on the vertical amplitude and the surface mass density (sec. \ref{sec:integrability}). Finally we briefly address these issues in the realm of modified theories of gravity (sec. \ref{sec:MOND}).
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We analyzed the stability of equatorial circular orbits in thin disks surrounded by smooth axisymmetric structures. The presence of the thin disk does not allow us to proceed in the same way as in the smooth case because of the delta-like singularity. In particular, the vertical stability criterion for smooth potentials, $\nu^{2}>0$, is not applicable anymore. We developed a consistent vertical stability criterion for circular orbits, which together with the (unchanged) radial stability criterion, ensures Liapunov stability of the corresponding circular orbit. Based on this new formalism, we find that nearly equatorial orbits have a third integral of motion, which is given by eq. (\ref{ZZ'sigma}). This is supported by numerical simulations, which additionally reveal that orbits with great vertical amplitude can be described approximately by this integral (sec. \ref{sec:numerical}). It would be interesting to see if this dependence on the surface density is present in more realistic models, not described by a razor-thin disk, but incorporating a stellar distribution described by a thickened disk. The introduction of the new stability criterion leads to the conclusion that all of the thin disk models presented in \cite{gonzalez-reina,ngc,pedraza} are stable in a first approximation, contrary to the statements shown in such references. This fact urged us to obtain additional models in sec. \ref{sec:FTDModels} in order to show that Hunter's method, taking into account the stability criterion constructed here, is a powerful tool to model the maximum disk of a number of flat galaxies. Having tested the stability of orbits in this class of models, it would be interesting to carry out more conclusive stability analyses based on statistical mechanics considerations (i.e. perturbed solutions of Boltzmann equation, Toomre's criterion, etc). We also have to point out that the stability analyses performed in references \cite{pn} and \cite{javier} need to be corrected. We also briefly addressed the problem of the stability criterion in modified theories of gravity, such as MOND and RGGR, in sec. \ref{sec:MOND} (the same problem, in the realm of general relativity theory, is being studied and the results will be shown in a next paper). We point out that whenever is possible establish a Hamiltonian formulation of the motion, it is also possible to perform an analysis similar to the presented here, in the Newtonian gravity realm. In particular, for the two examples studied here, we find that the stability criteria are also given by the relations $\Sigma>0$ and $\kappa^{2}>0$ and that the assumption of adiabatic invariance leads to eqs. (\ref{ZMOND}) and (\ref{ZRGGR}), introducing deviations from the Newtonian relation (\ref{ZZ'sigma}). This fact can be used as an additional test of these theories, once we have at disposal the required observational data.
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1206.6501
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1206
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1206.4316.txt
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We study the secular, hierarchical three-body problem to first-order in a post-Newtonian expansion of General Relativity. We expand the first-order post-Newtonian Hamiltonian to leading-order in the ratio of the semi-major axis of the two orbits. In addition to the well-known terms that correspond to the GR precession of the inner and outer orbits, we find a new secular post-Newtonian interaction term that can affect the long-term evolution of the triple. We explore the parameter space for highly inclined and eccentric systems, where the Kozai-Lidov mechanism can produce large-amplitude oscillations in the eccentricities. The standard lore, i.e.,~that General Relativity effects suppress eccentricity, is only consistent with the parts of phase space where the General Relativity timescales are several orders of magnitude shorter than the secular Newtonian one. In other parts of phase space, however, post-Newtonian corrections combined with the three body ones, can excite eccentricities. In particular, for systems where the General Relativity timescale is comparable to the secular Newtonian timescales, the three-body interactions give rise to a resonant-like eccentricity excitation. Furthermore, for triples with a comparable-mass inner binary, where the eccentric Kozai-Lidov mechanism is suppressed, post-Newtonian corrections can further increase the eccentricity and lead to orbital flips even when the timescale of the former is much longer than the timescale of the secular Kozai-Lidov quadrupole perturbations.
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\label{intro} Triple stellar systems are believed to be very common in Nature \citep[e.g.,][]{T97,Eggleton+07}. From dynamical stability arguments, these systems must be hierarchical triples, in which the (inner) binary is orbited by a third body on a much wider orbit. Probably, more than 50\% of the bright stars we see are at least (double) binary systems \citep{T97,Eggleton+07}. Given the selection effects against finding faint and distant companions, we can be reasonably confident that the number of triple systems is actually substantially greater than that observed. \citet{T97} showed that $40\%$ of binary stellar systems with period $<10$~days, in which the primary is a dwarf ($0.5 -1.5\,M_{\odot}$), have at least one additional companion. He also found that the fraction of triples and higher multiples among binaries with period ($10-100\,$day) is $\sim10\%$. Moreover, \citet{Pri+06} have surveyed a sample of contact binaries, and noted that 42$\pm5\%$ of 151 of them brighter than 10 mag.~are at least triples. We can then conclude that many close stellar binaries with two compact objects are likely produced through triple evolution. Long-term stability of triple system requires hierarchical configurations: an ``inner'' binary (with masses $m_1$ and $m_2$) in a nearly Keplerian orbit with semi-major axis (SMA) $a_1$, and an ``outer'' binary in which $m_3$ orbits the center of mass of the inner binary, with SMA $a_2\gg a_1$. Another stability condition is that the perturber does not make close approaches to the inner binary orbit. In this stability regime, a highly inclined perturber can produce large-amplitude oscillations in the eccentricity and inclination of the system, the so-called {\emph{Kozai-Lidov mechanism}} \citep{Kozai,Lidov}. The Kozai-Lidov mechanism is an important example of a secular effect (i.e., coherent interaction on timescales long compared to the orbital period) that is common in hierarchical triple systems but absent from two-body dynamics. This process has been proposed as an important element in the evolution of triple stars \citep[e.g.,][]{Har69,Mazeh+79,1998KEM,Dan,PF09,Tho10,Naoz+11sec,Prodan+12,Sharpee+12} and extrasolar planetary systems with an additional distant stellar companion \citep[e.g.,][]{Hol+97,Dan,Wu+07,Takeda,Naoz+12bin}. In addition, the Kozai-Lidov mechanism has been suggested to play an important role in both the growth of black holes (BHs) at the centers of dense stellar clusters and the formation of short-period binary BHs \citep{Wen,MH02}. Furthermore, \citet{Iva+10} showed that the most important formation mechanism for BH X-ray binaries in globular clusters may be triple-induced mass transfer in a BH-white dwarf binary. Given the hierarchical galaxy formation paradigm, and the strong evidence that a high abundance of the local galaxies host supermassive BHs (SMBHs), one expects that major galaxy mergers should inevitably result in the formation of SMBH binaries or multiples \citep{Valtonen96,Loren,Kulkarni+12,Dotti+12}. \citet{Bla+02} showed that the Kozai-Lidov mechanism plays an important role in the evolution of SMBH triples, where high eccentricity induced by the outer perturber can lead to a more efficient merger rate, due to gravitational wave (GW) emission \citep[see also][]{Seto12}. Also, recently \citet{AP12} showed that secular three-body effects play an important role in the evolution of binary compact objects near SMBH. GWs emitted during Kozai-Lidov--induced, highly eccentric orbits of compact binaries might be detectable using LIGO\footnote{\url{http://www.ligo.caltech.edu/}} and VIRGO\footnote{\url{http://www.ego-gw.it/}} (e.g., \citealt{Wen} but see \citealt{Ilya08} and \citealt{OLeary_06}), pulsar timing arrays \citep[e.g.,][]{Finn+10,Amaro-Seoane+11,2012ApJ...752...67K}, and future space-based GW observatories, such as eLISA/NGO \citep{eLISA1,eLISA2}\footnote{\url{http://elisa-ngo.org/}}. In fact, GWs associated with eccentric orbits are stronger and have a very different spectrum relative to their circular counterparts for sources at the same distance and with the same mass and spin. This may allow for the GW detection of eccentric inspirals with higher masses, larger SMAs or farther away from Earth relative to their quasi-circular counterparts \citep{Arun:2007qv,Arun:2007hu,Yunes:2009yz,2009MNRAS.395.2127O,2011arXiv1109.4170K}. Using GW information emitted by the close binary, it might be possible to constrain the parameters of the third body, such as its mass or distance, provided the GW signal-to-noise ratio is sufficiently high \citep{Yunes:2010sm,Galaviz+11}. The Kozai-Lidov mechanism is therefore tremendously important and there is still much to be understood. Recently, \citet{Naoz11,Naoz+11sec} showed that an eccentric outer orbit (and even a circular one with comparable mass inner binary) can behave significantly differently than previously assumed, the so-called ``eccentric Kozai-Lidov mechanism". Specifically, they showed that the inner orbits can flip from prograde to retrograde and back, and can also reach extremely high eccentricities close to unity, and the system behaves chaotically \citep{LN}. Most previous secular three body dynamics studies that incorporated GR effects did so through a pseudoÐpotential, constructed mainly to model accretion disks and 1st post-Newtonian ($\PN$) shifts in the innermost stable circular orbit \citep{NW91,Artemova+96,MH02}. It has been shown that the 1PN precession of the inner body may play an important role in secular evolution \citep[e.g.][]{Ford00,MH02,Bla+02,Mardling07,Dan,Zhang+13}. Here we expand our investigation to include both the eccentric Kozai-Lidov mechanism and the three body $\PN$ effects. We show here (\S \ref{sec:evol} and Appendix \ref{sec:2body}) that although this pseudo--potential does capture some $\PN$ effects, such as the precession rate, the full $\PN$ three-body Hamiltonian introduces other corrections that cannot be modeled with this potential. In this paper, we study the consistent inclusion of $\PN$ terms in the secular dynamical evolution of hierarchical triple systems. We restrict attention to the $\PN$ approximation of the three-body Hamiltonian. While it is well established that the eccentricity and inclination are constant in the $\PN$ two-body problem \citep{DD85}, it is not true for hierarchical triples. In addition to the standard GR precession of the inner and outer orbits, the $\PN$ corrections lead to a new secular interaction between the inner and outer binaries that affects their long-term evolution. We find that the standard lore, i.e.,~that GR effects suppress eccentricity, is only true when the GR timescales are several orders of magnitude shorter than the secular Newtonian ones. When the GR timescales are comparable to the secular Newtonian ones, we show that three-body interactions generally give rise to a resonant-like eccentricity excitation \citep[see also][]{Ford00}. We will be using the term ``resonance" here to describe the rapid excitations of the inner orbit's eccentricity, which occurs when the $\PN$ timescales are comparable to the secular Newtonian timescales. We demonstrate that even for systems with comparable inner binary masses, where the Kozai-Lidov mechanism is suppressed, and even when the GR timescales are much longer than the secular Newtonian ones, $\PN$ corrections continue to excite the eccentricity. This paper is organized as follows. We begin with a definition of the parameters used to describe a hierarchical triple system based on Newtonian and $\PN$ three-body Hamiltonians (\S \ref{Sec:Ham}). We then show that three-body evolution is modified by $\PN$ effects (\S \ref{sec:evol}). We discuss the different time-scales corresponding to the $\PN$ effects, and identify the region in phase space where important deviations might arise due to these terms (\S \ref{sec:times}). We then show that $\PN$ terms can, in many cases, excite the eccentricity of the inner orbit instead of suppressing it (\S \ref{sec:eexcit}). We conclude with a discussion in \S \ref{sec:dis}.
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\label{sec:dis} The Kozai-Lidov mechanism \citep[][see below]{Kozai,Lidov}, has been shown to play an important role for highly inclined hierarchical triples, from planetary systems to stellar size and/or massive compact objects \citep[e.g.,][and references therein]{Naoz+11sec}. For an eccentric outer perturber, the eccentricity of the inner orbit can reach values extremely close to unity, and the inclination can flip from prograde to retrograde \citep{Naoz11,Naoz+11sec}. The quadrupole Kozai-Lidov oscillations between the eccentricity and inclination still persist at octupole order, but they are further modulated on long timescales. We have here studied how the Kozai-Lidov mechanism is affected by $\PN$ corrections to the three-body Hamiltonian, focusing on secular and hierarchical three body systems. We expanded the $\PN$ Hamiltonian in the ratio of SMAs ($\alpha$) to third order beyond leading, i.e.~the leading-order terms in the $\PN$ Hamiltonian perturbation scale here as $a_{1}^{-2}$ and we carried out an expansion up to relative ${\cal{O}}(\alpha^{3})$. We also averaged over the orbital timescale of the inner and outer binary to investigate the long-term secular evolution of the system (\S \ref{Sec:Ham}). We examined the effects of the different $\PN$ terms in this expansion: $\PN$ precession of the inner orbit due to $\bar{\Ham}_{a_1^{-2}}^{\PN}$ (Eq.~\ref{eq:1PNa1}); $\PN$ precession of the outer orbit due to $\bar{\Ham}_{a_2^{-2}}^{\PN}$ (Eq.~\ref{eq:1PNa2}); and a new $\PN$ interaction term between the two orbits, ${\bar\Ham}_{\rm int}^{\PN}$ (Eq.~\ref{eq:1PNint}), which introduces a new inclination and eccentricity dependent modulation (e.g., Fig.~\ref{fig:e1_excitGR}). We compared the different timescales associated with the secular Newtonian and different $\PN$ terms (see Fig.~\ref{fig:timescales}). If the timescales associated with the $\PN$ effects are much shorter than the timescales associated with the eccentric Kozai-Lidov mechanism, i.e.~the secular Newtonian timescales, the growth of the eccentricity in the inner orbit tends to be suppressed. We confirm that the excitation of the eccentricity is indeed suppressed for systems where the Kozai-Lidov timescale is many orders of magnitude longer than the $\PN$ timescales. However, if the timescales of the $\PN$ effects are comparable to the secular Newtonian ones (see Fig.~\ref{fig:timescales}), we found two interesting regimes that present qualitatively different behavior. The first regime is where the $\PN$ timescales are comparable but slightly shorter than the Newtonian Kozai-Lidov timescale. \citet{Ford00}, studying the PSR~B1620$-$26 triple system, noted that the inner eccentricity may be greatly increased around some critical value of the outer SMA, due to the $\bar{\Ham}_{a_1^{-2}}^{\PN}$ term and the octupole term. We extended this calculation by including all averaged $\PN$ terms up to $\mathcal{O}(\alpha^3)$ and the Newtonian octupole term \citep{Naoz+11sec}, as well as exploring a wide region of phase space. We confirmed \citet{Ford00} result and found a resonant-like behavior, where the inner orbital eccentricity is greatly increased compared to the Newtonian case. This behavior exists also when including all averaged $\PN$ terms and for a wide range of mass ratios and orbital parameters. We parameterized the location of the resonant peak in terms of the SMAs by defining a parameter, $\mathcal{R}$ in Eq.~(\ref{eq:R}), as the ratio of the leading-order $\PN$ and secular Newtonian terms. This parameter depends on the ratio of the mass of the outer perturber to the total mass of the inner binary. The presence of the octupole term is important for the resonant $\PN$ eccentricity excitation, which is most apparent in the examples with a small mutual inclination. For systems where either the inner or the outer binary shrinks, for example due to GW radiation-reaction, the triple may pass through this three-body $\PN$ resonance. The amplitude and location of the resonance changes due to $\PN$ terms as a function of $\mathcal{R}$. We found that lower mutual inclinations in the prograde regime cause a wider peak (in terms of $\mathcal{R}$), while a less massive outer body tends to produce wider and higher amplitude peaks. A detailed investigation of the properties of the resonance is beyond the scope of this paper, but could be the subject of future investigations. It is important to note that the outer orbit precession and the interaction term affect the overall time evolution (see Figure \ref{fig:e1_excitGR}). Since these terms are a result of the expansion of the three body $\PN$ Hamiltonian in $\alpha$, it is not surprising that the different terms affect the location of the resonant like behavior (e.g., Figure \ref{fig:emax}). It is interesting however, that they produce a qualitatively different time evolution of the system (e.g., bottom panels of Figure \ref{fig:e1_excitGR}). This suggests that a system evolved under GR effects in the {\bf presence} of a third body has richness to it that should be examined in more detail. This is the subject of future investigation in the framework of direct 3-body integration. The second regime that exhibits qualitatively different behavior from that obtained with a quadrupole Newtonian Kozai-Lidov treatment is when the quadrupolar secular Newtonian timescales are shorter than the $\PN$ ones and when the inner binary has comparable mass components. The eccentric Kozai-Lidov mechanism, neglecting $\PN$ effects, is suppressed when $m_1\to m_2$, since the outer orbit's potential is effectively quadrupolar. As we showed in this paper, $\PN$ effects can break symmetry and excite eccentricity, triggering the eccentric Kozai-Lidov mechanism. As long as $\PN$ precession occurs on a comparable timescale (or lower) than the Newtonian octupole precession, i.e, $t^{\rm N}_{\rm quad}\lesssim t^{\PN}_{a_1^{-2}}\sim t^{\rm N}_{\rm oct}$, the eccentric Kozai-Lidov mechanism will be triggered. Eccentricity excitations are particularly interesting in the context of possible GW detections \citep{Wen,2010PhRvD..81b4007B,2005ApJ...634..921A,2010ApJ...719..851S}. If such excitations were not present, the frequency of the GWs emitted by the inner binary would be typically too low for detection with LIGO \citep[see however][for eccentric binaries which form in the LIGO band]{2009MNRAS.395.2127O,2011arXiv1109.4170K}. However, if eccentricity is secularly excited through a three-body interaction, the frequency of the GWs is also increased during pericenter passage, thus bringing the signals into the detector's sensitivity band. Such large eccentricities would then lead to GW-driven inspiral and the eventual merger of binaries. Whether such eccentric signals can be detected or not will depend on how close such sources are to Earth. But if detections are made with sufficiently high signal-to-noise ratio, then GWs could be used to measure the eccentricity of the inner binary, and thus, distinguish between different source populations.
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1206.4316
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1206
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1206.5791.txt
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In this study, a new method is presented to classify flares derived from the photoelectric photometry of UV Ceti type stars. Using \textit{Independent Samples t-Test}, the method is based on statistical analysis. The data used in the analyses were obtained from four flare stars observed between the years 2004 and 2007. Total number of flares obtained in the observations of \object{AD Leo}, \object{EV Lac}, \object{EQ Peg} and V1054 Oph is 321 in the standard Johnson U band. As a result, flare can be separated into two types as \textit{slow} and \textit{fast} depending on the ratio of \textit{flare decay time} to \textit{flare rise time}. The ratio is below the value 3.5 for all slow flares, while it is above 3.5 for all fast flares. Also, according to the \textit{Independent Samples t-Test}, there are about 157 seconds difference between equivalent durations of slow and fast flares. In addition, there are significant differences between amplitudes and rise times of slow and fast flares.
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Flares and flare processes are very hard worked subtitles of astrophysics. Lots of studies on flares have been carried out since the first flare was detected on the Sun by R.C. Carrington and R. Hodgson in 1 September 1959. Flare processes have not been perfectly understood yet \citep{Ben10}. However, the researches indicate that the incidence of red dwarfs in our galaxy is 65$\%$. Seventy-five per cent of them show flare activity, these stars are known as UV Ceti type stars \citep{Rod86}. In this respect, it will be easier to understand the evolution of red dwarfs if the flare processes are well known. This is because flare activity dramatically affects the evolutions of the red dwarfs. In this respect, an attempt to classify flares by considering shapes of flare light variations observed in the UV Ceti type stars, flares have been tried to classify. It is believed that classification of flares makes the flares and flare processes more intelligible. Flares of the UV Ceti type stars were first detected in 1939 \citep{Van40}. Discovering the flare stars with high flare frequency such as UV Cet, YZ CMi, \object{EV Lac}, \object{AD Leo} and \object{EQ Peg}, the detected flare numbers and the variety of flare light variations increased. The light variation of each flare is almost different from each other. In the first place, it is seen that there are lots of shapes for flare light variations \citep{Mof74, Ger05}. On the other hand, when large numbers of flares are examined, it is seen that there are only two main shapes for flare light variations. One of them is called the fast flare. Fast flares have higher energy and the shapes of their light variations are similar to the shapes of solar hard X-ray flares. The second type flares are called slow flares. Unlike fast flares, slow flares exhibit lower energy. The rise times of slow flares are almost equal to their decay times. The terms of fast and slow flares were used for the first time in 1960's in astrophysics. If the rise time of a flare is below 30 minutes, \citet{Har69} called that flare \textit{a fast flare}. If its rise time is above 30 minutes, they called it \textit{a slow flare}. Like \citet{Har69}, considering the shapes of flare light variations, \citet{Osa68} described two types of flares. However, \citet{Osk69} separated flares into four classes. Like \citet{Osk69}, considering only light variation shapes of the flares, Moffett (1974) directly classified flares such as classical, complex, slow, and flare event. On the other hand, Kunkel had asserted another idea in his PhD thesis in 1967 \citep{Ger05}. According to Kunkel, the observed flare light variations must be some combinations of some slow and fast flares. According to this idea, there are two main flare types. The complex flares mentioned by \citet{Mof74} are actually a combination of some fast and slow flares. And also, both slow flares and flare events mentioned by \citet{Mof74} can be classified as the same type flares. \citet{Gur88} described two flare types to model the flare light curves. \citet{Gur88} indicated that thermal processes are dominated in the processes of slow flares. And these flares are ninety-five per cent of all flares observed in UV Ceti type stars. The non-thermal processes are dominated in the processes of fast flares, which are all the other flares. According to \citet{Gur88}, there is a large energy difference between these two types of flares. In this study, we introduce a new statistical method for classifying flares. Using statistical \textit{Independent Samples t-Test} (hereafter \textit{t-Test}) analysis, this method is based on the distribution of flare equivalent durations versus flare rise times. Considering the studies in \citet{Osa68}, Kunkels PhD thesis \citep{Ger05} and \citet{Gur88}, we assume that there are two flare types such as fast and slow flares types. We classify flares into two types and demonstrate the similarities and differences between these two types of flares. In respect of these analyses, we discuss the results obtained from analyses of 321 flares detected in U band observations of flare stars \object{AD Leo}, \object{EV Lac}, \object{EQ Peg} and \object{V1054 Oph} between 2004 and 2007. The programme stars were selected for this study due to their high flare frequencies \citep{Mof74}. The flare data obtained in this study are useful for such analysis. This is because the data were obtained with systematical observations, using the same method. The flare activity of \object{AD Leo} was discovered by \citet{Gor49} for the first time. The star is a red dwarf and a member of The Castor Moving Group, whose age is about 200 million years \citep{Mon01}. \citet{Cre06} found that the flare frequency of \object{AD Leo} was 0.71 $h^{-1}$. The variation of the flare frequency was investigated by \citet{Ish91}. They mentioned that there is no variation in the flare frequency of \object{AD Leo}. The other star in this study is \object{EV Lac}, which is one of the well known UV Ceti stars. According to the spatial velocities, \object{EV Lac} seems to be a member of 300 million years old Ursa Major Group \citep{Mon01}. It has been known since 1950 that \object{EV Lac} shows flares \citep{Lip52, Van53}. The largest observed flare amplitude is 6.4 mag in U band observations of \object{EV Lac}. \citet{And82} detected 50 flares in U and B band observations of \object{EV Lac}. The author indicated that about 42 flares of 50 flares were found in some groups, which were occurred in every 5-6 days. The seasonal flare frequencies were computed from 1972 to 1981 and these frequencies were compared with the seasonal averages of B band magnitudes. According to this comparison, it was found that the activity cycle is about 5 years for \object{EV Lac} \citep{Mav86}. On the other hand, \citet{Ish91} indicated that there was no flare frequency variation from 1971 to 1988. In another study, \citet{Let97} showed that flare frequencies of \object{EV Lac} increased from 1968 to 1977. \object{EQ Peg} is another active flare star, whose flare activity was discovered by \citet{Roq54}. \object{EQ Peg} is classified as a metal-rich star and it is a member of the young disk population in the galaxy \citep{Vee74, Fle95}. \object{EQ Peg} is a visual binary \citep{Wil54}. Both of the components are a flare star \citep{Pet83}. Angular distance between components is given as a value between 3$\arcsec$.5 and 5$\arcsec$.2 \citep{Hai87, Rob04}. One of the components is 10.4 mag and the other is 12.6 mag in V band \citep{Kuk69}. Observations show that flares of \object{EQ Peg} generally come from the fainter component \citep{Fos95}. \citet{Rod78} proved that 65$\%$ of the flares come from faint component and about 35$\%$ from the brighter component. The fourth star in this study is V1054 Oph, whose flare activity was discovered by \citet{Egg65}. V1054 Oph (= Wolf 630ABab, Gliese 644ABab) is a member of Wolf star group \citep{Joy47, Joy74}. Wolf 630ABab, Wolf 629AB (= Gliese 643AB) and VB8 (= Gliese 644C), are the members of the main triplet system, whose scheme is demonstrated in Fig.1 given in the paper of \citet{Pet84}. Wolf 630 and Wolf 629 are visual binary and they are separated 72$\arcsec$ from each other. Wolf 630AB is a close visual binary in itself. Wolf 629AB is a spectroscopic binary. B component of Wolf 629AB seems to be a spectroscopic binary. VB8 is 220$\arcsec$ far away from the other components. There is an angular distance about 0$\arcsec$.218 between A and B components of Wolf 630 \citep{Joy47, Joy74}. The masses were derived for each components of Wolf 630ABab by \citet{Maz01}. The author showed that the masses are 0.41 $M_{\odot}$ for Wolf 629A, 0.336 $M_{\odot}$ for Wolf 630Ba and 0.304 $M_{\odot}$ for Wolf 630Bb. In addition, \citet{Maz01} demonstrated that the age of the system is about 5 Gyr.
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We observed 321 flares in U band observations of \object{AD Leo}, \object{EV Lac}, \object{EQ Peg} and V1054 Oph. Examining 321 flares, 61 fast and 79 slow flares were identified for analyses. The \textit{t-Test} was used as an analysis method. Flare rise times were accepted as dependent variables, while flare equivalent durations were taken as independent variables. The results obtained from the \textit{t-Test} analyses of the data show that there are distinctive differences between two flare types. These differences are important properties because the models of white light flares observed in photoelectric photometry must support these properties to explain both flare types. The distributions of the equivalent durations were represented by linear fits given by Equations (3) and (4) for these flare types. The slope of linear fit is 1.109 for slow flares, which are low energy flares. And, it is 1.227 for fast flares, which are high energy flares. The values are almost close to each other. It seems that the equivalent durations versus rise times increase in similar ways. In the case of UV Ceti stars, when flare models are considered, it is seen that there are two main energy sources for flares \citep{Gur88, Ben10}. These depend on the thermal and non-thermal processes \citep{Gur88}. Flares with small amplitude are generally the flares with low energy. The thermal processes are commonly dominant for these type flares. On the other hand, the flares, which have sudden rapid increases, are more energetic events. Non-thermal processes are dominant for this type. And thus, there is an energy difference between these two types of flares \citep{Gur88}. When the averages of equivalent durations for two type flares were computed in logarithmic scale, it was found that the average of equivalent durations is 1.348 for slow flares and it is 2.255 for fast flares. The difference of 0.907 between these values in logarithmic scale is equal to 157.603 second difference between the equivalent durations. As it can be seen from Equation (2), this difference between average equivalent durations affects the energies in the same way. Therefore, there is 157.603 times difference between energies of these two type flare. This difference must be the difference mentioned by \citet{Gur88}. The slopes of linear fits are almost close. On the other hand, if the \textit{y-intercept} values of the linear fits are compared, it is seen that there is 0.703 times difference in logarithmic scale, while there is 0.907 times difference between general averages. Considering also Figure 5, it is seen that equivalent durations of fast flares can increase more than slow flare equivalent durations towards the long rise times. Some other effects should be involved in the fast flare process for long rise times. These effects can make fast flares more powerful than they are. When the lengths of rise times for both flare types are compared, it is seen that there is a difference between them. The lengths of rise times can reach to 1400 seconds for slow flares, but are not longer than 400 seconds for fast flares. In addition, when the flare amplitudes are examined for both type flares, an adverse difference is seen according to rise times. While the amplitudes of slow flares reach to 1.0 mag at most, the amplitudes of fast flares can exceed 4.0 mag. Finally, when the ratios of flare decay times to flare rise times are computed for two flares types, the ratios never exceed the value of 3.5 for all slow flares. On the other hand, the ratios are always above the value of 3.5 for fast flares. It means that if decay time of a flare is 3.5 times longer than its rise time at least, the flare is a fast flare. If not, the flare is a slow flare. Therefore, the type of an observed flare can be determined by considering this value of the ratio. In the studies like \citet{Osa68}, \citet{Osk69}, \citet{Har69} and \citet{Mof74}, considering directly the shapes of flare light variations, the flares have been classified into two types as fast and slow flares. For instance, according to \citet{Har69}, if the rise time of a flare is above 30 minutes, the flare is slow flare. If not, it is a fast flare. However it is shown in this study that there are some fast flares, whose rise times are longer than the rise times of some slow flares. This is clearly seen from Table 3. This case indicates that a classification by considering only the rise time may not be right. Nevertheless, \citet{Mof74} separated flares into more than two groups such as classic, complex, spike and flare events. On the other hand, according to our results of \textit{t-Test} analyses, neither only one parameter nor the shape of the light variation was enough to classify a flare. The flare equivalent durations and also one more parameter should be taken into consideration in order to make such a classification. The values 3.5, the ratio of flare decay times to flare rise times, can give an idea about the rate of energy emitting in a flare process. The rise times of flares are some limits for each type. Maximum flare rise time is about 400 seconds for fast flares, while it can reach the values over 1400 seconds for slow flares. However, the decay times can take any duration without any limited values. Consequently, the ratio of flare decay times to flare rise times depends on rise time more than decay times. In the case of rise time, the difference between two type flares must be caused by whether the flare processes are thermal or non-thermal. We computed the duration as a rise time from the phase in which the brightness increases. Increasing of the brightness is caused by increasing the temperature of some region on the surface of the star. The flare rise time is an indicator of heating this region on the surface. Therefore, the ratio of flare decay times to flare rise times, so the values of 3.5, must be a critical value between thermal or non-thermal processes. As it is seen from the models of \citet{Gur88}, the differences between flare durations and flare amplitudes are seen between two flare types derived from observed flares in this study. The difference between amplitudes of slow and fast flares was given by Equation (22) in the paper of \citet{Gur88}. In the case of flare amplitude, the result obtained in this study is in agreement with this equation. Providing that the value 3.5 is a limit ratio for flare types, fast flare rate is 63$\%$ of all 321 flares observed in this study, while slow flare rate is 37$\%$. It means that one of every three flares is a fast flare, while two of them are slow flares. This result diverges from what \citet{Gur88} stated. According to \citet{Gur88}, slow flares with low energies and low amplitudes are 95$\%$ of all flares. The remainder are fast flares. When looking individually over each star, the rate of flare types is changing from star to star. Detected flare number of \object{AD Leo} is 110 as it can be seen from Table 2. Slow flare rate of \object{AD Leo} flares is 78$\%$, while the rate of fast flares is 22$\%$. Detected flare number is 98 in observation of \object{EV Lac} and 40 in observation of V1054 Oph. Slow flare rates of both stars are 75$\%$, while fast flare rates are 25$\%$. \object{EQ Peg} flare number is 78. Slow flare rate of them is 63$\%$, and fast flare rate is 37$\%$. In this study, one of the remarkable points is the correlation coefficients of linear fits. As it is seen in Table 4, the correlation coefficient is 0.732 for linear fit of slow flare type and 0.476 for fast flares. Although the correlation coefficient of the linear fit to the distribution of equivalent durations versus rise times is in an acceptable level for slow flares, it is relatively lower for fast flares. Regression calculations show that the best fits are linear for the distribution of equivalent durations versus rise times in logarithmic scales. The correlation coefficients of other fits are not higher than linear correlation coefficients. Especially, the correlation coefficient is lower for the fast flares due to the distribution of their equivalent durations. As it is seen from Figure 5, the equivalent durations of fast flares can take values in a wide range towards the longer riser times. This must be owing to the same reason of differences between \textit{y-intercept} values and the mean averages of equivalent durations of two flare types. As it is discussed above, while the slopes of the fits are nearly close to each other, there is a considerable difference between \textit{y-intercept} values and mean average of equivalent duration for two flare types. Consequently, all these deviations are seen in fast flares. The magnetic reconnection is dominant in this type of flares. A parameter in magnetic reconnection process causes some fast flares to be more powerful than the expected values. Eventually, some fast flares are more powerful than they are, while some of them are at expected energy levels. On the other hand, this parameter in magnetic reconnection process is not dominated in slow flare processes. And so, distribution of their equivalent durations is not scattered. This must be why the correlation coefficient of the fit is relatively higher for slow flares. In this classification method, the complex flares are an exceptional case. These flares must be composed of some different flares. The complex flare should be separated into component flares before classification. If the fast and slow flares can be modelled, using these models, the complex flares can be decomposed into component flares. In conclusion, some parameters can be computed from flares observed in photoelectric photometry. And, if the behaviours between these parameters can be analysed by suitable methods, the flare types can be determined. In this study, we analysed the distributions of equivalent durations versus flare rise time by the statistical analysis method, \textit{t-Test}. Finally, it is seen that using the ratios of flare decay times to flare rise times, flares can be classified. Thus, flares are classified into two types as fast and slow flares. It is seen that there are considerable differences between these two types of the flares. The differences and the similarities between flare types are important to understand the flare processes. This gives new ideas to model white light flares of UV Ceti stars.
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1206.5791
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1206
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1206.6353_arXiv.txt
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The nearby (d = 7.7 pc) A3V star Fomalhaut is orbited by a resolved dusty debris disk and a controversial candidate extrasolar planet. The commonly cited age for the system (200\,$\pm$\,100 Myr) from \citet{Barrado97} relied on a combination of isochronal age plus youth indicators for the K4V common proper motion system TW PsA. TW PsA is 1$^{\circ}$.96 away from Fomalhaut, and was first proposed as a companion by Luyten (1938), but the physicality of the binarity is worth testing with modern data. I demonstrate that TW PsA is unequivocally a physical stellar companion to Fomalhaut, with true separation 0.280$^{+0.019}_{-0.012}$ pc (57.4$^{+3.9}_{-2.5}$ kAU) and sharing velocities within 0.1\,$\pm$\,0.5 km s$^{-1}$ -- consistent with being a bound system. Hence, TW PsA should be considered ``Fomalhaut B''. Combining modern HR diagram constraints with four sets of evolutionary tracks, and assuming the star was born with protosolar composition, I estimate a new isochronal age for Fomalhaut of 450\,$\pm$\,40 Myr and mass of 1.92\,$\pm$\,0.02 M$_{\odot}$. Various stellar youth diagnostics are re-examined for TW PsA. The star's rotation, X-ray emission, and Li abundances are consistent with approximate ages of 410, 380, and 360 Myr, respectively, yielding a weighted mean age of 400\,$\pm$\,70 Myr. Combining the independent ages, I estimate a mean age for the Fomalhaut-TW PsA binary of 440 $\pm$ 40 Myr. The older age implies that substellar companions of a given mass are approximately one magnitude fainter at IR wavelengths than previously assumed.
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Fomalhaut ($\alpha$ PsA, HD 216956, HIP 113368) is a famous nearby A3V star with a large resolved dusty debris disk \citep[e.g.][]{Holland03} and an imaged candidate extrasolar planet \citep{Kalas08}. The age of Fomalhaut is mainly of interest for predicting the infrared brightnesses of substellar companions \citep{Kenworthy09, Janson12}, calculations of the total mass of the parent bodies generating the dust \citep{Chiang09}, and placing the dusty debris disk in evolutionary context with other stars \citep[e.g.][]{Rieke05}. In general, accurate ages for host stars of substellar objects are useful for constraining not only the masses of the companions, but accurate ages for the youngest stars may help constrain the initial conditions for the substellar objects \citep{Spiegel12}. \citet{Kalas08} recently announced the discovery of a faint optical companion (likely $\lesssim$3 \mjup\,) at separation 12''.7 (96 AU) from Fomalhaut. While the companion has been imaged multiple times at optical wavelengths, \citep[][Kalas et al., in prep]{Kalas08}, it has eluded detection in the infrared \citep{Janson12}. Given the importance of Fomalhaut as a benchmark resolved debris disk system and possible planetary system, a detailed reassessment of its age is long overdue. This paper is split into the following sections: 1) a review of published age estimates for Fomalhaut, 2) estimation of a modern isochronal age for Fomalhaut, 3) demonstration of the physicality of the Fomalhaut-TW PsA binary system, 4) age estimates for TW PsA based on multiple calibrations, and 5) estimation of a consensus age for the Fomalhaut-TW PsA system. These results supersede the age analysis for the Fomalhaut-TW PsA system presented at the 2010 Spirit of Lyot meeting in Paris \citep{Mamajek10}. \begin{deluxetable}{ccccc} \tabletypesize{\scriptsize} \tablecaption{Previous Ages for Fomalhaut\label{tab:ages}} \tablewidth{0pt} \tablehead{ \colhead{Age (Myr)} & \colhead{Ref} & \colhead{Method}} \startdata 200$\pm$100 & 1 & isochrones (Fom), Li, X-ray, rotation (TW)\\ 224$^{+115}_{-119}$ & 2 & isochrones (Fom)\\ 156$^{+188}_{-106}$ & 3 & isochrones (Fom)\\ 290 & 4 & isochrones (Fom)\\ 480 & 5 & isochrones (Fom)\\ 220 & 6 & ``Fomalhaut''\\ 419$\pm$31 & 7 & isochrones (Fom) \enddata \tablecomments{``TW'' = TW PsA, ``Fom'' = Fomalhaut. References: (1) \citet{Barrado97}, \citet{Barrado98}, (2) \citet{Lachaume99}, (3) \citet{Song01}, (4) \citet{diFolco04}, (5) \citet{Rieke05}, (6) \citet{Rhee07}, (7) \citet{Zorec12}.} \end{deluxetable}
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The kinematic data are consistent with Fomalhaut and TW PsA comoving within 0.1\,$\pm$\,0.5 \kms, and separated by only 0.28 pc. Given their coincidence in position, velocity, and statistical agreement in velocities expected for a wide bound binary, and remarkable agreement among independent age estimates ($\sim$10\% agreement), I conclude that Fomalhaut and TW PsA constitute a physical binary. Therefore a cross-comparison of their ages is useful. The new age estimates for Fomalhaut and TW PsA are listed in Table \ref{tab:ages_new}. The new isochronal age for Fomalhaut (450\,$\pm$\,40 Myr) is in good agreement with two recent isochronal estimates: 480 Myr \citep{Rieke05} and 419\,$\pm$31 Myr \citep{Zorec12}. It is clear that more modern evolutionary tracks and constraints on the HR diagram position of Fomalhaut are leading to an age twice as old as the classic age \citep[200 Myr;][]{Barrado97}. Fig. \ref{hrd} (bottom) shows a pleasing overlap between the inferred age probability distribution for Fomalhaut (using the \citet{Bertelli08} tracks) and the gyrochronology and Li ages for TW PsA (the two estimates with the smallest uncertainties). Based on the 4 independent ages in Table \ref{tab:ages_new}, the rounded weighted mean age for the Fomalhaut-TW PsA system is 440\,$\pm$\,40 Myr. This new estimate has relative uncertainties $\sim$5$\times$ smaller than the age quoted by \citet{Barrado97} and \citet{Barrado98} (200\,$\pm$\,100 Myr), and is tied to the contemporary open cluster age scale and modern evolutionary tracks. A factor of 2$\times$ older age for Fomalhaut has consequences for the predicted brightnesses of substellar companions. Using the \citet{Spiegel12} evolutionary tracks, it appears that a factor of 2$\times$ older age indicates that a given brightness limit at 4.5 $\mu$m (or M band) corresponds to thermal emission from a planet roughly 2$\times$ as massive if it were 200 Myr. A 1 M$_{Jup}$ planet of age 200 Myr has absolute magnitude $M_M$ = 20.4, but at 440 Myr is approximately 1.2 magnitudes fainter ($M_M$ = 21.6). Future searches for thermal emission from exoplanets orbiting Fomalhaut should take into account this older age.
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1206.6353
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1206
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1206.6165_arXiv.txt
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{ We forecast combined future constraints from the cosmic microwave background and large-scale structure on the models of primordial non-Gaussianity. We study the {\it generalized} local model of non-Gaussianity, where the parameter $\fnl$ is promoted to a function of scale, and present the principal component analysis applicable to an arbitrary form of $\fnlk$. We emphasize the complementarity between the CMB and LSS by using Planck, DES and BigBOSS surveys as examples, forecast constraints on the power-law $\fnlk$ model, and introduce the figure of merit for measurements of scale-dependent non-Gaussianity.}
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\label{sec:intro} There has recently been a surge in interest to study departures in the distribution of primordial density fluctuations from the random Gaussian case predicted by standard inflationary models. The reason for this renewed interest lies in the fact that any observable departures from Gaussianity would essentially rule out the standard single-field, slow-roll inflationary picture, pointing to a more complicated dynamics during the epoch of inflation (see e.g.\ \cite{Chen_AA,Komatsu_CQG} for reviews). It is therefore important to consider how one could parametrize primordial non-Gaussianity. A much-studied model of primordial non-Gaussianity is the local (or squeezed) model, which characterizes non-Gaussianity through a single parameter $\fnl$ \cite{Salopek,Verde_CMBLSS,Komatsu_Spergel} \eqn{eq:localNG}{ \Phi(x)=\phi_G(x)+\fnl(\phi_G(x)^2-\langle \phi_G (x)^2 \rangle ). } Here, $\Phi$ denotes the primordial curvature perturbations (Bardeen's gauge-invariant potential), $\phi_G(x)$ is a Gaussian random field, and the constant $\fnl$ is the parameter describing deviations from Gaussianity. The local model has been much studied for its simplicity -- it contains the first two terms of the most general local form of non-Gaussianity \cite{Babich_shape}. In a recent paper (\cite{Becker2011}; hereafter BHK11), we introduced a generalization of this model to one where, in Fourier space, $\fnl=\fnl(k)$ is a function of scale \begin{equation} \Phi(k)=\phi_G(k)+\fnl(k)\int \frac{d^3 k'}{(2 \pi)^3}\phi_G(k')\phi_G(k-k'). \label{eq:fnlk_kspace} \end{equation} This is a natural extension of the popular 'local' model\footnote{Other models also display scale dependence of primordial non-Gaussianity; for example, the Dirac-Born-Infeld braneworld theory typically leads to scale-dependent equilateral $\fnl$, which has been constrained in Ref.~\cite{Bean_DBI}.} $\fnl={\rm const}$, and it is physically well-motivated; it can describe inflationary scenarios with multiple light fields, one of which is responsible for the generation of curvature perturbations \cite{chris5,Shandera2010,Barnaby_dilaton}. The bispectrum in this generalized local model is \begin{equation} B_{\phi}(k_1, k_2, k_3) = 2[\fnl(k_1) P_\phi(k_2)P_\phi(k_3) + \perm], \label{eq:fnlk_bispec} \end{equation} where $P_\phi$ is the power spectrum of potential fluctuations. This reduces to the familiar expression $B(k_1, k_2, k_3) = 2 \fnl (P_\phi(k_1)P_\phi(k_2) + \perm)$ when $\fnl$ is a constant. To parametrize this model while retaining its full generality, it is convenient to consider a parametrization in piecewise-constant bins in wavenumber: \begin{equation} \fnl^i\equiv \fnl(k_i) \label{eq:fnl_piecewise} \end{equation} where each $\fnl^i$ is the value of $\fnl(k)$ in the $i$-th wavenumber bin. In BHK11, we used this parametrization to project errors on $\fnl(k)$ from a hypothetical Stage III galaxy survey. As in BHK11, we adopt 20 bins in wavenumber distributed uniformly in $\log(k)$, which is easily sufficient to obtain the best-measured principal components accurately. In this work, we perform an analysis that is similar in spirit to that in BHK11, but extended in several respects. First of all, we develop formalism and work out forecasts for how well the CMB, in particular Planck \cite{Planck_paper}, can measure $\fnlk$. We combine this with the LSS forecasts updated to reflect three specific galaxy surveys. Having done that, we obtain a clearer picture of where we can expect good constraints on non-Gaussianity in $k$-space. Finally, we project our forecasts to the specific, power-law model in wavenumber, and thus clarify at which wavenumber the CMB, LSS and the combined surveys best determine non-Gaussianity. Therefore, this work complements not only BHK11 and studies that forecasted errors for (occasionally slightly different) models of scale-dependent non-Gaussianity \cite{Sefusatti2009,Shandera2010,Giannantonio:2011ya}, but also many previous forecasts for future constraints on {\it constant} $\fnl$ \cite{Sefusatti,Smith_Zaldarriaga,Yadav2007,McDonald,Carbone,Slosar:2008ta,Carbone2,Fergusson2010,Sartoris,Cunha_NG,Fedeli:2010ud,Joudaki,Giannantonio:2011ya,Pillepich:2011zz,Hazra:2012qz}. The structure of this paper is as follows: in Sec.~\ref{sec:LSS}, we briefly review the main result from BHK11 -- the signature of the generalized local model on large-scale structure through halo bias -- and explore the effects of an additional term in the modeling of non-Gaussian bias first pointed out by Desjacques et al.\ \cite{Desjacques2011}. In Sec.~\ref{sec:CMB}, we find the signature of the generalized local model on the CMB bispectrum, with particular emphasis on Planck. Finally, in Sec.~\ref{sec:Results} we combine the results for a set of joint constraints; we also perform a principal-component analysis, and project constraints on a power-law model of $\fnl(k)$. We conclude in Sec.~\ref{sec:concl}. Details of the computational work can be found in the Appendices.
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\label{sec:concl} This paper focused on the ability of upcoming LSS and CMB surveys to probe more general models of primordial non-Gaussianity. We concentrated in particular on the generalized local model where the parameter $\fnl$ is promoted to an arbitrary function of scale $\fnl(k)$. Our starting point were the piecewise constant parameters in $k$, constraints on which are shown in Fig.~\ref{fig:unmarg_errs}, and their principal components which are shown in Fig.~\ref{fig:lss_cmb_pcs} and constrained in Fig.~\ref{fig:pc_errs}. Comparison with theory is easiest, however, by using a simpler parametrization in terms of ``running'' of the spectral index, $\nfnl\equiv d\ln \fnl(k)/d\ln k$. Using the two-parameter description of non-Gaussianity in terms of amplitude $\fnlstar$ and running $\nfnl$, we studied the extent to which a combination of LSS and CMB observations can constrain the running (Table \ref{tab:running}) and $\fnlk$ as a whole (Figures \ref{fig:pretty_plot_30}, \ref{fig:pretty_plot_BB}, and \ref{fig:pretty_plot_0}). For the power-law $\fnl(k)$, we found that both the bispectrum measurement from the CMB Planck survey and power spectrum measurement from an LSS survey can constrain $\fnlk$ tightly in a relatively narrow range of wavenumbers around $k\simeq 0.1\hmpcinv$. The scale best constrained by the CMB is larger (i.e.\ at a smaller $k$) than the scale best constrained by LSS: we get complementary information about $\fnlk$ from the two data sets. The ability of LSS to constrain $\fnlk$ effectively at a wide range of scales depends on the survey parameters and the fiducial model of $\fnlk$ chosen, as is clear from Figures \ref{fig:pretty_plot_30}--\ref{fig:pretty_plot_0} and Table \ref{tab:surveys}. Nonetheless, large galaxy redshift surveys planned for the future may well be competitive with, or even better than, the constraints on the magnitude and running of $\fnlk$ expected from Planck. Beyond the simple power-law model, we find that the combination of CMB and LSS helps pin down the best-constrained few principal components of $\fnl(k)$ better than either probe alone. Figure \ref{fig:pc_errs} shows that the degree of complementarity significantly depends on the details of (and systematics in) the LSS survey. The constraints from the DES and BigBOSS, and other upcoming LSS surveys can turn out to be worse {\it or better} than those illustrated here, depending on how well the systematics can be controlled. While (for example) the photometric redshift errors \cite{Cunha_NG}, calibration errors \cite{calibration}, and assembly bias of galaxies \cite{Reid_assembly} can all introduce parameter biases and degrade constraints, accurate calibration of these effects from simulations and observations, as well as selection of the ``golden'' class of objects with well understood properties whose clustering to use to measure non-Gaussianity, can cancel out these degradations. Moreover, we have not considered information from the LSS {\it bispectrum} which, while somewhat notoriously difficult to theoretically estimate due to non-Gaussian contributions from the gravitational collapse at late times (though see \cite{Chan_bisp,Baldauf_tidal} for recent progress on the matter), is nevertheless a very potent probe of primordial non-Gaussianity (e.g.\ \cite{Sefusatti:2007ih,Jeong_Komatsu_bispec,Sefusatti_halo_bisp,Figueroa}). Overall, a full exploration of the LSS and CMB systematics is a herculean task beyond the scope of this paper; nevertheless, we think we captured a few key systematics with our choice of survey specifications and nuisance parameters. Finally, we introduced the figure of merit for measurements of non-Gaussianity, defined as the inverse area of the constraint region in the plane of non-Gaussian amplitude and running (see Eq.~(\ref{eq:fom})). We are very encouraged by the fact that future constraints of non-Gaussianity will improve current-data figure of merit \cite{nfnl_wmap7} by more than an order of magnitude, and thus shed interesting constraints on the physics of inflation.
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1206.6165
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1206
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1206.1059_arXiv.txt
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We study the thermodynamics of helium at densities relevant for white dwarf physics. We find evidence that, as the temperature is increased, there is first a first order transition between two superconducting phases followed by a second order transition to the normal state. These transitions occur, for realistic densities, at temperatures below the crystallization temperature and the crystalline state is likely to remain as the true ground state of the system. The calculations are performed with a screening but non-dynamical electron background and we comment on the impact of this and other approximations to our result.
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Helium white dwarfs (He WDs) are astrophysical objects which are composed predominantly of helium nuclei and degenerate electrons. At typical WD densities, the nuclei are much closer together than typical atomic sizes but are still widely separated compared to typical nuclear sizes. It has long been known that as WDs cool the nuclei crystallize, locked into position by their mutual Coulomb interactions\cite{abr60,kir60,sal60}. Recently, it was pointed out that in He WDs the temperature at which the helium nuclei form a Bose-Einstein condensate (BEC) might be higher than the crystallization temperature and an intermediate superconducting phase may exist between the plasma and the crystal phases \cite{Gabadadze:2008mx,Gabadadze:2007si,Gabadadze:2009jb,Gabadadze:2009dz,PhysRevLett.66.2915,Ashcroft:kx}. In this phase, it is the ions that are superconducting; the electrons form an ordinary Fermi liquid. The low temperature properties of this phase are dominated by the physics of an unusual ``phonon" excitation \cite{Bedaque:2011hs} and leads to a very small specific heat and enhanced neutrino emission \cite{Bedaque:2012mr}, with possible consequences for the cooling of He WDs \cite{Benvenuto:2011fj}. A similar phase could also exist in a deuterium layer in brown dwarfs \cite{Berezhiani:2010db} and be relevant for inertial confinement \cite{PhysRevLett.78.483,Silva:1997fk,Jeanloz29052007} as well as other kinds of experiments \cite{Badiei200970,Andersson20093067} where high densities are also achieved . It is guaranteed that at large enough densities there will be a range of temperatures where the BEC can exist while the Coulomb crystal cannot. This can be understood by simple scaling arguments: a BEC should form when the thermal de Broglie wavelength $\sqrt{2\pi/MT}$ (here $M$ is the ion mass and $T$ is the temperature) becomes comparable to the interparticle spacing $l$, so that the condensation temperature should scale as $\Tbec\sim 1/Ml^{2}$. A Coulomb crystal should melt when the thermal energy is comparable to the nearest-neighbor interaction, so that $\Tcoulomb\sim Z^{2}\alpha/l$, where $Z$ is the atomic number of the crystallized nuclei and $\alpha=e^{2}/4\pi\approx 1/137$, where $e$ is the size of the electron charge. Since the number density $n\sim l^{-3}$, $\Tbec\sim n^{2/3}/M$ while $\Tcoulomb\sim n^{1/3}$. Thus, at very high density, the crystallization temperature is markedly lower than the condensation temperature, and for intermediate temperatures, the system should be a BEC. The natural question is: are astrophysical densities in this interesting regime? To answer this quantitative question, one needs to know the numerical coefficients that specify these critical temperatures. Because the condensation temperature scales inversely with the ion mass, the density at which $\Tbec=\Tcoulomb$ and beyond which \mbox{$\Tbec>\Tcoulomb$} is smaller for lighter nuclei. Thus, if a nuclear condensate forms in WDs, it should be most easily established in He WDs and not carbon-oxygen WDs. Detailed studies have determined the crystallization temperature to be $\Tcoulomb \sim (Ze)^{2}/180 l$ \cite{1975ApJ...200..306L,PhysRevA.21.2087,1993ApJ...414..695C}, meaning $\Tcoulomb\sim (a_{0}/l) 7000$K, where $a_{0}$ is the Bohr radius. There are various suggestions for the proportionality constant in \Tbec. Simply equating the de Broglie wavelength to the interparticle spacing suggests $\Tbec = 2\pi/Ml^{2}\approx 6.2/Ml^{2}$. A free Bose gas has $\Tbec=T_{c}^{(0)}\equiv 2\pi (4\pi\zeta(3/2)/3)^{-2/3} / Ml^{2}\approx 1.27/Ml^{2}$, where $\zeta$ is the Riemman zeta function. The temperature \Tbec\ is expected to go up when one considers repulsive interactions \cite{Huang:1999zz}. A slightly more detailed estimate (see \reference{Gabadadze:2009jb}) suggests $\Tbec=4\pi^{2}/3Ml^{2}\approx 13.2/Ml^{2}$, which is qualitatively supported by the numerical calculations in \reference{Rosen:2010es}. It is the object of this paper to make a reliable estimate of \Tbec. A BEC composed of nuclei (and not whole atoms) with a background of degenerate electrons is a novel system with rich phenomenology. Because the condensed nuclei are charged, the substance is electrically superconducting. The electrons provide a neutralizing electric charge, and additionally the dynamical response of these electrons implies an unusual gapless quasiparticle\cite{Bedaque:2011hs}. These quasiparticles imbue the substance with a very small specific heat \cite{Bedaque:2011hs}. Moreover, these quasiparticles can annihilate into neutrinos, and the power emitted per unit volume scales like $T^{11}$, so that the phenomenological relevance of this annihilation for He WD cooling depends strongly on the critical temperature of the nuclear condensate, with higher temperatures corresponding to more relevant neutrino emission\cite{Bedaque:2012mr}. These considerations motivate the detailed study of the thermodynamics of such a nuclear condensate. For calculational simplicity, we will work in the regime of stiff electrons, so that they simply screen the Coulomb interaction with a screening mass. In this light, our investigation may be seen as an investigation of nonrelativistic charged spin-0 bosons interacting via a screened Coulomb (that is, Yukawa) interaction. Surprisingly, we find that the system is significantly more complex than expected and that we can merely set an upper bound for the first-order transition temperature: $\Tbec<T_{c}^{(0)}$. We conjecture that this low first-order transition temperature is accompanied by an unforeseen second-order transition at $T_{c}^{(0)}$. This paper is organized as follows: in \secref{sec:model} we discuss this model in detail and calculate its one-loop effective potential. In \secref{sec:phase-diagram} we establish the phase diagram described by this model and investigate its properties analytically and numerically. In \secref{sec:global} we demonstrate that the condensed phase is globally disfavored anywhere the usual uncondensed phase exists, and to resolve this puzzle conjecture that the phase transitions that this system undergoes are more complicated than previously appreciated. Finally, in \secref{sec:conclusion} we make some remarks about the phenomenological relevance of nuclear condensates and discuss priorities for more deeply understanding this model.
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\label{sec:conclusion} Despite the difficulty posed by the fact that the effective potential is complex for positive chemical potential we were able to extract some physical consequences from our one-loop calculation. The main one is the indication of a sequence of a first order transition followed by a second order phase transition. This conclusion was developed by looking at the free energy values computed within the range of validity of the approximations performed and, as such, is quite robust. On the other hand we were not able to compute the critical temperature where the stable (superconducting) and metastable (normal) states trade places. But if the double transition conjecture described above is correct, this temperature is below the free boson critical temperature $T_c^{(0)}$. We also found that the superconducting phase exists, as a metastable state, for temperatures up to about 8 times $T_c^{(0)}$. This conclusion agrees with that in \reference{Rosen:2010es}; this is not entirely a coincidence. Contrary to the present paper, \reference{Rosen:2010es} analyzes the unscreened model ($m_s=0$). But the High $T$ region of the $v=v(T)$ curve is in the plasmon-dominated region where the screening is not important. In addition, the methodological differences between this paper and in \reference{Rosen:2010es} lead to a change in \eqref{eq:dVdmu} that is numerically small and, just like here, the $\mu$ dependence of the dispersion relation is neglected, albeit for different reasons. Although the model we analyzed (bosons interacting through a screened Coulomb (Yukawa) potential) is interesting on its own merits, applications to high density physics require a proper treatment of the effects of a dynamical electron background. Technically the main effect is the inclusion of the contribution of the cuts of $\Pi(p_0,\mathbf{p})$ in the effective potential calculation. Physically they correspond to the fact that the actual force between two bosons presents an oscillating component (Friedel oscillations) \cite{Gabadadze:2009zz,Dolgov:2010gy,Gabadadze:2008pj}. A detailed of the influence of these effects on the thermodynamics of the system will be left for a future publication. We have also not discussed the metastability of the superconducting state in a quantitative fashion. In particular, we have not estimated its lifetime. This is due to the fact that a proper estimate would us require to compute the free energy $F$ as a function $v$ in order to understand the size of the potential barrier separating normal from superconducting phases. Until a deeper understanding of the resummations needed to make sense of the one-loop results is achieved, this calculation is impossible. In fact, all the results discussed here as a well as a confirmation or falsification of our conjectured phase diagram hinge on an understanding of this resummation and that should be viewed as the number one priority for further progress in this topic. Finally, we can use the results of this paper to assess where the idea of nuclear condensates in dense matter stands. Recall that, for the existence of a intermediate temperature regime where the nuclear condensate can exist it is necessary that the crystallization temperature be smaller than the condensation temperature. If one is willing to consider metastable states, our estimate that the superconducting state extends up to $\approx 8\ T_c^{(0)}$ is similar to the hypothesis made in \cite{Gabadadze:2007si}. A such, the estimate that, at densities around $10^5 g/cm^3$ relevant for white dwarf physics a nuclear condensate should exist, stands unaltered. Only an estimate of the decay time of the false superconducting ground state can decide whether the inclusion of the metastable state is appropriate but, considering the extreme slow evolution of white dwarfs and the fact that they start out at high temperatures, suggest the the metastable state is irrelevant. In that case, only at much higher densities ($\rho \agt 2.4\times 10^7\ g/cm^3$) and temperatures ($T\agt 10^6\ K$) can the nuclear condensate exist.
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1206.1059
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1206
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1206.2147_arXiv.txt
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{I review the properties of pulsators located on the upper main sequence in the HR diagram and discuss asteroseismic inferences on the internal structure of stars of spectral type A and B. Special attention is given to the problem of uncertainties in stellar opacities in modelling.} \FullConference{Frank N. Bash Symposium New Horizons In Astronomy,\\ October 9-11, 2011\\ Austin Texas} \begin{document} Pulsating stars on the upper main sequence in the HR diagram commonly exhibit a convective core (which appears if M > 1.2 M$_\odot$) and a radiative envelope with thin convection zones close to the surface. The evolutionary status of pulsators located in this region in the HR diagram can be manifold, however: beside common main sequence pulsators (in which hydrogen core burning takes place) we also have pre-main sequence stars (no efficient nuclear reactions) and post-main sequence stars (hydrogen shell burning). We distinguish different types of pulsators along the main sequence band: among B and late type O stars we have the so-called $\beta$~Cephei pulsators with periods of hours and masses of 8--20~M$_\odot$ and the long-period SPB oscillators (slowly pulsating B stars) with periods of days and masses of 3--12~M$_\odot$. Moving towards lower masses, there are the $\delta$~Scuti pulsators (M = 1.5--2.5~M$_\odot$), which are dwarfs or giants of spectral type A2--F5 located in the extension of the Cepheid instability strip with periods of 0.02--0.3d. Pulsating magnetic stars among A stars are known as roAp pulsators with periods of 5--15 minutes. Among F-type stars there are the $\gamma$~Dor pulsators with masses of 1.4--1.6~M$_\odot$ and periods of 0.3--3~d. Recent observational reviews on these pulsators based on satellite photometry can be found in \cite{2011uytterhoevenI,2011balonaIII,2011balonaguzik} for A and F pulsators and for B stars in \cite{2011balpig,2009degroote}. Prior to the discussion of recent asteroseismic results on the physics in these stars I will review some basics of stellar pulsation.
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1206.2147
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1206
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1206.1234_arXiv.txt
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We present simulations of the compact galaxy group Stephan's Quintet (SQ) including magnetic fields, performed with the N-body/smoothed particle hydrodynamics (SPH) code \textsc{Gadget}. The simulations include radiative cooling, star formation and supernova feedback. Magnetohydrodynamics (MHD) is implemented using the standard smoothed particle magnetohydrodynamics (SPMHD) method. We adapt two different initial models for SQ based on Renaud et al. and Hwang et al., both including four galaxies (NGC 7319, NGC 7320c, NGC 7318a and NGC 7318b). Additionally, the galaxies are embedded in a magnetized, low density intergalactic medium (IGM). The ambient IGM has an initial magnetic field of $10^{-9}$ G and the four progenitor discs have initial magnetic fields of $10^{-9} - 10^{-7}$ G. We investigate the morphology, regions of star formation, temperature, X-ray emission, magnetic field structure and radio emission within the two different SQ models. In general, the enhancement and propagation of the studied gaseous properties (temperature, X-ray emission, magnetic field strength and synchrotron intensity) is more efficient for the SQ model based on Renaud et al., whose galaxies are more massive, whereas the less massive SQ model based on Hwang et al. shows generally similar effects but with smaller efficiency. We show that the large shock found in observations of SQ is most likely the result of a collision of the galaxy NGC 7318b with the IGM. This large group-wide shock is clearly visible in the X-ray emission and synchrotron intensity within the simulations of both SQ models. The order of magnitude of the observed synchrotron emission within the shock front is slightly better reproduced by the SQ model based on Renaud et al., whereas the distribution and structure of the synchrotron emission is better reproduced by the SQ model based on Hwang et al..
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The comparison of our numerical simulations with observations of SQ yields some agreement with several observed properties of the gaseous and stellar components of SQ. However, the quality of the match with observations is different for our two SQ models. In Table 7 we present a brief listing of the achievement of our simulations in comparison with the results of the previous numerical studies of SQ by \citet{ReAp10} and \citet{HwSt12}. In general, both presented models of SQ are capable of reproducing its large-scale structure, but the spacial distribution of matter is different due to the different sizes of the participating galaxies. Furthermore, the enhancement of gaseous properties is naturally more efficient for SQ model A compared to SQ model B, which is mainly due to the larger total masses of the involved progenitor galaxies. Regions of high synchrotron intensity might correspond to regions of high density, however, within periods of intensive shock ejection, the magnetic field is amplified by turbulence and shocks. During the interactions, the magnetic field and correspondingly the total and polarized synchrotron intensities are enhanced behind the shocks and transported into the IGM. This behaviour is well displayed in our SQ simulations. The observed prominent shock within SQ develops within both of our SQ models. The shock region is found in both the synthetic X-ray and the radio maps of our simulations. The distribution of the total synchrotron intensity compares very favorably with observations \citep[cf. Figs. 2 and 12, see also][]{XuLu03}. The highest synchrotron values are reached within the region of the large shock in both SQ models. However, model A reproduces the strength of the synchrotron emission in the observed order of magnitude, whereas the extension of the shock region is slightly better reproduced by model B, which clearly underestimates the strength of the synchrotron emission. The final configuration in SQ model A is reached after the high-speed intruder NGC 7318b hits the system and finally collides with NGC 7318a. Nevertheless, we claim that the large shock within this system results from the collision of NGC 7318b with the IGM. Especially the direction of the magnetic field vectors (which are directing towards NGC 7318b) supports this origin of the shock, which is also consistent with general acceptance. However, it cannot be ruled out that the large shock in SQ model A results at least partially from a collision of NGC 7318b with NGC 7318a. On the contrary, the large shock in SQ model B is certainly a result of an interaction of NGC 7318b with the IGM as the interaction between NGC 7318a and NGC 7318b happened approximately 220 Myr before the present-day configuration is reached. Overall, we can recognize a general trend in our simulations: the enhancement and propagation of regions with higher values of the studied gaseous properties (i.e. temperature, X-ray emission, magnetic field strength and synchrotron intensity) resulting from interactions and associated shocks and outflows is much more efficient for the more massive galaxies of SQ model A. SQ model B shows qualitatively similar effects of enhancement and propagation, but to a smaller extent. This is not surprising, as higher masses lead to higher kinetic energies of the interacting galaxies, which results in higher equipartition values i.e. the thermal energy (temperature) or the magnetic energy (magnetic field strength). These general findings and the good agreement of the synchrotron intensity in model A with observations may be interpreted as an indication for larger progenitor galaxies within SQ. Therefore, the promising extension of the shock region in both X-ray and synchrotron emission motivating the underlying formation scenario of SQ model B in combination with the total masses of the SQ model A would provide a very good starting basis for further studies.
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We have presented simulations of Stephan's Quintet including magnetic fields, radiative cooling, star formation and supernova feedback. We have investigated different properties of the gaseous component for two different galaxy models based on \citet{ReAp10}, SQ model A, and based on \citet{HwSt12}, SQ model B, respectively. We have set the focus on the general morphology, on the distribution of star-forming regions and star formation rates, on the temperature and the corresponding X-ray emission and finally on magnetic fields and the resulting total and polarized radio emission. A brief listing of the achievement of our simulations in comparison with the previous studies by \citet{ReAp10} and \citet{HwSt12} is shown in Table 7. The main results of our simulations can be summarized as follows: \begin{itemize} \item The present-day configuration of SQ model A develops within 320 Myr. The morphology of the system agrees qualitatively well with observations, only the position of the galaxy pair NGC 7318a/b, its small-scale details and the inner and outer tails cannot be reproduced correctly. The outer tail is generated in this model but already too diffuse to be visible at the present-day configuration as already noted for the original model of \citet{ReAp10}. \item The present-day configuration of SQ model B develops within 860 Myr. Again, the morphology of the system agrees qualitatively well with observations, however, the position of NGC 7318a is slightly too southern. Also, the small-scale features such as the arms of NGC 7318b or the smaller-scale structure of NGC 7319 cannot be reproduced correctly and the outer tail is shorter compared to observations. \item Within SQ model A, the total masses of the galaxies are approximately Milky Way-like. In contrast, the galactic masses of SQ model B are roughly a factor of 10 smaller compared to SQ model A. As lower galactic masses imply lower equipartition energies, the enhancement of the gaseous properties is commonly lower for SQ model B. \item The regions of active star formation within SQ model A are found mainly in the discs of the galaxies, and also within the inner tail and between NGC 7319 and the pair NGC 7318a/b. The latter partly coincides with the region of the large shock. Within SQ model B, the regions of active star formation are found within the inner discs of NGC 7319 and the galaxies NGC 7318a/b, but there is no region of active star formation between these galaxies. \item The global SFR strongly depends on the initial masses of the galaxies and is significantly higher for the more massive SQ model A. \item In both models, the temperature of the gas within the galaxies is cooler compared to the IGM, which gets heated by shocks and outflows caused by the interactions. The mean temperature in SQ model B is significantly lower compared to SQ model A. \item The X-ray emission shows the highest luminosities in the region of the large shock between NGC 7319 and the pair NGC 7318a/b within both models, in good agreement with observations \citep{PiTr97,SuRo01}. The X-ray luminosity in the shock region within SQ model B is about one order of magnitude smaller compared to SQ model A. \item We find high values of the magnetic field strength in the region of the large shock and also within outflow regions in both SQ models. The values of the magnetic field strength within SQ model B are approximately a factor of 3 smaller compared to SQ model A. \item The temporal evolution of the mean total magnetic field reveals an amplification of the magnetic field strengths along with the different interactions. Thereby, the increase in the magnetic field strength is more efficient for SQ model A. \item The synthetic radio maps of both models show a high total and polarized synchrotron intensity within the large shock, within NGC 7319 and around and within NGC 7318a. This finding agrees well with observations \citep[cf.][]{XuLu03}. \end{itemize} The large shock revealed by observations of SQ is most likely the result of a collision of NGC 7318b with the IGM. The observed ridge of radio emission can therefore be ascribed to shock activity. The shock front in our simulations is clearly visible in the X-ray and synchrotron emission within both SQ models. We emphasize the importance of shocks for the magnetic field amplification and the enhancement of the synchrotron emission. Whenever a high amount of synchrotron emission is detected in regions between interacting galaxies, it may be ascribed to shock activity. For future studies, a further development of the existing SQ models would be essential to draw more detailed conclusions on the extension and strength of the synchrotron emission within SQ. As the SQ model B results in a lower enhancement of the gaseous properties mainly because of the smaller masses, but displays the regions of enhanced X-ray and synchrotron emission quite well, it would be worthwile to use a different scaling of the total masses of SQ model B comparable to the total masses of the SQ model A. This would lead to a better comparability of the strengths of the gaseous properties of the present-day configuration of the two different models of SQ. Another particular focus in further studies should thereby be placed on the position and extension of the galaxy pair NGC 7318a/b, which we found to significantly affect the extension and structure of the large shock in SQ. As in our simulations the used particle masses are of the order of the mass of the largest molecular clouds, small-scale turbulence within the large shock region as recently observed by \citet{Gu12} cannot be modeled in our work. Therefore, further numerical simulations focusing on smaller scales would lead to a deeper understanding of the involved processes of shock activity, especially shocks wrapped around clouds and cloud like structures. Furthermore, observations of the radio emission ridge at different frequencies would be of particular interest in order to gain new insights into the shock region. This knowledge could then be used as a basis for further improvements of numerical SQ models.
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1206.1234
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1206
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1206.6078_arXiv.txt
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We study the dynamics of continuum driven winds from rotating stars, and develop an approximate analytical model. We then discuss the evolution of stellar angular momentum, and show that just above the Eddington limit, the winds are sufficiently concentrated towards the poles to spin up the star. A twin-lobe structure of the ejected nebula is seen to be a generic consequence of critical rotation. We find that if the pressure in such stars is sufficiently dominated by radiation, an equatorial ejection of mass will occur during eruptions. These results are then applied to $\eta$-Carinae. We show that if it began its life with a high enough angular momentum, the present day wind could have driven the star towards critical rotation, if it is the dominant mode of mass loss. We find that the shape and size of the Homunculus nebula, as given by our model, agree with recent observations. Moreover, the contraction expected due to the sudden increase in luminosity at the onset of the Great Eruption explains the equatorial ``skirt" as well.
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\label{Part1} We begin in \S\ref{sec:WindsTheory} by developing an approximate analytical model % of continuum driven winds from rotating stars. In \S\ref{sec:Spinup}, we show that at low mass loss rates, such winds are capable of spinning up the star, even though angular momentum is lost to the wind. In \S\ref{sec:SkirtTheory}, we study the consequences of eruptions to~the~evolution~of~the~stellar angular momentum. \subsection{Continuum-driven Winds from Rotating Stars} \label{sec:WindsTheory} A complete analysis of the dynamics of LBV eruptions is quite complicated. In general, one should solve the Navier-Stokes equations for the wind, coupled to the radiative transfer. Furthermore, as we argue that the star is rotating near breakup, its oblique shape must be taken into account. Such an analysis is amenable only by a numerical treatment. Since we wish to understand the underlying physics, this path will not be taken here. In order to encapsulate the relevant phenomena within a simple analytical model, we make the following simplifying assumptions. \begin{enumerate}[I.] \item The star is spherical. \item There is no latitudinal motion ($v_{\theta}\ll v_r)$. \end{enumerate} Under these approximations, the model is integrable. We first note that the typical wind velocity is of the order of magnitude of the escape velocity, which is much larger than the speed of sound. This implies that the kinetic energy density is significantly higher than the internal energy of the gas; mechanical pressure and viscosity will therefore be neglected. We focus on two fields: the velocity of the wind $\bf v$ and the radiative flux $\bf F$, and begin with conservation of momentum: \begin{equation} \rho\left(\frac{\partial}{\partial t}+{\bf v}\cdot {\nabla}\right){\bf v}=\sum{\bf f}. \label{eq:MomentumConservation} \end{equation} Here $\sum {\bf f} = \rho({\bf g}_\mathrm{grav} + {\bf g}_\mathrm{rad})$ is the sum of the different forces per unit mass, with \begin{equation} {\bf g}_\mathrm{grav}=-g\left(\frac{R}{r}\right)^{2}\hat{\bf r},\qquad {\bf g}_\mathrm{rad} =\frac{\kappa}{c}{\bf F}, \end{equation} where $g\equiv GM/R^{2}$ and $ \kappa$ is the opacity. At a steady state, the radial component of eq.~(\ref{eq:MomentumConservation}) is\footnote{Under these approximations the effective gravity is radial.} \begin{equation} v_{r}\frac{\partial}{\partial r}v_{r}-g_{\textrm{eff}}=\frac{\kappa}{c}F_{r}, \end{equation} where \begin{equation} g_{\textrm{eff}}\equiv\frac{v_{\phi}^{2}}{r}-\frac{GM}{r^{2}}. \end{equation} Conservation of angular momentum implies that \begin{equation} v_{\phi}=\left(\frac{R^{2}}{r}\right)\omega \sin\theta, \label{eq:AngularMomentumConservation} \end{equation} and hence \begin{equation} g_{\textrm{eff}}=-g\left[\left(\frac{R}{r}\right)^{2}-\left(\frac{R}{r}\right)^{3}{\Omega}^{2}\sin^{2}\theta\right], \end{equation} where % \begin{equation} \Omega\equiv\omega\sqrt{{R^{3}}/{GM}} \end{equation} is the spin in units of the classical breakup limit. The radiation field loses energy as it accelerates the wind, lifting it over the gravitational pull. Conservation of energy then implies \begin{equation} {\nabla}\cdot{\bf F}=-\frac{\kappa}{c}\rho{\bf F}\cdot{\bf v}, \end{equation} which reduces to \begin{equation} \frac{1}{r^{2}}\frac{\partial}{\partial r}\left(r^{2} F_{r}\right)=-\frac{\kappa}{c}\rho F_{r} v_{r}, \label{eq:flux} \end{equation} assuming axial symmetry and neglecting latitudinal motion. Conservation of mass, \begin{equation} \frac{\partial \rho }{\partial t} + \nabla \cdot \left( \rho {\bf v} \right)=0, \label{continuity} \end{equation} closes the set of equations. At a steady state with the assumed symmetries, it reduces to \begin{equation} \frac{\partial}{\partial r}\left(r^{2} \rho v_{r}\right)=0. \label{eq:Ftheta} \end{equation} For convenience, we define \begin{equation} f\left(\theta \right)\equiv \frac{4 \pi r^{2} \rho v_{r}}{-\dot M}, \label{fDef} \end{equation} where $\dot{M}$ is the mass loss rate and is therefore negative.\footnote{This choice was made to keep $f \left(\theta \right)=1$ in the case of no rotation.} Plugging this into eq.~(\ref{eq:flux}) and integrating gives \begin{equation} F_{r}= \frac{L}{4\pi r^{2}} e^{{-m f (\theta)}\left(1- {R}/{r}\right)} \chi(\theta), \label{eq:Fr} \end{equation} where the luminosity $L$ was introduced by dimensional analysis in order for the constant of radial integration $\chi(\theta)$ to be dimensionless, and % \begin{equation} m\equiv\frac{-{{\dot{M} \kappa}}}{{{4 \pi c R}}} \end{equation} is the ``photon tiring number".\footnote{This definition is slightly different from the original one defined by \cite{OwockiGayley1997}. Here the mass loss is compared with the Eddington luminosity, rather than the actual one.} Using von Ziepel's theorem \citep{vonZeipel1924} with the simplification of having just a radial flux, we can determine the radiative and mass fluxes at the base of the wind: \begin{equation} F_{r}\Big{|}_{r=R}\propto g_{\textrm{eff}}\Big{|}_{r=R}. \end{equation} The requirement $\int {\bf F}\cdot {\bf dS} = L $ gives \begin{equation} \chi\left(\theta \right)=\frac{1-{\Omega}^{2}\sin^{2}\theta}{1-\frac{2}{3}{\Omega}^{2}}. \end{equation} Unlike line driven winds, the mass flux expected in a SED continuum driven wind is inherently local. This is because the critical surface of the wind, where the gravitational and radiative forces balance, depends only on the size of the inhomogeneous structure formed by radiative hydrodynamic instabilities. Since the hydrostatic scale height, over which these instabilities operate, is much smaller than the stellar radius, mass loss is determined by local conditions. As the local atmospheric structure cannot distinguish between gravity and the centrifugal force, the mass loss will depend on the local {\em effective} gravity. Generalizing the local mass flux derived by \cite{ShavivNovae}, we have \begin{equation} {\bf \Phi} = \frac{\cal W} {c v_s}\left({\bf F} - {\bf F_{\textrm{crit}}}\right), \end{equation} where \begin{equation} {\bf F_{\textrm{crit}}}={\bf F_{\textrm{edd}}}\left(1-\Omega^{2}\sin^{2}\theta\right). \end{equation} This is analogous to the local wind flux in SED accretion disks \citep{DotanDisks}. Here $F_{\textrm{edd}} = g c / \kappa$ is the ``Eddington flux", for which the radiative flux $\bf F$ balances gravity in the case of no rotation; $\bf F_{\textrm{crit}}$ is the equivalent flux when rotation is added. Last, ${\cal W}$ is the wind constant, which encapsulates mostly geometrical features described by \cite{ShavivNovae}, but may have a weak dependence on $\Gamma \equiv F/F_{\textrm{edd}}$. \cite{Owocki2004} have shown that this dependence is stronger if the inhomogeneous structure comprising the unstable atmosphere is power law dominated by smaller scales. However, as we shall see below, this merely scales the mass loss by a constant factor since $\Gamma$ is latitudinally independent. Note that a weak latitudinal dependence of the mass loss should arise from the latitudinal dependence of $v_s$ at the base of the wind. However, $v_s$ is proportional to $\sqrt{T}$ which itself is a very weak function of the optical depth. Neglecting the aforementioned weak dependences, the flux depends on $\Gamma$ as \begin{equation} {\bf \Phi}=\rho{\bf v} \propto {\bf F-F_{\textrm{crit}}}\propto \left( \Gamma -1+\frac{2}{3}\Omega^{2} \right ) \chi\left(\theta \right). \label{eq:MassFlux} \end{equation} Under the above approximations, a comparison with eq.~(\ref{fDef}) reveals that \begin{equation} f\left(\theta \right)= \chi\left(\theta\right). \end{equation} A surface integration of ${\bf \Phi}$, \begin{equation} -\dot{M}=\int {\bf \Phi} \cdot {\bf dS}, \label{eq:MdotPhi} \end{equation} gives \begin{equation} m=\frac{1}{2}\mathcal{W}\frac{v_{\textrm{esc}}^{2}}{c\, v_{s}}\left(\Gamma-1+\frac{2}{3}\Omega^{2}\right)\equiv \tilde{\mathcal{W}} \left(\Gamma-1+\frac{2}{3}\Omega^{2}\right), \end{equation} where $\tilde{\mathcal{W}}$ is the scaled wind constant and $v^{2}_\textrm{esc}\equiv2GM/R$. An explicit form of $F_{r}$ (eq.~\ref{eq:Fr}) allows a direct integration of $v_{r}$ (eq.~\ref{eq:flux}): \begin{equation} \frac{1}{2}v_{r}^{2}\Big|_{R}^{r}=\int_{R}^{r}dr\left(g_{\textrm{eff}}+\frac{\kappa}{c}F_{r}\right). \end{equation} Neglecting the velocity at the base of the wind gives \begin{eqnarray} \left(\frac{v_{r}}{v_\textrm{esc}}\right)^{2} &=& \left(1-\frac{R}{r}\right)\left(\Gamma \chi\left(\theta \right) {\cal I}-1\right) \nonumber \\ &+& \frac{1}{2}{\Omega}^{2}\sin^{2}\theta\left[1-\left(\frac{R}{r}\right)^{2}\right] \label{Vr} \end{eqnarray} where \begin{eqnarray} {\cal I} &\equiv& \frac{1-e^{-m \chi\left(\theta \right) \left(1-{R}/{r}\right)}}{m\chi\left(\theta \right)\left(1-{R}/{r}\right)} \\ &=& 1 - \frac{m\chi(\theta)}{2} \left(1-\frac{R}{r}\right)+\mathcal{O}\left(m^{2}\right).\nonumber \end{eqnarray} Note that the wind approaches its terminal velocity after traversing just a few stellar radii: \begin{equation} \left(\frac{v_{r}}{v_{\infty}}\right)^{2} > 1-\frac{R}{r}. \label{eq:TerminalVelocity} \end{equation} This justifies our second assumption, that $v_{\theta}\ll v_r$. In principal, one could argue that even if the star were spherical, radiative diffusion would have given rise to latitudinal radiation and velocity components. However, since most of the acceleration takes place close to the surface, these components are bound to be small. \subsection{Self Spin-up by Continuum-driven Winds} \label{sec:Spinup} As the star blows wind, it loses angular momentum. However, specific angular momentum and thus the dimensionless spin $\Omega$ may increase, as we now show. We start with conservation of angular momentum. The angular momentum reduction of the star is the angular momentum taken by the wind: \begin{equation} \label{eq:lconserve} \dot{\ell}_s =-\dot{\ell}_w, \end{equation} where \begin{equation} \ell_s = \omega M R^2 \alpha_g^2, \end{equation} and \begin{eqnarray} \dot{\ell}_w&=&\omega \int R^2\sin^2\theta \,{\bf \Phi}\cdot {\bf dS} =-\omega \dot M R^2 \frac{1}{2}\int\chi\sin^3\theta d\theta \nonumber \\ &\equiv &-\omega \dot M R^2 \alpha^2_w. \end{eqnarray} Here $\alpha_g$ is the stellar radius of gyration. $\alpha_w$ is an effective radius of gyration of the wind, defined in the equation above. Dividing eq.~(\ref{eq:lconserve}) by $\ell_s$, we find \begin{equation} \label{eq:omegaevol} \frac{3}{2} \frac{\dot M}{M}+\frac{1}{2} \frac{\dot R}{R}+\frac{\dot \Omega}{\Omega}=\frac{\alpha_w^2}{\alpha_g^2} \frac{\dot{M}}{M}. \end{equation} In the upper part of the main sequence, where radiation pressure dominates, one roughly has that $R \propto \sqrt{M}$, such that $2 {{\dot R}/{R}} \approx {\dot M}/{M}$. Eq.~(\ref{eq:omegaevol}) can now be integrated to give \begin{equation} \log\frac{\Omega_f}{\Omega_i}=\left( \frac{\alpha_w^2}{\alpha_g^2}-\frac{7}{4} \right) \log{\frac{M_f}{M_i}}. \label{eq:loglog} \end{equation} This assumes that over the integration interval, $\alpha_w$ remains constant.\footnote{This assumption is applicable to both the present day wind with its low mass loss rate, and high load winds at a steady state. Eruptions are treated differently in the following section.} The condition for self spin-up is therefore \begin{equation} \alpha_w^2 < \frac{7}{4} \alpha_g^2. \label{eq:selfspinup} \end{equation} For a high load wind $\alpha^2_w\approx 2/5$, and the star spins down.\footnote{Massive stars typically have $\alpha_g^2\sim0.1$ \citep{Motz1952}. } However, at low mass loss rates, corresponding to $m\ll 1$, the flux is insufficient to push the wind over the effective potential at equatorial latitudes. The wind reaches infinity only for angles between the pole and \begin{equation} \theta_{\textrm{max}}\approx\sqrt{m\frac{2/{\tilde{\mathcal{W}}}-1}{\Omega^{2}\left(1-\frac{2}{3}\Omega^{2}\right)}} +\mathcal{O}(m^{3/2}). \label{eq:thetamax} \end{equation} We expect the rest of the wind to stagnate and fall back, thus taking no net angular momentum. Therefore, \begin{eqnarray} \alpha^2_w&=&\int_0^{\thetamax}\chi\sin^3\theta d\theta \\ &=&\frac{1}{4}\left( 1-\frac{2}{3}\Omega^2\right)^{-1}\thetamax^4+\mathcal{O}(\thetamax^6). \nonumber \label{eq:alphaw} \end{eqnarray} One should note that the photon tiring number should be computed using the mass loss at the surface of the star, not at infinity. This quantity relates to the observed ejecta by truncating the surface integration of $\bf \Phi$ (eq. \ref{eq:MdotPhi}) at $\theta_{\textrm{max}}$: \begin{eqnarray} m_\infty&=&m\int _{0} ^{\theta_{\textrm{max}}} \chi \sin \theta d \theta \\&=&m^2\frac{2/\tilde{\mathcal{W}}-1}{2\Omega^2\left(1-\frac{2}{3}\Omega^2 \right)^2}+\mathcal{O}\left(m^3\right). \nonumber \label{eq:minfty} \end{eqnarray} Using equations (\ref{eq:thetamax}) - (\ref{eq:minfty}) one finds \begin{equation} \alpha_w^2\approx m_\infty \frac{ 2/\tilde{\mathcal{W}}-1}{2\Omega^2\left(1-\frac{2}{3}\Omega^2 \right)}. \label{eq:alphaVSm} \end{equation} \subsection{Spin Evolution During Eruptions} \label{sec:SkirtTheory} In the previous section, we studied the effects of continuum driven winds on the evolution of angular momentum assuming a constant luminosity. As we shall now see, an abrupt change in luminosity has significant consequences regarding the evolution of the spin of the star. A sudden increase in luminosity can be understood as a result of an atmospheric phase transition, where the atmosphere becomes porous and the effective opacity drops \citep{ShavivNovae}. Consequently, the radiative flux at the surface becomes greater than the incoming convective flux. In order to compensate for the imbalance, the star must contract while radiating the difference in binding energy $U$:\footnote{Because of its low density, the atmosphere alone does not have sufficient gravitational binding energy. Note that the increase in luminosity is associated only with the release of binding energy due to contraction, which will eventually stop once nuclear reactions set in.} \begin{equation} \Delta L=L_{\textrm{erup}}-L_{\textrm{init}}=-\dot{U}\big|_M=-\frac{\partial U}{\partial R} \dot{R}, \end{equation} where $L_{\textrm{erup}}$ and $L_{\textrm{init}}$ are the luminosities during and before the eruption, respectively. A more convenient choice of parameters is $L_{\textrm{erup}}$ and $\lambda\equiv 1-L_{\textrm{init}}/L_{\textrm{erup}}$; henceforth, $L_{\textrm{erup}}$ will be denoted by $L$. In order to keep this discussion as generic as possible, we parametrize the binding energy as \begin{equation} U\equiv-\frac{GM^{2}}{R}\mathcal{B} \end{equation} and assume a constant $\mathcal{B}$, which is generally the case. The time scale for contraction is then \begin{equation} -\frac{R}{\dot{R}}=T_{\textrm{KH}}\mathcal{B}\lambda^{-1}, \end{equation} where $T_{\textrm{KH}}\equiv G M^2 /L R$ is the Kelvin-Helmholtz time scale. Plugging this into eq.~(\ref{eq:omegaevol}) gives \begin{equation} \frac{\lambda}{2}\left(T_{\textrm{KH}}\mathcal{B}\right)^{-1}+\left(\frac{\alpha_{w}^{2}}{\alpha_{g}^{2}}-\frac{3}{2}\right)\frac{\dot{M}}{M}=\frac{\dot{\Omega}}{\Omega}. \label{eq:AngMomCont} \end{equation} The contraction terminates when the temperature build up at the core becomes sufficient to generate the required luminosity through nuclear reactions. In order to estimate how long does the star need to contract, we assume that it behaves homologously and consider the scaling of the luminosity. If the specific energy production can be written as $\epsilon \propto \rho^p T^q$, we get \begin{equation} L = \int \epsilon \rho\, dV \propto \rho_c^{1+p} T_c^{q} R^3. \end{equation} In the following, we assume that $p=1$, which is suitable for most nuclear reactions, Hydrogen burning included. The temperature scales as $T \propto P/\rho$ when gas pressure dominates, and $T\propto P^{1/4}$ when radiation pressure dominates. If the dynamical time scale, $T_{\textrm{dyn}}\equiv 1/\sqrt{G\rho}$, is much shorter than the contraction time, one can safely assume mechanical equilibrium.\footnote{ For the case of \etacar, $T_{\textrm{dyn}}\approx$ 2 weeks.} Hydrostatics then give that $P\propto M^2/R^4$. We thus find $L \propto M^{2+q/\nu} R^{-3-q}$ with $\nu=1 ( \mathrm{or}~ 2)$, corresponding to the case where gas (or radiation) pressure dominates. Because $q$ is typically very large, the star needs to contract only by a relatively small amount \begin{equation} \frac{\Delta R_{\textrm{cont}}}{R} = 1 - \left( L_{\textrm{init}} \over L_{\textrm{erup}} \right)^{1/(3+q)}\left(1-\frac{\Delta M}{M}\right)^{(2+q/\nu)/(3+q)}. \label{eq:ContractionLength} \end{equation} The contraction will therefore take place for a duration \begin{equation} T_{\textrm{cont}}=\int \frac{dR}{\dot R}\approx T_{\textrm{KH}}\mathcal{B}\lambda^{-1}\frac{\Delta R_{\textrm{cont}}}{R}. \label{eq:ContractionTime} \end{equation}
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\label{sec:Conclusions} The twin lobe structure of the Homunculus nebula of \etacar\ strongly suggests that either fast rotation or binary interaction played an important role in the process of mass loss. In this work, we have chosen to pursue the idea that it is near critical rotation that sculpted the Homunculus, and that the mass loss was in the form of a continuum driven wind. The formation of the equatorial skirt is then associated with the evolution of angular momentum before the system settled in a steady state. As we have seen in \S\ref{sec:WindsTheory}, the structure of such winds can be obtained analytically, if one allows several simplifying assumptions. In particular, it was assumed that there is no latitudinal interaction between mass and radiative flux elements on different radial trajectories. This assumption later proved to be self consistent, as it turns out that most of the acceleration of the wind takes place near the surface.\footnote{Eq.~(\ref{Vr}) implies that the winds of \etacar\ reach half their terminal velocity already at $r \approx 1.15 R$. For comparison, line driven winds reach half their terminal velocity further out, at $r \approx 1.73R$ \citep[e.g.,][]{Lamers1999}.} Namely, the relevant radial scale is smaller than the latitudinal one; this acts to suppress lateral fluxes. By further simplifying the geometry and assuming a spherical star, the set of equations describing all conserved quantities became both algebraically closed and integrable. Obviously, this assumption cannot hold since rotation breaks spherical symmetry. Even so, we find that this simple model captures the gross features of the system, as depicted in figure~\ref{VofTheta}. Mainly, a twin-lobed structure emerges as a generic property of winds from such stars. The main ingredient required in order to form this structure is that of gravity darkening, as was already pointed out by \cite{Owocki1998} and \cite{Maeder1999} for line driven winds. Nevertheless, we do expect the oblate geometry to introduce distortion. For example, the predicted velocity may be lower near the equator, because the wind is launched from larger radii. This may in fact improve the agreement with the observations. For the specific case of \etacar\ (\S\ref{sec:WindApplication}), our model provides a rough description of the shape and size of the Homunculus nebula. In particular, continuum driven winds naturally explain the large wind momentum to photon luminosity observed in the Homunculus. Such high ratios cannot be explained by line driven winds. Note however, that within the approximations and the nominal parameters used, the typical mass of the Homunculus obtained is somewhat lower than the value recent observations imply \citep{Smith2003}, and the predicted luminosity is lower by a factor of $\sim 3$. These discrepancies may be reduced by a full numerical analysis. As a side note, we show in \S\ref{sec:spin} that the angular momentum gain due to a general spin-orbit coupling is correlated with the change in semi-major axis. As a consequence, in light of the high eccentricity of the orbit and stellar parameters assumed, the angular momentum \etacar\ may have received from interactions with its companion were insufficient to drive it to critical rotation. The fact that the rotation is very close to critical may seem like fine tuning, but it is quite possible that this is the natural state of the system. Since the star cannot sustain high load winds for long periods of time, it is reasonable that a low mass loss rate, as \etacar\ currently exhibits, is in fact the dominant mode of mass loss. Motivated by this reasoning, we analyzed in \S\ref{sec:Spinup} the explicit solutions for the wind velocity at low mass loss rates. We found that at equatorial latitudes, the wind stagnates and falls back to the surface, given that the star is rotating fast enough. The mass loss from pole centered winds, which do reach infinity, dominates over the loss of angular momentum and so the star spins up.\footnote{By ``spin-up" we refer to an increase of $\Omega\equiv\omega\,\sqrt{R^3/GM}$.} This is another aspect of continuum driven winds which is absent from the line driven winds model. We predict (\S\ref{sec:presentwind}) that in about $2\times10^3$ years at the current mass loss rate, \etacar\ will be back at critical rotation. This implies that if the previous large eruption cycle and subsequent wind were similar to the Great Eruption and the present wind, then the previous large eruption of \etacar\ must have taken place at least a few thousand years ago. We then showed in \S\ref{sec:SkirtTheory} that at the onset of LBV eruptions, the contraction that follows the atmospheric phase transition is fast enough to render equatorial latitudes unbound, if the star is critically rotating. The formation of an equatorial skirt is seen to be a generic consequence of the model as well. For the stellar parameters of \etacar\ (\S\ref{sec:SkirtApp}), we predict the skirt's mass to be of the order of $0.1\,M_\odot$. The effect of angular momentum diffusion, which is discussed in \S\ref{sec:AngMomDif}, complicates the estimation of the mass of the skirt. Taking it into account requires a more accurate knowledge of the mass loss to the wind, and specifically its time dependence. Essentially, the question is what happened first - did nuclear reactions kick in to stop the contraction or did the loss of angular momentum manage to diffuse inwards and render the formation of a skirt unnecessary? We await the appearance of new observational data regarding the time dependence of the Great Eruption.
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1206.6078
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1206
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1206.5692_arXiv.txt
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In this paper we describe and evaluate new spectral line polarisation observations obtained with the goal of mapping the surfaces of magnetic Ap stars in great detail. One hundred complete or partial Stokes $IQUV$ sequences, corresponding to 297 individual polarised spectra, have been obtained for 7 bright Ap stars using the ESPaDOnS and NARVAL high resolution spectropolarimeters. The targets span a range of mass from approximately 1.8 to 3.4~$M_\odot$, a range of rotation period from $2.56$ to $6.80$ days, and a range of maximum longitudinal magnetic field strength from 0.3 to over 4 kG. For 3 of the 7 stars, we have obtained dense phase coverage sampling the entire rotational cycle. These datasets are suitable for immediate magnetic and chemical abundance surface mapping using Magnetic Doppler Imaging (MDI). For the 4 remaining stars, partial phase coverage has been obtained, and additional observations will be required in order to map the surfaces of these stars. The median signal-to-noise ratio of the reduced observations is over 700 per 1.8~\kms\ pixel. Spectra of all stars show Stokes $V$ Zeeman signatures in essentially all individual lines, and most stars show clear Stokes $QU$ signatures in many individual spectral lines. The observations provide a vastly improved data set compared to previous generations of observations in terms of signal-to-noise ratio, resolving power and measurement uncertainties. Measurement of the longitudinal magnetic field demonstrates that the data are internally consistent within computed uncertainties typically at the 50 to 100$\sigma$ level. Data are also shown to be in excellent agreement with published observations and in qualitative agreement with the predictions of published surface structure models. In addition to providing the foundation for the next generation of surface maps of Ap stars, this study establishes the performance and stability of the ESPaDOnS and NARVAl high-resolution spectropolarimeters during the period 2006-2010.
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The classification Ap identifies a (main sequence) A or B type star which displays peculiar chemical abundances, usually combined with an observable magnetic field. Although other classes of chemically peculiar stars exist (e.g. Am stars, Hg-Mn stars, He-weak stars), these stars have been demonstrated to lack strong, organised magnetic fields at their surfaces (e.g. Shorlin et al. 2002, Wade et al 2006, Makaganiuk et al. 2011). Ap stars appear to be the only class of middle main-sequence stars for which, in all cases, an observable magnetic field is present (Auri\`erie et al. 2007). Since their discovery by Babcock in 1947, the magnetic fields of Ap stars have been established through observation to have important global dipole components with polar strengths ranging from hundreds to tens of thousands of gauss. The symmetry axis of the dipole component is almost always significantly tilted relative the stellar rotation axis. In addition, Ap stars generally spin much more slowly than non-peculiar stars of similar masses (Stepie\'n 2000), and as they spin they exhibit line profile variations attributed to rotational modulation of patchy, non-axisymmetric lateral and vertical distributions of chemical abundance in their photospheres. The distributions of abundance vary significantly from element to element: some are distributed relatively uniformly, while others show strong contrast; some are distributed in relatively simple patterns, while others show complex distributions. While it is generally accepted that the fundamental mechanism responsible for the chemical peculiarities is microscopic chemical diffusion (as described by Michaud 1970), the origin of chemical patchiness, and the relationship to the magnetic field, is poorly understood. The earliest studies of the magnetic field geometries of Ap stars interpreted the rotational variations of their longitudinal magnetic fields in the context of Stibbs' Oblique Rotator Model assuming a simple magnetic dipole field (e.g. Babcock 1947, 1951; Stibbs 1950). However, with the acquisition of increasingly sophisticated diagnostic data (the mean surface field (or mean field modulus), and high-resolution line profiles), it became clear that the large-scale field topologies exhibited important departures from the simple dipolar model (e.g. Preston \& Sturch 1967, Preston 1969, 1970, Landstreet 1970, 1988, 1989). Leroy and collaborators (Landolfi et al. 1993, Leroy et al. 1993, 1994, 1995ab, 1996, Bagnulo et al. 1995, Wade et al. 1996) systematically studied Ap stars using broadband linear polarisation measurements and models, constraining the transverse component of the magnetic field. Importantly, they found that differences between the observed linear polarisation variations and those predicted by the simple dipole model could not be fully explained by abundance inhomogeneities alone (Leroy et al. 1996). With this work, they established a modified dipolar model with a trend toward an outward expansion of the field lines over some parts of the magnetic equator, and showed the potential of linear polarisation for diagnosing small-scale structure of the magnetic fields of Ap stars. Thus, the observations and modeling undertaken during the latter half of the 20th century allowed progress from a simple view of the magnetic fields of Ap stars to a relatively sophisticated picture in which fields were known to show both global-scale and local-scale departures from a simple dipole. Leroy et al. (1996) commented that high-resolution spectropolarimetry represented the next step in furthering the study of the magnetic field geometry of Ap stars. Four years later, Wade et al. (2000a) published the first compendium of phase-resolved high-resolution spectropolarimetric observations of Ap stars in both circular and linear polarisation. Using the MuSiCoS spectropolarimeter, $R=35,000$ Stokes $IQUV$ spectra with a median S/N of 300 (per 4.6~\kms\ pixel) were obtained for 14 Ap stars. While the quality of the spectra was sufficiently good to show the shape and phase variation of all Stokes parameters in mean Least-Squares Deconvolved (LSD) line profiles, measurement in individual spectral lines was restricted to a few particularly strong lines, principally those of Fe~{\sc ii} multiplet 42. Nevertheless, the Stokes profiles of 53 Cam were used by Bagnulo et al. (2001) and Kochukhov et al. (2004) to evaluate published magnetic models developed based on less sophisticated data. Those authors found that models based on so-called "magnetic observables" (e.g Bagnulo 2000) led to derivation of surface magnetic field characteristics that were not consistent with the detailed Stokes profiles, and that both circular and especially linear polarisation profiles were required for realistic reconstruction of the field. Following these conclusions, Kochukhov et al. (2002) for $\alpha^2$~CVn and Lueftinger et al. (2010) for HD 24712 employed the new Magnetic Doppler Imaging technique (MDI), described by Piskunov \& Kochukhov (2002) and Kochukhov \& Piskunov (2002), to construct high resolution maps of the surface vector magnetic field maps using Stokes $IV$ observations and by preferring a global low-order multipolar field structure. Maps using Linear polarisation profiles (Stokes $Q$ and $U$) made by Kochukhov et al. (2004) and Kochukhov \& Wade (2010) for the Ap stars 53 Cam and $\alpha^2$~CVn. These maps were distinguished from earlier models in that they were computed directly from the observed polarised line profiles, making no {\em a priori} assumptions regarding the large-scale or small-scale topology of the field. The MDI surface magnetic field maps of both stars revealed that their magnetic topologies depart significantly from low-order multipoles. In particular, both studies concluded that while the global topology of the magnetic field was reasonably smooth, the strength of the field was quite patchy, indicating complex structure on relatively small scales. Simultaneous mapping of the distributions of the surface chemical abundances of several elements was also performed, allowing a comparison between the local field properties and local photospheric chemistry. It is important to note that the observational material used in the MDI studies of 53 Cam and $\alpha^2$ CVn represented the best data sets obtained from several years of MuSiCoS observations. In those spectra the uniquely valuable Stokes $Q$ and $U$ Zeeman signatures were only clearly detectable in 3 strong lines, with a significance (i.e. amplitude divided by error bar) of 5 or less. The relatively low signal-to-noise ratio and resolving power achievable with the MuSiCoS instrument led to some ambiguity in the field reconstruction, and limited the useful sample of stars to those with bright apparent magnitudes, strong fields and sharp lines. Because MDI exploits the indirect resolution of the stellar disc due to stellar rotation, this means that those stars best suited to reconstruction (relatively rapidly-rotating stars, with consequentially weaker Stokes profiles) were inaccessible to MuSiCoS. As a result, only an extremely limited range of stellar properties (rotation, mass, temperature, magnetic field, etc.) which may influence the phenomena of interest could be studied using the MuSiCoS data. To address outstanding questions surrounding the detailed magnetic structure of Ap stars and the effect of the magnetic field on atmospheric chemical transport processes, we have acquired new higher-resolution and signal-to-noise Stokes $IQUV$ spectra of a small sample of well-studied magnetic Ap stars using the new generation of high-resolution spectropolarimeters. In this paper we describe the observations obtained. We demonstrate the stability of the instrumentation during the 5 years of observation by evaluating the internal and external agreement of the data. We illustrate the quality of the observed Stokes profiles, comparing with MuSiCoS results and demonstrating that they represent a qualitative step forward in our ability to diagnose the magnetic structure of Ap stars.
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The goal of this project was to obtain a new data set in all four Stokes parameters for a selection of well studied Ap stars, with the ultimate aim to map these stars using Magnetic Doppler Imaging. The target list contained stars which span a large part of the parameter space of interest, with sufficient signal-to-noise ratio to not only greatly improve on the previous observations, but to also be suitable for MDI mapping. The final selection was based primarily on stars already identified by Wade et al. (2000a) as promising candidates for such study. Early on in the project it was clear that both ESPaDOnS and NARVAL have greatly improved the level of detail at which Ap stars can be studied. The resulting dataset obtained for this study is far superior to that obtained previously with MuSiCoS, and represents some of the highest resolution phase-resolved observations of Ap stars acquired to date. This data set has more individual lines showing variation, much improved signal-to-noise and smaller error bars associated with measurements of the longitudinal field and net linear polarisation. The new data have been shown to be consistent with the previous observations of Wade et al. (2000a) and also those of Leroy at al. (1995), with most targets agreeing well between the different epochs. Surprisingly we found that even when the data are of such high signal-to-noise and when the magnetic fields are strong, the LSD analysis is sensitive to the normalisation and the measured magnetic field is rather sensitive to the integration ranges chosen, with variations of sometimes on the order of 100 gauss with very small changes of the integration range. A key conclusion of this work is that even with such high-quality data, extreme care must still be taken with all stages of analysis to ensure consistent results at this level of precision. Although crosstalk was originally a concern, through a series of experiments we have shown that it is at a level which should not have a significant impact on the results. By using observations of $\gamma$ Equ, we have seen that the highest level of crosstalk in Stokes $Q$ is still within the noise and slightly above the noise in Stokes $U$ (around the 5 \% level.). It should be noted that even at these levels, this effect will be less significant in the broader-lined stars studied here. Considering this, we believe the other uncertainties associated with the analysis techniques have a greater effect: normalisation, blending, line masks used for LSD and the choice of integration ranges used for longitudinal field measurements. But as regards the final impact of the crosstalk on MDI mapping, this will be discussed and addressed in a future paper (Silvester et al. in preparation). With these high quality observations, we suspect that the limitations for mapping will in fact come not from the data (with strong Stokes $Q$ and $U$ signatures seen in many individual lines), but the ability to deal with line blends within the MDI code. An important result of this study is the confirmed stability of the global properties of the magnetic fields of these Ap stars. Over multiple epochs of observations the fields have remained constant, with little variation in both longitudinal field and linear polarisation measurements. In some cases measurements separated by over a decade still agree with each other within the uncertainties. By comparing MDI maps produced from the new observations of $\alpha^2$ CVn with those of Kochukhov and Wade (2010), we can potentially test for evolution of the field geometry which may occur on small spatial scales. Considering the longitudinal magnetic field measurements, agreement was found between the new measurements and those of Wade et al. (2000b), with the exception of HD 71866 and $\alpha^2$ CVn which showed a slight discrepancy. By re-reducing the MuSiCoS data with the new masks, the observations were brought into agreement. Whilst the general shapes of the netlinear polarisation variations were in agreement, one needed to invoke free parameters such as scaling which is consistent with what was described in Wade et al. (2000b). For any study which requires observations over multiple semesters, it is imperative that the instrument is stable and consistent throughout the campaign. Both NARVAL and ESPaDOnS proved to be very stable instruments, with resolution and signal-to-noise being constant over the 4 years of data. Indeed we have also shown the two instruments are consistent with one another with close to identical result from similar phases. These facts demonstrate that ESPaDOnS and NARVAL are both very capable instruments, well suited to high-resolution four Stokes measurements of magnetic stars over multiple year time scales. One of the targets ($\alpha^2$ CVn) has already been mapped with MDI using MuSiCoS data by Kochukhov \& Wade (2010). HD 112413 is an ideal star for determining how much of an improvement the new polarimetric data could give to MDI mapping. To quantify this improvement and to further confirm consistency, the new observations of $\alpha^2$ CVn were compared with the profiles predicted by the model of Kochukhov \& Wade (2010). As shown in Fig. \ref{2cvnmapfit}, good general agreement between the new observations and the MuSiCoS-derived model is observed. However, the new profiles show more complexity than was present in the MuSiCoS data, which is not fully reproduced in the current model and would likely require a more complex magnetic field distribution. In addition, the Stokes V profile amplitude is significantly underestimated by the model at a number of phases. The complete phase coverage of both 49 Cam (Silvester et al. in preparation), $\alpha^2$ CVn and HD 32633 will allow the completion of 4 Stokes parameter MDI maps for these stars, doubling the number of Ap stars studied using this technique. Out of the remaining targets HD 4778, HD 71866 and HD 118022 would also be a worthwhile candidates for MDI; conversely HD 40312 has small linear polarisation signatures, relative to the noise in their spectra, making them less suitable for MDI analysis. The mapping of 49 Cam is well underway and the results will be presented in Paper II (Silvester et al. in preparation). \begin{figure*} \begin{center} \includegraphics[width=0.85\textwidth]{2cvn-esp-mus-mdi.eps} \caption{Comparison betwen the new ESPaDOnS/NARVAL data for the Fe~{\sc ii} 5018 line (shown by black points) with the final model profiles adopted in the mapping of HD 112413 (by solid blue lines) by Kochukhov \& Wade (2010) } \label{2cvnmapfit} \end{center} \end{figure*}
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1206.5692
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1206
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1206.0554_arXiv.txt
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We investigate whether the late-time (at $z\leq 100$) velocity dispersion expected in Warm Dark Matter scenarios could have some effect on the cosmic web (i.e., outside of virialized halos). We consider effective hydrodynamical equations, with a pressurelike term that agrees at the linear level with the analysis of the Vlasov equation. Then, using analytical methods, based on perturbative expansions and the spherical dynamics, we investigate the impact of this term for a $1$keV dark matter particle. We find that the late-time velocity dispersion has a negligible effect on the power spectrum on perturbative scales and on the halo mass function. However, it has a significant impact on the probability distribution function of the density contrast at $z \sim 3$ on scales smaller than $0.1 h^{-1}$Mpc, which correspond to Lyman-$\alpha$ clouds. Finally, we note that numerical simulations should start at $z_i\geq 100$ rather than $z_i \leq 50$ to avoid underestimating gravitational clustering at low redshifts.
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\label{Introduction} In the standard cosmological $\Lambda$ cold dark matter ($\Lambda$CDM) scenario, most of the matter content of the Universe is made of CDM particles, which are cold and collisionless. This means that they have a negligible velocity dispersion during the matter-dominated era and density fluctuations on almost all scales (except very small scales at early times) grow through gravitational instability. This leads to a hierarchical scenario for the formation of large-scale structures, as the amplitude of density fluctuations at early time (e.g., at the beginning of the matter-dominated era) is larger on smaller scales. Then, small scales turn nonlinear first and merge to build increasingly large and massive structures as larger scales become nonlinear in the course of time \cite{Peebles1980}. This scenario (with an extra dark energy component or cosmological constant) is in good agreement with a large variety of cosmological observations, such as the cosmic microwave background (CMB) \cite{Komatsu2011} and galaxy surveys \cite{Tegmark2006}. However, this CDM model may disagree from observations on small scales (below the size of galaxies). Thus, CDM simulations typically predict too many satellite galaxies around Milky-Way-sized central galaxies as compared with observations \cite{Moore1999,Springel2008,Trujillo-Gomez2011}. They also predict power-law density profiles, $\rho \sim r^{-1}$, in the center of virialized halos \cite{Navarro1997}, whereas dark-matter-dominated dwarf galaxies \cite{Burkert1995} and some disk galaxies \cite{Salucci2000} exhibit flat density cores. This is the so-called ``core-cusp problem'' \cite{de-Blok2010}. One possible solution to these small-scale problems is a warm dark matter (WDM) scenario, with dark matter particles of a mass on the order of $1$keV. This intermediate case between the ``cold'' and ``hot'' dark matter scenarios provides a non-negligible velocity dispersion and a significant free-streaming that erases density fluctuations on small scales (mostly during the period where the particles are relativistic). This helps to cure the small-scale problems of the CDM scenario, while being indistinguishable from CDM on large scales, which preserves its good agreement with large-scale observations such as the CMB and galaxy surveys \cite{Bode2001,Avila-Reese2001,Menci2012,Lovell2012}. This favors a mass on the order of $1$keV \cite{de-Vega2010,de-Vega2012}. For larger masses we recover the CDM scenario and for smaller masses we recover the hot dark matter scenario, where structure formation begins too late (in particular, this is ruled out by the Gunn-Peterson bound \cite{Gunn1965}: quasar spectra show that the Universe must have been reionized before $z \sim 6$, which requires galaxy formation by this time). Similar lower bounds on $m$ are also obtained from the observed velocity dispersion of dwarfs galaxies and from the Lyman-$\alpha$ forest \cite{Viel2005,Seljak2006,Abazajian2006a,Boyarsky2009}. We must note that these small-scale problems may also be cured by the physics of the baryonic component, within the CDM scenario. For instance, reionization of the intergalactic medium \cite{Bullock2000,Benson2002} or feedback from stars and supernovae \cite{Kauffmann1993} suppress star formation in small satellite halos. Only a small fraction of the low-mass dark matter satellites would then shine in the sky and appear in galaxy surveys. This would reconcile the observed abundance with the CDM prediction but there remains some discrepancy for the shape of the satellite luminosity function \cite{Koposov2008}. Supernovae explosions may also transform a cusp density profile into a cored one, within small dark matter halos \cite{Mashchenko2006,Governato2010,de-Souza2011}. However, it is a difficult task to check that such models can explain galaxy properties from massive to dwarf galaxies and from $z=0$ to higher redshifts \cite{Font2011}. Therefore, WDM scenarios remain interesting alternatives to CDM that are still being investigated in many works. As recalled above, a particle mass on the order of $1-10$keV is a good candidate and it may correspond, for instance, to sterile neutrinos \cite{Dodelson1994,Abazajian2001,Abazajian2006,Shaposhnikov2006,Boyarsky2009a,Kusenko2009,Abazajian2012} or to gravitinos \cite{Kawasaki1997,Gorbunov2008}. At early times and on large scales, the formation of large-scale structures within WDM scenarios is studied through the linearized Vlasov equation \cite{Boyanovsky2008,Boyanovsky2011,Boyanovsky2011a,de-Vega2012a,de-Vega2012b}. At low redshift and on small scales, the nonlinear regime of gravitational clustering is investigated through numerical simulations \cite{Bode2001,Schneider2011} and halo models based on such simulations \cite{Smith2011,Dunstan2011}. In practice, one often uses the same N-body codes as for CDM scenarios and the only difference comes from the density power spectrum that is set at the initial redshift $z_i$ of the simulations. This means that one takes into account the high-$k$ cutoff due to free-streaming during the relativistic era but neglects the nonzero velocity dispersion at low redshifts, $z\leq z_i$. This is legitimate because the relative importance of this late-time velocity dispersion decreases with time and the main difference between the CDM and WDM scenarios with respect to large-scale structures arises from the high-$k$ cutoff of the power spectrum built during the relativistic regime. Nevertheless, it would be interesting to have a quantitative check of this approximation. This is the goal of this paper, where we obtain a quantitative estimate of the impact of the late-time WDM velocity dispersion on the formation of large-scale structures. We also estimate the sensitivity of the gravitational clustering measured at low redshift on the initial redshift $z_i$ of the simulations. Here we do not consider the inner regions of virialized halos, where the finite velocity dispersion can have important effects because of Liouville theorem. Indeed, this implies an upper bound on the coarse-grained phase-space distribution function \cite{Tremaine1979}, which can lead to cored density profiles instead of cusps \cite{Hogan2000} (but the behavior in central regions remains difficult to predict \cite{Vinas2012}). In contrast, we consider the cosmic web, that is, moderate density fluctuations or large scales, as well as the halo mass function itself. Then, using perturbative methods or the spherical dynamics, we compare such statistics (the power spectrum on large perturbative scales, the halo mass function, and the probability distribution of the density contrast) between the CDM and WDM scenarios, where we neglect or take into account the late-time WDM velocity dispersion. To simplify the analysis and to go beyond the linear regime, we use effective equations of motion similar to standard hydrodynamics. They involve a simplified pressurelike term in the Euler equation, associated with the late-time velocity dispersion, that is chosen so as to agree with results from the Vlasov equation at linear order. This should be sufficient for our goal, which is only to estimate the order of magnitude of the impact of this late-time WDM velocity dispersion. Thus, our study is complementary to Ref.\cite{Boyanovsky2011a} who investigates the effects of the velocity dispersion at low redshift through the linearized Vlasov equation. Our approach is not exact, because we use a fluid approximation, but it allows us to consider nonlinear density fluctuations. In particular, our goal is not to study in accurate details a specific WDM model but to investigate the generic impact of a late-time velocity dispersion on the cosmic web. After describing these effective equations of motion and our approach in Sec.~\ref{motion}, we present our results for a $1$keV dark matter particle in Sec.~\ref{Results}. We also consider the impact of the choice of initial redshift in numerical simulations in Sec.~\ref{lower-zi} and we conclude in Sec.~\ref{Conclusion}.
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\label{Conclusion} Using an effective Euler equation, that agrees with the Vlasov equation at the linear level (except for subdominant memory terms), we have estimated the impact of a late-time WDM velocity dispersion on the formation of large-scale structures. We have only considered the ``cosmic web'', that is, large perturbative scales, moderate density fluctuations, and the number counts of virialized halos, which can be studied with analytic tools. We have focused on the case of a $1$keV dark matter particle, which is representative of current WDM scenarios (lower masses are excluded by observations, such as Lyman-$\alpha$ forest data, while higher masses become indistinguishable from the CDM limit). We find that on perturbative scales, the deviation of the density power spectrum from the CDM case is only on the order of $1\%$, at $z \leq 5$, even though it is slightly amplified by the nonlinear dynamics. This is below the accuracy of the standard perturbative expansion and requires efficient perturbative schemes. On the other hand, the effects of the late-time velocity dispersion are negligible over most of the perturbative range at $z \leq 5$ (so that one could use the same perturbative approaches devised for the CDM case). We also find that the late-time velocity dispersion has a negligible impact on the halo mass function at $z \leq 5$ (in any case, below the $10\%$ accuracy that can be guaranteed by simulations), although at very low mass and high redshift, the cutoff may become sharper \cite{Barkana2001,Benson2012}. On the other hand, it has a non-negligible effect on the probability distribution of the density contrast on scales $x \leq 0.1 h^{-1}$Mpc at $z=3$. This means it should have some impact on the probability distribution of the Lyman-$\alpha$ flux decrement, measured on the spectra of distant quasars. Finally, we note that numerical simulations should use a high initial redshift, $z_i \geq 100$, rather than a low value, $z_i \leq 50$. Indeed, such a low initial redshift can lead to a significant underestimation of the power spectrum on perturbative scales and of the large-mass tail of the halo mass function, which is larger than the true signal associated with the WDM scenario (but of course, on smaller scales and on the low mass tail of the mass function, one is again dominated by the actual WDM signal). A low initial redshift does not help much either to reduce the effect of the late-time velocity dispersion for the probability distribution of the density contrast on scales associated with Lyman-$\alpha$ clouds. To go beyond the effective hydrodynamical equations used in this work, one should use the nonlinear Vlasov equation itself. However, this is a difficult task because of the additional velocity coordinates, which makes numerical implementations significantly heavier already for the CDM scenario \cite{Valageas2004,Tassev2011}. An alternative would be to extend the fluid approximation to higher orders \cite{Shoji2010}, by including equations of motion for the velocity moments of the Vlasov equation up to some higher order $n \geq 3$. However, because most of the WDM signal arises from nonperturbative scales, such a task may not be very rewarding, unless one builds methods that can be applied to the Lyman-$\alpha$ forest clouds, for instance.
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Recent observational results for magnetic fields in molecular clouds reviewed by Crutcher (2012) seem to be inconsistent with the predictions of the ambipolar diffusion theory of star formation. These include the measured decrease in mass to flux ratio between envelopes and cores, the failure to detect any self-gravitating magnetically subcritical clouds, the determination of the flat PDF of the total magnetic field strengths implying that there are many clouds with very weak magnetic fields, and the observed scaling $B \propto \rho^{2/3}$ that implies gravitational contraction with weak magnetic fields. We consider the problem of magnetic field evolution in turbulent molecular clouds and discuss the process of magnetic field diffusion mediated by magnetic reconnection. For this process that we termed ``reconnection diffusion'' we provide a simple physical model and explain that this process is inevitable in view of the present day understanding of MHD turbulence. We address the issue of the expected magnetization of cores and envelopes in the process of star formation and show that reconnection diffusion provides an efficient removal of magnetic flux that depends only on the properties of MHD turbulence in the core and the envelope. We show that as the amplitude of turbulence as well as the scale of turbulent motions decrease from the envelope to the core of the cloud, the diffusion of the magnetic field is faster in the envelope. As a result, the magnetic flux trapped during the collapse in the envelope is being released faster than the flux trapped in the core, resulting in much weaker fields in envelopes than in cores, as observed. We provide simple semi-analytical model calculations which support this conclusion and qualitatively agree with the observational results. Magnetic reconnection is also consistent with the lack of subcritical self-gravitating clouds, with the observed flat PDF of field strengths, and with the scaling of field strength with density. In addition, we demonstrate that reconnection diffusion process can account for the empirical Larson (1981) relations and list a few other implications of the reconnection diffusion concept. We argue that magnetic reconnection provides a solution to the magnetic flux problem of star formation that agrees better with observations than the long-standing ambipolar diffusion paradigm. Due to the illustrative nature of our simplified model we do not seek quantitative agreement, but discuss the complementary nature of our approach to the 3D MHD numerical simulations.
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Interstellar media are known to be turbulent and magnetized, and both turbulence and magnetic field are important for star formation (see Armstrong et al. 1994, Chepurnov \& Lazarian 2009, Crutcher 2012). The existing star formation paradigm has been developed with the concept that near-perfect flux freezing holds, i.e. that the magnetic field is well coupled with ions and electrons in the media (Alfv\'{e}n 1942). This is assumed in magnetically-mediated star formation theory, which was founded by the pioneering studies by L. Mestel and L. Spitzer (see Mestel \& Spitzer 1956, Mestel 1966) and brought to a high level of sophistication by other researchers (see Shu, Adams \& Lizano 1987, Mouschovias 1991, Nakano et al. 2002, Shu et al. 2004, Mouschovias et al. 2006). According to the theory, magnetic fields slow down and even prevent star formation if the media are sufficiently magnetized. The theory {\it assumes} that the change of the flux to mass ratio happens due to {\it ambipolar diffusion}, i.e. to the drift of neutrals which do not feel magnetic fields directly, but only through ion-neutral collisions. Naturally, in the presence of gravity, neutrals get concentrated towards the center of the gravitational potential while magnetic fields resist compression and therefore leave the forming protostar (e.g. Mestel 1965). The rate of ambipolar diffusion for a cloud in gravitational equilibrium with the magnetic field depends only on the degree of ionization of the media. The existing theory makes star formation inefficient for magnetically dominated (i.e. subcritical) clouds. The low efficiency of star formation corresponds to observations (e.g. Zuckerman \& Evans 1974), which is usually interpreted as a strong argument in support of the above scenario. This however does not solve all the problems; at the same time, for clouds dominated by gravity, i.e. supercritical clouds, this scenario does not work as magnetic fields do not have time to leave the cloud through ambipolar diffusion. Therefore for supercritical clouds magnetic fields should be dragged into the star, forming stars with magnetizations far in excess of the observed ones (see Galli et al. 2006, Johns-Krull 2007). Magnetic fields are important at all stages of star formation. In many instances the ideas of star formation based exclusively on ambipolar diffusion have been challenged by observations (Troland \& Heiles 1986, Shu et al. 2006, Crutcher et al. 2009, 2010a, see Crutcher 2012 for a review). While the interpretation of particular observations is the subject of scientific debates (see Mouschovias \& Tassis 2009), it is suggestive that there may be additional processes that the classical theory does not take into account. The primary suspect is turbulence, which is ubiquitous in interstellar media and molecular clouds\footnote{The presence of turbulence is an observational fact, while its sources are debatable. Supernova explosions, different instabilities, e.g. magnetorotational instability, have been discussed in the literature. Flows induced by the gravitational instabiltiy (see Vazquez-Semadeni et al. 2011) are also expected to be turbulent.} (see Larson 1981, Armstrong et al. 1994, Elmegreen \& Falgarone 1996, Lazarian \& Pogosyan 2000, Stanimirovic \& Lazarian 2001, Heyer \& Brunt 2004, Padoan et al. 2006, 2009, Chepurnov \& Lazarian 2010, Burkhart et al. 2010, 2012). Turbulence has revolutionized the field of star formation (see Vazquez-Semadeni et al. 1995, 2000, Ballesteros-Paredes et al. 1999, 2007, Elmegreen 2000, 2002, McKee \& Tan 2003, Elmegreen \& Scalo 2004, MacLow \& Klessen 2004, McKee \& Ostriker 2007), but the treatment of the turbulent magnetic fields stayed within the flux freezing paradigm. The cornerstone concept of magnetic flux freezing has been challenged recently. On the basis of the model of fast magnetic reconnection in Lazarian \& Vishniac (1999, henceforth LV99), Lazarian (2005) claimed that the removal of magnetic fields from turbulent plasma can happen during star formation due to magnetic reconnection rather than slow ambipolar drift (see also Lazarian \& Vishniac 2009). The process that was later termed {\it reconnection diffusion} does not depend on the degree of ionization, but rather on the properties of turbulence. The numerical confirmation of the idea was presented for molecular clouds and circumstellar accretion disks in Santos-Lima et al. (2010, 2012). The numerical testing of the LV99 reconnection model that is at the core of the concept of reconnection diffusion was successfully tested in Kowal (2009, 2012). Compared to this testing of reconnection diffusion, the testing in Kowal et al. (2009, 2012) was performed with much higher numerical resolution of the reconnection region and much better control of turbulence and other input parameters. In addition, more recent formal mathematical studies aimed at understanding of magnetic field dynamics in turbulent fluids supported the LV99 quantitative conclusions (Eyink 2011, Eyink, Lazarian \& Vishniac 2011). Moreover, the state of observations of magnetic fields in regions of star formation has progressed significantly (e.g. Crutcher 2012), so that comparisons of observational results with the theory are possible. These developments motivate us to further study the implications of the reconnection diffusion concept for star formation. We note that the interstellar medium is collisional (Yamada et al. 2006), and therefore ideas of collisionless reconnection (see Shay \& Drake 1998, Shay et al. 1998, Bhattacharjee et al. 2005, Cassak et al. 2006) are not applicable. Even if they were, they would not increase the rate of reconnection diffusion, which is controlled by the eddy diffusivity. Similarly, the ideas related to the plasmodial instability (see Shibata \& Tanuma 2001, Lourreiro et al. 2007, Uzdensky et al. 2010, Huang et al. 2011) do not change the reconnection rates either.\footnote{In special circumstances when the initial configuration of magnetic field contains a magnetic reversal and is laminar, the plasmoid reconnection can induce turbulence in 3D (Karamabadi 2012). This may be important process for Solar flares (LV99, Lazarian \& Vishniac 2009, Eyink et al. 2011). However, observations (see Larson 1981, Armstrong et al. 1995, Chepurnov \& Lazarian 2010, Gaensler et al. 2011, see also Elmegreen \& Scalo 2006 and Lazarian 2009 for reviews) testify that interstellar media and molecular clouds are turbulent which makes LV99 mechanism directly relevant.} In what follows, in \S 2 we discuss the reconnection of magnetic fields in turbulent fluids, in \S 3 we discuss recent observational results reviewed by Crutcher (2012) that can be interpreted as contradicting the model of star formation mediated by ambipolar diffusion. We provide in \S 4 a very simple semi-analytical model of magnetic field diffusion out of magnetized clouds and discuss its predictions. In \S 5 we discuss other observational tests. Astrophysical implications of reconnection diffusion are briefly discussed in \S 6, while the discussion and summary are provided in \S 7 and \S 8 respectively.
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\subsection{Reconnection diffusion: simulations and estimates} This paper was motivated by the observational results reviewed by Crutcher (2012), who showed that several important characteristics were not compatible with the ambipolar diffusion theory. We have shown that reconnection diffusion can produce agreement with the observations. Reconnection diffusion is based on the numerically tested LV99 model of fast reconnection in turbulent media as well as more recent insights into the violation of the frozen-in condition in turbulent magnetized fluids (Eyink 2011, ELV11). It is also appeals to the direct 3D MHD simulations of the magnetic field diffusion in Santos-Lima et al. (2010, 2012) that showed consistency with the statistical description of the magnetic diffusion in turbulent fluids (see Lazarian 2006). Therefore in the present paper we use this statistical description directly. This allows us to get insight into the evolution of magnetic fields at the core and the envelope. The approach to studying reconnection diffusion that is based on 3D numerical simulation has its own limitations. In Santos-Lima et al. (2010) we used an external gravitational potential. In more recent simulations in Leao et al. (2012) self-gravity is used. However, using an isothermal code to make the turbulent cloud more well-defined in the absence of the confining pressure of the ambient gas at a different temperature, we imposed an additional gravitational potential. We are experimenting with the ideas of introducing the ambient pressure without compromising the resolution and the simplicity of the simulation interpretation, but this is work in progress. In this situation, it is advantageous to get estimates and analyze idealized situations with clear physical meaning the way we do in this paper. Naturally, our simplified calculations do not seek the quantitative agreement with the observations, but only illustrate the qualitative behavior. Indeed, within the adopted toy model we do not evolve density together with magnetic field, but consider the diffusion of magnetic field out of an idealized cloud of constant density\footnote{A possible initial set up with the cloud of gas density $\sim 60$ cm$^{-3}$, radius 2.5 pc and the initial magnetic field of $\sim 6\times 10^{-6}$G may be advantageous. Such a cloud would be subcritical, which agrees with the notion of diffuse H~I clouds being subcritical. Then the reconnection diffusion would remove flux such that the cloud could collapse, with density building up to $10^4$ cm$^{-3}$ in the $r=0.2$ pc core and up to $10^3$ cm$^{-3}$ in the $r=1$ pc envelope, which are the mean parameters of the clouds observed by Crutcher, Hakobian \& Troland (2009). This is what we will try in future.}. We can mention that it is possible to argue that the process of reconnection diffusion is present in some of high resolution numerical simulations. However, without clear identification of the role of turbulence in fast reconnection, one may not be sure when the results are due to physically motivated reconnection diffusion and when they are the consequence of the bogus effects of numerical diffusion. For instance, Crutcher, Hakobian \& Troland (2009) refer to the simulations in Luntilla et al. (2009) that produce, in agreement with observations, higher magnetization of the cloud cores. If these cores are of the size of several grid units across, numerical effects rather than reconnection diffusion may be dominant and turbulence is suppressed at these scales. Our results show that in the presence of reconnection diffusion it is natural to expect that the magnetic field diffuses faster from the area of the envelope and the magnetic field strength is larger at the cloud core. This conclusion does not really depend on the particular model of turbulence and its distribution in the molecular cloud as far as those are constrained by existing observations. We should also stress that the 3D numerical simulations like those in Santos-Lima et al. (2010, 2012) {\it by themselves} do not provide the description of reconnection diffusion in realistic astrophysical environments. The interpretation of the results requires the proper understanding of scaling of magnetic reconnection with the dimensionless combination called the Lunquist number $S\equiv (L_{cur. sh} V_A/\eta)$, where $L_{cur. sh.}$ is the extent of the relevant current sheet, $\eta$ is Ohmic diffusivity. The Lundquist numbers in molecular clouds and in the corresponding simulations differ by a factor larger than $10^5$. In this situation one can establish the correspondence between the numerical simulations and astrophysical reality only if the reconnection does not depend on $S$. The independence from $S$ of magnetic reconnection is the conclusion of LV99 model. This model has been tested in Kowal et al. (2009, 2012) via a set of dedicated numerical simulations that confirmed the scaling predictions in LV99. This work exemplifies the advantages of the synergy of scaling arguments with numerical simulations as opposed to the ``brute force numerical approach'', which may not be productive while dealing with turbulence. \subsection{Predictions of the reconnection and ambipolar diffusion models} Within the concept of ambipolar diffusion the explanation of observational results reviewed by Crutcher (2012) is extremely difficult. Our point is that ambipolar diffusion is not the only process that can be responsible for the removal of magnetic flux from the media. We show that reconnection diffusion can reduce the magnetic field strength in the envelope with respect to the core. On a more fundamental level, reconnection diffusion changes the mass to flux ratio allowing magnetic field to diffuse away from the center of the gravitational potential. The available observations provide a rather mild constraint, namely, that mass/flux does not increase as fast in the core as ambipolar diffusion predicts. In many ways, ambipolar diffusion is a very special type of diffusion, the efficiency of which drops towards more ionized outer regions of the envelope inducing more efficient flux loss from cores. Typically the observed column densities through the cores studied by Crutcher et al. (2009) are about twice those through the envelope regions. If mass/flux were constant, then the fields in the envelopes would be 1/2 those in the cores. So just saying that the fields in the envelopes are less than those in the cores does not require reconnection diffusion, or any kind of diffusion. Reconnection diffusion is required if the envelope fields are even weaker than half of the core fields; this is what is observed. Ambipolar diffusion would require that envelope fields are stronger than half of the core fields. \subsection{Reconnection diffusion and ambipolar drift in turbulent media} Within this study we do not appeal to ambipolar diffusion, which is acceptable when the reconnection diffusion provides larger diffusivity for magnetic fields. One may argue that this is a generic situation in the presence of turbulence. For instance, Heitsch et al. (2004, henceforth HX04) performed 2.5D simulations of turbulence with two-fluid code and examined the decorrelation of neutrals and magnetic field that was taking place as they were driving the turbulence. The study reported an enhancement of the ambipolar diffusion rate compared to the ambipolar diffusion acting in a laminar fluid. HX04 correctly associated the enhancement with turbulence creating density gradients that are being dissolved by ambipolar diffusion (see also Zweibel 2002). While in 2.5D simulations of HX04 the numerical set-up precluded reconnection from happening (as magnetic field was perpendicular to the mixing plane and magnetic fields never crossed each other at an angle), the authors reported an enhanced rate that is equal to the turbulent diffusion rate $L V_L$, which is the result expected as a limiting case of the reconnection diffusion prediction for the special set-up studied. While we agree with HX04 in terms of the importance of turbulence, we suggest that it is misleading to refer to the process as ``turbulent ambipolar diffusion''. We believe that ambipolar diffusion does not play a role in the enhancement measured and the observed effect is entirely due to turbulence\footnote{A similar process takes place in the case of molecular diffusivity in turbulent hydrodynamic flows. The result for the latter flows is well known: in the turbulent regime, molecular diffusivity is irrelevant for the turbulent transport. Indeed, in the case of high microscopic diffusivity, turbulence provides mixing down to a scale $l_1$ at which the microscopic diffusivity both suppresses the cascade and ensures efficient diffusivity of the contaminant. In the case of low microscopic diffusivity, turbulent mixing happens down to a scale $l_2\ll l_1$, which ensures that even low microscopic diffusivity is sufficient to provide efficient diffusion. In both cases the total effective diffusivity of the contaminant is given by the product of the turbulent injection scale and the turbulent velocity.}. We note that, in the presence of turbulence, the independence of the gravitational collapse from the ambipolar diffusion rate was reported in numerical simulations by Balsara, Crutcher \& Pouquet (2001). Ambipolar diffusion may control flux removal from the laminar cores. In the presence of turbulence, it is reconnection diffusion that dominates. The results of HX04 support this notion, as the authors in their set up report the transition to turbulent diffusivity in the presence of turbulence. The set-up in HX04 precludes the magnetic fields from reconnection as field lines perform mixing motions being kept absolutely parallel. In any realistic turbulent 3D flow magnetic reconnection will be essential. \subsection{Reconnection diffusion and hyper-restivity concept} Another process of diffusive nature is related to ``hyperresistivity'' or enhanced physical resistivity of turbulent fluids. To explain fast removal of magnetic field from accretion disks Shu et al. (2006) appealed to the hyperrestivity concept (Strauss 1986, 1988, Bhattacharjee \& Hameiri 1986, Hameiri \& Bhattacharjee, Diamond \& Malkov 2003). The studies introducing hyper-resistivity attempt to derive the effective resistivity of the turbulent media in the context of the mean-field resistive MHD. Using magnetic helicity conservation the authors derived the electric field. Then, integrating by part, they obtained a term which could be identified with effective resistivity proportional to the magnetic helicity current. There are several problems with this derivation. In particular, the most serious is the assumption that the helicity of magnetic field and the small scale turbulent fields are separately conserved, which erroneously disregard the magnetic helicity fluxes through open boundaries that is essential for fast stationary reconnection (see ELV11). In more general terms, hyper-resistivity idea is an incarnation of the mean-field approach to producing fast reconnection. As explained in ELV11, the problem of such approaches is that the lines of the actual astrophysical magnetic field should reconnect, not the lines of the mean field. Therefore the correct approach to fast reconnection should be independent of spatial and time averaging. All in all, we believe that the concept of hyper-resistivity should not be applied to astrophysical environments. \subsection{Reconnection diffusion and ``magnetic turbulent diffusivity'' concept} One may claim that the reconnection diffusion concept extends the concept of hydrodynamic turbulent diffusion to magnetized fluids. The physics of it is very different from the ``magnetic turbulent diffusivity'' idea discussed within the theories of the kinematic dynamo (see Parker 1979). Reconnection diffusion, unlike ``magnetic turbulent diffusivity'', deals with dynamically important magnetic fields, e.g. with subAlfv\'{e}nic turbulence\footnote{The concept of ``magnetic turbulent diffusivity'' assumes that magnetic fields can be passively mixed up on all scales, up to the Ohmic diffusion scale. At the latter scale Ohmic dissipation is trivial and the diffusion is hydrodynamic at all scales.}. Thus in the process of reconnection diffusion, magnetic fields are {\it not} passively mixed and magnetic reconnection plays a vital role for the process. One may show that the domain of applicability of the ``magnetic turbulent diffusivity'' concept is extremely limited. Numerical calculations (Cho et al. 2010, Beresnyak 2012) testify that around 5\% of turbulent energy flow is transferred to magnetic energy and that this value does not depend on the initial magnetization of the fluid. Therefore, even if the turbulent flow initially is not magnetized, it is expected to develop dynamically important magnetization on the time scale of 10-20 large eddy turnover times. As a consequence, we expect that star formation will happen in magnetized media even for the first generation of stars.\footnote{Incidentally, the same logic suggests that the proposals of bringing magnetic field and kinetic energy to a state of equipartition in molecular clouds during their formation (see Sur et al. 2012) may not work. Indeed, molecular clouds are not likely to survive 20 crossing times.} \subsection{Reconnection diffusion and numerical effects} One might argue that reconnection diffusion is automatically a part and parcel of 3D numerical simulations of star formation. As we argue below this is only partially true and a proper understanding of reconnection diffusion is required for the interpretation the numerical simulations. The LV99 model predicts that the reconnection rates in turbulent fluids are independent of the local physics, but are determined by the turbulent motions. This is great news for numerical simulations: parasitic numerical effects that induce poorly controlled small scale diffusivity of magnetic field lines are not important on the scales of turbulent motions. In other words, the low resolution numerics may provide an adequate representation of high Lundquist number astrophysical turbulence as far as the reconnection diffusion is concerned. On the contrary, if the structures studied in numerical simulations (e.g. cores, filaments, shells) lose turbulence due to numerical diffusivity effects, we predict that the diffusion of magnetic flux in those simulated structures differs significantly from the diffusion in the actual interstellar structures where turbulence persists. Note, that naive convergence studies based on increasing the numerical resolution several times cannot notice the problem unless the resolution increased to the degree that the aforementioned structures become turbulent. In terms of numerical studies of star formation it is encouraging that it is the largest scales of turbulent motions that are important for the reconnection diffusion. At the same time, in many cases in numerical simulations we expect to observe magnetic field evolution influenced by strongly damped turbulent motions, which do not demonstrate power law cascade. The differences of reconnection diffusion in such damped regimes and in the astrophysically important regime of well developed turbulence with a power law spectrum should be a subject of further studies. At the same time, the importance of understanding reconnection diffusion is surely not limited by the proper interpretation of numerical simulations of star formation. Reconnection diffusion provides a theoretical way to predict the dynamics of magnetic flux in turbulent fluids. \subsection{Magnetic star formation in the presence of reconnection diffusion} In spite of the fact that reconnection diffusion relaxes the conventional flux freezing condition, the role of magnetic fields for star formation should not be disregarded. For instance, magnetic fields solve the angular momentum problem in the process of star formation (Mestel 1965). At the same time, reconnection diffusion allows to prevent the ``magnetic breaking catastrophe'' that, contrary to observations, prevents the formation of circumstellar disks (Lazarian \& Vishniac 2009, Santos-Lima et al. 2012). \subsection{Comparison with earlier papers} The concept of removal of magnetic fields from clouds via the process of LV99 reconnection was first introduced in Lazarian (2005) (see also discussions in Lazarian \& Vishniac 2009). The numerical calculations of reconnection diffusion in application to molecular clouds and accretion disks have been performed in Santos-Lima et al. (2010, 2012). However, the physical picture of how magnetic fields and matter can decouple in turbulent magnetized plasma was not presented there. In comparison, this paper focuses on the physics of reconnection diffusion and presents a simple analytical model illustrating how the tendencies in magnetic field distribution revealed in the observations by Crutcher et al. (2009, 2010a) can be accounted for, as well as other observations reviewed by Crutcher (2012). In addition, the current paper presents the additional evidence of the importance of reconnection diffusion in various astrophysical environments. While alternative explanations may be possible for the cases discussed (e.g, Li et al. 2011) we feel that the consistency of those with the reconnection diffusion concept provides an additional evidence in favor of the reconnection diffusion scenario. This calls for the necessity to include the concept of reconnection diffusion within the star formation paradigm.
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We study a combined sample of 264 star-forming, 51 composite, and 73 active galaxies using optical spectra from SDSS and mid-infrared (mid-IR) spectra from the Spitzer Infrared Spectrograph. We examine optical and mid-IR spectroscopic diagnostics that probe the amount of star formation and relative energetic contributions from star formation and an active galactic nucleus (AGN). Overall we find good agreement between optical and mid-IR diagnostics. Misclassifications of galaxies based on the SDSS spectra are rare despite the presence of dust obscuration. The luminosity of the [NeII] 12.8 $\mu m$ emission-line is well correlated with the star formation rate (SFR) measured from the SDSS spectra, and this holds for the star forming, composite, and AGN-dominated systems. AGN show a clear excess of [NeIII] 15.6 $\mu m$ emission relative to star forming and composite systems. We find good qualitative agreement between various parameters that probe the relative contributions of the AGN and star formation, including: the mid-IR spectral slope, the ratio of the [NeV] 14.3 $\mu m$ to [NeII] $\mu m$ 12.8 fluxes, the equivalent widths of the 7.7, 11.3, and 17 $\mu m$ PAH features, and the optical ``D" parameter which measures the distance a source lies from the locus of star forming galaxies in the optical BPT emission-line diagnostic diagram. We also consider the behavior of the three individual PAH features by examining how their flux ratios depend upon the degree of AGN-dominance. We find that the PAH 11.3 $\mu m$ feature is significantly suppressed in the most AGN-dominated systems.
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It has been established that a strong link exists between supermassive black holes (SMBHs) and the galaxies in which they live \citep{FM, Gebhardt, m-sigma} and between accreting SMBHs (or active galactic nuclei, AGN) and star formation in particular \citep[e.g.][]{Kauffmann, Cid}. Clues to this link can be revealed by detailed analysis of the spectra of galaxies and AGN. Star formation processes imprint their signatures on the spectrum. Stellar populations of various ages produce strong continuum emission in the ultraviolet (UV) through near-IR region, while gas photo-ionized by the hot young stars produces prominent emission-lines in the UV, optical, and IR. The accretion disk surrounding the supermassive black hole in an AGN produces its own strong UV and optical continuum emission (the ``big blue bump \citep[]{Shields}). This emission can swamp that of the stellar population of the host galaxy when there is a direct view of the accretion disk (i.e. in Type 1 AGN). Dust in the interstellar medium of the host galaxy will absorb optical and UV light from both stars and the AGN and re-radiate this in the IR. In AGN, this reprocessing also occurs in the dusty obscuring torus \citep{Antonucci, Urry}. In type 2 AGN, the direct view of the AGN is blocked by the obscuring torus, and the presence of the AGN can be inferred by prominent optical and infrared emission lines from highly ionized gas located beyond the torus and the strong mid-IR continuum emission from the torus. In such objects the signatures of host galaxys young stars and star formation are also present in optical and mid-IR spectra. These type 2 AGN are therefore ideal laboratories for studying the connection between AGN and star formation, provided that the tracers of this activity can be disentangled. Multi-wavelength star formation indicators have been extensively studied in samples of quiescent star forming galaxies \citep[see][for a review]{Ken98}. Star formation rates (SFRs) have been calibrated in the optical based on H$\alpha$ emission, which results from recombination following ionization due to energetic photons from massive O and B stars, and in the UV based on the continuum which provides a window into emission from young stars. Dust in the galaxy reprocesses optical and UV photons and re-radiates in the infrared (IR), thus revealing obscured star formation. Other IR indicators include emission from fine structure lines excited by starburst activity, such as [NeII] 12.81$\mu$m and [NeIII] 15.56$\mu$m \citep{Ho}, as well as spectral emission features associated with polycyclic aromatic hydrocarbons (PAHs), since these macro-molecules/small-grains can be ionized by single optical or UV photons \citep{Li} and are prominent in IR spectra of star-forming galaxies \citep{PAHFIT}. Such host galaxy star formation diagnostics can be contaminated by the presence of an AGN. Photons emitted by the accretion disk enhance the optical and UV emission used to derive SFRs. The obscuring medium enshrouding the central engine contributes significantly to the infrared emission, constituting approximately 20\% of the AGN bolometric luminosity \citep{12m}. Disentangling the contribution from the AGN vs. star formation then becomes a necessary intermediate step in using continuum IR emission as a tracer of star formation in active galaxies, and such a process can be plagued with uncertainties. Previous studies have indicated that emission from PAHs tend to be suppressed in galaxies hosting an active nucleus, which is sometimes attributed to the harsh radiation field produced by the AGN accretion disk destroying PAHs \citep[e.g.][]{Voit, O'Dowd, Treyer, me, DS10}. A complementary view of the AGN and starburst connection is to study spectroscopic signatures that parameterize the relative importance of these two processes. The optical BPT diagram (a plot of [OIII]$\lambda$5007/H$\beta$ vs. [NII]$\lambda$6583/H$\alpha$) provides a useful diagnostic to differentiate between star-forming galaxies, composite systems (galaxies with comparable amounts of star-formation and AGN activity) and Type 2, or obscured, AGN in the local universe \citep{BPT}. The optical ``D'' parameter is the distance a source lies from the locus of star forming galaxies on the BPT diagram: a higher D value indicates greater AGN dominance \citep{Kauffmann}. Ratios of IR fine structure lines also parameterize the ionization field hardness. For instance, [OIV]26$\mu$m and [NeV]14.32$\mu$m are primarily ionized by AGN \citep{Rigby, DS09, Armus, Gould10} whereas [NeII]12.81$\mu$m is excited by star formation \citep{Ho}. The ratio of these lines can then indicate the relative importance of these two processes \citep{Genzel, me, PS}. The mid-infrared (MIR) spectral slope, $\alpha_{20-30\mu m}$, steepens as the amount of emission from cold dust heated by stars increases relative to the amount of emission from hotter dust heated by the AGN. It is thus another potentially useful tool to assess the relative amount of star formation in active systems \citep[e.g.][]{Buchanan, Deo}. The equivalent width (EW) of PAH grains is another diagnostic to probe the relative amount of star formation to AGN activity. An anti-correlation between ionization field hardness and the ratio of PAH features at 7.7$\mu$m and 11.3$\mu$m has been reported and is interpreted as due to an increasing contribution to the mid-IR continuum emission by AGN-heated dust. The selective destruction of smaller PAH grains may also play a role \citep[e.g.][]{PAHFIT, O'Dowd, Treyer, Wu10}, and so the flux ratios of the different PAH features can then also potentially trace the relative contribution of AGN and star formation in active galaxies. In this paper, we combine samples of star-forming galaxies, obscured AGN (Seyfert 2 galaxies, or Sy2s) and composite systems to study the interplay between AGN activity and star formation. Using quiescent star forming galaxies as a base-line, we investigate the effects of AGN activity on the following star formation diagnostics: optically derived SFRs from the Sloan Digital Sky Survey \citep{Brinchmann}, the luminosities of the IR fine structure lines [NeII] 12.81$\mu$m and [NeIII] 15.56$\mu$m \citep{Ho} and the luminosities of the polycyclic aromatic hydrocarbons (PAHs) at 7.7$\mu$m, 11.3$\mu$m and 17$\mu$m \citep{PAHFIT, Farrah, DS10}. Which of these star formation proxies agree the best among star-forming galaxies and AGN and are thus least affected by the presence of an AGN? We expand upon the work of \citet{me} to test the agreement among diagnostics that parameterize the relative contributions of AGN activity and star formation, including the equivalent width (EW) of PAHs at 7.7, 11.3 and 17$\mu$m, ratios of PAH fluxes \citep{PAHFIT, O'Dowd, Treyer, DS10}, ratios of IR fine structure lines \citep{Genzel, Armus, PS}, mid-IR spectral slope \citep[$\alpha_{20-30\mu m}$][]{Buchanan, Deo} and the optical D parameter \citep{Kauffmann}. In \citet{me} we found that Sy2s with stronger PAH emission tend to have a softer ionization field and that the MIR spectral slope was well correlated with PAH EW. By expanding this parameter space into the regime of quiescent star forming galaxies, we test if PAH features and the mid-IR spectral slope are dependent on the hardness of the radiation field, regardless of its source, or if ionization by an AGN is necessary to appreciably affected observed features. Finally, we use the results of this analysis to empirically decompose the mid-IR (MIR) emission into a star-forming and AGN component. \renewcommand{\thefootnote}{\arabic{footnote}}
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We have explored the infrared and optical parameter space where star-forming galaxies, composites and AGN live to analyze diagnostics that parameterize host galaxy star formation. Using star-forming galaxies as a control sample, we have investigated which star formation diagnostics are least affected by AGN activity. We have studied diagnostics that trace the interplay between AGN and starburst activity and tested whether a smooth transition exists over a range of radiation field hardness. Finally, using the results of this analysis, we present an empirical decomposition of the MIR flux into a star-forming and an AGN component. {\it Our overall result is that the optical and mid-IR diagnostics of star formation and of the relative importance of young stars and the AGN generally agree well.} Our more specific results are summarized as follows: \begin{description} \item{\bf{SFR diagnostics}} The SDSS derived SFR$_{fiber}$ and the [NeII] 12.81$\mu$m luminosity agree well, and are the most reliable SF proxies we have considered for use in star-forming galaxies, composites and AGN. The sum of [NeII] and [NeIII] 15.56$\mu$m is systematically offset to higher values for AGN due to the active nucleus boosting the [NeIII] flux. The aggregate PAH luminosity (L$_{\Sigma PAH 7.7,11.3,17\mu m}$) relative to L$_{[NeII]}$ is suppressed in AGN compared with both quiescent and composite galaxies. Though the (L$_{\Sigma PAH 7.7,11.3,17\mu m}$)/SFR$_{fiber}$ ratio agrees between star-forming galaxies and Sy2s, analysis of PAH flux ratios indicate abnormal behavior in some AGN dominated systems, such as suppression of the 11.3$\mu$m feature, suggesting diminished PAH emission in strong AGN rather than [NeII] enhancement. The disparate slopes of (L$_{\Sigma PAH 7.7,11.3,17\mu m}$)/SFR$_{fiber}$ between Sy2s and star-forming galaxies also hint that AGN may contaminate the emission from PAH macro-molecules. \item{\bf{Ionization Field Hardness}} We have used the optical D parameter, which indicates the distance a source lies from the star-forming galaxy locus on the BPT diagram \citep{Kauffmann}, and L$_{[NeV]}$/L$_{[NeII]}$ \citep{Armus, PS} as probes of the incident radiation field hardness. The infrared diagnostics that parameterize the relative contributions of AGN to star-forming activity that we have considered ($\alpha_{20-30\mu m}$, PAH EWs and PAH flux ratios) only show significant trends with ionization field hardness in the AGN population. This result suggests that rather than parameterizing incident radiation hardness in general, these IR diagnostics are only effective descriptors of ionization field hardness in AGN. \begin{description} \item{\boldmath{$\alpha_{20-30\mu m}$}} In AGN, the spectral slope between 20-30 $\mu$m (measured by the spectral index $\alpha_{20-30\mu m}$) has been shown to steepen in the presence of cold dust due to star formation \citep{Buchanan, Deo}. For the AGN in our sample, $\alpha_{20-30\mu m}$ is significantly anti-correlated with proxies of ionization field hardness, albeit with wide scatter, however this result does not hold when expanded into the regime of quiescent star forming galaxies. \item{\bf{PAH EW}} PAH EWs are significantly anti-correlated with the increasing hardness of radiation field, but this is only true when considering the AGN. When considering just the star-forming galaxies, no trend exists between incident radiation hardness and the strength of the resulting PAH feature, which agrees with the findings of \citet{O'Dowd} and \citet{Treyer}. This result is consistent with the hypothesis that the mid-IR continuum has a significant contribution from dust heated by the AGN only in objects where the ionizing radiation field is dominated by the AGN. For the 11.3$\mu$m PAH feature, our analysis of the PAH flux ratios suggests that destruction of the PAH macro-molecule in strong AGN could also contribute to suppression of the PAH 11.3$\mu$m EW. \item{\bf{PAH flux ratios}} We find that the luminosity of PAH 11.3 $\mu m$ feature is signifcantly suppressed relative to the luminosities of the PAH 7.7 $\mu m$ and 17 $\mu m$ features and relative to the optical and mid-IR derived star formation rates in the most AGN-dominated systems. Thus, some care is required in using the PAH 11.3$\mu$m luminosity as a proxy for the SFR in AGN-dominated systems. \end{description} \item{{\bf Composite Systems}} Galaxies that are optically classified as composite systems are more akin to quiescent star forming galaxies in terms of the mid-IR parameters that trace star formation: the smaller relative energetic importance of the AGN in composites does not seem to affect [NeIII] and PAH luminosities. \item{\bf{[NeV] as unambiguous AGN signature}} Due to the high ionization potential of [NeV] (97.1 eV), it is potentially an unequivocal signature of AGN activity \citep{Gilli, Gould10}. One galaxy that is optically classied as a quiescent star-forming system in our sample has a [NeV] detection, yet other IR parameters (i.e. PAH EW values) are normal for a pure star forming galaxy. The nature of this galaxy as a strong ``hidden'' AGN is not clear. \item{\bf{Disentangling MIR Emission}} Using multiple linear regression, we fit the relation L$_{MIR}$ = $\alpha$L$_{SFR}$ + $\beta$L$_{AGN}$ to our full sample to determine $\alpha$ and $\beta$. We set L$_{SFR}$ = L$_{[NeII]}$ and L$_{AGN}$=L$_{[OIII],AGN}$ where in L$_{[OIII],AGN}$, we have subtracted out the estimated starburst contribution to the [OIII] flux using Eq. 3 in \citet{WHC}. We find $\alpha$=89$\pm$1 and $\beta$=111$\pm$7. This decomposition can be useful in estimating the MIR emission due to the circumnuclear AGN obscuration or conversely host galaxy star formation in AGN. \end{description} \clearpage
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1206.5830
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1206
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1206.2417_arXiv.txt
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{ The origin of the heliospheric magnetic flux on the Sun, and hence the origin of the solar wind, is a topic of hot debate. While the prevailing view is that the solar wind originates from outside coronal streamer helmets, there also exists the suggestion that the open magnetic field spans a far wider region. Without the definitive measurement of the coronal magnetic field, it is difficult to resolve the conflict between the two scenarios without doubt. We present two 2-dimensional, Alfv\'enic-turbulence-based models of the solar corona and solar wind, one with and the other without a closed magnetic field region in the inner corona. The purpose of the latter model is to test whether it is possible to realize a picture suggested by polarimetric measurements of the corona using the Fe XIII 10747\AA\ line, where open magnetic field lines seem to penetrate the streamer base. The boundary conditions at the coronal base are able to account for important observational constraints, especially those on the magnetic flux distribution. Interestingly, the two models provide similar polarized brightness (pB) distributions in the field of view (FOV) of SOHO/LASCO C2 and C3 coronagraphs. In particular, a dome-shaped feature is present in the C2 FOV even for the model without any closed magnetic field. Moreover, both models fit equally well the Ulysses data scaled to 1~AU. We suggest that: 1) The pB observations cannot be safely taken as a proxy for the magnetic field topology, as often implicitly assumed. 2) The Ulysses measurements, especially the one showing a nearly uniform distribution with heliocentric latitude of the radial magnetic field, do not rule out the ubiquity of open magnetic fields on the Sun.
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\label{sec_intro} Identifying the source regions of the heliospheric magnetic flux, and hence those of the solar wind, is a long standing issue in solar physics~\citep[see e.g.,][]{Schwenn_06,YMWang_09}. The difficulties associated with this identification are due mainly to the difficulty of directly measuring the solar coronal magnetic field, which is essential in most of the schemes mapping the in situ solar winds to their coronal sources~\citep[][to name but a few]{Levine_etal_77, WangSheeley_90, Neugebauer_etal_98, Neugebauer_etal_02, WangSheeley_06, Abbo_etal_10}. Although advances on polarimetric measurements with coronal emission lines have been recently made~\citep{Habbal_etal_01, Habbal_etal_03, Lin_etal_04, YLiu_09}, the coronal magnetic field remains largely unknown~\citep[e.g.,][]{Cargill_09}. Without detailed, quantitative measurements, the coronal magnetic field has been commonly constructed via numerical extrapolation. While all available schemes use the photospheric magnetic field as boundary input, they differ substantially in how to treat the effects of electric currents on the global coronal magnetic field in a volume bounded by the photosphere and some outer boundary. The electric currents may be neglected altogether~\citep[e.g., the potential field source surface model by][]{Schatten_etal_69, AltschulerNewkirk_69}, they may be assumed to be purely horizontal (e.g., the current sheet source surface model by~\citeauthor{ZhaoHoeksema_95}~\citeyear{ZhaoHoeksema_95}), or flow exclusively along magnetic field lines (e.g., the force-free model by~\citeauthor{Tadesse_etal_09}~\citeyear{Tadesse_etal_09}), or both the volumetric and sheet currents are self-consistently computed as a product of the plasma properties (the magnetohydrodynamic (MHD) models by e.g.,~\citeauthor{Lionello_etal_09}~\citeyear{Lionello_etal_09}). Polarized brightness (pB) images of the solar corona, routinely obtained with space-borne and ground-based coronagraphs, often guide the extrapolation schemes such that the resultant magnetic field configuration matches available pB images. Implied here is that the density structures, as manifested in pB images, reflect the magnetic topology in the inner corona. And usually dome-shaped streamer helmets, the most prominent feature in pB images, are seen as comprising closed magnetic fields. It follows that the bulk of the solar wind originates from open field regions outside streamer helmets, even though the precise fraction by which coronal holes and the quiet Sun contribute to the solar wind is debatable~\citep[][and references therein]{Kopp_94, Hu_etal_03}. However, this scenario is not universally accepted. A distinct picture has been advocated in which the solar wind flows along the ubiquitous open magnetic field lines that are not limited to coronal holes or the quiet Sun but come from throughout the Sun~\citep{WooHabbal_99, WooHabbal_03, Woo_etal_04}~\citep[see also][]{WooDruck_08}. Interestingly, the arguments raised to support this picture initially also came from density measurements. By combining the near-Sun pB values with radio occultation measurements as well as in situ solar wind data, \citet{WooHabbal_99} showed that the signatures of coronal sources are preserved in the measured solar wind away from the Sun, contending that these density imprints are almost radially propagating. A further support for this scenario comes from the fact that the white light images of the Sun at total eclipses, when properly processed, exhibit a rich set of filamentary structures that extend almost radially from the solar surface~\citep[see Fig.1 in ][]{WooHabbal_03}. Supposing these fine structures trace the magnetic field lines, this would suggest that some coronal magnetic field lines penetrate the dome-shaped streamer base. More importantly, these apparently open magnetic fields were also seen in the polarimetric measurements of the inner corona using Fe XIII 10747\AA\ line~\citep[see figures in][]{Habbal_etal_01, Habbal_etal_03}. Carefully addressing observational complications such as collisional depolarization and the Van Fleck effect, \citet{Habbal_etal_01} argued that the largely radially aligned polarization vectors may indeed reflect a coronal magnetic field that is predominantly radial. Given the importance of addressing the origin of the open magnetic flux of the heliosphere, it is surprising to see that while the traditional scenario has been extensively incorporated into numerical studies originated by~\citet{PneumanKopp_71}~\citep[also see][and references therein]{Lionello_etal_09}, the scenario proposed by~\citet{WooHabbal_03} has not been modeled quantitatively. Here we wish to implement this scenario in a numerical model, thereby testing it against two fundamental constraints that the traditional scenario can readily satisfy: one is the appearance of a dome-shaped bright feature in pB images, and the other is the fact established by Ulysses measurements that the radial magnetic field strength $B_r$ is nearly uniform with latitude beyond 1~AU~\citep{SmithBalogh_95, Smith_etal_01}. Note that the latter fact was used to refute the suggestion of a largely radially expanding solar wind~\citep{Smith_etal_01}, as the magnetic flux near the Sun is obviously nonuniform~\citep[e.g.,][]{Svalgaard_78, Vasquez_etal_03}. Before proceeding, we note that from SOHO/EIT images, it is obvious that there are a myriad of low-lying loop-like structures in the corona. As suggested by~\citet{Habbal_etal_01}, a large-scale closed magnetic field (the ``nonradial'' component in their paper) may also help shape the large-scale corona. To simplify our treatment, we will simply try to answer one question: can a dome-shaped bright feature show up in a magnetic configuration where there is no closed field at all? In essence, this is equivalent to saying that we are interested in the region somewhere above the layer below which loops abound in SOHO/EIT images, and beyond which the contribution from closed magnetic fields is assumed to be negligible. In what follows, we will present two numerical models that differ in the magnetic field configuration in the inner corona, one with and the other without a closed field region. Both models are able to incorporate the essential observational constraints near the coronal base, especially the latitudinal dependence of the radial magnetic field. We then construct pB maps to see whether they display features similar to what is seen in white light observations. Moreover, model results are also compared with several crucial parameters observed in situ. The models are described in section~\ref{sec_nummodel}, results from the numerical computations are given in section~\ref{sec_numresults}, and section~\ref{sec_conc} concludes this paper.
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\label{sec_conc} Identifying the origin of the Sun's open magnetic flux is of crucial importance in establishing the connection of the in situ solar winds to their sources. In the absence of a definitive measurement of the solar coronal magnetic field, this identification problem is subject to considerable debate~\citep[e.g.,][]{Schwenn_06, YMWang_09}. This is true even at solar minima when the configuration of the solar corona is relatively simple, the most prominent feature being the bright streamer helmets in white light images. While the prevailing view is that the majority of the solar wind originates from outside streamer helmets~\citep{PneumanKopp_71}, there also exists the suggestion that the open magnetic field is ubiquitous on the Sun and not confined to coronal holes or the quiet Sun~\citep{WooHabbal_99, WooHabbal_03}. Implementing the former scenario in a numerical model has been a common practice~\citep[e.g.][]{Lionello_etal_09}, however, so far the latter has not been modeled quantitatively and hence tested quantitatively. Here we offer such an implementation. We have constructed two 2-dimensional, Alfv\'enic-turbulence-based models of the solar corona and solar wind, one with and the other without a closed magnetic field region in the inner corona. The purpose of the latter model is to minimize the contribution of the closed magnetic field, trying to mimic a corona permeated with open magnetic fields which may infiltrate the dome-shaped streamer base. In specifying the boundary conditions at the coronal base, we have taken into account some important observational constraints, especially those on the magnetic flux distributions. Interestingly, the two models provide similar polarized brightness (pB) maps in the field of view (FOV) of the SOHO/LASCO C2 and C3 coronagraphs. In particular, a dome-shaped feature is present in the C2 FOV even for the model without any closed magnetic field. Moreover, both models fit equally well the Ulysses data scaled to 1~AU. Hence we suggest that: 1) The pB observations cannot be safely taken as a proxy for the magnetic field topology, as often implicitly assumed. 2) The Ulysses measurements, especially the one indicating that the radial magnetic field strength is nearly uniformly distributed with heliocentric latitude, do not rule out the ubiquity of open magnetic fields on the Sun. We do not intend to resolve the conflict of the two distinct scenarios currently available for the origin of the heliospheric magnetic flux. Rather, the presented numerical results suggest the likelihood that the magnetic field structure of bright features (e.g., helmet streamers) in the corona may be more diverse than traditionally viewed: the magnetic flux therein can be either closed or open. To differentiate the scenarios, it is likely that more stringent constraints come from the SOHO/UVCS measurements. For instance, measurements based on the Doppler dimming technique have yielded that along the direction transverse to the streamer helmet, there exists a transition in the inferred plasma speed from unmeasurable to significant values, and this transition seems to trace the streamer legs~\citep{Strachan_etal_02, Frazin_etal_03}, identified by the enhancement of the intensity ratio of O VI $\lambda$1032\AA\ to H I Ly$\alpha$~\citep{Kohl_etal_97}. Therefore, it remains to be seen whether a model permeated with open magnetic fluxes can account for this feature. To do this, an obvious need is to incorporate O$^{5+}$ ions in a three-fluid model and test both scenarios. At the moment, in such models only the traditional partially open scenario has been adopted~\citep[e.g.,][]{Li_etal_06, Ofman_etal_11}. Instead of implementing a further construction, let us end here by noting that one may also ask whether these observational features~\citep{Kohl_etal_97, Strachan_etal_02, Frazin_etal_03} are universal for all streamers.
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1206.2417
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1206
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1206.0909_arXiv.txt
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The expected imaging capabilities of future Extremely Large Telescopes (ELTs) will offer the unique possibility to investigate the stellar population of distant galaxies from the photometry of the stars in very crowded fields. Using simulated images and photometric analysis we explore here two representative science cases aimed at recovering the characteristics of the stellar populations in the inner regions of distant galaxies. Specifically: case~A) at the center of the disk of a giant spiral in the Centaurus Group, ($\mu_{\rm B}\sim$ 21, distance of 4.6 Mpc); and, case~B) at half of the effective radius of a giant elliptical in the Virgo Cluster ($\mu_{\rm I} \sim$ 19.5, distance of 18 Mpc). We generate synthetic frames by distributing model stellar populations and adopting a representative instrumental set up, i.e. a 42 m Telescope operating close to the diffraction limit. The effect of crowding is discussed in detail showing how stars are measured preferentially brighter than they are as the confusion limit is approached. We find that (i) accurate photometry ($\sigma \sim 0.1$, completeness $\gsim 90$\%) can be obtained for case~B) down to $I\sim 28.5$, $J\sim 27.5$ allowing us to recover the stellar metallicity distribution in the inner regions of ellipticals in Virgo to within $\sim 0.1$~dex; (ii) the same photometric accuracy holds for the science case~A) down to $J\sim 28.0$, $K\sim 27.0$, enabling to reconstruct of the star formation history up to the Hubble time via simple star counts in diagnostic boxes. For this latter case we discuss the possibility of deriving more detailed information on the star formation history from the analysis of their Horizontal Branch stars. We show that the combined features of high sensitivity and angular resolution of ELTs may open a new era for our knowledge of the stellar content of galaxies of different morphological type up to the distance of the Virgo cluster.
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One of the key issues in modern astronomy concerns the Star Formation History (SFH) in the Universe. Direct observations of galaxies up to high redshift can be used to map the SFH, but since the integrated galaxy light is dominated by the contribution of the most recent generations of stars, the information that can be derived on the underlying older stellar population is severely limited. However, the SFH in galaxies can be uncovered by interpreting the Colour-Magnitude Diagrams (CMD) of their stars, which contain the fossil record of their SFH (e.g., \citealt{greggio}, \citealt{holtz}, \citealt{tom}, \citealt{cole}, \citealt{mcq}, \citealt{weisz}; see also \citealt{cigno} for a review). This kind of studies require accurate photometry down to faint magnitudes in crowded fields, which, with current instrumentation, is feasible only for the nearest galaxies. This implies a very limited sampling of the SFH in the Universe, with plenty of dwarfs, a few spirals and no giant elliptical. By the end of this decade this situation is expected to change significantly, as a number of extremely large telescopes, as the Giant Magellan Telescope (GMT; \citealt{johns}), the Thirty Meter Telescope (TMT; \citealt{szeto}) and the European Extremely Large Telescope (E-ELT; \citealt{gilmozzi}) could come into operation. The large collecting area of ELTs coupled with adaptive optics cameras able to deliver quasi-diffraction limited images will allow us to probe a wide volume, where we can access a significative sample of galaxies. More importantly, it will be possible to study dense stellar fields. Indeed, because of crowding, stellar photometry in external galaxies is currently feasible only in regions of relatively low surface brightness. For giant galaxies, this prevents us from deriving detailed SFH where most of the galaxy mass is. One of the key advantages of taking images with extremely large aperture telescopes is the exceptionally good image quality when the telescope can work close to its diffraction limit. This significant improvement has two fundamental effects : 1) it produces a dramatic reduction of the background light over the point spread function (PSF) area considered for the photometry; 2) it allows a significant improvement in the spatial resolution (proportional to the telescope aperture). These combined advantages offer the unique opportunity to carry out observations of faint targets in crowded and/or structured objects (galaxies) at large distance, that cannot be exploited by any other telescope of smaller aperture, neither ground nor space based. In summary, the excellent resolution capabilities of next generation large aperture Telescopes will enable us to study high surface brightness regions, i.e. the inner parts of galaxies, where star formation was more conspicuous, as well as to address directly age and metallicity gradients. The photometric performance of ELTs operating close to the diffraction limit was investigated by \cite{olsen} and by \cite{atul}. Both studies focus on the impact of crowding conditions on the photometric quality, as quantified by the 1 $\sigma$ width of the error distribution as a function of magnitude, for stellar fields with different surface brightness. In this approach errors are considered to be symmetric, with stellar luminosities having the same probability of being overestimated or underestimated. However, crowding induces asymmetrical photometric errors, with an excess of stars measured brighter than they are (e.g. \citealt{tosi}; \citealt{carma96}; \citealt{alviopix}). This results into artificially brightened features on the CMD which may induce a systematic error in those parameters (e.g. distances and stellar ages) which are derived from their luminosity. In addition, crowding may also affect systematically the distribution of stars across the CMD, jeopardizing its interpretation in terms of star formation history. The impact of this asymmetry is negligible under low-crowding conditions (e.g. when the 1$\sigma$ width is smaller than about 0.1 mag), but grows rapidly with crowding. In this paper we present the results of end-to-end simulations of two specific science cases aimed at investigating the SFH in galaxies, which fully exploit the unprecedented capabilities foreseen for ELTs. \begin{figure}[t] \plotone{f1.eps} \caption{Expected limiting magnitudes for point sources observed with MICADO@E-ELT (solid lines) compared to the magnitude of RGB Tip, red HB and 10 Gyr old MS TO stars at a distance of 4.6 Mpc (dot-dashed lines). The limiting magnitudes, relative to a S/N of 20 (upper curve) and of 5 (lower curve), are computed for a total of 5 hours integration time.} \label{limits} \end{figure} There are several interesting programs concerning the study of resolved stellar populations in external galaxies, like the analysis of multiple generations in Globular Clusters (see \citealt{piotto} and references therein); the determination of the initial mass function and its variations (e.g. \citealt{kroupa}); the recovery of the spatially resolved SFH in isolated and in interacting galaxies (e.g. \citealt{cioni} for the Magellanic system); tracing the galaxy assembly process through the counts of resolved stars and detection of streams (e.g. \citep{annette}). Here we focus on the possibility of deriving global information on the star formation history for a significative sample of galaxies with a modest time investment. This means targeting distant galaxies, so that the field of view samples a fair fraction of their mass, and observing galaxies in groups or clusters, so as to enable the study of a representative galaxy population. On these premises we have concentrated on the following science cases: A) deriving basic information on the SFH in the central region of a disk galaxy at the distance of the Centaurus group; and B) studying the metallicity distribution at half of the effective radius of an elliptical galaxy in the Virgo Cluster. For these cases we have developed simulations of stellar photometry, based on synthetic stellar populations and assuming the expected performances of the Adaptive Optics assisted near-IR camera (MICADO, \citealt{daviesmess}) for the E-ELT (see details in Sect. 3.1). The resulting CMDs were the analyzed to assess the impact of the photometric errors on the specific science goal. In section 2 we describe our selected science cases. Sect. 3 report detail of our simulation and the following photometric analysis. Results of this study are given in Sect. 4 with general conclusion summarized in Sect. 5. \begin{figure}[t] \plotone{f2.eps} \caption{Synthetic CMD obtained with a constant rate of star formation over the last 12~Gyr, a Salpeter initial mass function (IMF), and a simple age-metallicity relation which, starting from $Z=0$, goes through the solar metallicity ($Z_\odot$) at 4.5 Gyr ago, and reaches a value of $Z=1.1 Z_\odot$ at the current epoch. The simulation, computed with the YZVAR code by G.P. Bertelli with the 2002 version of the Padova tracks \citep{girardi}, contains 200000 stars brighter than $M_K = -2$, and corresponds to a total mass of formed stars of $2.9\cdot 10^8$ M$_\odot$. The four diagnostic boxes superimposed contain 122 (RSG), 598 (BSG), 340 (AGB) and 22198 (RGB) synthetic stars.} \label{cmd_young} \end{figure}
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In this paper we have explored the feasibility of two specific scientific applications for the study of the SFH in distant galaxies with the expected performance of the 42 m E-ELT equipped with the MICADO camera working close to the diffraction limit of the telescope. We have focussed on giant galaxies located in the Centaurus group and in the Virgo Cluster. These are regions of the nearby Universe where a variety of galaxy types and morphology can be found, so that a significative sampling of the SFH in the Universe can be gathered. In spite of the large collecting area of E-ELT, old MS turn-offs cannot be investigated at these distances; however interesting information on the SFH can be derived from the analysis of the luminous portion of the CMD. The crucial advantage offered by the E-ELT working close to the diffraction limit is related to the exquisite spatial resolution that enables accurate photometry in crowded stellar fields. This will allow us to study the SFH in high surface brightness regions of giant galaxies where most of the stellar mass is located, and to address directly stellar population gradients, from the outskirts down to the inner regions of galaxies. \begin{figure}[t] \plotone{f17.eps} \caption{Distribution of [Fe/H] as derived from the positional coincidence on the $J$ band frame compared to that derived from the nearest neighbor match on the $(J,I-J)$ CMD.} \label{zdist} \end{figure} We have investigated how photometric errors impact on two science cases: A) the study of the SFH at the center of the disk of a spiral galaxy at a distance of 4.6 Mpc ($\mu_{\rm B} \simeq 21$); B) the recovery of the metallicity distribution in the inner regions (at 1/2 of the effective radius) in an giant elliptical galaxy at a distance of 18 Mpc ($\mu_{\rm I} \simeq 19.6$). To this end we have produced synthetic frames in the $I, J$ and $K$ bands by distributing theoretical stellar populations constructed with suitable star formation histories, and adopting the PSF currently foreseen for the MICADO camera assisted by the MAORY adaptive optics module. The synthetic frames were measured with a standard package for crowded fields photometry to yield output catalogues of detected stars. The input and output stellar lists have been compared to determine the photometric quality, and the output CMDs have been analyzed to assess to which extent the scientific aims of the two cases can be realized. For both science cases, and in all photometric bands, blending of stellar sources leads to an asymmetrical error distribution, and a general migration of star counts along the luminosity function towards the brighter bins. Over a wide magnitude range we detect an excess of stars recovered brighter than what they truly are, and the distributions of the (output--input) magnitude difference is more spread out on the negative (with respect to the positive) side. This effect becomes particularly relevant as one approaches the confusion limit, where stars are measured \textit{only} if their luminosity is artificially boosted by blending \citep{book}. At the same time, as the confusion limit is approached, a progressively larger incompleteness affects the luminosity function, partly balancing the excess due to blending. As a consequence, the measured luminosity function does not appear very different from the input one (see Figs \ref{lf_g} and \ref{lf_e}). Nevertheless the asymmetry of the error distribution persists and can lead to misinterpretations of the CMD, e.g. distances from RGB tip (and HB) stars may be underestimated; too young ages may be inferred from Main Sequence turn off stars. Similarly, the age distribution from the core Helium burning/clump stars could result too skewed at the young end, and bright RGB stars overshooting the tip could be mistaken for relatively young AGB stars. Besides the systematic shift to brighter magnitudes of important diagnostic features, the asymmetry of the error distribution affects in a systematic way the star counts in the boxes used to derive the SFH. This effect, which is very important when trying to derive information from stars with a magnitude close to the confusion limit, may also hamper the results in less severe crowding conditions. Here we have checked that for our cases A) and B), the photometric accuracy is sufficient for their scientific aims. Specifically, our results show that in the very central parts of disks of giant spirals located in the Centaurus group (case A): \begin{itemize} \item accurate photometry ($\sigma\sim0.1$) can be obtained down to $J \sim 28.0$ and $K \sim 27.0$. At these magnitudes the luminosity functions are $\gsim$ 90 \% complete; \item the characteristic plumes of the input CMD are very well reproduced in the output CMD down to $J \sim 26.5$, and the star counts in bright diagnostic boxes which target specific age ranges are perfectly reproduced. This implies that the star formation history in the last $\sim$ 1 Gyr can be accurately derived, and that a robust estimate of the mass transformed into stars at ages older than $\sim 1$ Gyr can be obtained from the star counts on the upper RGB; \item the analysis of the clump/HB stars to recover the SFH prior to $\sim 1$ Gyr ago with a better age resolution requires the application of the full synthetic CMD method, which includes a careful modelling of the photometric errors. For the surface brightness considered, the smearing of the distribution of the clump stars is significant, hampering a robust solution. This case will be better performed in less crowded regions of disks. \end{itemize} The simulations of the case B) stellar population shows that for giant ellipticals in the Virgo cluster: \begin{itemize} \item accurate photometry ($\sigma\sim0.1$) can be obtained down to $I \sim 28.5$ and $J \sim 27.5 $ at half of the effective radius ($\mu_{\rm B}=21.6$). At these magnitudes the luminosity functions are $\gsim$ 90 \% complete; \item the observed CMD well reproduces the width of the RGB in the upper 2 magnitudes, and the color distribution of bright RGB stars is only mildly altered by the photometric errors, so that its interpretation in terms of metallicity distribution of the stellar population is only slightly affected. We estimate that the photometric errors introduce an uncertainty of $\simeq 0.1$ dex on [Fe/H] when using the $I-J$ color combination as a tracer for the effective temperature. \end{itemize} Given the good photometric quality in the bright RGB region for science case B), we argue that it will be possible to study the SFH in the central parts of disks in spiral galaxies in Virgo from star counts in bright boxes, as those shown on Fig. \ref{cmd_young}. This will allow us to contrast the old (prior to $\sim 1$ Gyr) to the more recent star formation activity over entire galaxies and across the Hubble sequence for all cluster members. The good photometric quality in the upper RGB will also allow us to derive the metallicity distribution almost all over giant galaxies, thereby tracing metallicity and population gradients. For disks galaxies in the Centaurus group, the analysis of the Red Clump/HB feature will provide a more detailed information on old star formation episodes. The central regions of disks appear too crowded for this purpose, but the results presented here have been obtained with a standard package for the photometry, which was not designed to work with such a structured PSF as the one provided by MAORY. Likely, the photometric accuracy will improve when using reduction packages more suitable for complex PSF that is produced by adaptive optics systems. Further simulations that include a variable PSF shape over the whole field of view and other effects (non homogeneous background, field distortions, etc.) are required in order to fully characterize the capabilities of ELT imaging under different conditions. At the same time, simulations constructed with different choices for the stellar populations and crowding conditions are necessary to adequately explore the scientific return. These issues will be the subject of a forthcoming paper. \appendix
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1206.0909
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1206
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1206.4188_arXiv.txt
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Primordial Black Holes (PBH's) can form in the early Universe from the collapse of large density fluctuations. Tight observational limits on their abundance constrain the amplitude of the primordial fluctuations on very small scales which can not otherwise be constrained, with PBH's only forming from the extremely rare large fluctuations. The number of PBH's formed is therefore sensitive to small changes in the shape of the tail of the fluctuation distribution, which itself depends on the amount of non-Gaussianity present. We study, for the first time, how quadratic and cubic local non-Gaussianity of arbitrary size (parameterised by $\fnl$ and $\gnl$ respectively) affects the PBH abundance and the resulting constraints on the amplitude of the fluctuations on very small scales. Intriguingly we find that even non-linearity parameters of order unity have a significant impact on the PBH abundance. The sign of the non-Gaussianity is particularly important, with the constraint on the allowed fluctuation amplitude tightening by an order of magnitude as $\fnl$ changes from just $-0.5$ to $0.5$. We find that if PBH's are observed in the future, then regardless of the amplitude of the fluctuations, non-negligible negative $\fnl$ would be ruled out. Finally we show that $\gnl$ can have an even larger effect on the number of PBH's formed than $\fnl$.
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Primordial Black Holes (PBH's) play a very special role in cosmology. They have never been detected but this very fact is enough to rule out or at least tightly constrain many cosmological paradigms. Convincing theoretical arguments suggest that during radiation domination they can form from the collapse of large density fluctuations~\cite{ch}. If the density perturbation at horizon entry in a given region exceeds a threshold value, of order unity, then gravity overcomes pressure forces and the region collapses to form a PBH with mass of order the horizon mass. There are tight constraints on the abundance of PBH's formed due to their gravitational effects and the consequences of their evaporation (for recent updates and compilations of the constraints see Refs.~\cite{Josan:2009qn,Carr:2009jm}). These abundance constraints can be used to constrain the primordial power spectrum, and hence models of inflation, on scales far smaller than those probed by cosmological observations (e.g.~Refs.~\cite{Carr:1994ar,Green:1997sz, Peiris:2008be,Josan:2010cj}). These calculations usually assume that the primordial fluctuations are Gaussian. However since PBH's form from the extremely rare, large fluctuations in the tail of the fluctuation distribution non-Gaussianity can potentially significantly affect the number of PBH's formed. Therefore PBH formation probes both the amplitude and the non-Gaussianity of the primordial fluctuations on small scales. Bullock and Primack~\cite{Bullock:1996at} and Ivanov~\cite{Ivanov:1997ia} were the first to study the effects of non-Gaussianity on PBH formation, reaching opposite conclusions on whether non-Gaussianity enhances or suppresses the number of PBH's formed (see also Ref.~\cite{Yokoyama:1998xd}). Refs.~\cite{Hidalgo:2007vk,Saito:2008em} used a non-Gaussian probability distribution function (pdf) derived from an expansion about the Gaussian pdf~\cite{Seery:2006wk,LoVerde:2007ri} to study PBH formation. However, since PBH's form from rare fluctuations in the extreme tails of the probability distribution, expansions which are only valid for typical size fluctuations can not reliably be used to study PBH formation. Ref.~\cite{Lyth:2012yp} studied the constraints from PBH formation on the primordial curvature perturbation for the special cases where it has the form $\zeta = \pm (x^2- \langle x^2 \rangle)$, where $x$ has a Gaussian distribution, see also Ref.~\cite{PinaAvelino:2005rm}. The minus sign is expected from the linear era of the hybrid inflation waterfall (see also Ref.~\cite{Bugaev:2011wy}), while the positive sign might arise if $\zeta$ is generated after inflation by a curvaton-type mechanism. In this paper we go beyond this earlier work and calculate the constraints on the amplitude of the primordial curvature fluctuations, $\zeta$, from black hole formation for both the quadratic and cubic local non-Gaussianity models (parameterised by $\fnl$ and $\gnl$ respectively). In the process we calculate the probability distribution function of the curvature perturbation for these models. Our results are valid for arbitrary values of these non-linearity parameters, and we recover the known limiting results for very small or large non-Gaussianity. In Sec.~\ref{sec-PBH} we review the calculation of the PBH abundance constraints in the standard case of Gaussian fluctuations, before calculating the constraints for quadratic and cubic local non-Gaussianity in Sec.~\ref{sec:quad} and \ref{sec:cube} respectively. We conclude with discussion in Sec.~\ref{sec:conc}.
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\label{sec:conc} PBH formation probes the extreme tail of the probability distribution function of the primordial fluctuations. This is the region of the pdf which is most sensitive to the effects of any non-Gaussianity that may be present. We have, for the first time, calculated joint constraints on the amplitude and non-Gaussianity of the primordial perturbations, for arbitrarily large local non-Gaussianity. We have studied both quadratic and cubic local non-Gaussianity, parameterised by $\fnl$ and $\gnl$ respectively. On the scales associated with the cosmic microwave background and large scale structure, the constraints on primordial non-Gaussianity are approximately $|\fnl|\lesssim10^2$ \cite{Komatsu:2010fb} and $|\gnl|\lesssim10^6$ \cite{Smidt:2010sv}. In contrast we have shown that on much smaller scales non-linearity parameters of order unity can have a significant effect on the number of PBH's formed. This is because the non-linearity parameters have a larger effect on the tails of the fluctuation distribution than on the more moderate fluctuations probed by cosmological observations. We expect most other forms of non-Gaussianity to also have a significant effect on PBH production, since in general non-Gaussianity generates a skewness which affects the tails of the pdf. The signs of the non-linearity parameters are particularly important. If positive they always make the constraints tighter by acting in the same direction as the linear contribution to $\zeta$. A negative quadratic term tends to cancel the effect of the linear term, thereby reducing the abundance of large PBH forming fluctuations. The constraints on the amplitude of the power spectrum therefore become much weaker, of order unity for $\fnl \lesssim -0.5$. In practice this means that the amplitude of fluctuations will either be too small to form any PBH's at all, or so large that almost every horizon region collapses to form a PBH, which is already observationally ruled out. We hence conclude that a future detection of PBH's would rule out a negative value of $\fnl$ unless its value is tiny, $|\fnl|\ll1$. The case of negative $\gnl$ is different. For $\gnl\simeq-1$ the constraints are weakened as in the negative $\fnl$ case. However as $\gnl$ becomes more negative the constraints quickly become tighter again. In the limit of very large $\gnl$ the constraints are independent of its sign and very tight, approximately the cube of the constraints in the Gaussian case. We have also studied and plotted the probability distribution functions, showing that although the pdfs can diverge, all physical quantities, such as the PBH abundance, remain finite. The PBH constraints have previously been calculated for a pure $\chi^2$ pdf~\cite{PinaAvelino:2005rm,Lyth:2012yp} and for quadratic non-Gaussianity in the limits that $|\fnl|\ll1$ and the linear term dominates~\cite{Seery:2006wk,Hidalgo:2007vk}. We have shown that we recover these limiting cases, however, of particular significance is the fact that our calculations are valid for arbitrarily large quadratic and cubic local non-Gaussianity. A bispectrum in the squeezed limit is in general generated by all single field models of inflation, with an amplitude which is related to the spectral index by $\fnl=-5(n_s-1)/12$ \cite{Maldacena:2002vr,Creminelli:2004yq}. Although this value is too small to be seen on CMB scales, it might be important on PBH scales, since firstly the spectral index might be larger as the slow-roll parameters potentially become larger towards the end of inflation and secondly since the constraints are sensitive to smaller values of $\fnl$. An important issue, which goes beyond the scope of this work, is the calculation of the secondary non-Gaussianities generated through the effects of gravity being non-linear and through horizon re-entry after inflation, during which time $\zeta$ is no longer conserved. These calculations have been carried out on CMB scales and these effects generally cause an order of unity change to the non-linearity parameters (although the effect is scale dependent)~\cite{Bartolo:2010qu,Creminelli:2011sq,Bartolo:2011wb,Junk:2012qt,Lewis:2012tc}. Such small values of the non-linear parameters can have a significant effect on the number of PBH's formed, therefore it would be interesting to carry out a similar analysis valid for the much smaller scales on which PBH's form. If these effects generically lead to $\fnl\sim-1$ then this would suggest that PBH's are unlikely to have formed, unless inflation generated a larger and positive primordial $\fnl$ on the same scales. The current calculation could also be extended by allowing for simultaneous non-zero values of $\fnl$ and $\gnl$ or by studying the effects of higher order non-linearity parameters. The possibility of PBH formation constrains both the amplitude and degree of non-Gaussianity associated with the primordial density perturbations over a wide range of scales, smaller than those probed by cosmological observations. In this paper we have shown that even relatively small values of the non-linearity parameters, of order unity, can have a significant effect on the PBH constraints on the amplitude of the primordial perturbations. Therefore non-Gaussianity should be taken into account when calculating PBH constraints on inflation models. We have also shown that the observation of PBH's would rule out non-negligible negative $\fnl$, re-emphasizing the constraining power of PBH's.
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1206.4188
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1206
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1206.4694_arXiv.txt
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During their main sequence lifetimes, the majority of all Galactic Disc field stars must endure at least one stellar intruder passing within a few hundred AU. Mounting observations of planet-star separations near or beyond this distance suggest that these close encounters may fundamentally shape currently-observed orbital architectures and hence obscure primordial orbital features. We consider the commonly-occurring fast close encounters of two single-planet systems in the Galactic Disc, and investigate the resulting change in the planetary eccentricity and semimajor axis. We derive explicit 4-body analytical limits for these variations and present numerical cross-sections which can be applied to localized regions of the Galaxy. We find that each wide-orbit planet has a few percent chance of escape and an eccentricity that will typically change by at least $0.1$ due to these encounters. The orbital properties established at formation of millions of tight-orbit Milky Way exoplanets are likely to be disrupted.
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After leaving their birth clusters, most stars undertake a potentially harrowing multi-Gyr journey through the Galactic Disc. The stars are continuously perturbed by global Galactic phenomena and are periodically nudged by individual stellar encounters. Occasionally, an encounter is close enough to cause major disruption to any planets orbiting in the approaching systems. The currently observed exoplanet population may be shaped in part by these encounters. \subsection{Typical Closest Encounter Distances} Using simple arguments \citep[e.g. from Pg. 3 of][]{bintre2008}, one can crudely estimate an upper bound for the typical encounter distance, $r_{enc}$, over a main sequence lifetime. If $n$ denotes the space density of stars in the Galactic Disc, and $v_{\rm ran}$ is the random velocity of stars, then $r_{enc} \approx (4 \pi n v_{\rm ran } t_{\rm MS})^{-1/2}$, where $t_{\rm MS}$ is the main sequence lifetime. This estimate is conservatively large because gravitational focusing is not included. We can estimate $t_{MS}$ through simulations from the SSE stellar evolution code \citep{huretal2000}. Doing so yields Fig. \ref{closeenc}, which plots the closest encounter main sequence distance as a function of progenitor mass from $1 {\rm M}_{\odot} - 2 {\rm M}_{\odot}$, which represents a common range of exoplanet host masses. The majority of stars drawn from a standard stellar initial mass function \citep[see e.g.][]{paretal2011} will have masses under 1 ${\rm M}_{\odot}$, further suggesting that the typical encounter separations in Fig. \ref{closeenc} represent overestimates. The solid and dashed lines represent Solar and very low (1/200th of Solar) metallicities, respectively. The metallicity of a star helps dictate its main sequence lifetime, and hence the expected close encounter distance. The plot partially illustrates that differences in the metallicity of stars have little (indirect) effect on the close encounter distance. The figure demonstrates that the majority of all stars will suffer a close encounter of just a few hundred AU for a reasonable range of $n$ and $v_{\rm ran}$ values. Even in sparse environments, like the Solar neighborhood (with $n \approx 0.1$ pc$^{-3}$), Sun-like stars will approach one another at least once within a few hundred AU. This estimate corroborates the rough estimate of $500$ AU given by \cite{zaktre2004}, who consider only a 5 Gyr encounter timescale. \begin{figure} \centerline{ \psfig{figure=CloseEnc.eps,width=8.5cm,height=6.5cm} } \caption{ Upper bound estimates for typical main-sequence closest encounter distances, $r_{enc}$, between exosystems in the field of the Galactic Disc. Solar metallicity stars and very low metallicity stars are represented by solid and dashed lines, respectively. The random velocity of stars is $v_{\rm ran}$ and the space density of stars is $n$. Fiducial values of $n$ ($0.6 {\rm pc}^{-3}$) and $v_{\rm ran}$ ($50 {\rm km/s}$) are represented by the black curves. Because the majority of Galactic Disc stars are less massive than $1 {\rm M}_{\odot}$, they will have $r_{enc}$ values less than hundreds of AU. } \label{closeenc} \end{figure} \subsection{Stellar Encounter Orientations with Respect to Galactic Centre} Given that close encounters within hundreds of AU will typically occur, we can now attempt to characterize the orientations of the collisions with respect to the Galactic Centre. As outlined by \cite{quietal2011}, the distribution of velocities in the Galactic Disc is affected by a multitude of factors. Potential perturbers include Galactic Lindblad resonances \citep[e.g.][]{yuakuo1997,lepetal2011}, stellar streams from past mergers and interactions with satellite subhaloes \citep[e.g.][]{bekfre2003,gozetal2010}, and transient spiral density waves \citep[e.g.][]{deetal2004}. Stellar velocities may also be highly dependent on the phase and pattern speed of the Milky Way's spiral arms \citep{antetal2011}, suggesting drastic differences in the velocity distribution in different regions. These factors might help explain why the velocity components of the stars in the Solar neighborhood are neither isotropic nor Gaussian \citep{binetal2000,naketal2010}. Generally, the orbits of Disc stars are modulated vertically and epicyclically \citep[e.g. pgs. 164-166,][]{bintre2008}, and may undergo significant radial migration \citep[e.g.][]{schoenrich2011}. Further, the amplitude of the epicyclic and vertical oscillations are of the same order of magnitude \citep[e.g., Pg. 18,][]{bintre2008}, and are orders of magnitude longer than then the physical radii of the stars themselves. Therefore, we should expect that stars suffer close encounters with each other at random orientations with respect to the Galactic Centre. \subsection{Planetary Orbit Orientations with Respect to the Galactic Disc} Now we assess whether the planes of the planetary orbits should have a preferential orientation to the Galactic Disc. The severe misalignment of the Solar System's invariable plane with the Galactic plane at $\approx 60^{\circ}$ \citep[][]{huawad1966,dunetal1987} foreshadows the likely answer. Observations constrain the distribution of exoplanetary orbital planes poorly because most extrasolar planets have been discovered by the Doppler radial velocity technique, which alone does not provide any information about line-of-sight inclinations. Similarly, the stellar rotational axis orientation -- which is suggestive of planetary orbit orientation -- of the vast majority of non-exoplanet host stars is unknown. However, in cases where this information has been obtained, \cite{abt2001} and \cite{howcla2009} find that the these axes are orientated randomly. For exoplanet-host stars that harbour transiting planets, we {\it do} have line-of-sight inclination information. According to the Exoplanet Data Explorer\footnote{See the Exoplanet Data Explorer at http://exoplanets.org/}, as of 15 January, 2012, there are 141 transiting exoplanets. At most, the orbital plane of any of these planets is misaligned with our line-of-sight by $\approx 13.4^{\circ}$. However, the median misalignment angle is just $\approx 2.79^{\circ}$ and the standard deviation is $\approx 2.73^{\circ}$. Therefore, effectively we observe transiting planets edge-on, and the locations of these planets on the sky might suggest a relation between the Galactic plane and planetary orbital planes. In Fig. \ref{transit}, we plot the declination versus right ascension of the host stars of these 141 transiting planets from this database. In order to help assuage the strong observational bias in the plot, the plot markers are colored and shaped according to the planet names, which are often indicative of the program or collaboration who first discovered the planet. For example, the planets with names containing {\it ``Kepler''} or {\it ``KOI''} (Kepler Object of Interest) are all clustered in the same region on the plot. This is due to the the fixed field the {\it Kepler} space mission is observing. The plot definitively illustrates that observed planetary orbital planes are known to encompass a wide range of orientations with respect to the Galactic Disc. \begin{figure} \centerline{ \psfig{figure=RaDec.eps,width=8.5cm,height=6.5cm} } \caption{ Approximate line-of-sight exoplanetary orbital plane orientations. Plotted are the spatial coordinates of stars which host transiting planets. All data is taken from the Exoplanet Data Explorer, as of 15 January, 2012. Plot markers are determined based on whether the orbiting planet's name includes ``WASP'' (blue filled circles), ``HAT-P'' (downward-pointing brown filled triangles), ``OGLE'' (hollow gray squares) ``Kepler'' or ``KOI'' (hollow red diamonds), ``TrES'' (purple filled diamonds) or ``XO'' (upward-pointing orange filled triangles). Other transiting planets are given by filled black squares. The plot demonstrates that planetary orbital planes are known to encompass a wide range of alignments with respect to the Galactic Disc. } \label{transit} \end{figure} These considerations lead us to treat close encounters between two planetary systems in arbitrary directions with orbital planes that are arbitrarily aligned with each other. However, we must sensibly restrict the vast phase space of these encounters. We do so first by reviewing some published literature related to this topic. \subsection{Extending Previous Scattering Studies} The three-body problem which includes a star-planet pair experiencing a perturbation from an intruder star has been the subject of several studies, and is well-characterized in many regions of phase space. Most studies, however, treat these interactions in the context of cluster encounters \citep{hegras1996,davsig2001,freetal2006,spuetal2009}, which typically have a higher $n$, smaller $v_{\rm ran}$, and much shorter lifetime (tens of Myr) than in the field. An exception is \cite{zaktre2004}, who do consider perturbations in the field from a stellar intruder, but on a multi-planet system. They treat the perturbation as a superposition of three-body interactions, neglecting the contribution to the potential from the planets. Also, they treat the velocity vector of the intruder and the orbital plane as coplanar. Among their several useful results are i) about $10\%$ of all stars experience close encounters within 200 AU, ii) planetary eccentricities may be excited up to $0.1$ in the field, and iii) the extent of the excitation is strongly dependent on system size and phase. Here, we provide a multi-tiered extension to that work. First, we consider the potential of all four bodies in the close encounter of two one-planet systems, as most Milky Way stars are now thought to have planets \citep{casetal2012}. Previous studies of the 4-body problem often consider the more general case of the interaction of two stellar binaries \citep{mikkola1984,hut1993,bacetal1996,heggie2000,giespu2003,freetal2004,pfamut2006,sweatman2007} or a planet-less intruder perturbing a multi-planet exosystem \citep{maletal2011,boletal2012}. However, none of these studies consider the close encounter of two single-planet systems. Second, because field encounters are fast, we develop an analytical method based on impulses that can determine the change in orbital parameters without resorting to numerical simulations. We consider two extremes in phase for our analysis, although the method can in principle be generalized to arbitrary phases, and even arbitrary numbers of planets. Third, we do perform numerical simulations, here specifically for the purpose of obtaining normalized cross sections. These quantities then enable one to determine the overall rate of encounters and eccentricity excitations over a main sequence lifetime in localized patches of the Milky Way. As already argued earlier, we consider encounters of all mutual orientations, independent of their locations with respect to the Galactic Centre. \begin{table*} \centering \begin{minipage}{180mm} \caption{Variables Used in this Paper} \begin{tabular}{@{}ll@{}} \hline Variable & Explanation \\ \hline $a_h$ & Hyperbolic semimajor axis for a star \\ $a_{k0}^{(*)}$ & Initial semimajor axis for planet $k$ in the {\tt far} ($*=f$) and {\tt close} ($*=c$) cases \\[4pt] $a_{kf}^{(*)}$ & Final semimajor axis for planet $k$ in the {\tt far} ($*=f$) and {\tt close} ($*=c$) cases \\[4pt] $a_{\chi}^{(*)}$ & Contribution to Planet \#2's semimajor axis variation due to Planet \#1 alone in the {\tt far} ($*=f$) and {\tt close} ($*=c$) cases \\[4pt] $\alpha$ & Number of planetary orbital periods to numerically integrate before the close encounter \\ $b$ & Impact parameter of both stars \\ $b_{{\rm eje}}^{(f)}$ & {\tt far} case impact parameter value at which a planet escapes \\[3pt] $b_{{\rm eje,1}}^{(c)}$ & Maximum {\tt close} case impact parameter value separating planetary escape from boundedness \\[3pt] $b_{{\rm eje,2}}^{(c)}$ & Middle {\tt close} case impact parameter value separating planetary escape from boundedness \\[3pt] $b_{{\rm eje,3}}^{(c)}$ & Minimum {\tt close} case impact parameter value separating planetary escape from boundedness \\[3pt] $b_{\rm max}$ & Maximum impact parameter used in the numerical simulations \\ $b_{\rm min}$ & Impact parameter which causes a planet-planet collision \\ $b_{p1p2}$ & Impact parameter of both planets \\[2pt] $b_{p1s2}$ & Impact parameter of Planet \#1 and Star \#2 \\ $b_{s1p2}$ & Impact parameter of Star \#1 and Planet \#2 \\ $b_{{\rm stat,<}}^{(c)}$ & {\tt close} case lower impact parameter value at which there is no net perturbation on the planets \\[3pt] $b_{{\rm stat,>}}^{(c)}$ & {\tt close} case upper impact parameter value at which there is no net perturbation on the planets \\[3pt] $\beta$ & Factor by which $(a_{10} + a_{20})$ is multiplied to obtain $q$ for the numerical simulations \\ $\gamma$ & Fraction of the innermost planetary orbit used as a numerical integration timestep bound \\ $\delta$ & Dimensionless planet/star mass ratio for each system when both are physically equivalent \\ $\delta_k$ & Dimensionless planet/star mass ratio for system $k$ \\ $e_{\rm ext,max}^{(c)}$ & {\tt close} case local eccentricity maximum, for $ (b_{{\rm stat,>}}^{(c)}) < b $ \\[3pt] $e_{\rm ext,min}^{(c)}$ & {\tt close} case local eccentricity minimum, for $ b_{{\rm eje,1}}^{(c)} < b < b_{{\rm eje,2}}^{(c)}$ \\[3pt] $e_h$ & Hyperbolic eccentricity of a star \\ $e_{kf}^{(*)}$ & Final eccentricity for planet $k$ in the {\tt far} ($*=f$) and {\tt close} ($*=c$) cases \\[3pt] $e_{\chi}^{(*)}$ & Contribution to Planet \#2's eccentricity variation due to Planet \#1 alone in the {\tt far} ($*=f$) and {\tt close} ($*=c$) cases \\[4pt] $E_h$ & Hyperbolic anomaly of a star \\ $\epsilon$ & Dimensionless ratio equal to $a_{10}/a_{20}$ \\ $\eta$ & Dimensionless ratio equal to $V_{\infty}/V_{\rm crit}$ \\ $M_{pk}$ & Mass of planet $k$ \\ $M_{sk}$ & Mass of star $k$ \\ $M_{\rm tot}$ & Total mass of the 4-body system \\ $\mu$ & Sum of both stellar masses, times the Gravitational Constant \\ $n$ & Space density of stars \\ $N$ & Number of experiments \\ $\mathcal{N}$ & Number of times over a main sequence lifetime that $\left|\Delta e_1 \right| > \Upsilon$ occurs \\ $q$ & Pericenter of the star-star hyperbolic orbit \\ $r_{\rm enc}$ & Typical closest encounter distance for two stars in the Galactic Disc \\ $r_{\rm start}$ & Separation used to initialize numerical integrations \\ ${\rm RAND}$ & Low-discrepancy quasi-random Niederreiter number between 0 and 1 \\ $\sigma$ & Cross section \\ $\sigma_{\rm norm}$ & Normalized cross section \\ $t_{\rm enc}$ & Timescale of close encounter between both planetary systems \\ $t_{\rm integrate}$ & Numerical integration timescale \\ $t_{\rm MS}$ & Main Sequence lifetime \\ $T_k$ & Orbital period of planet $k$ about star $k$ \\ $\Upsilon$ & Given extent of an eccentricity perturbation \\ $v_{\rm ran}$ & Random stellar velocity \\ $V_{\infty}$ & Velocity of Star \#1 with respect to Star \#2 at an infinite separation \\ $V_{{\rm circ},k}$ & Circular velocity of planet $k$ about star $k$ \\ $V_{{\rm circ},k0}$ & Circular velocity of planet $k$ about star $k$ assuming $M_{pk}=0$ \\ $V_{{\rm circ},0}$ & Circular velocity of either planet for equal planetary masses and semimajor axes \\ $V_{\rm crit}$ & Velocity at which the total energy of the 4-body system equals zero \\ $V_{\rm peri}$ & Pericenter velocity of the star-star hyperbolic orbit \\ $|\Delta \vec{V}_{\bot}|$ & Magnitude of the velocity kick perpendicular to the direction of motion \\ \hline \end{tabular} \end{minipage} \end{table*} \subsection{Plan for Paper} We outline some of the key quantities in the hyperbolic 4-body problem in Section 2 before our analytical (Section 3 and the Appendix) and numerical (Section 4) explorations. Of particular note are the two specific orientations we model without numerical integrations (Sections \ref{farsec} and \ref{closesec}) and the eccentricity excitation frequencies arising from our numerical integrations (Section \ref{exfreq}). In Section 5, we interpret the results. Section 6 discusses related topics, and Section 7 provides a short conclusion. \subsubsection{Variables used} Table 1 delineates the variables applied throughout this paper. The subscript $k$ takes the values ``1'' and ``2'' and is used to describe the planet and star belonging to the different planetary systems taking part in the encounter. Primed and double-primed values are explained in the text where necessary.
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We have modeled the close encounter of two single-planet exosystems in the Galactic Disc, which mimicks a common occurence during middle-aged planetary evolution. We obtained analytical formulae and numerical cross sections which may be useful for future population studies of exoplanets in specific regions of the Milky Way. The resulting change in orbital parameters for wide-orbit ($a \approx 100-1000$ AU) planets is significant (with a typical $\Delta e$ of several hundredths to over a tenth) and potentially measurable, suggesting that these planets are highly unlikely to retain a static orbit during main sequence evolution. Although tight-orbit planets (with $a \lesssim 10$ AU) are more resistant to orbital changes, millions in the Milky Way will be affected, and lose their primordial orbital signatures. The most dynamically excited Milky Way exoplanets are likely to reside in the densest Galactic regions.
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1206.4694
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1206
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1206.6691_arXiv.txt
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We present compelling evidence for the complexity of the Fornax dwarf spheroidal. By disentangling three different stellar subpopulations in its red giant branch, we are able to study in detail the dependence between kinematics and metallicity. A well-defined ordering in velocity dispersion, spatial concentration, and metallicity is evident in the subpopulations. We also present evidence for a significant misalignment between the angular momentum vectors of the old and intermediate-age populations. According to the HST measurement of Fornax's proper motion, this corresponds to counter-rotation. These ingredients are used to construct a novel evolutionary history of the Fornax dwarf spheroidal, characterized as a late merger of a bound pair.
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As the size and quality of datasets on the local population of dwarf Spheroidals (dSphs) increase, evidence for the complexity of these systems is intensifying. The observed variety of properties and structures gives each dwarf its own distinctive character. Their star formation histories, for instance, are well known to differ significantly. These often give rise to multiple generations of stars, but the mechanisms that drive an intermittent star formation must be investigated on a case-by-case basis. Being able to untangle the properties of such stellar subpopulations is crucial in advancing our understanding of how they came in place. After the Sagittarius dwarf, Fornax is the second most luminous dSph orbiting the Milky Way and the only one (in addition to Sagittarius) to posses an associated Globular Cluster system~\citep{vdB98}. This is not Fornax's only peculiarity: for example, significant asymmetries in the isophotes were recorded as early as \citet{Ho61}, \citet{Esk88} and then confirmed in \citet{Ir95}. Recent deep photometric surveys \citep[][B06 in the following]{Ste98, Sav00, Bat06} have revealed a rich and prolonged star formation, comprising old stars ($\gtrsim$ 10 Gyr), a dominant population of intermediate age stars, as well as stars of only a few hundred Myr. All systematic studies of Fornax's star formation history \citep{Col08, Pino11, deB12} record a significant starburst at approximately 4 Gyr ago. The trigger of this activity, though, is still debated. On the one hand, at least three stellar overdensities have been identified in Fornax \citep{Col04, Col05, deB12}, and interpreted as reminiscent of the shell features observed in elliptical galaxies \citep{Mal80}. In this picture, Fornax collided with a low-mass sub-halo in the relatively recent past and swallowed it. On the other hand, the details are not free from difficulties, especially caused by the inferred metallicity and age of the stars composing the innermost clump \citep{Ol06, Col08}. These are about 1.5 Gyr old, which makes the timing of the global starburst problematic. They are also relatively metal-rich, which in turn does not implicate an external origin for their gas. Furthermore, a recent collision is quite unlikely when the energetics of encounters in a virialized Milky Way halo is considered \citep{DeR04}. After the pioneering work of~\citet{Wal11} (WP11 in the following), in this {\it Letter}, we provide new evidence for Fornax's complexity by disentangling different stellar subpopulations in the red giant branch (RGB) and by studying in detail their kinematics. Rather than a division into two subpopulations \citep[][B06, WP11]{Sav00}, one into three is preferred by the data. The identified subpopulations show significant differences in line-of-sight (LOS) velocity dispersion and higher order moments. Even more striking is the incompatibility between their detected rotations, with the relative rotation vector inclined $\sim 40^{\circ}$ with respect to the isophotal major axis. Finally, we use these new ingredients to inform a novel formation scenario for the Fornax dSph.
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1206.6691
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1206
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1206.1770_arXiv.txt
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We report here results of spectropolarimetric observations of the classical T~Tauri star (cTTS) GQ~Lup carried out with ESPaDOnS at the Canada-France-Hawaii Telescope (CFHT) in the framework of the `Magnetic Protostars and Planets' (MaPP) programme, and obtained at 2 different epochs (2009~July and 2011~June). From these observations, we first infer that GQ~Lup has a photospheric temperature of $4,300\pm50$~K and a rotation period of $8.4\pm0.3$~d; it implies that it is a $1.05\pm0.07$~\msun\ star viewed at an inclination of $\simeq$30\degr, with an age of 2--5~Myr, a radius of $1.7\pm0.2$~\rsun, and has just started to develop a radiative core. Large Zeeman signatures are clearly detected at all times, both in photospheric lines and in accretion-powered emission lines, probing longitudinal fields of up to 6~kG and hence making GQ~Lup the cTTS with the strongest large-scale fields known as of today. Rotational modulation of Zeeman signatures, also detected both in photospheric and accretion proxies, is clearly different between our 2 runs; we take this as further evidence that the large-scale fields of cTTSs are evolving with time and thus that they are produced by non-stationary dynamo processes. Using tomographic imaging, we reconstruct maps of the large-scale field, of the photospheric brightness and of the accretion-powered emission at the surface of GQ~Lup at both epochs. We find that the magnetic topology is mostly poloidal and axisymmetric with respect to the rotation axis of the star; moreover, the octupolar component of the large-scale field (of polar strength 2.4 and 1.6~kG in 2009 and 2011 respectively) dominates the dipolar component (of polar strength $\simeq$1~kG) by a factor of $\simeq$2, consistent with the fact that GQ~Lup is no longer fully-convective. GQ~Lup also features dominantly poleward magnetospheric accretion at both epochs. The large-scale dipole component of GQ~Lup is however not strong enough to disrupt the surrounding accretion disc further than about half-way to the corotation radius (at which the Keplerian period of the disc material equals the stellar rotation period), suggesting that GQ~Lup should rapidly spin up like other similar partly-convective cTTSs. We finally report a 0.4~\kms\ RV change for GQ~Lup between 2009 and 2011, suggesting that a brown dwarf other than GQ~Lup~B may be orbiting GQ~Lup at a distance of only a few au's.
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\label{sec:int} It is now well recognised that magnetic fields can significantly modify the life of stars and in particular their rotation rates. Their impact is thought to be strongest throughout the formation stages, when stars and their planetary systems build up from the collapse of giant molecular clouds. More specifically, fields are likely efficient at slowing down the cloud collapse, at inhibiting the subsequent fragmentation and at dissipating the cloud angular momentum through magnetic braking and the associated magnetised outflows / collimated jets \citep[e.g.,][for reviews]{Donati09, Andre09}. At a later stage, the newly born protostars (called classical T~Tauri stars or cTTSs) are apparently capable of generating magnetic fields strong enough to disrupt the central regions of their accretion discs and to funnel some of the inner disc material onto the stellar surface, thereby drastically modifying the overall mass accretion process \citep[e.g.,][]{Bouvier07}. For some time, the strong magnetic fields of cTTSs could only be inferred through indirect proxies such as continuum or line emission throughout the whole electromagnetic spectrum, from X-rays to radio wavelengths. Directly detected for the first time about 2 decades ago through the Zeeman broadening they induce on spectral lines observed in unpolarized light \citep[e.g.,][for a recent overview]{Johns07}, magnetic fields of cTTSs can now be characterized by various means. In particular, their large-scale topologies - controlling how the fields couple the central protostars to the inner regions of their accretion discs and thereby how disc material is being accreted - can be thoroughly investigated thanks to the advent of sensitive high-resolution spectropolarimeters dedicated to the study of stellar magnetic fields \citep[e.g.,][]{Donati03, Donati09}. By measuring circularly-polarised Zeeman signatures of cTTSs and by monitoring their rotational modulation, one can reconstruct the parent large-scale magnetic topologies, thus offering the possibility of studying magnetospheric accretion processes in a much more quantitative way. The international MaPP (Magnetic Protostars and Planets) project was designed mostly for this purpose. The first main goal of MaPP is to investigate (through a first survey of $\simeq15$ targets, with some of them observed at several epochs) how the large-scale magnetic topologies of cTTSs depend on key stellar parameters such as mass, age, rotation and accretion rate \citep[e.g.,][]{Donati10b, Donati11b}. The second main goal of MaPP is to provide an improved theoretical description, using both analytical modelling and numerical simulations, of how magnetic fields of cTTSs are generated and how they modify mass accretion processes \citep[see, e.g.,][]{Gregory10, Romanova11}, and more generally how critically they impact the formation of low-mass stars. A total of 690~hr of time was allocated for MaPP on the 3.6~m Canada-France-Hawaii Telescope (CFHT) over a timescale of 9 semesters (2008b-2012b). Up to now magnetic Zeeman signatures were detected on all selected targets; several major discoveries were achieved regarding the two science goals mentioned above, that we will recall below in the light of the new results presented here. This new study focusses on the cTTS GQ~Lup, whose mass is a fair match to that of the Sun and whose age is typical of cTTSs. Located near Lupus~1, in the Lupus star formation region \citep[$150\pm10$~pc away from the Earth,][]{Wichmann99, Crawford00}, GQ~Lup recently attracted a lot of attention following the discovery of its low-mass companion \citep[most likely a brown dwarf, with a mass in the range 10--40~\mjup,][]{McElwain07,Lavigne09} in the outer regions of its accretion disc \citep[at a distance of $\simeq$0.7\arcsec\ or 100~au,][]{Neuhauser05}. Although the nature of this sub-stellar companion is still a matter of speculation, it makes GQ~Lup an obvious target of study, to investigate the properties of the central protostar on the one hand (and in particular the star-disc interaction in which the large-scale magnetic field of the protostar plays a crucial role) and to better understand how stellar / planetary systems form. In this respect, efforts at modelling the magnetic field of the protostar and associated activity are worthwhile. The large-scale field is obviously a key parameter to unravel the physics of the star-disc magnetospheric interaction, whereas the potential presence of closer companions (e.g., that could explain the ejection on an outer orbit of the very distant companion detected already) can only be revealed through high-precision radial velocity (RV) measurements when an accurate description of the magnetic activity, and an efficient way of filtering the associated RV jitter, become available. We start this paper by describing the spectropolarimetric observations of GQ~Lup we collected at 2 different epochs and from which Zeeman signatures are clearly detected (Sec.~\ref{sec:obs}). Following a fresh re-determination of the main characteristics of this cTTS (Sec.~\ref{sec:gq}), we outline the rotational modulation and intrinsic long-term variability that we observe in the data (Sec.~\ref{sec:var}). We then detail the modelling of these data with our magnetic imaging code (Sec.~\ref{sec:mod}), compare our new results with previous ones and outline how they improve our understanding of how magnetic fields impact the formation of Sun-like stars (Sec.~\ref{sec:dis}).
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\label{sec:dis} This paper presents the first spectropolarimetric analysis of the cTTS GQ~Lup, following previous similar studies of several cTTSs of various masses and ages; this analysis uses extensive data sets collected at two different epochs (2009~July and 2011~June), in the framework of the MaPP Large Program with ESPaDOnS at CFHT. From these data, we start by redetermining the fundamental characteristics of GQ~Lup, and in particular its photospheric temperature, found to be $\simeq$250~K warmer than usually quoted in the literature. We also obtain that the rotation period of GQ~Lup is $8.4\pm0.3$~d, in good agreement with the previous estimate of \citet{Broeg07}. We finally conclude that GQ~Lup is a $1.05\pm0.07$~\msun\ star with an age of 2--5~Myr and a radius of $1.7\pm0.2$~\rsun\ that has just started to build a radiative core (see Fig~\ref{fig:hrd}). Strong Zeeman signatures are detected at all times in the spectra of GQ~Lup, both in LSD profiles of photospheric lines and in emission lines probing accretion regions at the chromospheric level. We report longitudinal fields ranging from --0.1 to --0.6~kG in photospheric lines, from 0.7 to 1.9~kG in the emission core of \caii\ lines, and from 1 to more than 6~kG in the narrow emission profile of \hei\ $D_3$ lines, making GQ~Lup the cTTS with strongest magnetic fields known as of today. We find in particular that different field polarities are traced by photospheric lines and accretion proxies (as for several other cTTSs), indicating that they probe different spatial regions of GQ~Lup. Longitudinal field curves also unambiguously demonstrate that the parent large-scale field of GQ~Lup significantly evolved between our two observing runs, with magnetic intensities in accretion regions (where the field is strongest) dropping by as much as 50\% between 2009~July and 2011~June; this makes GQ~Lup the second cTTS on which temporal evolution of the large-scale magnetic topology has been unambiguously demonstrated \citep[after V2129~Oph,][]{Donati11}. Using our tomographic imaging tool specifically adapted to the case of cTTSs, we convert our two data sets into surface maps of GQ~Lup, of the large-scale vector magnetic field on the one hand, and of the photospheric brightness and of the accretion-powered excess emission on the other hand. We find in particular that the large-scale field of GQ~Lup is strong, and mostly poloidal and axisymmetric (about the rotation axis). More specifically, we find that the poloidal field is dominated by an octupolar component aligned with the rotation axis (within 10\degr) and whose strength weakens from 2.4~kG in 2009~July to 1.6~kG in 2011~June; we also find that the large-scale dipole of GQ~Lup is about half as strong as the octupole, with a polar strength equal to 1.1 and 0.9~kG in 2009 and 2011, and is tilted by $\simeq$30\degr\ to the rotation axis (and thus largely parallel, rather than anti-parallel, to the octupole). The large-scale magnetic topology that we reconstruct for GQ~Lup is fully compatible with previous results obtained on cTTSs, suggesting that large-scale fields of protostars hosting small radiative cores are mostly poloidal and axisymmetric, with a dominant octupolar component \citep[e.g.,][]{Donati11}. This is also similar to that found on main-sequence M dwarfs, where stars with radiative cores smaller than 0.5~\rstar\ host rather strong and mainly poloidal and axisymmetric fields \citep[e.g.,][Gregory et al.\ 2012, submitted]{Morin08b}. This new result further argues that the magnetic fields of cTTSs are generated through dynamo processes rather than being fossil fields initially present in the parent cloud from which the star has formed and that unexpectedly managed to survive the various turbulent episodes of the cloud-to-disc and disc-to-star contraction phases; clear observational evidence that the large-scale field significantly weakened in a 2~yr timescale independently confirms this conclusion, and indicates at the same time that the underlying dynamo processes are non-stationary. We also find that the visible pole of GQ~Lup (and presumably the invisible pole as well) hosts a cool dark spot at photospheric level and concentrates most of the accreted material from the disc, if we judge from the location of the accretion-powered area of excess \caii\ emission at chromospheric level (roughly overlapping with the cool photospheric spot). These near-polar regions are consistent with the observed low-amplitude RV rotational modulation of photospheric lines and accretion proxies. We also note that additional low-latitude features are detected in 2009~July, both in the photospheric brightness map and in the distribution of accretion-powered emission, suggesting that accretion may also occur (though marginally) at low latitudes at this epoch; a similar conclusion was proposed by \citet{Broeg07} to attempt reconciling the low inclination of GQ~Lup with the large amplitude of its light curve at some epochs. As proposed below, this may be related to the stronger octupolar component of the large-scale magnetic field in the first of our two observing runs. Given the emission fluxes of the conventional accretion proxies, we infer that the logarithmic mass accretion rate at the surface of GQ~Lup is equal to $-9.0\pm0.3$ (in \mspy), being slightly stronger (by $\simeq$0.2~dex) in 2009~July than in 2011~June. From this, we infer that GQ~Lup should be capable of magnetically disrupting its accretion disc up to a radius of $\rmag=4.6\pm1.0$~\rstar, or equivalently $0.037\pm0.008$~au \citep[assuming an average dipole strength of 1~kG and using the analytical formula of][]{Bessolaz08}; in particular, we find that \rmag\ is significantly larger than the radius at which the dipole field starts to dominate over the octupole field \citep[equal to $\simeq$1.2~\rstar, for an octupole about twice stronger than the dipole,][]{Gregory11}, in agreement with our observation that the accretion flow is mostly poleward on GQ~Lup. For larger octupole to dipole polar strength ratios (and when the dipole and octupole are parallel rather than antiparallel), one would naturally expect GQ~Lup to trigger increasingly larger amounts of low-latitude accretion \citep[e.g.,][]{Romanova11}; this may be what is happening to GQ~Lup in 2009~July, i.e., when the octupole component is strongest and low-contrast low-latitude accretion signatures are being detected. More observations are however necessary to confirm whether observations are consistent with theoretical expectations on this specific issue. When \rmag\ is compared to the corotation radius $\rcor\simeq10.4$~\rstar\ (or 0.083~au), at which the Keplerian period equals the stellar rotation period, we find that $\rmag/\rcor$ is equal to $0.45\pm0.10$, far below the value (of $\simeq1$) at which star/disc magnetic coupling can start inducing a significant outflow through a propeller-like mechanism \citep[e.g.,][]{Romanova04} and thus force the star to spin down. We thus speculate that, following the recent build-up of a radiative core and the corresponding change in its large-scale magnetic topology, GQ~Lup can no longer counteract the increase of its angular momentum (resulting from both contraction and accretion) through a star-disc braking torque, and has no other option than to undergo a rapid spin up such as that TW~Hya experienced already \citep{Donati11b}. Again, more observations are needed to check these speculations, and in particular to monitor how the magnetic field of the roughly solar-mass GQ~Lup is evolving on a longer term (at least on a timescale comparable to that of the solar cycle) and to validate whether the topology we reconstructed from our 2009 and 2011 data sets is indeed typical and adequate to predict the rotational evolution of GQ~Lup. We complete this section by briefly discussing the RV change that we report for GQ~Lup, whose average LSD photospheric profile was observed to shift from $-3.2\pm0.1$~\kms\ in 2009~July to $-2.8\pm0.1$~\kms\ in 2011~June, much larger in particular than spurious shifts potentially attributable to instrument stability problems. Being comparable to the semi-amplitude of the activity-induced modulation of the RV curve (and readily visible from the plots themselves, see Fig.~\ref{fig:var1}, upper panels), this change is also undoubtedly too large to be realistically attributed to uncorrected residuals of RV effects induced by cool spot patterns at the surface of GQ~Lup. Various possibilities thus remain to explain it. The first option is that it is caused by long-term changes of the surface convection / granulation pattern; although 0.4~\kms\ may sound fairly large (for a Sun-like star), little is known in practice about how much this effect can modify RVs in stars as active as cTTSs. The second option is that it is caused by an additional body orbiting GQ~Lup, on a timescale of months or years. We already know that GQ~Lup is orbited by a brown dwarf \citep{Neuhauser05, McElwain07, Lavigne09}, but its distance from GQ~Lup (about 100~au) is likely too large to induce a RV shift similar to the one we report here in a timescale of only 2~yr; this may suggest that GQ~Lup hosts a third body much closer to the central protostar than GQ~Lup~B, e.g., a brown dwarf of a few tens of Jupiter masses at a distance of a few au's. Additional observations of GQ~Lup similar to those presented here are obviously needed to confirm which option applies.
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1206.1770
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1206
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1206.6220_arXiv.txt
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{I shall review what has been learnt during 20 years of lithium observations in stars belonging to metal-poor globular clusters. The focus will be on little evolved main- sequence, turnoff-point (TOP) and subgiant-branch (SGB) stars expected to display Spite-plateau lithium abundances like those found in the majority of field stars of similar metallicities. But is the Spite plateau of globular clusters the same as those of field stars? What effect does, e.g., cluster-internal pollution have on lithium abundances in the now dominant second generation of stars? It will be shown that it is primarily our incomplete knowledge of the temperature scale of Population II stars which currently limits the diagnostic power of globular clusters as regards the stellar-surface evolution of lithium.
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Owing to the publish-or-perish policy, all astro\-nomers are good at presenting scientific re-sults in papers. However, I would argue we are less good at reporting the process of scientific inference: how do we formulate scientific questions and how do we go about addressing them? The answers to these questions consist of more than the capabilities of our theories, our simulations, our telescopes and instruments.\footnote{If our ``creations'' were self-aware, they would ask ``Why did you create me?''. Is our only answer ``Because it was possible.''? Cf. antiquity's Prometheus and Mary Shelley's Frankenstein.} There is human curiosity and scientific instinct, there are preconceptions and prejudice. I feel that review talks and proceedings can serve a purpose in this context. Here, one can talk about past developments and their influence on the present state of affairs and future directions. As philosopher George Santayana (1863-1952) put it: ``Those who cannot remember the past are condemned to repeat it.'' So let's take a look at 20 years of lithium studies in Galactic Globular Clusters.
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Globular cluster studies of lithium have made very significant contributions to our understanding of surface-lithium evolution in metal-poor stars. I summarize some relevant points: \begin{itemize} \item Atomic diffusion moderated by some form of extra mixing modifies the surface abundances of all Spite-plateau stars, not only but in particular in terms of lithium. The correction is at least +0.2\,dex. \item There is no convincing evidence in favour of systematically higher (or lower) lithium abundances of Spite-plateau stars residing in GCs, differences do not seem to exceed 0.05\,dex. Care should, however, be taken to measure lithium in first-generation stars whenever the GC under investigation shows prominent elemental anti-correlations (O-Na, Mg-Al). \item Remaining uncertainties in the absolute effective temperatures of warm halo stars can at most account for a 0.15\,dex shift in log $\varepsilon$(Li). However, the indiect effect of the relative effective-temperature scale on the inferred extra-mixing efficiency and thus on the diffusion correction for lithium can be as large as 0.4\,dex. Atomic diffusion moderated by efficient extra mixing (T6.25) on a hot absolute effective-temperature scale can fully bridge the gap to the WMAP-calibrated BBN prediction. MHD modelling may alleviate the need for high $T_{\rm eff}$ values.\footnote{I re-iterate my bet from 2010: If you are willing to wager a bottle of wine that the remaining $\leq$ 0.2\,dex between diffusion-corrected stellar lithium and precision-cosmology (WMAP-calibrated BBN) lithium are due to {\em new physics}, contact the author.} \item There is no good reason not to agree on departures from LTE for lithium within a given modelling framework (1D, 3D). For 1D analyses, the Lind et al. (2009b) corrections are recommended. \item GC are no direct help in addressing the paradigm-shaking meltdown of the Spite plateau at very low metallicities. \end{itemize} What a beautiful mess GC and stellar physics has become during the past decade!
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1206.6220
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1206
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1206.5824_arXiv.txt
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With the first metal enrichment by Population (Pop)~III supernovae (SNe), the formation of the first metal-enriched, Pop~II stars becomes possible. In turn, Pop~III star formation and early metal enrichment are slowed by the high energy radiation emitted by Pop~II stars. Thus, through the SNe and radiation they produce, Populations II and III coevolve in the early Universe, one regulated by the other. We present large (4 Mpc)$^3$, high resolution cosmological simulations in which we self-consistently model early metal enrichment and the stellar radiation responsible for the destruction of the coolants (H$_{\rm 2}$ and HD) required for Pop~III star formation. We find that the molecule-dissociating stellar radiation produced both locally and over cosmological distances reduces the Pop~III star formation rate at $z$ $\ga$ 10 by up to an order of magnitude, to a rate per comoving volume of $\la$ 10$^{-4}$ M$_{\odot}$ yr$^{-1}$ Mpc$^{-3}$, compared to the case in which this radiation is not included. However, we find that the effect of LW feedback is to enhance the amount of Pop~II star formation. We attribute this to the reduced rate at which gas is blown out of dark matter haloes by SNe in the simulation with LW feedback, which results in larger reservoirs for metal-enriched star formation. Even accounting for metal enrichment, molecule-dissociating radiation and the strong suppression of low-mass galaxy formation due to reionization at $z$ $\la$ 10, we find that Pop~III stars are still formed at a rate of $\sim$ 10$^{-5}$ M$_{\odot}$ yr$^{-1}$ Mpc$^{-3}$ down to $z$ $\sim$ 6. This suggests that the majority of primordial pair-instability SNe that may be uncovered in future surveys will be found at $z$ $\la$ 10. We also find that the molecule-dissociating radiation emitted from Pop~II stars may destroy H$_{\rm 2}$ molecules at a high enough rate to suppress gas cooling and allow for the formation of supermassive primordial stars which collapse to form $\sim$ 10$^5$ M$_{\odot}$ black holes.
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The gravitational collapse of primordial gas into the first, Population (Pop)~III stars at $z$ $\ga$ 20 marks the end of the cosmic dark ages and ushers in an era of rapidly increasing complexity in the early Universe (e.g. Barkana \& Loeb 2001; Ciardi \& Ferrara 2005). Expected to be typically much more massive than most stars formed today (e.g. Bromm \& Larson 2004; Glover 2005), the first stars emit copious high energy radiation that ionizes and heats the gas in and around their host dark matter (DM) haloes (e.g. Kitayama et al. 2004; Whalen et al. 2004; Alvarez et al. 2006; Abel et al. 2007; Johnson et al. 2007), and also destroys H$_{\rm 2}$ and HD molecules in the primordial gas over cosmological distancees (e.g. Haiman et al. 1997; Glover \& Brand 2001; Ahn et al. 2009; Holzbauer \& Furlanetto 2012). As these molecules are critical for the cooling of the primordial gas, it is expected that their destruction leads to diminished rates of gas collapse and Pop~III star formation (e.g. Omukai \& Nishi 1999; Machacek et al. 2001; Mesinger et al. 2006a; Wise \& Abel 2007; O'Shea \& Norman 2008; Trenti et al. 2009). At the end of their brief (10$^6$ - 10$^7$ yr) lives, a large fraction (e.g. Heger et al. 2003) of Pop~III stars explode as supernovae (SNe) and eject the first heavy elements into the intergalactic medium (e.g. Ferrara et al. 2000; Bromm et al. 2003; Kitayama \& Yoshida 2005; Greif et al. 2007; Vasiliev et al. 2008; Whalen et al. 2008). This sets the stage for the onset of metal-enriched, Pop~II star formation (e.g. Wise \& Abel 2008; Greif et al. 2010; Maio et al. 2011) in the first galaxies (e.g. Bromm \& Yoshida 2011). The lower surface temperatures of Pop~II stars compared to Pop~III stars lead to a lower efficiency (per stellar mass) of ionizing photon production (e.g. Oh et al. 2001; Tumlinson et al. 2001; Schaerer 2002), but also lead to only a modest decrease in the efficiency of molecule-dissociating, Lyman-Werner (LW) photon production (e.g. Greif \& Bromm 2006), as well as to an increased efficiency in the production of H$^-$-ionizing photons (e.g. Shang et al. 2010).\footnote{Because molecular hydrogen is formed in the primordial gas largely via the reaction H$^-$ + H $\to$ H$_{\rm 2}$ + e$^-$, the destruction of H$^-$ effectively slows the production of H$_{\rm 2}$ (see e.g. Chuzhoy et al. 2007).} Thus, the high energy radiation emitted from early generations of Pop~II stars can have a dramatic impact in slowing the rate of Pop~III star formation. This interplay between Pops~II and III star formation constitutes a feedback loop whereby Pop~II star formation can only take place in regions enriched by Pop~III stars and the pace of Pop~III star formation (and the subsequent metal enrichment) is regulated by the amount of radiation emitted by Pop~II stars. Therefore, in order to properly model the earliest episodes of star and galaxy formation, it is necessary to model the formation of both populations and their respective chemical and radiative feedback, as a coupled system. Much previous work has treated these processes, with many results gleaned from cosmological simulations of early SNe feedback and metal enrichment (e.g. Tornatore et al. 2007; Wise \& Abel 2008; Wiersma et al. 2009a; Greif et al. 2010; Maio et al. 2011; Wise et al. 2012) and of the build-up of the global background (e.g. Yoshida et al. 2003; Wise \& Abel 2005; Johnson et al 2008) or the locally generated (e.g. Dijkstra et al. 2008; Ahn et al. 2009; Hummel et al. 2011; Petkova \& Maio 2011; Wise et al. 2012) stellar LW radiation field. Previous authors have also modelled the impact of the LW background on star formation, both self-consistently (Ricotti et al. 2002; 2008; Trenti et al. 2009; Petri et al. 2012) and at fixed levels (e.g. Kuhlen et al. 2012; Safranek-Shrader et al. 2012), together with simplified treatments of early metal enrichment. Assembling all of these ingredients -- star formation, mechanical feedback and metal enrichment from SNe, and both local and cosmological stellar LW radiation fields -- self-consistently in a cosmological context is the next step in simulating the formation of the earliest galaxies. Here we present the results of cosmological simulations which accomplish this task, self-consistently accounting for early metal enrichment and mechanical feedback from Pop~II and Pop~III SNe, as well as the impact of both the locally-generated and the cosmological background stellar LW radiation from both populations. While our simulation is of high enough resolution to track even the first episodes of star formation in minihaloes, we simulate a relatively large cosmological volume in order to follow the assembly of galaxies down to $z$ $\simeq$ 6. Thus, as we model galaxy formation in detail from the epoch of the first stars through the entire epoch of reionization, our results offer arguably the most complete picture to date of galaxy formation in the early Universe. In the next Section, we begin by describing the simulations that we have carried out, with particular attention paid to our implementation of LW feedback. In Section 3 we present our results, highlighting the impact that LW radiation has on star formation and chemical enrichment. Finally, we give our conclusions and provide a brief discussion of our results in Section 4.
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We have presented the results of large-scale cosmological simulations of the earliest stages of galaxy formation, in which the feedback from both Pop~II and III stars plays a defining role. In modelling the mechanical, chemical, and radiative feedback processes that regulate the formation of both of these stellar populations, we have tracked their coevolution in a self-consistent manner. We have found that in the early Universe the Pop~III star formation rate is largely regulated by the global LW background, while Pop~II star formation rate is instead largely regulated by the pace of metal enrichment. As these feedback processes are linked, the two Populations coevolve in a complex manner. In particular, we find that the Pop~III SFR is regulated to be at a level very close to that at which LW feedback becomes effective in suppressing the cooling of primordial gas in minihaloes ($J_{\rm 21,bg}$ $\simeq$ 0.04), as also found previously in much more simplified simulations (e.g. Johnson et al. 2008). The total SFR is also decreased due to LW feedback at early times, both due to the lower Pop~III SFR and to the slower pace of chemical enrichment, which by definition must precede Pop~II star formation. At later times ($z$ $\la$ 11), however, most likely due to the decreased mechanical feedback from SNe in blowing away the gas in relatively small DM haloes, we find that the effect of LW radiation is to raise the total SFR by a factor of $\la$ 2 above the rate obtained in its absence. This leads to the counter-intuitive result that the total mass in stars formed by $z$ $\sim$ 6 is in fact increased due to LW radiation. While all of the feedback effects that we have included are expected to negatively impact the rate of Pop~III star formation, we find that down to $z$ $\sim$ 6 Pop~III stars still form at a rate per comoving volume of $\sim$ 10$^{-5}$ M$_{\odot}$ yr$^{-1}$ Mpc$^{-3}$, just one order of magnitude below its peak value at $z$ $\sim$ 10. This continuation of Pop~III star formation down to such low redshifts implies that $\sim$ 80 percent of primordial PISNe occur at $z$ $\la$ 10. However, we have also confirmed that their overall low rate of occurence will likely require many fields of view to be surveyed by the JWST in order even a single PISN to be discovered. While our simulations are some of the largest and most comprehensive to date, we have not fully self-consistently tracked every important physical process governing galaxy formation. In particular, we have taken a simplified approach to account for the impact of reionization on the heating of the IGM, and we have neglected the impact of ionizing radiation on the ISM surrounding stellar clusters. Nonetheless, we find that our simple assumption of global reionization beginning at $z$ = 12 yields a global SFR that is sufficient to completely reionize the IGM by $z$ $\simeq$ 6. We note, however, that not accounting for the photoionization of the ISM within star-forming haloes, especially in the first Pop~III star-forming haloes, likely makes our results lower limits for the efficiency of metal enrichment; as shown by e.g. Kitayama \& Yoshida (2005) and Whalen et al. (2008), photoheating lowers the density of the ISM and allows SN ejecta to more readily escape into the IGM. Incidentally, because this photoheating of haloes by internal stellar sources leads to more gas being blown out of low-mass Pop~III star-forming haloes, accounting for it would also likely result in a further enhancement of Pop~II star formation at late times in the case with LW feedback. Finally, because we have modelled the spatial and temporal variation of the LW radiation generated by stars, together with metal enrichment, in a sufficiently large cosmological volume, we are able to identify regions of dense, primordial gas exposed to very high levels of LW radiation. We find that such regions exist at $z$ $\ga$ 6 in our 64 Mpc$^3$ (comoving) cosmological volume. This result corroborates other recent work (e.g. Agarwal et al. 2012) in supporting the theory that these sites may play host to the formation of the supermassive ($\ga$ 10$^4$ M$_{\odot}$) stellar seeds of the black holes inhabiting the centres of galaxies today. To summarize, our most important results are the following: \begin{itemize} \item Despite the negative feedback from LW radiation, chemical enrichment and photoionizing radiation during reionization, significant Pop~III star formation continues down to at least $z$ $\simeq$ 6 (see Section 3.1.1). \item We find that LW feedback leads to an overall enhancement in Pop~II star formation, as compared to the case without LW feedback. We attribute this to the fact that Lyman-Werner feedback delays the onset of Pop~III star formation until haloes are larger and less susceptible to gas blow-out by SNe, which results in larger reservoirs of gas for Pop~II star formation (see Section 3.1.2). \item Sufficiently high LW fluxes are produced for the primordial gas to collapse into $\sim$ 10$^5$ M$_{\odot}$ black holes by direct collapse, even within our relatively small (4 Mpc)$^3$ comoving simulation volume (see Section 3.4). \item Due to the continuation of Pop~III star formation down to $z$ $\simeq$ 6, a majority of the primordial PISNe that may be detected in upcoming surveys will likely be found at $z$ $\la$ 10 (see Section 3.5). \end{itemize} We have focused in the present work on the role of the LW radiaton produced by stars in governing the global SFR and metal enrichment in the early Universe. Many other issues, including the progress of chemical enrichment and the transition from Pop~III to Pop~II star formation, the global and individual properties of galaxies and DM haloes formed by $z$ $\sim$ 6, and the role of star-forming galaxies in the inhomogeneous process of reionization, will be explored in greater detail in additional work in the FiBY project.
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1206.5824
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1206
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1206.6104.txt
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{}{}{}{}{} \abstract % context {The ESO Public Survey ``VISTA Variables in the V\'ia L\'actea'' (VVV) \new{provides deep multi-epoch infrared observations for an unprecedented 562 sq. degrees of the Galactic bulge, and adjacent regions of the disk.}} % aims {\new{The VVV observations will foster the construction of a sample of Galactic star clusters with reliable and homogeneously derived physical parameters (e.g., age, distance, and mass, etc.). In this first paper in a series, the methodology employed to establish cluster parameters for the envisioned database are elaborated upon by analysing 4 known young open clusters: Danks\,1, Danks\,2, RCW\,79, and DBS\,132. The analysis offers a first glimpse of the information that can be gleaned from the VVV observations for clusters in the final database.}} % methods {Wide-field, deep $JHK_s$ VVV observations, combined with new infrared spectroscopy, are employed to constrain fundamental parameters for a subset of clusters.} % results {\new{%Results inferred from the deep near-infrared photometry, which features small uncertainties (e.g. the precision of the photometry is better than 0.1~mag for $K_s<18$~mag), the wide field-of-view of the VVV survey, and numerous high quality low resolution spectra (typically more than 10 per cluster), are used to establish independent cluster parameters which enable existing determinations to be assessed. Results are inferred from VVV near-infrared photometry and numerous low resolution spectra (typically more than 10 per cluster). The high quality of the spectra and the deep wide--field VVV photometry enables us to precisely and independently determine the characteristics of the clusters studied, which we compare to previous determinations. An anomalous reddening law in the direction of the Danks clusters is found, specifically $E(J-H)/E(H-Ks)=2.20\pm0.06$, which exceeds published values for the inner Galaxy. The G305 star forming complex, which includes the Danks clusters, lies beyond the Sagittarius-Carina spiral arm and occupies the Centaurus arm. Finally, the first deep infrared colour-magnitude diagram of RCW\,79 is presented, which reveals a sizeable pre-main sequence population. A list of candidate variable stars in G305 region is reported.}} % conclusions {\new{This study demonstrates the strength of the dataset and methodology employed, and constitutes the first step of a broader study which shall include reliable parameters for a sizeable number of poorly characterised and/or newly discovered clusters.}} % This study reaffirms the importance of the latest generation of infrared surveys (e.g., VVV) for fostering insight into heavily obscured young clusters
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\begin{figure}[htbp] \centering \includegraphics[width=7.5cm]{fig1a.eps} \includegraphics[width=7.5cm]{fig1b.eps} \includegraphics[width=7.5cm]{fig1c.eps} \caption{$JHK_s$ true colour images of the three clusters. Stars marked by circles were observed using near infrared spectrographs. In green are the stars observed using SofI at NTT and in red are the stars observed with MMIRS at the Clay telescope. The stars marked with a yellow circle in RCW\,79 were already observed by Martins et al. (\cite{Ma10}) and, hence, were not re-observed by us.} \label{FigTrue} \end{figure} It is commonly accepted that the majority of stars with masses $\geq$ 0.50 M${_\odot}$ form in clustered environments (e.g. Lada \& Lada \cite{La03}, de Wit at al. \cite{Wi05}), rather than individually. Our location within our own Galaxy gives us a unique perspective from which we can study star clusters in great detail and such studies have important implications for our understanding of the formation of large galaxies in general. Estimates indicate that the Milky Way (MW) presently hosts 23000--37000 or more star clusters (Portegies Zwart et al. \cite{Po10}). However, only 2135 open clusters have been identified (according to the 26 Jan 2012 version of Dias et al. \cite{Di02}), which constitute a sample affected by several well known selection effects (as with globular clusters; Ivanov et al. \cite{Iv05}). Less than half of these clusters have actually been studied, and this subset suffers from further selection biases. Extending this current sample towards the Milky Way's highly obscured central region would provide unique insight into the formation, evolution, and dissipation of open clusters in the Galactic environment. To achieve this goal, we are using the unprecedented deep infrared data from the VISTA Variables in the V\'ia L\'actea (VVV) survey (Minniti et al. \cite{Mi10}, Saito et al. \cite{Sa12}), one of the six ESO Public Surveys operating on the new 4-meter Visible and Infrared Survey Telescope for Astronomy (VISTA). We are in the process of building a large sample of star clusters (including many discovered by our group; Borissova et al. \cite{Bo11}, Bonatto et al. in prep), that are practically invisible in the optical bands. The strength of this sample will lie in the homogeneity of the data (i.e. all observed with the same instrument and set-up) and analysis employed. From which, we will estimate clusters' physical parameters, including: angular sizes, radial velocities (RVs), reddening, distances, masses, and ages. Moreover, as pointed out by Majaess et al. (\cite{Maj12}), VVV photometry allows these parameters to be determined with unprecedented accuracy for highly obscured clusters. \new{As a first step}, we are focusing our efforts on young open clusters in their first few Myrs. During this period, which corresponds to Phase $I$ in the recent classification of Portegies Zwart et al. \cite{Po10}, stars are still forming and the cluster contains a significant amount of gas. The evolution of the cluster during this phase is governed by a complex mixture of gas dynamics, stellar dynamics, stellar evolution, and radiative transfer, and is currently not completely understood (Elmegreen \cite{El07}, Price \& Bate \cite{Pr09}). Thus many basic (and critical) cluster properties, such as the duration and efficiency of the star-formation process, the cluster survival probability and the stellar mass function at the beginning of the next phase are uncertain. In this paper, we present a first sub-sample of 3 known young open clusters, studying them with both VVV colour-magnitude diagrams (CMDs) and low resolution near-infrared spectroscopy of the brightest stellar members. These clusters are RCW\,79, already studied by Martins et al. (\cite{Ma10}), and Danks\,1 and Danks\,2, discussed by Davies et al. (\cite{Da11}). We aim to describe our approach and present a first glimpse of the data quality. We detail our approach for determining the physical parameters of clusters observed with VVV, using previous works as references. As DBS\,132 is located close to Danks\,1 and Danks\,2, we also examine its CMDs and derive some preliminary cluster parameters. We begin by presenting the data in Section\,\ref{obs}, and relating our method and evaluating the accuracy of our work in Section\,\ref{res}. Subsequently, we describe the stellar variability detected in the clusters and their surroundings in Section\,\ref{var}. Then, in Section\,\ref{Discussion} we compare our results with previous studies and briefly discuss the characteristics of the star-forming regions in which the clusters are situated. Before concluding by summarising this work in Section\,\ref{Summary}.
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\label{Discussion} \subsection{The G305 star forming complex} This study provides additional information about the G305 star forming complex. In addition to the diffuse population of massive stars mentioned in Davies et al. (\cite{Da11}, see also Shara et al. \cite{Sh09} and Mauerhan et al. \cite{Ma11}), our spectral observations reveal 12 early B stars with distances that are comparable with Danks\,1 and Danks\,2. Therefore, although not cluster members they are definitely members of the G305 star forming complex. At this stage, it is hard to confirm if they are or are not runaway stars, because we cannot yet measure the proper motions using VVV data. Alternatively, they could be part of a larger association of young stars surrounding Danks\,1 and 2, formed within the same molecular cloud. Taking advantage of the VVV wide field of view, we examined the stellar content of the regions outlined by Hindson et al. (\cite{Hi10}). There are a number of already catalogued clusters in this region, namely: Danks\,1, Danks\,2, DBS\,83, DBS\,84, DBS\, 130, DBS\,131 (IR cluster G305.24+0.204, Clark et al. \cite{Cl04}, Leistra et al. \cite{Le05}, Longmore et al. \cite{Lo07}), DBS\,132, DBS\,133, DBS\,134, and G305.363+0.179 (Clark et al \cite{Cl04}). In addition to these, we have found three new young star clusters and/or stellar groups: VVV CL023, VVV CL022, VVV CL021 (Borissova et al. \cite{Bo11}), and a new star forming region SFR1. The preliminary VVV CMD of G305.24+0.204 (DBS\,131) is shown in Figure\,\ref{Fig131}, where, in addition to a well defined MS, highlighted previously by Leistra et al. (\cite{Le05}) and Longmore et al. (\cite{Lo07}), the PMS population is readily apparent in the VVV data. The adopted parameters of the plot are $m-M=12.85$~mag (3.72~ kpc), $E(J-K)=2.25$~mag, age 3--5~Myr. DBS\,131, like all other clusters from the G305 complex listed here, is very young (less than 5~Myr) and less massive than Danks\,1 and Danks\,2. A more detailed analysis of these objects will be presented in the next paper in our series (Borissova et al., in preparation). \begin{figure} \centering \includegraphics[width=8cm]{fig12.eps} \caption{$(J-K_s)$ vs. $K_s$ diagram for DBS\,131, presented as in Figure\,\ref{FigCMD}. Dark dots are field stars and blue crosses are cluster stars.} \label{Fig131} \end{figure} \subsection{Milky Way's structure} The clusters' context within the Milky Way's structure is now assessed. The depth of the VVV photometry permits \new{the mapping} of crowded low latitude Galactic fields (Minniti et al. \cite{Mi11}). The youth of the stellar constituents inferred from the spectral types (Table\,\ref{table:2}) implies that Danks\,1 and Danks\,2 are viable tracers of the Galaxy's spiral structure. Georgelin et al. (\cite{Ge88}) first mentioned that the G305 region should be within the Scutum-Crux arm of the Galaxy. Later, Baume et al. \cite{Ba09} noted that Danks\,1 and Danks\,2 may belong to the Carina spiral arm. That conclusion is tied to their distance, which is significantly nearer ($>50$\%) than that found here. Their optical distance is acutely sensitive to variations in $R_V$, particularly since the total extinction in the optical $A_V=R_V \times E(B-V)$ is sizeable\new{, and not constrained by any spectroscopic observation. The present results} support the larger distance cited by Davies et al. (\cite{Da11}). The positions of Danks\,1 and Danks\,2 are plotted, in Figure\,\ref{FigGS}, on a hybrid map of the Galaxy's spiral structure, as delineated by long-period classical Cepheids and young ($<10$~Myr) open clusters. Long-period Cepheids are more massive younger stars than their shorter period counterparts, since the variables follow a period-luminosity relation. Cepheid variables and young open clusters define an analogous (local) spiral pattern (e.g. Majaess et al. \cite{Maj09}). Danks\,1 and Danks\,2 lie beyond the Sagittarius-Carina spiral arm and occupy the Centaurus arm, along with numerous young Cepheids and clusters (e.g., TW Nor, VW Cen, and VVV CL070). VVV CL070 was discovered in the comprehensive survey by Borissova et al. (\cite{Bo11}), who discovered 96 new clusters in the region sampled by the VVV survey. \begin{figure} \centering \includegraphics[width=9.0cm]{fig14.eps} \caption{Map of local spiral structure as delineated by long period classical Cepheids (dots) and young clusters (circled dots) (see also Majaess et al. \cite{Maj09}). The Carina (A) and Centaurus (E) spiral arms are indicated on the diagram. Danks\,1 and Danks\,2 (red dot) reside in the Centaurus spiral arm.} \label{FigGS} \end{figure}
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1206.6104
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1206
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1206.2545_arXiv.txt
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The stellar scintillation is one of the fundamental limitation to the precision of ground-based photometry. The paper examines the problem of correlation of the scintillation of two close stars at the focus of a large telescope. The derived correlation functions were applied to data of the long-term study of the optical turbulence (OT) in the Northern Caucasus with MASS (Multi-Aperture Scintillation Sensor) instrument to predict the angular correlation of the scintillation at the Sternberg institute 2.5~m telescope currently in construction. A median angular radius of the correlation as large as 20 arcsec was found for the case of Kolmogorov OT. On the basis of the obtained relations we also analyze the correlation impact in ensemble photometry and conjugate plane photometry. It is shown that a reduction of the scintillation noise up to 8 times can be achieved when using a crowded ensemble of comparison stars. The calculation of the angular correlation can be repeated for any large telescope at the site where the OT vertical profiles are known.
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Many astronomical problems require high-precision photometric data. Two clear examples are the determination of parameters of extraterrestrial planetary systems from eclipsing light curves \citep{Everett2001,Winn2009,Southworth2009}, and the study of low-amplitude stellar pulsations, reflecting inner structure of stars \citep{Heasley1996,Aerts2010book}. The success of the CoRoT and Kepler space missions clearly demonstrates the need for further increase of the accuracy of astronomical photometry. Despite the superiority of space projects, observations with ground-based telescopes also have some advantages. They are more accessible, the number of the telescopes is large and their service life is tens times larger than for space missions. The larger apertures of these telescopes push the intrinsic limit of the accuracy, caused by photon noise, towards fainter objects or improved temporal resolution. The main limitation to obtain high-precision photometric results from the ground is due to the Earth's atmosphere, in particular the atmospheric scintillation of stars. The stellar scintillation is a random temporal fluctuations of the radiation entering the telescope aperture and is related to the amplitude distortions of the light wave having passed through the turbulent atmosphere. Photometric errors induced by the scintillation ({\it scintillation noise}) have repeatedly been studied \citep[see, e.g.,][]{Young1969,Dravins1997a}. In comparison to the photon noise, the scintillation noise does not depend on the brightness of the stars and thus cannot be reduced by choosing brighter objects. For typical exposure times (few seconds and longer), the scintillation noise decreases more slowly with telescope diameter than the photon noise and thus becomes to dominate on large telescopes. Consequently exactly the scintillation noise sets a limit to the photometric accuracy \citep{Heasley1996}. This problem becomes even more relevant for exposures shorter than 1~s. In contrast to the flux variations due to changes in atmospheric transparency, the stellar scintillation is characterized by a lower spatial (angular) correlation \citep{Dravins1997a}. That leads to an increase of the errors in differential photometry when using a distant comparison star. It is assumed that the angular correlation for the scintillation is typically a few arcseconds but the special case of large telescopes remains virtually unexplored. This paper is a theoretical study of the angular correlation of the stellar scintillation on the basis of researches on the scintillation as a tool for probing OT above astronomical sites. We first consider the link between OT parameters and photometric scintillation noise. Then the formulas for the correlation at large telescopes in two limiting regimes, short (much less than 1~s) and long (greater than 1~s) exposures, are derived and analyzed. In Section~\ref{sec:prognoz}, the results are applied to the OT vertical distribution above Mt Shatdzhatmaz in the North Caucasus \citep{2010MNRAS} where the Sternberg Astronomical Institute (SAI) 2.5~m diameter telescope is to be installed. The angular correlation functions and typical angles of the correlation are calculated. Finally, the results are applied to two methods aimed to reduce the scintillation noise, the recently proposed photometry in a conjugated entrance pupil \citep{Osborn2010} and the more traditional ensemble photometry \citep{Gilliland1993}.
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\label{sec:conclusion} In this paper, we discussed the problem of correlation of the stellar scintillation from two point-like sources, closely spaced in the sky. Using the standard model of the optical turbulence (Kolmogorov spatial spectrum of refractive index fluctuations) in the approximation of the weak perturbations, two situations were considered for large telescope aperture (much larger than the Fresnel raduis $r_\mathrm{F}$): 1) observations with short exposures, 2) observations with long exposures (typical photometric practice). Analytic expressions obtained for the ACFs, have lengthy description, therefore a simple power-law approximation, well describing the functions, has been proposed. As the {\it measure} of the angular correlation, the angle $\theta_1$, at which the correlation coefficient is equal 1/2, was adopted. Based on previously obtained data on the vertical distribution of the OT in the atmosphere above Mt Shatdzhatmaz in the Northern Caucasus, the simulation experiments have been performed for the SAI 2.5~m telescope in order to predict the angular correlation of the scintillation for photometric observations at this observatory. All hereinafter referred numerical results relate to this telescope and this site. Since in case of a large telescope, the ACF depends only on one parameter: the normalised angle $\gamma = \theta\,z/D$, it is possible to scale the obtained estimates for telescope with different diameter. Central obscuration in a telescope affects the correlation of the scintillation for short exposures but not for long exposures. Simulation for short exposure regime (fast photometry) predicts the median of the $\theta_1$ as large as $7$ arcsec. In the long exposure regime, the ACF averaged over all wind directions, does not depend on wind speed and exposure. In this case, estimated median of the correlation angles $\theta_1$ amounts to $\approx 20$ arcsec. Study of the conjugate-plane photometry shows that the method depends critically on the intensity of low-altitude OT. When the lower atmosphere is calm, one can use the comparison star at distance greater than several arcminutes. Analysis of the ensemble photometry shows that the compact, but quite resolved, ensemble of four comparison stars is enough to get the {reduction in the scintillation noise power by} $\approx 8$ times. Photometric observation of the different ensemble configurations could provide the necessary experimental verification of the proposed formulae and the performed simulations.
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1206.2545
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1206
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1206.0295_arXiv.txt
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We examine galaxy formation in a cosmological AMR simulation, which includes two high resolution boxes, one centered on a $3\times 10^{14}\ \Msun$ cluster, and one centered on a void. We examine the evolution of $611$ massive ($M_* > 10^{10} \Msun$) galaxies. We find that the fraction of the final stellar mass which is accreted from other galaxies is between $15$ and $40\%$ and increases with stellar mass. The accreted fraction does not depend strongly on environment at a given stellar mass, but the galaxies in groups and cluster environments are older and underwent mergers earlier than galaxies in lower density environments. On average, the accreted stars are $\sim 2.5$ Gyrs older, and $\sim 0.15$ dex more metal poor than the stars formed in-situ. Accreted stellar material typically lies on the outskirts of galaxies; the average half-light radius of the accreted stars is {\rm $2.6$ times larger} than that of the in-situ stars. This leads to radial gradients in age and metallicity for massive galaxies, in qualitative agreement with observations. {\rm Massive galaxies grow by mergers at a rate of approximately $2.6\%/ \mathrm{Gyr^{-1}}$}. These mergers have a median (mass-weighted) mass ratio less than $0.26\pm 0.21$, with an absolute lower limit of $0.20$, for galaxies with $M_*\sim10^{12}\ \Msun$. This suggests that major mergers do not dominate in the accretion history of massive galaxies. All of these results agree qualitatively with results from SPH simulations by \citet{Oser10, Oser11}.
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\label{intro} In a \LCDM{} cosmology, galaxy growth is hierarchical; large galaxies form by accreting smaller systems of gas, stars, and dark matter. This model has been extensively tested by simulations and observations. However, it is not known whether stars or gas dominate the accretion history of galaxies, and what effects this accretion has on the observable properties of the galaxies. With modern hydrodynamical simulations, we can directly measure what fraction of the stars in a massive galaxy have been formed in-situ and what fraction have been made in other galaxies which were subsequently accreted by the parent galaxy. Similarly, we can map out when accretion occurs, what kind of accretion events dominate, and how the accretion affects the observable properties of galaxies. In this work, we explore these effects for a set of massive galaxies from a hydrodynamical cosmological simulation. \citet{Oser10} (hereafter Os10) present a useful framework for studying the hierarchical build-up of galaxies. {\rm Based on hydrodynamical simulations of early-type galaxy formation, they suggest that massive galaxy formation can be divided into two phases: an early, rapid in-situ star formation period followed by a late merger-dominated period. The former period is similar to the picture previously called ``monolithic collapse'' \citep{ELS62, Larson1975, Carlberg1984} and {\rm currently} labelled ``cold flow-driven'' star formation \citep[e.g.][]{Dekel06,Dekel09,Elmegreen09, Ceverino2010}, while the latter phase is similar to prior pictures of galactic cannibalism \citep{OstrikerHausman77, White1978, SearleZinn78, Kauffmann1993}.} In this scheme, individual stars can be classified according to whether they formed within the virial radius of the final system (in-situ) or outside of it (accreted). Os10 show that this distinction between stars is unambiguous; there is a large spatial separation between the birth places of stars formed in-situ and those added in accretion events. They further show that massive galaxies are dominated by accreted stars, which are typically older than stars formed in-situ. In this work, we find that the accreted stars in a galaxy are older, more metal-poor, and found at larger radii in the final galaxy than stars formed in-situ. The metallicity trends are consistent with observations of radial trends in colour \citep[e.g.][]{deVaucouleurs61, Tortora2010, Tal2011}, metallicity \citep{Spinrad1971, Faber1977, Davies1993, Brough2007, Rawle2008, Spolaor2010, Kuntschner2010, Coccato2011} and globular cluster metallicities \citep{Forbes2011, Arnold2011} in early-type galaxies. {\rm The differences between accreted and in-situ stars are also seen in disc galaxies. For example, in the Milky Way stellar halo, there has long been observational evidence for at least two stellar populations. The inner halo consists of stars with high $\alpha-$element abundances \citep[e.g.][]{Nissen2010} and rotates with the galactic disk \citep{Deason2011}, while the outer halo has low $\alpha$-element abundances and is not rotating, suggesting it was created by accretion of satellite galaxies. Simulations of disc galaxies show that the differences in dynamics and chemical abundances in the stellar halo are dues to the differences in the accreted and in-situ stellar populations \citep{Brook2004, Abadi2006, Zolotov2009, Zolotov2010, Font2011}. } The two phase model of galaxy formation (i.e., early, in-situ star formation and late accretion) is also useful in explaining the evolution of compact, massive ellipticals at $z \sim 2$ \citep{Trujillo07, vanDokkum08, Damjanov09, vanDokkum2010, Oser11}. Both direct profile measurements and velocity dispersions have shown that these systems are $\sim 100$ times denser (within $1$ effective radius) than present-day ellipticals of the same mass \citep{Daddi2005, vanDokkum08, vanderWel2008, vanderWel2011, vandeSande11}. However, if compared on same {\it physical} scale, the densities in the central portions of $z\approx2$ and present-day ellipticals are similar, suggesting that early-type galaxies have increased their size through minor, dry mergers which add stars to the outskirts of massive galaxies \citep[][but see \citealp{Newman2012}]{Naab07, Naab09, Bezanson09, Hopkins09a, Carrasco2010, Oser11, Tal2012}. {\rm By using close companions to estimate the merger rate of compact galaxies at $0.4<z<2.5$, \citet{Newman2012} observe that minor merging may not be sufficient to explain the size evolution of compact galaxies from $z\approx2$.} However, the size evolution due to minor mergers has been found in hydrodynamical simulations \citep[e.g.][]{Naab09, Oser11} and is consistent with the size growth found by observations. In this work, we make use of galaxy catalogs from cosmological simulations done with adaptive mesh refinement (AMR) \citep{Cen2010}. {\rm The use of an Eulerian grid-based code instead of a Lagrangian particle-based code is a notable distinction between this work and that of Os10. \citet{Scannapieco2011} find that different numerical hydrodynamics can yield differences of a factor of two in simulated galaxy properties. However, \citet{Scannapieco2011} find that changes in the feedback implementation can yield even larger differences in the simulation results. We will address the differences in sub-grid physics between our work and that of Os10 below. } The AMR simulation studied here contains two high resolution boxes, one centered on a galaxy cluster and another centered on a void. % These large volumes simulated at high resolutions allow us to study the merging histories of more than $600$ galaxies in a variety of environments, from void to group to cluster. Following Os10, we have divided the stars in each galaxy into accreted and in-situ, which allows us to examine the star formation histories for the galaxies as a function of stellar mass and environment. The qualitative similarities between our results and those of Os10, despite profound differences in numerical techniques and sub-grid physics, help substantiate the two-phase model of galaxy evolution for early type galaxies. The paper is divided into the following sections: \S\ref{sec:explainSim} details the simulation, the building of the merging histories and the tagging of stars as accreted or in-situ; \S\ref{sec:insitu_acc} and \S\ref{ssec:compareOser} describe the properties of the in-situ and accreted stars and compare our simulation results to those of Os10. \S\ref{sec:mergers} focuses on the merger histories and the types (major vs. minor) of mergers the galaxies in our simulation underwent and their effects on the final galaxy properties. For the simulation and throughout this work, we use the following \LCDM{} cosmology, consistent with WMAP-7 \citep{Komatsu11}: $\Omega_M = 0.28$, $\Omega_\lambda = 0.72$, $\Omega_b = 0.046$, $\sigma_8 = 0.82$, $n=92$, $H_0 = 100\h\kmps \Mpc^{-1}$, and $h=0.7$.
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\label{sec:summary} \figmhostz The two-phase picture of galaxy formation put forward by Os10 provides a useful framework for studies of galaxy evolution. {\rm We apply this framework to study the galaxy merger histories of more than $600$ simulated galaxies with $M_* > 10^{10}\ \Msun$. In this work, we corroborate the two-phase model for galaxy formation.} The first phase consists of a period of in-situ star formation, which occurs around $0.75<z<1$. This in-situ phase accounts for 60--90\% of the star formation in galaxies with stellar masses above $10^{10}\ \Msun$. The remainder of the stars are added in mergers. The peak of the merger activity occurs at $z\approx0.5$, after the majority of the in-situ star formation has taken place. {\rm Our results are in good agreement with those from Os10 and other simulations distinguishing between accreted and in-situ star formation in stellar halos, mainly around disk galaxies \citep[e.g.][]{Font2011, Abadi2006, Zolotov2009, Zolotov2010}. In particular, the differences in stellar age, metallicity, and spatial distribution shown in Figures \ref{fig:mGroup_ages}--\ref{fig:mGroup_profile} are in agreement with the differences found by \citet{Font2011} and \citet{Zolotov2009}.} Although the accretion occurs at late times, the accreted systems are old. The accreted stars are on average $2$ Gyr older than the stars formed in-situ. This age difference correlates with a median metallicity difference of $0.15$ dex. {\rm Since the accreted material comes from a variety of sources (i.e., a range of galaxy host masses), the distribution of stellar metallicities is larger for the accreted stars than for the in-situ stars.} However, the simulated galaxies do not conform to the mass-metallicity relation, suggesting that the stellar feedback model does not provide sufficient feedback, and that higher resolution is needed. Improved feedback prescriptions should lower the stellar metallicities for the less massive galaxies and the accreted systems, thus increasing the differences in metallicity between in-situ and accreted stars. This effect may be countered by a decrease in metallicity in the centers of galaxies, if higher resolution and better feedback prescription help shut off the excess late star formation currently seen in the simulations. Unlike in simulations, observations cannot easily distinguish between stars formed in-situ and those added by accretion. Nonetheless, we find indirect signs of the two phases of galaxy formation which are observable today. Because the accreted stars formed in small systems at early times, their stellar ages and metallicities are noticeably different from the stars formed in-situ. Furthermore, because the accreted stars preferentially reside in the outskirts of galaxies, we find strong gradients in the stellar populations of massive galaxies which have been observed beyond the half light radius of early-type galaxies \citep{Foster2009,Spolaor2010, Greene2012}. As explained above, our results are consistent with models of hierarchical formation of galaxies; large systems are assembled at late times from older, smaller building blocks. Figure \ref{fig:mhostz} is an illustration of the hierarchical growth. Here, we plot the stellar age against the stellar mass of the system (galaxy) in which each star was born, for the accreted and in-situ stars, separately. For a galaxy with no major mergers (central panel), the distribution of accreted and in-situ stars are well separated. On average, accreted stars form early in small systems, while in-situ star formation occurs later, when the host galaxy is more massive. On the other hand, the distributions of accreted and in-situ stars are very similar for a galaxy which underwent a major merger (left panel). This is unsurprising since two galaxies with roughly the same mass should have similar star formation histories. {\rm The overall trends in Figure \ref{fig:mhostz} suggest that galaxy stellar mass and stellar age are anti-correlated. However, the low mass galaxies shown in this figure formed early and were subsequently accreted, at which point they stopped forming stars. These galaxies are not representative of present-day low mass galaxies, but are rather the building blocks of more massive galaxies. Indeed, the top panel of Figure \ref{fig:m_agesISACC} shows stellar mass and stellar age are correlated for galaxies which survive to $z=0$, in qualitative agreement with observations. } Observations cannot directly reproduce Figure \ref{fig:mhostz}. However, there may be a close mapping from the relation between stellar ages and progenitor masses to the relation between metallicity and $\alpha$-element enhancement \citep{Thomas05, Johnston08}. Old stellar populations have higher alpha-to-iron abundance ratios than younger stellar systems, while massive galaxies typically have higher metallicities than less massive systems. {\rm Figure \ref{fig:mhostz} shows that accreted stars come from older, less massive systems, so we expect that accreted stars should be alpha-element enhanced compared to stars formed in-situ. This will yield radial gradients in abundance ratios in addition to the gradients in metallicity discussed above. } This stellar abundance space % is used to map the accretion history of the Galactic halo \citep{BlandHawthorn03,Robertson05,Font06a,Font06b,Johnston08,Zolotov2010}. Using simulations of a Milky Way analogue, \citet{Tissera2012} find that the accreted stars in the disc ($15\%$ of the disc stars) are alpha-enhanced and older than the stars formed in-situ, in broad agreement with our results shown in Figure \ref{fig:mhostz}. This type of analysis is being extended to galaxies besides the Milky Way by measuring $\alpha$-to-Fe ratios as a function of galaxy radius \citep{Kuntschner2010, Spolaor2010, Greene2012}. As detailed chemical and spatial information becomes available for the outer regions of more galaxies, it may be possible to differentiate between a minor merger and major merger driven galaxy formation history for individual galaxies. The results of this work agree broadly with the results of Os10 based on a set of $40$ galaxies simulated using GADGET, an SPH code, at higher mass resolution. We claim that most of the differences can be accounted for by the differences in star formation efficiency, feedback, and resolution between the simulations. The higher star formation efficiency in Os10 yields a higher accreted fraction because small halos have time to form stars during infall from the virial radius of larger halos. These small stellar systems tend to decrease the number-weighted mean merger ratio \citep{Oser11}. In our simulations, the lower resolution tends to increase the in-situ star formation rate by preventing gas clumps from fragmenting and forming stars before being accreted. These gas clumps survive to the centers of the galaxies, where they then form stars in-situ. The effects of resolution on the in-situ and accreted fractions will be checked in future AMR simulations run at higher resolution. Together with the differences in feedback, these effects account for the difference of a factor of two between the accreted fraction observed in the SPH and AMR simulations. Despite these differences, there is an excellent qualitative agreement between this work and the work in Os10. This shows that the two phase model for the formation of massive galaxies is robust and independent of the numerical methods and subgrid physics. {\rm Therefore, the simulations can be used to predict observable consequences of the two phases of galaxy formation. In this work, we show that the differences in age, metallicity, and metallicity dispersion between the accreted and in-situ stars yield radial gradients in these quantities for present-day massive galaxies. Currently, there is evidence for radial trends in metallicity in early-type galaxies \citep[cf.][]{Greene2012, Spolaor2010}, consistent with our predictions for two phase galaxy formation and further observational tests of this picture will be extremely useful. }
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1206.0295
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1206
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1206.2773_arXiv.txt
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We report the discovery of $16$ detached M-dwarf eclipsing binaries with $J<16$ mag and provide a detailed characterisation of three of them, using high-precision infrared light curves from the WFCAM Transit Survey (WTS). Such systems provide the most accurate and model-independent method for measuring the fundamental parameters of these poorly understood yet numerous stars, which currently lack sufficient observations to precisely calibrate stellar evolution models. We fully solve for the masses and radii of three of the systems, finding orbital periods in the range $1.5<P<4.9$ days, with masses spanning $0.35-0.50\msun$ and radii between $0.38-0.50\rsun$, with uncertainties of $\sim3.5-6.4\%$ in mass and $\sim2.7-5.5\%$ in radius. Close-companions in short-period binaries are expected to be tidally-locked into fast rotational velocities, resulting in high levels of magnetic activity. This is predicted to inflate their radii by inhibiting convective flow and increasing star spot coverage. The radii of the WTS systems are inflated above model predictions by $\sim3-12\%$, in agreement with the observed trend, despite an expected lower systematic contribution from star spots signals at infrared wavelengths. We searched for correlation between the orbital period and radius inflation by combining our results with all existing M-dwarf radius measurements of comparable precision, but we found no statistically significant evidence for a decrease in radius inflation for longer period, less active systems. Radius inflation continues to exists in non-synchronised systems indicating that the problem remains even for very low activity M-dwarfs. Resolving this issue is vital not only for understanding the most populous stars in the Universe, but also for characterising their planetary companions, which hold the best prospects for finding Earth-like planets in the traditional habitable zone.
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\label{sec:intro} M-dwarfs ($M_{\star}\lesssim0.6\msun$) constitute more than seventy per cent of the Galactic stellar population \citep{Hen97} and consequently, they influence a wide-range of astrophysical phenomena, from the total baryonic content of the universe, to the shape of the stellar initial mass function. Furthermore, they are fast becoming a key player in the hunt for Earth-like planets (e.g. \citealt{Nut08, Kopp09,Law11}). The lower masses and smaller radii of M-dwarfs mean that an Earth-like companion causes a deeper transit and induces a greater reflex motion in its host than it would do to a solar analogue, making it comparatively easier to detect Earths in the traditional habitable zones of cool stars. The inferred properties of exoplanet companions, such as their density, atmospheric structure and composition, currently depend on a precise knowledge of the fundamental properties of the host star, such as its mass, radius, luminosity and effective temperature at a given age. Yet, to date, no theoretical model of low-mass stellar evolution can accurately reproduce all of the observed properties of M-dwarfs \citep{Hil04,Lop05}, which leaves their planetary companions open to significant mischaracterisation. Indeed, the characterisation of the atmosphere of the super-Earth around the M-dwarf GJ 1214 seems to depend on the spot coverage of the host star \citep{Moo12}. Detached, double-lined, M-dwarf eclipsing binaries (MEBs) provide the most accurate and precise, model-independent means of measuring the fundamental properties of low-mass stars \citep{And91}, and the coevality of the component stars, coupled with the assumption that they have the same metallicity due to their shared natal environment, places stringent observational constraints on stellar evolution models. In the best cases, the uncertainties on the masses and radii measured using MEBs can be just $0.5\%$ \citep{Mor09,Kra11a}. However, since M-dwarfs are intrinsically faint, only a small number of MEBs have been characterised so far with suitable accuracy to calibrate low-mass stellar evolution models, and there are even fewer measurements below $\sim0.35\msun,$ where stellar atmospheres are thought to transport energy purely by convection \citep{Chab97}. More worryingly, existing observations show significant discrepancies with stellar models. The measured radii of M-dwarfs are inflated by $5-10\%$ compared to model estimates and their effective temperatures appear too cool by $3-5\%$ (see e.g. \citealt{Lop05, Rib06,Mor10, Tor10, Kra11a}). This anomaly has been known for some time but remains enigmatic. Bizarrely, the two discrepancies compensate each other in the mass-luminosity plane such that current stellar models can accurately reproduce the observed mass-luminosity relationship for M-dwarfs. Two different physical mechanisms have been suggested as the cause of this apparent radius inflation: i) metallicity \citep{Ber06,Lop07} and ii) magnetic activity \citep{Mul01,Rib06,Tor06,Chab07}. \citet{Ber06} and \citet{Lop07} used interferometrically-measured radii of single, low-mass stars to look for correlation between inflation and metallicity. Both studies found evidence that inactive, single stars with inflated radii corresponded to stars with higher metallicity, but this did not hold true for active, fast-rotating single stars and further studies could not confirm the result \citep{Dem09}. While metallicity may play a role in the scatter of effective temperatures for a given mass (the effective temperature depends on the bolometric luminosity which is a function of metallicity), it seems unlikely that it is the main culprit of radius inflation. The magnetic activity hypothesis is steered by the fact that the large majority of well-characterised MEBs are in short ($<2$ day) orbits. Such short period systems found in the field (i.e. old systems) are expected to be tidally-synchronised with circularised orbits \citep{Zah77}. The effect of tidal-locking is to increase magnetic activity and is a notion that is supported by observations of synchronous, rapid rotation rates in MEBs, a majority of circular orbits for MEBs, plus X-ray emission and H$\alpha$ emission from at least one of the components. It is hypothesised that increased magnetic activity affects the radius of the star in two ways. Firstly, it can inhibit the convective flow, thus the star must inflate and cool to maintain hydrostatic equilibrium. \citet{Chab07} modelled this as a change in the convective mixing length, finding that a reduced mixing length could account for the inflated radii of stars in the partially-radiative mass regime, but it had negligible effect on the predicted radii of stars in the fully-convective regime. However, \citet{Jac09} showed that the radii of young, single, active, fully-convective stars in the open cluster NGC 2516 could be inflated by up to $50\%$, based on radii derived using photometrically-measured rotation rates and spectroscopically-measured projected rotational velocities. This therefore suggests that inhibition of convective flow is not the only factor responsible for the radius anomaly. The second consequence of increased magnetic activity is a higher production of photospheric spots which has a two-fold effect: i) a loss of radiative efficiency at the surface, causing the star to inflate and ii) a systematic error in light curve solutions due to a loss of circular symmetry caused by a polar distribution of spots. \citet{Mor10} showed that these two effects could account for $\sim3\%$ and $0-6\%$ of the radius inflation, respectively, with any any remaining excess ($0-4\%$) produced by inhibition of convective efficiency. This however is only under certain generalisations, such as a $30\%$ spot coverage fraction and a concentration of the spot distribution at the pole. One would perhaps expect the systematic error induced by star spots to be wavelength dependant, such that radius measurements obtained at longer wavelength would be closer to model predictions. \citet{Kra11a} searched for correlation between the radius anomaly and the orbital periods of MEBs, to see if the data and the models converged at longer periods ($\sim3$ days) where the stellar activity is less aggravated by fast rotation speeds. They found tentative evidence to suggest that this is the case but it is currently confined to the realm of small statistics. Not long after their study, the MEarth project uncovered a 41-day, non-synchronised, non-circularised, inactive MEB with radius measurements still inflated on average by $\sim4\%$, despite a detailed attempt to account for spot-induced systematics \citep{Irw11}. They suggest that either a much larger spot coverage than the $30\%$ they assumed is required to explain the inflation, or perhaps that the equation of state for low-mass stars, despite substantial progress (see review by \citealt{Chab05}), is still inadequate. Clearly, a large sample of MEBs with a wide-range of orbital periods is key to defining the magnetic activity effect and understanding any further underlying physical issues for modelling the evolution of low-mass single stars. This in turn will remove many uncertainties in the properties of exoplanets with M-dwarf host stars. With that in mind, this paper presents the discovery of many new MEBs to emerge from the WFCAM Transit Survey, including a full characterisation to reasonable accuracy for three of the systems using 4-m class telescopes, despite their relatively faint magnitudes ($i=16.7-17.6$). In Section~\ref{sec:discovery}, we describe the WFCAM Transit Survey (WTS) and its observing strategy, and Section~\ref{sec:obs} provides additional details of the photometric and spectroscopic data we used to fully characterise three of the MEBs. In Section~\ref{sec:identify}, we outline how we identified the MEBs amongst the large catalogue of light curves in the WTS. Sections~\ref{sec:indices}-\ref{sec:RVs} present our analysis of all the available follow-up data used to characterise three of the MEBs including their system effective temperatures, metallicities, H$\alpha$ emission and surface gravities, via analysis of low-resolution spectroscopy, their size-ratio and orbital elements using multi-colour light curves, and their mass ratios using radial velocities obtained with intermediate-resolution spectra. These results are combined in Section~\ref{sec:absdim} to determine individual masses, radii, effective temperatures. We also calculate their space velocities and assess their membership to the Galactic thick and thin disks. Lastly, in Section~\ref{sec:discuss}, we discuss our results in the context of low-mass stellar evolution models and a mass-radius-period relationship, as suggested by \citet{Kra11a}.
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In this paper, we have presented a catalogue of $16$ new low-mass, detached eclipsing binaries that were discovered in the WFCAM Transit Survey. This is the first time dynamical measurements of M-dwarf EBs have been detected and measured primarily with infrared data. The survey light curves are of high quality, with a per epoch photometric precision of $3-5$ mmag for the brightest targets ($J\sim13$ mag), and a median RMS of $\lesssim1\%$ for $J\lesssim16$ mag. We have reported the characterisation of three of these new systems using follow-up spectroscopy from ground-based $2-4$ m class telescopes. The three systems ($i=16.7-17.6$ mag) have orbital periods in the range $1.5-4.9$ days, and span masses $0.35-0.50\msun$ and radii $0.38-0.50\rsun$, with uncertainties of $\sim 3.5-6.4\%$ in mass and $\sim 2.7-5.5\%$ in radius. Two of the systems may be associated with the young-old disk population as defined by \citet{Leg92} but our metallicity estimates from low-resolution spectra do not confirm a non-solar metallicity. The radii of some of the stars in these new systems are significantly inflated above model predictions ($\sim3-12\%$). We analysed their radius anomalies along with literature data as a function of the orbital period (a proxy for activity). Our error-weighted statistical analysis revealed marginal evidence for greater radius inflation in very short orbital periods $<1$ day, but neither a linear nor exponentially decay model produced a significant fit to the data. As a result, we found no statistically significant evidence for a correlation between the radius anomaly and orbital period, but we are limited by the small sample of precise mass and radius measurements for low-mass stars. However, it is clear that radius inflation exists even at longer orbital periods in systems with low (or undetectable) levels of magnetic activity. A robust calibration of the effect of magnetic fields on the radii of M-dwarfs is therefore a key component in our understanding of these stars. Furthermore, it is a limiting factor in characterising the planetary companions of M-dwarfs, which are arguably our best targets in the search for habitable worlds and the study of other Earth-like atmospheres. More measurements of the masses, radii and orbital periods of M-dwarf eclipsing binaries, spanning both the fully convective regime and partially convective mass regime, for active and non-active stars, across a range of periods extending beyond 5 days, are necessary to provide stringent observational constraints on the role of activity in the evolution of single low-mass stars. However, the influence of spots on the accuracy to which we can determine the radii from light curves will continue to impede these efforts, even in the most careful of cases (see e.g. \citealt{Mor10,Irw11}). This work has studied only one third of the M-dwarfs in the WFCAM Transit Survey. Observations are on-going and we expect our catalogue of M-dwarf eclipsing binaries to increase. This forms part of the legacy of the WTS and will provide the low-mass star community with high-quality MEB light curves. Furthermore, the longer the WTS runs, the more sensitive we become to valuable long-period, low-mass eclipsing binaries. These contributions plus other M-dwarf surveys, such as MEarth and PTF/M-dwarfs, will ultimately provide the observational calibration needed to anchor the theory of low-mass stellar evolution.
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1206.2773
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1206
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1206.3363_arXiv.txt
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The atmospheres of hot Jupiters and other strongly-forced exoplanets are susceptible to a thermal instability in the presence of ohmic dissipation, weak magnetic drag and strong winds. The instability occurs in radiatively-dominated atmospheric regions when the ohmic dissipation rate increases with temperature faster than the radiative (cooling) rate. The instability domain covers a specific range of atmospheric pressures and temperatures, typically $P \sim 3$--$300$ mbar and $T \sim 1500-2500$~K for hot Jupiters, which makes it a candidate mechanism to explain the dayside thermal ``inversions'' inferred for a number of such exoplanets. The instability is suppressed by high levels of non-thermal photoionization, in possible agreement with a recently established observational trend. We highlight several shortcomings of the instability treatment presented here. Understanding the emergence and outcome of the instability, which should result in locally hotter atmospheres with stronger levels of drag, will require global non-linear atmospheric models with adequate MHD prescriptions.
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A variety of gaseous exoplanets with strongly-forced atmospheres have been discovered by astronomers (e.g., Charbonneau 2009). These exoplanets, exemplified by the hot Jupiter class, receive extreme levels of irradiation from their stellar host and likely experience permanent day- and night-side forcing conditions, from being tidally-locked on very compact orbits. Observationally, two of the most interesting trends emerging from studies of well-characterized, transiting hot Jupiters have been the tendency for many such planets to exhibit radius inflation, well above expectations from standard planetary cooling models, and the inference that thermal excesses (inversions) are present on the dayside of some of the strongly-irradiated planets (see Deming \& Seager 2009; Baraffe et al. 2010; Burrows \& Orton 2010; Winn 2010 for reviews). It has recently been proposed that hot, strongly-forced exoplanet atmospheres are the site of significant magnetic induction when fast, weakly-ionized atmospheric winds cross the pre-existing planetary magnetic field. This induction takes the form of magnetic drag acting to brake atmospheric winds and associated ohmic heating in the planetary atmosphere and interior (Batygin \& Stevenson 2010; Perna, Menou \& Rauscher 2010a,b; Menou 2012). While other mechanisms have been proposed, ohmic heating is currently one of the leading contenders to explain the inflated radii of hot Jupiters (Batygin et al. 2011; Laughlin et al. 2011; Menou 2012; Wu \& Lithwick 2012). In this letter, we explore the possibility that thermal inversions in hot Jupiter atmospheres are not caused by extra absorption of stellar light at altitude, as they have traditionally been interpreted (Hubeny et al. 2003; Fortney et al. 2006, 2008; Burrows et al. 2008; Madhusudhan \& Seager 2010), but instead result from the same induction mechanism that may explain radius inflation. In this alternative interpretation, thermal inversions have their origin in a thermo-resistive instability that affects radiatively-dominated atmospheric regions under specific conditions of weak magnetic drag, strong ohmic dissipation and fast winds.
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Our argument for the existence of a thermo-resistive instability in the atmospheres of hot, strongly-forced exoplanets relies on a number of assumptions which are not all well justified. Most importantly, the induction formalism used here and in other studies of magnetic effects in hot Jupiter atmospheres and interiors assumes axisymmetry and steadiness. We already emphasized in Perna et al. (2010a) the clearly non-axisymmetric properties of strongly-forced exoplanet atmospheres, a situation that would only be aggravated by the onset of a localized thermal instability (e.g., on the dayside). Our study of an MHD-based instability with a steady-state induction equation is another important shortcoming. Dimensional analysis of the time-dependent axisymmetric induction equation (see Eq~[1] in Menou 2012) suggests that induced currents would lag changes in temperature and resistivity by a resistive diffusion time, $\tau_{\rm diff} \sim L^2/\eta$. Adopting a pressure scaleheight for the characteristic lengthscale $L$, we estimate that $\tau_{\rm diff} < \tau_{\rm rad}$ for the green and red instability domains in Fig.~\ref{fig:one}, which suggests that current-adjustment delays will not strongly affect the instability development. At high temperatures and low densities, however, around the blue instability domain and at higher temperatures in Fig.~\ref{fig:one}, resistivities become small and $\tau_{\rm diff} > \tau_{\rm rad}$, which may invalidate the implicit assumption of instantaneous current adjustment made in our steady-state evaluation of the ohmic dissipation term in Eq.~(\ref{eq:ohm}). A careful consideration of this issue likely requires a full MHD treatment and we shall simply note here that time-dependent current adjustments may impact the thermo-resistive instability of the hottest exoplanet atmospheres beyond the simple treatment adopted here. Despite such limitations, it is tempting to associate the thermo-resistive instability mechanism identified here with the dayside thermal inversions inferred for a number of hot Jupiters. As shown in Fig.~\ref{fig:one}, the instability domain is restricted to a specific range of pressures and temperatures. For magnetic field strengths $B \sim 3$-$30$~G, instability domains can match well the upper atmospheric dayside conditions of moderately hot exoplanets such as HD209458b, with low enough nightside temperatures $T \lsim 1100$~K to suppress the instability. On cooler hot Jupiters such as HD189733b, dayside conditions would barely achieve instability (see $T$-$P$ profiles in Showman et al. 2008 and Rauscher \& Menou 2012). By contrast, very hot exoplanets with $T_{\rm eff} > 1500$-$2000$~K may be too hot on their daysides for instability, and even if their nightsides were to meet the temperature requirements, the instability may be suppressed because of the very effective wind drag exerted on the dayside or possibly strong deviations from radiative equilibrium on the nightside. In principle, some hot exoplanets may be preferentially unstable near their limb, which will be an interesting issue to explore with improved instability models. The outcome of the instability is difficult to anticipate beyond qualitative expectations. Unstable regions should reach temperatures high enough for saturation of the instability in the strong drag regime, which can amount to temperature excesses of several hundreds Kelvin according to the width of domains shown in Figs.~\ref{fig:one}--\ref{fig:three}. Perhaps more importantly, the radiative response of the vertically-coupled atmospheric layers, together with the horizontal coupling caused by locally modified thermal and drag conditions, will result in a very non-linear response that is best studied with global models. While short instability growth times may prove numerically challenging, global circulation models with adequate MHD treatments offer a promising avenue for progress. They would help clarify the global energetics of the instability, which is ultimately powered by the same differential thermal forcing as the global circulation itself.
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1206.3363
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1206
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1206.1685_arXiv.txt
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In this paper, we show that if passive fluctuations are considered, primordial black holes (PBHs) can be easily produced in the framework of single-field, slow-roll inflation models. The formation of PBHs is due to the blue spectrum of passive fluctuations and an enhancement of the spectral range which exits horizon near the end of inflation. Therefore the PBHs are light with masses $\lesssim 10^{15}\mbox{g}$ depending on the number of e-folds when the scale of our observable universe leaves horizon. These PBHs are likely to have evaporated and cannot be a candidate for dark matter but they may still affect the early universe.
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Inflation \cite{Lyth:2009zz} is becoming a standard model for the very early universe. The inflationary scenario, in which the present universe is only a small local patch of a causally connected region at early times which underwent an exponential expansion driven by the inflaton potential, is generally accepted for explaining the observed spatially flat and homogeneous universe. In addition, its quantum fluctuations during inflation give rise to primordial Gaussian matter density fluctuations with a nearly scale-invariant power spectrum, which is consistent with recent astrophysical and cosmological observations such as structure formation and cosmic microwave background anisotropies~\cite{texas04}. Although the simplest single-field, slow-roll inflation model works well, some basic questions have yet to be answered. What is the origin of the inflaton potential? Do classical matter density imhomogenities that we observe today genuinely come from quantum fluctuations of the inflaton? Are the observed matter density fluctuations truly Gaussian? How robust are the predictions for a subdominant contribution of tensor modes to the metric fluctuations, a slightly broken scale invariance, and a negligible running spectral index of the power spectrum? Future cosmic microwave background measurements and mega-scale mappings of the large scale structure will definitely answer some of these questions or perhaps pose a challenge to the standard inflation scenario. There has been a lot of studies on inflationary models that go beyond the simplest single-field, slow-roll inflation. A class of models has considered a new source for generating inflaton fluctuations (so-called passive density fluctuations) during inflation through a direct or gravitational coupling between the inflaton and other quantum fields. This leads to very interesting results such as the so-called warm inflation~\cite{berera}, the suppression of large-scale density fluctuations~\cite{wu}, possible constraints on the duration of inflationary expansion~\cite{wood}. the bursts of particle production that result in infra-red cascading~\cite{barnaby}, the trapped inflation in which the inflaton rolls slowly down a steep potential by dumping its kinetic energy into light particles at the trapping points along the inflaton trajectory~\cite{green,lee}, and electromagnetic dissipation in natural inflation~\cite{sorbo,Barnaby:2010vf, Barnaby:2011vw,Meerburg:2012id}. In all of these papers, essentially, the generation of passive density fluctuations is originated from quantum fluctuations in the back reaction of the couplings to the inflaton perturbation. The nature of passive fluctuations is usually non-Gaussian and non-scale-invariant. A particular feature is that the power spectrum of the passive fluctuations can be very blue~\cite{wood,lee,sorbo,Barnaby:2010vf, Barnaby:2011vw,Meerburg:2012id}. These passive fluctuations cannot dominate the primordial density perturbation at large scales as confirmed by cosmological observations such as cosmic microwave background (CMB) anisotropies and the formation of large scale structures. However, depending on individual models, their significant contribution to the non-Gaussianity is still possible. This will be tested soon in the Planck CMB mission and in future large-scale-structure surveys. In this paper, we point out that the passive fluctuations with a blue spectrum can dominate the primordial density perturbation in the very small scales, seeding the formation of primordial black holes (PBHs) in the radiation-domination era after inflation.
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\label{con} A lot of effort has been put in studying the passive fluctuations during inflation and their imprints on the CMB and in particular the non-Gaussian features. However, in this paper we have proposed a novel mechanism that passive fluctuations can produce PBHs with scales exiting horizon near the end of inflation. This would result in light PBHs with mass $M \lesssim 10^{15} \mbox{g}$ that may have interesting cosmological consequences in the early universe. We have discussed three models in which passive fluctuations can easily generate PBHs, noting that the characteristic power spectrum is rather generic, namely, the spectrum is extremely blue at small scales.
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1206.1685
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1206
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1206.6770_arXiv.txt
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{\referee{The} \textit{Swift}-{XRT} observations of the early X-ray afterglow of GRBs show that it usually begins with a steep decay phase. } {A possible origin of this early steep decay is the high latitude emission that subsists when the on-axis emission of the last dissipating regions in the relativistic outflow has been switched-off. We \referee{wish to establish which of various} models of the prompt emission are compatible with this interpretation. } {We successively consider internal shocks, photospheric emission\referee{,} and magnetic reconnection and obtain the typical decay \referee{timescale} at the end of the prompt phase in each case. } {Only internal shocks naturally predict a decay \referee{timescale} comparable to the burst duration, as required to explain {XRT} observations \referee{in terms of} high latitude emission. \referee{The decay timescale of the high latitude emission} is much too short in photospheric models and the observed decay must then correspond to an effective and generic behavior of the central engine. Reconnection models require some \textit{ad hoc} assumptions to agree with the data, which will have to be validated when a better description of the reconnection process becomes available. } {}
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Thanks to its precise localization capabilities followed by rapid slewing, the \textit{Swift} satellite \citep{gehrels_2004} can quickly -- typically in less than two minutes after a Gamma-ray burst (GRB) trigger -- \referee{repoint} its X-Ray Telescope ({XRT}, \citealt{burrows_2005}) toward the source. This achievement has % \referee{helped to fill} the gap \referee{in observations} between the prompt and late afterglow emissions and revealed the complexity of the early X-ray afterglow \citep{tagliaferri_2005, nousek_2006, obrien_2006}. Despite this complexity, the early steep decay that ends the prompt emission appears to be a generic behavior, common to most long GRBs. During this phase\referee{,} the flux decays with a temporal index $\alpha \simeq 3-5$ (with $F_{\nu} \propto t^{-\alpha}$). It lasts for a typical duration $t_{\rm ESD}\sim 10^{2}-10^{3}$ s and is usually followed by the shallow or normal decay phases. The rapid gamma-ray light curve variability suggests that prompt emission has to be produced by internal mechanisms \citep{sari_1997}, while the later shallow and normal decay phases are usually attributed to % \referee{deceleration} by the external medium (e.g. \citealt{meszaros_1997, sari_1998, rees_1998}). As the backward extrapolation of the early steep decay \referee{connects reasonably} to the end of the prompt emission, it is often interpreted as its tail. It has moreover been shown that it is too steep to result from the interaction with the external medium (see e.g. \citealt{lazzati_2006}). One of the most discussed and natural scenario to account for the early steep decay had been described by \citet{kumar_2000} before the launch of \textit{Swift}. It explains this phase by the residual off-axis emission -- or high latitude emission -- that becomes visible when the on-axis prompt activity switches off. This scenario gives simple predictions and several studies have been dedicated to check whether it agrees with observations (see e.g. \citealt{liang_2006, butler_2007, zhang_2007, qin_2008, barniol_2009}). \referee{ On the basis of a realistic multiple pulse fitting approach proposed by \citet{genet_2009}, \citet{willingale_2010} confirmed} that the observed temporal slope of the early steep decay can be well explained in this context\referee{,} while the accompanying spectral softening is (at least qualitatively) also reproduced. In this letter, we first discuss in \sect{section_constraints} some constraints implied by the high latitude scenario on the typical radius $R_{\gamma}$ where the prompt emission ends. We then investigate in \sect{section_comparison} if they are consistent with the predictions of the most discussed models for the prompt phase: internal shocks, \referee{Comptonized} photospheric emission\referee{,} and magnetic reconnection. We summarize our results in \sect{section_conclusion}, which also \referee{presents our conclusions.} \vspace{-2ex}
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\label{section_conclusion} The results of this letter emphasize the importance of the steep decay phase revealed by \textit{Swift} observations of the early X-ray afterglow. An attractive way \referee{of explaining} this phase is to suppose that it is produced by the high latitude emission of the last contributing shells. We have shown that this assumption leads to strong constraints on the radius $R_{\gamma}$ where the prompt emission ends. In particular, the associated \referee{timescale}, $\tau_{\rm HLE}\sim R_{\gamma}/2c\,\Gamma^2$ must be comparable to the burst duration $t_{\rm burst}$ to guarantee that the high latitude contribution is correctly connected to the end of the prompt light curve. We have then checked \referee{whether} these constraints are satisfied by different models for the prompt phase, namely internal shocks, \referee{Comptonized} photospheric emission\referee{,} and magnetic reconnection. \begin{itemize} \item Internal shocks naturally fulfill the condition $\tau_{\rm HLE}\sim t_{\rm burst}$ even in highly variable bursts, as the radius of the last shocks is governed by the longest variability timescale\referee{,} which is \referee{on} the order of $t_\mathrm{burst}$. \item Conversely\referee{,} in photospheric models $\tau_{\rm HLE}\ll t_{\rm burst}$, since the photospheric radius is typically several orders of magnitude smaller than the radius of internal shocks. The only way to produce the observed decay is then to suppose that it corresponds to an effective behavior of the central engine, which moreover should be common to most GRBs. \item In magnetic reconnection models, satisfying the constraints on $R_{\gamma}$ might be possible but is not naturally expected. It still requires the ad hoc assumption that $R_{\gamma} \simeq \Gamma^2 c t_{\rm burst}$, which will have to be justified when a better description of the reconnection process becomes available. \end{itemize} \begin{figure} \begin{center} \includegraphics[scale=0.28]{fig_2a.eps}\\ \vspace*{-2ex} \includegraphics[scale=0.28]{fig_2b.eps} \end{center} \vspace*{-5ex} \caption{\textbf{Early steep decay from high latitude emission in the internal shock framework: a multi-pulse burst.} \textit{Top:} initial distribution of the Lorentz factor; \textit{Bottom:} bolometric light curve \referee{on a} logarithmic scale (solid line). The contribution of each individual pulse is plotted \referee{as a} dashed line. } \label{fig_is_complex_3pulses} \end{figure} \vspace*{-2ex}
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1206.6770
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1206
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1206.4319_arXiv.txt
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The present observational understanding of the evolution of the mass accretion rates (\macc) in pre-main sequence stars is limited by the lack of accurate measurements of \macc\ over homogeneous and large statistical samples of young stars. Such observational effort is needed to properly constrain the theory of star formation and disk evolution. Based on HST/WFPC2 observations, we present a study of \macc\ for a sample of $\sim$ 700 sources in the Orion Nebula Cluster, ranging from the Hydrogen-burning limit to $M_*\sim 2 M_\odot$. We derive \macc \ from both the $U$-band excess and the $H\alpha$ luminosity ($L_{H\alpha}$), after determining empirically both the shape of the typical accretion spectrum across the Balmer jump and the relation between the accretion luminosity ($L_{\rm acc}$) and $L_{H\alpha}$, that is $L_{\rm acc}/L_\odot = (1.31 \pm 0.03) \cdot L_{H\alpha}/L_\odot$ + (2.63$\pm$ 0.13). Given our large statistical sample, we are able to accurately investigate relations between \macc\ and the parameters of the central star such as mass and age. We clearly find \macc\ to increase with stellar mass, and decrease over evolutionary time, but we also find strong evidence that the decay of \macc\ with stellar age occurs over longer timescales for more massive PMS stars. Our best fit relation between these parameters is given by: $\log$(\macc/$M_\odot\cdot$yr)=(-5.12 $\pm$ 0.86) -(0.46 $\pm$ 0.13) $\cdot \log (t$/yr) -(5.75 $\pm$ 1.47)$\cdot \log (M_*/M_\odot)$ + (1.17 $\pm$ 0.23)$\cdot \log (t/$yr$) \cdot \log (M_*/M_\odot)$. These results also suggest that the similarity solution model could be revised for sources with $M_* \gtrsim 0.5 M_\odot$. Finally, we do not find a clear trend indicating environmental effects on the accretion properties of the sources.
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During the Pre-Main-Sequence (PMS) phase of stellar evolution, the interaction between the forming star and the surrounding disk is regulated by the accretion of disk material along the field lines of the stellar magnetosphere \citep{H98}. The gravitational energy released as the material falls along accretion columns and hits the stellar surface creates a characteristic shock spectrum \citep{Calvet}, with excess emission especially strong in the Balmer continuum and recombination lines \citep{Gullbring, Calvet}. The relative accretion luminosity (\lacc) can be measured with spectroscopic \citep{Valenti93,Herczeg08} or photometric \citep{Gullbring,Robberto,DeMarchi} methods. The mass accretion rate, \macc, is then estimated through the relation \citep{H98}: \begin{equation} \dot{M}_{\rm acc} = \frac{L_{\rm acc} R_*}{0.8 G M_*}, \label{Macc_eq} \end{equation} where $R_*$ and $M_*$ are the radius and the mass of the star, respectively, and the factor $0.8$ accounts for the assumption that the infall originates at a magnetospheric radius $R_{\rm m} = $5 \citep{Shu94}. The mass accretion rate generally decreases over time during the first few Myr of PMS stellar evolution, as the circumstellar disks disperse their gaseous content on a timescale of $\sim$ 3-5 Myr \citep{Haisch01,Dahm05,Fedele}. Moreover, \macc\ is expected to scale with stellar mass. The evolution of \macc \ vs. time as a function of the stellar and disk mass represents a key aspect of PMS evolution and planet formation. \citet{H98} showed that while typical \macc \ values, at any given stellar age, can differ star by star of up to about two orders of magnitudes, on average they decrease exponentially with the stellar age $t$, i.e. \macc $\propto t^{-\eta}$ with $\eta \sim 1.5$. Several studies, targeting different young stellar clusters, have confirmed this general trend (e.g. \citealt{Robberto}, \citealt{Aurora10}), but the uncertainties in the age estimate \citep{Hartm97,Baraffe10} and the scarcity of rich and homogeneous samples limit the accurate assessment of this dependence. For what concerns the stellar mass dependence of \macc, \citet{Muzerolle03} found that $\dot{M}_{\rm acc}\propto M_*^b$ with $b=2$, although recent studies \citep[e.g.,][]{Rigliaco} suggest smaller values of $b$. Due to the large spread in the \macc\ values for a given $M_*$ (up to two orders of magnitudes, \citealt{Rigliaco}), this second relation has also never been accurately constrained. Although observational uncertainties and intrinsic variability will always contribute in scattering the measured mass accretion rates, accurate measurements of \macc \ on a large sample of PMS stars may allow us to assess more precisely the value of the two power law exponents ($\eta$ and $b$). This is a critical step, as they can be tied to the theoretical model for the disk evolution and structure, providing unique constraints to the initial conditions of planet formation (e.g. \citealt{Alexander06}, \citealt{Lodato}). As the nearest ($d\sim 414 \pm 7$ pc, \citealt{distanzaONC}) site of massive star formation, the Orion Nebula Cluster (ONC) provides a standard benchmark for star formation studies. The ONC population has been studied in depth, and determination of the individual stellar parameters are available for a large fraction of members. In particular \citet{H97}, \citet{DaRio} and \citet{DaRio12} have used both spectroscopic and photometric techniques to derive the spectral type of more than 1700 ONC sources. The corresponding initial mass function, ranging from the Brown Dwarfs (BD) regime to the O6 star $\theta^1$Ori-C, peaks at 0.2-0.3 $M_{\odot}$ \citep{DaRio12}, while the cluster mean age is $\sim$ 2.2 Myr, with evidence for an age spread of the order of $\sim$ 2 Myr \citep{Maddi}. Given the wealth of information on the individual cluster members, the ONC is ideally suited to conduct an extensive study of the mass accretion process in PMS stars. In this paper we present the results of such a study based on the HST/WFPC2 survey of the Orion Nebula Cluster (GO 10246, P.I. M. Robberto). Both U-band and $H\alpha$ data are used to estimate \macc \ for $\sim$ 700 sources and to analyze the relations between this parameter and the main stellar parameters (age, $M_*$). In Sec.~\ref{observations_sec} we illustrate the observation, data reduction and analysis of our WFPC2 photometry. Sec.~\ref{method} presents our derived modeling of the photospheric colors for our sources and for the typical accretion spectrum, and the methods used to obtain the \lacc \ based on our data, both using our UBI Diagram method and from the $H\alpha$ photometry, while in \ref{results} we present the derived quantities (\lacc, \macc) obtained both from the measurement of the $U$-band or from $L_{H\alpha}$; in Sec.~\ref{analysis} we study the evolution of these parameters as a function the main stellar parameter. Finally, in Sec.~\ref{conclusion} we summarize our conclusions.
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\label{conclusion} We have presented a study of mass accretion rates (\macc) in the Orion Nebula Cluster, for an unprecedentedly large ($\sim$ 700 stars) sample of PMS stars. This allowed us to perform a thorough statistical analysis of the dependence of this quantity on the central stellar parameters, and investigate this phase of the stellar mass build-up and disc evolution. We based our study on HST/WFPC2 photometric data over a large field of view, and derived \macc\ using two different accretion tracers: the ($U-B$) excess and the $L_{H\alpha}$. We study the systematic dependence of \macc\ with respect the age of the sources, their mass and position within the region. In particular, the \macc\ is found to vary with age and mass altogether. Our final relation between all these quantities is given by Equation \ref{fit3D_all} and in Table~\ref{param_fit_table}, and our best fit relation using all the sources in the sample is given by: $\log$(\macc/$M_\odot\cdot$yr)=(-5.12 $\pm$ 0.86) -(0.46 $\pm$ 0.13) $\cdot \log (t$/yr) -(5.75 $\pm$ 1.47)$\cdot \log (M_*/M_\odot)$ + (1.17 $\pm$ 0.23)$\cdot \log (t/$yr$) \cdot \log (M_*/M_\odot)$. We clearly find that the \macc\ increases with stellar mass, and decreases over evolutionary time. Interestingly, we also find evidence that for more massive stars the decay of \macc\ with time is much slower than for lower stellar masses. Similarly, for older stars, the dependence of \macc\ with $M_*$ appears significantly steeper. This might imply that these objects are not in the asymptotic regime (i.e. when $t\gg t_\nu$, where $\dot{M}_{acc}\propto t^{-\eta}$), or that the hypothesis of a simple dependence of the viscosity $\nu\propto R^\gamma$ is probably not compatible with our observations. In particular, we suggest significant discrepancies of our results with respect to the self-similar parametrization for sources with masses higher than $\sim 0.5 M_\odot$.
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1206.4319
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1206
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1206.1827_arXiv.txt
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We have obtained $Spitzer$ IRS spectra and MIPS 24, 70, and 160 $\mu$m photometry for a volume-limited sample of 22 SDSS-selected Low-ionization Broad Absorption Line QSOs (LoBALs) at $0.5 < z < 0.6$. By comparing their mid-IR spectral properties and far-IR SEDs with those of a control sample of 35 non-LoBALs matched in $M_i$, we investigate the differences between the two populations in terms of their infrared emission and star formation activity. Twenty five percent of the LoBALs show PAH features and 45\% have weak 9.7$\mu$m silicate dust emission. We model the SEDs and decouple the AGN and starburst contributions to the far-infrared luminosity in LoBALs and in non-LoBALs. Their median total, starburst, and AGN infrared luminosities are comparable. Twenty percent (but no more than 60\%) of the LoBALs and 26\% of the non-LoBALs are ultra-luminous infrared galaxies (ULIRGs; L$_{IR}>10^{12}L_{\odot}$). We estimate star formation rates (SFRs) corrected for the AGN contribution to the FIR flux and find that LoBALs have comparable levels of star formation activity to non-LoBALs when considering the entire samples. However, the SFRs of the IR-luminous LoBALs are 80\% higher than those of their counterparts in the control sample. The median contribution of star formation to the total far-infrared flux in LoBALs and in non-LoBALs is estimated to be 40-50\%, in agreement with previous results for PG QSOs. Overall, our results show that there is no strong evidence from the mid- and far-IR properties that LoBALs are drawn from a different parent population than non-LoBALs.
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Supermassive black holes are found to be ubiquitously present at the centers of galaxies with bulges and several relationships between the mass of the black hole and properties of the spheroid strongly suggest co-evolution of the two \citep[\eg][]{Kormendy1995, Magorrian1998, Laor1998, Gebhardt2000, Ferrarese2000, Kormendy2001, McLure2002, Tremaine2002}. The mechanisms via which the galaxy and black hole regulate each other's growth are still unknown. Various types of outflows, such as supernova winds and AGN-driven winds, have been invoked as the plausible feedback processes responsible for quenching the star formation in the host and clearing the gas from the nuclear region, and, thus, halting the accretion onto the black hole and limiting its mass \citep[\eg][]{DiMatteo2005}. However, observational evidence of the extent of their influence is still sparse and uncertain (e.g., \citealt{Moe2009}; \citealt{Bautista2010}; \citealt{Dunn2010}). Observations of young, recently fueled QSOs are the key to testing this evolutionary model. Particular attention has been paid to studying the properties of ultraluminous infrared galaxies (ULIRGs; $L_{IR}>10^{12}L_\odot$) since they are believed to be powered by both AGN and starbursts, although starbursts are generally believed to be responsible for the bulk of the power output (\citealt{Sanders1988}; for review see \citealt{Sanders1996}). The connection between AGN and ULIRGs is suspected by the fact that they are some of the most luminous sources in the universe with comparable luminosities of $L_{bol}> 10^{12}L_\odot$. In addition, they are both associated with strong infrared emission from dust \citep[\eg][]{Haas2003}. The morphologies and dynamics of ULIRGs indicate that these galaxies are associated with galaxy mergers \citep{Armus1987, Sanders1988, Murphy1996, Veilleux2002, Dasyra2006}. Similarly, many QSO hosts at $z<0.4$ show signs of interaction, even some of those that had previously been classified as ellipticals \citep{Canalizo2001, Canalizo2007, Bennert2008}. If starbursts, ULIRGs, and AGN are connected in an evolutionary sequence which was initiated by a galaxy interaction, observations of the transition stages of this process are necessary to better understand this connection \citep[see \eg][]{Hopkins2007, Hopkins2008a, Hopkins2008b}. BAL QSOs are promising candidates for newly emerging optical QSOs. BAL QSOs are a subclass of QSOs characterized by broad absorption troughs of UV resonance lines, blueshifted relative to the QSO's rest frame, which are indicative of gas outflows with speeds of up to 0.2$c$ \citep{Foltz1983}. BALs were rigorously defined by \citet{Weymann1991} to include only objects with broad absorption lines wider than 2000 km s$^{-1}$, blueshifted past the first 3000 km s$^{-1}$; however, some studies more recently have been more inclusive of the wide range of absorption observed and have relaxed that criterion to lower limits on the absorption width of 1000 km s$^{-1}$ \citep[\eg][]{Trump2006}. Hydrodynamic models show that AGNs are capable of launching such high velocity winds \citep{Murray1995, Proga2000, Gallagher2012}. Based on the material producing BAL troughs, there are at least three subclasses of BAL QSOs. The high-ionization BAL QSOs (HiBALs) are identified via the broad absorption from \ion{C}{4} $\lambda$1549, but they might have absorption from other high-ionization species such as Ly$\alpha$, \ion{N}{5} $\lambda$1240, and \ion{Si}{4} $\lambda$1394 \citep{Hall2002}. The low-ionization BAL QSOs (LoBALs), in addition to the lines present in HiBALs, feature absorption lines from \ion{Mg}{2} $\lambda$$\lambda$ 2796,2803, \ion{Al}{3}, and \ion{Al}{2}. A very small fraction of LoBALs, called FeLoBALs, show absorption in the rest-frame UV from metastable, excited states of \ion{Fe}{2} \citep{Hazard1987}. It is not well understood why only 10\%$-$30\% of the optically selected QSOs have BALs \citep{Tolea2002, Hewett2003, Trump2006, Gibson2009}, and only about one tenth of these are LoBALs \citep{Reichard2003a}. Due to the highly obscured nature and much redder continua of these objects, optical identification omits a large fraction of BALs. Therefore, although LoBALs are observed in only 1-3\% of all optically-selected QSOs, they comprise a much higher fraction of the QSOs selected at longer wavelengths \citep{Urrutia2009, Dai2010}. \citet{Allen2011} estimate that the intrinsic fraction of BAL QSOs can be as high as $\sim$40\% when the spectroscopic incompleteness and bias against selecting BAL QSOs in the SDSS are taken into account. Hence, BAL QSOs may only be rare in optically-selected samples. Models attempting to explain their occurrence need to account for their obscured nature. Currently there are two competing interpretations of the BAL phenomena: orientation and evolution. On one hand, BAL and non-BAL QSOs are thought to derive from the same parent population because of the remarkable similarity in their SEDs \citep{Weymann1991, Gallagher2007}. Nonetheless, QSO continua appear to be increasingly reddened in a sequence going from non-BALs to HiBALs to LoBALs \citep{Reichard2003b, Richards2003}. This finding inspired efforts to explain the low occurrence of BAL QSOs within the framework of the AGN unification model (Antonucci 1993), suggesting that, due to orientation effects, BALs are seen in classical QSOs only when viewed along a narrow range of lines of sight passing through the accretion disk wind. In this picture, high column density accretion disk winds of ionized gas are driven via resonance line absorption \citep{Murray1995, Murray1998, Elvis2000}. This model explains the low occurrence of BALs as a natural consequence of the fact that BALs are observed only at a small range of viewing angles. Although BALs are predominantly radio quiet sources \citep{Stocke1984, Stocke1992}, radio observations of the few radio-detected BALs provide a test to the orientation of the BAL wind with respect to the radio jet. Radio-detected BALs are observed at a wide range of inclinations \citep{Becker2000, Gregg2000, Brotherton2006, Montenegro-Montes2008, DiPompeo2010} suggesting that the occurrence of BALs is not a simple orientation effect \citep[e.g.,][]{DiPompeo2012}. Currently it is not clear whether or not radio-loud (RL) and radio-quiet (RQ) QSOs arise from the same parent population, so it is certainly possible that RL and RQ BAL QSOs are different classes. An alternative model proposes that BAL QSOs are young QSOs caught during a short-lived phase in their evolution when powerful QSO-driven winds are blowing away a dusty obscuring cocoon \citep[\eg][and references therein]{Hazard1984, Voit1993, Hall2002}. This model appears to be particularly applicable to LoBALs since these objects are suspected to be young or recently refueled QSOs \citep{Boroson1992, Lipari1994} and might be exclusively associated with mergers \citep{Canalizo2001}. Observations by \citet{Canalizo2002} of the only four known LoBALs at $z < 0.4$ at the time showed that: (1) they are ULIRGs; (2) they have a small range of far-IR colors, intermediate between those characteristic of ULIRGs and QSOs; (3) their host galaxies show signs of strong tidal interactions, resulting from major mergers; (4) spectra of their hosts show unambiguous interaction-induced star formation with post-starburst ages $\leq$ 100 Myr. Similarly, studies of FeLoBALs, both at low \citep{Farrah2005} and high redshifts \citep{Farrah2007, Farrah2010}, suggest that they are associated with extremely star-forming ULIRGs. Most recent hydrodynamic simulations of major galaxy mergers by \citet{DeBuhr2011} show that an AGN-driven BAL wind with an initial velocity $\sim$ 10000 km s$^{-1}$ would lead to a galaxy-scale outflow with velocity $\sim$ 1000 km s$^{-1}$, capable of unbinding 10$-$40\% of the initial gas of the two merging galaxies. Such AGN feedback could possibly explain the observed high-velocity outflows in post-starburst galaxies \citep{Tremonti2007} and ULIRGs \citep[\eg][]{Chung2011, Rupke2011, Sturm2011}. Further, if the paradigm suggesting that AGN feedback is responsible for regulating the growth of galaxies is correct, LoBALs may be at a unique stage where strong outflows are present, yet, star formation is still in the process of being quenched. Previous studies of large samples of BAL QSOs addressing their SEDs \citep{Gallagher2007} and submillimeter properties \citep{Willott2003} find that BAL and non-BAL QSOs are indistinguishable, consistent with the model that all QSOs contain BAL winds, and their detection depends on viewing angle. However, those samples mainly comprise HiBALs and refrain from drawing conclusions about LoBALs. Even if the detection of BAL troughs in QSO spectra depends on viewings angle, compelling evidence suggests that LoBALs are linked to IR-luminous galaxies, with dominant young stellar populations and disturbed morphologies. To test this possibility, we have undertaken the first multiwavelength investigation of a volume-limited sample of LoBALs. In a series of three papers, we address the nature of low-redshift LoBALs and their relationship to the broader QSO population. In particular, we test the idea that LoBALs might be a short, evolutionary stage when the AGN has been recently fueled by a merger and the ensuing winds are in the process of quenching the star formation. In this first paper of the series, we present $Spitzer$ IRS spectroscopy and MIPS photometry at 24, 70, and 160 $\mu$m of a volume-limited, statistically-significant sample of low-redshift, optically-selected LoBALs. To study their star-forming histories, we model the infrared SEDs of LoBALs and measure their far-infrared luminosities and star formation rates. In the upcoming papers, we will study the detailed morphologies of LoBAL host galaxies via $HST$ imaging and the nature of their stellar populations via Keck LRIS spectroscopy. Our sample of LoBALs and a control sample are described in $\S$ 2. Details of the observations are explained in $\S$ 3. We present the analysis and results in $\S$ 4. Discussion and summary of results are given in $\S$ 5. The conclusion is presented in $\S$ 6. We assume a flat universe cosmology with $H_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\Omega_M$ = 0.27, and $\Omega_{\Lambda}$ = 0.73. All luminosities in units of the bolometric solar luminosity were calculated using $L_{\odot}$ = 3.839 $\times$ 10$^{33}$ erg s$^{-1}$.
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We investigate the mid- and far-IR properties of a volume-limited sample of 22 low-ionization broad absorption line QSOs within the redshift range $0.5 < z < 0.6$. We model their SEDs from the optical to the far-infrared in an effort to estimate total infrared luminosities, the relative contributions from the starburst and the AGN, starburst luminosities, and star formation rates corrected for the AGN contamination of the FIR emission. We compare this LoBAL sample to a control sample of non-LoBALs, matched by $M_i$ within the redshift $0.45 < z < 0.83$, to examine the possible connection between these two classes of QSOs. We find that LoBALs are indistinguishable from non-LoBAL type-1 QSO in terms of their MIR spectral properties and FIR luminosities.
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1206.1827
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1206
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1206.5545_arXiv.txt
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We present a photometric and spectroscopic study of the unique isolated nearby dSph galaxy KKR25. The galaxy was resolved into stars with HST/WFPC2 including old red giant branch and red clump. We have constructed a model of the resolved stellar populations and measured the star formation rate and metallicity as function of time. The main star formation activity period occurred about 12.6 to 13.7 Gyr ago. These stars are mostly metal-poor, with a mean metallicity $\feh \sim -1$ to $-1.6$ dex. About 60 per cent of the total stellar mass was formed during this event. There are indications of intermediate age star formation in KKR25 between 1 and 4 Gyr with no significant signs of metal enrichment for these stars. Long-slit spectroscopy was carried out using the Russian 6-m telescope of the integrated starlight and bright individual objects in the galaxy. We have discovered a planetary nebula (PN) in KKR25. This is the first known PN in a dwarf spheroidal galaxy outside the Local Group. We have measured its oxygen abundance $12+\log(\textrm{O/H}) = 7.60 \pm 0.07$ dex and a radial velocity $V_h = -79$ \kms{}. We have analysed the stellar density distribution in the galaxy body. The galaxy has an exponential surface brightness profile with a central light depression. We discuss the evolutionary status of KKR25, which belongs to a rare class of very isolated dwarf galaxies with spheroidal morphology.
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\begin{figure*} \begin{tabular}{cc} \includegraphics[height=0.45\textwidth]{fig1a.ps} & \includegraphics[height=0.45\textwidth]{fig1b.ps} \end{tabular} \caption{ The map of galaxies in the SuperGalactic coordinates is centred on KKR\,25. The left panel presents the projection of galaxies on the SuperGalactic plane, while the right panel shows the edge-on view on the `pancake' of galaxies. The filled circle size is proportional to absolute magnitude of galaxies. Galaxies are coded by a colour from red for early types ($T=-5$) to blue for late types ($T=10$), according to morphological type in de Vaucouleur's numerical scale \citep{RC3}. The zero-velocity surface is shown by big blue circle around the Local Group. Distance to nearby structure is marked by dash-dotted lines with corresponding distance written on it. The brightest galaxies in the volume under consideration are Milky Way (WM), Andromeda galaxy (M31) and M81. They are signed on the figures. KK\,230 is a closest galaxy to KKR\,25. Three associations of dwarf galaxies \citep{DwarfsAssociations} are shown on left panel. DDO\,190 is the brightest member of 14+08 association. DDO\,187 corresponds to `Dregs' association and NGC\,4214 to 14+07. The Local Void occupies the upper half of the right panel just above KKR\,25. } \label{f:XYZ} \end{figure*} The isolated dwarf spheroidal galaxy KKR\,25 was discovered by \citet{KKR1999} during the search of dwarf galaxies in the direction of the Local Void. Follow-up observations with the 100-m radio telescope at Effelsberg \citep{HI2,HI5} have shown \hi\ emission in the object with a radial velocity $V_h=-139.5$ \kms. Direct images of this low surface brightness galaxy were obtained with the 6-meter telescope of the Special astrophysical observatory (SAO) of the Russian Academy of Sciences \citep{KKR25+HST} and with the Hubble Space Telescope (\textit{HST}). Its colour-magnitude diagram (CMD) shows a red giant branch population and a trace of blue stars. \citet{KKR25+HST} classify the object as transition type galaxy (dIrr/dSph) at the distance of 1.86 Mpc. On the other hand, the spectral survey of nearby dwarf LSB galaxies with the Russian 6-m telescope failed to detect an optical velocity of KKR\,25 \citep{NearbyLSB+Vh}. Moreover, deep radio observations with the Giant Metrewave Radio Telescope (GMRT) \citep{KKR25+GMRT} did not show significant \hi\ emission in the range $-256<V_h<-45$ \kms\ at the level $M_{HI}=0.8\times10^5$ $M_{\sun}$. \citet{KKR25+GMRT} concluded that `the non-detection of \hi\ in KKR\,25 suggests that previous single-dish measurements were affected by confusion with the Galactic emission. Our stringent limits on the \hi\ mass of KKR\,25 indicate that it is a normal dSph galaxy'. KKR\,25 is one of the most isolated galaxies inside the sphere of 3 Mpc around us. It settles at the distance of 1.9 Mpc from the Milky Way and at 1.2 Mpc above the SuperGalactic plane in the front of the Local Void. KKR\,25 is far away from the zero-velocity surface of $R_0=0.96\pm0.03$ Mpc \citep{KKMT2009}, which separates the Local Group from the cosmic expansion. The Local Group is the nearest massive structure to KKR\,25. The second close massive group is the M\,81 at the distance of 2.56 Mpc. There are no galaxies closer than 1 Mpc to KKR\,25 (see Fig.\ \ref{f:XYZ}). The nearest neighbour is the dwarf galaxy KK\,230 ($M_B=-9.8$). The isolation of KKR\,25 was pointed out by \citet{KKR25+HST}. \citet{DwarfsAssociations} note that KKR\,25 is only isolated galaxy on the scale of 3 Mpc from us and `every object in this volume is associated with either a luminous group, an association of dwarfs, or the dregs evaporating association'. Two associations 14+08 (around DDO\,190) and `Dregs' (around DDO\,187) stand on the distance of 1.2 and 1.4 Mpc from KKR\,25, respectively. In spite of its isolation KKR\,25 has no gas and looks like a normal dwarf spheroidal system. This fact draws our attention, because we expect to find dSph galaxies in dense regions, like groups and clusters of galaxies. Obviously, that any kind of interaction with massive galaxy is not suitable to explain properties of KKR\,25. This galaxy can play a crucial role in testing of different scenarios of dSph's formation. \citet{KKR25+HST} have found a globular cluster candidate in the HST images of the galaxy. An apparent magnitude of the object $V_T=20.59$ corresponds to $M_V=-5.79$, which is typical for Galactic globular clusters. In the framework of the \Halpha{} survey of the Canes Venatici I cloud of galaxies with the Russian 6-m telescope \citet{CVnI+Halpha} also have found a faint \Halpha{} knot on northern side of KKR\,25. A measured flux of the knot is $\log F=-14.64$ erg\,cm$^{-2}$\,s$^{-1}$. These interesting objects were targeted for spectroscopic study with 6-m telescope in current study. The main parameters of KKR\,25 are presented in the Table~\ref{t:KKR25}. Coordinates are taken from HyperLEDA\footnote{\url{http://leda.univ-lyon1.fr/}} \citep{HyperLEDA}. Apparent sizes were published by \citet{KKR1999}. The colour $(V-I)_T$ was measured by \citet{KKR25+HST}. The central surface brightness $\Sigma_{V}$ was estimated from profiles published by \citet{KKR25+HST}. All other values, total magnitude $V_T$, axis ratio, scale length $h$, heliocentric velocity $V_h$ and distance modulus $(m-M)_0$, are derived in the current work. The $V_T$, $(V-I)_T$ and $\Sigma_{V}$ magnitudes are not corrected for Galactic extinction. \begin{table} \caption{Main parameters of KKR\,25.} \begin{tabular}{@{}lr@{\,\,}ll} R.A. (J2000) & \multicolumn{2}{c}{$16\,13\,47.6$} & HyperLEDA \\ Dec (J2000) & \multicolumn{2}{c}{$+54\,22\,16$} & HyperLEDA \\ $E(B-V)$, mag & 0.008 & & \citet{DustMap} \\[3pt] Size, arcmin & \multicolumn{2}{c}{$1.1\times0.65$} & \citet{KKR1999} \\ $h$, arcsec & 16.7 & $\pm 1.1$ & this work \\ $b/a$ & $0.51$ & $\pm0.03$ & this work \\[3pt] $V_T$, mag & $15.52$ & $\pm0.22$ & this work \\ $(V-I)_T$, mag & $0.88$ & & \citet{KKR25+HST} \\ $\Sigma_{V}$, $\textrm{mag}/\square^{\prime\prime}$ & $23.97$ & $\pm0.03$ & \citet{KKR25+HST} \\[3pt] $V_h$(stars), \kms & $-65$ & $\pm15 $ & this work \\ $V_h$(PN), \kms & $-79$ & $\pm9 $ & this work \\[3pt] $(m-M)_0$, mag & $26.42$ & $\pm0.07$ & this work \\ Distance, Mpc & $1.93$ & $\pm0.07$ & this work \\[3pt] $V_{LG}$, \kms & $128$ & & this work \\ $M_V$, mag & $-10.93$ & & this work \\ $L_V$, $10^6 L_\odot$ & $2.0$ & & this work \\ $\Sigma_V$, $L_\odot$\,pc$^{-2}$& $9.6$ & & this work \end{tabular} \label{t:KKR25} \end{table}
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We present a photometric and spectroscopic study of the unique isolated nearby dSph galaxy KKR\,25. Let us briefly summarize the results of our study. We have estimated the distance modulus of KKR\,25 $(m-M)_0=26.42\pm0.07$ mag using the TRGB method. It corresponds to a distance $D=1.93\pm0.07$ Mpc. The new value is in good agreement with all previous distance measurements. We have derived a quantitative star formation history of the isolated dwarf spheroidal galaxy KKR\,25. The star formation history was reconstructed using \textit{HST}/WFPC2 images of the galaxy and a resolved stellar population modelling. According to our measurements, 62 per\,cent of the total stellar mass were formed during the initial burst of star formation occurred about 12.6 -- 13.7 Gyr ago. There are indications of intermediate age star formation in KKR\,25 between 1 and 4 Gyr with no significant signs of metal enrichment for these stars. A distribution of the stars in the galaxy is well described by an exponential profile with central depression. The exponential scale length is $h=156^{+12}_{-11}$ pc. The profile extends up to 5 scale lengths. The size of the depression $R=170^{+22}_{-30}$ is about the exponential scale length. We did not confirm the presence of globular clusters in KKR\,25. In the fact, the previously selected candidates are background objects, S0 galaxy at $z=0.34$ and QSO at $z=0.75$. The spectroscopy of \Halpha{} object in KKR\,25 revealed that it is a planetary nebula with oxygen abundance $12+\log(\textrm{O/H}) = 7.60 \pm 0.07$. We have serendipitously found the first PN in the dwarf spheroidal galaxy outside the Local Group. The search of extraordinary blue stars on CMD of stellar population gives the perspective method for selection of PN candidates in distant galaxies. We have derived heliocentric velocity of KKR\,25 using PN emission lines $V_h=-79\pm9$ and using integrated light of the galaxy $V_h=-65\pm15$. Our study shows, that KKR 25 belongs to the population of highly isolated dwarf spheroidal galaxies, now rarely detected, but extremely important for our understanding of galaxy evolution theory. The `primordial scenario' of galaxy formation is preferable against tidal stripping mechanism to explain the isolation of KKR\,25 and its morphology. Existence of big number of dwarf spheroidals in the field could explain the overabundance problem in modern simulations. The search for the dwarf spheroidal in voids is a crucial test for models of formation and evolution of dwarf galaxies.
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1206.5545
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0112
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hep-ph0112253_arXiv.txt
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s{ After brief personal recollections of the author's long-time friendship with Misha Marinov the problem of particle production by classical time- varying scalar field is discussed. In the quasiclassical limit the calculations are done by imaginary time method developed, in particular, in Marinov's works. The method permits to obtain simple analytical expressions which well agree with the later found numerical solutions. The results are compared with perturbative calculations and it is argued that perturbation theory gives an upper limit for the rate of production. } \vspace{0.4cm} \tableofcontents \newpage
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} It is demonstrated that imaginary time method very well describes particle production by scalar field. It is very simple technically and permits to obtain physically transparent results. The calculations here were done for a particular case of periodic or quasiperiodic oscillations of the field but, as shows the experience with production of $e^+e^-$-pairs by electric field (for a review see e.g. third paper in ref.\cite{marinov72}), the method also works well in the opposite case of short pulse fields. The method is applicable in the quasiclassical limit. In the opposite case perturbation theory is applicable and hence one can obtain simple and accurate (semi)analytical estimates practically in all parameter range. The results of calculations in the quasiclassical limit are in a good agreement with subsequent numerical ones\,\cite{baacke98,green99}. An important difference between the latter papers and the initial one\,\cite{dolgov90} lays in the interpretation of the results. According to all these papers the occupation numbers of the produced particles quickly approaches unity but, in contrast to refs.\cite{baacke98,green99}, it is argued in the paper\,\cite{dolgov90} that the total production rate is nevertheless suppressed in comparison to perturbation theory and the production of fermions by the inflaton with Yukawa coupling to fermions is always weak. This conclusion is verified above. As is shown in this paper, the occupation numbers may quickly reach unity both in perturbation theory and in non-perturbative case. Still even the production rate of particles obeying Boltzmann statistics is very weak to ensure fast (pre,re)heating. In the case of fermion production the rate is evidently much weaker because the production must stop when the occupation number reaches unity and to continue the process the produced fermions should be eliminated from the band. As is argued in sec.~\ref{ss:small}, the non-perturbative effects can only diminish the production rate. The bosonic case is opposite: more bosons are in the final state, the faster is production. Thus even in perturbation regime the boson production can be strongly amplified because their occupation number may reach unity in much shorter time than $1/\Gamma$ and the energy may be transferred from the inflaton to the produced bosons much faster than is given by the original perturbative estimates\,\cite{dolgov82}, where the effect of stimulated emission was not taken into account. Of course to realize this regime the band should not be destroyed by expansion and scattering, as argued in ref.~\cite{dolgov90}. To summarize, we have shown that perturbation theory gives a good estimate of production of light fermions and bosons if Fermi exclusion principle or stimulated emission respectively are taken into account. The formally calculated production rate in perturbation theory is always larger than the non-perturbative one, at least in the simple cases that we have considered. So the results of perturbation theory may be used as upper bounds for production rates. Moreover, perturbation theory helps to understand physical meaning of the obtained results and to interpret them correctly. In many realistic cases (e.g. for large $g\phi_0$ or $m_0$) perturbation theory is not applicable and to calculate the real production rate (not just an upper bound) one has to make more involved non-perturbative calculations. In quasiclassical (anti-perturbative) limit imaginary time method permits to obtain accurate and simple results and to avoid complicated numerical procedure
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hep-ph0112253_arXiv.txt
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0112
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astro-ph0112416_arXiv.txt
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We image 19 quasars with 22 \dla \ (DLA) systems using the F160W filter and the Near-Infrared Camera and Multiobject Spectrograph aboard the {\it Hubble Space Telescope}, in both direct and coronagraphic modes. We reach 5$\sigma$ detection limits of \ab H=22 in the majority of our images. We compare our observations to the observed Lyman-break population of high-redshift galaxies, as well as Bruzual \& Charlot evolutionary models of present-day galaxies redshifted to the distances of the absorption systems. We predict H magnitudes for our DLAs, assuming they are producing stars like an L$_{*}$ Lyman-break galaxy (LBG) at their redshift. Comparing these predictions to our sensitivity, we find that we should be able to detect a galaxy around 0.5-1.0 L$_{*}$(LBG) for most of our observations. We find only one new possible candidate, that near LBQS0010-0012. This scarcity of candidates leads us to the conclusion that most \dla \ systems are not drawn from a normal LBG luminosity function nor a local galaxy luminosity function placed at these high redshifts.
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In recent years, our knowledge of the high redshift universe has finally grown to include a significant number of galaxies powered mainly by star formation, often referred to as Lyman-break galaxies because most were discovered by searching for the Lyman-limit discontinuity using broad band colors \citep{ste96}. There is still the concern, however, that these galaxies may not be representative of the typical galactic mass at that epoch. Dust, for example, could strongly affect the discovery rate of galaxies found using searches, like the Lyman-break method, that search in rest frame ultraviolet light \citep{row97}. Quasar absorption line studies, on the other hand, have the advantage of being able to follow the neutral hydrogen content of the universe regardless of whether it emits light or not. While absorption line studies may also suffer from selection effects, with dustier intervening gas creating dimmer quasars that are either less well studied or not discovered at all \citep{pei95, car98}, they represent a potentially less biased tracer of the content of the universe than the Lyman-break galaxy. Examinations of \dla \ (DLA) systems, which make up the bulk of all absorption line neutral hydrogen, show that at high redshifts, the mass of neutral hydrogen seen in DLAs per unit comoving volume is roughly the same as the mass density seen in luminous matter (i.e. stars) in present-day spiral galaxies \citep{lan91}. Since this is a much larger quantity of neutral gas than we observe in galaxies today \citep{rao93, zwa98}, this suggests that much of the gas we see in high z DLAs has ended up in stars. In fact, \cite{wol95} identify the DLA systems as the likely progenitors of current spiral galaxies. However, some recent work has challenged this idea, showing no drop in the cosmological mass density of neutral hydrogen from redshifts of 3 down to 0.5 \citep{rao00}, instead of the expected gradual decline seen in other work \citep{smi96,jan98,sto00}. Another indication that DLAs may not be forming today's galaxies is the lack of any metal enrichment in lower redshift DLAs, with the mean Zn/H ratio remaining around 1/10 solar from z=1-3 \citep{pet99}. It should be noted that the statistical uncertainties involved in these measurements remain significant, since all work on DLAs at low redshifts relies on a small number of objects (23 with z$<$1.65 for \cite{rao00}), with the handful with highest column densities dominating the statistics. These low redshift DLA studies are also the most vulnerable to bias from any possible obscuration of quasars by dust in foreground absorbers \citep{sto00, pei95}. Modeling of velocity profiles has not settled the question. Supporting the contention that DLAs form spiral galaxies, \cite{pro97} demonstrate how models of rotating disks, not unlike today's spirals, fit the observed double trough velocity profiles seen in some DLA metal lines. However, there are many other viable models that also produce similar velocity profiles, such as merging protogalactic clumps \citep{hae98}, randomly moving clouds in spherical halos \citep{mcd99}, and multiple gaseous discs in a single halo \citep{mal00}. Another problem with the DLAs as spiral galaxies hypothesis is the number of DLA systems that are being discovered, compared with the number of spiral galaxies around at the present day. Based on the present day density and size of spiral galaxies, one can estimate the number that should intercept a random line of sight to a quasar at some high redshift. While the exact prediction depends on the assumed cosmology, \cite{lan91} found that the number of DLA systems being discovered is at least twice that predicted, even under the more favorable cosmological assumptions. This means that if DLAs are high-redshift spiral galaxies, then there must be strong evolution in the spiral population with redshift. Either spiral galaxies were more numerous in the past, or they were larger or both. Another way around this discrepancy is to assume that our understanding of local galaxies and the measured fiducial parameters used for these predictions are flawed. For instance, a large population of low surface brightness galaxies \citep{imp89} or low luminosity dwarf galaxies \citep{cen98} could dominate the absorption cross sections responsible for \dla \ absorption lines. This, however, is not the case for low-redshift Mg II absorption line systems for which visible counterparts have been found, where the majority are essentially bright galaxies with typical luminosities just under L$_{*}$ \citep{ber91}. It would be valuable to find some connection between DLAs and high-redshift galaxies, our two largest sources of information on the non-AGN high-redshift universe. Unfortunately, while we know a great deal about the distribution, metallicity, and evolution of the neutral gas in high-redshift DLAs, we know almost nothing about their associated stars, galactic sizes, or morphologies. Only by imaging starlight from these DLA systems can we measure how much, if any, star formation is actually occurring in these objects. Lyman break galaxies produce stars at a steady pace of about 10 M$_{\sun }$ yr$^{-1}$, depending heavily on cosmology and the large dust extinctions assumed \citep{see96}, not that different from the \ab 10 M$_{\sun }$ yr$^{-1}$ seen in H$\alpha$ emission line galaxies today \citep{gal95}. If DLAs are progenitors of present day galaxies and/or the same population as the Lyman break galaxies, then they should also be producing stars and starlight. This is true whether the DLAs are galactic disks, spheroids, or some other assembled structure. Searches for visible candidates for low to medium-redshift DLA systems have had mixed results. Often candidates are found, but less than half of them are suspected to be reasonably bright spiral galaxies. The rest are a non-uniform collection of low surface brightness galaxies, bright compact objects, and dim dwarfs \citep{ste94,leb97,lan97,rao98,pet00}. Again the statistics are poor (less than 15 objects) and most have not been spectroscopically confirmed, leaving the nature of these DLAs still murky. High-redshift systems have been even more problematic, with much lower candidate discovery rates. A major reason is that this is a difficult measurement, considering the large distances to these objects and their close projected proximity to a bright quasar. A common search technique has been to look for emission lines, Lyman \a \ or H\a , where the emission line flux ought to stand out from the background and the quasar light may be depressed. These searches have found at least two fairly unambiguous detections, near quasars PKS0528-250 \citep{mol93} and 2231+131 \citep{djo96}, but broader searches of DLA samples have generally found nothing at all \citep{low95,bun99}. Another strategy for imaging DLAs has been to go into the infrared, where K-corrections would be more favorable and dust, often cited as a possible explanation for the failure of Lyman \a \ searches, would likely be only a minor effect. For instance, \cite{ara96} observed ten quasars with DLA systems in the near-infrared, discovering two candidates (in front of quasars 0841+129 and 1215+333) after subtracting off the PSF of quasar. Since both candidates lie about an arcsecond away from the quasar, the PSF subtraction is critical and an error could lead to a false detection. Higher resolution imaging would greatly help. We have obtained imaging of 19 quasars with confirmed DLA systems using the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) onboard the Hubble Space Telescope (HST). This gives us the advantage of both the near-infrared, with its favorable K-corrections and low dust dependence, and high spatial resolution, which will separate any candidates from the quasar's light, even at distances of less than arcsecond. In addition, the majority of the observations were done using the NICMOS coronagraph, which greatly decreases scattered light from the quasar, improving chances of candidate detection. If DLAs are high-redshift disk galaxies producing stars, missed only because of their close proximity to the quasar, a survey of this type can uncover them. Another advantage of a large survey of a sample like ours is that it presents the failures as well as the successes, something that is often missing from papers on individual candidate discoveries. A rate of DLA system discoveries gives us a statistically significant result for the entire DLA class of objects, something an individual candidate can not. Section 2 describes our sample and the observations, while $\S$3 describes the data reduction and analysis. Section 4 discusses the sensitivity of our survey to these hypothetical L$_{*}$(LBG) galaxies. Section 5 presents the results of our candidate search. Finally, $\S$ 6 presents our discussion and conclusions. (A list of notes on individual DLAs is provided in the Appendix).
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We are not seeing galaxies near quasar lines of sight that contain \dla \ absorption. This result is very similar to the negative results of groups surveying DLA systems searching for the emission lines of Lyman \a \ \citep{low95} and H\a \ \citep{mal95,bun99}, suggesting that that their difficulties may not have been because the DLA systems do not produce strong emission lines, but because they do not produce much light at all. Our ability to look much closer to the quasar, 0.5 arcseconds or better, also weakens the argument that DLAs are not being found because they lie at such small impact parameters. In fact, the candidate seen by \cite{ara96} near Q1215+333 at a distance of 1.3 arcseconds is not seen in our data. With a K magnitude of 20.1, we should have picked it up, even with a red color (H-K $\geq$ 1.5). We tried models creating all stars in a single burst around z=10, but none were redder than H-K = 1.1 at z=2 \citep{bru96}, with most of our exponential models predicting H-K = 0.8-0.9. Colors redder than H-K = 1.2 required redshifts of z=3-3.5 and produced correspondingly dimmer galaxies. This makes it likely the former detection was a result of the very difficult subtraction of a bright PSF so near the quasar. It does not seem likely that we can continue to hide the majority of DLA systems by pushing them all into the tiny area (less than 0.8 arcsec$^2$ for our cornagraphic images) directly in front of the quasar, especially considering the extreme number evolution of galaxies that would require. Also not seen are the \cite{ste92} candidate for [HB89] 0000-263 (DLA \# 1), at a distance of 2.8$\arcsec$, and the \cite{ste95} and \cite{leb97} candidate for [HB89] 0454+039 (DLA \# 6), measured at two different distances from the quasar, 2.1$\arcsec$ and 0.8$\arcsec$ respectively. These last two may have been missed if the candidate were blue enough, R-H $\leq$ 3.4 for the [HB89] 0000-263 candidate and R-H $\leq$ 2.3 for [HB89] 0454+039. Models of high-redshift colors \citep{bru96} suggest that the former is likely (predicted R-H = 1.3-1.5 for exponential star formation models at z=3.4), while the latter is possible, but difficult to fit (predicted R-H = 2.3-3.5 for exponential star formation models at z=0.86). Galaxies as blue as R-H = 2.3 at z=0.86 require either strong continuous star formation or a strong starburst just prior to observation. All this leads us to speculate as to why we are not detecting light from almost any of these DLA systems. One possibility is that they not producing much light because they have not started making many stars by the period at which we are observing them. Most of our DLA sample comes from around z$\sim$2, so if most star formation occurred after that we would not be able to see it. One possible problem with this scenario is the previously observed decrease in comoving mass density of neutral gas seen in DLAs from z$\sim$3.5 to z=2 and then down into the present day. If DLAs are spiral galaxy progenitors, this is explained as the transformation of their neutral gas into stars. Initial work by \cite{wol95} indicated that as much one half of the neutral gas was depleted before z=2, meaning half of all stars should have been created before then and would therefore likely be visible to our survey. However, further work by \cite{smi96}, using some of the highest redshift DLAs known, showed a less dramatic neutral gas evolution, with only 20\% of the neutral gas processed by z=2. Also, as noted in the introduction, \cite{rao00} puts the whole idea of the decline of neutral gas in DLAs with time into doubt, showing no decline down into low redshifts. Neutral gas evolution aside, if DLAs are to become today's galaxies they must form stars at some point and a large amount of star formation is already measured to be underway before z=2 \citep{mad98}. Waiting until low redshifts to produce the vast majority of a galaxy's stars would create a problem with the number of highly luminous galaxies required at z$<$2. Dust could also play a factor, absorbing light even at rest wavelength optical wavelengths. However there is no evidence for such large quantities of dust seen in the quasar spectra \citep{pei91,pet97}, requiring a most unlikely conspiracy of much higher dust quantities surrounding the star formation regions, but not along the quasar line of sight. Another possibility is that the DLA systems are not compact, high surface brightness objects like the spiral galaxies we study, but are instead diffuse, low-surface-brightness galaxies as suggested by \cite{jim99}. These galaxies would be inefficient at transforming their gas into stars, so would take longer to start forming stars, and would be harder to detect, even if their total magnitude equalled our predictions. If this is the case, the majority of DLAs are not like the LBGs discovered so far, and either these low surface brightness galaxies have to evolve into the high surface brightness spirals we see around us today, or these DLA systems will never become today's L$_{*}$ galaxies. This possibilty -- that DLAs are not high-redshift versions of present day disk galaxies -- could by itself explain all of our upper limits. They might be protogalactic blobs that will eventually fall into galaxies, helping them form or grow, or they might never become part of standard galaxies at all, being something entirely separate. While there may be a handful of DLA systems that we can detect in emission, the vast majority do not seem to be taken from the distributions of either $z=0$ or $z=3$ galaxies. Either they are not producing as many stars, something is absorbing their light, or their surface brightness is unusually low. From the evidence gathered so far, it appears that the distribution of DLA systems is inconsistent with both the evolution of Lyman-break galaxies forward in time and present day galaxies backwards.
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astro-ph0112416_arXiv.txt
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0112
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astro-ph0112550_arXiv.txt
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An overview of the current status of WIMP direct searches is presented, emphasizing strategies, achievements and prospects.
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Experimental observations and robust theoretical arguments have established that our universe is essentially non-visible, the luminous matter scarcely accounting for one per cent of the energy density of a flat universe. The distribution of a flat universe ($\Omega=\Omega_{M}+\Omega_{\Lambda}=1$) attributes to the dark energy about $\Omega_{\Lambda}\sim70\%$, whereas the matter density takes the remaining $\Omega_{M}\sim30\%$, consisting of, both, visible ($\Omega_{l}\sim0.5\%-1\%$) and non-visible (dark) matter. This dark component consists of ordinary baryonic matter ($\Omega_{B}\sim 4-5\%$), (possibly made by machos, jupiters, dust, black holes, etc.) and a large fraction (up to $\Omega_{NB}\sim25\%$) of non baryonic dark matter, supposedly made by non-conventional, exotic particles. The minimal requirements to be fulfilled by the non-baryonic dark particles are to provide the right relic abundance, to have non-zero mass, zero electric charge and a weak interaction with ordinary matter. There are several candidates to such species of matter provided by schemes beyond the Standard Model of Particle Physics. The galactic DM axions, the SUSY WIMPs (like neutralinos) and the light neutrinos (of some non standard models) are particularly attractive. The WIMPs, Weak Interacting Massive Particles, are favorite of experimentalists ($\sim20$ experiments for WIMPs vs. 2 for axions) and so of this talk. This talk deal with the recent efforts done in its direct search illustrated by a few, selected experiments. \begin{table*}[htb] \caption{WIMP Direct Detection in underground facilities experiments currently running (or in preparation) } \label{table:2} \begin{tabular}{lll} \hline LABORATORY & EXPERIMENT & TECHNIQUE \\ \hline BAKSAN (Russia) & IGEX & 3$\times$1 Kg Ge-ionization \\ BERN(Switzerland) & ORPHEUS &(SSD) Superconducting Superheated Detector, 0.45 Kg Tin \\ BOULBY & NaI & NaI scintillators of few Kg (recently completed) \\ (UK)& NAIAD & NaI unencapsulated scintillators (50 Kg)\\ & ZEPLIN & Liquid-Gas Xe scintillation/ionization I: 4 Kg single phase \\ & & II: 30 Kg Two phases \\ & DRIFT & Low pressure Xe TPC (in preparation) 1 m$^{3}\rightarrow$ 10 m$^{3}$ \\ CANFRANC & COSME & 234 g Ge ionization \\ (Spain) & IGEX & 2.1 Kg Ge ionization \\ & ANAIS & 10$\times$10.7 Kg NaI scintillators \\ & ROSEBUD & 50g Al$_{2}$O$_{3}$ and 67g Ge thermal detectors \\ & & CaWO$_{4}$ 54g and BGO 46g scintillating bolometers \\ FREJUS/MODANE & SACLAY-NaI & 9.7 Kg NaI scintillator (recently completed) \\ (France) & EDELWEISS I &70 g Ge thermal+ionization detector \\ & EDELWEISS II & 4$\times$320 g Ge thermal+ionization detectors \\ GRAN SASSO &H/M &2.7 Kg Ge ionization \\ (Italy) &HDMS & 200g Ge ionization in Ge well \\ & GENIUS-TF & 40$\times$2.5 Kg unencapsulated Ge (in preparation) \\ & DAMA & NaI scintillators (87.3 Kg) \\ & LIBRA & NaI scintillators 250 Kg (in preparation) \\ & Liquid-Xe & Liquid Xe scintillator (6 Kg) \\ & CaF$_{2}$ & Scintillator \\ & CRESST I & (4$\times$260g) Al$_{2}$O$_{3}$ thermal detectors \\ & CRESST II & Set of 300g CaWO$_{4}$ scintillating bolometers (up to 10 Kg) \\ & MIBETA & 20$\times$340g TeO$_{2}$ thermal detector \\ & CUORICINO & 56$\times$760g TeO$_{2}$ thermal detector (being mounted) \\ & CUORE &1000$\times$760g TeO$_{2}$ (in preparation) \\ RUSTREL (France)& SIMPLE & (SDD)Superheated Droplets Detectors (Freon) \\ STANFORD UF/ & CDMS - I & 100g Si; 6$\times$165g Ge thermal+ionization detectors \\ SOUDAN(USA)& CDMS - II & 3$\times$250g Ge and 3$\times$100g Si Thermal+Ionization \\ SNO (Canada) & PICASSO &(SDD)Superheated Droplets Detectors (1.34g of Freon) \\ OTO & ELEGANTS-V & Large set of massive NaI scintillators \\ (Japan) &ELEGANTS-VI & CaF$_{2}$ scintillators \\ \hline \end{tabular}\\[2pt] \end{table*} Galactic halo WIMPs could be directly detected by measuring the nuclear recoil produced by their elastic scattering off target nuclei in suitable detectors at a rate which depends of the type of WIMP and interaction. In the case of WIMPs of $m\sim GeV~\mathrm{to}~\textit{TeV}$ and $v\sim10^{-3}c$ the nuclear recoil in the laboratory frame $E_{R}=\frac{\mu^{2}}{M}v^{2}(1-cos\theta)$ is in the range from 1 to 100 \textit{KeV}. $M$ is the nuclear mass, $\mu$ the ($m, M$) reduced mass and $\theta$ the WIMP-nucleus (c. of m.) scattering angle. Only a fraction $QE_{R}=E_{vis} (\equiv E_{eee})$ of the recoil energy is visible in the detector, depending on the type of detector and target and on the mechanism of energy deposition. The so-called Quenching Factor Q is essentially unit in thermal detectors whereas for the nuclei used in conventional detectors it ranges from about 0.1 to 0.6. The energy delivered by the WIMP results in a small signal (1-100 KeV) which shows up even smaller (for Q$<$1). Moreover this signal falls in the low energy region of the spectrum, where the radioactive and environmental background accumulate at much faster rate and with similar shape. That makes WIMP signal and background practically undistinguishable. On the other hand, the smallness of the neutralino-matter interaction cross-section implies that the process looked for is very rare. Customarily, one compares the predicted event rate with the observed spectrum. If the former turns out to be larger than the measured one, the particle which would produce such event rate can be ruled out as a Dark Matter candidate. That is expressed as a contour line $\sigma$(m) in the plane of the WIMP-nucleon elastic scattering cross section versus the WIMP mass. That excludes, for each mass m, those particles with a cross-section above the contour line $\sigma$(m). The level of background sets, consequently, the sensitivity of the experiment in eliminating candidates or in constraining their masses and cross sections. However, this simple comparison will not be able to identify the WIMP. A convincing proof of its detection should be provided by a distinctive signature characteristic of WIMPs. Such distinctive labels do exist: they are originated by the motion of the Earth in the galactic halo \cite{Drukier,Spergel}. These signatures are an annual modulation of the rate and a directional asymmetry of the nuclear recoil. Narrowing first the window of the possible WIMP existence and looking then for its identification is the purpose of the experimental searches. Table \ref{table:2} gives an overview of the experiments on direct detection of WIMPs currently in operation or in preparation, which is the subject of this talk. General reviews for WIMP dark matter are given in Ref. \cite{Gri}. WIMP direct detection is reviewed, for instance in Ref. \cite{Mor2,Mor}. WIMPs can be also looked for, indirectly, in the galactic halo, looking for its presence in cosmic ray experiments in terms of antiprotons, positrons or gamma rays produced by WIMP annihilation in the halo. One can also search in underground, underwater or under-ice detectors, looking (also indirectly) for WIMPs through the high energy neutrinos emerging as final products of the WIMP annihilation in celestial bodies (Earth or Sun).
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The direct search for WIMP dark matter proceeds at full strength. More than twenty experiments on direct detection illustrate the effort currently being done. New, dedicated experiments are focusing now in the identification of WIMPs, discriminating the nuclear recoils from the background, rather that in constraining or excluding their parameters space. Their current achievements and the projections of some of them have been reviewed in this talk. \begin{figure}[b] \centerline{\includegraphics[height=4.5cm]{fig13.eps}} \caption{} \label{fig13} \end{figure} The present experimental situation can be summarized as follows: the rates predicted for SUSY-WIMPs extend from 1-10 c/Kgday down to $10^{-4}-10^{-5}$ c/Kgday, in scatter plots, obtained within MSSM as basic frame implemented in various alternative schemes. A small fraction of this window is testable by some of the leading experiment. The rates experimentally achieved stand around 1 c/Kgday (0.1 c/Kgday at hand) (CDMS, EDELWEISS) and differential rates $\sim0.1-0.05$ \textit{c/KeV Kg day} have been obtained by IGEX and H/M, in the relevant low energy regions. The deepest region of the exclusion plots achieved stands around a few $\times10^{-6}$pb, for masses 50-200 GeV (DAMA, CDMS, EDELWEISS, IGEX). The current status of the best exclusion plots is depicted comparatively in Fig. \ref{fig13}. There exists an unequivocal annual modulation effect (see Fig. \ref{fig4}) reported by DAMA (four yearly periods), which has been shown to the compatible (DAMA) with a neutralino-WIMP, of m$\sim50-60$GeV and $\sigma^{Si}_{n}\sim 7\times10^{-6}$pb. Recent experiments exclude at greater or lesser extend (CDMS, EDELWEISS, IGEX) the DAMA region. To reach the lowest rates predicted ($10^{-5}$ c/Kgday) in SUSY-WIMP-nucleus interaction, or in other words, to explore coherent interaction cross-sections of the order of $10^{-9}-10^{-10}$pb, substantial improvements have to be accomplished in pursuing at its best the strategies reviewed in this talk, with special emphasis in discriminating the type of events. These strategies must be focussed in getting a much lower background (intrinsic, environmental, ...) by improving radiopurity and shieldings. The nuclear recoil discrimination efficiency should be optimized going from above 99.7\% up to 99.9\% at the same time that the energy $E_{vis}$ at which discrimination applies should be lowered. The measurement of the parameters used to discriminate background from nuclear recoils should be improved and finally one needs to increase the target masses and guaranty a superb stability over large exposures. With these purposes various experiments and a large R+D activity are under way. Some examples are given in Table \ref{table:1}. The conclusion is that the search for WIMPs is well focused and should be further pursued in the quest for their identification. \begin{table*}[htb] \caption{WIMP Direct Detection Prospect} \label{table:1} \begin{tabular}{ll} \hline &BEING INSTALLED/OR PHASE II EXPERIMENTS~(To start 2001-2002) \\ \hline CDMS-II & (Ge,Si) Phonons+Ioniz 7 Kg, B$\sim 10^{-2}-10^{-3}$ c/Kgd, $\sigma \sim 10^{-8}$ pb \\ EDELWEISS-II & (Ge) Phonons+Ioniz 6.7 Kg, B$\sim 10^{-2}-10^{-3}$ c/Kgd, $\sigma \sim 10^{-8}$ pb (40-200 GeV) \\ CUORICINO & TeO$_{2}$ Phonons 42 Kg, B$\sim 10^{-2}$ dru, $\sigma \sim 0^{-7}$ pb \\ CRESST-II & CaWO$_{4}$ Phonons+light, $B<10^{-2}-10^{-3}$ dru (15 KeV), $\sigma \sim 10^{-7}-10^{-8}$ pb \\ & (50-150 GeV) \\ IGEX & Ge Ioniz 2.1 Kg, B$< 10^{-1}-10^{-2}$ dru, $\sigma \sim 2\times 10^{-6}$ pb (40-200 GeV) \\ HDMS & Ge Ioniz 0.2 Kg, $\sigma \sim 6 \times 10^{-6}$ pb (20-80 GeV) \\ ANAIS & NaI Scintillators 107-150 Kg, B(PSD)$\leq 0.1$ dru, $\sigma \sim 2 \times 10^{-6}$ pb \\ NAIAD & NaI Scintillators 10-50 Kg, B(PSD)$\leq 0.1$ dru, $\sigma \sim 10^{-6}$ pb (60-200 GeV) \\ \hline & IN PREPARATION~(To start 2002-2003) \\ \hline LIBRA (DAMA) & NaI Scintillators 250 Kg \\ GENIUS-TF & Ge Ioiniz 40 Kg, B$<10^{-2}$ dru, $E_{Thr}$=10 KeV $\rightarrow \sigma \sim 10^{-6}$ pb (40-200 GeV), \\ & $E_{Thr}$=2 KeV $\rightarrow \sigma \sim 10^{-7} pb$ (20-80 GeV) \\ ZEPLIN-II & Xe-Two-phase 40 Kg, NR discrim$>$99\%, B$<10^{-2}$ dru, $\sigma \sim 10^{-7}$ pb \\ DRIFT-I & Xe TPC 1 m$^{3}$, B$<10^{-2}$ dru, $\sigma \sim 10^{-6}$ pb (80-120 GeV) \\ \hline & IN PROJECT~($>$2003-2005) \\ \hline CUORE & TeO$_{2}$ Phonons 760 Kg, $E_{Thr}\sim$2.5 KeV, B$\sim 10^{-2}-10^{-3}$ dru, $\sigma \sim 5\times 10^{-8}$ pb \\ GENIUS 100 & Ge ioniz 100 Kg, $E_{Thr}\sim$10 KeV \\ (GENINO) & B$\sim 10^{-3}-10^{-5}$ dru, $\sigma \sim 5\times 10^{-8}-2\times 10^{-9}$ pb \\ GEDEON & Ge ioniz 28-112 Kg, B$\sim 2\times 10^{-3}$ dru ($>$10 KeV) $\sigma \sim 10^{-7}-10^{-8}$ pb (40-200 GeV) \\ \hline & THE FUTURE~($>$2005-2007) \\ \hline DRIFT 10 & Xe 10 m$^{3}$ TPC, $\sigma \sim 10^{-}8$ pb \\ ZEPLIN-MAX & Xe Two-Phase, $\sigma \sim 10^{-10}$ pb \\ GENIUS & Ge ioniz 1-10 Tons, $\sigma \sim 10^{-9}-10^{-10}$ pb \\ DRIFT-1 ton & Xe 1 Ton TPC, $\sigma \sim 10^{-10}-10^{-11}$ pb \\ \hline \end{tabular}\\[2pt] \end{table*}
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astro-ph0112550_arXiv.txt
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astro-ph0112146_arXiv.txt
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We present spectroscopic and high speed photometric data of the eclipsing polar V895 Cen. We find that the eclipsed component is consistent with it being the accretion regions on the white dwarf. This is in contrast to Stobie et al who concluded that the eclipsed component was not the white dwarf. Further, we find no evidence for an accretion disc in our data. From our Doppler tomography results, we find that the white dwarf has $M$\gtae 0.7\Msun. Our indirect imaging of the accretion stream suggests that the stream is brightest close to the white dwarf. When we observed V895 Cen in its highest accretion state, emission is concentrated along field lines leading to the upper pole. There is no evidence for enhanced emission at the magnetic coupling region.
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Polars are interacting binaries consisting of a red dwarf secondary and a strongly magnetized ($\sim$10-200 MG) white dwarf primary. In these systems the secondary fills its Roche lobe. Material falls under gravity from the secondary towards the primary, initially along the binary orbital plane before the magnetic field of the primary forces it to leave the orbital plane and eventually impacts quasi-radially onto the white dwarf. Unlike non-magnetic cataclysmic variables (CVs) the strong magnetic field of the white dwarf is high enough to prevent the formation of an accretion disc around the primary. Although our understanding of the accretion flow near the white dwarf is now relatively well understood (eg Wu 2000 and references therein) the region where accretion flow first interacts with the magnetic field of the white dwarf is not. The flow interacts with the magnetosphere in a complex manner and it is not easy to isolate stream emission from other emission sources in the system (which are generally brighter). Eclipsing systems provide an opportunity to study the accretion flow as a separate and distinct source for a short period of time, when the emission from the bright accretion region on the white dwarf is blocked by the secondary. Light curves of the eclipse contain information about the structure and the brightness distribution along the stream. The stream brightness distribution can be retrieved using an indirect imaging technique which can reconstruct the brightness of the region between the primary and the secondary. One such technique is that of Hakala (1995) who devised an indirect imaging method based on Maximum Entropy to deduce the brightness distribution along the accretion stream of HU Aqr. This technique has been developed further by Harrop-Allin et al (1999a, 1999b, 2001) who used a more physically realistic stream trajectory, improving the model's optimizing algorithm and including the projection effects. V895 Cen was discovered serendipitously using the {\sl EUVE} satellite (Craig et al 1996). Craig et al observed strong line emission and low/high brightness states and concluded V895 Cen was likely to be a polar. Its orbital period of $P_{orb}$=4.765 h (Stobie et al 1996) is the second longest period currently known for a polar and the system is found to alternate frequently between high and low accretion states. Stobie et al (1996) concluded that it is an eclipsing polar but that the secondary minimum of the ellipsoidal variation was offset with respect to the eclipse. They suggested the eclipsed component was a hot compact source which appeared to be distinct from the white dwarf, probably associated with the accretion stream. Howell et al (1997) also concluded the eclipse to be of an extended object much larger than the white dwarf, perhaps a partial accretion disc which forms during the high state. Stobie et al (1996) found no evidence for significant levels of polarization in a low accretion state. To investigate the nature of the eclipsed component, we have applied the techniques of indirect imaging and Doppler tomography to this system to study the spatial and temporal changes in the stream and to determine if the eclipse is associated with the white dwarf or the stream.
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\subsection{The eclipsed source} Stobie et al (1996) suggested that the eclipsed component was not the white dwarf. This was based on the fact that the eclipse occurred before the secondary minimum in the light curve. Assuming the secondary minimum was the true marker of inferior conjunction, and taking the timings of the eclipse ingress and egress relative to this, they concluded that the eclipsed source was $\sim$30 white dwarf radii from the white dwarf. We have found that the phasing of the eclipse ingress and egress occurs at exactly the same orbital phase in both low, intermediate and high state data (cf Figure 5). If the eclipsed source was the stream as suggested by Stobie et al (1996) we would not expect to observe the stability in the eclipse features as we do. We now address the fact that the eclipse appears off-set from the secondary minimum in the low state light curves of Stobie et al (1996). Many polars show evidence for heating of the trailing face of the secondary by the accretion region on the white dwarf. It is expected that even if the irradiation is sharply reduced or switched off, the trailing face of the secondary will remain heated for some duration. Szkody et al (1999) estimate that in the case of the polar AR UMa it takes around 5 months for the secondary to cool down to the temperature of the unheated part of the star. Indeed, from our Doppler tomography results (\S \ref{tomo}) we find that the secondary in V895 Cen is still heated when we observed it in a low accretion state. This has the effect of increasing the optical flux between $\phi$=0.5--0.9 compared to a secondary star with no irradiation. We suggest that the apparent offset between the eclipse and the secondary minimum is due to asymmetric irradiation of the secondary star. \subsection{The indirect stream mapping results} Our results show that in the intermediate/low accretion state the stream brightens rapidly as it nears the white dwarf. In our brightest state data, we find that stream emission is concentrated mainly along the field lines leading to the upper pole. At face value this suggests that as the system reaches a high enough mass transfer rate, the accretion mode goes from a two-pole to a one-pole model. It is not clear why this would be the case, but it suggests that with an increase in the mass transfer rate, the upper pole is now the more favorable pole to accrete. In all our model fits, there is no evidence for a brightening of the accretion stream at the magneto-spheric interaction region. These results are similar to that of the low accretion state data of HU Aqr (Harrop-Allin et al.~ 2001). However, there was some indication that the stream brightened at the interaction region in the $U$ and $B$ bands. Using emission line data of HU Aqr in a high accretion state and a different technique to that used here, Vrielmann \& Schwope (2001) derived stream brightness maps and found a brightening of the stream in the magneto-spheric interaction region of HU Aqr. It is possible that the accretion state (and hence amount of irradiation) plays an important role in determining whether the stream is found to brighten at the interaction region. It is interesting to compare the coupling radius that we derive from our model fits with that of HU Aqr (Harrop-Allin 1999b, 2000). They find that in a high accretion state, $R_{\mu}\sim$0.18a. We find a mean value of $R_{\mu}\sim$0.25 from our fits. HU Aqr has a magnetic field strength of 36 MG (Schwope, Thomas \& Beuermann 1993). Since we expect the accretion flow to interact with the magnetic field when the magnetic pressure equals the ram pressure of the flow, we predict that the magnetic field strength of V895 Cen is significantly larger than that of HU Aqr, other things (such as \Mdot) being equal.
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astro-ph0112364_arXiv.txt
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Recent stellar population and chemical abundance studies point to an accreted origin of $\omega$ Cen. In this light, and given the retrograde, small size orbit of $\omega$ Cen, we search for a kinematical signature left by its hypothetical parent galaxy in the Solar neighborhood. We analyze the largest-to-date sample of metal poor stars (Beers {\it et al.} 2000) and we find that, in the metallicity range $-2.0 <$ [Fe/H] $\le -1.5$, a retrograde signature that departs from the characteristics of the inner halo, and that resembles $\omega$ Cen's orbit, can be identified.
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Recent advances in understanding the nature and origin of the highly unusual globular cluster $\omega$ Centauri (see Majewski {\it et al.} 2000 for a summary of properties), are due primarily to the following findings: the multiple-peak metallicity distribution seen in the structure of the giant branch (Lee {\it et al.} 1999; Pancino {\it et al.} 2000; Frinchaboy {\it et al.} 2001), the correlation between age and metallicity (e. g., Hughes \& Wallerstein 2000, Hilker \& Richtler 2000), and the s-process enhanced enrichment in cluster stars compared to halo stars of similar metallicity (Smith {\it et al.} 2000; Vanture, Wallerstein \& Brown 1994). These findings suggest that $\omega$ Cen underwent self-enrichment with at least three primary enrichment peaks (Pancino {\it et al.} 2000, Frinchaboy {\it et al.} 2001), over a period of at least 3 Gyr (Hughes \& Wallerstein 2000). The s-process heavy-elements are primarily synthesized in low-mass (1.5 to 3.0 M$_{\odot}$) asymptotic giant branch (AGB) stars (see e.g., Travaglio {\it et al.} 1999 and references therein). In order to enrich the cluster in s-process elements, the ejecta from low-mass stars that evolve on timescales of $10^{9}$ years had to be retained by the cluster and incorporated in the next generations of stars. This long and complex star formation history is inconsistent with the cluster originating on its current orbit, which is of low energy and confined to the disk. With a period of only 120 Myr (Dinescu, Girard, \& van Altena 1999 - hereafter DGvA), the frequent disk crossings would have certainly swept out all of the intracluster gas soon after its formation, and the result would be a single-metallicity system that would resemble most of the Galactic globular clusters. It appears thus that $\omega$ Cen evolved somewhere away and independently from the Milky Way, in a system that was massive enough to retain ejecta from previous generations of stars, and to undergo multiple episodes of star formation. Its current orbit can be reconciled with the complex star formation history only if it represents a strongly decayed orbit. This, in turn, requires a massive enough system such that dynamical friction was able to drag it to the inner regions of the Galaxy. This system must have also been rather dense in order to survive the tidal field of the Milky Way and continue to loose orbital energy due to dynamical friction down to an orbit with an apocenter of the order of the Solar circle radius. The current mass of $\omega$ Cen (5 $10^{6}$ M$_{\odot}$; Meylan {\it et al.} 1995) can not generate sufficient dynamical friction to modify its orbit to its current small size (DGvA). Following these arguments, the debris from the massive putative parent galaxy of $\omega$ Cen may be expected to imprint a kinematical feature in large samples of local, metal-poor stars. The purpose of this investigation is to search for such a feature in the kinematically hot halo. We have used three data sets: the largest, kinematically unbiased sample of metal poor stars ($\sim 1200$) provided by Beers {\it et al.} (2000) (hereafter B2000), the sample of globular clusters with measured absolute proper motions (DGvA updated with new distances from Harris 1996, and with a few more clusters; see Dinescu {\it et al.} 2001), and a small sample of stars with complete kinematics and abundance measurements for O, Na, Mg, Si, Ca, Ti, Cr, Fe, Ni, Y and Ba (Nissen \& Schuster 1997, hereafter NS97).
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We have shown that a distinct population of stars with a metallicity range that inludes the mean metallicity of $\omega$ Cen, and with $\omega$ Cen-like phase-space characteristics emerges from the B2000 data. Choosing a metallicity and orbital-parameter range (Section 4) such that we maximize the ``signal'' of this population with respect to the ``noise'' of the halo, we obtain an excess population at 2.3-$\sigma$ level. By considering the RR Lyrae stars in B2000, we also see this population. We find that the excess RR Lyrae population is predominantly of RRab type, with periods of 0.5 days, and [Fe/H] $\sim -1.5$. The candidates to have been torn from the system that once contained/was $\omega$ Cen have one main orbit property: they have a larger eccentricity ($e \sim 0.8$) (i. e. orbital energy) than that of $\omega$ Cen ($e = 0.67$). Using the disruption model developed by Johnston (1998) for the orbit of $\omega$ Cen, we find that trailing tidal debris with orbit characteristics of those of the candidates are found in the Solar neighborhood. We also find that HD 106038, a single, main sequence star with a chemical abundance pattern very similar to that in $\omega$ Cen stars, in particular enhanced s-process elements (NS97), has $\omega$ Cen-like orbital properties. Similarly, V716 Oph in B2000 (a BL Her-type variable found in globular clusters of which $\omega$ Cen is most abundant) has $\omega$ Cen-like orbital properties, and a metallicity close to the mean metallicity of $\omega$ Cen. We identify two globular clusters as candidates for belonging to the system that once contained/was $\omega$ Cen, NGC 362 and NGC 6779. The more metal rich cluster, NGC 362 shows a deficiency in [Cu/Fe] when compared to globular clusters of similar metallicity, a deficiency seen so far only in $\omega$ Cen stars. {\bf Acknowledgments}. I am grateful to M\'{a}rcio Catelan for his suggestions regarding the RR Lyrae stars, and to both M\'{a}rcio Catelan and Terry Girard for numerous helpful discussions concerning this work. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France.
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astro-ph0112364_arXiv.txt
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astro-ph0112478_arXiv.txt
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We present the first calculations to follow the evolution of all stable nuclei and their radioactive progenitors in stellar models computed from the onset of central hydrogen burning through explosion as Type II supernovae. Calculations are performed for Pop I stars of 15, 19, 20, 21, and 25\,\Msun using the most recently available experimental and theoretical nuclear data, revised opacity tables, neutrino losses, and weak interaction rates, and taking into account mass loss due to stellar winds. A novel ``adaptive'' reaction network is employed with a variable number of nuclei (adjusted each time step) ranging from $\sim700$ on the main sequence to $\gtrsim2200$ during the explosion. The network includes, at any given time, all relevant isotopes from hydrogen through polonium ($Z=84$). Even the limited grid of stellar masses studied suggests that overall good agreement can be achieved with the solar abundances of nuclei between $^{16}$O and $^{90}$Zr. Interesting discrepancies are seen in the 20\, \Msun model and, so far, only in that model, that are a consequence of the merging of the oxygen, neon, and carbon shells about a day prior to core collapse. We find that, in some stars, most of the ``$p$-process'' nuclei can be produced in the convective oxygen burning shell moments prior to collapse; in others, they are made only in the explosion. Serious deficiencies still exist in all cases for the $p$-process isotopes of Ru and Mo.
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\label{sec:intro} The nucleosynthetic yields of massive stars are important to many areas of astronomical research. Besides the inherent interest in understanding our nuclear origins, the abundances made in supernovae are used to diagnose models for the explosion and as input to still grander models for the formation and chemical evolution of galaxies and the intergalactic medium. They are the target of x-ray observations of supernova remnants and gamma-ray studies of radioactivities in the interstellar medium. Some can be used as cosmochronometers, others power the light curves, still others appear as anomalous abundances found in tiny meteroitic grains in our own solar system. For these reasons, nucleosynthesis calculations have a long history and a sizable community that carries them out. Most recently, nucleosynthesis in massive stars has been studied by \citet[WW95]{WW95}; \citet{TNH96,lim00} and others. With this paper, we embark on a new survey, similar to WW95, that will ultimately include stars of many masses and initial metallicities. The characteristics of this new study are improvements in the stellar physics (mass loss rates, opacities, reaction network, etc., \Sect{comp_proc}) and revisions to nuclear reaction rates (\Sect{mods}) that have occurred during the last eight years. This first paper particularly addresses recent improvements in nuclear physics. For elements heavier than about silicon, the nuclear level densities are sufficiently high (provided the particle separation energies are not too small) that the statistical - or ``Hauser-Feshbach'' - model can be used. Here, in their maiden voyage, we use rates calculated using the NON-SMOKER code \citep{rtk97,rt98}. The reaction library, from which the network is drawn, includes all nuclei from the proton-drip line to the neutron-drip line and elements up to and including the actinides \citep{RATH}. For elements lighter than silicon, where they have been measured, results are taken from the laboratory. Several different compilations are explored. The most critical choices are the rates for $^{12}$C($\alpha,\gamma)^{16}$O, $^{22}$Ne($\alpha$,n)$^{25}$Mg, and $^{22}$Ne($\alpha,\gamma)^{26}$Mg. In order to facilitate comparison, we have chosen a constant value equal to 1.2 times that of Buchmann (1996) for the $^{12}$C($\alpha,\gamma)^{16}$O rate in {\sl all} our calculations. For our {\sl standard} models (defined in $\S$3.1) we further adopt the lower bound of \citet{kaepp94} for $^{22}$Ne($\alpha$,n)$^{25}$Mg \citep{HWW01}. In future publications we will explore, in greater depth, the consequences of different choices for these rates (for $^{12}$C($\alpha,\gamma)^{16}$O, see also \citealt{WW93,BHW02}). A novel reaction network is employed, unprecedented in size for stellar evolution calculations. The network used by WW95, large in its day, had about 200 nuclides and extended only to germanium. Studies using reaction networks of over 5000 nuclei have been carried out for single zones or regions of stars in order to obtain the $r$-process, e.g., \citet{CCT85,fre99,kra93}, but ``kilo-nuclide'' studies of nucleosynthesis in complete stellar models (typically of 1000 zones each for 20,000 time steps) have not been done before. We describe in \Sect{dynet} a dynamically evolving network that adds and subtracts nuclides as appropriate during the star's life to ensure that all significant nuclear flows are contained. Our present survey uses a network that has the accuracy of a fixed network of 2500 isotopes. Section 4 discusses aspects of the stellar evolution that are critical to the nucleosynthesis and \Sect{results} gives the main results of our survey. We find overall good agreement of our nucleosynthesis calculations with solar abundances for intermediate mass elements (oxygen through zinc) as well as the ``weak component'' of the $s$-process (A $\ltaprx$ 90), and most of the $p$-process isotopes. However, there is a systematic deficiency of $p$-process isotopes below A $\approx$ 125 that is particularly acute for Mo and Ru, and around A $\approx$ 150. Possible explanations are discussed in \Sect{gamma}. We also find that the nucleosynthesis is at least as sensitive to the stellar model as to the nuclear physics and, in particular, find unusual results for a 20\,\Msun model (in the sense that the results differ greatly from both the sun and those at either 19 or 21\,\Msun). This is because of the merging of convective oxygen, neon, and carbon shells that occurred well before collapse in that model and not in the others (\Sect{results}).
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Using a nuclear reaction network of unprecedented size, nucleosynthesis has been investigated in several stellar models in the mass range 15 \Msun to 25 \Msun. The models include the best currently available nuclear and stellar physics. For the first time, it was also possible to self-consistently follow the $\gamma$-process up to Bi. Overall good agreement can be achieved with the solar abundances of nuclei between $^{16}$O and $^{90}$Zr. This good agreement is, to first order, independent of the reaction rate set employed; our current standard, \citet{ang99} or \citet{HWW01}, though several key nuclear uncertainties are identified. In addition to the well-known need for greater accuracy in the rate for $\alpha$-capture on \I{12}C, the rates for \I{22}{Ne}($\alpha$,n)\I{25}{Mg} and $^{22}$Ne($\alpha,\gamma)^{26}$Mg are critical. We also urge a re-examination of some of the neutron capture cross sections for the isotopes of nickel. For the $p$-isotopes, two regions of atomic mass are found where those isotopes are underproduced, $92 \leq A \leq124$ and $150\leq A\leq 165$. It remains unclear whether this deficiency is due to nuclear cross sections, stellar physics, or if alternative (additional) $p$-process scenarios have to be invoked. However, we find that part of the $p$-nuclides may be produced in convective oxygen shell burning during the last hour of the star's life. The remainder is made explosively. Interesting and unusual nucleosynthetic results are found for one particular 20 \Msun model due to its special stellar structure. This effect, a merging of heavy element shells late in the stars evolution, seems to be confined to a narrow range of masses. In particular it is not seen in 19 and 21 \Msun models. However, we have explored a very limited set of masses and those only in one spatial dimension (for caveats see Bazan \& Arnett 1994).
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astro-ph0112478_arXiv.txt
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astro-ph0112152_arXiv.txt
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We present optical observations of the fields of two X--ray sources located near the center of the shell-like supernova remnant G266.1--1.2. No objects brighter than R$\sim$22.5 and B$\sim$23 are present within the small \textit{Chandra} error region of \axj, besides a R$\sim$17 star that has already been excluded as a possible counterpart. A bright diffuse H$_{\alpha}$ nebula is present close to the position of the candidate neutron star.
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The supernova remnant G266.1--1.2 has been reported as a possible $\gamma$-ray source in the 1.156 MeV line of $^{44}$Ti (Iyudin et al. 1998). The short lifetime ($\sim$90 yrs) of this isotope, and the relatively small angular size of the remnant would imply an age of only $\sim$680 years and a small distance d$\sim$200 pc (Aschenbach et al. 1999). Thus G266.1--1.2 could be the remnant of the closest supernova event to have occurred in recent historical times. However, \textit{ASCA} observations showed that the X--rays from the SNR shell have a non-thermal spectrum and the fits require a high absorption value (Slane et al. 2001), favoring a distance of $\sim$1-2 kpc that would place G266.1--1.2 well beyond the Vela SNR (see also Mereghetti \& Pellizzoni 2001). The \textit{ASCA} data revealed also a central point source, \axj , surrounded by diffuse X--ray emission, that was interpreted as the neutron star associated to G266.1--1.2. A \textit{BeppoSAX} observation (Mereghetti 2001) of the central region of G266.1--1.2 showed the presence of a second source about 3$'$ north of that detected by \textit{ASCA} and with a harder spectrum. The northern source was named SAX~J0852.0--4615. Since the \textit{BeppoSAX} error circle of \axj contained two bright early type stars that might have produced the observed X--ray flux, while no optical counterparts brighter than V$\sim$15 were visible for \sax, it was unclear which of the two sources was the most likely neutron star candidate. The puzzle has been recently solved by a \textit{Chandra} observation that provided an arcsecond position for \axj (Pavlov et al. 2001). The new error box is incompatible with the two early type stars that were previously considered as possible counterparts, thus confirming that \axj is the most likely neutron star candidate. \sax was not detected in the 3 ks long \textit{Chandra} observation reported by Pavlov et al. (2001). This might be due to variability, or to the hardness of this source, that was detected with \textit{BeppoSAX} only above 5 keV. A deeper observation with \textit{XMM-Newton} confirmed the existence of \sax, with a flux about ten times fainter than that of \axj (Aschenbach, this conference). Here we present optical observations of the fields of these two X--ray sources.
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The deep optical limits for the possible counterparts of \axj confirm that this is most likely the neutron star remnant associated with G266.1--1.2. An interesting H$_{\alpha}$ nebula has also been discovered in the data presented here. Emission in the H$_{\alpha}$ has been detected around a few radio pulsars and is thought to originate in the interstellar medium shocked by the relativistic pulsar wind. These nebulae have either a ''cometary'' shape with the axis of symmetry along the direction of the pulsar transverse motion (e.g., PSR B2224+65 (''Guitar Nebula'', Cordes et al. 1993) or PSR B0740--28 (Stappers et al. these proceedings)) or an arc-like shape (e.g., PSR J0437--4715, Bell et al. 1996). The morphology of the diffuse emission shown in Fig.1 and Fig.2 does not present any obvious connection with the location of the candidate neutron star as determined with \textit{Chandra}. It is more likely that the nebula is related to the B[e] star Wray 16-30, which is located at the southern end of the nebula. However, its peculiar morphology and the location close to the center of G266.1--1.2 make this nebula a potentially interesting target for more detailed investigations.
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astro-ph0112152_arXiv.txt
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astro-ph0112222_arXiv.txt
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In this paper I investigate what factors -- both observational and physical -- can change the measured slope of the observed 21cm HI power spectrum. The following effects can make the observed turbulence appear two dimensional rather than three dimensional: 1) if the turbulence is contained in a thin filament or slab; 2) if the medium has a high optical depth; and 3) if any method of observation or analysis is used which effectively limits the emission from the medium under study to a thin slab, for example, by analyzing an individual channel map. Straightforward analysis of data can give misleading or incomplete results if these effects are not taken into account.
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The 21cm HI line has a column density power spectrum whose slopes are consistent with 2--Dimensional turbulence ($\alpha \sim -8/3$) on large spatial scales ($> 0.01$~pc) and narrow velocity ranges (Green 1993; Dickey \& Crovisier 1983; Lazarian \& Stanimirovic 2001; Dickey et al. 2001). The slope of this power spectrum is closer to a 3--D, Kolmogorov--like spectrum ($\alpha \sim -11/3$) for wider velocity ranges but can still be significantly different than the Kolmogorov value. However, the electron density spectrum is consistent with Kolmogorov--like turbulence. This suggests a) that electrons and neutrals have different turbulent characteristics, or b) that there are effects which change the measured power spectrum slopes. Lazarian \& Pogosyan (2000) have discussed how turbulent velocity may effect the observed power spectra of HI. However, there are other factors which Lazarian \& Pogosyan did not discuss which may effect the power spectra. Among these are opacity and filamentary structures in the HI.
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astro-ph0112222_arXiv.txt
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astro-ph0112291_arXiv.txt
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We report upper limits on CO J=2--1 and CO J=5--4 emission from the quasar SDSSp J104433.04-012502.2 at $z=5.73$ from observations made with the Berkeley-Illinois-Maryland-Association Array. Previously reported limits on CO J=6--5 emission (Iwata et al. 2001) were obtained at $z=5.80$, which is now thought to be off by 1\%, and the observations likely missed the relevant redshifts for molecular gas. The new $3\sigma$ upper limits on the line luminosities are $L^{'}_{CO}(2-1) < 5.1 \times 10^{10}$ K~km~s$^{-1}$~pc$^2$ and $L^{'}_{CO}(5-4) < 3.0 \times 10^{10}$ K~km~s$^{-1}$~pc$^2$, assuming 200~km~s$^{-1}$ linewidth. The CO J=5--4 observations place an upper limit on warm, dense molecular gas mass comparable to amounts derived for some other high redshift quasar systems from detections of this line. The limit on CO J=2--1 emission suggests that excitation bias does not affect this conclusion. In addition, no molecular gas rich companion galaxies are found in a $\sim1.4$~Mpc field surrounding the quasar.
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The study of the star formation properties and gas content of galaxies at far cosmological distances is one important step toward understanding galaxy formation and evolution. The quasar SDSSp~J104433.04-012502.2 at $z\approx5.8$ (hereafter SDSS~1044-0125), discovered by Fan et al. (2000) using the Sloan Digital Sky Survey, is among the highest redshift objects known. Recent observations of this quasar with SCUBA on the JCMT at 850~$\mu$m detect thermal continuum emission (reported by Iwata et al. 2001), which suggests a large reservoir of dust and therefore also molecular gas. The presence of a substantial gas mass gains support from the apparent X-ray weakness of the quasar, which likely results from heavy intrinsic absorption (Brandt et al. 2001, Mathur 2001). Molecular gas has been detected from at least a dozen $z>2$ objects through CO lines at millimeter wavelengths, and these observations provide important clues to the formation history of galaxies and their relationship to supermassive black holes. A recent search for CO J=6--5 emission from SDSS~1044-0125 by Iwata et al. (2001) reported an upper limit on the inferred molecular gas mass comparable to the detections for some high redshift quasars. Unfortunately, the search was centered at $z=5.80$, the initial redshift estimated by Fan et al. (2000). Recent spectroscopic studies of SDSS~1044-0125 give a more accurate value of $z=5.73\pm0.01$ (Djorgovski et al. 2001, see also Goodrich et al. 2001), about 1\% off from the initial estimate. Because of the narrow instantaneous bandwidth available to current millimeter interferometers, the revised redshift falls outside window that was searched for CO J=6--5 emission, and the observations likely missed the relevant redshifts for molecular gas in the quasar host. Since SDSS~1044-0125 has several properties in common with high redshift quasars where CO emission has been detected, (enumerated by Iwata et al. 2001), the revision of the optical redshift determination gives impetus to a new search for molecular gas. A potentially important limitation of searching for emission from CO lines with high rotational quantum numbers, like the J=6--5 transition, is that prevailing physical conditions may be insufficiently extreme to excite these lines. The surprising detection of extended emission in the low excitation CO J=2--1 line towards the quasar APM~08279+5255 at $z=3.91$ (Papadopoulos et al. 2001) suggests that low excitation CO lines can reveal molecular mass reservoirs that are one or two orders of magnitude larger than suggested by observations of high excitation CO lines. In this short paper, we present results of searches for CO J=2--1 and J=5--4 emission from SDSS~1044-0125 using the Berkeley-Illinois-Maryland Array (BIMA)\footnote{The BIMA array is operated by the Berkeley-Illinois-Maryland Association under funding from the National Science Foundation.} (Welch et al. 1996) that provide new limits on the amount of molecular gas associated with this luminous high redshift quasar. The BIMA 1~cm band receiver system, which was developed primarily for observations of the Sunyaev-Zeldovich effect (Carlstrom, Joy \& Grego 1996, Grego et al. 2000), provides a unique facility to search for highly redshifted low lying CO lines. For SDSS~1044-0125, the CO J=2--1 line is redshifted to 34 GHz, within the accessible tuning range. The standard digital correlator allows for several times larger velocity coverage than generally available at shorter wavelengths, sufficient to span the uncertainty in the quasar redshift determined from optical lines, as well as the typical kinematic offsets of molecular gas from the redshift derived optically. In addition, the small BIMA array antennas provide a large field of view, which enables imaging the quasar environs over Mpc scales at 34 GHz in a single pointing.
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Figure~\ref{fig:spectrum} shows the CO J=2--1 and CO J=5--4 spectra obtained at quasar position. The velocity binnings for the two lines are 200 and 162~km~s$^{-1}$, respectively. The noise in the CO J=2--1 spectrum is not uniform, and empirical $\pm1\sigma$ error bars derived from the rms noise measured from the images are shown for each velocity bin. There are some tantalizing hints of signal in adjacent channels close to the expected velocities, but features with similar (low) significance are present elsewhere in the data, and we do not consider any of these features to be reliable line detections. Figure~\ref{fig:channels} shows a series of maps with 200~km~s$^{-1}$ width that span the full half power field of view ($6\farcm6$) for the J=2--1 line. Various attempts at smoothing in both space and frequency did not uncover any significant CO emission in either of the observed transitions. Two mechanisms have been suggested for heating large masses of dust and gas in high redshift quasar systems: (1) high energy photons emitted from gases accreted onto a massive black hole and (2) bursts of star formation. If the dust is heated by the activity of a massive black hole, then bright emission may be expected from high excitation CO lines in a compact region close to the exciting source, from large amounts of warm, dense gas involved in fueling and accretion. The CO J=5--4 line, whose upper energy level lies 88~K above the ground state, requires warm gas ($>30~K$) at high densities ($>10^3$~cm$^{-3}$) to be populated significantly by H$_2$ collisions. Consequently, the upper limits on CO emission from the J=5--4 line constrains primarily the amount of molecular gas with these conditions close to the massive black hole or other powerful heating sources. On the other hand, such extreme physical conditions are not necessarily appropriate for starbursts, which are likely to be distributed over larger spatial scales and involve cooler, more diffuse molecular gas. If the dust is heated by primarily by star formation, then emission in CO J=2--1 line may be a more appropriate tracer of molecular gas content, given that the upper energy level lies just 16~K above ground and the excitation requirements are significantly less stringent. Following Solomon, Downes \& Radford (1992), we calculate upper limits to CO line luminosities with the expression \begin{equation} L^{'}_{CO} = 3.25\times10^{7} S_{CO} \Delta v \nu_{obs}^{-2} D_L^2 (1+z)^{-3}~~~~ {\rm K~km~s}^{-1}~{\rm pc}^2, \end{equation} where $S_{CO} \Delta v$ is the limit on the velocity integrated line flux in Jy~km~s$^{-1}$, $\nu_{obs}$ is the observing frequency in GHz, and $D_L$ is the luminosity distance in Mpc. The choice of cosmological parameters enters in $D_L$, and we adopt $H_0 = 75$ km~s$^{-1}$, $\Omega = 1$ and $\Omega_{\Lambda} = 0$ for consistency with most work in this field. (An alternative cosmology with $H_0 = 75$ km~s$^{-1}$, $\Omega = 1$ and $\Omega_{\Lambda} = 0.7$ results in $D_L$ larger by a factor of 1.54 for this redshift.) The effective linewidth is not known, but it likely falls in the range 150 to 550~km~s$^{-1}$ found for a large sample of ultraluminous galaxies in the local universe (Solomon et al. 1997). For the $3\sigma$ flux limit obtained in the more sensitive part of the CO J=2--1 spectrum, assuming a linewidth of 200~km~s$^{-1}$, $L^{'}_{CO}(2-1) < 5.1 \times 10^{10}$ K~km~s$^{-1}$~pc$^2$. For the $3\sigma$ flux limit obtained for CO J=5--4, again assuming a linewidth of 200~km~s$^{-1}$, $L^{'}_{CO}(5-4) < 3.0 \times 10^{10}$ K~km~s$^{-1}$~pc$^2$. If the assumed linewidth were two times larger, then these luminosity limits would be $\sqrt{2}$ times higher. Conversion of these CO luminosity limits to molecular gas mass limits is fraught with uncertainties. But a simple conversion factor from CO luminosity to H$_2$ mass is commonly taken to be $4.5~M_{\odot}$ (K~km~s$^{-1}$~pc$^2$)$^{-1}$, the value determined for Milky Way molecular clouds (Sanders, Scoville \& Soifer 1991). There is evidence from comparisons of luminosity based mass estimates with dynamical mass estimates that the conversion factor may be perhaps five times lower in ultraluminous objects (Downes \& Solomon 1998). Additional corrections of order unity are also needed to account properly for excitation from the elevated cosmic background radiation at high redshift. Adopting the Galactic conversion factor for CO J=2--1 line luminosity gives a limit on the {\em cold or diffuse} molecular gas mass of $\sim2.3\times10^{11}~M_{\odot}$ in the SDSS~1044-0125 system. Using the same conversion factor for the CO J=5--4 line luminosity gives a limit on the {\em warm and dense} molecular gas mass of of $\sim1.3\times10^{11}~M_{\odot}$ in the SDSS~1044-0125 system. These mass limits are comparable to the mass indicated from the detection of CO J=5--4 emission from some $z>4$ quasars, including at least two thought not to be amplified by gravitational lensing. In particular, observations of CO J=5--4 emission from BR1202-0725 at $z=4.7$ (Omont al. 1996, Ohta et al. 1996) and BRI1335-0417 at $z=4.4$ (Guilloteau et al. 1997) indicate molecular gas masses in excess of $10^{11}$~M$_{\odot}$ (adjusted for the cosmology and CO to H$_2$ conversion factor adopted here). There is no clear physical argument to explain why some quasar environments show CO emission at this sensitivity level while others do not (Guilloteau et al. 1999). In any case, the CO J=2--1 and J=5--4 luminosity limits suggest that the environment of SDSS~1044-0125 does not possess an enormous mass reservoir of either low excitation or high excitation molecular gas. The CO J=2--1 limit is comparable to the amount of molecular gas detected toward the lensed quasar APM~08279+5255, where Papadopoulos et al. (2001) found several CO J=2--1 emission features with total luminosity $6.6\pm3.1 \times 10^{11}$ K~km~s$^{-1}$~pc$^2$ attributed to (unlensed) molecular gas rich companion galaxies to the quasar host. For the SDSS~1044-0125 observations, such features would have been contained within one synthesized beam (together with any nuclear emission). The luminosity limit suggests that no comparable population of nearby massive companions is present. Moreover, no significant CO J=2--1 emission features are found within the entire field of view that spans $\sim1.4$~Mpc, which suggests that such massive cold molecular gas concentrations are rare. Observations of SDSS~1044-0125 with better sensitivity are needed to explore whether smaller but still significant concentrations of low excitation molecular gas are present in the environment of this high redshift quasar.
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astro-ph0112291_arXiv.txt
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astro-ph0112258_arXiv.txt
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The reduced proper motion diagram (RPMD) for a complete sample of 819 faint ($B \leq 22.5$) stars with high accuracy proper motions ($\sigma_{\mu}\sim1$ mas yr$^{-1}$) in an area of 0.3 deg$^2$ in the North Galactic Pole field SA57 is investigated. Eight stars with very large reduced proper motions are identified as faint white dwarf candidates. On the basis of larger than 6$\sigma$ measured proper motions and the lack of photometric variability over a twenty year baseline, we discriminate these white dwarf candidates from the several times more numerous QSOs, which can potentially occupy a similar location in the RPMD. For comparison, less than 4$\sigma$ proper motions and photometric variability are found in all but one of 35 spectroscopically confirmed QSOs in the same field. While spectroscopic confirmation of their status as white dwarfs is a necessary, but difficult, outstanding task, we discuss the implausibility that these stars could be any kind of survey contaminant. High quality proper motions lend confidence in our ability to separate white dwarfs from subdwarfs in the RPMD. If {\it bona fide} white dwarfs, the eight candidates found here represent a portion of the white dwarf population that hitherto has remained uninvestigated by previous surveys by virtue of the faint magnitudes and low proper motions of the stars. This faint, low velocity sample represents an increase in the white dwarf sky surface density to $B=22.5$ by an order of magnitude over that found in the previously most complete surveys to this depth. However, because the majority of the stars discovered here are at projected distances of more than a disk scaleheight above the Galactic midplane, their existence does not affect significantly the typical estimates of the local white dwarf density. On the other hand, as distant white dwarf candidates with low, typically thin disk-like transverse velocities ($< 40$ km s$^{-1}$), the newly discovered stars suggest a disk white dwarf scaleheight larger than the values of 250-350 pc typically assumed in assessments of the local white dwarf density (and thought to characterize the Galactic old thin disk in stellar population models). Both a $<V/V_{max}>$ and a more complex maximum likelihood analysis of the spatial distribution of our likely thin disk white dwarfs yield scaleheights of 400-600 pc while at the same time give a reasonable match to the {\it local} white dwarf volume density found in other surveys (although this good match is a result of the dominance of the one relatively nearby white dwarf in the $1/V_{max}$ density calculation). A high scaleheight persists even if the relatively small sample is pruned of any potential thick disk or halo white dwarfs. While our work is not optimized toward the study of halo white dwarfs as potential MACHO objects, our results do have interesting implications for this hypothesis. We can place some direct constraints (albeit weak ones) on the contribution of halo white dwarfs to the dark matter of the Galaxy. Moreover, the elevated scale height that we measure for the thin disk could alter the interpretation of microlensing results to the extent of making white dwarfs untenable as the dominant MACHO contributor.
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In Paper I of this series (Majewski 1992), proper motions were determined for nearly a thousand stellar objects in Selected Area 57 (SA57) at the North Galactic Pole to photographic $B_J \sim 22.5$ and $V_F \sim 21.5$.\footnote{$B_J$ is the passband produced by the combination of IIIa-J emulsion and GG385 filter, whereas $V_F$ is the combination of IIIaF + GG495.} Photometric parallaxes were determined for a subsample of 250 stars with $0.3 \leq B-V \leq 1.1$ and $U \leq 21.5$ based on photographic ultraviolet excess measurements. Since, in general, no direct measurement of the surface gravity of each star was readily available, a basic premise of the adopted analysis in Paper I was that the survey stars are on the main sequence. However, it is possible to exploit proper motions to discriminate luminosity classes of some stars through use of the reduced proper motion diagram (RPMD; see also Luyten 1922, Jones 1972a,b, Chiu 1980b, Evans 1992, Knox et al. 1999, hereafter K99; Cooke \& Reid 2000; Oppenheimer et al. 2001, hereafter O01). White dwarfs, on account of their very high reduced proper motions, should be readily identifiable in the RPMD. This technique confers certain advantages over color searches for white dwarfs; e.g., it is possible to identify cool white dwarfs that are not distinguishable from the more numerous late type field stars using colors alone. Deep searches for faint, cool white dwarfs are important for testing white dwarf cooling models into the regime of Debye crystallization, and, by applying cooling theory in conjunction with the white dwarf luminosity function, to set limits on the star formation history and age of Galactic stellar populations. White dwarfs can also be used as tracers of the density laws of old populations, and white dwarfs are proposed as potentially significant contributors to the dark matter component represented by gravitational microlensing events. Numerous studies have attempted to establish the local density and/or luminosity function of Population I white dwarfs, and especially, recently, at the red end of the white dwarf sequence, due to the interest in cool white dwarfs for both age dating the Galaxy and as a primary source for microlensing candidates. Results for the derived local white dwarf density found among the different surveys still range by a factor of two (Fleming et al. 1986; Jahrei{\ss} 1987; Liebert, Dahn \& Monet 1988, hereafter LDM; Boyle 1989; Ruiz \& Takamiya 1995; Oswalt et al. 1996; Festin 1998; K99; Reid et al. 2001; Ruiz \& Bergeron 2001). The question of completeness lingers when considering the results of these various surveys. So too does a proper understanding of the density laws appropriate to the samples garnered, since the conversion from a survey list to a local density requires an understanding of the {\it effective} volume surveyed, i.e., the volume foreshortened by the drop-off in density with distance from the Galactic midplane. Typically, white dwarf studies have {\it adopted} a standard value for an exponential scaleheight of the disk in such calculations, rather than attempted to {\it solve} for the density law from their white dwarf samples. This understandable reluctance derives from the relatively limited range of distances probed by complete samples (typically one third to one half of the traditional old disk scaleheight), which limits sensitivity to the form of the density law. Survey incompleteness at distances comparable to a disk scaleheight derives both from photometric {\it and} astrometric limitations, since proper motions provide the most commonly used means by which to identify white dwarfs (especially those redward of the field star main sequence turn off (MSTO)). Table 1 summarizes the major {\it astrometric} white dwarf surveys to date (not including studies made from archival survey data, such as the Lowell Proper Motion or Luyten Half-Second catalogues, which we represent by the work of LDM). With the exception of the deep, small area study by Chiu (1980b), these surveys are focused on stars with fairly large proper motions ($\gtrsim40$ mas yr$^{-1}$). Such a limitation progressively excludes white dwarf populations with ever larger ranges of transverse velocity as a function of distance. Figure 1 shows the limiting distances that are imposed on the detection of white dwarfs as a function of various apparent magnitude and proper motion limits. Figure 2 plots the sky density of detected white dwarfs against both photometric and astrometric limits for the surveys listed in Table 1. Table 1 and Figure 2 demonstrate that the detected white dwarf sky density appears to be more directly correlated with proper motion limits than survey depth. This is an important point, one worth considering given the new emphasis on properly accounting for the total white dwarf density in the foreground of lensed sources. The impressive K99 study, as the deepest, large area survey with the best proper motions to date, provides a benchmark for the present discussion. K99 claim to find no evidence of incompleteness in their survey sample, and that this survey represents the most complete large area sample to date is evidenced by their finding the largest sky surface density of white dwarfs for such a survey to date, 2.07 deg$^{-2}$. Through a variety of arguments, K99 suggest that their proper motion limit of 50-60 mas yr$^{-1}$ provides a reasonable compromise between minimizing spurious detections and maintaining completeness. For example, for bright ($R \sim 14$) white dwarfs, K99 argue that such a limit is more than enough to detect stars ``having a ({\it conservative}) tangential velocity of 40 km s$^{-1}$, [which] would have a proper motion of $\sim 80$ mas yr$^{-1}$" (emphasis added); clearly for fainter, more distant ($> 100$ pc) examples of stars with similar transverse velocities, however their survey quickly becomes incomplete (Figure 1). It is worth noting at this point that the results of the present analysis will focus on the discovery of white dwarf candidates that are primarily {\it slower} than K99's ``conservative" estimate. An important means by which K99 attempt to build their case for completeness is by establishing that their sample mean $(V/V_{max})$ statistic is nearly 0.5 (see \S 4), where here the $V$'s represent effective volumes {\it under the assumption of a 300 pc scaleheight}. While K99's investigation of variations in the assumed scaleheight show no significant alteration in their derived luminosity function, the authors do not state how varying the scaleheight affects their assumption of completeness. $<V/V_{max}>$ is primarily a test of uniformity that K99 and others have adapted to a test of completeness. Even K99 admit that obtaining $<V/V_{max}>=0.5$ cannot be regarded as proof of completeness, since an incomplete sample can yield a similar result. While this caveat might be considered all the more prescient if it can be shown that 300 pc is not a proper scaleheight to assume in the effective volume calculations in the first place, in the end, it should be noted that incompleteness in typical magnitude- and proper motion-limited samples may not significantly bias the derivation of the local luminosity function according to the Monte Carlo simulations of Wood \& Oswalt (1998) and Mendez \& Ruiz (2001). In this article, we explore the question of the disk white dwarf population density distribution using a low limit proper motion selected sample. We employ reduced proper motions to separate degenerate star candidates from subdwarfs and Population I main sequence stars within the deep proper motion sample of Paper I. Of course, these proper motion techniques have been used in many previous studies, some that explore to similar depth and/or much larger areas of sky than the Paper I sample (Table 1); given our small survey area (0.3 deg$^2$), the volume we probe is much smaller than that explored by most previous surveys, even considering our $B_J=22.5$ magnitude limit. Within the last few years, several new surveys have also reached $B\sim22.5$, and the K99 survey in particular has the potential, {\it based on photometric considerations}, to probe much larger volumes than our survey. In practice, however, differences in {\it astrometric quality} puts the present analysis in a unique niche in parameter space compared to all previous studies: No study in the literature has an astrometric precision comparable to that afforded by the 16 year baseline study using deep photographic exposures on fine-grained emulsions with good plate scale discussed in Paper I. The resultant precision ($\sim1$ mas yr$^{-1}$ at $B_J=21.5$ and $\sim1.6$ mas yr$^{-1}$ at $B_J=22.5$) allows for (1) an RPMD that is relatively ``clean" of astrometric error-induced scatter at the white dwarf locus, minimizing contamination problems, and (2) orders of magnitude smaller proper motion limits. Thus, we can find low proper motion white dwarf candidates that would be missed by all previous surveys, look for white dwarfs at larger distances, and ensure a much higher level of completeness than could be claimed before. The transverse velocity limits of a 1 mas yr$^{-1}$ survey to $B_J=22.5$ as a function of distance are shown in Figure 1; our astrometric advantage for probing to much larger distances than other surveys is clear. As we shall show, this advantage allows us to find a white dwarf sky density at $B_J=22.5$ that is likely an order of magnitude larger than that found by K99, which itself had a density that was larger than found by most previous surveys (see Table 1). The primary contribution to our sky density seems to be from an extended distribution of distant white dwarfs with cold kinematics typical of the thin disk. That the majority of the white dwarfs have projected distances larger than 300 pc, and a third of them beyond 900 pc, suggests the need for a revision of the normally assumed 250-350 pc scaleheight for the thin disk white dwarf population. Thus, a principal endeavor in this contribution is to use this admittedly small, but much less kinematically biased, sample of white dwarfs to define better the vertical distribution of the old disk, a task to which our unusually distant sample has particular leverage.
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Under the assumption that all proposed candidates are {\it bona fide} white dwarfs, the results of the present RPMD analysis of the Paper I survey has yielded a sky density of likely white dwarfs higher by an order of magnitude over the previous most complete samples. Our candidates represent the (expected) low velocity component of the disk white dwarf population excluded by previous proper motion searches. Several of our candidate white dwarfs are fainter and redder than the disk and halo MSTO, so that previous photometric (e.g., UV excess) surveys would not easily have found them. Our results suggest substantial astrometric and photometric incompleteness in previous surveys. Although Wood \& Oswalt (1998) and Mendez \& Ruiz (2001) have shown that incompleteness does not significantly bias derivation of the luminosity function or local density (and we concur that the incompleteness we describe here does not likely bias derivation of the local density), they do note that deriving star formation histories from biased samples is highly susceptible to error when the proper motion errors are large ($>$ 100 mas yr$^{-1}$). The tenfold gain in completeness in a large area survey with the photometric depth and astrometric precision of ours would improve the resolution with which star formation histories could be delineated from the white dwarf luminosity function. Such a survey, however, would require long time baseline observations at good plate scale. It is possible that many repeat observations over the course of the Sloan Digital Sky Survey could provide this level of precision for $V > 20$. Of course, HST can achieve such precision, but not over a large area, while the planned FAME and GAIA astrometric missions will deliver the proper motions, but not the depth. Our distant white dwarf candidates, while a small sample, provide leverage on the density law well above the Galactic midplane and suggest a higher white dwarf scaleheight than typically assumed, where the ``old disk scaleheight" of 250-350 pc falls at the very low end of the ``reasonable" range of scaleheights derived from the entire candidate white dwarf sample. It is noted that the lowest luminosity white dwarf candidates contribute the highest $V'/V'_{max}$ on average, and when they are excluded the derived scaleheight is lowered, though it remains high compared to white dwarf studies at brighter magnitudes. That we should find white dwarfs to have higher scaleheights than non-degenerate stars seems, at first blush, consistent with white dwarf cooling theory: One might expect the proportion of old stars among white dwarfs to be higher than among unevolved late-type stars and we might expect the highest vertical velocity dispersions for the oldest stars due to secular dynamical heating processes that progressively increase vertical velocity dispersions of stars with time. Yet, the notion of a relatively more heated white dwarf population appears to be at odds with the actual kinematics we measure for our sample of candidate white dwarfs. One might ask what such a dynamically cold population is doing at such large $z$.\footnote{Note that postulating overestimates of the distances to our stars does not fix the problem because moving the stars closer also make them dynamically colder in the derived transverse velocities.} We note that our velocity data do not fall along the Str\"omberg asymmetric drift relation (Binney \& Merrifield 1998) expected for secularly heated disk stars. A larger sample is needed to check this apparent contradiction. We have found that the method of maximum likelihood provides lower scale heights that are more consistent with (although still generally higher than) those found in the existing literature. Maximum likelihood is also far less sensitive to small number fluctuations. We propose use of this more elegant method of analysis to supplement or replace $<V/V'_{max}>$ methods in the future, especially when dealing with small samples. Our statements here must be tempered by two shortcomings of our survey. First of all, it is clear that a larger sample of faint white dwarf candidates with low velocities (requiring more precise proper motions) is needed to better constrain the white dwarf scaleheight, and we hope to increase our sample when additional fields with similar plate material are analyzed. In addition, spectroscopic confirmation of the present and any future deep samples of astrometrically identified white dwarfs would provide much stronger confidence in the interpretation of our results. Radial velocities, if obtainable, will be critical to verifying the kinematics of this sample, but require 6-10 meter class telescopes to obtain. It is hoped that these additional data will help resolve outstanding questions on the spatial distribution of disk white dwarfs.
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astro-ph0112258_arXiv.txt
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astro-ph0112544_arXiv.txt
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We report on spectral and timing observations of the nearest millisecond pulsar J0437--4715 with the \chan\/ X-ray Observatory. The pulsar spectrum, detected up to 7~keV, cannot be described by a simple one-component model. We suggest that it consists of two components, a nonthermal power-law spectrum generated in the pulsar magnetosphere, with a photon index $\gamma\approx 2$, and a thermal spectrum emitted by heated polar caps, with a temperature decreasing outwards from 2 MK to 0.5 MK. The lack of spectral features in the thermal component suggests that the neutron star surface is covered by a hydrogen (or helium) atmosphere. The timing analysis shows one X-ray pulse per period, with a pulsed fraction of about 40\% and the peak at the same pulse phase as the radio peak. No synchrotron pulsar-wind nebula is seen in X-rays.
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Millisecond (recycled) radio pulsars are distinguished from ordinary pulsars by their very short and stable periods, $P\la 10$~ms, $\dot{P}\sim 10^{-21}-10^{-19}$~s~s$^{-1}$. It is generally accepted that they are very old objects, with spin-down ages $\tau=P/2\dot{P}\sim 10^9-10^{10}$~yr and low surface magnetic fields $B\propto (P\dot{P})^{1/2}\sim 10^8-10^{10}$~G (e.g., Taylor, Manchester, \& Lyne 1993). Similar to ordinary pulsars, a millisecond pulsar can emit nonthermal X-rays from its magnetosphere, with a hard power-law spectrum and sharp pulsations. In addition to this nonthermal radiation, thermal X-rays can be emitted from the neutron star (NS) surface, provided the surface is hot enough. According to the models of NS thermal evolution (albeit rather uncertain at these old ages), recycled pulsars are too cold (surface temperature $T\la 0.1$~MK --- see, e.~g., Tsuruta 1998) to be detectable in X-rays. However, their polar caps can be heated up to X-ray temperatures by relativistic particles impinging onto the magnetic poles from the acceleration zones in the magnetosphere. The radio pulsar models (e.g., Cheng \& Ruderman 1980; Arons 1981; Michel 1991; Beskin, Gurevich, \& Istomin 1993) predict polar cap radii $\rpc\sim (2\pi R^3/Pc)^{1/2}$ (where $R\approx 10$~km is the NS radius), i.e., $\rpc\sim 1$--5~km for millisecond pulsars, although different models predict quite different polar cap temperatures, in the range of 1--10~MK. Detection of the polar cap thermal radiation would allow one to discriminate between various models of radio pulsars, study the properties of NS surface layers, and constrain the NS mass-to-radius ratio (Pavlov \& Zavlin 1997; Zavlin \& Pavlov 1998 [ZP98]). However, just as in the case of ordinary pulsars, this radiation is detectable only if it is not buried under stronger nonthermal radiation. The current theoretical models are not elaborate enough to predict in which (if any) of millisecond pulsars the thermal component can be brighter than the nonthermal one (in particular, both the thermal and nonthermal luminosities are expected to increase with spin-down energy loss $\dot{E}$, perhaps with different rates). Therefore, we have to rely upon the analysis of X-ray observations to distinguish the thermal and nonthermal components. The X-ray observatories \ros, \asca, and \sax\/ have detected 11 millisecond pulsars (nearly 1/3 of all X-ray-detected rotation-powered pulsars --- see Becker \& Pavlov 2001 for a recent review). Five of these pulsars are identified in X-rays only by positional coincidence with the radio pulsars and, due to the low number of recorded counts, provide only crude flux estimates. The radiation from 3 pulsars --- B1821--24 (Saito et al.~1997), B1937+21 (Takahashi et al.~2001), and J0218+4232 (Mineo et al.~2000) --- is clearly nonthermal: their power-law spectra, detected with \asca\/ and \sax\/ up to energies of 5--10~keV, are very hard, with photon indices $\gamma\sim 1$, and their pulse profiles show sharp peaks. Interestingly, these 3 pulsars are characterized by particularly large $\dot{E}$ values, $\dot{E}=(2-20)\times 10^{35}$~erg~s$^{-1}$, and their magnetic fields at the light cylinder, $B_{\rm lc}=B(R/R_{\rm lc})^3\sim 10^6$~G, are close to that of the Crab pulsar. The case for the other 3 pulsars --- J0437--4715 (Becker \& Tr\"umper 1993, 1999 [BT93, BT99]; ZP98), J2124--3358 (BT99), and J0030+0451 (Becker et al.~2000) --- is less certain. These pulsars show broad peaks of X-ray pulsations, but it does not necessarily mean that their radiation is thermal because broad peaks can be produced by nonthermal emission at some viewing angles. High-quality spectra have been recorded for the brightest of these pulsars, J0437--4715, but their interpretation has been controversial --- e.g., ZP98 suggest that the radiation detected with \ros\/ and \euv\/ can be interpreted as thermal radiation from hot polar caps, whereas BT99 argue that the radiation is nonthermal (see \S2). To resolve this controversy, the pulsar needed to be observed at energies above the soft \ros\/ and \euv\/ bands ($E\ga 2$~keV), and with high spatial resolution to avoid contamination of the pulsar emission by a nearby AGN which compromised the \asca\/ and \sax\/ data. The {\sl Chandra} X-ray Observatory provides both the superb spatial resolution and high throughput at higher energies, together with timing capability. In this paper we present the results of our observations of \psr\ with \chan. We start from a summary of the previous results on \psr\ in \S2. The spectral and timing analyses of the \chan\/ data are presented in \S3 and \S4. Implication of the results are discussed in \S5.
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The \chan\/ observations of \psr\ have allowed us to perform the spatial, timing and spectral analyses of the new data collected with high angular and spectral resolution in an extended energy range. The principal new result of the ACIS observation is the measurement of the pulsar's X-ray spectrum at higher energies, up to 7~keV. Among the formally acceptable fits of the combined \ros\/ PSPC and \chan\/ ACIS spectra, only the broken PL model corresponds to a purely nonthermal pulsar radiation in the X-ray band. However, the break energy of about 1.1 keV is 5--6 orders of magnitude lower than those observed in other pulsars. Such a difference looks too large to be explained by a lower magnetic field in the radiating region --- it would also require a lower break energy in the spectrum of radiating electrons/positrons, compared to ordinary pulsars. Moreover, extrapolation of the broken PL model to the optical $B$ and $V$ bands (assuming the same slope as in the X-ray range below 1.1~keV) predicts optical magnitudes, $m_B=19.6\pm 1.4$ and $m_V=19.1\pm 1.4$ (for extinction coefficients $A_B=0.3$, $A_V=0.2$ --- Danziger, Baade, \& della Valle 1993), much brighter than those detected from the white dwarf companion, $m_B=22.1\pm 0.1$ and $m_V=20.9\pm 0.1$ (Bailyn 1993; Danziger et al.~1993; Bell et al.~1993). Thus, we do not consider the broken PL model to be a plausible interpretation. The alternative model involves {\sl two components} of different origin --- thermal and nonthermal. The nonthermal component originates in the pulsar magnetosphere\footnote{Another source of unresolvable nonthermal radiation could be the shocked pulsar wind near the white dwarf companion (e.g., Arons \& Tavani 1993). However, at the distance of $a_p=1.5\times 10^{11}$ cm from the pulsar (van Straten et al.\ 2001), the companion intercepts only a fraction $\sim 6\times 10^{-4}$ of the wind (assuming the wind is approximately isotropic), too small to explain the observed nonthermal luminosity.}, whereas the thermal component is emitted from hot polar caps on the NS surface. Depending on assumption about the polar cap temperature distribution, one gets different relative contributions from these two components. In the model with uniformly heated polar caps, the nonthermal component described as a PL of photon index $\gamma=2.7-2.9$ provides about 80\% in the X-ray flux and dominates at energies below 0.6 keV and above 2.7~keV. However, this model encounters the same problems in the EUV and optical bands as the broken PL (the steeper slope of this PL component predicts even higher fluxes in the $B$ and $V$ bands). If we assume a more plausible polar cap model with temperature decreasing outwards from the cap center, then the thermal component becomes dominant between 0.06 keV and 2.5 keV, providing some 75\% of the X-ray flux, while the PL component of $\gamma=1.6-2.5$ dominates outside this band. In addition to a more realistic temperature distribution, the latter model is well consistent with the \euv\/ data and yields estimates on the hydrogen column density in agreement with the indirect measurements. As this PL component is fainter than in the two other models, its extension to the optical falls below the observed radiation of the white dwarf companion for photon indices $\gamma < 1.9$. This allows us to predict that the PL component should be observable in the UV (particularly, far-UV) range where it is brighter than the Wien tail of the white dwarf spectrum, assuming there is no turnover of the nonthermal spectrum between the UV and soft-X-ray energies. On the other side of the X-ray band, extrapolation of the PL component to the gamma-ray energies above 100~MeV predicts a photon flux $f<2\times 10^{-8}$~s$^{-1}$~cm$^{-2}$ ($\gamma>1.6$), below the upper limit, $f < 1.5\times 10^{-7}$~s$^{-1}$~cm$^{-2}$, obtained from the {\sl CGRO} EGRET observations (Fierro et al.~1995). We emphasize that these models require the thermal radiation to be emitted from hydrogen (or helium) NS atmosphere. The high spectral resolution of the ACIS data rules out an atmosphere comprised of heavier chemical elements. The HRC-S observation of \psr\ has demonstrated the \chan\/ timing capability at a millisecond level. The HRC-S pulse profile looks narrower, and the pulsed fraction is somewhat higher, than those obtained in the earlier \ros\/ and \euv\/ observations at lower energies, which could be explained by the properties of the thermal radiation from polar caps covered with a hydrogen or helium atmosphere. On the other hand, the shape of the profile is clearly asymmetric, with a longer rise and faster decay, which cannot be explained by a simple axisymmetric temperature distribution. Relativistic effects (particularly, the Doppler boost) should lead to a different asymmetry --- a faster rise and longer trail (Braje \& Romani 2000; Ford 2000). The analysis of HRC-S data demonstrates, for the first time, that the phase of the X-ray pulse virtually coincides with that of the radio pulse. If, as we suggest, the main contribution to the HRC-S band is due to the thermal polar cap radiation, and if the pulsar radio beam is directed along the magnetic axis, then the radio emission must be generated close to the NS surface --- e.g., the time difference of $<0.1$ ms between the X-ray and radio phases corresponds to a distance of $<30$ km, much smaller than the light cylinder radius, $R_{\rm lc}=275$ km. Alternatively, if the radio emission is generated at a higher altitude, the combination of field-line sweepback and aberration must contrive to cancel the radial travel-time difference. The \chan\/ observations show no sign of an X-ray PWN that could accompany the bow-shock revealed by the H$_\alpha$ observations. Three-sigma upper limits on the PWN brightness (in counts~arcsec$^{-2}$) can be estimated as $3(b/A)^{1/2}$, where $b$ is the background surface brightness ($b=0.51$ and 0.28 counts~arcsec$^{-2}$ for the ACIS-S and HRC-I images, respectively), and $A$ is the PWN area (we will scale it as $A=1000 f_A$ arcsec$^2$, assuming that a typical transverse size of the PWN somewhat exceeds the stand-off distance, $10''$, of the bow shock). For a power-law PWN spectrum with a photon index $\gamma=1.5$--2 (similar to those observed from other PWNe), these upper limits correspond to the PWN intensities $I_{x,{\rm pwn}}<(1.3$--$1.8)\times 10^{-17} f_A^{-1/2}$ and $I_{x,{\rm pwn}}<(3.6$--$5.7)\times 10^{-17} f_A^{-1/2}$ erg cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$, for the ACIS-S and HRC-I, respectively, in the 0.1--10 keV range. The corresponding upper limits on the PWN X-ray luminosity, $L_{x,{\rm pwn}}\approx 4\pi d^2 A I_{x,{\rm pwn}}$ are much smaller than the rotational energy loss rate, $\dot{E}=3.8\times 10^{33}$ erg s$^{-1}$ --- e.g., $L_{x,{\rm pwn}} < (3.0$--$4.2)\times 10^{28} f_A^{1/2}$ erg s$^{-1}$ for the more sensitive ACIS-S limit. The low upper limits on the PWN luminosity in X-rays can be simply explained by a low magnetic field in the PWN region, expected for a particle-dominated pulsar wind. The shock in the relativistic pulsar wind should be located just interior to the observed H$_\alpha$ bow shock (Arons \& Tavani 1993). When the wind electrons pass through the shock, their directions of motion become ``randomized'', and their synchrotron radiation may result in an X-ray nebula, provided the electron energies and the magnetic field are high enough in the post-shock region. The pre-shock magnetic field can be estimated as $B_1=[\dot{E}/(f_\Omega r_s^2c)]^{1/2} [\sigma/(1+\sigma)]^{1/2}=18\,f_\Omega^{-1/2} [\sigma/(1+\sigma)]^{1/2}~\mu$G, where $r_s\approx 2\times 10^{16}$ cm is the stand-off distance corresponding to $10''$ at $d=140$ pc, $f_\Omega=\Delta\Omega/(4\pi)\leq 1$ is the collimation factor of the wind, and the ``magnetization parameter'' $\sigma$ is the ratio of the Poynting flux to the kinetic energy flux. The maximum value of the post-shock magnetic field, $B\simeq B_1\simeq 18\, f_\Omega^{-1/2}~\mu$G, is obtained for $\sigma\gg 1$. However, according to Kennel \& Coronity (1984; KC84 hereafter), a significant fraction of the total energy flux upstream can be converted into (observable) synchrotron luminosity downstream only if $\sigma \lapr 0.1$ (e.g., these authors estimate $\sigma\approx 0.003$ for the Crab pulsar). For $\sigma\ll 1$, the post-shock magnetic field is $B\simeq 3(1-4\sigma) B_1\simeq 53\,f_\Omega^{-1/2} \sigma^{1/2}(1-4.5\sigma)~\mu$G (e.g., from eqs.\ [4.8] and [4.11] of KC84, $B$ increases from 3\,$\mu$G to 12\,$\mu$G when $\sigma$ increases from 0.003 to 0.1, at $f_\Omega=1$). Such low magnetic fields in the shocked wind strongly limit the maximum energy, $m_ec^2 \Gamma_{\rm max}$, of radiating electrons and, consequently, the maximum frequency $\nu_{\rm max}$ of the synchrotron radiation. Since the Larmor radius of most energetic electrons, $r_{L}=1.7\times 10^8\, \Gamma_{\rm max} B_{-5}^{~~-1}~{\rm cm}$, cannot exceed the shock radius $r_s$ substantially, we obtain $\Gamma_{\rm max} < 10^8\, f_s B_{-5}$, $h\nu_{\rm max} \sim (heB/4\pi m_e c)\Gamma_{\rm max}^2 < 0.6\, f_s^2 B_{-5}^{~~3}$ keV, where $B_{-5}=B/(10\, \mu{\rm G}$), $f_s\equiv r_L/r_s\sim 1$. If, for instance, $B<5 f_s^{-2/3}~\mu$G (i.e., $\sigma < 0.01\, f_s^{-4/3}$ in the KC84 model), the synchrotron emission at the bow shock is not expected to be seen in X-rays. Despite the superior quality of the \chan\/ data, which have allowed us to detect the hard tail of the pulsar's spectrum and pinpoint the absolute phase of the X-ray pulse, there still remain some open problems. Although our analysis strongly favors the thermal+nonthermal interpretation, it is still unclear which of the two components dominates in the X-ray radiation of \psr. To establish the relative contributions of these components, energy-resolved timing and time-resolved spectral analysis are needed, which, hopefully, will be possible with the forthcoming {\sl XMM-Newton} data.
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astro-ph0112544_arXiv.txt
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astro-ph0112402_arXiv.txt
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We consider the metallicities and kinematics of nearby stars known to have planetary-mass companions in the general context of the overall properties of the local Galactic Disk. We have used Str\"omgren photometry to determine abundances for both the extrasolar-planet host stars and for a volume-limited sample of 486 F, G and K stars selected from the Hipparcos catalogue. The latter data show that the Sun lies near the modal abundance of the disk, with over 45\% of local stars having super-solar metallicities. Twenty of the latter stars (4.1\%) are known to have planetary-mass companions. Using that ratio to scale data for the complete sample of planetary host stars, we find that the fraction of stars with extrasolar planets rises sharply with increasing abundance, confirming previous results. However, the frequency remains at the 3-4\% level for stars within $\pm0.15$ dex of solar abundance, and falls to $\sim1\%$ only for stars with abundances less than half solar. Given the present observational constraints, both in velocity precision and in the available time baseline, these numbers represent a lower limit to the frequency of extrasolar planetary systems. A comparison between the kinematics of the planetary host stars and a representative sample of disk stars suggests that the former have an average age which is $\sim60\%$ of the latter.
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astro-ph0112402_arXiv.txt
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astro-ph0112128_arXiv.txt
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This paper is the second in a series devoted to examining the multi-wavelength properties of supernova remnants (SNRs) located in nearby galaxies. We consider here the resident SNRs in the nearby Sculptor Group Sd galaxy NGC 7793. Using our own Very Large Array (VLA) radio observations at 6 and 20~cm, as well as archived $\it{ROSAT}$ X-ray data, previously published optical results and our own H$\alpha$ image, we have searched for X-ray and radio counterparts to previously-known optically-identified SNRs, and for new previously unidentified SNRs at these two wavelength regimes. Consistent with our prior results for NGC 300, only a tiny minority of the optically-identified SNRs have been found at another wavelength. The most noteworthy source in our study is N7793-S26, which is the only SNR in this galaxy that is detected at all three wavelengths (X-ray, optical and radio). It features a long ($\sim$ 450 parsecs) filamentary morphology that is clearly seen in both the optical and the radio images. N7793-S26's radio luminosity exceeds that of the Galactic SNR Cas A, and based on equipartition calculations we determine that an energy of at least 10$^{52}$ ergs is required to maintain this source. Such a result argues for the source being created by multiple supernova explosions rather than by a single supernova event. A second optically-identified SNR, N7793-S11, has detectable radio emission but is not detected in the X-ray. A radio-selected sample of candidate SNRs has also been prepared by searching for coincidences between non-thermal radio sources and regions of H$\alpha$ emission in this galaxy, and this search has yielded five new candidate radio SNRs, to be added to the 28 SNRs that have already been detected by optical methods. A complementary search for new candidate X-ray SNRs has also been conducted by searching for soft-spectrum sources ($\it{kT}$ $<$ 1 keV) that are coincident with regions of H$\alpha$ emission. That search has yielded a candidate X-ray SNR which is coincident with one (and possibly two) of the candidate radio SNRs, but considerable diffuse X-ray emission throughout the disk of NGC 7793 reduces the efficacy of the search. Like NGC 300, very little overlap in identifications is seen between the SNRs found through X-ray, optical and radio methods, and such a result argues for the role played by distance-dependent selection effects in determining the detectability of SNRs. In addition, we find that the density of the ambient interstellar medium (ISM) surrounding SNRs significantly impacts the spectral characteristics of the SNRs in this galaxy, consistent with surveys of the SNR populations in other galaxies.
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This paper is the second in a series devoted to the multi-wavelength study of supernova remnants (SNRs) in nearby galaxies. In our previous paper \citep*[hereafter referred to as Paper I]{Pannuti00} we analyzed observations made at the X-ray, optical and radio wavelengths of the nearby Sculptor Group Sd galaxy NGC 300. We sought to determine the X-ray and radio properties of the 28 SNRs identified previously in that galaxy through optical search techniques, namely H$\alpha$ and [\ion{S}{2}] narrow-band imaging \citep*[hereafter referred to as BL97]{DDB80, BL97}, and in addition, we searched for new candidate SNRs at the X-ray and radio wavelengths to complement this prior optical work. Our search yielded sixteen new candidate X-ray and radio SNRs, later reduced to fifteen by the recent work of \citet{RP01}, and the total number of SNRs and candidate SNRs in NGC 300 is now 43. We found very little overlap between the three sets of selected candidates, and we interpret this to indicate that a multiple-wavelength approach is necessary to detect a maximum number of candidate SNRs in a particular galaxy. We also hypothesized that the limited overlap between the selected sets of candidate SNRs indicated selection effects inherent in each type of survey: optical surveys are biased toward the detection of SNRs located in regions of low density and corresponding low optical confusion, while X-ray and radio surveys have the opposite bias and favor the detection of SNRs located in regions of high density. \par In this paper, we examine another Sculptor Group Sd galaxy, NGC 7793, and once again we consider observations made at the X-ray, optical and radio wavelengths. Following the paradigm set by our previous work (\citeauthor{Pannuti00}), our intent is to determine the X-ray and radio properties of the previously-known optically-identified SNRs as well as search for new candidate SNRs at those two wavelengths. Salient properties of NGC 7793 are listed in Table \ref{7793Props}: \citet{PC88} measured a distance to this galaxy of 3.38 Mpc and classified it as a member of the Sculptor Group of galaxies. Because of its proximity and its low inclination angle of 50$^{\circ}$ \citep{T88}, this galaxy makes an excellent choice for the study of galactic properties. These studies include an optical survey for resident SNRs (\citeauthor{BL97}), analyses of its HI content \citep{CP90}, its surface photometry \citep{C85} and its radio continuum properties \citep{H86}. Its X-ray properties have also been the subject of X-ray analysis, based on observations performed with the $\it{Einstein}$ satellite \citep{F92} and later observations performed with the $\it{ROSAT}$ satellite \citep*[hereafter referred to as RP99]{RP99}. In Section \ref{NGC7793ObsSection}, we describe the observations of this galaxy and data reduction at each wavelength, beginning with the radio (Section \ref{NGC7793RadioObsSubSection}), followed by the optical (Section \ref{NGC7793OptObsSubSection}) and concluding with the X-ray (Section \ref{NGC7793XrayObsSubSection}). We discuss the multi-wavelength properties of the optically-selected SNRs in Section \ref{NGC7793OptSNRsSection}, and then present the new candidate SNRs selected at the radio and X-ray wavelengths in Sections \ref{NGC7793RadioSNRsSection} and \ref{NGC7793XraySNRsSection}. A discussion of our findings in this work is presented in Section \ref{DiscussionSection} and finally our conclusions are given in Section \ref{ConclusionsSection}.
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} We have presented a multi-wavelength search and analysis of the SNR population in NGC 7793. The results and conclusions of this work can be summarized as follows: \par 1) NGC 7793, a nearby spiral galaxy, has been observed in the X-ray, optical and radio wavelengths to analyze its resident SNR population. This analysis has examined X-ray observations of the galaxy made by the ${\it ROSAT}$ ${\it PSPC}$ instrument, an H$\alpha$ image of the galaxy and new 6 and 20~cm observations made with the VLA. We have analyzed both the X-ray and radio spectral properties of the SNRs previously identified in the optical by \citeauthor{BL97}, and in addition we have searched for new candidate X-ray and radio SNRs. \par 2) N7793-S26 is the only optically-identified SNR that possesses both X-ray and non-thermal radio emission. The extreme radio luminosity of this object and its large filamentary structure have led to speculation about its creation. A study of the energetics of this SNR suggest that it was created by multiple supernova explosions rather than a single supernova event. One other optically-identified SNR, N7793-S11, possesses non-thermal radio emission but no X-ray emission. Consistent with prior studies of the SNR populations in the galaxies NGC 300 and NGC 6946, X-ray and radio emission from optically-identified SNRs are in general not detected. This lends more support to the hypothesis that searches conducted for SNRs at each wavelength -- X-ray, optical and radio -- all possess inherent biases. We also find additional evidence that the density of the ambient ISM surrounding an SNR plays a critical role in dictating the SNR's spectral characteristics. \par 3) A search for non-thermal radio sources at 6~cm and 20~cm (with a minimum detection level of 3$\sigma$) that are close to or within HII regions has yielded five candidate radio SNRs. Of these five sources, one (NGC 7793-R3) possesses an X-ray counterpart at the 3$\sigma$ level, namely P10, from the listing of X-ray sources in this galaxy that was prepared by \citeauthor{RP99}. Another candidate radio SNR, NGC 7793-R4, may also contribute X-ray flux to P10. \par 4) A search for candidate X-ray SNRs has not revealed any new sources in addition to the SNRs already found through radio and optical surveys. The search for SNRs at this wavelength is complicated by the presence of considerable diffuse X-ray emission throughout the entire disk of this galaxy. \par 5) The multi-wavelength campaign has added five new candidate SNRs to the 28 previously identified through optical methods. While the number of new detected SNRs is noticeably lower than the number of new SNRs found in our study of NGC 300 as presented in \citeauthor{Pannuti00}, we feel that this result is linked to the larger distance to NGC 7793 than NGC 300. Because the X-ray and radio observations of NGC 7793 did not have improved sensitivities to compensate for the increased distance to this galaxy compared to NGC 300, both surveys can only sample more luminous portions of the X-ray and radio SNR population in NGC 7793 compared to NGC 300.
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\PRE{\vspace*{.1in}} If extra spacetime dimensions and low-scale gravity exist, black holes will be produced in observable collisions of elementary particles. For the next several years, ultra-high energy cosmic rays provide the most promising window on this phenomenon. In particular, cosmic neutrinos can produce black holes deep in the Earth's atmosphere, leading to quasi-horizontal giant air showers. We determine the sensitivity of cosmic ray detectors to black hole production and compare the results to other probes of extra dimensions. With $n \ge 4$ extra dimensions, current bounds on deeply penetrating showers from AGASA already provide the most stringent bound on low-scale gravity, requiring a fundamental Planck scale $M_D > 1.3 - 1.8~\tev$. The Auger Observatory will probe $M_D$ as large as 4 TeV and may observe on the order of a hundred black holes in 5 years. We also consider the implications of angular momentum and possible exponentially suppressed parton cross sections; including these effects, large black hole rates are still possible. Finally, we demonstrate that even if only a few black hole events are observed, a standard model interpretation may be excluded by comparison with Earth-skimming neutrino rates.
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Tiny black holes (BHs) can be produced in particle collisions with center-of-mass energies above the fundamental scale of gravity~\cite{Amati:1987wq,'tHooft:rb}, where they should be well-described semi-classically and thermodynamically~\cite{Hawking:1975sw}. In conventional 4-dimensional theories, {\em viz.}, where the Planck scale $\sim 10^{19}~\gev$ is fundamental and the weak scale $\sim 1$ TeV is derived from it via some dynamical mechanism, the study of such BHs is far beyond the realm of experimental particle physics. Over the last few years, however, physicists have begun exploring an alternative approach to the longstanding gauge hierarchy problem, wherein the weak scale becomes the fundamental scale of nature and the Planck scale is derived from this, with the hierarchy in scales a consequence of large or warped extra dimensions~\cite{Antoniadis:1990ew,Randall:1999ee}. If this is the case, the fundamental scale of gravity can be ${\cal O}$(TeV), and BH production and evaporation may be observed in collisions of elementary particles~\cite{Banks:1999gd,Emparan:2000rs,% Giddings:2000ay,Giddings:2001bu,Dimopoulos:2001hw}. If gravity indeed becomes strong at the TeV scale, ultra-high energy cosmic rays provide a powerful opportunity to probe BH production at super-Planckian energies~\cite{Feng:2001ib}. Cosmic rays with energies $\sim 10^{19}~\ev$ have been observed~\cite{Nagano:ve}. They interact in the Earth's atmosphere and crust with center-of-mass energies $\sim 100~\tev$, far beyond the reach of present and planned man-made colliders. These cosmic rays may therefore produce BHs, allowing cosmic ray detectors to test the existence of TeV-scale gravity and extra dimensions by searching for evidence of BH production~\cite{Feng:2001ib,Anchordoqui:2001ei,% Emparan:2001kf,Ringwald:2001vk,Uehara:2001yk}. A particularly promising signal is provided by ultra-high energy cosmic neutrinos, which may produce BHs with cross sections two or more orders of magnitude above their standard model (SM) interactions. These BHs will decay promptly in a thermal distribution of SM particles. Of the order of a hundred BH events may be detected at the Auger Observatory~\cite{Feng:2001ib} as quasi-horizontal, deeply penetrating showers with distinctive properties~\cite{Anchordoqui:2001ei}. The possibility of BH production by cosmic rays supplements possible sub-Planckian signatures of low-scale gravity~\cite{Nussinov:1999jt,Jain:2000pu,Tyler:2001gt,% Alvarez-Muniz:2001mk,Sigl:2001th}. In this article we extend previous work to derive bounds from the non-observation of BH-initiated showers in current data at the Akeno Giant Air Shower Array (AGASA). We also extend previous analyses of BH discovery prospects at Auger, and discuss in detail the possibility of distinguishing BH events from SM events. A preliminary version of some of these results was presented in Ref.~\cite{GAP}. We begin in Sec.~\ref{sec:limits} with an overview of TeV-scale gravity. We collect and review existing bounds on the fundamental Planck scale in a uniform convention. In Sec.~\ref{sec:BH} we discuss semiclassical BH production, including the effects of angular momentum and the production of Kerr BHs, as well as the proposed exponential suppression advocated by Voloshin~\cite{Voloshin:2001vs,Voloshin:2001fe}. This is followed in Secs.~\ref{sec:flux} and \ref{sec:acceptance} by detailed discussions of cosmogenic neutrino fluxes and ground array experiments, respectively. Our results for event rates and new limits on the scale of higher-dimensional gravity are presented in Secs.~\ref{sec:AGASA} and \ref{sec:Auger}. We begin with current data from AGASA. The AGASA Collaboration has already reported no significant signal for neutrino air showers during an observation time (live) of 1710.5 days~\cite{agasa}. Given the standard assumption of a geometric black hole cross section, we find that this data implies the most stringent bound on the fundamental Planck scale to date for $n \ge 4$ extra dimensions, exceeding limits derived~\cite{Ringwald:2001vk} from Fly's Eye data~\cite{Baltrusaitis:mt} and also more stringent than the constraints from graviton emission and exchange obtained by the LEP~\cite{Pagliarone:2001ff} and D\O~\cite{Abbott:2000zb} Collaborations. In Sec.~\ref{sec:Auger} we then consider the prospects for BH production at the Auger Observatory. Tens of black holes may be observed per year; conversely, non-observation of BHs will imply bounds as large as 4 TeV on the fundamental Planck scale. In Sec.~\ref{sec:skimming} we note that comparison to Earth-skimming neutrino event rates~\cite{Bertou:2001vm,Feng:2001ue,Kusenko:2001gj,Domokos:1997ve} allows one to distinguish BH events from SM events. This point was noted already in Ref.~\cite{Feng:2001ib}, but was not considered in Ref.~\cite{Ringwald:2001vk}, leading to weaker conclusions. Here, we consider this point quantitatively and find that, even with a handful of BH events, a SM explanation may be excluded based on event rates alone. If seen, black holes created by cosmic rays will provide the first evidence for extra dimensions and TeV-scale gravity, initiating an era of detailed study of black hole properties at both cosmic ray detectors and future colliders, such as the LHC~\cite{Giddings:2001bu,Dimopoulos:2001hw,Cheung:2001ue,% Park:2001xc,Rizzo:2001dk,Dimopoulos:2001qe,Hossenfelder:2001dn}. Our conclusions are collected in Sec.~\ref{sec:conclusions}.
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\label{sec:conclusions} In this work we have shown that cosmic ray observations in the recent past (AGASA) and in the near future (Auger) provide extremely sensitive probes of low-scale gravity and extra dimensions. We have focused on the production of TeV-scale BHs resulting from collisions of ultra-high energy cosmic neutrinos in the Earth's atmosphere, and have considered the impact of various theoretical issues in the determination of the BH production cross section. In particular, mass shedding, the production of BHs with non-zero angular momentum, and a possible enhancement of the BH cross section can be expected to give minor perturbations. The exponential suppression proposed by Voloshin is more significant, but large and observable BH event rates are still possible. More specifically, in the case of $n$ extra spatial dimensions compactified on an $n$-torus with a common radius, we have found the following: \begin{itemize} \item{Present bounds on atmospheric BH production imply 95\% CL lower limits on the fundamental Planck mass of $\md\ge 1.3-1.5~\tev$ for $n=4$, rising to $\md\ge 1.6-1.8~\tev$ for $n=7$. These bounds follow from the non-observation of a significant excess of deep, quasi-horizontal showers in 1710.5 days of running recently reported by the AGASA Collaboration~\cite{agasa}. The absence of a deeply-penetrating signal in the Fly's Eye data~\cite{Baltrusaitis:mt} also implies lower bounds on $M_D$. These are consistently weaker, however. For example, for $n=6$, $\xmin=1$, and the same (PJ) flux we have used, Ringwald and Tu find $M_D > 900~\gev$~\cite{Ringwald:2001vk}. We find this difference to be significant: the AGASA and Fly's Eye constraints rely on identical theoretical assumptions, and given the scaling in \eqref{scaling}, a factor of 2 difference in $\md$ bounds corresponds to a factor of more than 4 in acceptance or, equivalently, running time. The AGASA limits derived here exceed the D\O\ bound $\md\agt 0.6-1.2$ TeV, where the variation reflects uncertainty from the choice of ultraviolet cutoff for graviton momenta transverse to the brane. The cosmic ray limits are subject to a separate set of uncertainties, discussed at length above, but follow from conservative evaluations of the neutrino flux and experimental aperture, and $\xmin=1$. For $\xmin=3$, these limits are somewhat reduced, but still generally exceed the Tevatron bounds. The cosmic ray bounds from AGASA therefore represent the best existing limits on the scale of TeV-gravity for $n\ge 4$ extra spatial dimensions. A summary of the most stringent present bounds on $\md$ for $n \ge 2$ extra dimensions is given in Fig.~\ref{fig:summary}. } \begin{figure}[tbp] \postscript{summary.eps}{0.56} \caption{Bounds on the fundamental Planck scale $\md$ from tests of Newton's law on sub-millimeter scales, bounds on supernova cooling and neutron star heating, dielectron and diphoton production at the Tevatron, and non-observation of BH production at AGASA. Future limits from the Auger ground array, assuming 5 years of data and no excess above the SM neutrino background, are also shown. The range in Tevatron bounds corresponds to the range of brane softening parameter $\Lambda/\md =0.5-1$. The range in cosmic ray bounds is for $\xmin=1-3$. See text for discussion. } \label{fig:summary} \end{figure} \item{The reach of AGASA will be extended significantly by the Auger Observatory. If no quasi-horizontal extended air shower events are observed in 5 years (beyond the expected two SM neutrino events supplemented by as many as 10 hadronic background events), Auger will set a limit of $\md\agt 3~\tev$, at 95\% CL, for $n \ge 4$. Even in the case where the cross section is decreased by the exponential suppression factor in \eqref{sigmasupp}, a bound $\md\agt$ 2 TeV may be found under the same background assumptions. } \item{Conversely, given the large reach of Auger, tens of BH events may be observed per year. We have discussed in some detail how combined measurements of quasi-horizontal air showers and Earth-skimming $\nu_{\tau}\rightarrow \tau$ events may be used to identify new neutrino interactions beyond the SM, even with complete uncertainty about the incident neutrino flux. In the case of BH production, the quasi-horizontal event rate is enhanced, while the Earth-skimming rate is suppressed, since BH production in the Earth acts as an absorptive channel, depleting the SM rate. With counting experiments alone, one can therefore exclude a SM interpretation of BH events, and may distinguish BH events from almost all other possible forms of new physics.} \end{itemize} In conclusion, in the next several years prior to the analysis of data from the LHC, super-Planckian BH production from cosmic rays provides a promising probe of extra dimensions. Searches for BH-initiated quasi-horizontal showers in the Earth's atmosphere at AGASA provide the most stringent bounds on low-scale gravity at present, and the Auger Observatory will extend this sensitivity to fundamental Planck scales well above the TeV scale.
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I conclude that it is practicable to measure $p\overline{p}$ annihilation spectra separately from underlying power-law backgrounds if the appropriate selection criteria are applied. By hypothesis, any source emitting this spectral feature must be artificial; it is extremely difficult to imagine any other possibility except new laws of physics (\S 1). When applied to the EGRET data the method would have detected a steady $p\overline{p}$ annihilation spectrum down to levels $\sim 2 \times 10^{-8}$ photon/(cm$^{2}$ s), and transients on time-scales ranging down to $\sim 100$ d at levels ranging up to $\sim 10^{-7}$ photon/(cm$^{2}$ s), outside the Galactic plane, both numbers being about a factor 10 higher in the plane. Variable emission detected from the known extragalactic source QSO 2206+650 is presumably not related to ETI activity. These results, limited though they are, are the first ever obtained in this field. They exclude the presence of "human-scale" antimatter-powered space probes (such as might be constructed by humans in this century [11]) within a radius of $\sim 10$ AU, and more ambitious human-like crewed interstellar craft out to several thousand AU. They will be greatly improved by future high energy $\gamma$-ray missions such as GLAST.
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{ We present the results of 1.3~mm observations of the Crab Nebula, performed with the MPIfR bolometer arrays at the IRAM 30-m telescope. The maps obtained, of unprecedented quality at these wavelengths, allow a direct comparison with high-resolution radio maps. Although the spatial structure of the Crab Nebula does not change much from radio to millimetre wavelengths, we have detected significant spatial variations of the spectral index between 20~cm and 1.3~mm. The main effect is a spectral flattening in the inner region, which can be hardly explained just in terms of the evolution of a single population of synchrotron emitting electrons. We propose instead that this is the result of the emergence of a second synchrotron component, that we have tried to extract from the data. Shape and size of this component resemble those of the Crab Nebula in X rays. However, while the more compact structure of the Crab Nebula in X rays is commonly regarded as an effect of synchrotron downgrading, it cannot be explained why a similar structure is present also at millimetre wavelengths, where the electron lifetimes far exceed the nebular age. Our data, combined with published upper limits on spatial variations of the radio spectral index, also imply a low-energy cutoff for the distribution of electrons responsible for this additional synchrotron component. Although no model has been developed so far to explain the details of this component, one may verify that the total number of the electrons responsible for it is in agreement with what predicted by the classical pulsar-wind models, which otherwise are known to fail in accounting for the number of radio emitting electrons. This numerical coincidence can give indications about the origin of this component. We have also detected a spectral steepening at millimetre wavelengths in some elongated regions, whose positions match those of radio synchrotron filaments. The steepening is taken as the indication that magnetic fields in synchrotron filaments are stronger than the average nebular field.
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The Crab Nebula is the prototype of synchrotron nebulae powered by a spinning-down pulsar, also known under the name of ``plerions'' (Weiler \& Panagia \cite{wei78}). This is an extensively studied object, and a wealth of information on the synchrotron nebula comes from detailed observations performed in various spectral ranges, like in radio, infrared, optical, UV and X rays. Modelling all the available data in a comprehensive frame represents a formidable task for the theory. Classical approaches to the modelling of the Crab synchrotron emission, like Pacini \& Salvati (\cite{ps73}) and Kennel \& Coroniti (\cite{kc84a}, \cite{kc84b}), got some success. But more elaborate models can hardly get any substantial improvement with respect to the original approaches, partly because the geometric structure of the Crab Nebula is very complex, but probably also because the processes involved are not fully understood. When more quantitative and detailed modelling will be possible with other plerions we expect to face similar problems: in these respects a large part of the results on the Crab Nebula are likely to be exported to other objects. Considering just the total luminosity spectrum, Pacini \& Salvati (\cite{ps73}) successfully reproduced it from radio to optical, by simply assuming a pure power-law distribution for the injected ``particles'' (hereafter used to indicate relativistic electrons, as well as positrons): but in order to explain by their model the further spectral steepening in the X-ray range and beyond, an {\it ad hoc} spectrum for the injected particles is required. On the other hand Kennel \& Coroniti (\cite{kc84b}) successfully reproduced the spectrum from optical to gamma rays, just assuming a power-law distribution (over a range of energies) for the particles accelerated at the termination shock of the pulsar relativistic wind. However their model fails in explaining the observed radio emission: the problem is that the best-fit wind model implies also an estimate of the total number of radio emitting particles injected into the nebula, which is at least a factor 100 lower than what measured. Up to now this discrepancy has been cured only by introducing some {\it ad hoc} assumptions (see e.g.\ Atoyan \cite{ato99}). As far as the spatial structure of the nebula is concerned, it is not difficult to explain qualitatively its behaviour with frequency, namely the shrinking of the nebular size with increasing frequency: in fact the latter corresponds to increasing particle energy, and therefore decreasing synchrotron lifetimes. However quantitative approaches fail to reproduce the observed profiles, both in the Kennel \& Coroniti (\cite{kc84b}) and in the Pacini \& Salvati (\cite{ps73}) frameworks: the implications of the assumptions in the latter paper on the nebular spatial extent have been investigated by Amato et al.\ (\cite{aea00}). A common characteristic of the above models is that the particles are advected outwards with the magnetic field, following the MHD equations. Somehow better results are for instance obtained by including also diffusive processes, but only when an {\it ad hoc} diffusion coefficient is taken. An alternative to the above scenarios relies on assuming the coexistence of two (or more) components of injected particles, with different spectra as well as with different spatial locations. But one may be unwilling to increase the complexity of the models, unless a stringent evidence in that sense comes out from the observations. Since adiabatic losses preserve the slope of the particles distribution, the most direct test on the presence of multiple components of the injected particles consists into measuring spatial variations of the synchrotron spectral index that cannot in any way result from a synchrotron downgrading (i.e.\ the spectral softening consequent to synchrotron evolution of the emitting particles). This can be done observing at frequencies so low that the related particles are subject to negligible synchrotron losses, and therefore whose distribution retains the slope which had at the injection. Beforehand this kind of test had been done only at radio wavelengths (below with the term ``radio wavelengths'' we roughly indicate wavelengths above 1~cm). Previous claims of spatial variations of the radio spectral index (Velusamy et al.\ \cite{vea92}) have been then contradicted (Bietenholz et al.\ \cite{bea97}). Possibly some variations of the spectral index are present in the very central region, with scales of a few arcsec and associated to the ``wisps'' structures (Bietenholz \& Kronberg \cite{bk92}): however such result could be a mere artifact, originated by comparing data taken at different epochs, in the presence of rapidly moving wavy structures (Bietenholz et al.\ \cite{bea01}). In fact recent radio observations (Bietenholz et al.\ \cite{bea97}) strongly support the idea of a single injected distribution, by putting a tight upper limit, 0.01, to spatial variations of the spectral index, at least on scales larger than 16\arcsec. In this paper we will show that millimetric wavelengths represent the most appropriate spectral range to investigate this issue in the Crab Nebula, by providing new pieces of information with respect to the radio. The map presented in this paper, with a 10.5\arcsec\ resolution, is by far better than the only map of the Crab Nebula previously published at these wavelengths (Mezger et al.\ \cite{mea86}, with only 120\arcsec\ resolution). The paper is organized as follows: in Sect.~2 we report on the observation parameters and on the data reduction; Sect.~3 describes the procedure by which our 1.3~mm map has been compared with a 20~cm radio map; features coming out from this comparison, namely the emergence of a second component in the inner regions and a general bending in filaments spectra, are respectively discussed in Sect.~4 and Sect.~5; Sect.~6 shows that the nature of the emission from the inner component is synchrotron; the morphology of the new synchrotron component is compared in Sect.~7 with maps at other wavelengths; in Sect.~8 we comment on possible spurious effects on our results deriving from time variability of the source; Sect.~9 concludes.
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We have measured inhomogeneities in the spectral index between radio and mm wavelengths, and we have shown that they could be better explained in terms of the emergence of a further synchrotron component, undetected at radio wavelengths, which is located in the inner part of the nebula. The transition in size and shape of the Crab Nebula, moving from the radio to the X-ray spectral range, is qualitatively explained in terms of synchrotron downgrading. But the fact that size and shape of component {\bf B} at mm wavelengths resemble those seen in X rays cannot be explained in that way, since mm emitting particles are subject to only minor synchrotron losses. We suggest instead that two different synchrotron components coexist in the Crab Nebula and that the morphological transition taking place from radio to X rays requires a change of the relative importance of the two components. In order to fit the data, the energy distribution of particles emitting in component {\bf B} requires to have a low-energy cutoff. With this, the total number of particles in component {\bf B} are in agreement with what predicted by Kennel \& Coroniti (\cite{kc84b}) model. Finally, in synchrotron filaments we have found the evidence for a spectral break at a frequency lower than that averaged over the whole nebula. Although there may be some effects related to particle diffusion through filaments, we take a magnetic field in filaments higher than in their surroundings as the main cause of this effect.
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We analytically investigate the formation of an HII region in the accreting envelope of a newborn star. Special care is taken to examine the role of ionizing radiation force. This effect modifies velocity and density distributions and thereby affects the expansion of the HII region. As a result, the upper limit of the stellar mass imposed by the growth of an HII region around a forming star is increased by a larger factor than the previous estimate. In particular, for a star forming out of metal-free gas, this mechanism does not impose a firm upper limit on its mass.
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Stars are supposed to be born in the dense regions of gas clouds. The theory of the gravitational fragmentation of the gas clouds describes how these dense regions are formed inside the parental gas clouds. The typical mass scale of the dense region can be predicted by the Jeans mass of the cloud at the time of the fragmentation. However, a perturbation with a mass scale larger than this Jeans mass is always unstable. Therefore a dense region that is more massive than the typical Jeans mass can form depending on the initial fluctuation spectrum in the cloud. In this sense, the upper limit for the mass of stars cannot be obtained from the analysis of the gravitational fragmentation. Once, the origin of the observed upper limit of stellar mass (around $100M_{\sun}$) was attributed to the instability of massive stars due to the $\varepsilon$-mechanism (e.g., Schwarzshild \& H\"{a}rm 1959). Since later studies revealed that this mechanism leads to moderate mass-loss rather than disruption of the stars (e.g., Appenzeller 1970), formation process, instead of the stability of massive stars, has been believed to limit the upper bound for the stellar mass (e.g., Nakano, Hasegawa, \& Norman 1995). The star formation process is an accretion of ambient matter onto a protostar (stellar core) forming inside a protostellar cloud. In this scenario, the final mass of a star depends on how much mass the protostar can acquire. Although it is not yet clear what mechanism stops the accretion, increasingly strong feedback from a massive forming star is likely to halt the accretion. Larson \& Starrfield (1971) were first to propose that a firm upper limit on the stellar mass is provided by the formation of an HII region around a forming star, as well as by the radiation force acting on dust grains in the accreting envelope. The formation of an HII region prevents further accretion in the following way: when the HII region reaches outer layers of a protostellar cloud, the temperature and then the pressure support surge by some large factor, so that the further infall of material is immediately halted. Larson \& Starrfield (1971) concluded that both of those mechanisms set an upper limit of about 50 M$_{\sun}$ on the masses of Pop I stars forming out of the present interstellar medium. The radiation force onto dust grains was found by subsequent studies to set more stringent mass upper limit (Appenzeller \& Tscharnuter 1974; Kahn 1974; Yorke \& Kr\"{u}gel 1977; Wolfire \& Cassinelli 1987). Among them, Wolfire \& Cassinelli (1987) suggested that even stars as massive as about 10 M$_{\sun}$ cannot form by accretion owing to this mechanism unless dust is significantly depleted. However, for so-called Pop III stars forming out of metal-free gas, this mechanism of the radiation force does not work because of the absence of dust grains. The $\varepsilon$-mechanism drives even milder mass-loss for those stars than Pop I massive stars (Baraffe, Heger, \& Woosley 2001). According to recent simulations of the fragmentation of primordial clouds, the mass scale of fragments (protostellar clouds) is as large as $10^{3} M_{\sun}$ (Bromm, Coppi, \& Larson 1999; Abel, Bryan, \& Norman 2000; Tsuribe 2001). Therefore, the growth of the HII region plays a crucial role in limiting the maximum mass of a Pop III star, if the upper limit by this mechanism is less than the mass scale of fragments. This motivates us to study here the upper mass limit of Pop III stars due to the formation of an HII region. In discussing the formation of an HII region, often the free-fall assumption has been imposed on the flow in the accreting envelope (e.g., Yorke 1986). The role of the momentum transfer to the gas due to ionizing radiation has been neglected. In the context of galaxy formation, however, its significance in dynamical and thermal evolution of the intergalactic medium has been pointed out by Haehnelt (1995). He showed, for the collapse of subgalactic clouds of $\la 10^{10} M_{\sun}$, even a radiation-pressure-driven bounce is possible. Taking this effect into account, in this paper, we study the formation of an HII region in the accreting envelope of a forming massive star. Although this effect is negligible as long as the radius of the HII region is small, it will turn out that inclusion of the ionizing radiation force alters the later evolution of the HII region. In the course of the expansion of the HII region, the flow becomes slower than the free-fall rate owing to the ionizing radiation force. The growth of the HII region is strongly suppressed by this effect. Consequently, the mass upper limit imposed by the expansion of the HII region is increased by a large factor. In particular, the formation of an HII region does not set any firm mass upper limit for Pop III stars. In Section 2, we briefly summarize the relevant aspects of the formation of an HII region in a free-falling envelope around a newborn star. The effect of ionization radiation force is included in Section 3. Finally, we provide a summary and discussion in Section 4.
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To estimate the mass upper limit of metal-free stars imposed by the formation of an HII region, we have found solutions for spherically symmetric and steady accretion flow onto a star emitting ionizing radiation. The behavior of the flow is determined by the velocity of the flow entering the HII region. If the flow velocity at the edge of the HII region goes beyond (below) a critical value, the flow is always accelerating (decelerating, respectively) in the HII region. When this value is equal to the critical value, the velocity remains constant in the HII region. In this critical flow, the radiation force due to ionizing photons exactly balances the gravity. We applied those solutions for a compact HII region forming around an accreting star. As the HII region grows in radius, the accreting flow evolves from an accelerating solution to a critical one. Soon after the critical flow is reached, accretion is halted by radiation force due to the ionizing radiation. However, even stars as massive as $10^{3}M_{\sun}$ are unable to emit ionizing photons sufficient to halt the accretion. The halting by gas pressure is even more difficult. Therefore, contrary to the previous expectation, the formation of the HII region does not impose a stringent mass upper limit (at least up to $10^{3}M_{\sun}$) on metal-free stars. We have estimated the upper limit of the mass of Pop III stars in relation to the formation of an HII region. Note that other effects (e.g., stellar wind, mass outflow due to pulsational instability, etc.) could grow in importance and might decrease the upper mass limit of Pop III stars below the value obtained in this paper. We leave detailed studies on these topics for future work. Also, in this paper, we have adopted very simplified assumptions, for example spherical symmetry, steady accretion, etc. Here, we discuss other complexities and possible deviations from our picture. We have considered only photoionization and electron scattering as sources of radiation force. Here, we mention briefly the radiation force due to Ly$\alpha$ photons. Ly$\alpha$ photons are emitted from recombination in the HII region. The flow inside the HII region is exerted by the radiation force due to Ly$\alpha$ photons. We can extend our theory easily to include the Ly$\alpha$ pressure , since, in the HII region, the radiation force due to the Ly$\alpha$ photons is proportional to that due to ionization photons (equation 12 of Haehnelt 1995; see also Braun \& Dekel 1989; Bithell 1990). However, according to Haehnelt(1995), the radiation force due to Ly$\alpha$ photons is at most of marginal importance relative to that due to the ionizing photons. Thus, for the sake of simplicity, we have not included it. The HII region is surrounded by an HI layer, which is very optically thick to Ly$\alpha$ photons. Without dust grains, which absorb Ly$\alpha$ photons and reemit them as infrared photons, the Ly$\alpha$ photons must diffuse out through the HI layer. In the course of this, the HI layer is pushed outward by those photons. Doroshkevich \& Kolesnik (1976) argued that this mechanism expels the HI layer soon after the HII region is formed and thereby limits the mass of stars below 10 M$_{\sun}$. However, Harrington (1973) showed that, even without dust, the two-photon emission process decreases the Ly$\alpha$ photon density drastically, and the Ly$\alpha$ radiation force is not dynamically important in an HI layer surrounding an HII region. In our case, supersonic motion in the accreting envelope also decreases the Ly$\alpha$ photon density. Considering these facts, it is likely that the Ly$\alpha$ radiation force does not play a significant role in our case. Therefore, we chose to neglect it here and assumed the free-fall outside the HII region. We have identified the base of an HII region as the stellar surface, and have taken the typical value of the stellar radius $10R_{\sun}$ as the inner boundary radius $R_{\rm in}$ of the HII region. However, when the accretion rate is high, a photosphere may be formed in the accreting flow (Wolfire \& Cassinelli 1986; Stahler, Palla, \& Salpeter 1986). In this case, $R_{\rm in}$ should be taken as the radius of the photosphere, since we assume the HII region is optically thin to continuum absorption except for photoionization outside $R_{\rm in}$. In spite of higher $R_{\rm in}$, this effect generally works towards the smaller HII region, since $Q$ drops (Wolfire \& Cassinelli 1986). Hence, our conclusion of higher upper mass limit than the former estimate remains the same. The spherical symmetry is clearly an oversimplification if the accreting matter has large angular momentum and an accretion disk is formed. Even in this case, the spherical symmetric flow is a good approximation outside the centrifugal radius. Suppose here that inside the centrifugal radius, there is a ``cavity'' around the disk. In this case, $R_{\rm in}$ should be taken as the radius of the cavity. If $Q$ is unattenuated inside the cavity, the halting of the flow outside becomes easier because of higher $R_{\rm in}$ and then lower $Q_{\rm crit}$ (see eq. [\ref{eq:qc_0}]). The accretion through the disk might continue, however (Nakano 1989). Those issues are still too speculative and beyond the scope of this paper. Although the accelerating solution is Rayleigh-Taylor stable, it might be unstable to % perturbations: if a portion of the flow becomes slightly slower than the rest, it becomes slower and slower relative to the average flow owing to the increased radiation force. In this case, density irregularities or blobs could be formed in the flow, and our spherically symmetric solution might be regarded as an approximate description of the average flow. Further study of the stability and dynamical evolution of the flow will be interesting.
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astro-ph0112345_arXiv.txt
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astro-ph0112035_arXiv.txt
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{This paper is concerned with small angular scale experiments for the observation of cosmic microwave background anisotropies. In the absence of beam, the effects of partial coverage and pixelisation are disentangled and analyzed (using simulations). Then, appropriate maps involving the CMB signal plus the synchrotron and dust emissions from the Milky Way are simulated, and an asymmetric beam --which turns following different strategies-- is used to smooth the simulated maps. An associated circular beam is defined to estimate the deviations in the angular power spectrum produced by beam asymmetry without rotation and, afterwards, the deviations due to beam rotation are calculated. For a certain large coverage, the deviations due to pure asymmetry and asymmetry plus rotation appear to be very systematic (very similar in each simulation). Possible applications of the main results of this paper to data analysis in large coverage experiments --as PLANCK-- are outlined.
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Many experiments are being designed for the observation of the Cosmic Microwave Background (CMB) anisotropies. From the maps of a given experiment operating with a non-circular (asymmetric) rotating beam, a certain angular power spectrum ($C_{\ell}$ quantities) can be extracted. Different rotations can lead to distinct $C_{\ell}$ coefficients and, the question is: How different are these coefficients? In other words, how relevant is the effect of the rotation strategy on the resulting angular power spectrum? In a previous paper (Arnau \& S\'aez 2000), it was shown that, in the absence of rotation and when the level of instrumental noise is low enough, the effect of a non-circular beam can be subtracted --namely, the beam can be deconvolved-- using the Fourier transform. This subtraction can be performed in such a way that the resulting spectrum, after deconvolution, is very similar to the true one. That is possible if the number of pixels inside the beam, $N_{\mbox{in}}$, is not too great. Indeed, $N_{\mbox{in}}$ cannot be much greater than 10; however, if the beam rotates, the deconvolution is not possible. Nobody has described either the importance of beam rotation or a method to eliminate its effects. The main goal of this paper is the estimation of the effects due to rotation. In Arnau \& S\'aez 2000 (and also in S\'aez, Holtmann \& Smoot, 1996, and S\'aez\& Arnau, 1997). a sort of modified angular power spectrum was used. Here, we extract the standard $C_{\ell}$ quantities from a certain number of squared patches of the sky. Recently, Wu et al. (2001) have proposed a method for data analysis in the case of asymmetric beams. This method is based on an optimal circular beam associated to the asymmetric one. The effects of beam rotation are not studied at all by these authors. Although our methods apply to CMB anisotropy experiments in general, we will pay particular attention to PLANCK mission (scheduled by ESA for 2007). As it was emphasized in Burigana et al. (1998), {\em beam responses are typically nonsymmetric for detectors de-centred from the telescope focus}. Taking into account that CMB anisotropy experiments require observations at different frequencies, various detectors are necessary, which must be distributed as close as possible from the focus; for instance, in the PLANCK mission, around one hundred of detector (bolometers and radiometers) must be distributed in the focal plane. If the focal plane rotates (rotation of the telescope around the spin axis), the beams do. The effect of this rotation deserves attention. Furthermore, there are various identical detectors for each frequency, which are located at different positions in the focal plane and, consequently, the deformations of these beams would be different (identical) if they are located at different (the same) distances from the optical focus; nevertheless, even for identical deformations, the orientations of the resulting asymmetric beams would be different. The motion of the line of sight through the sky also produces a beam asymmetry. The effective beam diameter $\theta_{_{\mbox{FWHM}}}$ appears to be enhanced in the direction of this motion (see Hanany, Jaffe \& Scannapieco 1998). This small effect is due to the beam displacement during the measurement process. It is not taken into account in this paper. Beam rotation depends on the particular experiment under consideration. Given a pixelisation, the beam centre points towards a given pixel a certain number of times, $N_{\mbox{p}}$, and, then, the temperature assigned to this pixel is an average of the temperatures corresponding to each of the $N_{\mbox{p}}$ measurements. The fact that measurements from various beam orientations are averaged could be important. In the case of PLANCK mission, a rough estimate of number $N_{\mbox{p}}$ is given in S\'aez \& Arnau (2000). Here, it is worthwhile to improve a little on that estimate. The satellite has been designed in such a way that: (i) it will cover the full sky in seven months, with a coverage which can be considered as uniform in most part of the sky, (ii) its line of sight will move around a big circle on the sky completing a turn each minute and, (iii) it will move around the same circle for two hours (120 turns). On account of these facts, if the pixel size is $\Delta$ and the angle subtended by the motion of the line of sight between two successive measurements is $\Delta \alpha = \zeta \theta_{_{\mbox{FWHM}}}$, where $\theta_{_{\mbox{FWHM}}}$ is the beam diameter, then, the average number of measurements per pixel (in a seven months observing period) is $N_{\mbox{p}} = 42 \Delta^{2} / \zeta \theta_{_{\mbox{FWHM}}}$, where all the angles are given in arc-minutes (see S\'aez \& Arnau, 2000, for comparison) and, furthermore, the average number of measurements per pixel performed while the line of sight turns 120 times around a given circle is $N_{\mbox{pc}} = 120 \Delta / \zeta \theta_{_{\mbox{FWHM}}}$. From these formulae, it follows that the average number of circles passing by a pixel --during seven months-- is $N_{\mbox{c}} = N_{\mbox{p}}/N_{\mbox{pc}} \simeq \Delta/3$, this result is consistent with the fact that, for a given observational strategy, the number $N_{\mbox{c}}$ is expected to be dependent only on the pixelisation. Of course, it is independent on beam asymmetry. The number of measurements corresponding to different orientations could be important for the effect we are looking for, which is produced by the rotation of asymmetric beams. The larger the pixel size, the better the situation (the greater $N_{\mbox{c}}$). Since the detectors are rigidly attached to the focal plane, any beam has almost the same orientation each time it crosses a given pixel during its motion (120 turns) along a given circle; however, this orientation changes from circle to circle. From the above comments and estimates, it follows that the average number of measurements per pixel corresponding to different beam orientations is $N_{\mbox{c}}$. If the full sky is covered two times and, the second coverage is not identical to the first one, this average number would be $2 \Delta/3$. For $5^{\prime } < \Delta < 10^{\prime}$, this number ranges from 3.3 to 6.6. Nevertheless, there are various detectors in the focal plane for each frequency and, by assuming that all the beams have the same shape but different orientations, the above $N_{\mbox{c}}$ number can be multiplied by the number of beams.
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In the absence of beam, the pixelisation effect and the REPC have been disentangled to conclude that pixelisation produces very systematic deviations with respect to the true angular power spectrum. This conclusion has been obtained using two very different methods for simulations and data analysis (see Sect. 3) We have studied the deviations in the angular power spectrum produced by the rotation of an asymmetric beam. Two rotation strategies have been considered. One of them (SR) is similar to that of future experiments as PLANCK. The second strategy (RR) is very different from SR, and it has been introduced for comparisons. Maps with and without the dust and synchrotron radiations from the Milky Way (at $100 \ GHz$) have been considered. In Sect. 4, the rotation effects corresponding to different cases have been described and compared, now let us present some general comments. If radiation from the galaxy is not considered, the most important conclusion is that rotation effects are very systematic for any rotation strategy and $f_{\mbox{sky}} = 0.39$. They are so systematic that we can subtract the deviations appeared in a 50-simulation, from the spectrum of another one, to recover very well the spectrum corresponding to the nonrotating beam (except for small deviations which seem to be essentially due to the REPC). Furthermore, the resulting effects depend on the rotation strategy, in particular, for large $\ell$ values and, consequently, they must be estimated --using simulations-- in each particular case. Radiation from the galaxy --which can be seen as a non homogeneous an non isotropic statistical field-- contributes significantly to the observable signal, except in the case of the polar galactic regions (G5--G6). In the G1--G2 and G3--G4 cases, the effect of beam rotation is significantly, but not dramatically, different from that obtained in the absence of foregrounds. After the deviations in the angular power spectrum due to beam asymmetry and rotation have been estimated and characterized (the main goal of this paper) and, after proving that beam effects are very systematic, some practical applications can be easily outlined. Take the CMB power spectrum corresponding to a certain theoretical model of structure formation in a given universe, take also a model for the foregrounds, a pixelisation, the asymmetric beam for a given frequency, and the rotation strategy of an experiment with a large enough coverage (i.e. PLANCK), and then, use a simulation --as the 50-simulations of this paper or similar-- to find the spectrum $C_{\ell}^{^{\mbox{CMB}}}$ after smoothing with the asymmetric rotating beam. Repeat the simulation a large enough number of times and verify that the resulting $C_{\ell}^{^{\mbox{CMB}}}$ spectra are similar in all cases (systematic character). Finally, use the deviations among the resulting spectra to assign an error bar to $\langle C_{\ell}^{^{\mbox{CMB}}} \rangle$. Use these data --obtained from simulations-- to answer the following question: Is the theoretical model under consideration compatible with the observational data from the experiment? In order to find the answer, the observational data could be analyzed as follows: (i) Eliminate a part of the instrumental noise using an appropriate method (wavelets, Fourier transform, and so on), (ii) Separate components (CMB, synchrotron from our galaxy, and so on) taking into account the frequency dependences, but keeping beam smoothings unaltered (usually, the beams are eliminated at this stage under simplifying assumptions and without considering rotation), (iii) use the map of the CMB component --which has already been separated from foregrounds-- to extract the experimental spectra, $C_{\ell}^{^{\mbox{CMB}}}(exp)$, and finally (iv) compare $C_{\ell}^{^{\mbox{CMB}}}$ with $C_{\ell}^{^{\mbox{CMB}}}(exp)$ and study if these spectra can be identified taking into account the error bars. If they can, the theoretical model is compatible with observations. Note that --at the last step of the process-- we compare a simulated spectrum with an observational one, and note also that both spectra are obtained from maps which have been smoothed with the same rotating asymmetric beam; hence, the proposed method for data analysis includes beam rotation, treating it (after verification) as the source of a very systematic effect. Of course, this method has been only outlined, and much more work would be necessary before implementation.
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astro-ph0112035_arXiv.txt
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hep-th0112128_arXiv.txt
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By considering a simplified but exact model for realizing the ekpyrotic scenario, we clarify various assumptions that have been used in the literature. In particular, we discuss the new ekpyrotic prescription for passing the perturbations through the singularity which we show to provide a spectrum depending on a non physical normalization function. We also show that this prescription does not reproduce the exact result for a sharp transition. Then, more generally, we demonstrate that, in the only case where a bounce can be obtained in Einstein General Relativity without facing singularities and/or violation of the standard energy conditions, the bounce cannot be made arbitrarily short. This contrasts with the standard (inflationary) situation where the transition between two eras with different values of the equation of state can be considered as instantaneous. We then argue that the usually conserved quantities are not constant on a typical bounce time scale. Finally, we also examine the case of a test scalar field (or gravitational waves) where similar results are obtained. We conclude that the full dynamical equations of the underlying theory should be solved in a non singular case before any conclusion can be drawn.
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Modern ideas of particle physics, such as superstring, $M-$theory~\cite{sugrastring} or quantum gravity~\cite{QG}, cannot in general be subject to experimental constraints because of the enormous energies (usually of the order of the Planck mass) at which they are supposed to become effective. According to recent theoretical developments~\cite{largeD,RS}, there is hope that space possesses more than three large dimensions and that these extra dimensions might turn out to be observable in a not too distant future. The majority of the theoretical models that have been built so far are however based on extremely high energy extensions of the standard particle physics model, and thus currently need to be tested by the yardstick of cosmology, the latter being the only playground at which those theories could have acted. According to the now standard paradigm that describes the early universe and that is expected to stem from such high energy particle models, a phase of superluminal accelerated expansion known as inflation~\cite{inflation} preceded the radiation-dominated epoch. Up to now, no model has come as close to being a reasonable challenger to solve the standard cosmological puzzles (flatness, homogeneity and monopole excess). The extra bonus provided by the inflationary phase is that it leads naturally to a scale-invariant density fluctuation spectrum that seems to be in agreement with the observations. Inspired by the recent developments of $M-$theory~\cite{heterotic}, in particular through Ref.~\cite{HW}, and invoking brane cosmology, recent work~\cite{ekp,ekpN,perturbekp} claimed to be able to solve all the aforementioned problems as well, including a new way of producing primordial cosmological perturbations. Although the model, both in its ``old''~\cite{ekp} and ``new''~\cite{ekpN} versions is plagued with many difficulties~\cite{pyro}, as a potential alternative to inflation (see also Ref.~\cite{PBB} in that respect), it is worth examining in detail, would it be only to re-enforce the confidence we may have in the latter. In both the original and most recent versions, the universe is supposed to consist of a four dimensional (visible) brane evolving in a higher (in practice 5) dimensional bulk. By assuming the brane to be a Bogomolnyi-Prasad-Sommerfield (BPS) state~\cite{BPS}, one ensures that the curvature $\Ka$ vanishes, thus addressing the flatness problem. To begin with, another brane, that can be either a light bulk brane~\cite{ekp}, or the other (hidden) boundary brane~\cite{ekpN}, moves freely in the bulk until it collides with the visible brane. The collision time is interpreted as the hot big bang at which point the model is made to match the standard cosmological model. Apart from the collision time, the theory, which can be seen as effectively four dimensional in the long wavelength limit, relies on the General Relativity (GR) theory together with some extra fields. In this effective 4D model, the Universe collapses, experiences a bounce at some instant in time, and starts expanding. As far as cosmological perturbations are concerned, only GR calculations have been discussed up to now. The pre-impact phase has been the subject of many tentative calculations of the perturbation spectrum that would be generated by quantum perturbations of the brane~\cite{perturbekp}. A general agreement has now been reached~\cite{perturbekp,Lyth,BF,Hwang} that the curvature perturbation spectrum $P_{\zeta }$ has spectral index $n_\zeta=3$, while that of the Bardeen potential $P_{\rm \Phi }$ ends up with $\nS=1$, i.e., a scale invariant spectrum. On the other hand, the spectra of $\Phi$ and $\zeta$ are identical in the post-impact phase, and enter the Cosmic Microwave Background Radiation (CMBR) multipole moments. It is therefore of utmost importance to obtain full knowledge of these spectra not only in the pre-impact phase, but also after the bounce has occurred, i.e., at times that are observable now. In other words, the fate of $\Phi$ and $\zeta$ through the bounce is the main issue before any conclusion regarding the model can be drawn. Only a few definite statements can be done about the bounce epoch. The first, which was advocated by many authors, is that GR does hold during it, or, stated differently, that it lasts sufficiently little that corrections to GR can be regarded as negligible. Lacking the actual theory, this is the only statement that can be endowed with a predictive power. To begin with, it implies that there was no singularity, and, if the null energy condition is to be satisfied, that space is positively curve, i.e., $\Ka=1$. Under these conditions, ordinary perturbation theory~\cite{Bardeen,perturb} can be applied. By assuming continuity of the Bardeen potential and the well known conserved quantity $\zeta$ (defined below), it was then found~\cite{BF} that the scale invariant spectrum does not survive the bounce, with the actual resulting spectrum being much lower than the observed one. The temporary conclusion of this fact is that in order that the ekpyrotic model be still compatible with the observation, a new procedure must be applied to the bounce. Arguing against GR during the bounce epoch sounds natural, as in particular either the real theory is at least 5-dimensional, or, worst indeed, in the case of the new scenario~\cite{ekpN}, the manifold becomes (curvature) singular there, obviously leading to a breakdown of ordinary GR across the bounce. In this case, a new criterion should be derived to replace the ordinary junction conditions. Such a criterion was provided in Ref.~\cite{perturbekp}, although without a physical motivation, leading to the recovery of the observationally correct spectrum. The very exhibition of junction conditions leading to a scale invariant spectrum could then be seen as a hint that constructing a realistic theory satisfying observational constraints was not impossible. Even if one is prepared to accept such drastic changes in the standard cosmological picture, one might wonder as to the use of perturbation theory on top of an otherwise singular background~\cite{Lyth}. Moreover, it should be mentioned that although the old scenario, because describable as an effective bounce occurring at a low enough temperature, was avoiding the over-production of grand unified scales monopoles~\cite{monop}, the new model, being singular, poses this problem in a way which is as acute as it was before the advent of inflation. Finally, the puzzle of trans-Planckian scales~\cite{transPl}, quoted in Ref.~\cite{reply} as a caveat for inflation, can be transposed in the new ekpyrotic model in the same words. This article is organized as follows. After a brief reminder of the ekpyrotic model of the universe (Sec.~\ref{sec:ekp}), we examine in detail the junction conditions suggested in Ref.~\cite{perturbekp} (Sec.~\ref{sec:pert}). We concentrate in particular on the fact that this proposed criterion rests on an altogether arbitrary (hence unphysical) normalization function, so that whatever spectrum can be obtained: obtaining a scale invariant spectrum in this model thus turns out to be equivalent to imposing it from the outset. We also demonstrate that the new prescription leads to an incorrect prediction in the exact case of a radiation to matter domination transition. We then consider a second possibility, i.e., we examine an effective bounce in a context where the linear perturbation theory is still valid. We therefore considered first, in section~\ref{sec:hydro}, the simplest case in which not only does GR apply, but also in which all the calculations can be performed analytically and consistently (indeed providing a nice textbook example for cosmological perturbation theory illustration), namely that of a ${\cal K}=1$ bouncing universe with hydrodynamic perturbations~\cite{ppnpn1}. Then, using this toy model, we examine how the relevant perturbed quantities behave through the bounce. We pay special attention to the ``short time bounce limit'' (this is related to the question ``how sharp is sharp'' evoked in Ref.~\cite{Lyth}) and study whether, in this limit, the bounce can be considered as a surface where the equation of state jumps. If so this would allow us to use the standard junction conditions. The second example that one can treat completely is by considering a test scalar field. Indeed, in this case, one does not need to specify what the origin of the background evolution is. In section~\ref{sec:scalar}, we calculate the spectrum of a spectator scalar field in such a bouncing background. Assuming no strong deviation from GR at the perturbed level (we remind that such deviations are necessary in the bounce region), and $\Ka=0$, this also gives the tensor perturbation spectrum. The description of a bouncing universe with ${\cal K}=0$ requires special care, as GR does not allow for such a configuration to take place unless the Null Energy Condition (NEC) is violated. Although this case is clearly contrived, it provides at least an example where some arguments presented recently in the literature can be implemented concretely, at the level of equations.
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The conclusions that can be drawn from this work is that it seems impossible to apply any known and well motivated criterion to pass through a bounce, whether regular or singular, in a model independent way as all quantities of interest explicitly depend on the details of the underlying model. The ekpyrotic model, although a potentially interesting alternative to the inflationary paradigm, does pass through such a bounce. Therefore, if one really wants to calculate the spectrum in the ekpyrotic universe then it seems necessary, first, to consider a situation where there is no divergence and, second, to provide us with the actual (maybe five-dimensional) equations of motion during the bounce, knowing that these equations cannot be those of GR.
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hep-th0112128_arXiv.txt
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astro-ph0112565_arXiv.txt
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In this paper we propose that the accelerating expansion of the present matter-dominated universe, as suggested by the recent distance measurements of type Ia supernovae, is generated along with the evolution of space in extra dimensions. The Einstein equations are first analyzed qualitatively and then solved numerically, so as to exhibit explicitly these patterns of the accelerating expansion in this scenario. A fine-tuning problem associated with such a scenario is also described and discussed.
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The recent distance measurements of type Ia supernovae suggest an accelerating expansion of the present universe \cite{Perlmutter:1999np,Riess:1998cb}. In many of the current cosmological models, the present accelerating expansion is driven by an energy source called ``dark energy'', with a positive cosmological constant \cite{Lambda models} or the so-called ``quintessence'' (a slowly evolving scalar field \cite{Caldwell:1998ii,ComplexQ}) as a possible candidate. Instead of attributing this acceleration to the mysterious dark energy, we consider in this paper the possible existence of extra spatial dimensions and explore the feasibility of producing the present accelerating expansion via the evolution of these extra dimensions. The application of extra dimensions is a general feature in theories beyond the standard model, especially in theories for unifying gravity and other forces, such as superstring theory. These extra dimensions should be ``hidden'' for consistency with observations. Various scenarios for ``hidden'' extra dimensions have been proposed, for example, a brane world with large compact extra dimensions in factorizable geometry proposed by Arkani-Hamed \emph{et al.} \cite{Arkani-Hamed} (see also \cite{Antoniadis:1990ew}), and a brane world with noncompact extra dimensions in nonfactorizable geometry proposed by Randall and Sundrum \cite{Randall&Sundrum}. In this paper, we employ the simplest scenario: small compact extra dimensions in factorizable geometry, as introduced in the Kaluza-Klein theories \cite{Kaluza&Klein}. We study spatially homogeneous, isotropic, perfect-fluid cosmological models in $(1+3+n)$ dimensions where $n$ is the number of extra dimensions. In Sec.\ \ref{general features}, we first obtain, from the Einstein equations, some general features of the evolution of the higher-dimensional space, especially for a radiation-dominated universe and a (nonrelativistic) matter-dominated universe. In Sec.\ \ref{accel 3-space}, we then explore the possibility of producing an accelerating expansion of the ordinary three-space via the evolution of `extra space' for a matter-dominated universe. We analyze the Einstein equations to show qualitatively the behavior of this evolution and obtain numerical solutions which illustrate explicitly the accelerating expansion of the ordinary three-space along with the collapse of the extra space. We note that, while the Kaluza-Klein cosmology and inflation in higher dimensional space-time in connection with the early universe were discussed widely in the 1980s \cite{Freund:1982pg,KK-cosmology} (for a review, see \cite{Kolb:1990vq}; for recent work, see \cite{Arkani-Hamed:1999gq,Mongan:2001cr}), the focus of this paper has to do with the present accelerating matter-dominated universe.
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We have investigated the scenario of producing the accelerating expansion of the present universe via evolving small extra dimensions. For a radiation-dominated universe, such as our early universe, we obtain a stable solution with static extra dimensions. Accordingly, the existence of extra dimensions may have no significant influence on the evolution of the ordinary three-space. This is a good feature which we need for preserving the concordance between observations and current theories regarding the early (radiation-dominated) universe, especially for primordial nucleosynthesis. On the contrary, such a solution with static extra dimensions does not exist for the present matter-dominated universe. The features of the evolution can also be read off from Eq.\ (\ref{alpha-volume relation}), or Eqs.\ (\ref{alpha-volume for RD}) and (\ref{alpha-volume for MD}), which are derived from Eq.\ (\ref{alpha-volume relation}). Equation (\ref{alpha-volume for RD}) shows the decreasing expansion rate of the extra space along with the increase of the $(3+n)$-dimensional volume $V_{3+n}$. This implies the stability of the solution with static extra dimensions in the radiation-dominated universe for the case of $k_a=k_b=0$ as already mentioned above. On the other hand, Eq.\ (\ref{alpha-volume for MD}) shows two possible evolution patterns of the matter-dominated universe: (i) The expansion rates of the ordinary three-space and the extra space tend to catch up with each other along with the increase of the (3+n)-dimensional volume $V_{3+n}$. (ii) One of these two expansion rates is positive and increasing, while the other is negative and decreasing, along with the decrease of the (3+n)-dimensional volume $V_{3+n}$. We note that an increasing positive expansion rate represents an accelerating expansion. A quantitative analysis of the matter-dominated case with $k_a=k_b=0$ leads to Fig.\ \ref{accel-decel plot}, which indicates four classes of evolution path. A universe that decelerates first and then accelerates is included in one of them. Therefore the accelerating expansion of the present universe may be appropriately described in this scenario. In addition, the case with two extra dimensions is analyzed in detail. The five resultant evolution paths we draw demonstrate the existence of a critical value for the initial condition $\eta_0$, which divides two classes of path: the one in which the universe decelerates first and then accelerates and the other in which the universe always decelerates. We note that the critical value $\eta_{cr}$ is exactly the parameter $K_{rep}$, a ``repeller'' in the flow diagram. However, the existence of the critical value (or the ``repeller'') also implies a fine-tuning problem, i.e., the initial value of $\eta$ has to be chosen delicately so that it is close enough to the critical value $\eta_{cr}$ in order to possess a long enough decelerating epoch followed by an accelerating epoch. The existence of extra dimensions is a general feature in theories beyond the standard model in particle physics. It may manifest itself as a source of energy in the ordinary three-space, such as ``effective'' dark energy or even ``effective'' dark matter. The geometrical structure and the evolution pattern of extra dimensions therefore may play an important role in cosmology. In this work we study a simple scenario of extra dimensions that is subject to a fine-tuning problem. Nevertheless, other scenarios with richer structures, such as those in \cite{Arkani-Hamed} and \cite{Randall&Sundrum}, may also provide suitable evolution patterns and are worthy of being further investigated. {\em Note added}. For the sake of simplicity, we have in this paper considered only the case with $k_a=k_b=0$, i.e., both our ordinary three-space and the extra space are flat. The general situations with nonzero $k_a$ or $k_b$ clearly may offer many interesting possibilities and are currently under serious investigation.
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astro-ph0112565_arXiv.txt
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astro-ph0112086_arXiv.txt
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We present new {\it HST} I-band images of a sample of 77 nearby, late-type spiral galaxies with low inclination. The main purpose of this catalog is to study the frequency and properties of nuclear star clusters. In 59 galaxies of our sample, we have identified a distinct, compact (but resolved), and dominant source at or very close to their photocenter. In many cases, these clusters are the only prominent source within a few kpc from the galaxy nucleus. We present surface brightness profiles, derived from elliptical isophote fits, of all galaxies for which the fit was successful. We use the fitted isophotes at radii larger than $2\as$ to check whether the location of the cluster coincides with the photocenter of the galaxy, and confirm that in nearly all cases, we are truly dealing with ``nuclear'' star clusters. From analytical fits to the surface brightness profiles, we derive the cluster luminosities after subtraction of the light contribution from the underlying galaxy disk and/or bulge.
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Over the past decade, high dynamic range observations with modern CCD detectors have shown that compact stellar nuclei are a common feature of spiral galaxies of all Hubble types. For example, \cite{mat97} found 10 objects with compact nuclear star clusters in a survey of 49 southern, very late-type spirals. However, as one progresses along the Hubble sequence towards earlier types, the increasingly luminous bulge component with its steeply rising surface brightness profile makes the identification of an additional, unresolved cluster extremely difficult. It therefore took the unique spatial resolution of the {\it Hubble Space Telescope (HST)} to demonstrate that \nc s are a common phenomenon also in earlier Hubble types \citep*[e.g.][]{car98}. The {\it HST} currently provides the only means to investigate the structural properties of nuclear star clusters, as demonstrated by \cite{mat99}, and to cleanly separate their emission from the underlying galaxy disk/bulge. Despite the recent progress, the formation mechanism of nuclear star clusters remains largely a mystery. Intuitively, there are good reasons to expect matter accumulation in the deep gravitational wells of galaxies with massive bulges, and hence active star formation in their nuclei. In contrast, the gravitational force all but vanishes in the centers of pure disk galaxies with shallow surface brightness profiles and without any discernible bulge component. In these galaxies, the dynamical center is not a ``special'' place and it is far from obvious how a massive stellar cluster could have formed there. The shallow gravitational potential might provide a natural explanation for the fact that spirals of late Hubble type are not known to contain super-massive black holes. On the other hand, nuclear star clusters can be extremely compact: the nucleus of M33, for example, has likely undergone core collapse and is as compact as any known globular cluster \citep{kor93}. So far, no satisfying explanation has been put forward to explain the high gas densities that must have been present in the nuclei of these shallow disk galaxies to enable the formation of such massive and compact objects. It is also unknown whether nuclear star clusters form repeatedly or only once - a question with important implications for the dynamical and morphological evolution of their host galaxies. To make progress along this line, it is essential to obtain the age distribution of nuclear star clusters. So far, reliable age estimates exist for only a handful of nuclear star clusters. Interestingly, most of them appear to be rather young: our Galaxy has a central stellar cluster with an age of only $\sim 3$~Myrs \citep{kra95}, and both M31 and M33 have blue nuclei that are very likely young star clusters \citep{lau98}. More recently, we have published \nc\ ages derived from ground-based spectroscopy for IC~342 \citep*[$\rm \leq 60\>$Myrs,][]{boe99}, and NGC~4449 \citep*[$6-10\>$Myrs,][]{boe01,gel01}. In addition, the dominant stellar population of the nuclear cluster in NGC 3227 is less than 50~Myrs old \citep*{sch01}. However, it is possible (and in fact quite likely) that ground-based observations predominantly target the brightest and hence youngest clusters. In order to get a more representative picture of nuclear star clusters, it is important to study a galaxy sample which is free from selection effects that favor the high end of the \nc\ luminosity range. In this paper, we describe the results of an {\it HST} I-band imaging survey of an unbiased sample of nearby, face-on, very late-type spirals (Scd or later). The main goals of the survey are (a) to determine the frequency of nuclear star clusters in very late-type spirals, (b) to derive their luminosity and size distribution, (c) to compare their properties to those of nuclear star clusters in earlier Hubble types which have been more extensively studied with {\it HST} \citep{car97,car98,car01,car99}, and (d) to provide a source catalog for follow-up spectroscopic observations to age-date their stellar populations. The main purpose of this paper is to present the complete dataset. In a companion paper (B\"oker \ea\ 2002, in preparation), we describe the statistics of the full sample and investigate whether the properties of nuclear star clusters correlate in any way with those of their host galaxies. This paper is organized as follows: in \S~\ref{sec:obs}, we describe our sample selection criteria, the observational strategy, and the data reduction procedure, and we present the final images as well as the results of the isophotal analysis. In \S~\ref{sec:disc}, we discuss whether the clusters indeed occupy the nuclei of their host galaxies, and how they compare to other luminous star clusters observed in a variety of starburst environments. We conclude in \S~\ref{sec:concl}.
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\label{sec:concl} We have presented a catalog of {\it HST}/WFPC2 I-band images of an unbiased sample of 77 nearby, late-type spiral galaxies with low inclination. From an isophotal analysis of the images, we demonstrate that about 75\% of the sample galaxies host a compact, luminous stellar cluster at or very close to their photocenter. These clusters often are completely isolated from other comparable structures, emphasizing that even in the relatively shallow potential wells of late-type galaxy disks, the center is well-defined, and has a unique star formation history. From analytical fits to the surface brightness profiles, we determine the flux attributable to the cluster. The distribution of absolute cluster luminosities has a FWHM of 4 magnitudes, and a median value of $\rm M_I = -11.5$, comparable to young super star clusters in starbursting galaxies. Together with initial estimates of their size distribution, this suggests that \nc s in spiral galaxies of the latest Hubble types are a fairly homogenous class of objects. The dataset is a representative survey of late-type spiral galaxies in the local universe, and as such yields a valuable source catalog for spectroscopic follow-up observations which are needed to further constrain the star formation history of \nc s. We have begun such a follow-up program both with {\it HST} and ground-based observatories.
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astro-ph0112086_arXiv.txt
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astro-ph0112423_arXiv.txt
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One of the most important quantities in theoretical models of AGNs is the black hole mass (M$_{BH}$) that, together with the total luminosity, defines the fraction of the Eddington luminosity at which the AGN is emitting. Determination of M$_{BH}$ in AGN is difficult mainly because of the bright emission from the nucleus and their large distance. The main method that has proved to be successful in AGN is reverberation mapping, which is extremely time consuming and gives results on M$_{BH}$ that depend on the assumed geometry of the accretion disk. Therefore, only for a few well studied quasars and Seyfert galaxies M$_{BH}$ is known (see e.g. [7], [11] and references therein). This method cannot obviously be employed for BL Lac objects because they lack prominent emission lines. Therefore other methods need to be applied to infer M$_{BH}$ for BL Lacs. The discovery of a relation between M$_{BH}$ and the luminosity of the bulge in nearby early-type galaxies offers now a new tool for estimating the mass of the central BH (see e.g. review [10]). This has been done for two samples of nearby quasars [9] and BL Lacs [12]. Recently, a tighter correlation was found relating M$_{BH}$ with the central stellar velocity dispersion $\sigma$ of the spheroidal component in nearby galaxies [5,2], that can also be used to estimate M$_{BH}$ in AGN. The relationship appears to predict more accurately [10] M$_{BH}$, but requires the measurement of $\sigma$ in the host galaxies of AGN that is difficult to obtain, in particular for objects at moderately high redshift and with very luminous nuclei. On the other hand, for BL Lacs that have relatively fainter nuclei than quasars, this measurement (at least for low redshift objects) can be secured with observations at medium-sized telescopes. We present here the first estimates of stellar velocity dispersion of BL Lacs from our ongoing program aimed specifically at deriving M$_{BH}$ from the M$_{BH}$ -- $\sigma$ correlation. We selected a sample of nearby (z$<$0.2) BL Lacs for which high quality images were obtained either from the ground using the Nordic Optical Telescope (NOT) [3] or with HST+WFPC2 [13,4]. From the images a characterization of the host galaxies and of the nuclear luminosity are obtained. This allows us to compare M$_{BH}$ with the mass (and the luminosity) of the host galaxy and also to evaluate the Eddington ratio, provided that the nuclear emitted power is corrected for the beaming factor. Moreover, a comparison of M$_{BH}$ for BL Lacs with different jet/ disk luminosities can be used to test the hypothesis (see e.g. [8]) that the accretion rate changes from largely sub-Eddington, for low luminosity weak-lined sources, to near-Eddington for high luminosity, strong-lined sources. If the accretion rate in terms of Eddington ratio were the same in both classes, the BH masses should differ almost by three orders of magnitude.
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astro-ph0112423_arXiv.txt
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astro-ph0112437_arXiv.txt
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We study the dynamics of phase transitions in the interstellar medium by means of three-dimensional hydrodynamic numerical simulations. We use a realistic cooling function and generic nonequilibrium initial conditions to follow the formation history of a multiphase medium in detail in the absence of gravity. We outline a number of qualitatively distinct stages of this process, including a linear isobaric evolution, transition to an isochoric regime, formation of filaments and voids (also known as ``thermal'' pancakes), the development and decay of supersonic turbulence, an approach to pressure equilibrium, and final relaxation of the multiphase medium. We find that 1\%-2\% of the initial thermal energy is converted into gas motions in one cooling time. The velocity field then randomizes into turbulence that decays on a dynamical timescale $E_k\propto t^{-\alpha}$, $1\lsim\alpha\lsim2$. While not all initial conditions yield a stable two-phase medium, we examine such a case in detail. We find that the two phases are well mixed with the cold clouds possessing a fine-grained structure near our numerical resolution limit. The amount of gas in the intermediate unstable phase roughly tracks the {\em rms} turbulent Mach number, peaking at 25\% when ${\cal{M}}_{rms}\sim8$, decreasing to 11\% when ${\cal{M}}_{rms}\sim0.4$.
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Thermal instability (TI) has many implications in astrophysics (e.g., a clumpy interstellar medium [ISM], stellar atmospheres, star formation, globular cluster and galaxy formation, etc.; see Meerson 1996 for a recent review). The instability may be driven by radiative cooling of optically thin gas (radiation-driven TI) or by exothermic nuclear reactions \citep{schwarzschild.65}. Linear stability theory for a medium with volumetric sources and sinks of energy in thermal equilibrium was developed by \citet{field65}, who identified three unstable modes: the {\em isobaric} mode (the pressure-driven formation of condensations not involving gravitation) and the two {\em isentropic} modes (the overstability of acoustic waves propagating in opposite directions). Hunter (1970, 1971) extended these results to an arbitrary nonstationary background flow, showing that cooling-dominated media are potentially more unstable than those in equilibrium, while heating provides stabilization. The most common applications of thermal instability to the ISM and star formation deal with the isobaric mode that was employed to explain the observed multiphase structure of the ISM \citep{pikel'ner68,field..69,mckee.77,mckee90,wolfire....95}. Analysis of infinitesimal perturbations gives two characteristic length scales for the isobaric condensation mode: (1) a cooling scale $\lambda_{\rmn{p}}=c/\omega_{\rmn{p}}$ (where $c$ is the adiabatic sound speed and $\omega_{\rmn{p}}$ is the growth rate) and (2) a critical scale $\lambda_{\kappa}=c\sqrt{t_{d}/\omega_{\rmn{p}}}$ (where $t_{d}$ is the characteristic thermal diffusion time). These two length scales define short-, intermediate-, and long-wavelength limits \citep{meerson96,kovalenko.99}. In the short-wavelength limit, small isobaric perturbations are inhibited by heat conduction, so that $\omega_{\rmn{p}}<0$ for $\lambda<\lambda_{\kappa}$. In the long-wavelength limit, the perturbations cannot grow isobarically because of the finite sound speed effects, and thus $\omega_{\rmn{p}}\rightarrow0$ for $\lambda/\lambda_{\rmn{p}}\rightarrow\infty$. This is only true if the gas is {\em isochorically} stable \citep{parker53,field65,shchekinov78}, otherwise the growth rate remains finite: $\omega_{\rmn{p}}\rightarrow\omega_{\rmn{\rho}}>0$ for $\lambda/\lambda_{\rmn{p}}\rightarrow\infty$, but only large-scale temperature perturbations are growing, thus resulting in pressure variations and the formation of shock waves. The growth rates and characteristic scales depend on the heating and cooling properties of a given medium. Under the ISM conditions, if one assumes thermal equilibrium (i.e., an exact balance between cooling and heating), isochoric instability manifests itself only at relatively high temperatures, $T \gsim 10^5$~K. However, in cooling-dominated regimes, it can develop at temperatures as low as $10^3$~K. \begin{figure*} \epsscale{2.0} \plotone{fig1.ps} \caption{Snapshots of the gas density field (perspective volume rendering): (a) First condensation at $t=0.07$~Myr, (b) thermal pancakes at $t=0.1$~Myr, (c) collapse and turbulization of cellular structure at $t=0.17$~Myr, (d) two-phase medium at $t=1.5$~Myr (5~pc box, $256^3$ grid points). The log density color coding is as follows: The dense blobs at the intersections of the filaments, $\rho>10^{-22}$~g cm$^{-3}$, are light blue; the stable cold phase, $\rho\in[10^{-23}, 10^{-22}]$~g cm$^{-3}$, is blue; the unstable density regime, $\rho\in[10^{-23.7}, 10^{-23.0}]$~g cm$^{-3}$, is yellow to brown; and the low-density gas, including stable warm phase ($\rho<10^{-23.7}$~g cm$^{-3}$), is a transparent red. The figure is also available as an mpeg animation in the electronic edition of the Astrophysical Journal. \label{fig1}} \end{figure*} A specific feature of TI in the ISM is a large ($\gsim1$ dex) gap in gas densities between the two stable phases. The density range of interest in a galaxy formation context is even larger. This implies the importance of nonlinear effects in the dynamics of phase transitions. Nonlinearity brings into play nonequilibrium effects. Already weakly nonlinear development of condensations in an initially homogeneous gas in thermal equilibrium drives the system away from equilibrium. The mean pressure drops since $\bar{\rho^2} > \bar{\rho}^2$ and cooling overcomes heating globally \citep{kritsuk85}. Later, on a timescale of $\sim\omega_{\rmn{p}}^{-1}$, as condensations get denser and cooler, the isobaric condition $\lambda\ll\lambda_{\rmn{p}}$ becomes violated locally within them, so the system departs from pressure equilibrium. These effects are essential for the isobaric mode of TI in the ISM and star formation contexts. Therefore, analytical nonlinear solutions to ``isobarically'' reduced TI equations are insufficient to describe the radiative stage of the phase transition ($t\sim\omega_{\rmn{p}}^{-1}$). During this strongly nonlinear stage {\em large-scale} condensations form in such a way that gas moves almost inertially and its kinetic energy dominates thermal energy ($p\ll \rho v^2$). Accordingly, the gas velocities in these condensations are of the order of the sound speed in the unperturbed state. The situation here is directly analogous to the long-wave gravitational instability, so that results concerning the origin of cellular structure and \citet{zel'dovich70} ``pancakes'' can be entirely carried over to the case of long-wave TI \citep{meerson.87}. \citet{sasorov88} gave an elegant proof that qualitatively the same result applies to {\em small-scale} TIs; i.e., the onset of TI produces voids and highly flattened condensations along certain two-dimensional surfaces. These are also called ``thermal'' pancakes \citep{meerson96}. The formation of filaments was simultaneously noticed in two-dimensional numerical simulations of TI in the solar transition region \citep{dahlburg....87,karpen..88}. Ever since, thermal pancakes are being rediscovered both analytically and in numerical simulations (e.g., Lynden-Bell \& Tout 2001). Thermal pancakes are transient. However, what happens next, before the evolution turns to a conductive relaxation stage \citep{meerson96}, until recently has remained the ``terra incognito'' of TI theory. The problem of ``postradiative'' mechanical relaxation toward a static multiphase medium requires a solution for the full set of hydrodynamic equations that can only be obtained numerically. One-dimensional simulations pioneered by \citet{goldsmith70} demonstrated that TI develops large motions in the ISM (see also Hennebelle \& P{\'e}rault 1999 for an example of how large motions can trigger TI). For some time, progress in this direction was precluded by numerical difficulties in modeling convergent cooling flows with high Mach numbers and high-density contrasts (e.g., V{\'a}zquez-Semadeni, Gazol, \& Scalo 2000). Recent multidimensional numerical simulations of the ISM evolution in disklike galaxies include effects of gravity, differential rotation, star formation and supernova feedback (de Avillez 2000; V{\'a}zquez-Semadeni et al. 2000; Wada \& Norman 2001; Wada 2001). However, it is hard to determine the role of TI in shaping the ISM structures found in these models, partly because the simulations still do not resolve length scales important for TI and partly because of the additional physical effects. The purpose of this Letter is to report on results of three-dimensional numerical simulations of classical TI that fill the gap in theory, exploring in detail the late radiative stage and postradiative relaxation toward a multiphase medium. Our major result is that formation of thermal pancakes induces turbulence in the ISM that serves as a nonlinear saturation mechanism for TI. As a consequence of a turbulent cascade, (1) information about initial perturbations is lost, including the imprints of heat conduction in the density power spectrum during the linear stage, and (2) turbulent diffusion becomes the dominant transport mechanism during the postradiative relaxation stage. \begin{figure*} \epsscale{1.8} \vspace{-1.8cm} \plottwo{fig2a.ps}{fig2b.ps} \caption{Snapshots of phase diagrams (timing and labels correspond to those in Fig. \ref{fig1}): (a) $t=0.07$~Myr, (c) $t=0.17$~Myr, (e) $t=0.5$~Myr, (d) $t=1.5$~Myr. The black dots show scatter plots of pressure vs. density. The ``dash'' at $P=4.55\times10^{-10}$~dyn~cm$^{-2}$ in (a) shows the isobaric initial conditions. Background yellow-filled contours specify the part of the phase plane where isobaric mode is unstable; overlaid magenta contours are the regions of isochoric instability. The thick solid line shows thermal equilibrium curve. Density PDFs are plotted at the bottom of each panel (see scale to the right). \label{fig2}} \end{figure*}
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Our fiducial case was constructed to produce a stable two-phase medium because of its relevance to the Galactic ISM. We are interested in TI over a wide range of conditions as might be found in the ISM of high-redshift protogalaxies. We have simulated other cases with different parameter choices that do not produce stable two-phase media. However, we find they all develop turbulence in the nonlinear radiative stage of TI. Here we briefly discuss how the turbulence and asymptotic phase structure depend on initial conditions, deferring a more complete discussion to a future paper (A. Kritsuk \& M. Norman 2002, in preparation). The level of induced turbulence is determined by the efficiency of conversion of the initial thermal gas energy into kinetic energy of turbulent flow by nonlinear development of TI: $E_k^{max}={\cal{C}}(\rho_0, T_0, \varepsilon, Q, L)E_{th}(0)$ (see Fig. \ref{fig3}). The conversion factor ${\cal{C}}$ is a complex function of its variables. It varies from about 2\% to $\lsim1$\% in our models. In general, higher $\varepsilon$ and/or $T_0$ values provide higher conversion; a lower heating level ($Q<1$) supports TI and therefore works in the same direction. Larger boxes, as a rule, also produce more turbulence. The turbulence is induced on the initial cooling time and decays on a dynamical timescale, which is typically much longer. The turbulent Mach number at $t\sim\omega^{-1}_{\rmn{p},0}$ depends on the mean temperature at this time. For nonequilibrium initial conditions ($Q\ll1$ or $Q=0$), this is much lower than the initial temperature. In our fiducial case, ${\cal{M}}_{rms}$ peaks at 8, dropping to 0.4 after 20 initial cooling times. For equilibrium initial conditions ($Q=1$), turbulent velocities remain subsonic (${\cal{M}}_{rms}\sim0.3$). But we would expect higher Mach numbers if the bistable range of pressure were wider than provided by our adopted cooling function. The initial gas density determines the number and mass fractions of thermal phases in the {\em relaxed} state depending on its position relative to the valleys and hills on the thermal equilibrium curve. This is consequence of our choice of constant volume boundary conditions, which means that the mean density in the box remains constant. After the rapid cooling stage, our models with low initial densities $\rho_0=1$-$5\times10^{-25}$~g~cm$^{-3}$, high temperature $T_0=2\times10^6$~K, and $Q\in\{0.3, 1\}$ generate turbulence, evolve through a transient three-phase stage, and then relax to a single-phase low-pressure warm ISM. While turbulence is a generic feature of nonlinear saturation of TI, our simulations show that {\em detailed} turbulent properties and the nature of emerging multiphase medium do depend sensitively on the Mach number and effective equation of state controlled by heating and cooling; this will be discussed elsewhere. Two identical simulations, except that cutoffs in initial power spectra were different ($k_{max}=8$ and 32 on a $128^3$ grid, $L=100$~pc, $Q=0$), demonstrated considerable structural differences in density distributions at the thermal pancake stage, $t_{tp}$, and surprisingly similar ``chaotic'' density structures and identical velocity power spectra at $\sim6\,t_{tp}$, when turbulent mixing covered the whole computational domain. This implies that the imprints of heat conduction in the density power spectrum during the linear stage could be erased later by the developing turbulent cascade. TI is certainly not the only potential source of turbulence in the ISM, but it cannot be ignored at least in those scenarios that actively employ TI to explain the origin and properties of observed objects. We suggest a paradigm shift concerning the role of thermal instability in the ISM and the nature of multiphase ISM. The idea of ``static'' two-phase ISM introduced in late 1960s (pressure-confined thermally stable dense clouds embedded in rarefied intercloud gas forming as a result of TI and subject to phase exchange due to cloud evaporation/condensation) must give way to the notion of a dynamic multiphase ISM, in which TI induces slowly decaying turbulence and in which turbulent diffusion regulates phase exchange processes. In this new emerging picture, the dense clouds are shapeless random aggregations of cold Lagrangian gas parcels; the clouds do not preserve their identity in real space on their sound crossing timescale until self-gravity tightens the fragments up into a self-gravitating cloud to form stars. Our results may suggests modifications to the scenario of a three-phase ISM \citep{mckee.77,mckee90,heiles01a} that are yet to be understood.
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astro-ph0112437_arXiv.txt
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astro-ph0112092_arXiv.txt
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{ A systolic algorithm rhythmically computes and passes data through a network of processors. We investigate the performance of systolic algorithms for implementing the gravitational $N$-body problem on distributed-memory computers. Systolic algorithms minimize memory requirements by distributing the particles between processors. We show that the performance of systolic routines can be greatly enhanced by the use of non-blocking communication, which allows particle coordinates to be communicated at the same time that force calculations are being carried out. The performance enhancement is particularly great when block sizes are small, i.e. when only a small fraction of the $N$ particles need their forces computed in each time step. Hyper-systolic algorithms reduce the communication complexity from $O(Np)$, with $p$ the processor number, to $O(N\sqrt{p})$, at the expense of increased memory demands. We describe a hyper-systolic algorithm that will work with a block time step algorithm and analyze its performance. As an example of an application requiring large $N$, we use the systolic algorithm to carry out direct-summation simulations using $10^6$ particles of the Brownian motion of the supermassive black hole at the center of the Milky Way galaxy. We predict a 3D random velocity of $\sim 0.4$ km s$^{-1}$ for the black hole. } \begin{article}
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\label{sec:intro} Numerical algorithms for solving the gravitational $N$-body problem have evolved along one of two lines in recent years. Direct-summation codes compute the complete set of $N^2$ interparticle forces at each time step; these codes were designed for systems in which the finite-$N$ graininess of the potential is important, and are limited by their $O(N^2)$ scaling to moderate ($N\lap 10^5$) particle numbers. The best-known examples are the {\tt NBODY} series of codes introduced by Aarseth \cite{Aarseth:99a}. These codes typically use high-order schemes for integration of particle trajectories and avoid the force singularities at small interparticle separations either by softening, or by regularization of the equations of motion \cite{KS:65}. A second class of $N$-body algorithms replace the direct summation of forces from distant particles by an approximation scheme. Examples are tree codes \cite{Barnes:89}, which reduce the number of direct force calculations by collecting particles in boxes, and algorithms which represent the large-scale potential via a truncated basis-set expansion (e.g. \cite{Allen:90}) or on a grid (e.g. \cite{Miller:68}). These algorithms have a milder, $O(N\log N)$ scaling for the force calculations and can handle much larger particle numbers although at some expense in decreased accuracy \cite{Spurzem:99}. The efficiency of both sorts of algorithm can be considerably increased by the use of individual time steps for advancing particle positions, since many astrophysically interesting systems exhibit a ``core-halo'' structure characterized by different regions with widely disparate force gradients. An extreme example of a core-halo system is a galaxy containing a central black hole \cite{Merritt:99}. The efficiency of individual time steps compared with a global time step has rendered such schemes standard elements of direct-summation codes (e.g. \cite{Aarseth:99b}). Here we focus on direct-summation algorithms as implemented on multi-processor, distributed-memory machines. Applications for such codes include simulation of globular star clusters, galactic nuclei, or systems of planetesimals orbiting a star. In each of these cases, values of $N$ exceeding $10^5$ would be desirable and it is natural to investigate parallel algorithms. There are two basic ways of implementing a parallel force computation for $O(N^2)$ problems. 1. {\bf Replicated data algorithm.} All of the particle information is copied onto each processor at every time step. Computing node $i$, $1\le i\le p$, computes the forces exerted by the entire set of $N$ particles on the subset of $n_i=N/p$ particles assigned to it. 2. {\bf Systolic algorithm.} At the start of each time step, each computing node contains only $N/p$ particles. The sub-arrays are shifted sequentially to the other nodes where the partial forces are computed and stored. After $p-1$ such shifts, all of the force pairs have been computed. \noindent (The term ``systolic algorithm'' was coined by H. T. Kung \cite{Kung:78,Kung:82} by analogy with blood circulation.) Both types of algorithm exhibit an $O(Np)$ scaling in communication complexity and an $O(N^2)$ scaling in number of computations. The advantage of a systolic algorithm is its more efficient use of memory: since each processor stores only a fraction $1/p$ of the particles, the memory requirements are minimized and a larger number of particles can be integrated. Other advantages of systolic algorithms include elimination of global broadcasting, modular expansibility, and simple and regular data flows \cite{Kung:82}. The performance of a systolic algorithm suffers, however, whenever the number of particles on which forces are being computed is less than the number of computing nodes. This is often the case in core-halo systems since only a fraction of the particles are advanced during a typical time step. As an extreme example, consider the use of a systolic algorithm to compute the total force on a {\it single} particle due to $N$ other particles. Only one processor is active at a given time and the total computation time is \begin{equation} N\tau_f + p(\tau_l + \tau_c) \end{equation} where $\tau_f$ is the time for one force calculation, $\tau_l$ is the latency time required for two processors to establish a connection, and $\tau_c$ is the interprocessor communication time. Thus the algorithm is essentially linear and no advantage is gained from having multiple processors. An efficient way to deal with the problem of small group sizes in systolic algorithms is via {\it nonblocking communication}, a feature of MPI that allows communication to be put in the background so that the computing nodes can send/receive data and calculate at the same time \cite{MPI:98}. In a nonblocking algorithm, the time per force loop for a single particle becomes \begin{equation} {N\tau_f\over p} + p(\tau_l + \tau_c). \end{equation} The second term is the waiting time for the last computing node to receive the particle after $p$ shifts. The first term is the time then required to compute the forces from the subset of $N/p$ particles associated with the last node. As long as the calculation time is not dominated by interprocessor communication, the speedup is roughly a factor of $p$ compared with the blocking algorithm. Here we discuss the performance of systolic algorithms as applied to systems with small group sizes, i.e. systems in which the number of particles whose positions are advanced during a typical time step is a small fraction of the total. Section 2 presents the block time step scheme and its implementation as a systolic algorithm. Section 3 discusses the factors which affect the algorithm's performance, and Section 4 presents the results of performance tests on multiprocessor machines of blocking and nonblocking algorithms. Section 5 presents a preliminary discussion of ``hyper-systolic'' algorithms with block time steps, which achieve an $O(N\sqrt{p})$ communication complexity at the cost of increased memory requirements. Finally, Section 6 describes an application of our systolic algorithm to a problem requiring the use of a large $N$: the gravitational Brownian motion of a supermassive black hole at the center of a galaxy.
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We have introduced two variants of a systolic algorithm for parallelizing direct-summation $N$-body codes implementing individual block time step integrators: the first with blocking communication, and the second with non-blocking communication. Performance tests were carried out using $N$-body models similar to those commonly studied by dynamical astronomers, in which the gravitational forces vary widely between core and halo and for which the particle block sizes are typically very small. The nonblocking scheme was found to provide far better performance than the blocking scheme for such systems, providing a nearly ideal speedup for the force calculations. By engaging a sufficient number of computing nodes, particle numbers in excess of $10^6$ are now feasible for direct $N$-body simulations. For parallel machines with very large processor numbers, we describe the implementation of a hyper-systolic computing scheme which provides a communication scaling of $O(\sqrt{p})$ at the expense of increased memory demands. The codes used to write this paper are available for download at:\\ \verb+http://www.physics.rutgers.edu/~marchems/download.html+ \begin{acknowledgment} This work was supported by NSF grant 00-71099 and by NASA grants NAG5-6037 and NAG5-9046 to DM. We thank Th.~Lippert, W.~Schroers, and K.~Schilling for their encouragement and help. We are grateful to the NASA Center for Computational Sciences at NASA-Goddard Space Flight Center, the National Computational Science Alliance, the Center for Advanced Information Processing at Rutgers University, the John von Neumman Institut f\"ur Computing in J\"ulich, and the H\"ochstleistungsrechenzentrum in Stuttgart for their generous allocations of computer time. \end{acknowledgment}
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astro-ph0112092_arXiv.txt
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astro-ph0112571_arXiv.txt
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We have computed quasiequilibrium sequences of synchronously rotating compact binary star systems with constant rest masses. This computation has been carried out by using the numerical scheme which is different from the scheme based on the conformally flat assumption about the space. Stars are assumed to be polytropes with polytropic indices of $N=0.5$, $N=1.0$, and $N=1.5$. Since we have computed binary star sequences with a constant rest mass, they provide approximate evolutionary tracks of binary star systems. For relatively stiff equations of state ($N < 1.0$), there appear turning points along the quasiequilibrium sequences plotted in the angular momentum --- angular velocity plane. Consequently secular instability against exciting internal motion sets in at those points. Qualitatively, these results agree with those of Baumgarte {\it et al.} who employed the conformally flat condition. We further discuss the effect of different equations of state and different strength of gravity by introducing two kinds of dimensionless quantities which represent the angular momentum and the angular velocity. Strength of gravity is renormalized in these quantities so that the quantities are transformed to values around unity. Therefore we can clearly see relations among quasiequilibrium sequences for a wide variety of strength of gravity and for different compressibility.
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Binary neutron stars are very interesting objects. From the observational point of view, we will have a chance to get new eyes for the Universe by detecting gravitational waves in the first or second decade of this century. It is highly possible that the first signal may be that from compact binary stars, such as binary neutron stars, a black hole --- neutron star binary system, or binary black holes. On the other hand, theoretically, our understanding of evolution of compact binary stars is far from complete because it is considerably difficult to treat the ``2-body'' problem from a state with a wide separation to a merging stage consistently in the framework of general relativity. However, recent investigations have found a new approach to this problem. Since the time scale of the orbital change due to gravitational wave emission is rather long compared with the orbital period except for in the final few milliseconds of the coalescing stage, we can neglect gravitational wave emission for most stages of evolution. In other words, we can treat the system in ``quasiequilibrium'' (see e.g., \cite{WMM96,BCSST98a,BGM99}). Following this idea, several groups have obtained quasiequilibrium sequences of binary neutron stars~\cite{WMM96,BCSST98a,BGM99,UE00,UUE00}. Most of them adopted the assumption that the spatial part of the metric is conformally flat (the conformally flat condition: hereafter, CFC)~\cite{WMM96,BCSST98a,BGM99,UE00}. For axisymmetric rotating polytropes results of the scheme with the CFC were compared with those obtained by the numerically exact code and found to be reasonably accurate~\cite{CST96}. However, since there are no exact numerical solutions for binary configurations, one could not know the accuracy of the results obtained by the scheme with the CFC. Therefore, it is desirable to develop different schemes from that with the CFC and to compare results of different schemes for nonaxisymmetric configurations. As one of those alternatives, in the previous paper, we developed a new numerical scheme to obtain quasiequilibrium structures of nonaxisymmetric compact stars as well as the space time around those stars in general relativity and obtained quasiequilibrium sequences of synchronously rotating binary polytropes~\cite{UUE00}. In that scheme, the Einstein equations are solved directly without assuming the CFC. The obtained results, however, could not be compared with those of \cite{BCSST98a} because different polytropic relations were used. In this paper, we have used the same polytropic equation as that used in~\cite{BCSST98a} and computed quasiequilibrium sequences of synchronously rotating polytropes with a constant rest mass. Therefore, we can directly compare our results with those with the CFC. In actual computations, we have improved our numerical scheme and succeeded in making our scheme more robust.
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In this paper we have solved quasiequilibrium sequences of synchronously rotating binary star systems in general relativity without assuming the CFC. We have constructed the constant rest mass sequences and shown that for stiff equations of state ($N < 1.0$), evolutionary curves have turning points so that synchronous rotation of the system breaks down at that point and that the internal motion will be excited. Our results can be compared with those of Baumgarte {\it et al.} who employed the CFC~\cite{BCSST98a}. Quantitatively, there are some differences between two results as seen from Figs.~\ref{J-Omega N=1.0} and \ref{J-Omega N=1.5}. These differences may come from different choices of the metric. However, it should be noted that qualitative features are very similar, i.e., the dependency on the polytropic index of the appearance of the turning points and so on. Therefore, although it is hard to give exact values of the angular velocity and/or the angular momentum at the turning points from quasiequilibrium approach, the occurrence of the instability could be correctly predicted. Nevertheless, since there are no exact solutions for the binary neutron star systems, we should keep in mind that there is a possibility that both of the two results might not represent the exact solutions. In Figs.~\ref{dj-dm N=0.5} and \ref{dj-dm N=1.0}, the nondimensional gravitational mass and the nondimensional angular momentum of each star are plotted against the nondimensional angular velocity. From these figures, it can be seen that turning points, i.e., the minima of each value, of two curves coincide. It implies that secular stability of binary systems can be found by investigating either the gravitational mass or the angular momentum. This is a nice feature that agrees the requirement between the change of the gravitational mass and that of the angular momentum as follows: \begin{equation} dM = \Omega dJ \ , \label{dj-dm} \end{equation} where $dM$ and $dJ$ are the changes of the gravitational mass and the angular momentum of two configurations with the same rest mass, respectively. This relation can be reduced from the first law of thermodynamics, which is shown below, for the binary systems for which the rest mass, entropy, and vorticity of each fluid element are conserved~\cite{FKM01}: \begin{equation} dE = \Omega dJ \ , \label{dj-de} \end{equation} where $E$ means the half of the total energy of the system. It should be noted that these requirements can be checked if we can obtain highly accurate models. As seen from Figs.~\ref{dj-dm N=0.5} and \ref{dj-dm N=1.0}, changes of the gravitational mass and the angular momentum are three or four orders of magnitude smaller than the corresponding quantities. Unfortunately, since we cannot insist that our values have such high accuracy, we do not show our results here. As seen from Figs.~\ref{J-Omega N=1.0}, \ref{J-Omega N=1.5} and \ref{J-Omega N=0.5}, the ranges of the values of $\bar{M}_0 \bar{\Omega}$ and $\bar{J}$ are considerably wide for the values of $N$. Even for the sequences with the same $N$, the values of $\bar{M}_0 \bar{\Omega}$ and $\bar{J}$ range widely. Thus it is not easy to understand the effects of the strength of gravity and/or the equation of state. In order to see the features of the evolutionary sequences at a glance, we introduce the following two nondimensional quantities, one of which can be considered to represent the angular velocity and the other of which corresponds to the angular momentum: \begin{eqnarray} \label{New J} \hat{j} &\equiv& \frac{J}{J_0} \ , \\ \label{New omega} \hat{\omega} &\equiv& \frac{\Omega}{\Omega_0} \ , \end{eqnarray} where \begin{eqnarray} \label{normalization J} J_0 &\equiv& \frac{7}{5} G^{1/2} M^2 \left(\frac{M}{R}\right)^{-1/2} = \frac{7}{5} M \frac{GM}{c^2} c \left( \frac{GM}{c^2R} \right)^{-1/2} \ , \\ \label{normalization Omega} \Omega_0 &\equiv& \frac{1}{2} G^{1/2} M^{-1} \left(\frac{M}{R}\right)^{3/2} = \frac{1}{2} c \left(\frac{GM}{c^2} \right)^{-1} \left(\frac{GM}{c^2 R} \right)^{3/2} \ . \end{eqnarray} Here, $R$ means the radius of the star on the major axis measured in the Schwarzschild-like coordinate. Our coordinate system in this paper is a kind of isotropic one and so $R$ is defined as follows: \begin{eqnarray} R &=& r_{\rm AB} \left( 1 + \frac{M}{2r_{\rm AB}}\right) ^ 2 \ , \\ r_{\rm AB} &=& \frac{r_{\rm B} - r_{\rm A}}{2} \ . \end{eqnarray} The meaning of these quantities, ${\hat j}$ and ${\hat \omega}$, can be roughly understood if we consider a system in Newtonian gravity which consists of two identical rigid spheres of uniform density in a contact phase. For such a system, $\hat{j} = 1$ and $\hat{\omega} = 1$. Another property of these quantities can be seen from the definition of the normalization factors, Eqs.~(\ref{normalization J}) and (\ref{normalization Omega}). In these expressions, the differences originating from the different strength of gravity are ``renormalized" by introducing the term related to the quantity $(M/R)_{\infty}$. Thus we will call $\hat{\omega}$ a renormalized angular velocity and $\hat{j}$ a renormalized angular momentum. In Figs.~\ref{J-Omega New}, the renormalized angular velocity $\hat{\omega}$ is plotted against the renormalized angular momentum $\hat{j}$ for several sequences of $N$ and $(M/R)_{\infty}$. As seen from this figure, the values of $\hat{\omega}$ and $\hat{j}$ for all evolutionary sequences with constant rest masses are scaled to values around unity. For the Newtonian sequences, the position of contact phases for smaller values of $N$ approaches $(1.0, 1.0)$ but never reaches that point because configurations are not rigid bodies and deformed from spheres by the tidal force from the companion star. Several characteristic features can be seen in this figure. First, if we compare the sequences with the same value of $(M/R)_{\infty}$, sequences with stiffer equations of state locate at the upper-right region. This can be explained as follows. If we choose models which have the same values of the gravitational mass and the angular velocity, the radii are the same so that the inertial moment is larger for the stiffer polytropes. It implies that the angular momentum is larger for stiffer equations of state. Concerning the value of the renormalized angular velocity at the contact stage, the gravitational force is stronger for the stiffer polytropes because of the distribution of the matter inside the star. Thus larger angular velocity is required for configurations with stiffer equations of state. Second, if we compare the sequences with the same value of $N$, more relativistic sequences locate at the larger values of $\hat{j}$. This is explained as follows. If we choose the models which have the same values of $N$, $M$ and $\hat{\omega}$, we obtain the following relation: \begin{equation} \hat{j} \propto \frac{< (r \sin \theta)^2 >}{R^2} \ , \end{equation} where $<F>$ is the mass average of the quantity $F$ defined as follows: \begin{equation} < F > \equiv \frac{\int F dm}{\int dm} \ . \end{equation} Here $dm$ is the mass element of the configuration. Since, in general, the change of the averaged quantity is smaller than the change of the quantity itself, the change of $\hat{j}$ is affected mainly by the change of $R$ which decreases as $(M/R)_{\infty}$ increases. Therefore the value of $\hat{j}$ increases and the curves are shifted towards right in the plane. It should be noted that if we consider sequences with the same value of $(M/R)_{\infty}$ but different values of $N$, differences due to different values of $N$ are amplified for configurations with the larger value of $(M/R)_{\infty}$. Since real neutron stars cannot be approximated by a single polytropic relation all through the whole star, the evolutionary sequences cannot be approximated by the assumptions adopted in this paper, i.e., the assumption that $N$ and $K$ are conserved. Therefore, in order to get information about real evolutions, we need to construct evolutionary sequences with realistic equations. \\ We would like to thank Dr. K\=oji Ury\=u for his helpful discussions. FU is a Research Fellow of the Japan Society for the Promotion of Science (JSPS) and is grateful to JSPS for the financial support. This work was partially supported by the Grant-in-Aid for Scientific Research (C) of JSPS (12640255).
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astro-ph0112571_arXiv.txt
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astro-ph0112217_arXiv.txt
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The RossiXTE mission provided us with an unprecedentedly large database of X-ray observations of transient black-hole candidates. These systems are crucial for the understanding of the physical properties of mass accretion onto black holes. Here I review the results on a selected sample of systems and describe their behavior in a purely phenomenological way. From these results, we can derive a better classification of the spectral and timing characteristics of black hole candidates in terms of basic states.
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Transient black hole candidates (TBHC) are the most important laboratories where we can study accretion onto black holes. The main reasons are two: first, there is a very limited number of bright persistent systems; second, transients go through a large range of mass accretion rate during their outbursts, therefore allowing us to study how the accretion properties change with accretion rate. In the case of persistent sources, they show rare state transitions (if any), and even when they do it is not clear what their dependence on accretion rate is (see e.g. \cite{zhan1997}). Before the launch of the Rossi X-ray Timing Explorer (RXTE), a relatively small number of systems was known (see \cite{tana1995,vdk1995} for a review). From these sources, a general ``canonical'' paradigm for the spectral/timing properties of TBHC had emerged (see \cite{tana1995,vdk1995,miya1993}). Although there were notable exceptions, this paradigm was a good starting point for theoretical modeling. Four separate states were identified, which will be described here according to the behavior in the 2-20 keV band: \begin{itemize} \item Low/Hard State (LS): the energy spectrum can be described by a single power law with a photon index $\Gamma\sim$1.6. In addition, sometimes a weak disk-blackbody (DBB) component with kT$<$1 keV is observed, contributing less than a few percent to the detected flux. The Power Density Spectrum (PDS) is characterized by a strong band-limited noise with a break frequency below 1 Hz and a large fractional variability (30-50\%). A low-frequency $<$1 Hz QPO is sometimes observed. \item Intermediate State (IS): the energy spectrum can be decomposed in two components: a power law with $\Gamma\sim$2.5 and a clearly detectable DBB with kT$\leq$1 keV. The PDS shows a band-limited noise with a break frequency higher than the LS (1-10 Hz) and a fractional variability of 5-20\%. Sometimes a 1-10 Hz QPO is observed. \item High/Soft State (HS): the energy spectrum is dominated by a DBB component with kT$\sim$1 KeV, with the power-law component either below detection or extremely weak and steep ($\Gamma\sim$ 2--3). Very weak noise is detected in the PDS, in the form of a power-law with a few \% of fractional variability. \item Very High State (VHS): the energy spectrum is a combination of a DBB (kT$\sim$1-2 keV) and a power law ($\Gamma\sim$ 2.5). The PDS can be of two types: either a band-limited noise similar to that of the IS, or a power-law, sometimes with a QPO. This state was observed only in two sources: the transient GS 1124-68 and the persistent source GX 339-4. \end{itemize} The dependence of these states and their transitions on increasing accretion rate was determined mostly by the only transient system that had shown all four of them, GS 1124-68 \cite{miya1994,ebis1994}: LS--IS--HS--VHS. As one can see from the description above, the IS is very similar to the VHS, both in energy and timing characteristics. What was believed to be different between them is the value of the accretion rate: the VHS was observed at very high accretion rates, while the IS was observed much later in the outburst of GS 1124-68, after a long period of HS, and therefore at lower accretion rate (see \cite{bell1997}). Despite the exceptions, a transient was expected to have a fast-rise/exponential-decay light curve (reflecting the time history of accretion rate), lasting a few months, possibly with one or more re-flares, undergo a number of state transitions in the sequence outlined above following accretion rate changes, and return to quiescence after a period of a few weeks to months, until the next outburst. Once again, the prototypical source would be GS 1124-68.
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Since the launch of RXTE, thanks to the presence of the ASM and the operational flexibility of the mission, the available X-ray data on black hole candidates, especially transient systems, as increased substantially. The picture that emerges from the analysis of this large database is at the same time simpler and more complex than the paradigm that existed before. It is more complex because the phenomenology became quite complicated. The variety of QPOs observed in XTE J1550-564 and GRS 1915+105, and the extreme structured variability of the latter are perhaps the best examples. However, it also became somewhat simpler, as the number of basic spectral/timing states is now reduced to three, despite the tremendous diversity of some timing features, and the presence of a second parameter governing state transitions (although its physical nature needs to be addressed) might yield a direct measurement of a fundamental parameter of these systems. Among all sources, GRS 1915+105 shows more variability and state-transitions than all other sources together. it might be the way to solving basic problems of accretion, but it might also turn out to be an endless complication that leads away from the solutions. However, it is important that theoretical models address the general picture described above in addition to trying to describe in extreme detail the spectral distribution of a particular state in a particular source, especially since any spectral fit to low-resolution data involves by definition a strong a priori bias. \small
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astro-ph0112021_arXiv.txt
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We are undertaking an extensive X-ray monitoring campaign of the two Crab-like pulsars in the Large Magellanic Clouds, PSR~B0540-69 and PSR J0537-6910. We present our current phase-connected timing analysis derived from a set of 50 pointed X-ray observations spanning several years. From our initial 1.2 yr monitoring program of the young 50 ms pulsar PSR~B0540-69, we find the first compelling evidence for a glitch in its rotation. This glitch is characterized by $\Delta \nu / \nu = (1.90 \pm 0.05) \times 10^{-9}$ and $\Delta \dot \nu / \dot \nu = (8.5 \pm 0.5) \times 10^{-5}$. Taking into account the glitch event, we derive a braking index of $n = 1.81 \pm 0.07$, significantly lower than previous reported. For the 16 ms pulsar, PSR~J0537-6910, we recorded 6 large glitch events during a period of nearly 3 years, the highest rate of all known Crab-like systems. Despite the extreme timing activity, the long term spin-down of this pulsar continues to average $-1.9743 \times 10^{-10}$ Hz/s.
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A characteristic signature of young rotation-powered pulsars is the phenomena of ``glitches'', sudden discontinuous changes in their spin periods (e.g., see Lyne \& Graham-Smith 1998). The physical causes of these glitches are not understood. Suggestions include sudden changes in the neutron star (NS) crust configuration (``starquakes''), abrupt reconfiguration of the magnetic field, or perhaps to the sudden unpinning of vortices in the superfluid neutrons in the inner part of the NS crust. For the latter, the amplitude of the glitch provides an estimate of the fractional part of the moment of inertia carried by superfluid neutrons (Lyne et al. 1996). The largest glitches have relative amplitudes ($\Delta\nu/\nu$) of several parts per million, but the range of amplitudes covers many orders of magnitude. Often there is a partial recovery back toward the pre-glitch rotation rate on a time scale of $\sim 100$ days, however, the spin-down rate may be permanently altered. Lyne (1995) noted that the amount of recovery in the rotation rate tends to be inversely proportional to the characteristic age of the pulsar. For some pulsars (e.g. Crab) the glitches are accompanied by a persistent increase in the spin-down rate with a relative amplitude of a few $\times 10^{-4}$. This increase may be caused, for example, by changes in the alignment of the magnetic field because of starquakes (Allen \& Horvath 1998; Link et al. 1998). In this paper we present preliminary results on the first detection of glitches from the two LMC Crab-like pulsars.
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astro-ph0112021_arXiv.txt
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astro-ph0112351_arXiv.txt
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{ We present the results of $B$, $V$, $R$ surface photometry of ESO~603-G21 -- a galaxy with a possible polar ring. The morphological and photometric features of this galaxy are discussed. The central round object of the galaxy is rather red and presents a nearly exponential surface brightness distribution. This central structure is surrounded by a blue warped ring or disk. The totality of the observed characteristics (optical and NIR colors, strong color gradients, HI and H$_2$ content, FIR luminosity and star-formation rate, rotation-curve shape, global mass-to-luminosity ratio, the agreement with the Tully-Fisher relation, etc.) shows that ESO~603-G21 is similar to late-type spiral galaxies. We suppose that morphological peculiarities and the possible existence of two large-scale kinematically-decoupled subsystems in ESO~603-G21 can be explained as being a result of dissipative merging of two spiral galaxies or as a consequence of a companion accretion onto a pre-existing spiral host.
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The past several years have been very rich in observational studies of galaxy formation and evolution. One of the most interesting conclusions of these studies is the continuous assembly of galaxies (see Ellis 2001 for a recent review). Among the best local examples of delayed galaxy formation are the so-called multi-spin galaxies -- objects with more than one axis of rotation (Rubin 1994). Polar-ring galaxies (PRG) are probably the best known instance of multi-spin objects (see Whitmore et al. 1990, hereafter PRC, for definition and catalog of such objects ). PRG probably represent products of merger or external accretion phenomena (PRC, Reshetnikov \& Sotnikova 1997). In this article we present the results of photometric observations of ESO~603-G21 -- a good PRG candidate according to Whitmore et al. (1990) (see Fig.~3f in Whitmore et al. 1990 and contour maps in our Fig.~2). This galaxy resembles an early-type galaxy with a well-developed bulge and an extended warped and edge-on disk/ring. A dust lane can be seen at the intersection of the bulge and the disk/ring. The spectroscopic data for this object indicate a complex scenario. The rotation curves for \mbox{ESO~603-G21} show that the gas and stars in the disk/ring revolve around the minor axis (PRC, Arnaboldi et al. 1994). At P.A.~=~24$^{\rm o}$ (minor axis), the spectra show no motion of the gas perpendicularly to the disk/ring. In contrast, the absorption line rotation curve indicates the existence of stellar motion along this axis (Arnaboldi et al. 1994). There stellar kinematics possibly indicate that the underlying stellar body is triaxial (Arnaboldi et al. \cite{a1}, Arnaboldi et al. \cite{a}).
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The global kinematical structure of ESO~603-G21 -- stellar rotation along two orthogonal position angles (Arnaboldi et al. \cite{a1}) -- suggests that the object is a polar-ring galaxy. The host galaxy is probably an early-type galaxy with an exponential-like surface brightness distribution. The central galaxy is surrounded by a warped star-forming ring or disk. In general, ESO~603-G21 looks similar to other classic PRG (e.g. NGC~4650A). There are, nonetheless, several facts complicating such an interpretation. First, the central round component shows very low surface brightness which may indicate that the central object is not an early-type galaxy like in most classic PRG (PRC). Second, in the near-infrared region ($K$ passband) most of the stellar light comes from a bright nearly-exponential disk. Third, the central round object, clearly visible in the optical images (Fig.~2), is quite faint in the $K$ passband (Arnaboldi et al. \cite{a}). The totality of the observed characteristics (optical and NIR colors, color gradients, HI and H$_2$ content, FIR luminosity and star-formation rate, rotation-curve shape -- Fig.~6 --, the agreement with the Tully-Fisher relation, etc.) suggests that ESO~603-G21 could be an unusual late-type spiral galaxy with a kinematically-decoupled extended "bulge". Therefore, it may be similar in some respects to NGC~4672 and NGC~4698, which are early-type disk galaxies with geometric and kinematical orthogonal decoupling between the bulge and disk (Bertola et al. \cite{b}, Sarzi et al. \cite{sa}), or to NGC~2748, which is a late-type spiral galaxy with possible ongoing accretion of a dwarf companion onto the central region of the galaxy (Hagen-Thorn et al. 1996). The bulge-like feature of ESO~603-G21 can be a "true" polar ring that is formed during almost perpendicular accretion of an early-type companion onto central region of a pre-existing disk galaxy. Another interesting interpretation of the observed ESO 603-G21 peculiarities is that the galaxy may be the result of a dissipative merger event (this scenario was proposed recently by Iodice et al. 2001 to explain the NGC~4650A puzzles). According to Bekki (1997, 1998), dissipative galaxy merging with a near polar orbit orientation can transform two late-type spirals into one PRG. In this scenario, a spiral galaxy falling from the polar axis of the target galaxy triggers the outwardly propogating density wave in the gaseous disk of the victim galaxy. Then, gaseous dissipation and star formation transform the victim disk into polar ring or disk. The central object is the intruding galaxy that has been turned into an early-type-like galaxy during the merging. Figure ~\ref{draw} depicts the various internal substructures of ESO 603-G21 as revealed in this work (see item 3.2). Such complex, non-settled, fine structure of the galaxy supports our supposition about relatively late formation of the ''bulge'' due to external accretion or a merger. \begin{figure} \psfig{file=fig7.ps,width=8.7cm,clip=} \caption[7]{Sketch of ESO 603-G21 as seen in careful visual inspection of the images in the various filters. North is up and East is on the left.} \label{draw} \end{figure} Interestingly enough, the companion of ESO 603-G21 is \mbox{ESO 603-G20}, an edge-on disk-galaxy without any ``explicit" evidence of interaction. The relative velocity between both objects is 65 kms$^{-1}$ (see NED and references therein); this suggests that both objects may form a bound system! There is, nonetheless, a third faint object between them, which seems to bear a very faint and narrow bridge to ESO~603-G21. The coordinates of the centroid (J2000) of this object are $\alpha$ = 22$^{\rm h}$ 51$^{\rm m}$ 10.4$^{\rm s}$ and $\delta$ = -20$^{\rm o}$ 14$'$ 59.5$\arcsec$ within 1$\arcsec$ of error, as calculated from the Digitized Sky Survey (DSS) image (see Fig.~8). Therefore, we have denoted this object Anon~J225110.4-201459.5. This is probably a low-surface brightness galaxy. On the basis of the DSS image we have found that the $B$-band total magnitude of the galaxy is $B_T=18.0\pm0.5$. \begin{figure} \psfig{file=fig8.ps,width=8.7cm,clip=} \caption[8]{In spite of the DSS resolution, a high-pass filtering has slightly enhanced the probable bridge between ESO 603-G21 (on the left) and Anon~J225110.4-201459.5 at the center (see text). ESO 603-G20 is on the right. North is to the top, east on the left.} \label{dss} \end{figure} It is essential to note that the disks of \mbox{ESO 603-G21} (Fig.1) and \mbox{ESO 603-G20} (Fig.~\ref{dss}) are strongly warped. This feature, as well as the probable bridge, may be an indication of ongoing interaction in the \mbox{ESO 603-G21}--\mbox{ESO 603-G20}--Anon~J225110.4-201459.5 triple system (e.g. Reshetnikov \& Combes 1999). So \mbox{ESO 603-G21} is not an isolated object, but a member of a group of galaxies (like, for instance, NGC~4650A). Such dense spatial environment supports the idea that \mbox{ESO 603-G21} may represent the result of a merging event. To test this scenario, detailed numerical simulations are needed.
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astro-ph0112398_arXiv.txt
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{ In inhomogeneous optically thick synchrotron sources a substantial part of the electron population at low energies can be hidden by self-absorption and overpowered by high energy electrons in optically thin emission. These invisible electrons produce Faraday rotation and conversion, leaving their footprints in the linear and circular polarized radiation of the source. An important factor is also the magnetic field structure, which can be characterized in most cases by a global magnetic field and a turbulent component. We present the basic radiative transfer coefficients for polarized synchrotron radiation and apply them to the standard jet model for relativistic radio jets. The model can successfully explain the unusual circular and linear polarization of the Galactic Centre radio source Sgr A* and its sibling M81*. It also can account for the circular polarization found in jets of more luminous quasars and X-ray binaries. The high ratio of circular to linear polarization requires the presence of a significant fraction of hidden matter and low-energy electrons in these jets. The stable handedness of circular polarization requires stable global magnetic field components with non-vanishing magnetic flux along the jet, while the low degree of total polarization implies also a significant turbulent field. The most favoured magnetic field configuration is that of a helix, while a purely toroidal field is unable to produce significant circular polarization. If connected to the magnetosphere of the black hole, the circular polarization and the jet direction determine the magnetic poles of the system which is stable over long periods of time. This may also have implications for possible magnetic field configurations in accretion flows.
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The detection of circular polarization (CP) in the continuum of radio sources is believed to be a powerful tool to test physical source models (Hodge \& Aller \cite{HA79}). But CP in extragalactic radio sources is extremely elusive (Roberts et al. \cite{RC75}; Ryle \& Brodie \cite{RB81}; Weiler \& de Pater \cite{WdP83}) and is detected in only a few sources. A more recent ATCA-survey (Rayner, Norris \& Sault \cite{RNS2000}) for CP in radio-loud Quasars, BL Lacs and Radio Galaxies with improved sensitivity of $0.01\%$, has shown a clear correlation of fractional CP with spectral index, in the sense that CP is stronger in flat and inverted spectrum sources. Circularly polarized radiation is therefore preferentially produced in self-absorbed radio cores. The fractional CP at 5 GHz is found to be between 0.05\% and 0.5\% in 11 out of 13 inverted spectrum sources at the ATCA spatial resolution of $2$ arcsec. At higher VLBA-resolution ($\sim 0.5$mas) Homan \& Wardle (\cite{HW99}) report localized CP of 0.3\%-1\% in the jet-cores of 3C273, PKS 0528+134, and 3C279, which in a few cases may be as high as the local linear polarization. It is also found, that intraday variable sources are circularly polarized (Macquart et al. \cite{MK00}), and that LP (linear polarization) and CP are both variable on timescales below 1 day. Recently CP was also found in X-ray binaries (Fender et al. \cite{Fender2000}\&\cite{Fender2002}). While the handedness of CP is remarkably stable over several years (Komesaroff et al. \cite{K84}; Homan \& Wardle \cite{HW99}; Fender et al. \cite{Fender2002}) for individual sources, no preferred handedness of CP in general is found. An even more challenging situation than observed in radio-loud extragalactic jet sources presents itself in the centre of our galaxy. The compact radio source Sgr A$^*$ (see Melia \& Falcke \cite{MF01}), believed to be coincident with the dynamical centre of the Milky Way with a mass of $2.6\,10^6 M_\odot$ (Eckart \& Genzel \cite{EG96}; Ghez et al. \cite{G98}) presumably in a single black hole, exhibits consistently larger circular than linear polarization in the range of $1.4$ to $15$ GHz (Bower et al. \cite{BFB99}; Sault \& Macquart \cite{SM99}) with CP between $0.2$\% and $1$\%. LP is small and below the detection limits (Bower et al. \cite{BB99}, \cite{BW99}) with the exception of sub-mm measurements, which possibly shows LP at a level of $10\%$ in the range $750-2000 \mu$m (Aitken et al. \cite{A2000}). The beam size of the sub-mm observations is $\sim 10$ arcsec. The flux is dominated by extended dust emission or free-free emission and the synchrotron source is comparably weak at these wavelength. The inverted radio spectrum of Sgr A$^*$ ($S_\nu \propto \nu^\alpha,\, \alpha \approx 0.3$) can be interpreted as either optically thin synchrotron emission (Beckert et al. \cite{Beckert96}) or self-absorbed synchrotron emission from a jet-like outflow (Falcke et al. \cite{FMB93}; Falcke \& Markoff \cite{FM00}). The idea of synchrotron emission by thermal electrons from Sgr A$^*$ was briefly considered by Reynolds \& McKee (\cite{RMcK80}) and revived for mildly relativistic electrons in the self-absorbed ADAF models for accretion in the galactic centre (Narayan et al. \cite{N98}). The first ADAF models under-predicted the radio flux between $1$--$100$\,GHz, which can be attributed to an outflow or jet. The upper limits for Sgr A$^*$ in the infrared require a sharp high energy cut-off for the electron distribution below $\gamma_{\mathrm{max}}$ of a few\,$\times 10^2$. Therefore thermal or quasi-monoenergetic electrons are responsible for the radio emission (Beckert \& Duschl \cite{BD97}), which distinguishes Sgr A$^*$ from high-luminosity, radio-loud AGNs. A close relative of Sgr A$^*$ is found in the centre of the normal spiral M81. The radio source M81$^*$ exhibits an elongated jet-like structure (Bietenholz et al. \cite{BBR00}), has a similiar radio spectrum (Reuter \& Lesch, \cite{RL96}), a slightly larger luminosity, still below the AGN level, and has recently be found to be circularly polarized (Brunthaler et al. \cite{Brun01}) without detectable LP. The fractional variability of CP is usually stronger than of LP, which in turn is stronger than for the total intensity. Together with the preserved handedness this poses servere constrains on possible scenarios for CP production and its variability (Komesaroff et al. \cite{K84}). The suggested mechanisms are ({\em a}) intrinsic cyclo-synchrotron emission from low-energy electrons or from electrons with small pitch angles seen close to the magnetic field direction (Legg \& Westfold \cite{LW68}), conversion from LP to CP as a propagation effect induced by ({\em b}) low energy electrons inside the relativistic plasma (Hodge \& Aller \cite{HA77}) or ({\em c}) by a magnetized cold plasma surrounding the synchrotron source. This requires either Faraday rotation (not possible in pure electron/positron jets) or changing (e.g., turbulent) B-field directions along the line of sight in the source. A further possiblity for CP production are ({\em d}) inhomogeneous rotation measures in intervening cold plasma either close to the source or in our galaxy (Macquart \& Melrose \cite{MM2000}). The existence of these plasma screens can be infered from interstellar scattering believed to be the cause for intraday variability in some sources (Rickett et al. \cite{R95}; Dennett-Thorpe \& de Bruyn \cite{DB00}; Macquart et al. \cite{MK00}; Beckert et al. \cite{BCore02}). This model predicts variable CP with a time averaged mean of $<$CP$> = 0$. In this paper we consider propagation effects like Faraday rotation and cyclic conversion of LP to CP and back (Pacholczyk \cite{P73}) in turbulent, self-absorbed jets or outflows. First results were already published in Falcke et al. (\cite{FB02}). We rederive some of the basic radiation transfer coefficients which, for example, could also be used for anisotropic particle distributions. The application of conversion to compact radio jets has been explored perviously by Jones (\cite{J88}) using different techniques and without focusing on sources with large circular polarization and the role of globally ordered magnetic fields. Here we investigate the standard jet model with respect to the new polarization data placing some emphasis on the role of turbulence, the ratio of low- to high-energy particles, and the magnetic field confirguation. The paper is organized as follows: In Sec.\,\ref{emitrans} we review the basic production channels for CP. The outfow/jet model and the possible turbulence in the $B$-field is presented in Sec.\,\ref{outflowmod}. The consequences of Faraday rotation and conversion are discussed in Sec.\,\ref{dpolcon} followed by a detailed model of Sgr A$^*$. Polarization variability is the topic of Sec.\,\ref{polvar} and we close with a discussion of our results in Sec.\,\ref{discus}.
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\label{discus} Recent observations of radio circular polarization in AGN, X-ray binaries, and the Galactic Centre black hole suggest that CP at the 0.3\%-1\%-level is common to many self-absorbed synchrotron sources. Faraday rotation and conversion in a magnetized and therefore bi-refringent plasma produce enhanced circular and reduced linear polarization. Both processes are sensitive to the presence of low-energy electrons and the orientation of the global magnetic field. The standard jet model for compact radio cores with a helical plus a turbulent magnetic field can well reproduce the circular and linear polarization spectrum of sources like Sgr A* and M81* with their high CP-to-LP ratio. The suppression of LP is achieved by the presence of a significant number of low-energy electrons in the source and an absence of an optically thin power-law extending to higher energies. The same model can also explain the typical level of circular polarization in blazars and the CP-to-LP ratio observed in blazars and X-ray binary jets. In this case a number of low-energy electrons is reduced with respect to the Sgr A* model and a power-law in the electron distribution exists. For Sgr A* the number of low-energy electrons producing conversion and depolarization needs to be significantly higher (by 2-3 orders of magnitude) than the number of radiating hot electrons. This means that a large fraction of the outflowing jet material is in the form of hidden matter shielded by self-absorption. This increases the estimates of the total jet power, which can be 5 orders of magnitude higher than the synchrotron luminosity. If one presumes that this power has to be provided by an accretion flow, the minimum accretion rates of $10^{-9..-8} M_{\odot}$/yr, previously estimated from ``maximally-efficient'' jet models for Sgr A* (Falcke et al.~\cite{FMB93}; Falcke \& Biermann \cite{FB99}) need to be raised to about $10^{-6} M_{\odot}/$yr. This is quite consistent with recent estimates of Bondi-Hoyle accretion rates onto Sgr A* (Baganoff et al. \cite{Baganoff2002}) and with suggestions for a coupled jet plus ADAF model (Yuan et al. \cite{Yuan2002}), where the emission from the accretion disk is highly suppressed with respect to the jet. It is also interesting to note that to fit the CP with conversion one requires an asymmetry in the magnetic field components. This is naturally achieved by a helical magnetic field as is presumed to exist in jets. A symmetric configuration, e.g. a tightly wound helix or even a toroidal magnetic field structure would have difficulties to produce the observed level of CP. The stable handedness of CP over 20 years also implies a long-term stable component of the unidirectional field along the line-of-sight. This indicates that the polarity of the magnetic field (the ``magnetic north pole'') has remained constant over the last two decades. In view of the rather short accretion time scale in Sgr A$^*$ one could also speculate that this polarity is related to the accretion of a stable large-scale magnetic field which is accreted and expelled via the jet. The same can be said about blazars and X-ray binaries, where the stability found in GRS1915+105 by Fender et al. (\cite{Fender2002}) is particularly interesting since the intrinsic accretion time scales in X-ray binaries are much shorter than those in supermassive black holes. Another important aspect of CP measurements is the question of the matter content of jets. We find that the constraints from CP of individual jet components for the jet power in blazars are not quite as severe as previously claimed and a statement in support of a pure electron/positron jet has to viewed with caution. For Sgr A* or M81* the situation may be different. If the depolarization is indeed intrinsic to the jet and not a surrounding medium (Agol \cite{Agol2000}, Quataert \& Gruzinov \cite{Quataert2000}), one needs a high Faraday optical depth in the source, which can only be achieved by an excess of ``warm'' ($1\la\gamma\la100$) electrons in an electron/proton plasma. While we have here assumed that all electrons are distributed in a single power-law, the actual situation may be quite different. For Sgr A* a power-law is actually not needed and we could obtain rather similar results with a two-temperature electron distribution, with temperatures corresponding to $\gamma_{\rm min}$ and $\gamma_{\rm max}$ respectively. This is not quite possible in blazars or bright X-ray binary jets, where extended electron power-laws are directly observed in the optically thin regime. It could well be that the radiative inefficiency of Sgr A* is due to the lack of effective shock acceleration that would increase the number of high-energy electrons with respect to the number of low-energy electrons (and in turn decrease the CP-to-LP ratio). The origin of these different electron distributions and their role for the radio-loudness of jet sources should be a very exciting question for further research. By improving our sensitivity and imaging all four Stokes parameters at multiple frequencies in the future, it will be possible to construct models of the entire emission and transfer processes in the source and determine the composition and energy spectrum of the relativistic plasma within jets. \appendix
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astro-ph0112398_arXiv.txt
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astro-ph0112167_arXiv.txt
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Observations using the Australia Telescope Compact Array at a wavelength of 6~cm have uncovered the radio counterpart to the compact X-ray nebula surrounding the Vela pulsar. Two lobes were found oriented about the spin axis of the pulsar, starting at the edge of X-ray emission, they extend to three times the size. The northern lobe has a bright, defined edge and an integrated flux of 0.14 Jy, while the southern lobe of 0.12 Jy is more diffuse.
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High resolution images from Chandra X-ray observations of the Vela pulsar have shown (Helfand, Gotthelf, \& Halpern 2001; Pavlov et al. 2001) an X-ray pulsar wind nebula (PWN), which has been modelled in detail by Helfand et al. (2001) with an alternative interpretation by Radhakrishnan \& Deshpande (2001). X-ray and synchrotron emissions are closely linked as highly charged particle flows drive both, so radio emission is a natural comparison. This targeted observation was optimised at 6~cm to match the Chandra image resolution and to improve on previous radio knowledge. We imaged a compact radio counterpart about the pulsar (figure 1). Previous radio studies (Bietenholz, Frail, \& Hankins 1991; Frail et al. 1997a; Bock et al. 1998a; Bock, Turtle, \& Green 1998b) have focussed on larger scales of filaments or wisps and the apparent connection to the Vela-X region of the supernova remnant. Reprocessed ATCA archive data showed diffuse extended emission in the region of the lobes and indicated a lack of emission in the X-ray PWN region, but it was not clear if this was due to sensitivity limits or a signature of radio emission.
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astro-ph0112241_arXiv.txt
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2MASSI J1315309$-$264951 is % an L3 dwarf with strong H$\alpha$ emission discovered in the course of a color-selected survey for active galactic nuclei using % the Two-Micron All-Sky Survey (2MASS). The strength of its % H$\alpha$ emission % decreased by about a factor of two between two epochs separated by 137 days. This is the first time that variable \ha\ emission has been reported in an L dwarf, and is probably the first observation % of an \ha\ flare in an L dwarf. The value of $\log(L_{\rm H\alpha}/L_{\rm bol}) > -4.17$ observed at the discovery epoch is larger than that of any other L dwarf but comparable to that of % 2MASSI J1237392+652615, the only reported T dwarf with \ha\ emission. The observed variability indicates that the \ha\ emission of 2MASSI J1315309$-$264951 is powered either by magnetic fields or by accretion in a binary system. Spectroscopic or narrow-band \ha\ monitoring of L and T dwarfs on timescales of hours to days would be the most useful step toward a better understanding of their \ha\ emission mechanism(s). %
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\label{INTRO} For many years the coolest, lowest-mass stars known were M dwarfs, but in recent years the L and T dwarfs have extended the stellar sequence to even lower temperatures and masses \markcite{kir99}({Kirkpatrick} {et~al.} 1999). These three classes of low-mass stars and brown dwarfs exhibit distinctly different spectral features due to the decrease in photospheric temperature from M through L to T. In addition, M dwarfs often show H$\alpha$ in emission. Down to the early M dwarfs, \ha\ activity correlates with rotation and thus decreases with age as stars lose angular momentum over time via stellar winds. Beyond objects of spectral type M8, however, it appears that \ha\ activity is stronger in more massive objects, even if they are older \markcite{giz00}({Gizis} {et~al.} 2000). The frequency of \ha\ emission peaks around type M7 and declines for later-type L and T dwarfs, reaching zero at L5 \markcite{giz00}({Gizis} {et~al.} 2000). However, contrary to this trend, \markcite{bur00}{Burgasser} {et~al.} (2000) % reported the discovery of a T dwarf with strong H$\alpha$ emission. Here I report the discovery of an L3 dwarf with similarly strong H$\alpha$ emission.
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\label{CONCL} 2MASSI\,J1315309$-$264951 is an L3 dwarf with strong H$\alpha$ emission which decreased in strength by about a factor of two between two epochs separated by 137 days, the first reported variable \ha\ emission in an L dwarf. The \ha\ emission in 2MASSI J1315309$-$264951 must be powered either by magnetic activity or by accretion in a binary system. Accreting binaries are rare, so that hypothesis is unlikely. The spectra presented here do not rule out a slow variation in \ha\ strength, but slow variations of the observed amplitude are rare among M dwarfs. Since flaring powered by reconnection of magnetic fields is common in M dwarfs and a radio flare has been detected in the L3.5 dwarf 2MASSW\,J0036159+182110 \markcite{ber02}({Berger} 2002), a flare is the logical explanation for the \ha\ variability in 2MASSI\,J1315309$-$264951, The value of $\log(L_{\rm H\alpha}/L_{\rm bol}) > -4.17$ observed at the discovery epoch is larger than that of any other L dwarf but comparable to the value of $-4.3$ observed for 2MASSI J1237392+652615, the only reported T dwarf with \ha\ emission. However, both these values lie well below the average $\log(L_{\rm H\alpha}/L_{\rm bol})=-3.8$ observed in M dwarfs. Only two L dwarfs and one T dwarf are known to exhibit \ha\ emission of strength $\log(L_{\rm H\alpha}/L_{\rm bol}) > -5$. Thus perhaps two percent of L or T dwarfs exhibit \ha\ emission this strong at any given time. Given the small number statistics, this is consistent with the duty cycles observed for \ha\ flares ($\sim$7\%; \markcite{giz00}{Gizis} {et~al.} 2000) and radio flares (2-10\%; \markcite{ber02}{Berger} 2002) among late-type dwarfs. Spectroscopic or narrow-band \ha\ monitoring of L and T dwarfs on timescales of hours to days is needed to determine if the frequency of strong \ha\ emission is governed by flaring, by the upper envelope of the magnetic field strength distribution, or by accretion in binary systems.
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astro-ph0112241_arXiv.txt
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astro-ph0112288_arXiv.txt
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We measured the bar pattern speed, $\om$, of the SB0 galaxy NGC 1023 using the Tremaine-Weinberg (1984) method with stellar-absorption slit spectroscopy. The morphology and kinematics of the \hi\ gas outside NGC 1023 suggest it suffered a tidal interaction, sometime in the past, with one of its dwarf companions. At present, however, the optical disc is relaxed. If the disc had been stabilized by a massive dark matter halo and formed its bar in the interaction, then the bar would have to be slow. We found $\om = 5.0 \pm 1.8$ \kmsa, so that the bar ends near its co-rotation radius. It is therefore rotating rapidly and must have a maximum disc.
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\label{sec:introduction} Strong bars are seen in optical images of roughly 30\% of all high surface brightness (HSB) disk galaxies (Sellwood \& Wilkinson 1993) and this fraction rises to 50\%-75\% in the near IR (Knapen 1999; Knapen \etal\ 2000; Eskridge \etal\ 2000). Understanding the structure and dynamics of barred (SB) galaxies is, therefore, an issue of some importance. The principal dynamical quantity for SB galaxies is the rotation frequency/pattern speed of the bar, $\om$. This is usually parametrized by the distance-independent ratio $\vpd \equiv \lag/\len$, where $\len$ is the semi-major axis of the bar and $\lag$ is the radius to the Lagrangian point, where the gravitational and centrifugal forces cancel out in the bar's rest frame. (The Lagrangian radius is therefore the generalization to strong bars of the corotation radius.) Contopoulos (1980) argued that a self-consistent bar is possible only if $\vpd \geq 1$. A bar is termed fast when $1.0 \leq \vpd \ltsim 1.4$, while, for a larger value of $\vpd$, a bar is said to be slow. A variety of methods have been used to attempt measurent of bar pattern speeds (see, for example, the review of Elmegreen 1996). Most measurements of $\vpd$ rely on hydrodynamical simulations. These usually try to match the gas flow in the region of the bar, particularly at the shocks, which works because the location of the shocks depends on $\vpd$, moving further ahead of the bar as $\vpd$ increases. A bar needs to be fast for the shocks to remain curved with their concave sides towards the bar major axis, as is usually observed (van Albada \& Sanders 1982; Athanassoula 1992). Detailed simulations of gas flows in individual galaxies also result in fast bars; examples include: NGC 1365 ($\vpd = 1.3$, Lindblad \etal\ 1996), NGC 1300 ($\vpd = 1.3$, Lindblad \& Kristen 1996), and NGC 4123 ($\vpd = 1.2$, Weiner \etal\ 2001). Hydrodynamical simulations can also recover $\vpd$ by matching morphological features in \hi; some examples are: NGC 7479 ($\vpd =1.22$, Laine 1996), NGC 1073 ($\vpd =1 - 1.2$, England \etal\ 1990), NGC 3992 ($\vpd =1$, Hunter \etal\ 1989), and NGC 5850 $(\vpd =1.35$, Aguerri \etal\ 2001). A direct method for measuring $\om$, using a tracer population which satisfies continuity, was developed by Tremaine \& Weinberg (1984). Since gas is subject to phase changes, it is not well-suited for this application. Old stellar populations in the absence of significant obscuration, on the other hand, are ideal for the Tremaine-Weinberg (TW) method. This has permitted application of the method to a small number of early type SB galaxies: NGC 936 ($\vpd = 1.4 \pm 0.3$, Kent 1987 and Merrifield \& Kuijken 1995), NGC 4596 ($\vpd = 1.15^{+0.38}_{-0.23}$, Gerssen \etal \ 1998) and NGC 7079 ($\vpd = 0.9 \pm 0.15$, Debattista \& Williams 2001). The observational evidence, therefore, favors fast bars. The perturbation theory calculations of Weinberg (1985) predicted that a fast bar would be slowed down rapidly in the presence of a massive dark matter (DM) halo. Such slow-down has been seen in various simulations (Sellwood 1980; Little \& Carlberg 1991; Hernquist \& Weinberg 1992; Athanassoula 1996). The fully self-consistent, high resolution $N$-body simulations of Debattista \& Sellwood (1998) also confirmed this prediction; however they showed that, for a maximum disc (here taken to mean a disc which dominates the rotation curve throughout the inner few disc scale-lengths, \cf\ van Albada \& Sancisi 1986), a fast bar can survive for a large fraction of a Hubble time. Subsequently, Tremaine \& Ostriker (1999) suggested that bars manage to remain fast not because discs are maximal but rather because the inner parts of DM halos are flattened and rapidly rotating. However, Debattista \& Sellwood (2000) showed that rapid slow-down persists even then unless the halo angular momentum is very large relative to that of the disc. Thus they concluded that SB galaxies must be maximal, and argued that the same must be true for all high surface brightness disc galaxies. This conclusion rests on a small number of pattern speed measurements; in view of the fact that maximum discs are in conflict with the predictions of cold dark matter (CDM) cosmologies (\eg\ Navarro \etal\ 1997), enlarging the sample of measured pattern speeds is desireable. In this paper, we report observations of NGC 1023, for which we applied the TW method. The rest of this paper is organized as follows. The TW method is described briefly in \S\ref{sec:twmethod}. Then, in \S\ref{sec:properties} we give an overview of the previously known properties of NGC 1023. The photometric observations, reduction and results, including $\len$, are presented in \S\ref{sec:photometry}, while \S\ref{sec:spectroscopy} presents the spectroscopic observations and results. We derive the rotation curve, corrected for the asymmetric drift, from which we obtain $\lag$. With these results at hand, we then apply the TW method in \S\ref{sec:pattern_speed}. We present our conclusions in \S\ref{sec:conclusions}.
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\label{sec:conclusions} We have found that the bar in NGC 1023 is fast, as are all bars which have been measured to date. Debattista \& Sellwood (1998, 2000) showed that fast bars can persist only if the disc is maximal. Following Ostriker \& Peebles (1973), it is sometimes thought that the unbarred (SA) galaxies are stabilized by massive DM halos. However, massive DM halos are not necessary for stabilizing discs; a rapidly rising rotation curve in the inner disc, such as when a massive bulge is present, is also able to inhibit bar formation (Toomre 1981; Sellwood \& Evans 2001). Debattista \& Sellwood (1998) argued that unbarred HSB galaxies must also be maximal for, if HSB disc galaxies form a continuum of DM halo masses spanning massive DM halo-stabilised SA galaxies to maximal SB galaxies, then slow bars must also be found in the intermediate range of halo masses. If we seek to avoid intermediate halo masses and slow bars by postulating (for whatever reason) a bimodal DM halo mass distribution for HSB galaxies, then we are left with the possibility that tidal interactions can still form bars, which would be slow (Noguchi 1987; Salo 1991; Miwa \& Noguchi 1998). Thus Debattista \& Sellwood (2000) concluded that the absence of slow bars requires that all HSB disc galaxies are maximal. However, it is possible that no such slow bars have been found because of an observational bias against SB systems with evidence of tidal interactions. We have chosen to study NGC 1023, in part, because it shows signs of a weak interaction in its past, without being at present significantly perturbed. The fast bar we found indicates that NGC 1023 has a maximal disc. If SA galaxies are stabilized by massive halos, we should find slow bars in that fraction of SB galaxies in which the bar formed through the interaction. While it is not possible to reach a general conclusion on the DM content of SA galaxies based on our measurement for a single galaxy, a large enough sample of similar SB galaxies with mild interactions in the past will be able to address this question. \bigskip \noindent {\bf Acknowledgements.} \noindent V.P.D. and J.A.L.A. acknowledge support by the Schweizerischer Nationalfonds through grant 20-56888.99. V.P.D. wishes to thank the Dipartimento di Astronomia dell'Universit\`a di Padova for hospitality while preparing for the observations. E.M.C. acknowledges the Astronomisches Institut der Universit\"at Basel for the hospitality while this paper was in progress. We wish to thank the staff of the JKT telescope, particularly to the support astronomer J. C. Vega-Beltr\'an and the staff of the TNG. We are indebted to R. Bender and R. Saglia for providing us with the FCQ package which we used for measuring the stellar kinematics. We thank the anonymous referee for suggestions that helped improve this paper. This research has made use of the Lyon-Meudon Extragalactic Database (LEDA) and of the NASA/IPAC Extragalactic Database (NED). \bigskip \noindent
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astro-ph0112077_arXiv.txt
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In this paper I argue that, far from necessarily hindering bar formation in disc galaxies, inner haloes may stimulate it. This constitutes a new instability mechanism by which bars can grow. To show this I use a number of $N$-body simulations whose initial conditions have identical discs and more or less concentrated haloes. They show that the bar that grows in the more halo-dominated environment is considerably stronger than the bar that grows in the more disc-dominated environment. This result is obtained from simulations with live haloes, i.e. composed of particles which respond to the disc and take part in the evolution. On the other hand, if the halo is rigid, it hinders or quenches bar formation, as expected. Comparison of two simulations which are identical in everything, except that the halo is live in the first one and rigid in the second one, leads me to suggest that the halo response can help the bar grow. Following the orbits of the stars in the halo, I find that a considerable fraction of the halo particles are in resonance with the bar. The halo may thus take angular momentum from the bar and stimulate its growth. I finally discuss whether and how the results of the $N$-body simulations can be applied to real galaxies.
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It is by now well established that galactic discs can be bar unstable (e.g. Miller, Prendergast \& Quirk 1970; Hohl 1971). In the quest for stability three main stabilising mechanisms have been proposed (see e.g. reviews by Athanassoula 1984, Sellwood \& Wilkinson 1993 and references in either) : i) The disc could be immersed in a massive spheroid, e.g. a bulge and/or an inner halo (Ostriker \& Peebles 1973). ii) The disc could be hot or have a hot center (Athanassoula 1983, Athanassoula \& Sellwood 1986). iii) The galaxy could be sufficiently centrally concentrated to stop the initially linear wave from reaching the center (Toomre 1981; Sellwood \& Evans 2001). Each one of these mechanisms has generated a lot of discussion, both regarding its efficiency and the way it can be applied to real galaxies. Here I will address the first of them, which was also historically the first to be introduced, and show that, contrary to what has been argued so far, inner haloes can, at least in some cases, enhance the bar. For this I will use a series of numerical simulations of disc-halo systems, described by Athanassoula \& Misiriotis (2002, hereafter AM02) and Athanassoula (2002b in preparation, hereafter A02b). I describe the simulations and their results in section 2, and in section 3 I discuss the role of the halo. Finally in section 4 I give a general discussion and address the applicability of my results to real galaxies.
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\label{sec:discuss} In the above I have compared three simulations starting off with identical discs, but different halo components. Any differences in their dynamical evolution should thus be attributed to the haloes. The strongest bar forms in the most halo dominated case, provided this is live, followed by the one in the disc-dominated case. In the simulation with the rigid halo there is only a very weak bar, or mild oval distortion, in the inner part. I thus reach the interesting conclusion that haloes can, at least in some cases, stimulate the bar instability and lead to stronger bars. This can be understood by a frequency analysis of the halo orbits, which reveals a large number of resonant orbits. Since these can exchange energy and angular momentum with stars at other resonances (Lynden-Bell \& Kalnajs 1972) they can stimulate the bar instability, contrary to previous beliefs. The evolution of the galaxy leads to considerable concentration of the disc material in the central areas. Thus model MD starts off as disc dominated in the central parts, and, with time, the disc further enhances its superiority. Model MH starts off quite differently. Initially the halo is slightly more important than the disc within the radius at which the disc rotation curve is maximum, and considerably more so at larger radii, as witnessed from its circular velocity curve, shown in Fig.~\ref{fig:inrotcur}. The central concentration of the disc increases considerably with time, so that, after the bar has grown, the disc dominates in the inner region. This may contribute an additional argument to the long standing debate of whether galactic discs are maximum or sub-maximum (e.g. Athanassoula, Bosma \& Papaioannou 1987; Bosma 1999, 2000; Bottema 1993; Courteau \& Rix 1999; Kranz, Slyz \& Rix 2001; Sellwood 1999, Weiner, Sellwood \& Williams 2001). Sackett (1997) and Bosma (2000) give a simple working definition to distinguish between maximum and sub-maximum discs, based on the value of $\gamma = V_{d,max} / V_{tot}$, where $V_{d,max}$ is the circular velocity due to the disc component and $V_{tot}$ is the total velocity, both calculated at a radius equal to 2.2 disc scale lengths. According to Sackett (1997) this ratio has to be at least 0.75 for the disc to be considered maximum or maximal. In the simulations it is not easy to define a disc scale length after the bar has formed, so I will calculate $\gamma$ at the radius at which the disc rotation curve is maximum, which is a well defined radius and is roughly equal to 2.2 disc scale lengths in the case of an exponential disc. Model MD starts off with $\gamma > 0.75$, so that the disc starts maximum and stays so all through the simulation. In fact the value of $V_{d,max} / V_{tot}$ increases somewhat with time. Model MH has initially a value of $\gamma$ around 0.68, i.e. close to the value of 0.63 advocated by Bottema (1993), and is therefore initially sub-maximum. This value, however, increases abruptly after the bar has formed, so that the disc can be considered maximum well before the time shown in Fig.~\ref{fig:basic}, with a value of $\gamma$ roughly equal to 0.86. Thus the formation of the bar leads the disc to evolve from sub-maximum to maximum, and hence strongly argues for maximum discs in disc galaxies with strong bars. This means that if we observe a strong bar in a disc galaxy the above simulations argue strongly and quantitatively that the underlying disc is maximum. The existence of gas should not alter this result. Indeed if the gas leads to a density distribution with a weak bar or no bar, then the above argument will be irrelevant, since it applies only to galaxies with strong observed bars. On the other hand, if the resulting bar is strong, then it should have rearranged the disc material sufficiently for the above argument to hold. We can reach similar results about the disc-to-halo mass ratio if we use the criterion of Athanassoula, Bosma \& Papaioannou (1987), who examined what spiral perturbations can grow in a given disc/halo decomposition of an observed rotation curve. In a similar way I can calculate the $m$ component that will be strongest amplified via the swing amplification mechanism (Toomre 1981) in my simulations at or around the radius at which the disc rotation curve reaches its maximum. I find that, for model MD, it is the $m$ = 2 all through the simulation, as expected. The initial disc for model MH is certainly not maximum. I find that at $t$ = 0 higher $m$ components will be the most strongly amplified. The evolution, however, changes this, so that after the bar has grown it is the $m$ = 2 component that is the strongest amplified at or around the radius at which the disc rotation curve reaches its maximum. Both MD and MH models thus have, after the bar has grown, a disc which is intermediate between the ``no $m$ = 1'' and ``no $m$ = 2'' limits advocated by Athanassoula, Bosma \& Papaioannou (1987) for real galaxies. Indeed these authors made a link between the structure present in a disc at a given time and the underlying halo mass at that time (not the initial halo mass) and thus their results are in good agreement with the above simulations, and many other similar ones (e.g. AM02 and A02b) . Apart from the halo, several other parameters can influence the formation of a bar. In particular let me stress the importance of the velocity dispersion of the disc particles, the effect of gas, and the effect of the velocity and mass distribution in the halo, as well as the existence of a gaseous companion. A complete description of all these effects is well beyond the scope of the present contribution and will be presented elsewhere. Let me just add a few preliminary words about the effect of the disc velocity dispersion. A sequence of MH-type galaxies shows that for larger initial velocity dispersion of the disc particles the bar is less strong and for sufficiently high values, becomes oval, or quasi-circular, in good agreement with what was found for 2D models (Athanassoula 1983). A sequence of MD-type models is more complicated. In these models the bar grows faster and becomes very long and strong. At that time, however, a strong buckling instability develops which leads to a considerable decrease of the bar amplitude. The final amplitude of the bar is a result of the competition between these two effects and this may be close. Only at sufficiently large values of the velocity dispersion can we be sure that the resulting oval will be very thick, as in the MH sequence. All the above are rather preliminary and will be discussed at length elsewhere. Finally the mass and velocity distribution of the halo component, together with the bar pattern speed and its time dependence, should influence how each of the resonance regions is populated and how responsive it is, and therefore influence its ability to exchange energy and angular momentum with the bar. Since very little is known on the composition of the halo, let alone about the distribution of the matter in it, it is very difficult to pursue this issue further. Nevertheless the arguments in section~3 lead to the prediction that at least some of the stars of the visible halo should be in resonance, in as much as they trace the relatively inner parts of the halo. Testing this would necessitate accurate information on the six phase space coordinates of a sufficiently large number of halo stars, as well as a sufficiently accurate description of the halo potential and the bar pattern speed. Our own Galaxy, which is barred, is the only place where advances with future astrometric satellites may make this possible, if we concentrate in areas which could have a high fraction of resonant stars.
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astro-ph0112077_arXiv.txt
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0112
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astro-ph0112307_arXiv.txt
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We discuss the CM diagram of the galactic cluster NGC2420 to the light of current theoretical predictions. By relying on the most recent updating of the physical input, one finds too luminous theoretical He burning stars together with the evidence for a misfitting of the lower portion of the MS. Moreover one finds two well known overshooting signatures, as given by i) the large extension of the ``hook" preceding the overall contraction gap, and ii) the scarcity of stars just at the end of the gap. We show that the overluminosity of He burning stars appears as a constant prediction of models based on updated physics, whereas alternative assumptions about the Equation of State can account for the MS fitting. Moreover, due to the scarse statistical significance of the observational sample, one finds that overshooting signatures can be present also in canonical (without overshooting) predictions. We conclude that, unfortunately, NGC2420 does not keep the promise to be of help in costraining the actual dimensions of convective cores in H burning MS stars, suggesting in the meantime that using clumping He burning stars as theoretical standard candle is at least a risky procedure. In this context the need for firmer constraints about the reddening of galactic clusters is shortly discussed.
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The beautiful CM diagram presented in 1990 by Anthony-Twarog et al. for the intermediate age open cluster NGC2420 has been in the last ten years a favourite target for all the people concerned with the evolution of low to intermediate mass stars. The occasion for revisiting this cluster has been given to us by the recent paper by Pols et al. (1998, hereinafter P98), who presented new evolutionary tracks carefully discussing the fit of a selected (and well chosen) sample of galactic clusters. As discussed in that paper, NGC2420 seems to give a good chance to put firm constraints on the efficiency of the mechanism of core overshooting. Owing to the relevance of such an issue, we decided to go deep into the matter, hoping eventually to settle down such a long debated argument. However, as we will discuss in the following, the situation is far from being assessed.
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According to the previous sections, one finds that three out of the four theoretical misfitting of the cluster CM diagram can be accounted for within current evolutiony scenarios. On the contrary, no assumption appears able to reconcile the predicted luminosity of He burning stars with observation. To explore all the possibilities one may guess that a given amount of mass loss could account for this discrepancy. One generally assumes that mass loss occurs in the advanced phase of H shell burning, so that the internal structure of the He burning star is not affected by such an occurrence, which only decreases the amount of envelope surrounding the central He core. Under this assumption, the effect of mass loss on He burning models can be easily computed by simply decreasing the envelope of the constant-mass model. Numerical simulations shows that to reach the agreement between theory and observation for the clump luminosity one needs to decrease the He burning mass from the standard value of $\approx$ 1.5 M$_{\odot}$ to 1.0 M$_{\odot}$. This very high amount of mass loss seems very unlikely to us. At present time the discrepancy between theory and observation in the luminosity of He burning stars with degenerate RGB progenitors it is not a surprising result; it appears as a constant prediction of models based on updated physics (see e.g. P98, Castellani et al. 2000), whereas similar suggestions for the overluminosity of theoretical models have been also derived from the pulsational properties of RR Lyrae (see e.g. Caputo et al. 2000). We conclude that using clumping He burning stars as theoretical standard candles, as we did (Castellani et al. 1999), is at least a risky procedure. We note that in the above quoted work, one could recognize a signature of the overluminosity in the need of assuming rather large reddenings for all the clusters. In this context, it follows that firm constraints about the reddening of galactic clusters will help in solving this problem.
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astro-ph0112307_arXiv.txt
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0112
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astro-ph0112131_arXiv.txt
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We present a robust method to derive the duty cycle of QSO activity based on the empirical QSO luminosity function and on the present-day linear relation between the masses of supermassive black holes and those of their spheroidal host stellar systems. It is found that the duty cycle is substantially less than unity, with characteristic values in the range $3-6\times 10^{-3}$. Finally, we tested the expectation that the QSO luminosity evolution and the star formation history should be roughly parallel, as a consequence of the above--mentioned relation between BH and galaxy masses.
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The discovery of remarkable correlations between the masses of supermassive BHs hosted at the centers of galaxies and the global properties of the parent galaxies themselves~\cite{8,9,12} leads to a natural link between the cosmological evolution of QSOs and the formation history of galaxies~\cite{2}. The investigation of such interesting correlations looks promising not only to better understand how and when galaxies formed, but also to obtain information on the QSO population itself~\cite{4}. Here we focus on two specific points raised by the general remarks above: 1) The use of the ``Magorrian relation'' to determine the QSO duty cycle at redshift $z=0$; 2) The expected relation between the cosmological evolution of the total luminosity emitted by star--forming galaxies and that of the total luminosity emitted by QSOs. As we will see, an interesting consequence of this last point is the possible existence of a physical process limiting gas accretion onto BHs at high redshifts. The {\it observational} inputs of our analysis are the Magorrian relation~\cite{12}, the galaxy mass--to--light ratio (from the Fundamental Plane)~\cite{5,6}, the present-day luminosity function of spheroids~\cite{14}, the present-day and the integrated QSO cosmological (light) evolution~\cite{13}, and finally the star formation history~\cite{11}. A possible alternative to the use of the mass-to-light ratio is the use of the Faber-Jackson~\cite{7} relation coupled with the so-called $\Mbh-\sigma$ relation~\cite{8,9}. The technical details will be given elsewhere~\cite{4}.
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astro-ph0112131_arXiv.txt
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0112
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astro-ph0112461_arXiv.txt
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The BaseL Stellar Library (BaSeL) is a library of synthetic spectra which has already been used in various astrophysical applications (stellar clusters studies, characterization and choice of the COROT potential targets, eclipsing binaries, ...). This library could provide useful indications to 1) choose the best photometric system for the GAIA strategy by evaluating their expected performances and 2) apply systematically the BaSeL models for any sample of GAIA targets. In this context, we describe one of the future developments of the BaSeL interactive web site to probe the GAIA photometric data: an automatic determination of atmospheric parameters from observed colours.
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The Basel Stellar Library (BaSeL) is a library of theoretical spectra corrected to provide synthetic colours consistent with empirical colour-temperature calibrations at all wavelengths from the near-UV to the far-IR (see Lejeune et al. 1997, 1998 for a complete description, and Westera et al. 1999 for the most recent version). These model spectra cover a large range of fundamental parameters (2000 $\leq$ T$_{\rm eff}$ $\leq$ 50,000 K, $-$5 $\leq$ [Fe/H] $\leq$ 1 and $-$1.02 $\leq$ log g $\leq$ 5.5) and hence allow to investigate a large panel of multi-wavelength astrophysical questions, as briefly reviewed in the next section. Since they are based on synthetic spectra, they can in principle be used in many photometric systems taken either individually or collectively, and this is another major advantage of these models. The "BaSeL interactive server" is the web version of the BaSeL models ({\tt http://www.astro.mat.uc.pt/BaSeL/}). This server is under development and the photometric systems presently available in interactive mode are: Geneva, Washington, Johnson-Cousins, Str\"omgren, HST-WFPC2, photographic RGU, and EROS. All details about this server will be given elsewhere.
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Covering a large spectral domain, extending from the UV to the far-infrared, the BaSeL models are adapted to perform simulations with the proposed GAIA photometric systems and should help to choose the most efficient one. We present and discuss a proposition to develop an automatic method, already used with success for COROT potential targets (Lastennet et al. 2001a), for a systematic determination of fundamental parameters from BaSeL synthetic multi-photometry. This new tool should be publicly available in 2002 on the following web site {\tt http://www.astro.mat.uc.pt/BaSeL/}. \small
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astro-ph0112461_arXiv.txt
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0112
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astro-ph0112182_arXiv.txt
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I review standard big bang nucleosynthesis and some versions of nonstandard BBN. The abundances of the primordial isotopes D, He-3, and Li-7 produced in standard BBN can be calculated as a function of the baryon density with an accuracy of about 10\%. For He-4 the accuracy is better than 1\%. The calculated abundances agree fairly well with observations, but the baryon density of the universe cannot be determined with high precision. Possibilities for nonstandard BBN include inhomogeneous and antimatter BBN and nonzero neutrino chemical potentials.
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Big bang nucleosynthesis (BBN) is among the main observational evidence for big bang. The discovery of the cosmic microwave background (CMB) provided us with the temperature scale of the early universe, and allowed the calculation of the primordial nuclear abundances produced in the big bang. The four light isotopes, $\D$, $\EHe$, $\UHe$, and $\ZLi$ are mainly produced in the big bang, and the calculated abundances agree fairly well with astronomical observations. Standard big bang nucleosynthesis (SBBN) has a single free parameter, the baryon-to-photon ratio, \begin{equation} \eta \equiv \frac{n_b}{n_\gamma} = 10^{-10}\ldots10^{-9}, \end{equation} which is related to the present baryonic contribution to the critical density $\Omega_b$ via the Hubble constant $H_0 \equiv h 100$~km/s/Mpc by \begin{equation} \et \equiv 10^{10}\eta = 274\Omega_b h^2. \end{equation} For decades, BBN has provided the best determination of the amount of baryonic matter in the universe. The agreement with observations is obtained in the range $\et = 1.5\ldots6$. Despite optimistic claims from time to time, BBN has not really progressed towards a much more precise determination of $\eta$. Observers claim higher precision from determinations of primordial abundances of single isotopes, but disagree with each other or, within the context of SBBN, with primordial abundances of other isotopes. Difficult questions about systematic errors in observations and chemical evolution relating the present abundances to primordial abundances have prevented further progress. During the past year, a competing method for estimating the amount of baryonic matter has appeared. In the angular power spectrum of the anisotropy of CMB, the relative heights of the even and odd acoustic peaks are sensitive to the baryon-to-photon ratio. The first preliminary estimates from the Boomerang \cite{Boom00} and Maxima-1 \cite{Maxima00} experiments appeared to be in conflict with BBN, giving a higher baryon density, $\Omega_b h^2 \sim 0.03$, or $\et \sim 8$ \cite{BooMax00}. The Boomerang collaboration has since revised their estimate downward, to $\Omega_b h^2 = 0.022^{+0.004}_{-0.003}$ \cite{Boom01}, which agrees with SBBN, but the Maxima-1 estimate has been revised upward to $\Omega_b h^2 = 0.0325\pm0.0125$ ($95\%$ c.l.) \cite{Maxima01}. With the coming satellite experiments CMB may surpass BBN as the method for estimating $\eta$. BBN will then become a tool for understanding the astrophysics of chemical evolution, by telling us the primordial abundances. While SBBN is simple and natural, and is at present in reasonable agreement with observations, there is interest in studying nonstandard variants of BBN. For one thing, BBN is a sensitive probe of the physics of the early universe. If we change something about our assumptions regarding the conditions in the early universe, or the physics relevant at that time, we are likely to change the primordial abundances and ruin the agreement with observations. Thus for many things BBN provides the strongest constraint. On the other hand, from time to time there have been suggestions for disagreement between the estimated primordial abundances of the different isotopes, and/or other ways of estimating $\eta$. If such disagreements persist, nonstandard BBN (NSBBN) may be the solution. I shall begin with a review of the physics of SBBN, and then discuss a few NSBBN scenarios.
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Standard BBN is compelling in its simplicity. While there is controversy among the observers and some apparent discrepancy between the estimated primordial abundances of the different isotopes and SBBN, these are probably not serious, and most likely represent difficulties in making observations and estimating primordial abundances based on observed ones. SBBN is thus not in trouble. Unfortunately, because of these difficulties, it is not able to pin down the baryon-to-photon ratio very precisely. It is somewhere in the range $\et = 1.5$--$6$, or $\Omega_b h^2 = 0.005$--$0.022$. The high redshift deuterium measurements point towards the upper end of this range, $\et \sim 5$--$6$, or $\Omega_b h^2 \sim 0.02$. The recent estimates from CMB anisotropy, $\Omega_b h^2 = 0.022^{+0.004}_{-0.003}$ from Boom\-erang \cite{Boom01} and $\Omega_b h^2 = 0.0325\pm0.0125$ ($95\%$ c.l.) from Maxima-1 \cite{Maxima01} also favor this upper end of the range. We are eagerly waiting for more precise CMB measurements in the coming years. While standard BBN is in good shape, there is interest in studying nonstandard BBN: to assess the robustness of SBBN, to constrain possibilities for nonstandard physics or cosmology, and to be ready to provide relief if observational discrepancies turn out to be serious for SBBN. We discussed here four NSBBN scenarios: 1) electron neutrino degeneracy, 2) electron neutrino degeneracy combined with a speed-up of the expansion rate due to extra energy density, 3) inhomogeneous BBN, and 4) antimatter BBN. All these scenarios are able to relieve the tension between the $\D$ and $\UHe$ observations. The three last ones may also allow a larger baryon density than SBBN, but with some difficulty: 2) is constrained in that respect by other cosmological constraints and 3) and 4) cannot do much for the $\ZLi$ constraint. I thank Elina Sihvola for permission to reproduce figures from \cite{Sihvola01,Sihvola01a}.
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astro-ph0112182_arXiv.txt
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0112
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astro-ph0112527_arXiv.txt
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The structure of dust spirals in the nuclei of the SAab-type Liner galaxies NGC 4450 and NGC 4736 is studied using archival HST PC images. The spirals are typically only several hundredths of a magnitude fainter than the neighboring disks, so unsharp mask techniques are used to highlight them. The ambient extinction is estimated to be less than 0.1 mag from the intensity decrements of the dust features and from the spiral surface filling factor, which is about constant for all radii and sizes. The nuclear dust spirals differ from main-disk spirals in several respects: the nuclear spirals have no associated star formation, they are very irregular with both trailing and leading components that often cross, they become darker as they approach the center, they completely fill the inner disks with a constant areal density, making the number of distinct spirals (the azimuthal wavenumber $m$) increase linearly with radius, and their number decreases with increasing arm width as a power law. Fourier transform power spectra of the spirals, taken in the azimuthal direction, show a power law behavior with a slope of $-5/3$ over the range of frequencies where the power stands above the pixel noise. This is the same slope as that found for the one-dimensional power spectra of HI emission in the Large Magellanic Cloud, and also the slope expected for a thin turbulent disk. All of these properties suggest that the dust spirals are a manifestation of acoustic turbulence in the inner gas disks of these galaxies. Such turbulence should dissipate orbital energy and transfer angular momentum outward, leading to a steady accretion of gas toward the nucleus.
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Dust spirals in the inner kpc regions of galaxies reveal a source of compression that can affect the angular momentum distribution of the gas and possibly drive accretion to an AGN. Barred galaxies with an inner Lindblad resonance tend to have two long and symmetric dust spirals near the resonance that are a continuation of the leading-edge dust lanes in the bar (Athanassoula 1992). Many barred galaxies have ILR rings too (Buta \& Crocker 1993; P\'erez-Ram\'irez et al. 2000; Knapen et al. 2000). Non-barred galaxies, galaxies without an ILR (e.g., late Hubble types), and regions of barred galaxies inside their ILRs can have more irregular dust spirals. Whether the nuclear spirals are regular or irregular, their presence in optical images suggests a density variation that is at least a factor of 2 and therefore likely to involve shocks. These shocks are oblique for azimuthal flows, so the gas will experience a torque when it enters a spiral and an opposite torque when it leaves. In cases where the gas moves faster than the spirals, the net torque is negative and the gas loses angular momentum and energy at the shock, causing it to spiral inward. If the spirals move faster than the gas, as might be the case outside a fast nuclear bar, then the gas gains angular momentum and may gradually move out. We are interested in the source of compression for the irregular nuclear dust spirals that appear in non-barred galaxies or inside the ILRs of some barred galaxies. We have proposed that some of these are caused by random sonic noise that amplifies weakly as it propagates toward the center (Elmegreen et al. 1998; Montenegro, Yuan, \& Elmegreen 1999; Englmaier \& Shlosman 2000). A signature of this process is an irregularity of structure, a wide range of pitch angles reflecting different times of origin and different wave propagation directions, a tendency to fill the volume available with a spiral separation comparable to the epicyclic radius for motions at the sound speed, and a general trend of increasing density with decreasing radius for each spiral. Because the velocities inside these spirals cannot yet be measured, there is no way to be sure they are driving inflow. Nevertheless, if any of the spirals come close to the nucleus, then such accretion would seem to be likely. Nuclear dust spirals and clouds have been studied previously using Hubble Space Telescope (HST) data. Van Dokkum \& Franx (1995) found them on WFPC V-band images of early type galaxies. Malkan, Gorjian, \& Tam (1998) did a large snapshot survey of active galaxies and classified the nuclear morphology including the dust. Elmegreen et al. (1998) noted dust spirals in WFPC2 images of the interacting galaxy NGC 2207 and proposed they might drive accretion to the nucleus. Regan \& Mulchaey (1999) found nuclear dust spirals in 6 AGN galaxies using WFPC2 and NICMOS images, and Martini \& Pogge (1999) found them with HST data in 20 Seyferts; both studies concluded that the spirals could drive accretion. Ferruit, Wilson \& Mulchaey (2000) studied 12 early type Seyfert galaxies in more detail and noted the dust spirals too. Tomita et al. (2000) used HST archival images to study dust features in E and S0 galaxies, but did not comment on spirals specifically. Tran et al. (2001) studied nuclear dust in elliptical galaxies. Here we measure the extinction and structural properties of irregular dust spirals in the central kpc of two LINER galaxies, NGC 4450 and NGC 4736. Their Hubble types are about the same, SA(s)ab and RSA(r)ab (de Vaucouleurs et al. 1991), and their distances are taken to be 16.8 and 4.3 Mpc (Tully 1988).
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The nuclear spirals in NGC 4450 and NGC 4736 vaguely resemble the outer spirals, which are somewhat flocculent in each case, but the nuclear spirals do not continue smoothly from the outer spirals and there are important structural differences. The B-band image of NGC 4450 in the Carnegie Atlas of Galaxies (Sandage \& Bedke 1994) shows two long dust spirals in the main disk, along with some flocculent structure; the stellar spiral arms are smooth. In contrast, the nuclear region has no stellar arms and at least 7 prominent dust arms, some with pitch angles as high as $45^\circ$ and some crossing each other. The eastern side of the nuclear region shows more dust than the western side because of the galaxy's inclination. Some small dust feathers extend nearly radially from the center toward the south, reminiscent of jets. HST spectral observations by Ho et al. (2000) reveal double-peaked line profiles with high velocity wings, characteristic of accretion disk activity observed in other LINERs. NGC 4736 is an early-type galaxy with an outer ring and a circumnuclear starburst ring. Its main disk structure is flocculent and defined primarily by the dust. Sandage \& Bedke's B-band print shows the inner disk structure as composed of many tightly-wrapped arms, but the central region is saturated in the reproduction. Waller et al. (2001) present UIT UV and ground-based R-band images of the central regions, including an unsharp-masked image showing the complicated flocculent structure of the main disk. They also show an HST FOC image of the main nuclear dust arms. Ground-based NIR observations by Mollenhoff, Matthias, \& Gerhard (1995) suggested a weak stellar bar with a length of $20^{\prime\prime}$, which was also noted by Maoz et al. (1995) from an HST FOC image. The bar was observed in CO by Sakamoto et al. (1999) and Wong \& Blitz (2000). In the HST image, the region corresponding to the bar shows up as an elongated disk with a position angle nearly perpendicular to the major axis of the galaxy. The structure inside the circumnuclear ring, which is really two tightly wrapped arms, consists of a dozen dust arms within a radius of 50 pc from the center, branching to dozens more dust spirals out to 200 pc. The nuclear dust spirals are not attached to the main inner disk dust spirals. The nuclear dust in NGC 4450 and NGC 4736 has several characteristics that differ from spiral arms and dust clouds in main galaxy disks. These are: \begin{itemize} \item The nuclear dust spirals shown here have no associated star formation. Other nuclear spirals in different galaxies have star formation (e.g., Coma D15 in Caldwell, Rose, \& Dendy 1999), so the gaseous nature here is not universal. The lack of star formation suggests that the inner gas disks in NGC 4450 and NGC 4736 are not strongly self-gravitating. The same was true for NGC 2207 (Elmegreen et al. 1998) and for several other inner disks in the study by Martini \& Pogge (1999). Our opacity estimate in Section 2.2 is also consistent with this. \item The nuclear dust is in the form of spiral arms of various pitch angles, widths, and lengths. Some of the arms are trailing, a few are leading, and many cross each other. This pattern is generally more irregular than main disk flocculent arms (see atlas in Elmegreen 1981). Main disk flocculent arms are rarely leading. They generally do not cross each other; if they branch into spurs, then this branching is toward larger radii (Elmegreen 1980). They also have star formation that gives them a thicker, more patchy quality, rather than a filamentary quality. \item The nuclear dust spirals in NGC 4450 and NGC 4736 have decreasing contrast with increasing radius. Ambient dust extinctions generally decrease with galactocentric distance because of the exponential distribution of gas column density. The flocculent galaxy NGC 5055 has such a decrease, for example, as measured by the extinctions of OB associations (Acarreta et al. 1996). Nuclear dust spirals are not just ambient extinctions, however. They are morphologically more similar to main disk spiral arms than diffuse cloud extinctions because they are organized and most likely formed by compressive processes in the presence of shear. From this point of view, nuclear dust spirals should be compared to main disk spirals, and then the radial decrease in nuclear spiral amplitude is unusual. Density wave spirals in non-barred galaxies tend to get stronger with increasing radius, out to at least the corotation zone (Elmegreen \& Elmegreen 1984; Elmegreen et al. 1996). The unusual result that nuclear dust spirals get weaker with radius is presumably the result of crowding near the center for waves that move inward, as predicted for solutions to Bessel's wave equation (Elmegreen et al. 1998; Montenegro et al. 1999). This is a different dynamical situation than for main disk spiral arms, for which the curvature terms ($\propto 1/kr$ for wavenumber $k$) in the wave equations can usually be ignored (e.g., Bertin et al. 1989). \item The number density of dust features is about constant with radius, indicating that the inner disk is completely filled with structure. This is unlike the situation for main galaxy disks which often have a small number of arms (e.g., 2-5) that get further apart with radius. The nuclear spirals also have some indication of a hierarchical or fractal structure because of a non-integer slope of the size distribution, examined with unsharp masks. This size distribution is approximately a power law with a slope in the range from 0.5 to 0.8. This power law is reminiscent of other properties of interstellar clouds formed by turbulence, but complicated in this case by the effects of shear, which make spirals rather than clumps, and by the Coriolis force, which resists turbulent motions on large scales. The dust features are also very weak, and the smaller clouds, as well as those further from the center, are difficult to see above the pixel noise. Fractal structure in the dust of another galaxy was also found by Keel \& White (2001) using a background illumination source. Fractal structure in galactic clouds is well known (e.g., Falgarone, Phillips, \& Walker 1991). \item Fourier transform power spectra in the azimuthal direction show the characteristic signature of turbulence compression, which is a power law slope of $-5/3$ for one-dimensional structures that are larger than the line-of-sight thickness of the galaxy disk. This makes the nuclear dust features studied here resemble the HI clouds in the LMC, with the important difference that the nuclear clouds are spiral filaments, presumably affected by shear, and the LMC clouds are globular and shell-like in a low shear environment. \end{itemize} There is no direct evidence in our observations for accretion driven by the dust spirals. Radial velocities will have to be measured to determine this. However, the increase in dust opacity with decreasing radius for the main spirals is consistent with the amplification that is expected for inward motions. In that case, the spirals could drive nuclear accretion.
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astro-ph0112527_arXiv.txt
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